ISO Science Legacy: a Compact Review of ISO Major Achievements

ISO SCIENCE LEGACY

A Compact Review of ISO Major Achievements Cover ﬁgures: Background

ISOCAM image of the Rho Ophiuchi Cloud, Abergel et al. Astronomy and Astrophysics 315, L329

Left inserts, from top to bottom: 170 μm ISOPHOT map of the Small Magellanic Cloud (40 pixel size, 1 resolution) from Wilke et al., A&A 401, 873–893 (2003). 2–200 micron composite spectrum of the Circinus galaxy obtained with the SWS and LWS spectrometers showing a plethora of atomic, ionic and molecular spectral, along with various solid-state features from dust grains of different sizes in Verma et al. this volume. Water vapour spectral lines detected in the atmospheres of all four giant planets and Titan, in Cernicharo and Crovisier, this volume. Cristalline silicates detected by ISO in different environments, in stars (young and old) and in comet Hale-Bopp in Molster and Kemper, this volume. Pure rotational hydrogen lines observed towards the molecular hydrogen emission peak of the Rho Ophiuchi ﬁlament in Habart, this volume. ISO SCIENCE LEGACY

A Compact Review of ISO Major Achievements

Edited by

CATHERINE CESARSKY European Southern Observatory, Garching, Munich, Germany

and

ALBERTO SALAMA European Space Agency, Madrid, Spain

Reprinted from Space Science Reviews, Volume 119, Nos. 1–4, 2005 A.C.I.P. Catalogue record for this book is available from the Library of Congress

ISBN: 1-4020-3843-7

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Printed in the Netherlands TABLE OF CONTENTS

Foreword vii GENERAL FRANK MOLSTER and CISKA KEMPER / Crystalline Silicates 3–28 JOSE´ CERNICHARO and JACQUES CROVISIER / Water in Space: The Water World of ISO 29–69 EMILIE HABART, MALCOLM WALMSLEY,LAURENT VERSTRAETE, STEPHANIE CAZAUX, ROBERTO MAIOLINO, PIERRE COX, FRANCOIS BOULANGER and GUILLAUME PINEAU DES FORETSˆ / Molecular Hydrogen 71–91 DAVID ELBAZ / Understanding Galaxy Formation with ISO Deep Surveys 93–119 SOLAR SYSTEM THIERRY FOUCHET, BRUNO BEZARD´ and THERESE ENCRENAZ / The Planets and Titan Observed by ISO 123–139 THOMAS G. MULLER,¨ PETER´ ABRAH´ AM´ and JACQUES CROVISIER / Comets, Asteroids and Zodiacal Light as Seen by ISO 141–155 STARS and CIRCUMSTELLAR MATTER BRUNELLA NISINI, ANLAUG AMANDA KAAS, EWINE F. VAN DISHOECK and DEREK WARD-THOMPSON / ISO Observations of Pre-Stellar Cores and Young Stellar Objects 159–179 DARIO LORENZETTI / Pre-Main Sequence Stars Seen by ISO 181–199 MARIE JOURDAIN DE MUIZON / Debris Discs Around Stars: The 2004 ISO Legacy 201–214 JORIS A. D. L. BLOMMAERT, JAN CAMI, RYSZARD SZCZERBA and MICHAEL J. BARLOW / Late Stages of Stellar Evolution 215–243 INTERSTELLAR MEDIUM ALAIN ABERGEL, LAURENT VERSTRAETE, CHRISTINE JOBLIN, RENE´ LAUREIJS and MARC-ANTOINE MIVILLE-DESCHENESˆ / The Cool Interstellar Medium 247–271 ELS PEETERS, NIEVES LETICIA MARTIN-HERN´ ANDEZ,´ NEMESIO J. RODRIGUEZ-FERN´ ANDEZ´ and XANDER TIELENS / High Excitation ISM and Gas 273–292 EMMANUEL DARTOIS / The Ice Survey Opportunity of ISO 293–310 OUR LOCAL UNIVERSE... MARC SAUVAGE, RICHARD J. TUFFS and CRISTINA C. POPESCU / Normal Nearby Galaxies 313–353 APRAJITA VERMA, VASSILIS CHARMANDARIS, ULRICH KLAAS, DIETER LUTZ and MARTIN HAAS / Obscured Activity: AGN, Quasars, Starbursts and ULIGs Observed by the Infrared Space Observatory 355–407 . . . AND BEYOND SEB OLIVER and FRANCESCA POZZI / The European Large Area ISO Survey 411–423 LEO METCALFE, DARIO FADDA and ANDREA BIVIANO / ISO’s Contribution to the Study of Clusters of Galaxies 425–446 FOREWORD

Building upon pioneering work in the 1960’s and 1970’s using ground-based, rocket- and balloon-borne systems, the realm of infrared astronomy was fully opened by the first cryogenic telescope in space – IRAS, launched in 1983. Over its ten-month lifetime, IRAS surveyed almost the whole sky in four broad infrared bands. This survey permitted the first evaluations of the total energy emitted by various systems in our galaxy and in the local universe. However, it could not address the detailed mechanisms and processes responsible for the emission detected, nor the exploration of the distant universe. IRAS results graphically illustrated to astronomers the need for sensitive infrared observatories, allowing detailed spatial and spectroscopic study of specific targets. All over the world, high priority was assigned to cooled space infrared telescopes. Following the Japanese IRTS mission, the first major satellite of this type to fly was ESA’s Infrared Space Observatory (ISO). Launched in November 1995, ISO completed almost 30 000 scientific observations in its 2.5-year operational lifetime. Making use of its four sophisticated and versatile scientific instruments (a camera, a photopolarimeter and two spectrometers), ISO provided astronomers with a wealth of data of unprecedented sensitivity at infrared wavelengths from 2.5 to 240 μm. ISO has made, and continues to make, lasting contributions to all areas of astronomy, from the solar system to the frontiers of cosmology, unravelling the history of the universe. Between 1996 and 2004, over 1200 papers appeared in the refereed literature based on ISO data. NASA’s Spitzer Space Telescope, launched eight years later, has enhanced capabilities compared to ISO and, once again, infrared astronomers are offered matchless observing opportunities. However, the published ISO results and the ISO archive remain a valuable resource for research work. They provide guidelines for studies not only with Spitzer but also with future facilities, such as ASTRO- F, SOFIA, Herschel, JWST and ALMA. With the Spitzer Space Telescope now in full operations, we thought that it would be beneficial to the astronomical community to have at hand, in a single volume, a review of the main discoveries owed to the ISO satellite. We did not ask the ISO founding fathers and mothers to write the articles, but instead turned mostly towards younger astronomers whose careers have been strongly influenced by ISO. The articles have been refereed by ourselves or by other scientists at our request. The book is organised as follows: first, overviews of four major themes investigated with ISO (crystalline silicates, molecular hydrogen, deep surveys; water in the universe), and then thirteen chapters reviewing ISO science from the solar system to the distant universe. It is not possible to gather in one book all the advances due to ISO, but we hope that this compendium of over 480 pages will give the essence of the original results obtained by the first full-fledged space infrared observatory. Catherine Cesarsky and Alberto Salama Guest editors GENERAL CRYSTALLINE SILICATES

FRANK MOLSTER1,∗ and CISKA KEMPER2 1ESTEC/ESA, Keplerlaan 1, 2201 AZ Noordwijk, The Netherlands 2(Spitzer Fellow) Department of Physics and Astronomy, UCLA, 475 Portola Plaza, Los Angeles, CA 90095-4705; Present address: University of Virginia, Department of Astronomy, P.O. Box 3818, Charlottesville, VA 22903-0818, USA (∗Author for correspondence: E-mail: [email protected])

(Received 16 July 2004; Accepted in ﬁnal form 2 November 2004)

Abstract. One of the big surprises of the Infrared Space Observatory (ISO) has been discovery of crystalline silicates outside our own Solar system. It was generally assumed before that all cosmic silicates in space were of amorphous structure. Thanks to ISO we know now that crystalline silicates are ubiquitous in the Galaxy (except for the diffuse ISM) and sometimes even in very large quantities (>50% of the small dust particles). The evolution of the crystalline silicates is still not completely clariﬁed, but the combination of theoretical modeling and observations have already shed light on their life-cycle. The absence of crystalline silicates in the diffuse ISM provides us with information about the dust amorphization rate in the ISM.

Keywords: crystalline silicates, infrared astronomy

1. Introduction

Before the Infrared Space Observatory (ISO) opened the mid- and far-infrared range for high-resolution spectroscopy, it was generally assumed that cosmic dust silicates were of amorphous structure. The crystalline silicates, the highly ordered counterparts of the amorphous silicates, were only known to be present on earth, in the solar system in comets (Hanner et al., 1994; Hanner, 1996), Interplanetary Dust Particles (IDPs) (MacKinnon and Rietmeijer, 1987; Bradley et al., 1992) and in the dust disk of β-Pictoris (Knacke et al., 1993; Fajardo-Acosta and Knacke, 1995), also a crystalline olivine feature was reported in the polarized 10 μm spectrum of AFGL2591 (Aitken et al., 1988). Apart from the crystalline silicates in the IDP’s, that were found with the aid of transmission electron microscopy and only later conﬁrmed by infrared spectroscopy (Bradley et al., 1992), in all other cases the crystalline silicate features were found by infrared spectro(polari)metry around 10 μm. With the present day knowledge it is relatively easy to understand why crystalline silicates were only discovered to be ubiquitous after ISO was operational. Before ISO was launched the primary MIR/FIR window for observations was around 10 μm. And although the crystalline silicates do have strong features in this area, they are in general overwhelmed by emission from the much more abundant and Space Science Reviews (2005) 119: 3–28 DOI: 10.1007/s11214-005-8066-x C Springer 2005 4 F. MOLSTER AND C. KEMPER warmer amorphous silicates. Furthermore, most of the crystalline silicates have a (relatively) low temperature (<150 K), which suppresses the intrinsically strong crystalline silicate features in the 10 μm region. Thanks to the extended wavelength range (up to 200 μm) of the spectrographs on board ISO, the composition of the cold (<150 K) dust, which has the top of its SED at wavelengths above 15 μm, could be studied in detail for the ﬁrst time. With ISO, crystalline silicates have been found around young stars (Waelkens et al., 1996), comets (Crovisier et al., 1997), and evolved stars (Waters et al., 1996) (see Figure 2), but not convincingly so in the interstellar medium. The presence of the crystalline silicates in the different galactic environments will be discussed in Section 2. The properties of the crystalline silicates will be discussed in Section 3. In Section 4 we will discuss the crystalline silicate (trans-)formation and destruction processes based on the ISO (and other astronomical) observations in comparison with laboratory measurements.

2. The Ubiquitous Presence of Crystalline Silicates

2.1. WHAT ARE SILICATES?

Silicates are the most common form of minerals in the solar system, and probably also beyond. Chemically, they consist of silica tetrahedras (SiO4) which are combined with metal cations, such as Mg2+ or Fe2+ in a lattice structure. In ordered (crystalline) lattice structures, the tetrahedras can share their oxygen atoms with other tetrahedras to form different types of silicates: olivines: (Mg, Fe,...)2SiO4, pyroxenes: (Mg, Fe,...)SiO3 or quartz: SiO2. In case of unordered (amorphous) structures, the number of shared oxygen-atoms may vary for each silica-anion (see Figure 1). The reader is referred to for instance Klein and Hurlbut (1993) for more background reading on mineralogy. Spectroscopically, all silicates, amorphous and crystalline alike, will show resonances around 10 and 20 μm in the mid-infrared, due to the Si O stretching and the O Si O bending mode arising from the silica-tetrahedras. Alignment of the tetrahedras may cause sharp peaked resonances, whereas amorphous silicates will show a broad feature which can be seen as a blend of such sharp resonances. Crystalline silicates can also be distinguished from amorphous silicates due to the presence of lattice modes at λ 25 μm, see Figure 1. The formation temperature determines whether a silicate grain becomes crystalline or amorphous. If silicates are formed above the glass temperature (Tglass), then there is enough mobility in the silicate to from the crystalline, energetically most favorable, lattice structure. However, when the silicates condense at lower temperatures, such mobility is not present, and the grain solidiﬁes in amorphous form. Amorphous grains can become crystalline by annealing (crystallization through heating) or vaporization and recondensation above the glass temperature. Hence, CRYSTALLINE SILICATES 5

Figure 1. A possible atomic structure of a disordered (or amorphous) silicate and that of an ordered (or crystalline) silicate together with their typical infrared emission spectra. The tetrahedras are four oxygen atoms around a silicon atom and the big circles are metal atoms. Note the many sharp features in the crystalline silicate spectrum and the two broad bumps at 10 and 20 μm for the amorphous silicate spectrum. the presence of crystalline silicates traces the occurrence of high-energy processes. On the other hand, damage caused by cosmic ray hits or grain–grain collisions can amorphitize silicates.

2.2. OBSERVATIONAL INVENTORY

Crystalline silicates have been found around evolved stars with an oxygen-rich dusty outﬂow: AGB-stars, post-AGB stars and Planetary Nebulae; (Waters et al., 1996; Sylvester et al., 1999; Molster et al., 2001a; Molster et al., 2002b). Surprisingly however, there are also several examples of crystalline silicates being present around stars with a carbon-rich chemistry (e.g., Waters et al., 1998; Molster et al., 2001b). It is expected that in such cases, the crystalline silicates originate from (a) previous mass-loss episode(s), when the mantle of the star was still oxygen-rich. The fraction of silicates that is in crystalline form is modest; using full radiative transfer calculations it is found that the typical abundance of crystalline silicates in the winds of low mass post-MS stars is of the order of 10% or less of the total 6 F. MOLSTER AND C. KEMPER

Figure 2. Crystalline silicates are detected in many different environments. These environments include young stars such as HD 100546 (upper left), evolved stars such as IRAS 09425-6040 (upper right) and also comet Hale–Bopp (lower left). A featureless dust continuum has been subtracted from the spectra to enhance the appearance of the crystalline silicates. For reference, laboratory spectra of the crystalline silicates enstatite and forsterite, multiplied with 85 K blackbodies, are shown in the lower right panel. silicate mass (Kemper et al., 2001). It should be noted that this number does depend on the laboratory spectra used. Using different sets of laboratory data, can make a difference in the abundance sometimes up to a factor 2. A simple model fit to the spectra reveals that enstatite is about three times as abundant as forsterite in these outflows (Molster et al., 2002c). The above mentioned values for the abundances −5 −1 are only derived for stars with a relatively high mass loss rate (M˙ 10 M yr ). To date, there is no evidence for the presence of crystalline silicates around low −6 −1 mass-loss-rate AGB stars (M˙ 10 M yr ; Waters et al., 1996; Cami et al., 1997; Sylvester et al., 1999). It is well possible that this is a temperature effect. Kemper et al. (2001) show that the temperature difference between the crystalline and amorphous silicates will cause the emission from the crystalline silicates to be overwhelmed by the warm amorphous silicates in the optically thin dust shell around low mass-loss-rate stars. This temperature difference is caused by the difference in optical properties (see also Section 3). Although most post-main-sequence stars have modest degrees of crystallinity, in some peculiar objects, likely as a result of binary interaction, the crystalline CRYSTALLINE SILICATES 7 silicate abundance can be very high, up to 75% of the small grains (Molster et al., 2001b). Despite the deposition of crystalline silicates into the interstellar medium, we do not have strong evidence for the presence of crystalline silicates in this environment. From the absorption profile in the direction of the galactic center an upper limit of about 0.2 ± 0.2% for the degree in crystallinity of silicates in the diffuse interstellar medium has been derived (Kemper et al., 2004). Crystalline silicates are suggested to be present in the dense ISM in Carina, by the detection of the 65 μm band ascribed to diopside (CaMgSiO3; Onaka and Okada, 2003), although the intrinsically strong mid-infrared features are not detected. Cesarsky et al. (2000) suggest that several emission bands in the ISO spectra of the Orion bar may be due to crystalline silicates. We note, however, that the peak position, shape and width of the bands seen in Orion differ from those of crystalline silicates in other environments. Re-analysis of the ISO spectrum of this line-of-sight shows that there is no evidence for crystalline silicates in the Orion Bar (Kemper et al., in preparation). The ISO data show that the abundance of crystalline silicates in Young Stellar Objects (YSOs) is higher than the upper limits that are established for the interstellar medium. The first Herbig Ae/Be star observed by ISO-SWS, HD 100546, shows a remarkably rich crystalline silicate spectrum (Waelkens et al., 1996; Malfait et al., 1998). A thorough analysis of the infrared spectrum of this star indicates that the crystalline silicates dominate in the 100–200 K temperature range (Bouwman et al., 2003). In this temperature range the crystalline silicates are the dominant small dust grains, even more abundant than the small amorphous silicates. As it later turned out, this star does not have a typical dust composition. Most other young stars show much more modest crystalline silicate abundances (Meeus et al., 2001), but still higher than the ISM. The relatively high abundance of crystalline silicates in the circumstellar dust shells and disks around young stars compared to the interstellar medium points to an in situ formation mechanism. For most comets only ground-based spectra of the 10 μm region are available. This limits the determination of dust abundances. These abundances are very sensitive to the temperature of the different dust components, which is very difficult to determine from a small wavelength range (8–13 μm). Fortunately, while ISO was operational the comet Hale–Bopp passed by. The long wavelength coverage of ISO made it possible to constrain the temperature distribution of the individual dust components much better than before. Several authors have tried to determine the crystallinity of the silicates in Hale–Bopp (Wooden et al., 1999; Brucato et al., 1999; Bouwman et al., 2003). Depending on additional grain properties, such as grain size, shape and whether the grains are thermal contact with each other, the derived degree of crystallinity of the silicates ranges between more than 90% (Wooden et al., 1999) and about 7% (Bouwman et al., 2003) of the total dust mass. Recent work by Min et al. (private communication) suggests that the crystallinity of the silicates in Hale–Bopp may be even lower. This illustrates the still remaining difficulty in determing the dust composition. The main differences come from the 8 F. MOLSTER AND C. KEMPER use of different laboratory spectra with different Fe contents for the silicates and different grain size and shape distributions. A high resolution high S/N spectrum of the dust features might help to determine the composition and even the average shape of the dust grains in future comets (see also Section 3), but the abundance of the crystalline silicates in this comet will remain uncertain. An interesting correlation has been found between the abundance of the crystalline silicates and the geometry of the dust distribution around a star. Figure 3 shows the correlation between the strength of the 33.6 μm feature (attributed to

Figure 3. Correlation between the IRAS 60 μm over 1.1 mm (stars) or 1.3 mm (pentagons) flux ratio and crystalline silicate band strength, measured from the crystalline forsterite band at 33.6 μm ([F33.6 μm peak − Fcont]/Fcont). The selected sources have dust color temperatures above ≈100 K, so that the Planck function peaks at wavelengths shorter than 60 μm. The emptiness of the upper right corner of this diagram is an indication that the presence of highly crystalline dust is correlated with disks and grain-growth. All stars above the line (F33.6 μm peak − Fcont)/Fcont = 0.2 have relatively small 60 μm over mm-flux ratios and have indications for the presence of a disk (e.g., by direct imaging and/or the spectral energy distribution). The stars below this line are predominantly normal outflow sources, without any evidence for a disk, although exceptions exist (e.g., Roberts 22). NGC 6302 and BD+30 3639 have been shifted by a factor 2 along the x-axis since it is estimated that only half of the mm continuum flux is due to dust emission while the other half is due to free–free emission. Figure taken from (Molster et al., 1999b). CRYSTALLINE SILICATES 9 forsterite) and the IRAS 60 μmmm−1 flux ratio in post-MS stars. The first value gives an indication of the abundance ratio between amorphous and crystalline silicates in the dust, while the second ratio gives an indication of the average grain size. It is interesting to note that those sources which seem to have a large abundance of crystalline silicates, also show evidence for the presence of a disk-like structure and for grain coagulation (Molster et al., 1999b). Note that the opposite is not true, if stars do have a disk (and large grains) it does not automatically imply that they have a high fraction of crystalline silicates. A similar correlation between degree of crystallinity and grain size is seen in the disks around pre-MS stars (see Figure 4) (Bouwman et al., 2001). The origin of this correlation remains unclear. Whether the crystallization of amorphous silicates is related to grain coagulation, or that we

Figure 4. Dust processing in planet forming systems. The panels show ISO-SWS continuum subtracted spectra of the 10 μm regions, ordered by peak position, the first spectrum (of the ISM) shows the bluest feature. A model fit to each spectrum is included in the plot. The model includes grain growth, which is reflected in the width of the feature, and the shift towards longer wavelengths, as well as crystallization with the appearance of the 11.3 μm feature. From this plot it becomes clear that there is a correlation between grain growth and crystallinity in the various phases of planet formation. Relatively young and unprocessed silicates are seen in the top first few spectra (including the spectra of the diffuse ISM and a red giant, μ Cep) whereas the more evolved planetary systems are found further down in the sequence. Figure adopted from (Bouwman et al., 2001). 10 F. MOLSTER AND C. KEMPER are simply dealing with two different processes, which both require long timescales has not yet been established (see also Section 4).

3. The Properties of Crystalline Silicates

The sharp infrared features (see e.g., Jäger et al., 1998 for an assignment of the bands of forsterite and enstatite) of the crystalline silicates allow a quite accurate identification of the crystalline materials. The two most abundant crystalline silicates that have been found are forsterite (Mg2SiO4) and enstatite (MgSiO3, which can have a monoclinic and orthorhombic crystallographic structure, called respectively clino and ortho-enstatite). As an example, a fit to the continuum subtracted spectrum of MWC922 with only these two species is shown in Figure 5. The high resolution spectroscopy of ISO-SWS and LWS made it not only possible to determine the type of crystalline silicates, it also made it possible for some stars to determine their exact mineralogical composition of some circumstellar dust components. The presence of Fe2+ in crystalline silicates not only changes the opacity, it also reduces the strength and increases the wavelength of the crystalline silicate features significantly (see e.g., Jäger et al., 1998). As long as the amount

Figure 5. (Left): The continuum subtracted spectrum of MWC922 (solid line; Molster et al., 2002b), compared with the calculated emission spectrum of forsterite (at 90 K; dashed line), and clino and ortho-enstatite (at 100 K; dotted line). The temperatures have been derived by fitting the continuum subtracted spectrum with only forsterite and enstatite, 50% ortho and 50% clino enstatite (Molster et al., 2002c). The fit is shown as the dashed-dotted line. The vertical lines denote the diagnostic features of forsterite (dashed lines) and enstatite (dotted lines). (Right): The continuum subtracted spectrum of NGC6302 (solid line; Molster et al., 2002b), compared with the calculated (using optical constants derived from laboratory experiments) emission spectrum of diopside (at 70 K; dashed line), and crystalline water ice (at 40 K; dotted line). The temperature of diopside is chosen the same as the temperature found for the other pyroxenes (Molster et al., 2002c), and water ice has been chosen to fit both the 40 and 60 μm complex. The combined result is shown as the dashed-dotted line. The vertical lines indicate the wavelengths of diagnostic features of diopside (dashed lines) and crystalline water-ice (dotted lines). Note that no effort has been made here to fit the forsterite features at 33 μm and the enstatite feature at 40 μm. CRYSTALLINE SILICATES 11 of Fe2+ in the crystal is not too high, a first order estimate is that the crystalline silicates only contribute to the features and not to the continuum. The subtraction of a continuum (caused by other dust components) from the ISO spectra, gives you therefore a reasonable first impression of the crystalline silicates. In Figure 5 we show two continuum subtracted spectra of observed stars (MWC922, and NGC6302). The diagnostic features are indicated with dashed and dotted lines. Note that the difference between clino- and ortho-enstatite only becomes apparent after 40 μm. A comparison of the strength of the individual features of ortho and clino-enstatite shows that around most stars the abundance is about equal. For some stars, with very high mass loss rates, ortho-enstatite may be more abundant than clino-enstatite (Molster et al., 2002c). Besides the above mentioned Mg-silicates, there is also evidence for Ca- pyroxenes such as diopside (Koike et al., 2000). In Figure 5, we show the calculated emission spectra of diopside and crystalline water ice. Both have a rather broad feature near the peak of the 60 μm band. However, the observed 60 μm band is broader than the individual H2O ice and diopside bands, while a sum of both materials gives a satisfactory fit to the 60 μm region (note that the addition of the carbonate dolomite improves the fit even further Kemper et al., 2002b). Unfortu- nately, the low abundance of diopside in combination with blending of the short wavelength diopside bands with those of forsterite and enstatite make it hard to unambiguously identify diopside based on only the shorter wavelength bands. This implies that we can only clearly identify this material in systems which have very cool dust (T < 100 K), such as OH/IR stars and planetary nebulae. Another aspect of the chemical composition of the crystalline silicates, in particular the olivines and pyroxenes, is the Fe/Mg ratio in the lattice. The silicate composition can be rather accurately determined because the wavelength and strength of the bands are very sensitive to differences in the Fe/Mg ratio. Laboratory studies show that a simple relation exists between the position of the crystalline silicate bands and the Fe/Mg ratio in the lattice. The peak position shift in the frequency space due to the inclusion of Fe goes roughly linear with the percentage of [FeO] in the silicate (Jäger et al., 1998). x =− . 100ν 1 8 for olivines (Mg(2−2x)Fe2x SiO4) and (1) x =− . 100ν 1 5 for pyroxenes (Mg(1−x)Fex SiO3), (2) with 0 ≤ x ≤ 1 and ν = νx − ν0 where νx and ν0 are respectively the wavenum- ber of the feature for composition x and x = 0. This implies that the shift in the wavelength domain is proportional to λ2 and therefore clearest at the longest wavelengths. This is the reason why it is very difficult to determine the exact Mg/Fe ratio from observations limited to only the 10 μm region. Forsterite has a weak band at 69 μm. This band is indeed very sensitive to changes in the iron abundance in the crystal (see e.g., Koike et al., 1993; Jäger et al., 1998). Figure 6 shows the 12 F. MOLSTER AND C. KEMPER

Figure 6. The observed FWHM and peak wavelength of the 69.0 μm feature in the spectra of the dust around stars (open diamonds for the sources with a disk and open circles for the sources without a disk) and in the laboratory at different temperatures (filled triangles – forsterite; Fo100 (Bowey et al., 2001), and filled squares – olivine; Fo90 (Mennella et al., 1998)). The temperatures are indicated at each point, and within the resolution the 24 K and 100 K for Fo90 are similar. Note that the measurements were not corrected for the instrumental FWHM (≈0.29, for the ISO observations, and 0.25 and 1.0 μm for the laboratory observations of respectively Fo100 and Fo90). Figure taken from (Molster et al., 2002c). evidence for the high Mg/Fe ratio of the crystalline olivines in the outflows of evolved stars. In fact the data is even consistent with the absence of iron in the matrix (i.e., forsterite). From this plot it also becomes clear that the forsterite crystals are very cold. A similar result holds for the pyroxenes based on the 40.5 μm feature. The crystalline olivines and pyroxenes found in the dusty winds of evolved stars invariably show evidence for very Mg-rich crystals (forsterite and enstatite). The best results in determining the composition of the silicates are achieved for evolved – post-MS – stars, where due to stellar evolution constraints often an oxygen-rich chemistry prevails in the outflow. The determination of the exact composition of the crystalline silicates around young stars is more complicated. Only for the peculiar star HD100546 a 69 μm feature due to forsterite has been found and even that one is rather noisy (Malfait et al., 1998). In all other stars such band has not been found in that wavelength region. This is likely due to a combination of sensitivity, too high dust temperature (a higher temperature creates a broader and weaker feature) and/or too low abundance. In young stars, the characteristic features of crystalline silicates are found at wavelengths below 45 μm, especially in the 10 μm complex region. It cannot be excluded that some Fe is present in the matrix and this will also reduce the strength of the feature. Although the position of the features at wavelengths below 45 μm seem to indicate that the crystalline silicates in young stars are also relatively Fe-poor. CRYSTALLINE SILICATES 13

The large number of bands of the crystalline silicates makes it possible to determine the temperature of the different species of crystalline silicates independent from each other. In general the forsterite and enstatite grains have similar temperatures, but are usually cooler than the amorphous grains (Molster et al., 2002c). This implies that the crystalline and amorphous silicates are not in thermal contact, and either they are spatially distinct, or they have significantly different optical properties. Although it is difficult to exclude the first option, the temperature difference can be explained in a straightforward way if we take into account the differences in the Fe/Mg ratios for the amorphous and crystalline silicates. As shown above, the crystalline silicates are very Fe-poor. In contrast, it has been argued that amorphous silicates contain Fe, either in the matrix or as a metal inclusion (Jones and Merrill, 1976; Ossenkopf et al., 1992; Kemper et al., 2002a). Adding a modest amount of iron already increases the opacity in the near-IR significantly (Dorschner et al., 1995). Radiative transfer modeling shows that, if one assumes that the amorphous silicates have Fe/Mg ≈1 and that the crystalline silicates are Fe-free, the temperature differences found are fully explained by the difference in near-IR opacity. Both the amorphous and crystalline silicates can then be co-spatial (see e.g., Molster et al., 1999a; Kemper et al., 2001; Hoogzaad et al., 2002).

4. The Evolution of Crystalline Silicates

Although the crystalline lattice structure is energetically the most favorable state for silicates, the observations show that most of the silicate grains are amorphous or glassy. This requires an explanation. First of all, some threshold energy has to be overcome before an amorphous silicate will crystallize. At temperatures above 1000 K the crystallization takes place on a timescale of seconds to hours, while below 900 K it will take years to more than a Hubble timescale (Hallenbeck et al., 1998). If silicate dust condenses much quicker than the crystallization timescale, it will be amorphous. This implies that if the silicates form or quickly cool off below 900 K, they will remain amorphous. Since the majority of the silicates that forms is amorphous, this may imply a formation below the glass temperature. An alternative theory assumes that the silicates condense as crystals but their crystal structure will get destroyed in time. Tielens et al. (1997) propose a scenario to explain the difference between crystalline and amorphous silicates by this destruction mechanism. They propose that gas phase Fe diffuses into the Mg-rich crystals around 800 K (which is for all practical purposes below the crystallization temperature) and destroys the crystal structure during its intrusion. In this scenario all grains would form crystalline and only later the bulk would turn amorphous as Fe is adsorbed. Very high spatial resolution far-IR observations of the dust around AGB stars should indicate if there is indeed a rather sudden decrease of the crystalline silicate abundance in the dust shell at a temperature of roughly 800 K. 14 F. MOLSTER AND C. KEMPER

Finally, the velocity difference between dust and gas particles can be high enough for destruction and sputtering. For example, for silicate grains the velocity difference with a helium atom should be about 30 km s−1 for sputtering, and these velocities are not unlikely (Simis, 2001). Colliding crystalline silicate grains might be heated to high enough temperatures to melt or even evaporate. In the ﬁrst case, the very quick cooling timescale (depending on the grain size), might result in amorphous silicates. Although it remains difﬁcult to explain the different opacities for amorphous and crystalline silicate grains in this scenario.

4.1. DUST FORMATION IN EVOLVED STARS

There are two main classes of evolved stars that produce significant amounts of dust: (i) the evolved low-mass stars, with their slow expanding winds during the AGB phase, and (ii) the high-mass stars that go through a red supergiant phase during which they have a stellar wind with properties somewhat similar to those of the low-mass AGB stars. There are also some less common stars that may produce dust, such as Luminous Blue Variables, Novae, Supernovae and some pre-main sequence stars. Here we will discuss the dust formation in AGB stars; a similar description holds for the red supergiants. It is generally believed that the stellar wind from an AGB star increases with −7 −4 −1 time, from M˙ ≈ 10 to 10 M yr (see e.g., Habing, 1996). In M giants with −7 −1 low mass loss rates (M˙ < 10 M yr ) the dominant dust species that form are simple metal oxides (Cami, 2002; Posch et al., 2002); amorphous silicates are not prominent in the spectra. The low density in the wind likely prevents the formation of relatively complex solids like the silicates, because that requires multiple gas– solid interactions. It is not clear at present whether these simple oxides also form in the innermost regions of AGB winds with higher mass loss rates. The most abundant dust species in the winds of oxygen-rich stars with intermediate and high mass loss rates are the silicates. The formation of silicates is not well understood. Homogeneous silicate dust formation directly from the gas phase is not possible, it requires first the formation of condensation seeds. It seems reasonable to search for high temperature condensates that can form through homogeneous nucleation that could serve as condensation seeds for other materials. Gail and Sedlmayr (1998) propose that in O-rich winds TiO2 (rutile) clusters can act as condensation nuclei for silicates. Quantum mechanical calculations seem to confirm this possibility (Jeong et al., 2000). However, no evidence has been found yet for the presence of rutile in the spectra of outflows of evolved stars; probably the low gas-phase abundance of titanium prevents detection. One of the main discoveries of ISO in the field of solids in space has been the detection of crystalline silicates in the outflows of AGB stars with high mass loss rates. It should be kept in mind that the abundance of crystalline silicates compared to the amorphous silicates is modest, typically 10–15%, and only in very CRYSTALLINE SILICATES 15 special cases the crystalline silicates dominate. Crystalline silicates features are only detected in AGB stars with a fairly high mass loss rate, usually a threshold −5 −1 value of 10 M yr is observed (Cami et al., 1997; Sylvester et al., 1999). As discussed before it remains unclear whether the lack of crystalline silicate features in the spectra of low mass-loss rate AGB stars indicate that the dust produced by these stars is completely amorphous, or that still a large fraction of crystalline dust can be present, found at lower temperatures than the amorphous silicates (Kemper et al., 2001). Sogawa and Kozasa (1999) derive that for low mass loss rates (M˙ < 3 × 10−5 −1 M yr ) the growing silicate grains do not heat to high enough temperatures to become crystalline, whereas for higher mass-loss rates, the thermal processing of the grains makes them entirely crystalline. Even though these authors are able to explain a threshold value for the mass-loss rate below which crystalline silicates do not exist, they are unable to provide sufficient explanation for the co-existence of crystalline and amorphous grains in many sources. Dust nucleation theories are now challenged to explain the co-existence and difference in properties of crystalline and amorphous silicates in AGB stars: (i) crystalline silicates are Fe-poor, amorphous ones contain Fe in some form, (ii) crystalline and amorphous silicates have different temperatures and thus are separate populations of grains. “Partially crystalline” grains do not seem to exist, which can be explained from a thermodynamic point of view: thermal crystallisation timescales are a steep function of temperature, and it requires extreme fine-tuning in the thermal history of a particle to observe it partially crystalline. As mentioned above, the crystalline olivines and pyroxenes formed in the winds of AGB stars with high mass loss are Fe-free (Molster et al., 2002c). Interestingly, it is also the Fe-free olivines and pyroxenes that are the first silicates that are expected to condense. This will take place at temperatures high enough so they will crystallize very quickly. Thermodynamic calculations show that the Fe-containing silicates will condense at a somewhat lower temperature (Gail and Sedlmayr, 1999), and since there is no evidence for Fe-containing crystalline silicates, they apparently remain amorphous. A possible explanation for this phenomenon is that as soon as some dust forms, radiation pressure will accelerate the dust (and through the drag- force also the gas) and the material will quickly be pushed to cooler regions. So, only the very first condensed silicates, the Fe-free silicates, might have a chance to crystallize, while the later condensed silicates (the Fe-containing silicates) are formed in an environment where the annealing timescale is much longer than the accretion timescale. Weshould note here that it is not clear that a ferromagnesiosilica will condense at all from the gas-phase. Gas phase condensation experiments of a Fe-Mg-SiO-H2-O2 vapor show only condensates of a magnesiosilica and of a ferrosilica composition but not of a ferromagnesiosilica composition (see Figure 7). While in most cases crystalline silicate dust is only a minor component in circumstellar dust shells, there are some evolved stars with a very high abundance of crystalline silicates. Remarkably, these always have a disk-like dust density 16 F. MOLSTER AND C. KEMPER

Figure 7. Ternary diagram MgO-FeO-SiO2 (oxide wt%) with the chemical compositions of gas to solid condensed grains from a Fe-Mg-SiO-H2-O2 vapor. The “average bulk solid” composition (the big dot) is roughly the gas phase composition which might have been somewhat less SiO2-rich. Figure taken from Rietmeijer et al. (1999). distribution (Molster et al., 1999b). Because the temperature in these disks is well below the annealing temperatures of amorphous silicates, Molster et al. (1999b) suggested that a low-temperature crystallization process is responsible for the higher fraction of crystalline silicates. Several processes were discussed by Molster et al. (1999b), but no conclusive answer could be given. Recently, partial crystallization at room temperature due to electron irradiation has been detected (Figure 8 and Carrez et al. (2001)). Although these results look promising, there are still open questions. It is not sure that this process can create highly crystalline material in suf- ﬁcient quantities. Investigations are still necessary to ﬁnd out whether the particles can become completely crystallized or will simply become partially amorphous and partially crystalline. Note, that electron irradiation only causes local crystallization unlike thermal crystallization, which acts on the whole grain at the same time. Observations indicate that the amorphous and crystalline silicates are two different grain populations which are not in thermal contact (Molster et al., 2002c) and it is unclear that this can be achieved by this process. Finally, the compositional implications (with or without Fe) of this process are also not well studied yet.

4.2. DUST FORMATION IN SUPERNOVAE

Unfortunately, little is known about the dust production rate in supernovae (SNe). Generally, SNe are assumed to produce only modest amounts of dust compared to the dust production by AGB stars. Yet,recent submillimeter continuum observations obtained with SCUBA on the JCMT suggest that a much larger amount of dust may CRYSTALLINE SILICATES 17

Figure 8. The infrared absorption spectrum of an amorphous silicate smoke before (dashed line) and after (solid line) a 10 min irradiation by the electron beam of a transmission electron microscope. Note the upcoming of the spectral features. be produced or processed by SNe, up to a fraction of 75% of the dust injected into the interstellar medium (Dunne et al., 2003). On the other hand, (Dwek, 2004) argues that if the dust causing the submm continuum were present in the form of metallic iron needles, the dust mass reduces to a fraction of the dust mass derived by (Dunne et al., 2003). Hence, determining the composition of dust of supernova origin proves to be of crucial importance. Only few infrared spectra of SN dust ejecta are available and they are not always conclusive. A 22 μm feature is observed in the ISO SWS spectrum of an area centered on a fast moving knot in supernova remnant Cas A. (Arendt et al., 1999) identify this dust component as proto-silicates, which appear to have similar properties as amorphous silicates. However, (Douvion et al., 2001) reject this identification based on a SWS spectrum taken very close to position N3, which shows much similarities (but is not exactly equal). They fit their spectrum with a mixture of three dust species (MgSiO3, SiO2 and Al2O3) with each having two different temperatures together with some synchrotron emission. The presence of amorphous MgSiO3 was predicted by Kozasa et al. (1991) and fitted the ISOCAM spectra reasonably well (Douvion et al., 1999). However, we would like to note that 18 F. MOLSTER AND C. KEMPER the rather artificial temperatures (all dust species have different temperatures both for the hot and cold phases which are not directly related to opacity differences), the arbitrary dust sizes and the questionable feature identifications, leave some room for alternative interpretations. In addition, the large differences in physical and chemical conditions between AGB winds and SN ejecta make it difficult to assume similarities in the types of dust that may condense in SN ejecta.

4.3. THE PROCESSING OF DUST IN THE ISM

The average lifetime of a silicate dust grain in the ISM is roughly 4 × 108 years, while the replenishment rate of ISM dust is about 2.5 × 109 years (Jones et al., 1996). This means that interstellar dust grains are destroyed and re-formed about six times before finally entering a star forming region to be incorporated into a new generation of stars. Thus, new grains must be formed in the ISM. In addition, silicates may be processed in the diffuse ISM. Below, we discuss dust formation, processing and destruction processes that may occur in the ISM. The composition of the silicates in the interstellar medium has been subject to several different studies. Based on infrared spectroscopy with SWS and other instruments it is found that the silicates in the ISM are largely amorphous. Olivines are the dominant form, and only some very small fraction may be crystalline (Kemper et al., 2004). Evidence for crystallinity seems to be absent. Some tentative detections of crystalline silicates are reported in the Orion Bar (Cesarsky et al., 2000), and in the Carina star forming region (Onaka and Okada, 2003), although re-analysis of the SWS data of the Orion Bar indicate that there are no detectable features due to crystalline silicates in these spectra. The absorption line-of-sight towards the Galactic Center observed with ISO-SWS, shows that the silicates in the diffuse ISM are completely amorphous (see Figure 9). An upper limit of 0.2% (±0.2%) can be placed on the degree of crystallinity in the diffuse ISM (Kemper et al., 2004). Using ground-based 10 μm spectroscopy, (Bowey and Adamson, 2002) argue that a broad feature due to silicates, which appear to be amorphous upon first examination, may in fact be caused by a combination of a large number of different crystalline silicates. Based on only the 10 μm spectroscopy the distinction between purely amorphous silicates and a cocktail of crystalline silicates cannot be easily made, however, all these crystalline silicates will cause sharp resonances in the mid-infrared. The presence of these resonances could be easily checked with ISO-SWS data, and there are not cases known where such resonances exist, while the 10 μm feature appears smooth. Hence, we conclude that the silicates in the ISM are predominantly amorphous. It is evident that, because the silicates around post- and pre-main-sequence stars have a significant degree of crystallinity, while the dust in the ISM is predominantly amorphous, processing of the silicate dust in the ISM plays an important role. The dust formation process in the ISM is not clear, however it seems inevitable that it takes place in the dark molecular clouds, where high densities prevail. The CRYSTALLINE SILICATES 19

Figure 9. Optical depth in the 10 μm silicate feature observed towards the Galactic Center, with ISO-SWS. The four panels show fits to the feature for 0, 0.2, 0.5 and 1.0% respectively of the silicates in the form of crystalline silicates around the line of sight. The bottom part of each of the panel shows the residuals to each fit. The best fit is found for a composition containing 0.2% of the silicates in crystalline form, while 0 and 0.4% mark the limits of acceptable results. Figure adopted from Kemper et al. (2004).

C/O ratio in the ISM is in general smaller than unity, which normally would result in the formation of oxygen-rich material. However, the conditions for dust formation are totally different. The most important difference between the dust formation process in the interstellar medium and the dust around stars is the temperature and density of the gas. Around stars dust forms at temperature of about 1000 K, while in the ISM dust temperatures between 10 and ≈100 K are found. These low temperatures are far below the glass temperature, therefore the grains formed in the ISM will be amorphous. This is consistent with the fact that no crystalline silicates have been found in the ISM. From depletion pattern differences between cool clouds in the Galactic disk and warm clouds in the galactic halo, it seems that the dust that forms in the ISM has an olivine stoichiometric ratio (Jones, 2000). Clearly, this does not imply that they are in fact olivines, let alone crystals. 20 F. MOLSTER AND C. KEMPER

On the other side of the balance, dust destruction and processing play an important role in the ISM as well. The harsh conditions that prevail in the ISM can result in the destruction or modification of interstellar dust. Interstellar shocks may shatter grains into smaller fragments, thus changing the grain size distribution by producing more small grains, or leading to complete evaporation (Tielens et al., 1994). In addition, the degree of crystallinity goes down in the ISM. We observe these grains in the outflows of evolved stars, but not along lines of sight through the ISM. The exact abundance of the crystalline silicates brought into the ISM by stars is not known yet. This depends on how much cold crystalline silicates are hidden in the warm amorphous silicate profiles, but it is likely in the order of 15% of the silicates (Kemper et al., 2004). But the lack of crystalline silicate features in the ISM, makes it plausible that the crystalline silicates are amorphitized very quickly in the ISM, assuming destruction mechanisms are equally efficient on crystalline and amorphous silicates. In fact both destruction and amorphization are expected to occur in the ISM. It has been suggested that the amorphization might be explained by lattice defects caused by cosmic ray hits. If the ambient temperature is below the crystallization temperatures there will be not enough internal energy to overcome the threshold energy necessary to fix the defects in the crystal, and an amorphous particle track forms. Not all cosmic rays are useful though, only a select energy range can cause the amorphization of crystalline silicate grains. (Day, 1977) shows that MeV protons hardly have any influence on the crystalline silicates and electron irradiation might even make grains more crystalline (Carrez et al., 2001). The stopping power of keV ions in silicate material is high enough to cause significant damage to the lattice structure however. The main source of acceleration for these ions are SN shock waves. Laboratory experiments of the irradiation of silicates with keV ions of different nature show that the small crystalline silicates (1 μm) are likely to be amorphitized in the ISM (Demyk et al., 2001; Jäger et al., 2003; Brucato et al., 2003). This amorphization process works best on small grains since there is a limited penetration depth, only the first micron will be completely amorphitized. Larger grains might therefore survive complete amorphization, but they will suffer more from sputtering and evaporation in these shocks (Jones et al., 1996). Unfortunately, the cosmic flux of ions in the keV energy range seem to be too small to explain the fast amorphization observed in the diffuse ISM (Kemper et al., 2004), but amorphization by Fe2+ ions with keV energies or higher may be feasible on such time scales, due to the larger stopping power (E. Bringa, priv. comm.). To summarize, it is expected that the silicate grains, especially the crystalline ones, will not survive a long stay in the ISM: larger grains will be eroded by sputtering and evaporation, while smaller grains may be amorphitized and destroyed completely. The final structure of silicates exposed to He+ irradiation shows some resemblance with the GEMS (glass with embedded metals and sulfides) in IDPs (Bradley 1994b; Demyk et al., 2001). In this respect it is also interesting to note CRYSTALLINE SILICATES 21 that sometimes inside a GEMS a so-called “relict forsterite” grain is found, which shows evidence for heavy irradiation damage, much more than what is expected during their stay in the solar nebula. Furthermore, these irradiation experiments showed that O and Mg are preferentially sputtered away from the first 100 nm of the test samples. This might explain the evolution of the amorphous silicate dust which looks more like an amorphous olivine around the evolved stars and more like an amorphous pyroxene around young stars (Demyk et al., 2001). A similar selective sputtering, but then by solar wind irradiation has been observed in solar system material (Bradley, 1994a).

4.4. DUST PROCESSING IN STAR-FORMING REGIONS

Without doubt the solar system is the best evidence available for the dramatic changes in dust composition that occur during star and planet formation. The vast literature on the mineralogy of solar system objects, ranging from planets to IDPs demonstrates that large differences exist between the properties of solids in the solar system and in interstellar space. In recent years it has become evident that these differences also exist for dust in primitive solar system objects such as comets. Clearly, understanding these differences is one of the main challenges of the field of star- and planet formation. Apart from an increase in average size, infrared spectroscopy has also shown that the composition of dust in proto-planetary disks has changed compared to that in the ISM. Probably the clearest example is the crystallization of the amorphous silicates around Herbig Ae/Be stars. As said above, there is no convincing evidence for the presence of crystalline silicates in the ISM. The crystalline silicates found around young stars therefore have to be formed in situ. They will be formed both by gas-phase re-condensation, upon evaporation in the inner solar system, as well as by annealing of amorphous silicates, as is also seen in our own solar system (Molster and Bradley, 2001). Both processes will take place close to the young star. This seems to be supported by the observations. In most of these stars there are indications for the presence of crystalline silicates, but only in the 10 μm region, indicating that the crystalline silicates are indeed hot and close to the star (e.g., Meeus et al., 2001; Bouwman et al., 2001, see also Figure 4). The gas phase condensation process will produce very Fe-poor crystalline silicates, similar to what has been found around evolved stars, while the annealing process will very likely produce Fe-containing crystalline silicates, because the original amorphous silicates are likely to contain iron. Since the amorphous silicates will in general not have the stoichiometric composition of one single mineral (such as forsterite or enstatite), an annealed grain will usually be a grain consisting of two or more different minerals (see e.g., Figure 9 of Rietmeijer et al., 1999). The presence of silica (SiO2) in the dust emission spectra of young stars supports this scenario. This material is very transparent and will normally be very cold relative to the 22 F. MOLSTER AND C. KEMPER other minerals. The fact that the silica features in the 10 μm region are detected, indicates that silica is much warmer than expected. It should therefore be in thermal contact with the other minerals, as is expected when it is a product of annealing of amorphous silicates (Bouwman et al., 2001). Recently, nebular shocks have been suggested as an alternative mechanism to provide the temperatures necessary for annealing of the amorphous silicates in the proto-solar nebula at larger distances (Harker and Desch, 2002). Although, this mechanism fails to explain some of the aspects of the crystalline silicates in the solar system, and will therefore not be applicable to all crystalline silicates found, it cannot be excluded as a competing mechanism. This mechanism might also play a role in the disks around evolved stars. Unfortunately, the sensitivity of ISO was not high enough to detect crystalline silicate features in the spectra of pre-main-sequence stars of solar mass, and only the much brighter and more massive Herbig AeBe stars are known to have crystalline components based on their ISO spectroscopy. Recent ground based studies however have shown that crystalline silicates are indeed present around T Tau stars, pre- main-sequence stars with a mass comparable to that of the Sun, indicating that processing occurs in these systems as well (Honda et al., 2003; Meeus et al., 2003), a conclusion confirmed by the first results from the Spitzer Space Telescope (Uchida et al., 2004). A spectacular result of the ISO mission has been the remarkable similarity between the spectral appearance of the solar system comet Hale–Bopp (Crovisier et al., 1997) and that of the Herbig Ae/Be star HD 100546 (Malfait et al., 1998). In both objects the emission bands of crystalline silicates, as usual Fe-poor, stand out (Figure 10). The origin of the high abundance of crystalline silicates in both objects

Figure 10. The ISO SWS spectra of Hale–Bopp (Crovisier et al., 1997) and HD100546 (Malfait et al., 1998). The crystalline forsterite features are indicated. CRYSTALLINE SILICATES 23

Figure 11. The 33 μm complex of the young star HD100546 (Malfait et al., 1998), the comet Hale– Bopp (Crovisier et al., 1997) and an average of 33 μm complexes of evolved stars with evidence for an excretion disk (Molster et al., 2002a), together with two laboratory measurements of forsterite one measured in Jena (J¨ager et al., 1998) and one in Japan (Koike et al., 1999). Note the difference in width of the 33.6 μm feature. Figure taken from (Molster et al., 2002c).

is not clear. In HD 100546 the abundance of forsterite increases with distance from the star (Bouwman et al., 2003), which is not expected in the case of radial mixing of annealed silicates (Nuth et al., 2000; Harker and Desch, 2002; Bockel´ee-Morvan et al., 2002). Bouwman et al. (2003) propose a local production of small forsterite grains as a result of the collisional destruction of a large parent body. Most considerations about the formation of crystalline silicates come from theory. However, also laboratory experiments sometimes shed new light on this subject. There are indications that the width of some features are an indication of the internal structure of a crystal, and therefore of their formation history. Figure 11 shows two laboratory spectra of forsterite: the forsterite sample from Jena has been made from a melt, and is likely polycrystalline, while the sample from Japan is a single crystal. Note the difference in band width between both spectra. It is interesting to compare these band shapes to those of some characteristic bands seen in the spectra of evolved and young stars. The evolved stars, with freshly made stardust, show a band shape which resembles that of the Japan sample, i.e., corresponding to a single crystal. This may be expected, since gas-phase condensation would naturally lead 24 F. MOLSTER AND C. KEMPER to particles with a single crystal structure. On the other hand, the ISO spectrum of HD 100546 shows a resemblance to that of the polycrystalline sample from Jena. This is an indication that the forsterite in HD100546 has a polycrystalline structure. Both fractionation of a differentiated body (as proposed by Bouwman et al. (2003) for HD100546), as well as annealing of amorphous silicates (as seem to occur around many other young stars) will lead to polycrystalline forsterite and is therefore in agreement with the observations. The signal to noise of the Hale–Bopp spectrum, prevents to draw conclusions about the origin of the crystalline silicates in this comet. However, it is important to stress, that there are other effects that can inﬂuence the band shape, such as blends with other dust species, and grain size and shape effects.

5. Conclusions and Outlook

As often in astronomy, opening a new window to the universe results in a landslide of exciting discoveries, and often in directions that were not really anticipated. ISO is no exception to this: before ISO was launched, it was expected that most discoveries related to dust(formation) would be made in carbon-rich environments. It was thought that the oxygen-rich environments would only show the broad, smooth amorphous silicates. The detection of crystalline silicates in circumstellar environments came therefore as a big surprise. Of course, the crystalline silicates are just a part of the exciting world of astromineralogy, it includes many more mineral species. This area of research has gained a tremendous momentum with the launch of ISO. The spectral resolution and coverage of this satellite was perfect for the study of dust at infrared wavelengths. ISO went on where IRAS had stopped. The opening of the electromagnetic spectrum beyond 23 μm proved to be very fruitful for astromineralogy. Dust identifications were no longer based on single features, but on complete spectral mapping. It is now even possible to determine for some dust species the exact mineralogical composition, which is a big step forward in the field of astromineralogy. The ISO discovery of the crystalline silicates triggered many follow-up laboratory studies on silicate formation and processing. These studies indicate that thermal equilibrium calculations cannot explain the diversity and abundances of the amorphous and crystalline silicates. Progress can be made, especially in combination with new laboratory measurements of the infrared properties of dust species. How- ever, the requirements for the cosmic dust analogues, are quite severe (appropriate size, no contaminations and preferably at low temperatures) and these experiments are therefore not easy to perform. Laboratory experiments are an important tool on our path to understand the dusty Universe. Also radiative transfer models were suddenly confronted with the increase of spectral detail. The infrared spectra are now accurate enough that real dust species can (and should) be incorporated. The adding of more and realistic dust species leads CRYSTALLINE SILICATES 25 to more complex calculations, but also to a significant increase of our knowledge about the location (and therefore evolution) of the different dust species, a good example is the space distribution of forsterite around HD100546 (Bouwman et al., 2003). One of the most exciting new venues of dust studies is the possibility to quantitatively compare the dust in our own solar system to that seen in space. Laboratory techniques are becoming increasingly powerful and provide important information about the formation history of solar system and stardust samples. High spatial resolution observations at infrared wavelengths will allow us to study the (crystalline) silicate formation and growth in red giants and its time dependence, thus putting tight constraints on dust nucleation models. Sensitive spectrometers will allow studies of dust in stars outside our galaxy, giving insight into dust formation in external galaxies with different metallicities. Observationally, ISO is only the beginning of a development involving new missions such as Spitzer, Herschel, SOFIA and JWST, as well as improved infrared facilities at ground-based observatories such as Subaru and the VLT, which will tremendously increase our knowledge of astromineralogy, and in particular of the mineralogy of proto-planetary systems. We will be able to study the processing of silicates in newly forming planetary systems and compare directly to the history of the formation of our solar system as it is recorded in planets, asteroids and comets. Recent discoveries with Spitzer (see ApJS vol. 154, 2004) show that this telescope is capable of much deeper studies, and will detect weaker crystalline silicate features, in fainter objects. This will be particularly important for the understanding of dust processing in planet forming systems akin the solar system. The higher spatial resolution offered by the state-of-the-art ground based facilities allows for a spectral mapping of circumstellar environments (e.g., van Boekel et al., 2004). Especially with the mid-infrared interferometry available, it will be possible to determine how the degree of crystallinity varies within the circumstellar environment, which will put constraints on the dust processing mechanisms, and help disentangle the physical properties of planet forming systems. The ISO archive will, however, remain a valuable source of information about astromineralogy in general. This is because the wide spectral coverage of ISO is currently unparalleled. It is precisely this property of ISO, together with its spectral resolution, which made it possible to have so much progress in the field of astromineralogy in general and the field of crystalline silicates in particular.

Acknowledgments

We thank Rens Waters for the many useful discussions on crystalline silicates. Furthermore, we thank Jeroen Bouwman for providing Figure 4. Support for this work was provided by NASA through the Spitzer Fellowship Program, under award 011 808-001. This work is based on observations with ISO, an ESA project with 26 F. MOLSTER AND C. KEMPER instruments funded by ESA Member States (especially the PI countries: France, Germany, the Netherlands and the United Kingdom) and with the participation of ISAS and NASA.

References

Aitken, D. K., Smith, C. H., James, S. D., Roche, P. F., and Hough, J. H.: 1988, Mon. Not. R. Astron. Soc. 230, 629. Arendt, R. G., Dwek, E., and Moseley, S. H.: 1999, Astrophys. J. 521, 234. Bockelée-Morvan, D., Gautier, D., Hersant, F., Huré,J.-M., and Robert, F.: 2002, Astron. Astrophys. 384, 1107. Bouwman, J., de Koter, A., Dominik, C., and Waters, L. B. F. M.: 2003, Astron. Astrophys. 401, 577. Bouwman, J., Meeus, G., de Koter, A., Hony, S., Dominik, C., and Waters, L. B. F. M.: 2001, Astron. Astrophys. 375, 950. Bowey, J. E. and Adamson, A. J.: 2002, Mon. Not. R. Astron. Soc. 334, 94. Bowey, J. E., Lee, C., Tucker, C., Hofmeister, A. M., Ade, P. A. R., and Barlow, M. J.: 2001, Mon. Not. R. Astron. Soc. 325, 886. Bradley, J. P.: 1994a, Meteoritics 29, 447. Bradley, J. P.: 1994b, Science 265, 925. Bradley, J. P., Humecki, H. J., and Germani, M. S.: 1992, Astrophys. J. 394, 643. Brucato, J. R., Colangeli, L., Mennella, V.,Palumbo, P., and Bussoletti, E.: 1999, Planetary and Space Science 47, 773. Brucato, J. R., Strazzulla, G., Baratta, G., and Colangeli, L.: 2003, Astron. Astrophys. 413, 395. Cami, J.: 2002, Molecular gas and dust around evolved stars. PhD thesis, University of Amsterdam, The Netherlands. Cami, J., de Jong, T., Justtannont, K., Yamamura, I., and Waters, L. B. F. M.: 1997, Astrophys. Space Sci. 255, 339. Carrez, P., Demyk, K., Leroux, H., and Cordier, P.: 2001, Meteoritics & Planetary Science, Supple- ment, p.A36, 36, 36. Cesarsky, D., Jones, A. P., Lequeux, J., and Verstraete, L.: 2000, Astron. Astrophys. 358, 708. Crovisier, J., Leech, K., Bockelee-Morvan, D., Brooke, T. Y., Hanner, M. S., Altieri, B., et al.: 1997, Science 275, 1904. Day, K. L.: 1977, Mon. Not. R. Astron. Soc. 178, 49P. Demyk, K., Carrez, P., Leroux, H., Cordier, P., Jones, A. P., Borg, J., et al.: 2001, Astron. Astrophys. 368, L38. Dorschner, J., Begemann, B., Henning, T., Jäger,C., and Mutschke, H.: 1995, Astron. Astrophys. 300, 503. Douvion, T., Lagage, P. O., and Cesarsky, C. J.: 1999, Astron. Astrophys. 352, L111. Douvion, T., Lagage, P. O., and Pantin, E.: 2001, Astron. Astrophys. 369, 589. Dunne, L., Eales, S., Ivison, R., Morgan, H., and Edmunds, M.: 2003, Nature 424, 285. Dwek, E.: 2004, Astrophys. J.. in press. Fajardo-Acosta, S. B. and Knacke, R. F.: 1995, Astron. Astrophys. 295, 767. Gail, H.-P.and Sedlmayr, E.: 1998, Chemistry and Physics of Molecules and Grains in Space, Faraday Discussions No.: 109, p. 303. Gail, H.-P. and Sedlmayr, E.: 1999, Astron. Astrophys. 347, 594. Habing, H.: 1996, Astron. Astrophys. Rev. 7, 97. Hallenbeck, S. L., Nuth, J. A., and Daukantas, P. L.: 1998, Icarus 131, 198. CRYSTALLINE SILICATES 27

Hanner, M. S.: 1996, Composition and optical properties of cometary dust. In ASP Conf. Ser.: 104: IAU Colloq.: 150: Physics, Chemistry, and Dynamics of Interplanetary Dust, page 367. Hanner, M. S., Lynch, D. K., and Russell, R. W.: 1994, Astrophys. J. 425, 274. Harker, D. E. and Desch, S. J.: 2002, Astrophys. J. Lett. 565, L109. Honda, M., Kataza, H., Okamoto, Y. K., Miyata, T., Yamashita, T., Sako, S., et al.: 2003, Astrophys. J. 585, L59. Hoogzaad, S. N., Molster, F. J., Dominik, C., Waters, L. B. F. M., Barlow, M. J., and de Koter, A.: 2002, Astron. Astrophys. 389, 547. Jäger, C., Fabian, D., Schrempel, F., Dorschner, J., Henning, T., and Wesch, W.: 2003, Astron. Astro- phys. 401, 57. Jäger, C., Molster, F. J., Dorschner, J., Henning, T., Mutschke, H., and Waters, L. B. F. M.: 1998, Astron. Astrophys. 339, 904. Jeong, K. S., Sedlmayr, E., and Winters, J. M.: 2000, Circumstellar Dust Shells around Oxygen-rich Long-period Variable Stars. In Astronomische Gesellschaft Meeting Abstracts, volume 17, page 30. Jones, A. P.: 2000, J. Geoph. Res. 105, 10257. Jones, A. P., Tielens, A. G. G. M., and Hollenbach, D. J.: 1996, Astrophys. J. 469, 740. Jones, T. W. and Merrill, K. M.: 1976, Astrophys. J. 209, 509. Kemper, F., de Koter, A., Waters, L. B. F. M., Bouwman, J., and Tielens, A. G. G. M.: 2002a, Astron. Astrophys. 384, 585. Kemper, F., Jäger,C., Waters, L. B. F. M., Henning, T. H., Molster, F. J., Barlow, M. J., et al.: 2002b, Nature 415, 295. Kemper, F., Vriend, W. J., and Tielens, A. G. G. M.: 2004, Astrophys. J. 609, 826. Kemper, F., Waters, L. B. F. M., de Koter, A., and Tielens, A. G. G. M.: 2001, Astron. Astrophys. 369, 132. Klein, C. and Hurlbut, Jr. C. S.: 1993, Manual of Mineralogy (After James D. Dana), 21st edition. John Wiley & Sons, Inc. Knacke, R. F., Fajardo-Acosta, S. B., Telesco, C. M., Hackwell, J. A., Lynch, D. K., and Russell, R. W.: 1993, Astrophys. J. 418, 440. Koike, C., Shibai, H., and Tuchiyama, A.: 1993, Mon. Not. R. Astron. Soc. 264, 654. Koike, C., Tsuchiyama, A., Shibai, H., Suto, H., Tanabé,T., Chihara, H., et al.: 2000, Astron. Astrophys. 363, 1115. Koike, C., Tsuchiyama, A., and Suto, H.: 1999, The spectra of crystalline silicates in infrared region. In Proceedings of the 32nd ISAS Lunar and Planetary Symposium, pages 175–178. Kozasa, T., Hasegawa, H., and Nomoto, K.: 1991, Formation of dust grains in the ejecta of SN 1987A. II. Astron. Astrophys. 249, 474. MacKinnon, I. D. R. and Rietmeijer, F. J. M.: 1987, Reviews of Geophysics 25, 1527. Malfait, K., Waelkens, C., Waters, L. B. F. M., Vandenbussche, B., Huygen, E., and de Graauw, M. S.: 1998, Astron. Astrophys. 332, L25. Meeus, G., Sterzik, M., Bouwman, J., and Natta, A.: 2003, Astron. Astrophys. 409, L25. Meeus, G., Waters, L. B. F. M., Bouwman, J., van den Ancker, M. E., Waelkens, C., and Malfait, K.: 2001, Astron. Astrophys. 365, 476. Mennella, V., Brucato, J. R., Colangeli, L., Palumbo, P., Rotundi, A., and Bussoletti, E.: 1998, Astro- phys. J. 496, 1058. Molster, F. J. and Bradley, J. P.: 2001, Meteoritics & Planetary Science, Supplement 36, A140. Molster, F. J., Lim, T. L., Sylvester, R. J., Waters, L. B. F. M., Barlow, M. J., Beintema, D. A., et al.: 2001a, Astron. Astrophys. 372, 165. Molster, F. J., Waters, L. B. F. M., and Tielens, A. G. G. M.: 2002a, Astron. Astrophys. 382, 222. Molster, F. J., Waters, L. B. F. M., Tielens, A. G. G. M., and Barlow, M. J.: 2002b, Astron. Astrophys. 382, 184. 28 F. MOLSTER AND C. KEMPER

Molster, F. J., Waters, L. B. F. M., Tielens, A. G. G. M., Koike, C., and Chihara, H.: 2002c, Astron. Astrophys. 382, 241. Molster, F. J., Waters, L. B. F. M., Trams, N. R., Van Winckel, H., Decin, L., van Loon, J. T., et al.: 1999a, Astron. Astrophys. 350, 163. Molster, F. J., Yamamura, I., Waters, L. B. F. M., Nyman, L. A., Käufl,H. U., de Jong, T., et al.: 2001b, Astron. Astrophys. 366, 923. Molster, F. J., Yamamura, I., Waters, L. B. F. M., Tielens, A. G. G. M., de Graauw, T., de Jong, T., et al.: 1999b, Nature 401, 563. Nuth, J. A., Hill, H. G. M., and Kletetschka, G.: 2000, Nature 406, 275. Onaka, T. and Okada, Y.: 2003, Astrophys. J. 585, 872. Ossenkopf, V., Henning, T., and Mathis, J. S.: 1992, Astron. Astrophys. 261, 567. Posch, T., Kerschbaum, F., Mutschke, H., Dorschner, J., and Jäger,C.: 2002, Astron. Astrophys. 393, L7. Rietmeijer, F. J. M., Nuth, J. A., and Karner, J. M.: 1999, Astrophys. J. 527, 395. Simis, Y. J. W.: 2001, Mass loss modulation in dust forming stellar winds. PhD thesis, Leiden Obser- vatory, The Netherlands. Sogawa, H. and Kozasa, T.: 1999, Astrophys. J. 516, L33. Sylvester, R. J., Kemper, F., Barlow, M. J., de Jong, T., Waters, L. B. F. M., Tielens, A. G. G. M., et al.: 1999, Astron. Astrophys. 352, 587. Tielens, A. G. G. M., McKee, C. F., Seab, C. G., and Hollenbach, D. J.: 1994, Astrophys. J. 431, 321. Tielens, A. G. G. M., Waters, L. B. F. M., Molster, F. J., and Justtanont, K.: 1997, Astrophys. Space Sci. 255, 415. Uchida, K. I., Calvet, N., Hartmann, L., Kemper, F., Forrest, W. J., Watson, D. M., et al.: 2004, Astrophys. J. Suppl. 154, 439. van Boekel, R., Waters, L. B. F. M., Dominik, C., Dullemond, C. P., Tielens, A. G. G. M., and de Koter, A.: 2004, Astron. Astrophys. 418, 177. Waelkens, C., Waters, L. B. F. M., de Graauw, M. S., Huygen, E., Malfait, K., Plets, H., et al.: 1996, Astron. Astrophys. 315, L245. Waters, L. B. F. M., Cami, J., de Jong, T., Molster, F. J., van Loon, J. T., Bouwman, J., et al.: 1998, Nature 391, 868. Waters, L. B. F. M., Molster, F. J., de Jong, T., Beintema, D. A., Waelkens, C., Boogert, A. C. A., et al.: 1996, Astron. Astrophys. 315, L361. Wooden, D. H., Harker, D. E., Woodward, C. E., Butner, H. M., Koike, C., Witteborn, F. C., et al.: 1999, Astrophys. J. 517, 1034. WATER IN SPACE: THE WATER WORLD OF ISO

JOSE´ CERNICHARO1,∗ and JACQUES CROVISIER2 1CSIC, Instituto de Estructura de la Materia, Department of Molecular and Infrared Astrophysics, C/Serrano 121, 28006 Madrid, Spain 2Observatoire de Paris, 5 place Jules Janssen, F-92195 Meudon, France (∗Author for correspondence: E-mail: [email protected])

(Received 8 April 2005; Accepted in ﬁnal form 14 April 2005)

Abstract. In this review we present the main results obtained by the ISO satellite on the abundance and spatial distribution of water vapor in the direction of molecular clouds, evolved stars, galaxies, and in the bodies of our Solar System. We also discuss the modeling of H2O and the difﬁculties found in the interpretation of the data, the need of collisional rates and the perspectives that future high angular and high spectral resolution observations of H2O with the Herschel Space Observatory will open.

Keywords: water, molecular clouds, evolved stars, galaxies, comets, Mars, giant planets, Titan

1. Introduction

The determination of the abundance of water vapor is a long standing problem in astrophysics. Because water vapor is predicted to be an abundant molecule in the gas phase, a determination of its spatial extent, its distribution, and its abundance is crucial in modeling the chemistry and the physics of molecular clouds. In warm molecular clouds, water vapor could play a critical role in cooling the gas (Neufeld and Kaufman, 1993; Neufeld et al., 1995) and, hence, in the evolution of these objects. In spite of the important astrophysical output that could be obtained from near, mid and far-infrared observations of water vapor (see e.g., Cernicharo, 1998; van Dishoeck, 2004), the wavelength range 2–200 μm was poorly studied prior to ISO due to the large absorption at H2O frequencies produced by water vapor in the Earth’s atmosphere. Some ground-based and airborne observations of H2O have been performed, but most of the observed transitions are maser in nature making the analysis of the data very difficult (Waters et al., 1980; Cernicharo et al., 1990, 1994, 1996a; Menten and Melnick, 1991; González-Alfonso et al., 1995, 1998a; Cernicharo et al., 1999). 18 Even the observation of the H2 O isotope gives only poor estimates on x(H2O) due to contamination of the observed transitions by other molecular lines (Jacq et al., 1988; Gensheimer et al., 1996), or to the observation of only one line which avoids the correct determination of the excitation conditions of the 18O isotopic species of water vapor (Zmuidzinas et al., 1995). HDO has also been observed (see e.g., Jacq Space Science Reviews (2005) 119: 29–69 DOI: 10.1007/s11214-005-8058-x C Springer 2005 30 J. CERNICHARO AND J. CROVISIER et al., 1990; Helmich et al., 1996a; Pardo et al., 2001b; Comito et al., 2003), but it is also difficult to derive the H2O abundance from this rare isotopic species. Some efforts have been made as well to derive the water abundance in molecular clouds + from the observation of related species like H3O (Phillips et al., 1992; Goicoechea and Cernicharo, 2001a). However, the determination of the H2O abundance from these observations is not straightforward. The 616–523 masing transition of H2O at 22 GHz, which arises from levels around 700 K, has been used since its detection by Cheung et al. (1969), to trace high excitation gas around star-forming regions and evolved stars. The size of the emitting regions at this frequency is typically of the order of a few milliarcsec- onds (a few 1013 cm at the distance of Orion). Hence, no information has been obtained on the role of water vapor at large spatial scales from this line. Other H2O lines have been detected from ground or airborne based telescopes like the 313–220 transition at 183 GHz (Waters et al., 1980; Cernicharo et al., 1990, 1994, 1996a; González-Alfonso et al., 1995, 1998a), the 414–321 transition at 380 GHz (Phillips et al., 1980), the 1029–936 transition at 321 GHz (Menten et al., 1990a) and the 515–422 transition at 325 GHz (Menten et al., 1990b; Cernicharo et al., 1999). Among these lines only the 313–220 transition at 183 GHz, for which the terrestrial atmosphere can be partially transparent under excellent weather conditions (see Cernicharo, 1985, 1988; Cernicharo et al., 1990; Pardo, 1996; Pardo et al., 2001a for details on atmospheric transmission at mm and submm wavelengths), has been used to map the emission of water at very large spatial scales (Cernicharo et al., 1994, 1996a, 1997b,c). As an example, the map of the Orion molecular cloud shown in Cernicharo et al. (1994) is 6 orders of magnitude larger than the size of the spots detected at 22 GHz. In that work the water abundance was estimated for the different large-scale components of the Orion molecular cloud for the first time. However, due to the maser nature (although relatively weak) of this water transition, the interpretation of the data depends critically on the assumed kinetic temperature of the gas. In Orion, Cernicharo et al. (1994) estimate that changing the temperature of the emitting gas from 60 K (the assumed value) to 150 K the water abundance required to fit the data would be reduced by 1 order of magnitude in the extended ridge molecular cloud. The ISO mission (Kessler et al., 1996), and especially, the Long Wavelength Spectrometer, LWS (Clegg et al., 1996), and the Short Wavelength Spectrometer, SWS (de Graauw et al., 1996), have provided a unique opportunity to observe several H2O lines in a large variety of astronomical environments. The SWS spectrometer covers the region 2–45 μm and is well adapted to study ro-vibrational transitions (stretching and bending modes). On the other hand, the LWS spectrometer provides a powerful tool to observe the pure rotational transitions of light molecular species like H2O. The majority of the observations of H2O pure rotational lines were performed at the low spectral resolution of the LWS grating mode (∼1000 km s−1), which produces strong spectral dilution in the search for molecular features in most ISM sources. Nevertheless, some molecular clouds, such as Orion and SgrB2, have WATER IN SPACE 31 been observed in many water lines using the much higher resolution of the Fabry- Pérot coupled to the SWS and LWS instruments (∼10 km s−1 for the SWS/FP and ∼35 km s−1 for the LWS/FP; see Figures 2 and 3). The instruments on board the Infrared Space Observatory (ISO) have provided the most complete opportunity to study the emission/absorption of water and other important molecular species through transitions which are inaccessible from the ground or airborne platforms. Other space observatories, like Odin (see Hjalmarson et al., 2003), and the Submillimeter Wave Astronomy Satellite (SWAS; Melnick et al., 2000), have also provided data on the ground transition of o-H2O at 557 GHz. These two satellites provide a higher spectral resolution than ISO, with the drawback of having access to only one line of water vapor and a lower angular resolution. The higher angular resolution provided by ISO and the possibility to cover the full 2–200 μm range, i.e., the pure rotational spectrum of H2O and the ro-vibrational transitions associated to its bending and stretching modes, makes ISO the most powerful instrument available so far for the study of water vapor. In this review, we present the main results on gas phase H2O obtained with the ISO-SWS and ISO- LWS spectrometers in the direction of molecular clouds, evolved stars, galaxies and solar system bodies.

2. Water in the ISM

ISO has provided observations of H2O in a large variety of objects of the ISM (see e.g., van Dishoeck, 2004, for a review on ISO spectroscopy). The low spectral resolution data from the LWS have been used to study the water emission in shocks, photo dissociation regions, and high and low-mass protostars. Figure 1 shows the far-infrared spectrum at several positions of the class 0 source L1448 (Nisini et al., 2000). These data show the presence of H2O, OH, CO and atomic ﬁne structure lines. The far-infrared cooling is mainly due to the emission of these species, together with H2 in the mid and near infrared. The ISO observations probe a gas of very different physical conditions to those determined from ground-based telescopes. The water molecules and H2 seem to be excited to temperatures between 500 and 1200 K. Probably this emission arises from very thin layers of shocked gas along the jet emerging from L1448-mm. In spite of the low spectral resolution of the LWS grating spectrometer, the sensitivity of ISO allows to study this gas in detail. The low and medium-J lines of CO that have been observed from ground-based observatories trace a colder gas arising from different regions of the shocks. The number of low-mass protostars studied with the LWS is rather impressive. In most cases the observed water emission has been interpreted as arising from shocks along the molecular jet (see above for L1448; Moro-Martin et al., 2001, for CepE; Ceccarelli et al., 1999a, for IRAS16293-2422; Nisini et al., 1996, and Liseau et al., 1996, for HH54; Benedettini et al., 2000, for HH24-26; etc., see the 32 J. CERNICHARO AND J. CROVISIER

Figure 1. ISO/LWS continuum-substracted spectra in the direction of L1448-IRS3 and L1448-mm and the red lobe of the outﬂow associated to L1448-mm. The lines from OH, H2O and CO, together with atomic ﬁne structure lines, are indicated in the different panels (from Nisini et al., 2000).

review by Nisini, 2003, and Nisini et al., 2002). Water vapor emission was also found in the direction of IC1396N, a bright rimmed globule, by Saraceno et al. (1996). The LWS water emission measured toward young low-mass stars have also been interpreted, using sophisticated models, as arising from the envelopes that surround the protostars (see e.g., Ceccarelli et al., 1999b). Maret et al. (2002) have modeled the water, CO and OH lines in the far-infrared in the direction of NGC1333-IRAS4 and conclude that the observed emission is reasonably well reproduced by their models. However, the lack of angular resolution and the poor sensitivity of most water lines to the physical parameters, avoid a clear assignment of the observed emission to the very small regions surrounding the newly formed protostars. Her- schel will provide a more detailed view of the origin of water vapor in these objects. Nevertheless, the 183.3 GHz observations of water by Cernicharo et al. (1996a) toward HH7-11 and L1448 seem to indicate that H2O is present in both the shocks along the molecular jets and the central source. WATER IN SPACE 33

High spectral resolution (from 7 to 40 km s−1) with ISO is available for just a few molecular clouds. The SWS/FP instrument has been used mainly to observe the bending mode of water in the direction of high-mass protostars and to observe pure rotational lines of H2O in Orion (see e.g., Wright et al., 2000). The LWS/FP has been used to observe pure rotational lines of water in a large variety of molecular clouds. Figure 3 shows the LWS/FP results in the direction of Sgr B2(M). Lines from OH, H2O, C3, NH, NH2,NH3, CH, HD, and CH2 and atomic ﬁne structure features have been detected (Cernicharo et al., 2000; Goicoechea and Cernicharo, 2001a,b, 2002; Goicoechea et al., 2004; Ceccarelli et al., 2002; Polehampton et al., 2002, 2005). These observations provide a spectral resolution of 35 km s−1, which is enough to probe molecular clouds with broad lines, but precludes the study of more quiescent clouds with narrow lines. The Herschel Space Observatory will get around this problem thanks to its high spectral resolution heterodyne receiver HIFI. Relevant observations of water with the SWS/FP instrument have been reported by Wright et al. (2000, see Figure 2). Water ice has been observed toward a large number of molecular clouds. Gibb et al. (2004) give an overview of the stretching and bending water ice features in clouds of the ISM. The review by van Dishoeck (2004) provides a detailed view on the gas and ice spectroscopy that ISO has performed.

Figure 2. SWS Fabry-Perot spectra toward IRc2 of a selection of H2O absorption lines. The ﬁnal panel displays the 5.3–7.0 μm grating spectrum of the H2O ν2 = 0 → 1 (see Section 2.2 for the analysis of this feature; from Wright et al., 2000). 34 J. CERNICHARO AND J. CROVISIER and atomic ﬁne structure lines are 3 , and NH 2 , CH, NH, NH + O 3 ,H 3 O, C 2 (2004). et al. ISO/LWS/FP observations of Sgr B2(M). The lines from OH, H Figure 3. indicated from Goicoechea WATER IN SPACE 35

We are going to present in more detail selected ISO data on prototypical ISM objects such as Orion, SgrB2 and cold dark clouds. The goals are to show how different the water line proﬁles could be depending on the physical conditions of the cloud, the difﬁculties in getting physical parameters (a benchmarking of radiative transfer codes for water vapor was held in 2004 at the Lorentz Center in Leiden; van der Tak et al., 2005), and the need to carefully select the water lines we could wish to observe with future space observatories, like Herschel, in order to have a correct view of the excitation conditions of water and to get a good estimate of the physical conditions of the clouds.

2.1. THE PURE ROTATIONAL LINES OF WATER

2.1.1. Rotational Water Vapor Lines in Emission: Orion The Orion Kleinmann-Low nebula (Orion–KL) is the nearest (at a distance of ∼450 pc) and probably the most studied star-forming region in the sky (see review, e.g., Genzel and Stutzki, 1989). The core of Orion–KL is associated with several places of massive star formation, and thus, the cloud is heavily influenced at large scales by violent phenomena such as outflows and shocks. The brightest objects are BN and IRc2, which are separated ≤10. Among these sources, IRc2 has the 5 largest intrinsic luminosity (L ∼ 10 L). The fact that multiple gas components with different physical conditions and velocity fields overlap along the line of sight complicates the interpretation of data from Orion–KL. This also results in a segregation of the gas and dust chemistry: different species and different reaction patterns dominate in the different regions (e.g. Blake et al., 1987). Many molecular observations from the millimeter to IR domains have contributed to the characterization of the three main gas components towards Orion IRc2. The Hot Core, centered ∼2 south of IRc2 with a diameter of ∼10 is probably an ensemble of warm and dense (≤107 cm−3) clumps. Spectral lines arising −1 −1 from the hot core show line widths of v = 5–15 km s at vLSR 6kms . The outflows or the plateau, with a diameter of ≤40, give a highly anisotropic emission from gas flows at high- (∼100 km s−1) and low-velocity (∼18 km s−1) with −1 −1 v = 30–150 km s and centered vLSR 9kms . These flows encounter inhomogeneous density media in their way out, producing shocked regions, and finally plunge into the extended cloud. In particular, shocks produced by the interaction of the high-velocity flow and the surrounding cloud at ∼30–40 from IRc2, produce the bright H2 regions known as Peak1 and Peak2 (first detected by Beckwith et al., 1978). Finally the whole region is embedded in a quiescent and extended component (the molecular Ridge) with narrow line-widths v = 3–5 km s−1. All these inhomogeneous components are spatially present in the ISO LWS and SWS beams, and may thus contribute to the emerging mid- and far-IR spectrum of water vapor. Several ground-based observations of water maser lines have been performed in the Orion–KL region. From VLBI observations of the o-H2O line at ∼22 GHz, Genzel et al. (1981) determined the kinematics and velocity of the 18 km s−1 flow, 36 J. CERNICHARO AND J. CROVISIER while the widespread nature of the water vapor around IRc2 has been probed with maps of water at ∼183 GHz (Cernicharo et al., 1990, 1994). This was the first time that water vapor was detected, and its abundance estimated, in the different large-scale components of Orion IRc2. In addition, the high excitation of the plateau gas has also being revealed by observations of water at ∼325 GHz (Menten et al., 1990b; Cernicharo et al., 1999). Still, the interpretation of the excitation mechanisms of these maser lines is not obvious. The HDO species has also been detected and modeled towards Orion (Jacq et al., 1990; Pardo et al., 2001b), but is not straightforward to obtain water abundances without taking into account the chemistry of each gas component. ISO spectrometers have provided the unique opportunity of observing many IR thermal lines towards Orion. Figure 4 shows the LWS data in the direction of Orion. Several positions were observed between 45 and 198 μm providing information on atomic fine structure lines as well as on CO and H2O pure rotational lines (see Sempere et al., 2000 for an analysis of the CO data). This figure shows that water vapor is present over a region of several arcminutes in size (see also the 183.3 GHz observations of Cernicharo et al., 1994). The intensity of the water lines decreases from the center to the surrounding positions by a factor of 10, which corresponds to a variation in both the water abundance and the physical conditions of the gas. The large spatial distribution of water emission makes the interpretation of the water lines extremely difficult. In addition, this large-scale map has been obtained in the grating mode of the LWS spectrometer, i.e., with a spectral resolution of 300–600. Hence, no information is available on line profiles. Harwit et al. (1998) have published the results on 8 lines of water observed with the LWS/FP between 71 and 125 μm (see Figure 5) and conclude that these lines arise from a molecular cloud subjected to a magnetohydrodynamic C-type shock. The lines indicate a water abundance of the order of 5 × 10−4. However, the fine interpretation of these lines is not obvious as radiative pumping from IR photons could be playing an important role in the excitation of water. Moreover, the collisional rates for the temperatures prevailing in the shock are poorly known. One will need a larger set of lines, including those from the isotopic species of water. Some additional lines were observed with the ISO/SWS/FP spectrometer by Wright et al. (2000; see Figure 2). Opposite to the long wavelength water line profiles, the lines appear in absorption below 40 μm. The authors have obtained a water abundance of 2–5 × 10−4 for a kinetic temperature of 200–350 K in the warm shocked gas. The turnover point between absorption and emission is an important clue in interpreting the data (opacity, line strength, spatial distribution of the gas, ...). Cernicharo and co-workers have observed all water lines in the LWS 18 17 domain with the FP including several lines of the isotopic species H2 O and H2 O. In addition, they have made a map with the LWS/FP in two water lines in order to constrain the models of the different excitation processes of water. These authors have detected more than 60 pure rotational thermal lines with the LWS Fabry-Pérot 18 17 on board ISO, including also several lines of the H2 O and H2 O isotopologues WATER IN SPACE 37

Figure 4. LWS raster of the Orion molecular cloud between 140 and 197 μm. Offsets (in arcseconds) are indicated in each box (in declination for the vertical strip and in right ascension for the horizontal one). Line identiﬁcation are given in the zoom to the central position in the top-left panel. The CO data, including the LWS/FP data, have been presented and analyzed by Sempere et al. (2000).

of H2O. Figure 6 shows a selection of the observed lines together with the low- resolution spectrum of Orion. Water lines with upper energy levels below ∼1500 K have been detected by ISO towards Orion IRc2. Nevertheless, the spatial/spectral resolution of these thermal water IR lines is not enough to resolve the complexity of the observed regions. As an example, it is difﬁcult to separate narrow from broad water line components or to have the sensitivity to trace the line-wing emission. Besides, several radiative 38 J. CERNICHARO AND J. CROVISIER

Figure 5. Water line proﬁles of the observed transitions in the direction of Orion (from Harwit et al., 1998).

transfer effects (e.g., self-absorption and scattering) may possibly occur at velocity scales that cannot be resolved by ISO. Due to the complexity of the region and the different excitation conditions prevailing along the line of sight, many water lines have to be taken into account in the 16 μ models. The H2 O lines with wavelengths above 100 m are seen in emission. How- ever for shorter wavelengths, lines arising from energy levels below ∼700 K and with large line strengths show P-Cygni profiles and have a large emission/absorption velocity range, while those with weak line strengths, or arising from higher energy levels, are observed in pure emission. Note that most water lines detected by ISO below ∼45 μm (van Dishoeck et al., 1998; Wright et al., 2000) are observed in pure absorption. The peak LSR velocity of the emission lines occurs around +20 km s−1, while that of pure absorption lines occurs at −8kms−1 (Wright et al., 2000). This means that the filling factors for water vapor as seen by the LWS and SWS beam apertures are different, i.e., the water emission observed by LWS-FP is likely arising from a larger surface than that observed with the SWS-FP.In fact, the far-IR thermal emission in the water line at ∼179.5 and ∼174.6 μm, respectively, is clearly more extended than the cloud position represented by Orion IRc2 (Cernicharo et al.,in press). The widths of the far-IR water lines are broader than the instrumental width determined by the Lorentzian instrument response, thus, the water lines are partially resolved. WATER IN SPACE 39 veral water lines in the (in press). The central panel shows the grating et al. by Cernicharo O lines. The individual panels show the LWS/FP observations of se 2 μ m. Water line profiles of the observed transitions in the direction of Orion low-resolution spectrum with a zoom showing some CO and H Figure 6. range 45–180 40 J. CERNICHARO AND J. CROVISIER

The angular extent and the spatial origin of the far-IR water lines are the main drawbacks when trying to model the spatially unresolved ISO observations. Besides, any detailed fit to the data requires a detailed knowledge of the geometry and of the relative distribution of dust and water. Fortunately, the high angular resolution maps of water obtained from ground-based observations of the 183.3 and 325 GHz lines do reveal the spatial and velocity distribution of water vapor (Cernicharo et al., 1994, 1999). These authors detected at least four different water components from which the millimeter and submillimeter water lines arise: the extended molecular Ridge, the high-velocity outflow, the low-velocity outflow and very narrow and strong features which are also associated with the small bullets observed at 22 GHz (Genzel et al., 1981). The interpretation of these lines requires detailed non-LTE and non-local statistical equilibrium calculations and radiative transfer models including both gas and dust (Daniel et al., in press). Figure 7 shows these models (including gas and dust) for different water vapor abundances. The original code has been described in detail by González-Alfonsoand Cernicharo (1993, 1997, 1999). Models assume that most of the water emission/absorption arise from the low-velocity Plateau, although the effect of the lower excitation extended Ridge has also been investigated. A spherical geometry for an outflow expanding at 18 km s−1 consisting of two cloud components of different temperatures and densities has been adopted. These gas and dust layers surround a 10 continuum IR source (the hot core), with color temperatures between 200 and 300 K. The effect of this continuum radiation and the effect of the dust opacity, together with collisions, has a deep impact in the population of H2O rotational levels (Cernicharo et al., in press; Daniel et al.,in press). The models reproduce quantitatively and qualitatively the observed profiles and intensities of the water lines and the IR continuum detected by ISO. However, some water vapor lines are sensitive to the water abundance, while others are sensitive to the adopted geometry. Very high abundances, 3×10−4, are needed to explain most lines below 40 μm, while abundances between 3 × 10−5 and 3 × 10−4 are required for the long wavelength lines. In addition, some lines are difficult to be reproduced by the models. Calibration errors and/or inaccuracy in the collisional rates could be responsible for these effects. The water abundance, as derived from ISO and ground-based observations, is consistent with that derived by Wright et al. (2000) and Harwit et al. (1998), a few 10−4 for the shocked regions, and decreases to 10−5–10−6 for the large-scale components. Any realistic model of water for clouds such as Orion has to properly take into account the phenomena described above. The high angular and spectral resolution that will be provided by Herschel will allow us to improve the models for water vapor in Orion.

2.1.2. Rotational Water Vapor Lines in Absorption: Sgr B2 7 Sgr B2 is the most massive cloud in the Galaxy, ∼10 M (Lis and Goldsmith, 1990). This cloud is a paradigmatic object in the Galactic Center region as its geometrical properties, physical conditions and chemical characteristics resemble WATER IN SPACE 41

Figure 7. Some of the observed water lines (black thick line) together with the models discussed in the text which correspond to water abundances (1, 2.5, 5, 10, 30) 10−5. See text for explanation of the effect of water abundance on the line profiles (from Cernicharo et al., in press and Daniel et al., in press). 42 J. CERNICHARO AND J. CROVISIER a miniature galactic nucleus with a ∼15 extend (Goicoechea et al., 2004). Opposite to what is found toward other star-forming regions, such as Orion (see above and Cernicharo et al., 1998a), the observations of the 212–101 line at ∼179.5 μm in the direction of Sgr B2 show that the line appears in absorption rather than in emission (Cernicharo et al., 1997a). Later, the launch of SWAS allowed the observation 16 18 of the 110–101 fundamental transition of both H2 O at 557 and H2 O at 548 GHz (first detected by the Kuiper Airborne Observatory, KAO; see Zmuidzinas et al., 1995). Although the velocity resolution is ∼1kms−1, the large beam of SWAS (∼4) could make these observations mostly sensitive to the extended cold and less dense gas. SWAS observations have provided an estimate of the H2O abundance in the low excitation clouds located in the line of sight toward Sgr B2 (Neufeld et al., 2000, 2003). However, the fact that only the ground-state absorption line is detected makes difficult a detailed study of the water vapor excitation mechanisms in Sgr B2. Cernicharo et al. (in press) get around this problem with their analysis of the far-IR observations of 14 thermal lines of water vapor toward Sgr B2(M) with the ISO/LWS/FP spectrometer, and the first mapping of the 183.3 GHz maser emission of p-H2O around Sgr B2 main condensations. The far-IR spectrum of Sgr B2(M) is dominated by the absorption produced by H2O, NH3,C3, OH, and NH (see Cernicharo et al., 1997a, 2000; Ceccarelli et al., 2002; Goicoechea et al., 2001b, 2004). The low-resolution spectrum obtained with the LWS, together with the available data from the LWS/FP, have been analyzed by Goicoechea et al. (2004; see Figure 21 for a comparison of the grating spectrum of Sgr B2(M) and Arp220). The dust emission could be fitted with a Tdust 30–40 K with a peak opacity at 80 μmof 5. A detailed view of the absorption lines found in SgrB2 in the 40–200 μm range is shown in Figure 3, while those of water vapor are shown in Figure 8. Except the 212–101 ground state line at ∼179.5 μm, all water lines have the same profile and are centered at Sgr B2(M) velocities. The ∼179.5 μm line appears saturated and in absorption from −150 to +100 km s−1 indicating that it is tracing both cold water vapor located in the foreground gas toward the GC and the warm gas around Sgr B2(M). The widespread absorption produced by the 212–101 line has been previously presented by Cernicharo et al. (1997a), while the 110–101 absorption has been mapped with SWAS (Neufeld et al., 2003). These observations probe that low excitation H2O is present in the clouds along the line of sight of Sgr B2. The averaged velocity from all H2O lines observed with the ISO/LWS–FP is +50 ± 5kms−1, in agreement with the velocity of other + related O-bearing species such as H3O or OH (Goicoechea and Cernicharo, 2001a, 2002). Taking into account the wavelength calibration accuracy of the LWS/FP instrument, this velocity is in reasonable agrement with the expected +65 km s−1 cloud seen in radio observations (Hüttemeister et al., 1995). Note, however, that velocities close to +50 km s−1 are also associated with gas surrounding most of the Sgr B2 continuum sources, so that the bulk of the H2O absorption can arise from this specific LSR component (note also that the 183.3 GHz line appears in some positions at this velocity; see Figure 9). Possible overlaps with other molecular WATER IN SPACE 43

16 18 Figure 8. ISO/LWS/FP observations of H2 O and H2 O toward Sgr B2(M) and rotational energy 16 diagram of o-H2 O. The ordinate scale corresponds to F/Fc and the abscissa to the wavelength in microns. All the far-IR water lines appear in absorption. The emission line corresponds to the 313–220 16 maser transition of p-H2 O at 183 GHz (observed with the IRAM-30 m telescope at Pico Veleta, Spain) (from Cernicharo et al., in press). species occur at some wavelengths. In particular, the o-H2O432–423 line is blended = μ 18 with HF J 2–1 at 121.697 m (Neufeld et al., 1997) and the o-H2 O212–101 line + at 181.053 μm has a small contribution from the H3O Q(1, 1) line (Goicoechea and Cernicharo, 2001a). 44 J. CERNICHARO AND J. CROVISIER

16 Figure 9. Map of the 313–220 line of p-H2 O at 183 GHz around Sgr B2 main condensations. Contours (in K km s−1) are indicated in the figure. Different line profiles at significant positions of ∗ −1 the map are also shown. The intensity scale is in TA and the abscissa is the LSR velocity in km s . The map is centered at the position of Sgr B2(M) (from Cernicharo et al., in press).

Observations of the H2O313–220 line (El 200 K) at ∼183 GHz toward Sgr B2 are presented in Figure 9 (integrated line intensity over the observed region and several spectra at relevant positions). The emission appears at the LSR velocities of Sgr B2 with no contribution from other clouds along the line of sight. Note the different line shapes and intensities of the ∼183 GHz emission for positions in front and around the main condensations. Different models (LVG and non-local) for the water vapor absorption have been presented in Cernicharo et al. (in press). The authors conclude that the LVG method is not well suited to treat the radiative transfer of H2O in Sgr B2. To take into account the radiative coupling between regions with different physical and/or excitation conditions, the RT has to be treated with non-local techniques. In order to constrain the physical parameters they have used the available spectroscopic ground based and ISO data of CO and other species, including the extended maser emission from H2O at 183.3 GHz. The models include all the rotational levels of H2O involving WATER IN SPACE 45 transitions with wavelengths longer than 40 μm. The level population is computed considering collisional excitation and radiative excitation by line and continuum photons emitted by the dust. This has been computed consistently assuming that the water molecules and the dust grains are coexistent. The geometry is a core+shell structure. Obviously all water transitions in the core are thermalized to the dust temperature due to the large opacity in the far-IR. The computed line profiles are a result of convolving with the angular resolution of the LWS detectors (∼80). To test the sensitivity of the models against the physical parameters, N(H2O) and Tk were computed for 1.8 × 1016,9× 1017 and 1.8 × 1018 cm−2, and 40, 100, 200, 300 and 500 K respectively, while densities were explored in the range from 5 × 103 to 5 −3 17 6.4×10 cm by steps of 2. The different non-local models for N(H2O) = 9×10 cm−2 are shown in Figure 10. The main problem to interpret the H2O absorption toward Sgr B2 is due to the large opacities of the far-IR water lines, typically ∼103 and even ∼104 for the o-H2O212–101 line at ∼179.5 μm. Under these conditions, many weak water lines have to be detected in order to constrain the physical conditions and the column density. In addition, as the radiative excitation dominates the population of the far- IR levels, lines with a weak dependence on the dust excitation should be searched. Taking into account the 14 water lines detected in the far-IR, the non-local results imply that these lines are not very sensitive to the temperature and that the only indication about the column density has to be searched in weak far-IR H2O lines or in the 183.3 GHz one. Even so, the far-IR absorption arises from the low density external layers of gas, while it is very likely that the 183.3 GHz line could have an important contribution from inner and denser regions. Another problem arises from the fact that models with low N(H2O) values and Tk 200 K, indicated that it is also difficult to distinguish between different H2 densities. Limits to Tk and n(H2) have to be searched in weak lines below 70 μm. The models for high water column densities (Figure 10) predict absorption lines in the LWS range at ∼56.3, ∼57.6, ∼58.7, ∼78.7, ∼99.5, ∼125.4 and ∼136.5 μm. At the spectral resolution and sensitivity of the LWS/FP, none of these lines can be 18 −2 detected (and they have not). This implies N(H2O) < 10 cm toward the warm envelope of Sgr B2. The ∼136.5 μm line, also predicted by the models with large N(H2O), is contaminated by the absorption produced by the C3 Q(8) ro-vibrational line, which is also predicted by the models for tri-atomic carbon absorption in Sgr B2 (Cernicharo et al., 2000). In addition, Cernicharo et al. (in press) have not detected any line of water vapor at wavelengths shorter than 60 μm. Some of these lines are predicted (even in emis- 16 sion) by the models with large N(H2O). Models considering N(H2O) = 1.8×10 cm−2 are compatible with ISO detections and non-detections. Only the ∼67.1 μm line is weaker than observed. Taking into account the difficulties related to the H2O models in the far-IR, Cernicharo et al. (in press) conclude that (4 ± 2) × 1016 cm−2 is the best H2O column density to fit the ISO observations of the gas surrounding Sgr B2. 46 J. CERNICHARO AND J. CROVISIER . Models for − 2 cm 17 et al. , in press). 10 × O column density is 9 . 0 ) are shown (from Cernicharo 2 − 3 cm 5 10 × to 1 . 6 3 10 × O rotational lines in Sgr B2(M). The total H densities (from 5 . 0 2 2 Results from selected non-local models of H Figure 10. different kinetic temperatures (from 40 to 500 K) and H WATER IN SPACE 47

The array of models in Figure 10 shows the difﬁculty in interpreting the far-IR data for H2O. However, such variety of models has to be computed in order to test the sensitivity of the water lines to the physical parameters and to collisional and/or radiative pumping. Future observations with the Herschel Space Observatory will provide data at high angular resolution, similar in fact to that achieved in ground based 183.3 GHz observations. However, the correct selection of the water lines sample to be observed will be critical to interpret the data and to derive the physical conditions of the gas and the water abundance.

2.1.3. Water Vapor Lines in Absorption: Dark Clouds toward Sgr A The abundance of water vapor in dark clouds has been analyzed through emission observations by SWAS (see e.g., Snell et al., 2000). The upper limits they have −8 obtained, x(H2O) ≤ 7 × 10 , correspond to very low H2O abundances and indicate that getting positive results when searching for water emission in very low kinetic temperature regions is difﬁcult. Although some weak emission from such cold clouds could be expected, one has to take into account that the surrounding gas, even if the density is not enough to pump the rotational levels of H2O, will absorb and re-emit over a large volume the photons arising from the densest regions in the cores of these clouds (see Cernicharo and Gu´elin,1987 for an analysis of this effect in the high dipole moment molecule HCO+). The result is that extremely weak emission could arise from dark clouds. Consequently, emission measurements are not well suited to derive the abundance of H2O. However, absorption measurements are better suited for this task in these objects when a continuum background source is available. Moneti et al. (2000) have used ISO to derive the water abundance in the dark clouds in front of the galactic center Sgr A∗. By observing CO and H2O gas and ice features with the SWS and LWS spectrometers (see Figure 11) they have been able to characterize the different components in the spectra: the warm gas surrounding Sgr A∗, and the cold gas and dust in the dark clouds in front of this cloud. However, water line observations are not enough to constrain the physical properties of the gas. Therefore, they complemented their study with an analysis of the pure rotational high-J lines of CO, also observed by ISO, to provide some additional key parameters for the interpretation of the water lines (see Figure 12). The use of the ro-vibrational and rotational lines of H2O, which have very different opacities, and the information provided by other molecular species has been extremely useful in deriving the water abundance in dark clouds. Plume et al. (2004), and references therein, have used the SWAS satellite to observe water absorption produced by the clouds placed in the line of sight to- −8 −6 ward W49A. They have derived H2O abundances of 8.2 × 10 to 1.5 × 10 , i.e., very similar values to those obtained in the dark clouds observed by Moneti et al. (2000). Where these lines are formed, i.e., in the innermost regions of the cold clouds or in their external layers (photochemistry) is not clear yet. Future observations with the Herschel Space Observatory in several lines of water and its 48 J. CERNICHARO AND J. CROVISIER

Figure 11. ISO/SWS spectra of the stretching mode of CO (a, b, c) and the bending mode of water vapor (d, e, f) in the direction of Sgr A∗. Panels a and d show the observations and a model for the continuum (including ices); panels b and e show the contribution of the ices and of the warm CO; panels c and f show the normalized spectrum and model spectrum (shifted for clarity), and panel g shows the details of the ortho- and para-H2O transitions (from Moneti et al., 2000). isotopic species (and other molecular lines, like CO lines to constrain models and the interpretation of water emission/absorption) will provide a more detailed view of the water abundance and chemical processes leading to its formation in cold dark clouds.

2.2. THE BENDING MODE OF WATER

Helmich and collaborators (Helmich et al., 1996b) have detected many ro- vibrational lines of water vapor due to absorption from the ground level up to the bending mode, ν2, opening the way to study the role of water vapor in star-forming regions. They detected more than 30 infrared absorption lines toward AFGL 2591 (see Figure 13) arising from gas at high kinetic temperature, TK 300 K. van Dishoeck and Helmich (1996) observed the bending mode of water toward massive young stars covering a large range of physical conditions. Hot water with excitation temperature >200 K was found toward GL 2136 and GL 4176 (in addition to GL 2591) with an abundance as large as 2–3 × 10−5 and comparable to the abundance of H2O in ices in the coldest parts of these clouds (see Figure 13). WATER IN SPACE 49

Figure 12. (a) J = 3–2 and J = 7–6 lines of CO observed at the CSO toward Sgr A∗. The downward arrows indicate the central position of the cold gas absorption features. The upward arrows indicate ∗ the velocity of the Sgr A molecular cloud. (b) LWS-FP spectra of the 212–101 and 221–110 lines of ∗ o-H2O toward Sgr A . The drawing at bottom left indicates the transition levels involved. The 101–212 line shows three unresolved components at v = 20, 5 and 55 km s−1 while in the other transition only the 55 km s−1 feature is present. The inset panel shows the LWS-grating spectrum; the CO – warm gas – and H2O lines are indicated. The upward arrows indicate the velocities of the cold absorbing gas and the downward arrow that of the Sgr A∗ molecular cloud (from Moneti et al., 2000).

Boonman and van Dishoeck (2003) have analyzed the available observations of the bending mode of H2O toward a significant number of molecular clouds and found water abundances in the range 5 × 10−6–6 × 10−5 with a clear tendency to have larger water abundances as the gas temperature increases, which suggests that grain-mantle evaporation is important. They have compared these results with chemical models showing that three different chemical processes, ice evaporation, high-T chemistry, and shocks, can reproduce the high gas-phase H2O abundances inferred. Of course, the results concern the inner regions of the cloud that surrounds these massive protostars. Although the observation of the bending mode of H2O provides an excellent tool to derive water abundances in high-mass star-forming regions, the main drawback is that these observations are limited to the directions with strong infrared emission at 6 μm, which are spatially limited to a few arcseconds. Nevertheless, the study of the bending mode of water could provide good insights on the role of infrared photons in pumping of water molecules to high energy levels. One of the nicest surprises provided by observations of the bending mode of H2O with ISO occurred in Orion: toward its central continuum source the R-branch of the band was observed in absorption, while the P-branch was found in emission. These results were correctly interpreted by González-Alfonso et al. (1998b) and González-Alfonso 50 J. CERNICHARO AND J. CROVISIER

Figure 13. Normalized spectra of NGC7538 IRS9, GL4176, GL2136 and GL2591. A model H2O 18 −2 −1 spectrum for a column density of 2 × 10 cm , Tex = 300 K and v = 5kms is shown for comparison (from van Dishoeck et al., 1996; see also Helmich et al., 1996b; see Boonman and van Dishoeck, 2003 for observations of a large number of sources in the ν2 mode of H2O). and Cernicharo (1999; see Figure 14; an overview of the H2O ν2 band is shown in the bottom-right panel of Figure 2). In fact, theoretical models made for evolved stars predicted such a behavior under some specific physical and geometrical conditions (González-Alfonsoand Cernicharo, 1999). These models were done prior to the observations and the excellent agreement found provided strong confidence in the physical parameters obtained in the interpretation of the ro-vibrational lines of water (high water abundances, the role of dust in the IR pumping, etc., see e.g. González-Alfonso et al., 2002). Finally, the bending mode of water could also be observed in absorption in dark clouds when a strong infrared source is located behind the cloud (see Section 2.1.3).

3. Water in Evolved Stars

3.1. O-RICH STARS

O-rich evolved stars, i.e., stars with C/O < 1, are the best water vapor factories in space. In the innermost layers of their circumstellar envelopes the physical WATER IN SPACE 51

Figure 14. Continuum normalized spectrum of Orion-BN/KL between 5.5 and 7.2 μm. The dotted line is the result of the “radiative excitation” H2O model described by González-Alfonso et al. (1998b). The solid line is the result of adding to the line fluxes of that model the fluxes of a collisionally excited model. The strongest lines are labeled (ortho: lower labels; para: upper labels), the first three quantum numbers correspond to the level in the ν2 = 1 bending state and the second set of three quantum numbers belongs to the level in the ground ν = 0 state (from González-Alfonso et al., 1998b). conditions (high density and temperature) allow a chemistry under thermodynamical equilibrium. As CO is the most stable molecule in space, most carbon atoms are locked in CO leaving oxygen atoms available to build other gas phase species such as CO2 and H2O. González-Alfonsoand Cernicharo (1999) have modeled the water spectrum of evolved stars. The pure rotational lines and the bending mode were analyzed. Figure 15 (left panel) shows these models for two different mass loss rate regimes. 52 J. CERNICHARO AND J. CROVISIER

Figure 15. Left panel: Predicted ISO LWS spectrum of water vapor and CO for two models with −5 −1 −1 mass loss rates and expansion velocities of 1.3 × 10 M yr Vexp = 20 km s (upper subpanel) −7 −1 −1 and 2 × 10 M yr Vexp = 9kms (lower subpanel). The water vapor and CO abundances are −4 −4 14 2.7 × 10 and 6 × 10 respectively. The internal radius of the H2O envelope is 10 cm, while the stellar radius is 3.7 × 1013 cm. The source is placed at 500 pc and the models take into account the excitation of molecules to the bending mode. The lines of CO are indicated, all the remaining lines belong to H2O. Right panel: Predicted ISO SWS spectrum of water vapor at 6 μm in the two models quoted above (panels (a) and (b)). Model c is similar to b except that the inner radius of the H2O shell 14 9 coincides with the stellar radius. The region between r = R∗ and r = 10 cm has n(H2) = 2 × 10 −3 −1 cm and Vtur = 5kms (from Gonz´alez-Alfonsoand Cernicharo, 1999).

In the high end of mass loss rates, the far-infrared spectrum is fully dominated by the pure rotational lines of H2O (the CO lines are also present but weaker than those of water). The low-mass loss rate model in the same figure shows that CO and water lines have similar intensities but the densest water spectrum dominates the far-infrared emission. The models shown in Figure 15 have been convolved to the spectral resolution of the LWS grating spectrometer and indicate that water will dominate all cooling processes in the envelope (at least in the regions where water is not yet photodissociated). The predictions for the bending mode are rather surprising (Figure 15, right panel). Depending on the mass loss rates, the size of the envelope, and the gas density (radiation versus collisions), the bending mode will present a peculiar behavior WATER IN SPACE 53 with the R-branch in absorption and the P-branch in emission for low-mass loss rates (panel b), or with the P-branch in absorption/emission for high-mass loss rates (panel a). Changing the size of the envelope has important effects on the shape of the bending band as shown in panel c, where some lines appear in absorption while other lines are in emission in each branch (see figure caption). This kind of behavior was nicely verified in Orion (see above; González-Alfonso et al., 1998a; van Dishoeck et al., 1998) and shows that the observation of the whole bending band provides important constraints when modeling the water excitation. The first ISO spectrum showing the pure rotational lines of H2O was obtained by Barlow et al. (1996) in the direction of W Hya (seeFigure 16). The spectrum is dominated by a forest of water lines confirming that H2O molecules are the dominant coolants of the gas in the circumstellar envelope. They have find a water abundance of 8 × 10−4 for r ≤ 4.5 × 1014 cm and 3 × 10−4 at large radii. The star −7 −1 mass loss rate is 6 × 10 M yr . Neufeld et al. (1996) also observed this object −5 with the SWS/FP spectrometer. They obtained a mass loss rate of (0.5–3)×10 M yr−1, i.e., two orders of magnitude larger than the value obtained by Barlow et al. (1996) or that obtained from dynamical considerations. This difference is probably arising from the different physical parameters assumed in the interpretation and to the lack of precise collision rates for H2O–H2 at the temperatures prevailing in these warm objects. Justtanont et al. (2004) have analyzed all the available data

Figure 16. Continuum-substracted LWS spectrum of W Hya together with the best fit model for water emission. The dominant transitions are labeled (from Barlow et al., 1996). 54 J. CERNICHARO AND J. CROVISIER for W Hya and conclude that the only way to put in agreement the mass loss rates calculated from different molecular tracers is to increase the water abundance. The −8 −1 best value they obtain from CO lines is (3.5–8) ×10 M yr which is a factor 10 below that of Barlow et al. (1996) and three orders of magnitude below the mass loss rate derived by Neufeld et al. (1996). The main drawback of water vapor observations in O-rich stars is the large line opacities through the whole envelope which make extremely difficult the selection of rotational lines sensitive to the physical conditions. Many other O-rich evolved stars have been observed with ISO. R Cas has been modeled by Truong-Bach et al. (1999) getting a water abundance of 1.1 × 10−5. More detailed models have been developed to interpret the high-mass loss rate stars VY CMa (Neufeld et al., 1999; they also report the detection of rotational lines inside the ν2 level) and NML Cyg (Zubko et al., 2004). Most AGB stars are bright in the mid infrared which favor the observation of the bending mode of water in these objects. However, the lack of spectral resolution with ISO makes difficult the interpretation of the data because low excitation lines form throughout the whole envelope while high excitation lines, showing probably P-Cygni profiles, arise from the innermost layers of the envelope. Justtanont et al. (1996) observed the ν2 H2O band toward NML Cyg (see Figure 17). Although the column density of water they derive is high the excitation temperature is rather low implying that most lines in the SWS spectrum are formed in the external envelope. They also observed the ν1 and ν3 bands of H2Oat2.7μm. OH/IR objects are O-rich stars with a high-mass loss rate. They are characterized by the presence of strong OH masers and strong emission in the infrared. ISO observed a large number of these objects with the LWS and SWS spectrometers.

Figure 17. Continuum-divided spectrum of the ν2 band of H2O in NML Cyg (bottom). The modeled 19 −2 spectrum with Tex = 100 K and N(H2O) = 10 cm is shown at the top (from Justtanont et al., 1996). WATER IN SPACE 55

Sylvester et al. (1999) analyzed these data and found a wealth of broad spectral features due to crystalline silicates and crystalline water ice in emission and absorption. Crystalline ice at 43 and 60 μm was detected ﬁrst toward these objects by Omont et al. (1990) using the Kuiper Airborne Observatory (KAO). The full frequency coverage of the SWS and LWS spectrometer permits a more detailed study of these bands. Sylvester et al. (1999) concluded that there is a certain onset value for the mass loss rate above which these features appear in the spectrum. Water vapor is also present in the spectra of these objects (see Figure 18). Molster et al. (2001, and references therein) have analyzed the ISO spectrum of NGC6302, one of the ﬁrst objects observed with ISO showing the crystalline water ice bands. These bands have also been detected toward the C-rich planetary nebula CPD-56◦8032 by Cohen et al. (1999).

Figure 18. Continuum-divided spectra of several OH/IR stars in the 2–7 μm region. The central wavelengths of the absorption features due to H2O (gaseous and ice phases), and of gaseous CO, CO2 and SiO are indicated (from Sylvester et al., 1999). 56 J. CERNICHARO AND J. CROVISIER

3.2. C-RICH STARS

In carbon-rich stars C/O > 1 and all oxygen atoms will be locked into carbon monoxide. Hence, oxygen will be unavailable to build other molecules. Nevertheless, Melnick et al. (2001) have reported the detection of H2O with SWAS in the direction of IRC+10216, the prototype of carbon-rich AGB stars. The far-IR spectrum of this object as observed by ISO was reported by Cernicharo et al. (1996b) and no clear indication of H2O could be found as the spectrum is dominated by a forest of emission lines from CO and HCN. Melnick et al. (2001) have suggested that the observed water vapor is arising from the evaporation of cometary bodies orbiting IRC+10216. More sensitive observations with Herschel of several water lines having different excitation conditions could provide a clearer picture about the origin of H2O in this star. The situation is different for post-AGB carbon-rich stars. Herpin and Cernicharo (2000) have presented ISO LWS observations of the proto-planetary nebula CRL 618, a carbon rich object in a very fast evolutionary phase prior to reaching the planetary nebula stage. The far-infrared spectrum is essentially dominated by CO lines (see Figure 19). All the other species have much lower intensities contrary to that found in the AGB star IRC+10216 where HCN lines (from the ground and vibrationally excited states) are as strong as those of CO (see Cernicharo et al., 1996b). The main difference between both objects has to be found in the physical structure of their CSEs. That of CRL 618 has a central hole in molecular species filled by a bright H II region. In addition to the lines of CO, 13CO, HCN and HNC, Herpin and Cernicharo (2000) reported the detection of H2O and OH emission together with the fine structure lines of [O I] at 63 and 145 μm. The abundances of these species relative to CO are 4 × 10−2,8× 10−4 and 4.5 in the regions where they are produced. The authors have suggested that O-bearing species other than CO are produced in the innermost region of the circumstellar envelope. UV photons from the central star photodissociate most of the molecular species produced in the AGB phase and allow a chemistry dominated by standard ion-neutral and neutral-radical reactions. Not only these reactions allow the formation of O-bearing species but also modify the abundances of C-rich molecules like HCN and HNC for which the authors have found an abundance ratio of 1, much lower than what is typical in AGB stars. Herpin et al. (2002) have made a comparative study of three carbon-rich post-AGB objects, CRL618, CRL2688 and NGC7027. In the early stages of the AGB to PN evolution (represented by CRL2688) the far-infrared spectrum is dominated by CO lines. In the intermediate stage, e.g., CRL618, very fast outflows are present which, together with the strong UV field from the central star, dissociate CO. The released atomic oxygen is seen via its atomic lines and allows the formation of new O-bearing species, such as H2O, H2CO and OH. At the planetary nebula stage, e.g., NGC7027, a large fraction of the old CO AGB material has been reprocessed. The far-infrared spectrum is dominated by atomic and ionic lines. New species, WATER IN SPACE 57

13 Figure 19. ISO/LWS continuum-substracted spectra of CRL618. The lines of CO, CO, HCN, H2O and OH are indicated by arrows, while those of HNC are indicated by vertical lines (from J = 22–11 at 150.627 μmtoJ = 17–16 at 194.759 μm). The ﬁgure is from Herpin and Cernicharo (2000) and the result of their model is shown as a solid line.

such as CH+ (see Cernicharo et al., 1997d), appear. The water vapor formed during the protoplanetary nebula stage has been photodissociated. The chemical evolution during the protoplanetary nebula stage has been modeled, using only neutral and neutral-radical reactions, by Cernicharo (2004). The radicals are the products of the photodissociation of molecules formed during the AGB phase of the star. The author has shown that H2O, H2CO (a molecule detected in CRL618 by Cernicharo et al., 1989), and CO2 are easily formed in the warm and dense gas of the photodissociation region surrounding the hot central star. These models also predict very large abundances for other carbon chains and carbon clusters as those found by Cernicharo et al. (2001a,b). Hence, photochemistry plays an important role in the chemical evolution of the envelopes of post-AGB stars making possible the presence of O-bearing species in a medium where the physical conditions and the atomic abundances would result in very little amounts of these species. 58 J. CERNICHARO AND J. CROVISIER

4. Water in Galaxies

Due to the small size of the ISO telescope the observation of water in gas phase toward extragalactic objects is strongly limited in sensitivity. Nevertheless, observation of water ice in 18 galaxies, from a sample of 103 galaxies observed with ISO, has been reported by Spoon et al. (2002). They have found that water ice is present in most ULIRGs, whereas it is weak or absent in the large majority of Seyferts and starburst galaxies. Only a few objects show gas phase molecular emission or absorption: Arp220, NGC1064 and NGC253. Goicoechea, Martin-Pintado and Cernicharo (2005), have analyzed the OH emission/absorption in NGC1064 and NGC253, while Arp220 has been studied in detail by Gonz´alez-Alfonso et al. (2004). Figure 20 shows the far-IR spectrum of Arp220 as observed by ISO. The spectrum looks very similar to that

Figure 20. ISO LWS spectrum of Arp220 where the most prominent line features are identiﬁed. The gray line shows the best model for the lines (from Gonz´alez-Alfonso et al., 2004). WATER IN SPACE 59

Figure 21. Continuum-normalized spectra of SgrB2 (M) and Arp220. The main carriers of some line features are indicated (from González-Alfonso et al., 2004). of Sgr B2 (see Figure 21). The modeling of the Arp220 far-infrared spectrum (see below), and of the OH lines in NGC1064 and NGC253, was done by comparison of their spectra with that of some well-studied Galactic sources, which provides important clues about the physical conditions in the regions where the lines are formed. In Arp220 González-Alfonso et al. (2004) have observed absorption in the lines of OH, H2O, NH, NH3 and CH in addition to [O I]at63μm and emission in the [C II] line at 158 μm. They have modeled the continuum and the emission/absorption of all observed features by using non-local radiative transfer codes. The continuum from 25 to 1300 μm corresponds to a nuclear region, optically thick in the far-infrared, with a size of 0.4 and a dust temperature of 106 K surrounded by an extended region (2) heated mainly through absorption of the nuclear infrared radiation. The OH column densities are high toward the nuclear region (2–6 × 1017 cm−2) and in the extended region ( 2 × 1017 cm−2) while the water column density is high toward the nucleus (2–10 × 1017 cm−2) and lower in the extended region. The column densities in a halo that accounts for the absorption in the lowest lying lines are similar to what is found in the diffuse clouds toward the star-forming regions in the Sgr B2 molecular cloud in the neighborhood of the Galactic Center. One of the most surprising results from González-Alfonso et al. (2004) is the large column 16 −2 density needed to explain the observations of NH and NH3 (1.5 × 10 cm and 16 −2 3 × 10 cm , respectively) while NH2, a molecule detected in Sgr B2 (M) by Goicoechea and Cernicharo (2001b), has a column density below 2 × 1015 cm−2. 60 J. CERNICHARO AND J. CROVISIER

The excellent ISO spectrum of Arp220 makes it a reference template for understanding the dusty interstellar medium of ULIRGs. The model of Gonz´alez-Alfonso et al. (2004) shows the important role of the dust photons in pumping the high- excitation lines of OH and H2O, i.e., a similar situation to than that found in Sgr B2(M) (Goicoechea and Cernicharo, 2002). The modeling of water vapor and OH in Arp220 indicates the prominent role of PDR molecular chemistry in the extended region of Arp220, while chemistry in hot cores and shocks in the nucleus is also contributing to the observed abundance of H2O. Herschel, with its three instruments (HIFI, PACS and SPIRE), and its large telescope (3.5 m), will permit to observe the molecular content of more distant objects. The full frequency coverage provided by PACS and SPIRE will allow to make more detailed and sensitive studies of H2O and OH, while HIFI will provide high spectral resolution of the line proﬁles of these molecules. It is worth noting that the H2O models for water in Arp 220, and for sources in our galaxy, have required the simultaneous observation and modeling of other species, in particular OH and CO. Moreover, the use of galactic sources as templates as initial input for the interpretation of molecular extragalactic observations could be a mandatory step. Goicoechea et al. (2005) have used Orion and Sgr B2 as templates to interpret the far-infrared spectra of NGC253 and NGC1064. Thanks to ISO, this information will be available when observing extragalactic sources with Herschel. It will be, without any doubt, the only way to interpret correctly the H2O lines, which will provide important information on the chemistry and physical conditions of the gas.

5. Water in the Solar System

Water is ubiquitous in the Solar System. It is present – as ice and liquid on the surface of the Earth; – as ice and – presumably, in the past – as liquid on Mars; – in various amounts in the atmospheres of most planets; – as the main constituent of cometary ices; – in primitive meteorites such as carbonaceous chondrites; – as ice on numerous planetary satellites and – presumably – on trans-Neptunian objects (TNO); – and even in the Sun’s atmosphere. Water is also a requisite for the apparition and evolution of life. It is thus important to know the cycle of water in the Solar System. Problems that are still open include: – to which extent cometary (and TNO) ices are coming from unprocessed interstellar ice; – what is the origin of water in the giant planets’ atmospheres; WATER IN SPACE 61

– to which extent did the infall of comets, asteroids and interstellar dust particles (IDPs) contributed to Earth (and Mars) water. The contribution of ISO to cometary and planetary studies has also been reviewed by Crovisier (2000), Encrenaz (2000), Fouchet et al. (2004) and M¨uller et al. (2004).

5.1. COMETS

Water is the main constituent of cometary ices. Its sublimation governs cometary activity close to the Sun (r ≤ 4 AU). However, the direct observation of cometary water is very difficult from the ground, where only lines from infrared hot bands can be observed. The water ro-vibrational lines are emitted by fluorescence excited by solar radiation. Their measurement allows us to determine the water production rates in comets. Water was best observed by ISO in the exceptionally bright comet C/1995 O1 (Hale-Bopp), as part of a target-of-opportunity program, in September–October 1996. The comet was then at r = 2.8 AU from the Sun and its water production rate 29 −1 was 3.3×10 molecules s . The ν1, ν3 and hot bands around 2.7 μm were resolved in rotation with the SWS (Figure 22; Crovisier et al., 1997a, 1999b). The ν2 band at 6.5 μm was also detected with the SWS, for the first time in a comet (Crovisier et al., 1997b). Several rotational lines (212–101 and 303–212 around 180 μm and weaker

Figure 22. The SWS spectrum of water in the 2.7 μm region observed in comet C/1995 O1 (Hale- Bopp) on 27 September and 6 October at 2.8 AU from the Sun (top). Line assignations are indicated. The synthetic ﬂuorescence spectrum which corresponds to the best ﬁt to the data is shown at the 29 −1 bottom. It corresponds to Q[H2O] = 3.6 × 10 molecules s , Trot = 28.5 K and Tspin = 28 K (adapted from Crovisier et al., 1997a). 62 J. CERNICHARO AND J. CROVISIER

Figure 23. Several rotational lines of water observed in comet C/1995 O1 (Hale-Bopp) with the LWS. The C II line is not cometary; it comes from the background (from Crovisier et al., 1999a). ones) were observed with the LWS (Figure 23). This was also the first detection of rotational lines of water in a comet, before the 110–101 line at 557 GHz line could be observed by heterodyne techniques with the SWAS and Odin satellites. In comet Hale-Bopp, water was only detected when the comet was at r = 2.9 and 3.9 AU. At farther distances (4.6 AU pre-perihelion and 4.9 AU post-perihelion), only CO and CO2 (which are much more volatile species) were detected, with PHT-S. This shows the evolution between CO-dominated (at large r’s) and H2O- dominated (at small r’s) regimes of sublimation of cometary ices. ISO also observed H2O in two short-period, Jupiter-family comets. In 103P/Hartley 2, which was observed close to its perihelion at r = 1 AU, the water production rate was 1.2 × 1028 molecules s−1 (Colangeli et al., 1999; Crovisier et al., 1999a). In 22P/Kopff, observed in a less productive state after perihelion at r = 1.9 AU, the water production rate was only 3.7 × 1027 molecules s−1 and the water lines at 2.7 μm were just detected by the SWS (Crovisier et al., 1999a). The SWS spectra obtained in the 2.6 μm region on comets Hale-Bopp and Hartley 2 allowed us to investigate the rotational distribution of water, and therefore its excitation conditions. Water excitation is governed by radiative excitation of the fundamental bands of vibration by the Sun IR radiation, collisions in the inner coma, and radiative trapping effects (Bockelée-Morvan, 1987). On the other hand, WATER IN SPACE 63 the kinetic temperature of the coma is governed by the balance between radiative cooling (through emission of water rotational lines) and photolytic heating (principally through water photodissociation). In the field of view of the SWS, cometary water is partly relaxed to its lowest rotational levels. As predicted by modeling, the observed rotational temperatures of water are low: 28 K for Hale-Bopp, 20 K for Hartley 2, < 11 K for Kopff (Crovisier et al., 1997a, 1999a,b). Water exists in two nuclear-spin states according to the spins of its hydrogen atoms: ortho (I = 1) and para (I = 0). Conversions between the two states are forbidden, so that the ortho-to-para population ratio is linked to the temperature (spin temperature Tspin) of the last re-equilibration of water. The SWS spectra allowed us to determine Tspin of water in comets Hale-Bopp and Hartley 2. Tspin = 28 and 35 K were derived respectively for the two comets (Crovisier et al., 1997a, 1999a,b). A similar value (29 K) was observed in the past for comet 1P/Halley (Mumma et al., 1993). The spin temperature of cometary ammonia (which is indirectly determined from the visible spectrum of NH2; Kawakita et al., 2004) has also similar values (25–32 K). These Tspin, all observed in a small range for comets of different orbits and different dynamical history, are still to be understood. It has been argued that Tspin could reflect the temperature of the grains where the molecules formed, but it seems hard to believe that the ortho-to-para ratio could be preserved over cosmological times without re-equilibration. Tspin could reflect the temperature of the inner nucleus ices. But then, one would expect different temperatures for comets with different orbits and history. It would be interesting to investigate by laboratory experiments if Tspin is preserved during the sublimation process, or if some fractionation occurs. It would also be interesting to compare the cometary water ortho-to-para ratio with those observed for interstellar water.

5.2. PLANETS AND SATELLITES

5.2.1. Mars Water on Mars is now studied by a wealth of means. ISO also contributed to the observation of H2O in Mars’ atmosphere. Ro-vibrational absorption lines were observed in the 2.6 and 6.5 μm regions with the SWS, while many rotational lines were observed all over the 20–200 region with the SWS and LWS (Figure 24; Encrenaz et al., 1999; Burgdorf et al., 2000; Lellouch et al., 2000). From these observations, the mean water concentration at the ground was 4 × 10−4 at that time (31 July 1997) and saturation was reached at an altitude of 10 km. This corresponds to 15 μm of precipitable water.

5.2.2. Giant Planets and Titan An important result of ISO was the discovery of water in the stratospheres of the four giant planets and Titan. These detections are based upon water rotational lines observed in emission with the SWS and LWS, using the Fabry-P´erot(Figure 25; 64 J. CERNICHARO AND J. CROVISIER

Figure 24. Lines of water in the 2.6 μm region of the spectrum of Mars, observed with the SWS grating. A modeled spectrum is superimposed (from Lellouch et al., 2000).

Figure 25. Examples of detections of H2O lines in all four giant planets and Titan, with the SWS Fabry-P´erot(from Feuchtgruber et al., 1999).

Feuchtgruber et al., 1997, 1999; Coustenis et al., 1998; Lellouch et al., 2002; Encrenaz, 2003). The disk-averaged water column densities are observed in the range ≈1014 cm−2 (Uranus, Neptune, Titan) to ≈2 × 1015 cm−2 (Jupiter, Saturn). Because water condenses at the tropopause, its presence cannot be explained by an internal source of these objects. An external inﬂux has to be invoked. This is the same for CO2, observed in the stratosphere of Jupiter, Saturn and Uranus.

5.3. THE CYCLE OF WATER IN THE SOLAR SYSTEM

Observations at infrared (ground-based and from ISO) and radio wavelengths point to the similarity between interstellar and cometary ices (Ehrenfreund et al., 1997; WATER IN SPACE 65

Bockel´ee-Morvan et al., 2000). This does not prove that comets accreted unprocessed interstellar ices, but rather suggests that similar chemical processes were at work in the primitive Solar Nebula and in interstellar clouds. Impacts are an important phenomenon for the redistribution of water in the Solar System. Shock chemistry following the collisions of large bodies (planetesimals, comets, asteroids) will reshufﬂe most of the original molecular species. The oxygen compounds of the impactor could then lead to the formation of H2O (and CO2) in the post-impact chemistry. On the other hand, for infall of IDPs, the original molecules could be preserved. Icy IDPs could thus be an important source of water. Several studies have considered asteroids, comets and meteoroids as plausible sources of the water accreted by the Earth (e.g., Morbidelli et al., 2000). Martian water could have a similar origin. The collision of comet Shoemaker-Levy 9 and Jupiter is a well-documented case. Unfortunately, there is still ambiguity on the nature of the impactor (comet or asteroid?); its chemical nature could not be precisely characterized from pre-impact observations (Crovisier, 1996). There is little doubt, however, that the chemical species observed after the impact result from shock chemistry, rather than coming from the impactor (Lellouch, 1996). In this context, three different sources have been invoked to explain the presence of oxygen species (water and carbon dioxide) discovered by ISO in giant planets (e.g., Encrenaz, 2003; Lellouch et al., 2002; Fouchet et al., 2004):

– infall of IDPs; – sputtering from icy rings and satellites; this is the preferred source for Saturn and Titan; – impact of small bodies (comets and asteroids); this is certainly the case for Jupiter, for which the fall of comet Shoemaker-Levy 9 in 1994 provided the main source of oxygen.

6. Conclusions

Water has been found everywhere in space by ISO. The interpretation of the data is not straightforward but, nevertheless, we had for the first time the opportunity to study the role of water in the chemistry of all kind of objects in space, from solar system bodies to star-forming regions, evolved stars and galaxies. Despite the limited spectral resolution of most water observations with ISO, the possibility to observe its full far and mid infrared spectrum has opened great possibilities in the interpretation of the data. Complex radiative transfer models able to treat the very large opacities of water have been developed, the chemistry of water on dust grains and gas phase has received an extraordinary and unique input from ISO, the presence of water in the upper layers of giant planets and their moons has told us about the chemistry of these objects, the detection of water in extragalactic sources 66 J. CERNICHARO AND J. CROVISIER could be used as a tracer of the physical and chemical processes prevailing in the nuclei of galaxies, etc. Future missions like Herschel will benefit from the huge information ISO has provided on water vapor. The selection of water lines according to criteria of sensitivity to the physical parameters of the clouds, their interpretation and the prediction of the full spectrum of water in several clouds will be a mandatory step to prepare the Herschel mission. However, little new information will be obtained if collisional rates adapted to the temperatures of the clouds in the ISM and CSM are not available. Several aspects have to be improved for future water observations: we urgently need collisional rates for H2O–H2 at temperatures going from 10 to 2000 K (dark clouds to evolved stars). For AGB stars collisional rates between the ground and the bending mode are also necessary. These data have to be provided by physicists and chemists doing ab initio calculations or laboratory experiments. A very close collaboration between astronomers and the world of quantum chemistry is needed in order to promote these calculations and to be sure that we will get the maximum scientific output from Herschel when observing water in space.

Acknowledgements

We would like to thank J.R. Pardo, F. Daniel and J.R. Goicoechea for useful comments and suggestions and a critical reading of the manuscript. They have kindly provided several ﬁgures and calculations for this paper. J. Cernicharo thanks the Spanish Ministry of National Education (MEC) for funding support under grants AYA2003-2785, AYA2002-12125E, ESP2002-11650, AYA2002-10113-E, and ESP 2002-01627. Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, the Netherlands and the United Kingdom) and with the participation of ISAS and NASA.

References

Barlow, M., et al.: 1996, Astron. Astrophys. 315, L241. Beckwith, S., Persson, S. E., Neugebauer, G., and Becklin, E. E.: 1978, Astrophys. J. 223, 484. Benedettini, M., et al.: 2000, Astron. Astrophys. 359, 148. Blake, G. A., Sutton, E. C., Masson, C. R., and Phillips, T. G.: 1987, Astrophys. J. 315, 621. Bockel´ee-Morvan, D.: 1987, Astron. Astrophys. 181, 169. Bockel´ee-Morvan, D., Lis, D. C., Wink, J. E., et al.: 2000, Astron. Astrophys. 353, 1101. Boonman, A. M. S. and van Dishoeck, E.: 2003, Astron. Astrophys. 403, 1003. Burgdorf, M. J., Encrenaz, T., Lellouch, E., et al.: 2000, Icarus 145, 79. Ceccarelli, C., et al.: 1999a, Astron. Astrophys. 331, 372. Ceccarelli, C., et al.: 1999b, Astron. Astrophys. 342, L21. Ceccarelli, C., et al.: 2002, Astron. Astrophys. 383, 603. Cernicharo, J.: 1985, ATM a Program to Compute Atmospheric Transmision Between 0–1000 GHz, IRAM Internal Report. WATER IN SPACE 67

Cernicharo, J.: 1988, Ph.D. Thesis, Universitéde Paris VII. Cernicharo, J. and Guélin,M.: 1987, Astron. Astrophys. 176, 299. Cernicharo, J., Guélin,M., Martin-Pintado, J., and Mauersberger, R.: 1989, Astron. Astrophys. 222, L1. Cernicharo, J., Thum, C., Hein, et al.: 1990, Astron. Astrophys. 231, L15. Cernicharo, J., González-Alfonso,E., Alcolea, J., et al.: 1994, Astrophys. J. 432, L59. Cernicharo, J., González-Alfonso,E., and Bachiller, R.: 1996a, Astron. Astrophys. 305, L5. Cernicharo, J., et al.: 1996b, Astron. Astrophys. 315, L201. Cernicharo, J., Lim, T., Cox, P., et al.: 1997a, Astron. Astrophys. 323, L25. Cernicharo, J.: 1997, The Far Infrared and Submillimeter Universe, 15–17 April 1997, Grenoble, France, ESA SP-401, p. 91. Cernicharo, J., González-Alfonso,E., and Lefloch, B.: 1997b, First Workshop on Analytical Spec- troscopy, 6–8 October 1997, ESA SP-419, p. 23. Cernicharo, J., X-Liu, González-Alfonso,E., et al.: 1997c, Astrophys. J. 483, L65. Cernicharo, J.: 1998, Astrophys. Space Sci. 263, 175. Cernicharo, J., González-Alfonso,E., et al.: 1999a, The Universe as Seen by ISO, Paris, 1999. Cernicharo, J., Pardo, E., Serabyn, et al.: 1999b, Astrophys. J. 520, L131. Cernicharo, J., Goicoechea, J. R., and Caux, E.: 2000, Astrophys. J. 534, L199. Cernicharo, J., et al.: 2001a, Astrophys. J. 546, L123. Cernicharo, J., et al.: 2001b, Astrophys. J. 546, L127. Cernicharo, J.: 2004, Astrophys. J. 608, L41. Cernicharo, J., Goicoechea, J. R., Daniel, F., et al.: in press, Astrophys. J. Cheung, A. C., Rank, D. M., Townes, C. H., et al.: 1969, Nature 221, 626. Clegg, P. E., et al.: 1996, Astron. Astrophys. 315, L38. Cohen, M., Barlow, M., Sylvester, R. J., et al.: 1999, Astrophys. J. 513, L135. Colangeli, L., Epifani, E., Brucato, J. R., et al.: 1999, Astron. Astrophys. 343, L87. Comito, C., Schilke, P., Gerin, M., et al.: 2003, Astron. Astrophys. 402, 635. Coustenis, A., Salama, A., Lellouch, E., et al.: 1998, Astron. Astrophys. 336, L85. Crovisier, J.: 1996, in: Noll, K. S., Weaver, H. A., and Feldman, P. D. (eds.), The Collision of Comet Shoemaker-Levy 9 and Jupiter, Cambridge University Press, Cambridge, p. 31. Crovisier, J., Leech, K., Bockelée-Morvan, D., et al.: 1997a, Science 275, 1904. Crovisier, J., Leech, K., Bockelée-Morvan, D., et al.: 1997b, in: Heras, A. M., Leech, K., Trams, N. R., and Perry, M. (eds.), First ISO Workshop on Analytical Spectroscopy, ESA SP-419, p. 137. Crovisier, J., Encrenaz, T., Lellouch, E., et al.: 1999a, in: Cox, P., and Kessler, M. F. (eds.), The Universe as Seen by ISO, ESA SP-427, p. 161. Crovisier, J., Leech, K., Bockelée-Morvan, D., et al.: 1999b, in: Cox, P., and Kessler, M. F. (eds.), The Universe as Seen by ISO, ESA SP-427, p. 137. Crovisier, J.: 2000, in: Mihn, Y. C., and van Dishoeck, E. F. (eds.), Astrochemistry: From Molecular Clouds to Planetary Systems (IAU Symposium 197), ASP, p. 461. Daniel, F., Goicoechea, J. R., Cernicharo, J., et al.: in press, Astrophys. J. de Graauw, T., et al.: 1996, Astron. Astrophys. 315, L49. Encrenaz, T., Feuchtgruber, H., Burgdorf, M., et al.: 1999, in: Cox, P., and Kessler, M. F. (eds.), The Universe as Seen by ISO, ESA SP-427, p. 173. Encrenaz, T.: 2000, in: Casoli, F., Lequeux, J., and David, F. (eds.), Infrared Space Astronomy, Today and Tomorrow (Ecole´ de Physique des Houches). EDP Sciences, Paris, p. 89. Encrenaz, T.: 2003, Planet. Space Sci. 51, 89. Ehrenfreund, P., d’Hendecourt, L., Dartois, E., et al.: 1997, Icarus 130,1. Feuchtgruber, H., Lellouch, E., de Graauw, T., et al.: 1997, Nature 389, 159. Feuchtgruber, H., Lellouch, E., Encrenaz, T., et al.: 1999, in: Cox, P., and Kessler, M. F. (eds.), The Universe as Seen by ISO, ESA SP-427, p. 133. Fouchet, T., Bézard,B., and Encrenaz, T.: 2004, The Planets Observed by ISO, this volume. 68 J. CERNICHARO AND J. CROVISIER

Gensheimer, P. D., Mauersberger, R., and Wilson, T. L.: 1996, Astron. Astrophys. 314, 281. Genzel, R. and Stutzki, J.: 1989, Annu. Rev. Astron. Astrophys. 27, 41. Genzel, R., Reid, M. J., Moran, J. M., and Downes, D.: 1981, Astrophys. J. 244, 884. Gibb, E. L., Whittet, D. C. B., Bougert, A. C. A., and Tielens, A. G. G. M.: 2004, ApJSS 151, 35. Goicoechea, J. R. and Cernicharo, J.: 2001a, Astrophys. J. 554, L213. Goicoechea, J. R. and Cernicharo, J.: 2001b, in: Pilbratt, G. L., Cernicharo, J., Heras, A. M., Prusti, T., and Harris, R. (eds.), The Promise of FIRST, ESA SP-460; Noordwijk: ESA/ESTEC, p. 413. Goicoechea, J. R. and Cernicharo, J.: 2002, Astrophys. J. 576, L77. Goicoechea, J. R., Rodriguez-Fernandez, N. J., and Cernicharo, J.: 2004, Astrophys. J. 600, 214. Goicoechea, J. R., Martin-Pintado, J., and Cernicharo, J.: 2005, Astrophys. J. 619, 291. González-Alfonso,E. and Cernicharo, J.: 1993, Astron. Astrophys. 279, 506. González-Alfonso,E., Cernicharo, J., Bachiller, R., and Fuente, A.: 1995, Astron. Astrophys. 293, L9. González-Alfonso,E. and Cernicharo, J.: 1997, Astron. Astrophys. 322, 938. González-Alfonso,E., Cernicharo, J., Alcolea, J., and Orlandi, M. A.: 1998a, Astron. Astrophys. 334, 1016. González-Alfonso,E., Cernicharo, J., van Dishoeck, E. F., et al.: 1998b, Astrophys. J. 502, L169. González-Alfonso,E. and Cernicharo, J.: 1999, Astrophys. J. 525, 845. González-Alfonso,E., Wright, C. M., Cernicharo, J., et al.: 2002, Astron. Astrophys. 386, 1074. González-Alfonso,E., Smith, H. A., Fischer, J., and Cernicharo, J.: 2004, Astrophys. J. 613, 247. Harwit, M. Neufeld, D. A., Melnick, G. J., and Kaufman, M. J.: 1998, Astrophys. J. 497, L105. Helmich, F. P., van Dishoeck, E. F., and Jansen, D. J.: 1996a, Astron. Astrophys. 313, 657. Helmich, F. P., van Dishoeck, E. F., Black, J. H., et al.: 1996b, Astron. Astrophys. 315, L173. Herpin, F. and Cernicharo, J.: 2000, Astrophys. J. 530, L129. Herpin, F., Goicoechea, J. R., Pardo, J. R., and Cernicharo, J.: 2002, Astrophys. J. 577, 961. Hjalmarson, A., et al.: 2003, Astron. Astrophys. 402, L39. Hüttemeister, S., Wilson, T. L., Mauersberger, R., Lemme, C., Dahmen, G. and Henkel, C.: 1995, Astron. Astrophys. 294, 667. Jacq, T., Jewell, P. R., Henkel, C., Walmsley, C. M., and Baudry, A.: 1988, Astron. Astrophys. 199, L5. Jacq, T., Walmsley, C. M., Henkel, C., et al.: 1990, Astron. Astrophys. 228, 447. Justtanont, K., de Jong, T., Helmich, F. P., et al.: 1996, Astron. Astrophys. 315, L217. Justtanont, K., de Jong, T., Tielens, A. G. G. M., et al.: 2004, Astron. Astrophys. 417, 625. Kawakita, H., Watanabe, J., Furusho, R., et al.: 2004, Astrophys. J. 601, 1152. Kessler, M. F., et al.: 1996, Astron. Astrophys. 315, L27. Lellouch, E.: 1996, in: Noll, K. S., Weaver, H. A., and Feldman, P. D. (eds.), The Collision of Comet Shoemaker-Levy 9 and Jupiter, Cambridge University Press, Cambridge, p. 213. Lellouch, E., Encrenaz, T., de Graauw, T., et al.: 2000, Planet. Space Sci. 48, 1393. Lellouch, E., Bézard,B., and Moses, J. I.: 2002, Icarus 159, 112. Lis, D. C. and Goldsmith, P. F.: 1990, Astrophys. J. 356, 195. Liseau, R., Ceccarelli, C., Larsson, B., et al.: 1996, Astron. Astrophys. 315, L181. Maret, S., Ceccarelli, C., Caux, E., Tielens, A. G. G. M., and Castets, A.: 2002, Astron. Astrophys. 395, 573. Melnick, G. J., et al.: 2000, Astrophys. J. 539, L77. Melnick, G. J., Neufeld, D. A., Ford, K. E. S., et al.: 2001, Nature 412, 160. Menten, K. and Melnick, G. J.: 1991, Astrophys. J. 377, 647. Menten, K. L., Melnick, G. J., and Phillips, T. G.: 1990a, Astrophys. J. 350, L41. Menten, K. L., Melnick, G. J., Phillips, T. G., and Neufeld, D. A.: 1990b, Astrophys. J. 363, L27. Molster, T. J., et al.: 2001, Astron. Astrophys. 372, 165. Moneti, A., Cernicharo, J., and Pardo, J. R.: 2000, Astrophys. J. 549, L203. Morbidelli, A., Chambers, J., Lunine, J. I., et al.: 2000, Meteoritics Planet. Sci. 35, 1309. WATER IN SPACE 69

Moro-Martin, A., Noriega-Crespo, A., Molinari, S., Testi, L., Cernicharo, J., and Sargent, A.: 2001, Astrophys. J. 555, 146. Müller, T. G., Abrahám,P.,´ and Crovisier, J.: 2004, this volume. Mumma, M. J., Weissman, P. R., and Stern, S. A.: 1993, in: Levy, E. H., and Lunine, J. I. (eds.), Protostars and Planets III, University of Arizona Press, Tucson, p. 1177. Neufeld, D. A. and Kaufman, M. J.: 1993, Astrophys. J. 418, 263. Neufeld, D. A., Lepp, S., and Melnick, G. J.: 1995, Astrophys. J. 100, 132. Neufeld, D. A., Chen, W., Melnick, G. J., et al.: 1996, Astron. Astrophys. 315, L237. Neufeld, D. A., Zmuidzinas, J., Schilke, P., and Phillips, T. G.: 1997, Astrophys. J. 488, L141. Neufeld, D. A., Feuchtgruber, H., Harwit, M., and Melnick, G.: 1999, Astrophys. J. 517, L147. Neufeld, D. A., Ashby, M. L. N., Bergin, E. A., et al.: 2000, Astrophys. J. 539, L111. Neufeld, D. A., Bergin, E. A., Melnick, G. J., and Goldsmith, P. F.: 2003, Astrophys. J. 590, 882. Nisini, B., Lorenzetti, D., Cohen, M., et al.: 1996, Astron. Astrophys. 315, L321. Nisini, B., Benedettini, M., Giannini, T., et al.: 2000, Astron. Astrophys. 360, 297. Nisini, B., Giannini, T., and Lorenzetti, D.: 2002, Astrophys. J. 574, 246. Nisini, B.: 2003, Astrophys. Space Sci. 287, 207. Omont, A., Moseley, S. H., Forveille, T., et al.: 1990, Astrophys. J. 355, L27. Pardo, J. R.: 1996, Ph.D. Thesis, UniversitéPierre et Marie Curie, Paris, France. Pardo, J. R., Cernicharo, J., and Serabyn, G.: 2001a, IEEE Trans. Antennas Propag. 49, 1683. Pardo, J. R., Cernicharo, J., Herpin, F., et al.: 2001b, Astrophys. J. 562, 799. Phillips, T. G., Kwan, J., and Huggins, P. J.: 1980, in: Andrew, B. H. (ed.), IAU Symposium 87: Interstellar Molecules, Reidel, Dordrecht, p. 21. Phillips, T. G., van Dishoeck, E., and Keene, J.: 1992, Astrophys. J. 399, 533. Plume, R., Kaufman, M. J., Neufeld, D. A., et al.: 2004, Astrophys. J. 605, 247. Polehampton, E. T., Baluteau, J. P., Ceccarelli, C., Swinyard, B. M., and Caux, E.: 2002, Astron. Astrophys. 388, L44. Polehampton, E. T., Menten, K., Brunken, S., Winnewise, G., and Baluteau, J. P.: 2005, Astron. Astrophys. 431, 203. Saraceno, P., Ceccarelli, C., Clegg, P., et al.: 1996, Astron. Astrophys. 315, L293. Sempere, M. J., Cernicharo, J., Lefloch, B., et al.: 2000, Astrophys. J. 530, L123. Snell, R. L., Howe, J. E., Ashby, M. L. N., et al.: 2000, Astrophys. J. 539, L101. Spoon, H. W. W., Keane, J. V., Tielens, A. G. G. M., et al.: 2002, Astron. Astrophys. 385, 1022. Sylvester, R. J., Kemper, F., Barlow, M. J., et al.: 1999, Astron. Astrophys. 352, 587. Truong-Bach, R. J., et al.: 1999, Astron. Astrophys. 345, 925. van der Tak, et al.: 2005, in: Wilson, A. (ed.), Proceedings of the Dusty and Molecular Universe: A Prelude to Herschel and ALMA, 27–29 October 2004, Paris, France, ESA SP-577, Noordwijk, ESA Publications Division, Netherlands, ISBN 92-9092-855-7, p. 431. van Dishoeck, E. and Helmich, F.: 1996, Astron. Astrophys. 315, L177. van Dishoeck, E., Wright, C. M., Cernicharo, J., et al.: 1998, Astrophys. J. 502, L173. van Dishoeck, E.: 2004, Annu. Rev. Astron. Astrophys. 42, 119. Waters, J. W., Gustincic, J. J., Kakar, R. K., et al.: 1980, Astrophys. J. 235, 57. Wright, C. M., van Dishoeck, E. F., Black, J. H., et al.: 2000, Astron. Astrophys. 358, 689. Zmuidzinas, J., Blake, G. A., Carlstrom, J., et al.: 1995, in: Haas, M. R., Davidson, J. A., and Erickson, E. F. (eds.), ASP Conference Series, Vol. 73, p. 33. Zubko, V., et al.: 2004, Astrophys. J. 610, 427–435. MOLECULAR HYDROGEN

EMILIE HABART1,∗, MALCOLM WALMSLEY1, LAURENT VERSTRAETE2, STEPHANIE CAZAUX1, ROBERTO MAIOLINO1, PIERRE COX2, FRANCOIS BOULANGER2 and GUILLAUME PINEAU DES FORETSˆ 2 1Osservatorio Astroﬁsico di Arcetri, INAF, Largo E. Fermi 5, I-50125 Firenze, Italy 2Institut d’Astrophysique Spatiale, Universite´ Paris-Sud, 91405 Orsay, France (∗Author for correspondence: E-mail: [email protected])

(Received 16 July 2004; Accepted in ﬁnal form 18 November 2004)

Abstract. Observations of H2 line emission in galactic and extragalactic environments obtained with the Infrared Space Observatory (ISO) are reviewed. The diagnostic capability of H2 observations is illustrated. We discuss what one has learned about such diverse astrophysical sources as photon- dominated regions, shocks, young stellar objects, planetary nebulae and starburst galaxies from ISO observations of H2 emission. In this context, we emphasise use of measured H2 line intensities to infer important physical quantities such as the gas temperature, gas density and radiation ﬁeld and we discuss the different possible excitation mechanisms of H2. We also brieﬂy consider future prospects for observation of H2 from space and from the ground.

Keywords: molecular hydrogen, infrared: spectroscopy, molecular processes, ISM: molecules, ISM: clouds, star formation, circumstellar matter, infrared: galaxies

1. Physics of H2

1.1. ASTROPHYSICAL IMPORTANCE OF H2

Molecular hydrogen is the most abundant molecule in the Universe and plays a fundamental role in many astrophysical contexts (e.g., Dalgarno, 2000). It is found in all regions where the shielding of the ultra-violet (UV) photons, responsible for its photo-dissociation, is sufficiently large, i.e. where AV 0.01–0.1 mag. H2 makes up the bulk of the mass of the dense gas in galaxies and could represent a significant fraction of the total baryonic mass of the Universe. Further on, H2 plays two main roles that render it key for our understanding of the interstellar medium, in particular of the processes regulating star formation and the evolution of galaxies. Firstly, the formation of H2 on grains initiates the chemistry of interstellar gas; secondly, H2 is recognized as a major contributor to the cooling of astrophysical media. H2 has also specific radiative and collision properties that make it a diagnostic probe of unique capability. Many competing mechanisms could contribute to Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, The Netherlands, and the United Kingdom), and with the participation of ISAS and NASA. Space Science Reviews (2005) 119: 71–91 DOI: 10.1007/s11214-005-8062-1 C Springer 2005 72 E. HABART ET AL. its excitation (see Section 1.2) and H2 can serve as probe of a wide range of physical environments. Thus, as we understand reasonably well its radiative and collision properties we can construct realistic models of the response of H2 to its surroundings.

1.2. EXCITATION AND FORMATION OF H2

The hydrogen molecule is highly symmetric as it has containing two identical hydrogen atoms. Due to this symmetry, the molecule has no dipole moment and all ro-vibrational transitions within the electronic ground state are quadrupolar with low spontaneous coefﬁcient A values. The molecule exists in two, almost independent states, namely ortho-H2 (spins of H nuclei parallel) and para-H2 (spins antiparallel). There are no radiative transitions between ortho- and para- H2 but ortho-para conversion may occur through proton exchange reactions between H2 and, for example, H0 and H+. Molecular hydrogen may be excited through several mechanisms: (1)Inthe presence of far-ultraviolet radiation (FUV, λ>912 A),˚ the molecule is radiatively pumped into its electronically excited states. As it decays back into the electronic ground state, it populates the high vibrational levels, and the subsequent cascade to v = 0 gives rise to optical and infrared ﬂuorescent emission, and a characteristic distribution of level populations (e.g., Black and van Dishoeck, 1987, Sternberg, 1989). This excitation mechanism is observed in photon-dominated regions (PDRs) where it dominates the excitation of ro-vibrational and high rotational levels. (2) If the gas density and temperature are high enough, inelastic collisions can be the dominant excitation mechanism, at least for the lower energy levels (e.g., Le Bourlot et al., 4 −3 1999). In dense PDRs (n H 10 cm ) and shocks, collisions maintain the lowest pure rotational levels in thermal equilibrium. (3)H2 formation in excited states can also contribute to the excitation of the molecule. (4) Finally, in environments such as active galactic nuclei (AGNs) or X-ray emitting young stellar objects where hard X-rays are capable to penetrate deeply into zones opaque to UV photons, X-rays can dominate H2 excitation (Maloney et al., 1996; Tine et al., 1997). Except in the early Universe, most H2 is thought to be produced via surface reactions on interstellar grains (e.g., Gould and Salpeter, 1963; Hollenbach and Salpeter, 1971), since gas-phase reactions are predicted to be too slow. But due to our limited understanding of the relevant properties of interstellar grains (composition, structure and hydrogen coverage) and hence of grain surface reactions, the H2 formation mechanism is not yet understood. Numerous theoretical and experimental studies have thus been dedicated to a better understanding of the H2 formation process (Gould and Salpeter, 1963; Hollenbach and Salpeter, 1971; Sandford and Allamandola, 1993; Duley, 1996; Parneix and Brechignac, 1998; Pirronello et al., 1997; Pirronello et al., 1999; Takahashi et al., 1999; Katz et al., 1999; Williams et al., 2000; Sidis et al., 2000; Biham et al., 2001; Joblin et al., 2001; Cazaux and MOLECULAR HYDROGEN 73

Tielens, 2002; Cazaux and Tielens, 2003). Another approach to this issue is to use observations to provide estimates of the H2 formation rate in different regions of the ISM in order to see how it depends on the local physical parameters. Recently, combined ISO observations of H2 line emission and dust emission as well as FUSE observations of UV pumping lines in absorption have highlighted the H2 formation process (see Section 3.1).

2. Observations of H2

Direct observation of H2 is difﬁcult. Electronic transitions occur in the ultraviolet region, to which the Earth’s atmosphere is opaque, and observations can only be made from a space-based platform. Ro-vibrational and rotational transitions are faint because of their quadrupolar origin. In the case of rotational lines which occur in the mid-IR, the Earth’s atmosphere is at best only partly transparent. Moreover, most of H2 may hide in cool, shielded regions (e.g., Combes, 2000) where the excitation is too low to be detected through emission lines and whose extinction is too high to allow the detection of the UV pumping lines. Space-based missions in the UV (Copernicus, ORPHEUS, FUSE) and the mid- IR (ISO, Spitzer) provide the most stringent tests of our current understanding of H2 in space.

2.1. UV ABSORPTION LINES

In the 1970s, the Copernicus satellite observed molecular hydrogen electronic lines in absorption towards nearby early type stars (e.g., Spitzer et al., 1974). This experiment yielded reliable column densities for diffuse clouds (AV ≤ 1), the ﬁrst measurements of the temperature of H2 in diffuse clouds (∼50–100 K) as well as an estimate of the formation rate of H2 on interstellar dust grains (Jura, 1975). Re- cently, thanks to the greater sensitivity of FUSE, fainter stars with higher extinctions could be observed allowing the study of translucent clouds (e.g., Rachford et al., 2002). However, such observations are limited by the small number of sufﬁciently luminous background sources against which H2 can be detected in absorption. In practice, UV observations only allow the study of molecular gas in the Solar Neighborhood and provide no spatial information on the distribution of H2.

2.2. INFRARED OBSERVATIONS

Prior to the advent of ISO, infrared emission of H2 was observed in the 2 μm ro- vibrational lines (most notably the 1-0 S(1) line at 2.12 μm) towards Galactic shocks and PDRs, planetary nebulae, and in some instances the nuclei of active galaxies (for the ﬁrst studies see e.g. Shull and Beckwith, 1982, Fischer et al., 1987). However, in PDRs and shocks only a small fraction of the H2 is vibrationally excited. 74 E. HABART ET AL.

The ﬁrst observations of the v = 0 pure rotational lines of H2 have been done from the ground towards a PDR – the Orion bar – by Parmar et al. (1991). The lowest rotational transitions of H2 produced by collisions (see Section 1.2) provide a thermometer for the bulk of the gas above ∼80 K. Also, due to the low A-values of these transitions, optical depth effects are usually unimportant and thus these lines provide accurate measures of the warm (T 80 K) gas mass. With the sensitivity of the ISO instruments (i.e., SWS and CAM) several pure rotational lines of H2 have been observed in a variety of galactic and extragalactic environments. Moreover for each type of environment, numerous objects with a wide range of physical conditions were observed allowing us to adress a number of outstanding questions concerning the physics of H2. Moreover, the complete ISO wavelength coverage (2–200 μm), which includes many lines due to gas-phase + 0 + 0 atoms (e.g., C ,O,Si ,S), molecules (e.g., H2, CO, H2O, OH and CO2) and bands of polycyclic aromatic hydrocarbons (PAHs), silicates and ices, allows us to study the links between H2, the dust and the other gas components. A review on ISO spectroscopic results from molecular clouds to disks by van Dishoeck (2004) illustrates the diagnostic capabilities of the various lines and bands detected by ISO. This article summarizes the ISO spectroscopic observations of H2. The bulk of the data comes from the Short Wavelength Spectrometer (SWS, 2–45 μm, de Graauw et al., 1996). Relevant spectra have also been obtained with the Camera- Circular Variable Filter (CAM-CVF, 2.3–16.5 μm, Cesarsky et al., 1996).

3. An Overview of ISO Observations of H2 in Galactic and Extragalactic Environments

Infrared Space Observatory has observed H2 emission from sources as diverse as photo-dominated regions (PDRs), shocks associated with outﬂows or supernovae remnants, circumstellar environment of young stellar objects, planetary nebulae, and external galaxies. Combined with the theoretical models of PDRs and shocks these observations have allowed the derivation of important physical quantities in these objects, namely, the hydrogen density nH, the intensity of the incident far- 1 ultraviolet (FUV,6 < hν<13.6 eV) radiation represented by χ, the H2 rotational excitation temperature Trot (a measure of the gas temperature, see below) and the H2 column density NH2 . In this section, these H2 observations and the inferred physical quantities are reviewed (see Table I for a summary).

3.1. PHOTON-DOMINATED REGIONS

Photon-dominated regions are one of the most important sources of H2 emission. PDRs formed at the surfaces of molecular clouds where FUV radiation encounters 1The scaling factor χ represents the intensity of the FUV field in units of the Draine (1978) average interstellar radiation field. MOLECULAR HYDROGEN 75 et al. (2000); 6 4124 ◦ in Habart 6 5 –10 3 –10 –10 et al. 4 4 13 150 10 10 10 470 7 6 225 BD + 40 –10 3 rotational levels is essentially 5 2 Starbursts Seyferts 13 150 10 1–100 1–100 YSOs 10 7 LkH α temperature but two components H (2000b); (8) Herpin et al. 4 10 500 720 6 0.1–1 0.069 0.055 5 –10 2 6 10 100 12 80 –10 4 688 10 5 0.18 Shocks (1999); (7) van den Ancker (1996); (12) Valentijn and van der Werf (1999); (13) Rigopoulou 5 6 et al. –10 –10 2 et al. 5 10 23 11 170 NGC 6946 NGC 891-outer disk Normal galaxy 630 10 4 1.4 TABLE I 4 (1994, 1995); (3) data from Bertoldi, private communication and see references 5 10 observations in various astrophysical objects. 10 2 × 6 6 × et al. (1996); (6) Cesarsky –10 –10 3 3 0 10–20 10 150 10 center Galactic (2001); (11) Valentijn 0.5–2.4 Orion Bar Orion Pk1 Cep A W IC 443 390 0.5–3 3 1 et al. rotational lines and assuming that the population distribution of low 2 of the Draine (1978) average interstellar radiation field. (1999) (note that the observations cannot be fitted by a single excitation et al. χ et al. , Jansen 5 Sample of ISO H et al. 5 pure rotational levels. 0.003 9 ≥ 10 41 2 PDRs 0.5–1 × 10 2 (2000); (5) Wright 7 (2003); (2) Thi et al. –10 6 4 0 0 et al. 8 750 900 CRL 618 Helix PPN PNe 250 650 Oph W IC 63 330 620 1 ∼ 620 K are needed) and Jansen ) 0.6 5 ) 68 (1998); (10) Rodriguez-Fernandez − 2 − 2 cm cm )1 )1 et al. 21 21 − 3 − 3 (2004); (4) Rosenthal column density inferred from the intensity of H (10 (K) (10 (K) (cm 2 2 (cm 2 ∼ 100 K and c H b a c H rot H b a Excitation temperature of the low H Incident FUV radiation field expressed in units H H rot N T n χ Objects Ref. a b c in LTE. References: (1) Habart at (2002). et al. (9) Cox N T Ref. n χ Objects 76 E. HABART ET AL. and dissociates the dense molecular gas. They are the transition zone between the dense, cold molecular gas and the tenuous, warm, ionized gas (for a recent review see Hollenbach and Tielens, 1999). The thermal and chemical structure of the PDR is governed by the FUV radiation2 and depends mainly on two parameters, namely, the incident FUV flux χ and the density nH. Infrared Space Observatory has observed a series of pure H2 rotational lines towards a number of PDRs sampling a wide range of excitation conditions: high excitation PDRs with χ>103 such as the Orion Bar (Bertoldi et al., private communication), NGC 2023 (Moutou et al., 1999; Draine and Bertoldi, 2000a), NGC 7023 (Fuente et al., 1999); moderately excited PDRs with 102 <χ<103 such as IC 63 (Thi et al., 1999), S 140 (Timmermann et al., 1996) and ρ Oph W (Habart et al., 2003); and low excitation PDRs such as the cool PDR Ced 201 (Kemper et al., 1999) or the S 140 extended region (Li et al., 2002). Detailed studies using ISO observations of both H2 and dust emission, in conjunction with ground-based NIR imaging, allows us to bring new insights mainly into (i) the temperature and density structure of PDRs, (ii) the dominant heating and cooling processes and (iii) the H2 formation process.

3.1.1. H2 as a Thermal Probe in PDRs 4 −3 At densities nH 10 cm of bright PDRs, collisions maintain the lowest rotational levels of H2 (v = 0, J 5) in thermal equilibrium. The population of these levels is therefore a good indicator of the gas temperature. This is reﬂected in their excitation diagrams (see Figure 1) where their populations are consistent with the Boltzmann law. We report in Table I the excitation temperature of the H2 pure rotational levels for some PDRs with different excitation conditions. All the derived Trot are found to be similar between 300 and 700 K. It should be noted that, in spite of their simple appearance, the interpretation of these diagrams has to take account of several factors. First, the gas temperature varies rapidly through the PDR layer from several hundred K at the edge close to the excitation source to less than 30 K when AV > 1. The pure rotational H2 lines arise primarily in the outer warm layer. Second, the excitation temperature of H2 pure rotational levels gives in principle an upper limit to the kinetic temperature as UV pumping may contribute to the excitation of H2 even for low J. For several PDRs observed by ISO, the H2 line intensities and the gas temperature are found to be higher than predicted by current models (Bertoldi, 1997; Draine and Bertoldi, 1999b; Thi et al., 1999; Habart et al., 2003; Li et al., 2002). The cause of this discrepancy is either that, in the models, the gas is not hot enough or alternatively that the column density of H2 is too low in the zones where the

2The FUV radiation ﬁeld sets the PDR chemical structure by ionizing carbon, silicon, etc., and by dissociating H2, CO and most other molecules. It also heats the gas mainly via the photoelectric effect on dust grains. MOLECULAR HYDROGEN 77

Figure 1. Upper panels. (a) Left: ISOCAM map, with the LW2 filter (5–8.5 μm), of the bright filament along the western edge of the ρ Ophiuchi main cloud (Abergel et al., 1996). The filament is immersed in the radiation field of the star HD147889 (see the star sign). The four SWS positions (small squares) and the CVF map positions are marked (big squares). (b) Right: ISOCAM-CVF brightness profile of the 0-0 S(3) H2 line (solid line) and PAH emission (dotted lines) along the cut going through the SWS pointings. Lower panels. (a) Left : Pure rotational H2 lines observed towards the H2 emission peak of the ρ Oph filament. The line close to 0 shows the rms noise. (b) Right : Excitation diagram (corrected for dust attenuation) for H2; Nu is the column density of the transition upper level, gu is the degeneracy of the upper level and T u is the upper level energy in Kelvin. The solid line is rotational temperature (Trot) thermal distribution for an H2 ortho-to-para equal at 1. gas is warm. Several possible explanations have been proposed. An increase of the grain photoelectric heating rate is one possibility. Draine and Bertoldi (2000a) have, in particular, shown that increased photoelectric heating rates based on an enhanced dust-to-gas ratio in the PDR due to gas-grain drift can reproduce the H2 observations in the reflection nebula NGC 2023. Another explanation, namely an increased H2 formation rate, is discussed below. Finally, non-equilibrium processes, e.g. propagation of the ionization and photodissociation fronts which will bring 78 E. HABART ET AL. fresh H2 into the zone emitting line radiation, can be invoked but detailed models are needed.

3.1.2. H2 Formation Rate Recently, Habart et al. (2004) combined ISO observations of rotational lines of H2 towards a sample of PDRs with observations of ro-vibrational lines made with ground-based telescopes in order to constrain the formation rate of H2 under the physical conditions in the layers responsible for H2 emission. Using as a diagnostic the 0-0 S(3)/1-0 S(1) line ratio, which strongly depends on the value of the H2 formation rate Rf, they find that, in regions of moderate excitation (χ ≤1000, such as Oph W, S140 and IC 63), Rf is about a factor 5 larger than the standard rate estimated in diffuse clouds. On the other hand, towards regions with higher χ (such as NGC 2023 and the Orion Bar), R f is found to be similar to the standard value. This result can provide at least a partial explanation of the discrepancy discussed 0 above since a larger Rf will move the H /H2 transition zone closer to the edge of the PDR and consequently increase the temperature of the H2 emitting gas as well as the absolute intensity of the H2 lines. This finding of efficient H2 formation at high gas temperatures (Tgas ≥ 300 K) has also fundamental implications for our understanding of the H2 formation process. Since the residence time of weakly bound H atoms on grains (also called physisorbed) at such high temperatures is very short, a process involving strongly bound H atoms is required. An indirect chemisorption process - where a physisorbed H-atom scans the grain surface to recombine with a chemisorbed H-atom (as recently discussed by Cazaux and Tielens, 2002; Cazaux and Tielens, 2003) – is capable of explaining the ISO data (Habart et al., 2004) and could be the most important mechanism in PDRs. Moreover, small grains (radii < a few 10 nm) which dominate the total grain surface and spend most of their time at relatively low (below 30 K for χ ≤ 3000) temperatures may be the most promising surface for forming H2 in PDRs (Habart et al., 2004). The relationship between aromatic dust particles (referred to as PAHs, radii < 1 nm) and H2 has been addressed both in the laboratory (Joblin et al., 2001) and observationally (Le Coupanec, 1998; Joblin et al., 2000; An and Sellgren, 2001; Habart et al., 2003; Habart et al., 2004). In particular, using ISOCAM and ground-based observations for a sample of PDRs, (Habart et al., 2004) find that the 3 Rf/[C/H]PAH ratio is constant. This suggests, but does not prove, that formation of H2 on PAHs could be important in PDRs. This result is supported by the observed spatial correlation between the H2 line emission and the PAH features (see Figure 1 and Le Coupanec, 1998; Joblin et al., 2000; An and Sellgren, 2001; Habart et al., 2003). In the H2 formation process, very small grains or VSGs (radii between a few nm and a few 10 nm) may also be important (e.g., Boulanger et al., 2004; Rappacioli et al., 2004). In the near future, Spitzer will allow us to examine the

3 [C/H]PAH is the abundance of carbon locked up in PAHs. MOLECULAR HYDROGEN 79

VSG emission between 15 and 40 μm towards a number of PDRs and we will be able to better understand the nature of VSGs and to determine their role in the H2 formation process.

3.1.3. Ortho-to-Para H2 Ratio In several cases, such as the ρ Oph-West (Habart et al., 2003), NGC 2023 (Moutou et al., 1999) or NGC 7023 PDR (Fuente et al., 1999), the H2 ortho-to-para ratio is significantly smaller than the equilibrium ratio of 3 expected in gas at the temperatures derived from excitation diagrams. This could indicate the presence of (i) dynamical processes, such as the advection of colder gas into the PDR, or (ii) fast conversion of ortho-H2 to para-H2 during the formation on the surface of non magnetic, non metallic grains (Le Bourlot, 2000). It must be emphasized here that, as pointed out by (Sternberg and Neufeld, 1999), care has to be taken with the interpretation of the data since a low ortho/para ratio in excited levels naturally results from the fact that the optical depth in the UV pumping lines is higher for the ortho- than for the para-H2 lines. Moreover, none of the ISO data measure the bulk of ortho- and para-H2 which resides in J = 0 and 1. We note however that for the physical conditions (χ and nH) of PDRs discussed here, population of the rotational levels by the cascade following the FUV pumping is generally negligible below S(4) and that the lower lying rotational levels are efficiently populated by collisions. In this latter case, the excitation temperature is well defined and the populations of the J = 0 and 1 levels can be derived from a Boltzmann law.

3.2. SHOCKS

As in the case of PDRs, shocks are luminous sources of H2. Shocks in molecular clouds can be caused by a large variety of processes including outflows and jets from young stars, supernovae and expanding H II regions (for theoretical reviews see Draine and McKee, 1993; Hollenbach, 1997). In a shock, the mechanical energy is converted into heat and radiated away through cooling lines of atoms and molecules in the IR. For a wide range of shock conditions, H2 molecules are not dissociated and the gas becomes warm enough for the lowest rotational lines to be excited by collisions. Infrared Space Observatory has detected many lines of H2 in a number of shocks associated with outflows as Orion OMC-1 (Rosenthal et al., 2000), Cep A West (Wright et al., 1996), DR 21 (Wright et al., 1997), HH 54 (Neufeld et al., 1998), HH 2 (Lefloch et al., 2003), L1448 (Nisini et al., 2000) and also in shocks associated with the supernova remnants IC 443 (Cesarsky et al., 1999), 3C 391, W 44 and W 28 (Rho and Reach, 2003). ISO’s main contributions to the study of shocks have been to (i) better characterize the shock structure and physical conditions, especially that found in the slower shocks; (ii) better measure the total cooling power. The ISO data have also convincingly shown that most of the broad-band emission from shocks is due to lines (e.g., Cesarsky et al., 1999). For example, the ISOCAM 80 E. HABART ET AL.

7- and 15-μm bands are dominated by the H2 S(5) and S(1) lines, whereas the IRAS 12, 25, and 60 μm emission is due to H2 S(2), S(0), and [O I] 63 μm, respectively. As for PDRs, the pure rotational H2 transitions dominate the mid-infrared spectra and are characterized by excitation temperatures of Tex ∼ 500–1000 K (see Table I and Figure 2). In contrast, the well-studied H2 ro-vibrational lines usually yield Tex = 2000–3000 K. The most impressive collection of H2 lines is seen for Orion (see Figure 2, Rosenthal et al., 2000), where 56 H2 lines have been seen including the v = 0 S(1)–S(25) lines with upper energy levels up to 42,500 K. Since Tex increases strongly with shock velocity, the observed range of Tex has generally been interpreted in terms of two steady state C-shocks; one with a low density, low velocity (<20 km s−1) and high covering factor to explain the warm, ∼700 K, gas, and another with a higher density, higher velocity (30–50 km s−1) and low covering range to explain the hot, ∼ 2000 K, gas (e.g., Wright et al., 2000). However, in some circumstances, shocks are unlikely to have attained a state of equilibrium, and a time-dependent approach must be considered. The model of non stationary C-shocks of (Pineau Des Forêtsand Flower, 2000) shows that, at early times, shocks exhibit both C- and J-type characteristics and the existence of both warm and hot components can be explained in terms of a single shock. The H2 excitation diagram can thus provide a measure of the age of the shock. In the case of the outflow observed in Cep A West (Pineau Des Forêtsand Flower, 2000) as well as the supernova remnants IC 443 (Cesarsky et al., 1999), the shocked gas is estimated to have an age of approximately 1.5 × 103 year. For most shocks, the inferred ortho/para ratio is 3, consistent with the warm conditions. Major exceptions are HH 54 (Neufeld et al., 1998; Wilgenbus et al., 2000) and HH 2 (Lefloch et al., 2003), where significantly lower ortho/para ratios of 1.2–1.6 are found. This low ratio is thought to be a relic of the earlier cold cloud stage, where the timescale of the shock passage was too short to establish an equilibrium ortho/para ratio. Thus, the ortho/para ratio can be used as a chronometer, indicating for HH 54 that the gas has been warm for at most 5,000 yr.

3.3. CIRCUMSTELLAR ENVIRONMENT OF YSOS

H2 is expected to be ubiquitous in the circumstellar environment of Young Stellar Objects (YSOs). It is the main constituent of the molecular cloud from which the young star is formed and is also expected to be the main component of the circumstellar disk. Most of this material however is at low temperatures and the pure rotational lines are difﬁcult to observe. However, certain zones may be heated to temperatures of a few hundred K and produce observable H2. The intense UV radiation generated by accretion as well as that from the central star itself creates a PDR in the surrounding material. Another way of producing warm H2 is in shocks caused by the interaction of an outﬂow with the surrounding molecular cloud or envelope. MOLECULAR HYDROGEN 81

Figure 2. Upper panel: SWS spectrum towards the brightest H2 emission peak of the Orion OMC-1 outﬂow of Orion Peak 1 showing a multitude of ro-vibrational and pure rotational lines (Rosenthal et al., 2000). Lower panel: Excitation diagram (corrected for dust attenuation) for H2; for the lowest energy levels the excitation temperature is about 600 K whereas at E(v, J)/K ≥14,000 K it is about 3000 K. 82 E. HABART ET AL.

3.3.1. Intermediate- and High-Mass YSOs: From Shocks to PDRs van den Ancker et al. (2000a,b) have studied H2 emission in a sample of 21 YSOs mostly of intermediate- and high-mass. Pure rotational H2 emission was detected in 12 sources. In Figure 3, we show the excitation diagrams of some sources at different stages of evolution. To determine the excitation mechanisms responsible for the H2 emission, van den Ancker et al. (2000a,b) compare these observations to PDR and shock models. However, for the physical conditions prevailing in these objects, there is generally a considerable overlap between the temperatures predicted by the different models. Thus, van den Ancker et al. (2000a,b) use additional tracers, such as PAH emission – indicative of the presence of a PDRy – or [SI] 25.25 μm emission indicative of shocks. They ﬁnd that the H2 lines of the embedded sources originate from shocks, whereas more evolved objects are characterised by H2 lines from a PDR. As the envelope disperses, the PDR component starts to dominate compared to the shock emission.

3.3.2. Circumstellar Disks Deep searches for the H2S(0) and S(1) lines at 28 and 17 μm have been performed with the SWS towards a dozen young stars (Thi et al., 2001). Detections have been claimed in several sources but contamination from surrounding cloud material can be important (e.g., Richter et al., 2002). Recently, using TEXES/IRTF with higher spectral (λ/λ ∼ 80, 000 at 12.3 μm) and spatial resolution (0.85 at 12.3 μm), (Richter et al., 2004) have detected the H2 S(1) and S(2) lines in circumstellar disks around several Herbig Ae and T Tauri stars. Preliminary analysis suggests that the gas temperatures are between 400 and 800 K. This is consistent with recent FUSE observations of two Herbig Ae stars, namely HD100546 and HD163296, suggesting the presence of dense and warm (∼400–700 K) H2 in the circumstellar disks (Lecavelier des Etangs et al., 2003). In the near future, observations from Spitzer and VISIR, accompanied by more sophisticated modelling, should be able to clarify the origin of the H2 emission from disks.

3.4. PLANETARY NEBULAE AND PPN

H2 is also an important component of the neutral material of envelopes expanding away from the hot central stars of Planetary nebulae (PNe). In these objects, the H2 emission may originate in the shock between the shell and the wind of the precursor red giant or in the PDR on the inside of the neutral shell. Moreover, the H2 excitation mechanism is expected to change with time (e.g., Natta and Hollenbach, 1998). Presently, the largest body of evidence comes from observations of the H2 2-μm ﬂuorescent transitions. The morphology and excitation of the H2 ﬂuorescent emission in AFGL 2688 and NGC 7027 are, in particular, consistent with an evolutionary scheme in which shock excitation dominate in the proto-planetary nebula (PPN) phase and FUV excitation from PDRs dominate in young PNe (e.g., Cox et al., 1997). MOLECULAR HYDROGEN 83

Figure 3. Excitation diagrams (corrected for dust attenuation) for H2 of young intermediate-mass stars in the BD + 40◦4124 group (van den Ancker et al., 2000b). From bottom to top, we ﬁnd in a rough evolutionary sequence LkHα 225, LkHα 224 and BD + 40◦4124.

Infrared Space Observatory allowed us to observe several rotational and ro- vibrational lines of H2 from PPN (such as CRL 618, Herpin et al., 2000), young nebulae (such as NGC 7027, Bernard Salas et al., 2001), and fully evolved PNe (such as the Helix or the Dumbbel, Cox et al., 1998; Bachiller et al., 2000). Anal- ysis of these data brings new perspectives on the different possible H2 excitation mechanisms at different evolutionary stages of PNe. As an illustration, we discuss in the following two extreme cases, namely the proto-planetary nebula CRL 618 and the evolved Helix nebula. 84 E. HABART ET AL.

3.4.1. The PPN CRL 618 Many rotational and ro-vibrational lines of H2 were detected by ISO-SWS towards the proto-planetary nebulae CRL 618 (Herpin et al., 2000). The analysis of these data shows that shocks in the lobes may collisionally excite the lowest pure rotational emission whereas the highest transitions (v ≥ 2) may come from fluorescence in the torus. A mix of fluorescence and collisions may produce the v = 1 to 0 transitions in an intermediate zone. Recent high spatial and spectral resolution spectro-imaging of the H2 1-0 S(0) line using BEAR reveal the presence of multiple, high-velocity, molecular outflows that align with the optical jets, and show how jets interact with circumstellar gas and shape the environment in which planetary form (Cox et al., 2003).

3.4.2. The Helix In the Helix, an evolved PNe, a detailed study of the H2 rotational lines (Cox et al., 1998) has been carried out with ISOCAM data. The CVF spectro-imaging data shows that the mid-IR emission is entirely caused by H2 and neon lines (see Figure 4); no dust bands are detected and the continuum emission is very weak. The H2 emission traces the individual cometary globules of the molecular envelope of the nebula (see Figure 4). Despite the fact that the globules in the Helix show a typical PDR morphology, (Cox et al., 1998) conclude that the far-UV radiation from the central star is of much too low intensity (χ ∼ 6) to account for the H2 line intensities and high temperatures (900 K) that are observed. The extreme-UV radiation from the star could perhaps play an important role in heating the H2 if it survives up to the ionization front. Time-dependent calculations including the effects of rapid changes in the stellar effective temperature, χ and nH as well as the advance or retreat of the ionization front with respect to the PDR gas are needed to better understand the physics of PNe. Comparison with future H2 observations at high angular resolution should help to further constrain the excitation of H2 in PNe.

3.5. GALACTIC CENTER REGION AND DIFFUSE GALACTIC MEDIUM

Infrared Space Observatory has also provided some clues in our understanding of warm gas in the galactic center region as well as in the diffuse interstellar medium.

3.5.1. Warm H2 in the Galactic Center Region Several pure-rotational lines have been observed towards a sample of molecular clouds distributed along the central 500 pc of the Galaxy (Rodriguez-Fernandez et al., 2001). The S(0) and S(1) lines trace gas with temperatures of 150 K while the S(4) and S(5) lines indicate temperatures of 600 K. The 150 K-H2 column density is typically 1–2 × 1022 cm−2. This is the ﬁrst direct estimate of the total column density of the warm molecular gas in the Galactic center region. These warm column MOLECULAR HYDROGEN 85

Figure 4. Upper panel: An ISOCAM map of the Helix Planetary Nebula in the LW2 ﬁlter dominated by the strong 0-0 S(5) H2 line (Cox et al., 1998). Lower panel : CVF spectra toward the H2 emission peak in the western rim of the Helix nebula. densities represent a fraction of 30% of the gas traced in CO. In order to explain all the observed H2 lines, (Rodriguez-Fernandez et al., 2001) ﬁnd that a combination of both high and low density PDRs or shocks is required. Moreover, non-equilibrium ortho-to-para ratios have been found in two clouds of the sample suggesting heating by low velocity shocks (∼10 km s−1, Rodriguez-Fernandez et al., 2000). 86 E. HABART ET AL.

3.5.2. Hot Gas in the Cold Diffuse Medium Most of the diffuse interstellar medium is cold, but it must have pockets of hot gas to explain the large observed abundance of molecules such as CH+, OH and HCO+. To test the existence and determine the characteristics of this component, (Falgarone et al., 2000) have searched using ISO-SWS for H2 rotational lines along a line of sight close to the galactic plane, carefully avoiding star forming regions and molecular clouds. The intensities of the S(2) and S(3) lines are about one order of magnitude above those predicted by radiative excitation. This result suggests the existence of a small fraction (few percent) of hot (a few 100 K) gas in the cold interstellar medium and for which UV photons cannot be the sole heating source. The possible excitation process could be the deposition of mechanical energy in shocks or in vortices formed by the viscous dissipation of turbulence (see in this review the chapter on the cool interstellar medium by Abergel et al. (1996)).

3.6. EXTRAGALACTIC SOURCES

Prior to ISO, H2 emission in external galaxies was detected in the near-IR ro- 4 vibrational lines, from typically 10 M of H2 heated to temperatures of ∼2,000 K, in the centers of starburst, active and (ultra)luminous infrared galaxies. However, −6 gas at these temperatures is a very small fraction (10 ) of the total amount of H2 gas (e.g., van der Werf et al., 1993). ISO-SWS offered the unique opportunity to detect pure rotational H2 emission in various galaxies, from normal to ultraluminous, and thus to study the amount of warm (T 80) gas in these galaxies (e.g., Valentijn et al. 1996; Valentijn and van der Werf, 1999; Genzel and Cesarsky, 2000; Rigopoulou et al., 2002; Lutz et al., 2003). These observations provide unique probes of the various heating mechanisms. For instance, in starburst galaxies, H2 emission may be generated by UV-pumping in star-forming regions or by shocks due to outﬂows, supernova remnants or large-scale streaming motions in spiral arms and bars; in addition X- ray excitation may be produced by multiple supernovae remnants or by an active galactic nucleus or cooling ﬂow.

3.6.1. Normal Galaxies The first extragalactic detection of the 0-0 S(0) and S(1) lines in the center of the 6 nearby Scd galaxy NGC 6946 revealed 5 × 10 M of warm, (T = 170 K) circumnuclear gas (Valentijn et al., 1996). Radiation from massive stars or shocks from e.g. supernovae explosions in the nuclear complex are the most obvious mechanisms for producing this warm component. Throughout the disk of the edge-on spiral galaxy NGC 891, such warm H2 components have also been detected (Valentijn and van der Werf, 1999). Additionally, a cooler Tgas ∼ 80 K component was observed in the outer disk. This can most likely be identified with extended low-density PDRs which form the warm envelopes of giant molecular clouds, heated by the local in- MOLECULAR HYDROGEN 87 terstellar medium. Moreover, this component may harbour a very significant mass and the implications for the galactic mass budget have been explored by (Valentijn and van der Werf, 1999). This finding might provide a solution to the missing mass problem.

3.6.2. Starburst and Seyfert Galaxies Using ISO-SWS, Rigopoulou et al. (2002) has performed a survey of H2 emission from active galaxies displaying a wide range in nuclear activity including starburst and Seyfert galaxies. A number of transitions from 0-0 S(0) to S(7) are detected in both starbursts and Seyferts. The data reveal different components at different temperatures. The temperature of the more abundant warm gas – probed by the S(0) and S(1) lines – is similar in starbursts and Seyferts with a value of around ∼150 K. This warm gas accounts for as much as 10% of the total galactic gas mass in starbursts and up to 35% in Seyferts. Also, by comparison with model predictions, (Rigopoulou et al., 2002) find that a combination of various PDRs could explain the line ratios observed in starbursts. Although such a combination of PDRs can also match the observations in Seyferts, it is likely that shocks and heating from the central X-ray source are also present. This is confirmed by the observed H2/PAH emission ratio. In fact, in starbursts and Seyferts with a strong starburst component, (Rigopoulou et al., 2002) find that the PAH and H2 emission are well correlated implying that these emissions have a common starburst origin. On the other hand the H2/PAH emission ratio is higher in pure “AGN-dominated” objects suggesting the presence of an extended circumnuclear component of warm gas heated by the nuclear X-ray emission in which the enhanced H2 emission originates. This warm gas component is either too far away for UV photons to reach or, acts as a shield to UV photons resulting in both cases in suppressed PAH emission.

4. Conclusions and Future Prospects

Infrared Space Observatory provided a fundamental step forward in that it enabled us for the first time to exploit the potential of the pure rotational series of H2.Itwas possible to estimate temperatures and densities in the important transition layers of the ISM separating neutral and ionised gas. Such measurements also allowed estimates of the H2 formation rate on grain surfaces and suggested that molecular hydrogen might preferentially be formed on PAHs and/or very small grains. ISO H2 data also allowed us to better characterize the shocks associated with outflows from young stars or supernova remnants. H2 lines can in particular be used to estimate the age of the shock. Moreover, new insights into the different possible H2 excitation mechanisms at different evolutionary stages of young stellar objects and planetary nebulae have been obtained. ISO also gave us for the first time the possibility to observe the pure rotational lines of H2 in galaxies, from normal to ultraluminous. This permitted the study of the amount of warm (T 80 K) gas in these galaxies 88 E. HABART ET AL. and provided unique probes of the dominant heating mechanisms (photon heating, shock, X-ray). An obvious question is that of what one can expect from future space missions as well as from the use of modern spectrometers on large ground-based instruments. It is perhaps worth stressing that our present capacity to use infrared observations of H2 as a probe of interstellar (galactic and extragalactic) conditions is limited by a variety of factors including sensitivity but also spatial and spectral resolution. The importance of high-spatial resolution is made clear both by ground-based images in the 2 μm 1-0 S(1) line as well as by the difficulties encountered both from the ground and space in separating diffuse emission and emission from point sources in crowded fields. Hardly less important is high-spectral resolution both because of the kinematic information which becomes available for R = λ/λ greater than 10,000 and because of the problems of blending of weak spectral features. The IRS instrument on Spitzer covering the 5–40 μm range is already demonstrating the advantages of increased sensitivity for the study of H2 lines in young star forming regions (e.g., Morris et al., 2004) and galaxies (e.g., Smith et al., 2004; Armus et al., 2004). Spitzer will extend medium resolution spectroscopy of H2 lines into regimes that ISO could not probe, and will provide a large number of new targets for further works. Sensitive ground-based studies are already possible with ISAAC at the VLT (wavelength range 1–5 μm, R ∼ 1,000–10,000) but the future presumably belongs to CRIRES (wavelength range 1–5 μm, R ∼ 20,000–100,000). CRIRES/VLT will allow the observation of high J H2 rotational transitions as well as a variety of ro-vibrational lines with both high spectral and spatial resolution. In addition, VISIR/VLT will permit at longer wavelengths observation of the 0-0 S(1) and S(3) H2 transitions with a spatial resolution of ∼0.3 and R ∼ 20,000 at 10 μm. Of course, the sensitivity of these ground-based instruments is limited by the atmosphere and thus the insight gained by such measurements (relative say to ISO) will be greatest for stellar and other compact sources of H2 emission. The potential of such high-resolution ground-based instruments has already been demonstrated by the results from the TEXES/IRTF spectrometer (5–25 μm, R ∼ 100, 000) observing the rotational transitions of H2 in disks (Richter et al., 2004) and the Orion Bar (Allers et al., 2004). Such ground-based instruments are hopefully the forerunners of high-class instrumentation in space or from airborne platforms. The spectrometers on SOFIA (covering 0.3–1600 μm) will soon observe several H2 lines including the 28 μm S(0) transition with diffraction limited resolution of 1.5 and R ∼ 104–105 at 20 μm. In the long term, MIRI on JWST (5–28 μm, ∼0.1 of spatial resolution) will allow high-class maps of the H2 rotational transitions and comparisons with the 2 μm 1-0 S(1) maps. The combination of these instruments will be very powerful and it will become a challenge for theorists to understand the images they produce. MOLECULAR HYDROGEN 89

References

Abergel, A., Bernard, J. P., Boulanger, F., Cesarsky, C., and Desert, F. X.: 1996, A&A 315, L329. Allers, K. N., Jaffe, D. T., Lacy, J. H., and Matthew, J. R.: 2004, astro-ph/0404077. An, J. H. and Sellgren, K.: 2001, Astronomy and Astrophysics Series, Ed. Pachart Publishing Hous Tucson 198, 5914. Armus, L., Charmandaris, V., Spoon, H. W. W., Houck, J. R., and Soifer, B. T.: 2004, Astronomy and Astrophysics Series, Ed. Pachart Publishing Hous Tucson 204, 3319. Bachiller, R., Cox, P., Josselin, E., Huggins, P. J., Forveille, T., Miville-Deschênes, M. A., and Boulanger, F.: 2000, ISO beyond the peaks, p. 171. Bernard Salas, J., Pottasch, S. R., Beintema, D. A., and Wesselius, P. R.: 2001, A&A 367, 949. Bertoldi, F.: 1997, First ISO Workshop on Analytical Spectroscopy,p.67. Biham, O., Furman, I., Pirronello, V., and Vidali, G.: 2001, ApJ 553, 595. Black, J. H. and van Dishoeck, E. F.: 1987, ApJ 322, 412. Boulanger, F., Rubio, M., Habart, E., and Verstraete, L.: 2004, in preparation. Cazaux, S. and Tielens, A. G. G. M.: 2002, ApJ 575, L29. Cazaux, S. and Tielens, A. G. G. M.: 2003, ApJ. Cesarsky, C. J., Abergel, A., Agnese, P., Altieri, B., and Augueres, J. L.: 1996, A&A 315, L32. Cesarsky, D., Cox, P., Pineau Des Forêts,G., van Dishoeck, E. F., Boulanger, F., and Wright, C. M.: 1999, A&A 348, 945. Combes, F.: 2000, H2 in Space, Cambridge University Press, p. 275. Cox, P., Maillard, J. P., Huggins, P. J., Forveille, T., and Simons, D.: 1997, A&A 321, 907. Cox, P., Boulanger, F., Huggins, P. J., Tielens, A. G. G. M., and Forveille, T.: 1998, ApJ 495, L23. Cox, P., Huggins, P. J., Maillard, J. P., Muthu, C., Bachiller, R., and Forveille, T.: 2003, ApJ 586, L87. Dalgarno, A.: 2000, H2 in Space, Cambridge University Press,p.3. de Graauw, T., Haser, L. N., Beintema, D. A., Roelfsema, P. R., and van Agthoven, H.: 1996, A&A 315, L49. Draine, B. T. and McKee, C. F.: 1993, ARA&A 31, 373. Draine, B. and Bertoldi, F.: 2000a, H2 in Space, Cambridge University Press, Cambridge, p. 131. Draine, B. T. and Bertoldi, F.: 1999b, Cox, P., and Kessler, M., (eds.), The Universe as seen by ISO, p. 553. Draine, B. T.: 1978, Astrophysical Journal Supplement Series 36, 595. Duley, W. W.: 1996, MNRAS 279, 591. Falgarone, E., Verstraete, L., Pineau Des Forêts, G., Flower, D., and Puget, J. L.: 2000, Rotational lines. H2 in Space, Cambridge University Press, Cambridge p. 37. Fischer, J., Smith, H. A., Geballe, T. R., Simon, M., and Storey, J. W. V.: 1987, ApJ 320, 667. Fuente, A., Martin-Pintado, J., and Rodriguez-Fernandez, N.: 1999, ApJ 518, L45. Genzel, R. and Cesarsky, C. J.: 2000, ARA&A 38, 761. Gould, R. J. and Salpeter, E. E.: 1963, ApJ 138, 393. Habart, E., Boulanger, F., Verstraete, L., Pineau Des Forêts,G., Falgarone, E., and Abergel, A.: 2003, A&A 397, 623. Habart, E., Boulanger, F., Verstraete, L., Walmsley, C. M., and Pineau Des Forêts,G.: 2004, A&A 414, 531. Herpin, F., Cernicharo, J., and Heras, A.: 2000, H2 in Space, Cambridge University Press, Cambridge, p. 24. Hollenbach, D. and Salpeter, E. E.: 1971, ApJ 163, 155. Hollenbach, D.: 1997, IAUS 182, 181. Hollenbach, D. J. and Tielens, A. G. G. M.: 1999, Reviews of Modern Physics 71, 173. Jansen, D. J., van Dishoeck, E. F., and Black, J. H.: 1994, A&A 282, 605. 90 E. HABART ET AL.

Jansen, D. J., van Dishoeck, E. F., Black, J. H., Spaans, M., and Sosin, C.: 1995, A&A 302, 223. Joblin, C., Maillard, J. P., De Peslouan, P., Vauglin, I., Pech, C., and Boissel, P.: 2000, H2 in Space, Cambridge University Press, Cambridge, p. 107. Joblin, C., Boissel, P., Pech, C., Armengaud, M., and Frabel, P.: 2001, Infrared and Submillimeter Space Astronomy, EDP Sciences. Jura, M.: 1975, ApJ 197, 575. Katz, N., Furman, I., Biham, O., Pirronello, V., and Vidali, G.: 1999, ApJ 522, 305. Kemper, C., Spaans, M., Jansen, D. J., Hogerheijde, M. R., van Dishoeck, E. F., and Tielens, A. G. G. M.: 1999, ApJ 515, 649. Le Bourlot, J., Pineau des Forêts,G., and Flower, D. R.: 1999, MNRAS 305, L802. Le Bourlot, J.: 2000, A&A 360, L656. Le Coupanec, P.: 1998, Ph.D thesis. University of Paris 6. Lecavelier des Etangs, A., Deleuil, M., Vidal-Madjar, A., Roberge, A., and Le Petit, F.: 2003, A&A 407, 935. Lefloch, B., Cernicharo, J., Cabrit, S., Noriega-Crespo, A., Moro-Martin, A., and Cesarsky, D.: 2003, ApJ 590, L41. Li, W., Evans, N. J. I., Jaffe, D. T., van Dishoeck, E. F., and Thi, W. F.: 2002, ApJ 568, 242. Lutz, D., Sturm, E., Genzel, R., Spoon, H. W. W., Moorwood, A. F. M., Netzer, H., and Sternberg, A.: 2003, A&A 409, 867. Maloney, P. R., Hollenbach, D. J., and Tielens, A. G. G. M.: 1996, ApJ 466, 561. Morris, P., Crowther, P., and Houck, J.: 2004, Astronomy and Astrophysics Series, Ed. Pachart Pub- lishing Hous Tucson in press. Moutou, C., Verstraete, L., Sellgren, K., and Léger,A.: 1999, The Universe as Seen by ISO, p. 727. Natta, A. and Hollenbach, D.: 1998, A&A 337, 517. Neufeld, D. A., Melnick, G. J., and Harwit, M.: 1998, ApJ 506, L75. Nisini, B., Benedettini, M., Giannini, T., Codella, C., Lorenzetti, D., Di Giorgio, A. M., and Richer, J. S.: 2000, A&A 360, 297. Parmar, P. S., Lacy, J. H., and Achtermann, J. M.: 1991, ApJ 372, L25. Parneix, P. and Brechignac, P.: 1998, A&A 334, 363. Pineau Des Forêts,G. and Flower, D.: 2000, H2 in Space, Cambridge University Press, Cambridge, p. 18. Pirronello, V., Liu, C., Shen, L., and Vidali, G.: 1997, ApJ 475, L69. Pirronello, V., Liu, C., Roser, J. E., and Vidali, G.: 1999, A&A 344, 681. Rachford, B. L., Snow, T. P., Tumlinson, J., Shull, J. M., and Blair, W. P.: 2002, ApJ 577, 221. Rappacioli, M., Joblin, C., and Boissel, P.: 2004, A&A, (submitted). Rho, J. and Reach, W.: 2003, Rev. Mex. AA, p. 263. Richter, M. J., Jaffe, D. T., Blake, G. A., and Lacy, J. H.: 2002, ApJ 572, L161. Richter, M. J., Lacy, J. H., Greathouse, T. K., Jaffe, D. T., and Blake, G. A.: 2004, astro-ph/0403349. Rigopoulou, D., Kunze, D., Lutz, D., Genzel, R., and Moorwood, A. F. M.: 2002, A&A 389, 374. Rodriguez-Fernandez, N. J., Martin-Pintado, J., de Vicente, P., Fuente, A., Huttemeister, S., Wilson, T. L., and Kunze, D.: 2000, A&A 356, 695. Rodriguez-Fernandez, N. J., Martin-Pintado, J., Fuente, A., De Vicente, P., Wilson, T. L., and Huttemeister, S.: 2001, A&A 365, 174. Rosenthal, D., Bertoldi, F., and Drapatz, S.: 2000, A&A 356, 705. Sandford, S. A. and Allamandola, L. J.: 1993, ApJ 409, L65. Shull, J. M. and Beckwith, S.: 1982, ARA&A 20, 163. Sidis, V., Jeloica, L., Borisov, A., and Deutscher, S.: 2000, H2 in Space, Cambridge University Press, Cambridge, p. 89. Smith, J. D. T., Dale, D. A., Draine, B. T., Hollenbach, D. J., and Armus, L.: 2004, Astronomy and Astrophysics Series, Ed. Pachart Publishing Hous Tucson 204, 8203. MOLECULAR HYDROGEN 91

Spitzer, L., Cochran, W. D., and Hirshfeld, A.: 1974, Astrophysical Journal Supplement Series 28, 373. Sternberg, A. and Neufeld, D. A.: 1999, ApJ 516, 371. Sternberg, A.: 1989, ApJ 347, 863. Takahashi, J., Masuda, K., and Nagaoka, M.: 1999, ApJ 520, 724. Thi, W. F., van Dishoeck, E. F., Black, J. H., Jansen, D. J., Evans, N. J., and Jaffe, D. T.: 1999, The Universe as Seen by ISO, p. 529. Thi, W. F., Van Dishoeck, E. F., Blake, G. A., Van Zadelhoff, G. J., and Horn, J.: 2001, ApJ 561, 1074. Timmermann, R., Bertoldi, F., Wright, C. M., Drapatz, S., Draine, B. T., Haser, L., and Sternberg, A.: 1996, A&A 315, L281. Tine, S., Lepp, S., Gredel, R., and Dalgarno, A.: 1997, ApJ 481, 282. Valentijn, E. A. and van der Werf, P. P.: 1999, ApJ 522, L29. Valentijn, E. A., van der Werf, P. P., De Graauw, T., and De Jong, T.: 1996, A&A 315, L145. van den Ancker, M., Wesselius, P. R., and Tielens, A. G. G. M.: 2000, H2 in Space, Cambridge University Press, Cambridge, p. 21. van den Ancker, M. E., Wesselius, P. R., and Tielens, A. G. G. M.: 2000, A&A 355, 194. van der Werf, P. P., Genzel, R., Krabbe, A., Blietz, M., Lutz, D., Drapatz, S., Ward, M. J., and Forbes, D. A.: 1993, ApJ 405, 522. van Dishoeck, E. F.: 2004, ARA&A, astro-ph/0403061. Wilgenbus, D., Cabrit, S., Pineau Des Forˆets,G., and Flower, D. R.: 2000, A&A 356, 1010. Williams, D., Clary, D. C., Farebrother, A., Fisher, A., and Gingell, J.: 2000, H2 in Space,p.13. Wright, C. M., Drapatz, S., Timmermann, R., van der Werf, P. P., Katterloher, R., and De Graauw, T.: 1996, A&A 315, L301. Wright, C. M., Timmermann, R., and Drapatz, S.: 1997, Proceedings of the First ISO Workshop on Analytical Spectroscopy, p. 311. Wright, C., Drapatz, S., Timmermann, R., van Dishoeck, E., and Bertoldi, F.: 2000, H2 in Space, Cambridge University Press,Cambridge, p. 29. UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS

DAVID ELBAZ CEA Saclay/DSM/DAPNIA/Service d’Astrophysique, Orme des Merisiers, F-91191 Gif-sur-Yvette Cedex, France (E-mail: [email protected])

(Received 26 July 2004; Accepted in ﬁnal form 12 September 2004)

Abstract. We present the results obtained through the various ISO extragalactic deep surveys. Al- though IRAS revealed the existence of galaxies forming stars at a rate of a few tens (LIRGs) or even hundreds (ULIRGs) solar masses in the local universe, ISO not only discovered that these galaxies were already in place at redshift one, but also that they are not the extreme objects that we once believed them to be. Instead they appear to play a dominant role in shaping present-day galaxies as reﬂected by their role in the cosmic history of star formation and in producing the cosmic infrared background detected by the COBE satellite in the far infrared to sub-millimeter range.

Keywords: deep surveys, infrared, galaxy formation

1. General Perspective

It is widely accepted that in the local universe stars form giant molecular clouds (GMCs) where their optical and mostly UV light is strongly absorbed by the dust which surrounds them. Whether extinction was already taking place in the more distant universe where galaxies are less metal rich was less obvious ten years ago. −1 Galaxies forming stars at a rate larger than about 20 M yr were known to radiate the bulk of their luminosity above 5 μm thanks to IRAS, the so-called luminous (LIRGs, 12 > log(LIR /L ) ≥ 11) and ultra-luminous (ULIRGs, log(LIR /L ) ≥ 12) infrared (IR) galaxies. In the past, when galaxies were more gaseous and formed the bulk of their present-day stars, it would have been logical to expect and detect the past star formation events of galaxies in the IR regime and to detect a large population of LIRGs/ULIRGs too. However, prior to the launch of the Infrared Space Observatory (ISO, Kessler et al., 1996), this idea was not widely spread. Partly because of a cultural reason: star formation rates (SFR) were commonly measured from optical emission lines and rest-frame UV light in galaxies. This may explain why the redshift evolution of the SFR density per unit co-moving volume computed by Madau et al. (1996) became famous. However, in the first presentation of the density of UV light per unit co-moving volume (Lilly et al., 1996) the authors were cautious to avoid converting their UV light into a SFR because of the unknown Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, The Netherlands, and the United Kingdom), and with the participation of ISAS and NASA. Space Science Reviews (2005) 119: 93–119 DOI: 10.1007/s11214-005-8060-3 C Springer 2005 94 D. ELBAZ factor to correct for extinction. Already IRAS observations indicated a rapid decline of the co-moving number density of ULIRGs since z ∼ 0.3 (Kim and Sanders, 1998, see also Oliver et al., 1996), but this was over a small redshift range and with small number statistics. In the few years that followed the launch of ISO, several observations showed that galaxy formation could not be understood, at least on an observational basis, without accounting for dust extinction as a major ingredient. The ISO deep surveys played a major role in this process, together with other results summarized below. They clearly established that extreme events such as those taking place in local LIRGs and ULIRGs must have been more common in the past, so much that they can now be considered as a standard phase that most galaxies experienced during their lifetime, at least once, but maybe even several times. The first result of the ISOCAM surveys, as well as the ISOPHOT ones, was the great difference of the counts measured at faint flux densities with respect to local ones from IRAS (Elbaz et al., 1999; Dole et al., 2001). The universe must have been much richer in IR luminous galaxies in the past, either because galaxies were more IR luminous, at fixed galaxy density, and/or because the number density of galaxies was larger in the past, which was partly expected due to the reverse effect of hierarchical galaxy formation through mergers. The strength of the excess of faint objects came as a surprise, but its consequences on the past star formation history of galaxies was confirmed by the convergence of other observations going in the same direction. – The nearly simultaneous discovery of the cosmic infrared background (CIRB, Puget et al., 1996; Fixsen et al., 1998; Hauser and Dwek, 2001 and references therein), at least as strong as the UV-optical-near IR one, whereas local galaxies only radiate about 30% of their bolometric luminosity in the IR above λ ∼ 5 μm. – The 850 μm number counts from the SCUBA sub-millimeter bolometer array at the JCMT (Hughes et al., 1998; Barger et al., 1998; Smail et al., 2002; Chapman et al., 2003, and references therein) which also indicate a strong excess of faint objects in this wavelength range, implying that even at large redshifts dust emission must have been very large in at least the most active galaxies. – The most distant galaxies, individually detected thanks to the photometric redshift technique using their Balmer or Lyman break signature showed the signature of a strong dust extinction. The so-called “β-slope” technique (Meurer et al., 1999) used to derive the intrinsic luminosity of these galaxies and correct their UV luminosity by factors of a few (typically between 3 and 7, Steidel et al., 1999; Adelberger and Steidel, 2000) was later on shown to even underestimate the SFR of LIRGs/ULIRGs (Goldader et al., 2002). – The slope of the sub-mJy deep radio surveys (Haarsma et al., 2000). It has now become clear that the cosmic history of star formation based on rest-frame UV or emission line indicators of star formation such as [0II] or [Hα] UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 95 strongly underestimate the true activity of galaxies in the past if not corrected by strong factors due to dust extinction. Although distant galaxies were less metal rich and much younger, they must have found the time to produce dust rapidly in order to efficiently absorb the UV light of their young stars. We have tried to summarize in the following the role played by the ISO deep surveys in establishing this new perspective (see also Genzel and Cesarsky, 2000). However, 10 years after ISO’s launch we are still trying to understand the consequences of these findings on galaxy formation scenarios. Are these distant LIRGs really similar to local ones? What do they teach us about the connection between star and galaxy formation on one side (IMF, triggering of star formation, conditions of star formation, ...) and between galaxy and large-scale structure formation on the other side (role of the environment in triggering star formation events, galaxy versus group and cluster formation, ellipticals vs. spirals, ...)? How much energy radiated by Compton thick embedded active nuclei remains undetected even by the Chandra and XMM-Newton X-ray observatories? These questions together with others that will be discussed in the following demonstrate the liveliness of this field that will continue to feed the next generation of telescopes and instruments to come such as Herschel, ALMA, the James Webb Space Telescope (JWST) or the Spitzer and GALEX space observatories presently in use.

2. The ISO Surveys

Extragalactic deep surveys with ISO were performed in the mid- and far-IR with ISOCAM (Cesarsky et al., 1996) and ISOPHOT (Lemke et al., 1996), respectively. In both wavelength ranges, the steep slopes of the number counts indicate that a rapid decline of the IR emission of galaxies must have taken place from around z ∼1toz = 0. As shown in Figure 1, ISOCAM could detect galaxies at 15 μm in the LIRG regime up to z ∼1.3, while ISOPHOT was limited to either nearby galaxies 12 or moderately distant ULIRGs such as the one of LIR ∼ 4 × 10 L at a redshift of z = 1 shown in the Figure 1 (normalized SED of Arp 220). The detection limits of ISOCAM and ISOPHOT are compared to SCUBA and Spitzer in the Figure 2. The deepest ISO surveys reached a completeness limit of 0.1 mJy at 15 μm (plain line) and a depth of S15 ∼ 40 μJy (incomplete, Aussel et al., 1999) in blank fields or a twice deeper completeness level in the central part of nearby galaxy clusters using lensing magnification (Altieri et al., 1999; Metcalfe et al., 2003), and 120 mJy at 170 μm (dotted line). The right axis on the plot shows the conversion of the IR detection limit into a SFR detection limit. Any galaxy forming stars at a rate −1 larger than 30M yr could be detected with ISOCAM up to z ∼ 1, assuming that the measured 15 μm flux density of a galaxy can be used to derive its “bolometric” IR luminosity, i.e. LIR = L(8–1000 μm ) (see Section 5.2). Number counts represent the first scientific result of an extragalactic survey. They can be used to constrain the models. Used alone, they leave some degeneracies 96 D. ELBAZ

Figure 1. Effect of the k-correction on the detection of distant LIRGs and ULIRGs by ISOCAM, ISOPHOT and SCUBA. Only ISOCAM at 15 μm can detect LIRGs up to z ∼ 1, while ISOPHOT and SCUBA are sensitive to distant ULIRGs of a few 1012 L .

Figure 2. Sensitivity limits of ISOCAM (15μm, 0.1 mJy), ISOPHOT (170 μm, 120 mJy), SCUBA (850 μm, 2 mJy) and Spitzer (24 μm, 0.45 mJy and 0.022 mJy corresponding to the expected detection limits of the SWIRE and GOODS Legacy Programs). This ﬁgure was generated assuming that distant galaxies SEDs are similar to local ones. We used the library of template SEDs constructed by Chary and Elbaz (2001). unsolved but at least they can demonstrate whether the distant universe was different from the local one in this wavelength range by comparing them to “no evolution” predictions assuming some spectral energy distributions (SEDs) for the k-correction. UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 97

3. The ISOCAM 6.75-μm Deep Surveys

Deep images of blank fields at 6.75- and 15-μm done with ISOCAM provide a different view on the distant universe. While the ISOCAM-LW3 filter, centered at 15μm, can probe dust emission up to z ∼ 2 for luminous objects, the redshift range to probe star formation from dust emission with the ISOCAM-LW2 band, centered at 6.75μm, is limited to the relatively nearby universe due to k-correction. However, the emission of the old stellar component, which peaks in the near IR, being brought to this wavelength range for high-z galaxies, their stellar masses are better constrained from this flux density. 11 Sato et al. (2004) derived stellar masses around M ∼ 10 M for galaxies with z ∼ 0.2–3 from the correlation between rest-frame near IR luminosity and stellar mass. The stellar mass-to-light ratios were derived from model fit of the set of observed magnitudes depending on the galaxies star formation histories. The contribution of the 6.75-μm selected galaxies to the cosmic stellar mass density of galaxies per unit co-moving volume as a function of redshift were estimated to be comparable to those inferred from observations of UV bright galaxies (see Figure 3). Given the narrow mass ranges, these estimates were obtained as simple summations of the detected sources and should therefore be considered as lower limits only. The faint 6.7 μm galaxies generally had red colour. A comparison with a particular population synthesis model suggests that they have experienced vigorous star formation at high redshifts. The derived large stellar masses for the faint 6.7 μm galaxies also support such star forming events at the past. However, these events can only be detected at longer wavelengths using the 15- μm ISOCAM or 90 and 170-μm ISOPHOT bands with ISO (or SCUBA at 850 μm and now the Spitzer Space Telescope). The next sections are devoted to these surveys and their consequences on our understanding of the history of star formation in the universe.

4. The ISOPHOT Far IR Deep Surveys

ISOPHOT was used to survey the sky at 90 and 170μm, mainly (plus a few surveys at 60, 120, 150 and 180 μm, see Juvela et al., 2000; Linden-Voernle et al., 2000; cf. review by Dole, 2002). At 90 μm, the 46 arcsec pixels and FWHM of ISOPHOT represented a signiﬁcant improvement with respect to IRAS, however due to k- correction and sensitivity limitations, only the local universe could be probed at such wavelengths. The 170-μm band was more favourable for the detection of distant ULIRGs as shown in the Figure 1 where the peak emission of a ULIRG of 4 × 1012 L and redshifted at z = 1 is shown (normalized SED of Arp 220). A major issue at these wavelengths with a 60-cm telescope is obviously the identiﬁcation of optical counterparts for the determination of a redshift and the separation of local and distant galaxies. 98 D. ELBAZ

Figure 3. (Left) ISOCAM 6.75-μm image of the SSA13 field (Figure 2 from Sato et al., 2003, 23 h observation). This image reaches an 80% completeness limit of 16 μJy in the central 7 arcmin2 region. North is top and east to the left in J2000. The center is (RA, Dec) = (13h12m26s, 42◦4424.8). The map shows signal-to-noise ratio (S/N) per 0.6 arcsec sub-pixel. The signal is an average, weighted by the inverse of the assigned variance, and the noise is a normalized standard deviation. (Right) Stellar mass density in the universe as a function of redshift (Figure 13 from Sato et al., 2004). The right axis shows densities normalized to the critical density of the universe. The contributions of the 6.75-μm galaxies are shown with solid and double circles for the “combined” (including photometric redshifts) and purely spectroscopic samples, respectively. The double circles are plotted at slightly lower redshifts. The horizontal bars represent the redshift ranges of the bins and the vertical bars show one sigma errors, taking account of Poisson noise and uncertainties in stellar mass and Vmax. Several other estimates are overlaid (see Sato et al., 2004 for references). The empty/solid squares and the diamond are obtained from full integration of a Schechter fit to their respective luminosity or stellar mass function at each redshift bin. The X marks and the empty circles are quasi-fully integrated ∗ . < −2 < . values with a finite integration range from 10 to 1/20 L , and 10 5 log(Mstar[h65 M]) 11 6, respectively. The triangles are simply summed values of the detected sources. The two dashed curves are deduced by integrating the star formation rate density in the universe, which is derived from the UV luminosity density as a function of redshift (Cole et al., 2001). The upper curve is an extinction corrected case for E(B − V ) = 0.15, and the lower one has no dust correction.

4.1. SOURCE COUNTS AND COSMIC IR BACKGROUND

Source counts at 170 μm (e.g. Kawara et al., 1998; Puget et al., 1999; Dole et al., 2001) exhibit a steep slope of α = 3.3 ± 0.6 between 180 and 500 mJy and, like in the mid-IR range, show sources in excess by a factor of 10 compared with no- evolution scenario. The brightness fluctuations in the Lockman Hole were used by Matsuhara et al. (2000) to constrain galaxy number counts down to 35 mJy at 90 μm and 60 mJy at 170μm, confirming the existence of a strong evolution down to these flux densities. Using a new data reduction method, (Rodighiero and Franceschini, 2004) extended the previous works of Kawara et al. (1998), Efstathiou et al. (2000) and Linden-Voernle et al. (2000) down to lower flux densities (30 mJy at 90 μm ) and found a clear excess of faint objects with respect to no evolution (see Figure 4a). However, the resolved sources account for less than 10% of the Cosmic UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 99

Figure 4. (a) Differential 90 μm counts dN/dS normalized to the Euclidean law (N ∝ S−2.5)(extracted from Rodighiero and Franceschini, 2004). Results compared with those from the preliminary analysis of the ISOPHOT ELAIS survey (Efsthatiou et al., 2000, open circles) and with those from Kawara et al. (2004, filled triangles). The long-dashed line shows the expected contribution of non- evolving spirals as in the model of Franceschini et al. (2001). (b) Fluctuations of the CIB in a power spectrum analysis of the FIRBACK/ELAIS N2 field at 170 μm by Puget and Lagache (2000). Ob- served power spectrum: diamond; straight continuous line: the best fit cirrus power spectrum; dash line: cirrus power spectrum deduced from Miville-Deschênes et al. (2002); continuous curve: detector noise.

Infrared Background at 170 μm, which is expected to be resolved into sources in the 1–10 mJy range. Sources below the detection limit of a survey create fluctuations. If the detection limit does not allow to resolve the sources dominating the CIB intensity, which is the case in the far IR with ISO, characterizing these fluctuations can constrain the spatial correlations of these unresolved sources of cosmological significance. An example of the modeled redshift distribution of the unresolved sources at 170 μm can be found in Figure 12 of Lagache et al. (2003); the sources dominating the CIB fluctuations have a redshift distribution peaking at z ∼ 0.9. After the pioneering work of Herbstmeierer et al. (1998) with ISOPHOT, Lagache and Puget (2000) discovered them at 170 μm in the FIRBACK data, followed by other works at 170 and 90 μm (Matsuhara et al., 2000; Puget and Lagache, 2000; Kiss et al., 2001). Figure 4b shows the CIB fluctuations in the FN2 field by Puget and Lagache (2000), at wavenumbers 0.07 < k < 0.4 arcmin−1.

4.2. NATURE OF THE ISOPHOT GALAXIES

Determining the nature of the far IR galaxies has been a longer process than in the mid-IR, mainly because of the difficulty to find the shorter wavelength counterparts in a large beam. Various techniques have been used to overcome this problem, one of the most successful being the identification using 20-cm radio data (e.g. Ciliegi et al., 1999). Another technique is the far IR multiwavelength approach (Juvela et al., 2000) that helps constraining the position and the SED; it also helps separate 100 D. ELBAZ the cirrus structures from the extragalactic sources. A variation is to use ISOCAM and ISOPHOT data, like the ELAIS Survey (Rowan-Robinson et al., 2004, see also Oliver in this book). Finally, the Serendipity Survey (Stickel et al., 1998, 2000), by covering large and shallow areas, allows to detect many bright objects easier to follow-up or already known. Far-IR ISO galaxies can be sorted schematically into two populations. First, the low-redshift sources, typically z < 0.3 (e.g., Serjeant et al., 2001; Patris et al., 11 2002; Kakazu et al., 2002), have moderate IR luminosities, below 10 L, and are cold (Stickel et al., 2000). Second, sources at higher redshift, z ∼ 0.3 (Patris et al., 2002) and beyond, z ∼ 0.9 (Chapman et al., 2002) are more luminous, typically 11 L > 10 L, and appear to be cold. Serjeant et al. (2001) derived the Luminosity Function at 90 μm, and started to detect an evolution compared to the local IRAS 100 μm sample.

5. The ISOCAM 15-μm Deep Surveys

A series of surveys were performed within the Guaranteed Time (IGTES, ISOCAM Guaranteed Time Extragalactic Surveys, Elbaz et al., 1999) and Open Time as summarized in the Table I, where the surveys at 7 (6.75 μm ) and 15 μm are sorted by increasing depth irrespective of wavelength. The major strength of ISOCAM is its spatial resolution (PSF FWHM of 4.5 at 15μm, Okumura, 1998) and sensitivity, which permitted to detect galaxies down to the LIRG regime up to z ∼1.3 in the deepest surveys (Figures 1 and 2), well above confusion. Cosmic rays were a stronger limitation than photon or readout noise by inducing ghost sources when they were not perfectly removed, especially those with long-term transients associated to them. Two techniques were developed to solve this issue, the so-called PRETI (Pattern REcognition Technique for Isocam data, Starck et al., 1999) and LARI (Lari et al., 2001) techniques. A third technique was developed by Désert et al. (1999), in which the mosaic data were analyzed using the beam switching approach, the Three Beam technique. PRETI consists in a multi-scale wavelet decomposition of the signal, while LARI tries to account for the physical processes taking place in the detectors, including the effect of neighboring pixels. The LARI technique was first applied to the ELAIS surveys (see Oliver et al., in these book) and more recently to the IGTES surveys of the Lockman Hole (Rodighiero et al., 2004; Fadda et al., 2004) and of the Marano field (Elbaz et al., in preparation). The quality of the ISOCAM images is shown with the case of the Marano FIRBACK deep field (30×30) in the Figure 6. The central part of this field was imaged at a deeper level (UDSF) and analyzed with both techniques. The resulting contours are overlayed on a VLT-FORS2 image of the field, demonstrating the consistency and robustness of both source detection algorithms (Figure 7). Lensing appeared to be a very powerful tool to extend the detection of faint objects by a factor 2–3 (see Metcalfe et al., 2003). However, studies such as the UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 101

TABLE I Table of the ISOCAM extragalactic surveys sorted by increasing depth.

Name λ (μm) Area (’2) Int. (min) depth (mJy) No. of objects

ELAIS N11 15 9612 0.7;0.7 1;0.7 490 ELAIS N21 7;15 9612;9612 0.7;0.7 1;0.7 628;566 ELAIS N31 7;15 4752;3168 0.7;0.7 1;0.7 189;131 ELAIS S11 7;15 6336;14256 0.7;0.7 1;0.7 304;317 ELAIS S21 7;15 432;432 0.7;0.7 1;0.7 40;43 Lockman Shallow2,3 15 1944 3 0.25 457 (80% compl.) 0.45 260 Comet Field4 12 360 10 0.5 37 CFRS14+525 7;15 100;100 18;11 0.3;0.4 23;41 CFRS03+006 7;15 100;100 6;22 0.5;0.3 – Lockman Deep2,7 7;15 500;500 18;11 0.3;0.4 166 Marano DSF2 7;15 900;900 15.4;15.4 0.19;0.32 180 A3708 7;15 31.3;31.3 42;42 0.052;0.21(u) 4;20 (80% compl.) 0.080;0.293(u) Marano UDSR2 7;15 85;90 120;114 0.18;0.14 –;142 Marano UDSF2 7;15 89;90 114 0.08;0.14 115;137 A22188 7;15 20.5;20.5 84;84 0.054;0.121(u) 18;46 (80% compl.) 0.079;0.167(u) HDFN+FF2,9 7;15 10;27 116;135 0.05;0.1 7;44 HDFS2,10 7;15 28;28 168;168 0.05;0.1 16;63 A239011,8 7;15 5.3;5.3 432;432 0.038;0.050(u) 10;28 (80% compl.) 0.052;0.092(u) Lockman PGPQ12 7 9 744 0.034 15 SSA1313 7 16 1264 0.006 65 7(80% compl.) 0.016

Notes. Col. (1) Survey name. Col. (2) wavelength of the survey with “7” for the LW2 filter centered at 6.75 μm and covering 5–8.5 μm, and “15” for the LW3 filter centered 15 μm and covering 12–18 μm. Col. (3) total area in square arcmin. Col. (4) integration time per sky position. Col. (5) depth (and 80% completeness limit when indicated). Col. (6) number of objects detected above this depth. References: (1) Oliver et al., 2000; Rowan-Robinson et al., 2004, (2) Elbaz et al., 1999 and in preparation, (3) Rodighiero et al., 2004, (4) Clements et al., 1999, (5) Flores et al., 1999, (6) Flores (private communication), (7) Fadda et al., 2004, (8) Metcalfe et al., 2003, (9) Aussel et al., 1999; Goldschmidt et al., 1997, (10) Oliver et al., 2002, (11) Altieri et al., 1999, (12) Taniguchi et al., 1997, (13) Sato et al., 2003, 2004. Superscript (u)indcates that the depth of the deep surveys in the field of lensing clusters does not include the correction for lensing amplification. relative clustering of IR galaxies versus field galaxies are better done in larger homogeneous fields, where the role of cosmic variance can also be quantified. As we will see below, the environment of distant LIRGs seems to play a major role in triggering their star formation activity. 102 D. ELBAZ

Figure 5. ISOCAM 15- μm differential counts, with 68% error bars. The counts are normalized to a Euclidean distribution of non-evolving sources, which would have a slope of α =−2.5 in such a universe. This plot is an extension of the Figure 2 from Elbaz et al. (1999) combining the data points from Elbaz et al. (1999), for the ISOCAM Guaranteed Time Extragalactic Surveys (IGTES, filled and empty dots for sources below and above 1 mJy respectively, see text), Gruppioni et al. (2003, bold empty circles) for the ELAIS-S1 field, Rodighiero et al. (2004, large open stars) and Fadda et al. (2004, large open squares), for the Shallow and Deep Surveys of the Lockman Hole, respectively (analysis of the IGTES data using the “Lari” technique, see Section 5). The hatched area materializes the range of possible expectations from models assuming no evolution and normalized to the 15 μm local luminosity function (LLF) from Fang et al. (1998) and the template SED of M51 (with an assumed uncertainty of ±20%; this 15 μm LLF was recently confirmed over a larger sample of local galaxies by H. Aussel, private communication).

5.1. SOURCE COUNTS AND COSMIC IR BACKGROUND

The source counts at 15 μm exhibit a strong excess of faint sources below S15 ∼ 2 mJy. This excess is usually defined by comparison with model predictions assuming that galaxies behaved similarly in the distant universe as they do today. Such “no evolution” behavior is represented by a shaded area in the Figure 5 (see figure caption). Galaxies above this flux density do fall within this region, as illustrated by the data points from the ELAIS-S1 field (Gruppioni et al., 2003, see also Oliver et al., in this book). In Figure 5, we have separated the data points from the IGTES (Elbaz et al., 1999) between those below and above S15 = 1 mJy, with filled and open dots, respectively. Data above this flux density from Elbaz et al. (1999) appear to be inconsistent with those derived from ELAIS-S1. Most of those points were derived from the Shallow Survey of the Lockman Hole within the IGTES, which suffered from having less redundant observations of a given sky pixel. At that time, ISOCAM data reduction methods were not optimized for such surveys, but since then, they have been improved to better deal with such shallow surveys. Among the techniques that are discussed below, the “Lari” technique is particularly suitable for UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 103

Figure 6. ISOCAM 15-μm image of the Marano FIRBACK Deep Survey (DSF) from the IGTES (Elbaz et al., 1999, and in preparation). Data reduction with PRETI. such low redundancy surveys (and even more for the very shallow ELAIS surveys, see Oliver et al., in this book) and a recent analysis of the Lockman Hole Deep (Fadda et al., 2004, large open squares) and Shallow (Rodighiero et al., 2004, large open stars) surveys from the IGTES provided new number counts at these flux densities perfectly consistent with those derived from ELAIS-S1 by Gruppioni et al. (2003) also using the same “Lari” technique. Note that the models designed to fit the ISOCAM number counts were constrained by the Elbaz et al. (1999) results, hence overproduce the number of sources above S15 ∼ 2 mJy. As a natural result, they have also overpredicted the number of sources detected in the high flux density regime at 24 μm with Spitzer (see Section 7 and Papovich et al., 2004; Chary et al., 2004). Above the Earth’s atmosphere, the 15-μm light is strongly dominated by the zodiacal emission from interplanetary dust and it has not yet been possible to make a direct measurement of the 15-μm background, or EBL. Individual galaxies contribute to this background and a lower limit to the 15-μm EBL can be obtained by adding up the fluxes of all ISOCAM galaxies detected per unit area down to a given flux limit. The resulting value is called the 15-μm integrated galaxy light (IGL). As in Elbaz et al. (2002), the differential number counts can be converted into a differential contribution to the 15 μm IGL as a function of flux density, estimated 104 D. ELBAZ

Figure 7. ISOCAM 15-μm contours overlayed on the VLT-FORS2 image (7×7, R-band) of the Marano FIRBACK Ultra Deep Survey (UDSF). The LARI (light contours, yellow on screen, grey on paper) and PRETI (dark contours, red on screen) detect the same objects. from the following formula: dIGL dN S = × 15 × ν (1) dS dS 1020 15

−1 where dN (sr ) is the surface density of sources with a flux density Sν[15 μm]= −20 −2 −1 S15 (mJy) over a flux density bin dS (mJy) (1 mJy = 10 nW m Hz ) and ν15 (Hz) is the frequency of the 15 μm photons. Below S15 ∼ 3 mJy, about 600 galaxies were used to produce the points with errors bars in Figure 8a. Figure 8b shows the 15-μm IGL as a function of depth. It corresponds to the integral of Figure 8a, where the data below 3 mJy were fitted with a polynomial of degree 3 and the 1-σ error bars on dIGL/dS were obtained from the polynomial fit to the upper and lower limits of the data points. The 15-μm IGL does not converge above a sensitivity limit of S15 ∼ 50 μJy, but the flattening of the curve below S15 ∼ 0.4 mJy suggests that most of the 15 μm EBL should arise from the galaxies already unveiled by ISOCAM. Franceschini et al. (2001) and Chary and Elbaz (2001), developed models which reproduce the number counts from ISOCAM at 15 μm, from ISOPHOT at 90 and 170 μm and from SCUBA at 850 μm, as well as the shape of the CIRB from 100 UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 105

(a) (b)

Figure 8. (a) Differential contribution to the 15 μm Integrated Galaxy Light as a function of flux density and AB magnitude. The plain line is a fit to the data: Abell 2390 (Altieri et al., 1999), the ISOCAM Guaranteed Time Extragalactic Surveys (IGTES, Elbaz et al., 1999), the European Large Area Infrared Survey (ELAIS, Serjeant et al., 2000) and the IRAS all sky survey (Rush et al., 1993). (b) Contribution of ISOCAM galaxies to the 15 μm extragalactic background light (EBL), i.e. 15 μm Integrated Galaxy Light (IGL), as a function of sensitivity or AB magnitude (AB = −2.5 log(SmJy) + 16.4). The plain line is the integral of the fit to dIGL/dS (a). The dashed lines correspond to 1-σ error bars obtained by fitting the 1-σ upper and lower limits of dIGL/dS. to 1000 μm. These models consistently predict a 15-μm EBL of: EBLmodels (15 μm) ∼ 3.3nWm−2 sr−1 (2) If this prediction from the models is correct then about 73 ± 15% of the 15 μm EBL is resolved into individual galaxies by the ISOCAM surveys. This result is consistent with the upper limit on the 15 μm EBL estimated by Stanev and Franceschini (1998) of EBLmax (15 μm) ∼ 5nWm−2 sr−1 (3) This upper limit was computed from the 1997 γ -ray outburst of the blazar Mkn 501 (z = 0.034) as a result of the opacity of mid-IR photons to γ -ray photons, which annihilate with them through e+e− pair production. It was since confirmed by Renault et al. (2001), who found an upper limit of 4.7 nW m−2 sr−1 from 5 to 15 μm. Nearly all the IGL is produced by sources fainter than 3 mJy (94 %) and about 70 % by sources fainter than 0.5 mJy. This means that the nature and redshift distribution of the galaxies producing the bulk of the 15 μm IGL can be determined by studying these faint galaxies only.

5.2. MID-IR AS A SFR INDICATOR

When normalized to the 7.7- μm PAH (polycyclic aromatic hydrocarbon) broad emission line, the spectra of different galaxies exhibit very different 15, 25, 60 or 106 D. ELBAZ

(a) (b)

Figure 9. IR luminosity correlations for local galaxies (from Elbaz et al., 2002). (a) ISOCAM- LW3 (15 μm) versus ISOCAM-LW2 (6.75 μm) luminosities (νLν ) (56 galaxies). (b) ISOCAM- LW3 (15 μm) versus IRAS-12 μm luminosities (45 galaxies). (c) LIR[8–1000 μm] versus ISOCAM- LW3 (15 μm) luminosity (120 galaxies). (d) LIR[8–1000 μm] versus LW2-6.75 μm luminosities (91 galaxies). Filled dots: galaxies from the ISOCAM guaranteed time (47 galaxies including the open squares). Open dots: 40 galaxies from Rigopoulou et al. (1999). Empty triangles: four galaxies from 10 Tran et al. (2001). Galaxies below LIR ∼ 10 L present a ﬂatter slope and have LIR/L B < 1.

100 μm over 7.7 μm ratio. This was often interpreted as an indication that measuring the monochromatic luminosity of a galaxy at one mid-IR wavelength was useless to determine its bolometric IR luminosity, LIR = L(8–1000 μm ). However, the variation of the far over mid-IR ratio is correlated with LIR and local galaxies do exhibit a strong correlation of their mid- and far-IR luminosities (Figure 9). These correlations can be used to construct a family of template SEDs or correlations from which the LIR, and therefore SFR, of a galaxy can be derived from its mid-IR luminosity (Chary and Elbaz, 2001; Elbaz et al., 2002). The LIR derived from this technique are consistent with those derived by the radio-far IR correlation, when radio-mid–far IR data exist (Elbaz et al., 2002; Garrett 2002; Gruppioni et al., 2003). In the Figure 10, we have reproduced the plot from Elbaz et al. (2002) complemented with galaxies detected within the ELAIS survey (Rowan-Robinson et al., 2004). Except at low luminosities where the contribution of cirrus to the IR luminosity becomes non negligible, the 1.4 GHz and 15 μm rest-frame luminosities are correlated up to z ∼ 1 and therefore predict very consistent total IR luminosities. UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 107

Figure 10. 15 μm versus radio continuum (1.4 GHz) rest-frame luminosities. Small ﬁlled dots: sample of 109 local galaxies from ISOCAM and NVSS. Filled dots with error bars: 17 HDFN galaxies (z ∼ 0.7, radio from VLA or WSRT). Open dots with error bars: 7 CFRS-14 galaxies (z ∼ 0.7, Flores et al., 1999, radio from VLA). Open diamonds: 137 ELAIS galaxies (z ∼ 0–0.4).

A similar result was later on obtained using the MIPS instrument onboard the Spitzer Space Observatory at 24 μm (Appleton et al., 2004). Several studies compared the SFR derived from the IR luminosity with the optical SFR derived from the Hα emission line (Rigopoulou et al., 2000; Cardiel et al., 2003; Flores et al., 2004; Liang et al., 2004). Rigopoulou et al. (2000) found a large excess of SFR(IR) versus SFR(Hα) even after correcting the latter for extinction. The Balmer decrement was only measured for limited number of objects in the sample and the extinction correction was derived from broadband photometry, which suffers from strong uncertainties in particular due to the degeneracy between age, metallicity and extinction. However, Cardiel et al. (2003) confirmed the direct measure of the SFR(IR) excess using the combination of NIRSPEC and LRIS, for the distant galaxies, and the Echelle Spectrograph and Imager (ESI), for the closer ones, at the Keck telescope. Using high resolution VLT-FORS2 spectra, Flores et al. (2004) and Liang et al. (2004) were able to measure directly the Balmer decrement (using Hα/Hβ or Hβ/Hγ ) and to subtract the underlying nebular emission lines with a fit of the stellar continuum. Although the SFR(IR) exceeds the estimate from emission lines for the most active objects, the data present a clear correlation between SFR(IR) and SFR(Hα) which suggests that the star formation regions responsible for the IR luminosity of distant LIRGs are not completely obscured. This correlations also confirms that the mid-IR is indeed a good SFR estimator.

5.3. NATURE OF THE ISOCAM GALAXIES

The Hubble Deep Field North and its Flanking Fields (HDFN+FF) provides the best coverage in spectroscopic redshifts and deep optical images for an ISOCAM deep 108 D. ELBAZ survey. We used the revised version of the Aussel et al. (1999) catalog for which 85% (71%) of the 40 (86) galaxies above 100 μJy (30 μJy) have a spectroscopic redshift to determine the average properties of ISOCAM galaxies summarized in the Figures 11 and 12. Their optical counterparts are relatively bright and their median-mean redshift is close to z ∼ 0.8 (Figure 11). Note the redshift peaks in which the ISOCAM galaxies are located, leaving wide empty spaces in between. Most ISOCAM galaxies are located within large-scale structures, here mainly those at z = 0.848 and z = 1.017, which might be galaxy clusters in formation where galaxy–galaxy interactions are amplified (Elbaz and Cesarsky, 2003, see discussion below). Thanks to the deepest soft to hard X-ray survey ever performed with Chandra in the HDFN, it is possible to pinpoint active galactic nuclei (AGNs) in this field including those affected by dust extinction. Only five sources were classified as

Figure 11. (Left) Histogram of the R(AB) magnitudes of the 15-μm ISOCAM galaxies detected in the HDFN+FF (revised catalog of Aussel et al.,downto∼30 μJy. (Right) Redshift distribution of the HDFN+FF ISOCAM galaxies.

Figure 12. (Left) Distribution of the log(LIR (8–1000 μm)/L ) of the HDFN+FF ISOCAM galaxies derived from their 15-μm ﬂux densities. (Right) Stellar mass as a function of redshift of ﬁeld HDFN proper galaxies (dark dots) and of ISOCAM galaxies (open circles). All stellar masses were derived using multi-bands SED modelling by Dickinson et al. (2003). UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 109

AGN dominated on the basis of their X-ray properties (Fadda et al., 2002). Hence, unless a large number of AGNs are so dust obscured that they were even missed with the 2 Megaseconds Chandra survey, the majority of ISOCAM galaxies are powered by star formation. This result is consistent with observations of local galaxies which indicate that only the upper luminosity range of ULIRGs are dominated by an AGN (Tran et al., 2001). Using the mid-far IR correlations (Chary and Elbaz, 2001; Elbaz et al., 2002, see also Section 5.2), the LIR distribution of the HDFN mid-IR sources is plotted in Figure 12a. Most of them belong to the LIRG and ULIRG regime, although when including flux densities below completeness down to 30 μJy, one finds also intermediate luminosities. Finally, their stellar masses are among the largest in their redshift range, when compared to the stellar mass estimates by Dickinson et al. (2003) in the HDFN (Figure 12b). In the local universe, both LIRGs and ULIRGs exhibit the typical morphology of major mergers, i.e. mergers of approximately equal mass (Sanders and Mirabel, 1996; Sanders et al., 1999). In the case of LIRGs, the merging galaxies show a larger separation than for ULIRGs, which are mostly in the late phase of the merger. Figure 13 (Elbaz and Moy, 2004) presents the HST-ACS morphology of a sample of z ∼ 0.7 LIRGs in the GOODSN field (extended HDFN). Less than half of these galaxies clone the morphology of local LIRGs, which implies that the physical processes switching on the star formation activity in distant LIRGs might be different than for local ones. The gas mass fraction of younger galaxies

Figure 13. HST-ACS images of LIR galaxies with 11 ≤ log(LIR/L) ≤ 12 (LIRGs) and z ∼ 0.7. The double-headed arrow indicates the physical size of 50 kpc. The IR luminosity increases from left to right and from top to bottom. 110 D. ELBAZ being larger, other types of interactions might generate a LIRG phase in the distant universe, such as minor mergers or even passing-by galaxies producing a tidal effect. The fact that such interactions are more frequent than major mergers could also explain the importance of the LIRG phase for galaxies in general and also the possibility for a galaxy to experience several intense bursts in its lifetime. The appearance of this phase of violent star formation could be facilitated during the formation of groups or clusters of galaxies. A striking example of this is given by the large fraction of LIRGs detected in the distant galaxy cluster J1888.16CL, located at a redshit of z = 0.56 (Duc et al., 2004). Among the 27 objects for which spectra were obtained, six of them belong to the cluster while an extra pair with slightly higher redshifts may lie in infalling groups. All eight galaxies exhibit weak emission lines in their optical spectra, typical of dust enshrouded star forming galaxies, none of these lines being broad enough to indicate the presence of type I AGNs. In this relatively young galaxy cluster, the −1 mechanism that may be triggering SFRs between 20 and 120 M yr for at least eight objects of the cluster could well be tidal collisions within sub-structures or infalling groups. On a more local scale, the mid-IR emission of the Abell 1689 (z = 0.181) galaxies exhibits an excess of the B-[15] colour with respect to richer and closer galaxy clusters, such as Coma and Virgo, which suggests the presence of a mid-IR equivalent to the Butcher–Oemler effect, i.e. the star formation activity of galaxies as reflected by their mid-IR emission increases with increasing redshift (Fadda et al., 2000). The high fraction of blue galaxies initially reported for this cluster by Butcher and Oemler (1984) was later on confirmed by Duc et al. (2002), who also found that the actual SFR for these galaxies was on average 10 times larger than the one derived from the [O II] emission line implying that 90% of the star formation activity taking place in this cluster was hidden by dust. LIRGs could then be a tracer of large-scale structures in formation as suggested by their redshift distribution (Figure 11b, see also Elbaz and Cesarsky, 2003). In order to test this hypothesis (Moy and Elbaz in preparation) compared the fraction of ISOCAM galaxies (above the completeness limit of 0.1 mJy) found in “redshift peaks” to the one obtained when randomly selecting field sources of equal optical and K-band magnitudes in the same redshift range. Redshift peaks of different strengths, measured as N-σ , were defined by smoothing the field galaxies redshift distribution by 15,000 km/s (Figure 14a) and measuring the peaks N-σ above the smoothed distribution. The Monte-Carlo samples of field galaxies to be compared to the redshift distribution of the ISOCAM galaxies were selected within the real redshift distribution and not the smoothed one. The distribution of the Monte-Carlo simulations are shown with error bars at the 68 and 90% level. When considering the whole ISOCAM catalog, i.e. including faint sources below completeness, it appears that ISOCAM galaxies are more clustered than field galaxies. Less than 32% of the simulations present a comparable clustering than the whole ISOCAM sample which contains a large fraction of non LIRGs. Strikingly, when only selecting the brightest galaxies above the completeness limit of 0.1 mJy, one finds that they are LIRGs UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 111

(a) (b)

Figure 14. (a) Redshift distribution of field galaxies in the GOODSN field. The continuous line is the smoothed distribution with a window of 15,000 km/s. (b) Differential fraction of sources within redshift peaks (see definition in the text) stronger than N-σ. Small filled squares: fraction of sources in peaks for the whole optical catalog of 930 field galaxies (Wirth et al., 2004). Small open squares: median of the fraction of sources in redshift peaks for a series of Monte-Carlo simulations of a sub-sample of the field galaxies corresponding to the same number of galaxies as in the ISOCAM catalog and within the same range of redshifts and optical-near IR magnitudes. Error bars contain 68 and 90% of the simulations. Large Open Circles: total sample of ISOCAM galaxies (75 sources with spectroscopic redshifts and optical-near IR magnitudes). Large open squares: sub-sample of ISOCAM galaxies above the completeness limit of 0.1 mJy (41 sources). and ULIRGs which fall in the densest redshift peaks, above 6-σ . The probability to randomly select a sample of galaxies from the field (with equal magnitudes and redshift range) in the redshift peaks of 6-σ and above is less than 1%. As a result, mid-IR surveys are very efficient in selecting over-dense regions in the universe, which in return are very efficient in producing a LIRG. In contrast, mid-IR selected galaxies are locally less clustered than field galaxies (Gonzalez-Solares et al., 2004), which could be a natural result of the fact that only the less clustered galaxies still produce IR luminous phases while more clustered galaxies lived their IR luminous phase in the past (see also Elbaz and Cesarsky, 2003). Finally, one question remains to be addressed about distant LIRGs: how long does this starburst phase last and how much stellar mass is produced during that time? Marcillac et al. (2004) used a Bayesian approach and simulated 200,000 vir- tual high-resolution spectra with the Bruzual and Charlot (2003) code to determine the recent star formation history of distant LIRGs as well as their stellar masses. These ISOCAM galaxies were observed using the VLT-FORS2 (λ/λ ∼ 2000 in the rest-frame) in three different fields. A prototypical LIRG at z ∼ 0.7 is found to 10 have a stellar mass of ∼5×10 M and to produce about 10% of this stellar mass within about 108 years during the burst. A remarkable result of this study is that the position of distant LIRGs in a diagram showing the value of the H8 Balmer absorption line equivalent width versus the strength of the 4000 A˚ break signs the presence of a burst of star formation within these galaxies, with an intensity of about −1 50M yr as also derived from their mid-IR emission. This result supports the 112 D. ELBAZ idea that distant LIRGs are not completely opaque to optical light and that one can learn something about their star formation history based on their optical spectra. Liang et al. (2004) compared the gas metallicity of the same sample of objects than Marcillac et al. (2004) with local galaxies of similar absolute magnitudes in the B band. Even accounting for an evolution in the B-band luminosity, the distant LIRGs turn out to be about twice less metal rich. This result suggests that between z ∼1 and today, LIRGs do produce a large fraction of the metals located in their host galaxies in agreement with the strong evolution of the cosmic star formation history found by the models fitting the ISO source counts.

6. Cosmic Evolution, Star-Formation Rate History

The combination of surveys at different wavelengths, from ISOCAM, ISOPHOT and SCUBA, together with the the shape and intensity of the CIRB, was used by several authors to constrain the parameters of their backward evolution models assuming a combination of luminosity and density evolution as a function of redshift of the IR luminosity function at 15 or 60 μm : Roche and Eales (1999), Tan et al. (1999), Devriendt and Guiderdoni (2000), Dole et al. (2000), Chary and Elbaz (2001), Franceschini et al. (2001, 2003), Malkan and Stecker (2001), Pearson (2001), Rowan-Robinson (2001), King and Rowan-Robinson (2003), Takeuchi et al. (2001), Xu et al. (2001, 2003), Balland et al. (2002), Lagache et al. (2003), Totani and Takeuchi (2002), Wang (2002). Being limited by the sensitivity of the extragalactic surveys, the major output of these models was to show that LIRGs and ULIRGs were much more common in the past than they are today. Chary and Elbaz (2001) derived that the co-moving IR luminosity due to LIRGs was about 70 times larger at z ∼1 than it is today (Figure 15). Quite logically, it resulted that the contribution of LIRGs to the cosmic star formation history was so large in the past that it dominated the integrated star formation that galaxies experienced in the past. Hence, LIRGs should now be considered not as a type of galaxies but instead as a common phase of intense star formation that any galaxy may have experienced. Elbaz et al. (2002) derived the contribution to the peak of the CIRB at 140 μm from the population of LIRGs around z ∼ 0.8 detected in the ISOCAM 15 μm deep surveys and found that these objects alone can explain more than two thirds of the peak and integrated intensities of the CIRB (see Figure 16). Hence the CIRB is the signature of the strong redshift evolution of LIRGs and the fossil record of star formation which took place in such burst phases.

7. From ISO to Spitzer, Herschel, the JWST and ALMA

On 23 August, 2003, NASA’s Spitzer space telescope (formerly SIRTF) was launched. Among its ﬁrst results came the source counts at 24 μmdownto∼20 μJy UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 113

Figure 15. Star-formation rate from Chary and Elbaz (2001). Upper panel: min and max range from their model, and observed UV/opt data; dots represent Xu et al. (2001) model. Lower panel: 3 different evolution scenarios from their model and data corrected for extinction. Line: pure luminosity; upper dash: pure density; lower dash: luminosity+density. which confirmed what ISO deep surveys already saw: a strong excess of faint sources indicating a rapid redshift evolution of IR luminous galaxies. When compared to models developed to fit the ISO counts, the faint end of the Spitzer counts are perfectly fitted as shown in the Figure 17 reproduced from Chary et al. (2004). On the high flux density range, around 1 mJy and above, less galaxies are found than predicted by those models (see also Papovich et al., 2004). This is partly, if not integrally, due to the fact that even ISOCAM-15- μm number counts were initially overestimated above S15 ∼ 1 mJy as discussed in Section 5.1, since the data reduction techniques were not optimized for the surveys with little redundancy over a given sky pixel. Although some refinement of the template SEDs used in the models might be considered (as suggested by Lagache et al., 2004), the 24-μm number 114 D. ELBAZ

Figure 16. Integrated Galaxy Light (IGL, filled dots) and Extragalactic Background Light (EBL, open squares, grey area) from the UV to sub-millimeter (from Elbaz et al., 2002). EBL measurements from COBE: 200–1500 μm EBL from COBE-FIRAS (grey area, Lagache et al., 1999), 1.25, 2.2, 3.5, 100, 140 μm EBL from COBE-DIRBE (open squares). IGL in the U,B,V,I,J,H,K bands from Madau and Pozzetti (2000). The upper end of the arrows indicate the revised values suggested by Bernstein et al. (2001, factor two higher). 6.75 μm (ISOCAM-LW2 filter) IGL from Altieri et al. (1999, filled dot). Hatched upper limit from Mkn 501 (Stanev and Franceschini, 1998). The ISOCAM 15 μm IGL (2.4 ± 0.5 nW m−2 sr−1) is marked with a star surrounded by a circle. The other star surrounded by a circle is the prediction of the contribution of the 15- μm sources to 140 μm (Elbaz et al., 2002). counts appear to be perfectly consistent with the up-to-date 15-μm counts. More- over, the high flux density range does not strongly affect the conclusions of the models based on previous 15-μm counts since bright objects do not contribute significantly to the CIRB and to the cosmic density of star formation. Hence these refinements are not strongly affecting the conclusions derived on galaxy formation and evolution based on the ISO deep surveys and summarized in the previous sections (see also Dole et al., 2004 for the Spitzer counts in the far IR). Another hint on the consistency of Spitzer MIPS-24-μm surveys with ISOCAM- 15 μm is given by the comparison of the images themselves. Galaxies detected at 15 and 24 μm are clearly visible in both images in Figure 18, although several 24-μm sources do not have a 15-μm counterpart. This results from the combination of the better sensitivity of MIPS, by a factor 2 or slightly more for the deepest surveys (galaxies are detected down to S15 ∼ 40 μJy in the HDFN, see Aussel et al., 1999, without lensing magnification), and of the k-correction. For galaxies above z ∼ 1, the PAH bump centered on the PAH feature at 7.7-μm starts to exit to 15-μm broadband while it remains within the MIPS-24-μm band up to z ∼ 2. The combination of ISOCAM and MIPS can be used to test whether the library of template SEDs that were used to derive “bolometric” IR luminosities from 8 to 1000 μm for the 15-μm galaxies are correct at least in the mid-IR range. One of the most important test is to check whether the 7.7-μm PAHbump is still present at z ∼ 1 UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 115

Figure 17. Completeness corrected galaxy counts in the MIPS 24- μm channel from Spitzer observations of the ELAIS-N1 field from Chary et al., (2004). The error bars reflect the Poissonian uncertainty. The horizontal bars represent the minimum and maximum flux density in that bin. The lines show four models for 24- μm counts: King and Rowan-Robinson (2003, KRR), Xu et al. (2001, Xu), Chary and Elbaz (2001, CE), Lagache et al. (2003, LDP). The symbols are plotted at the average of the flux densities of the detected sources in that bin for the data while the lines are plotted at the counts- weighted flux average for the models. The lower plot in the figure shows the histogram of the actual number of sources detected in each flux bin without any completeness correction.

Figure 18. ISOCAM 15-μm image (left) of the Ultra-Deep Survey in the Marano FIRBACK field (depth 140 μJy, 80% completeness) versus Spitzer MIPS-24 μm image (right; depth 110 μJy, 80% completeness). The crosses identify 16 galaxies detected at 15 and 24 μm for which VLT-FORS2 spectra were obtained and which SEDs were fitted in Elbaz et al. (2004). and whether the 24 over 15-μm flux density ratio is consistent with the template SEDs used by the models from which star formation histories were derived. The template SEDs designed by Chary and Elbaz (2001) provide a very good fit to the combination of both mid-IR values for a sample of 16 galaxies detected by 116 D. ELBAZ

ISOCAM and MIPS (crosses on the Figure 18). The LIR(8–1000 μm ) derived from either the 15- or 24- μm luminosities or the combination of both to constrain the SED fit present a 1-σ dispersion of only 20% (Elbaz et al., 2004). For galaxies located around z ∼ 1, the relative 15 and 24- μm luminosities clearly suggest the presence of a bump at 7.7 μm as observed in nearby galaxies and due to PAHs. Many questions remain unsolved that will be addressed by future missions, staring with Spitzer. Only when a fair sample of redshifts will have been determined for the distant LIRGs will we be able to definitely ascertain the redshift evolution of the IR luminosity function. Already for the brightest part of it, campaigns of redshift measurements have started, on ELAIS fields for the nearby objects, and on the Marano field with VIMOS and the Lockman Hole with DEIMOS for more distant objects. The fields selected for Spitzer legacy and Guaranteed time surveys were also carefully selected to be covered at all wavelengths and followed spectroscopically, so that this issue should be addressed in the very near future. Due to confusion and sensitivity limits, direct observations in the far IR will not reach the same depth than mid-IR ones until the launch of Herschel scheduled for 2007. The direct access to the far IR distant universe with Herschel will certainly bring major information on galaxy formation, together with the James Webb Space Telescope (JWST) up to 30 μm and the Atacama Large Millimeter Array (ALMA), which will bring an improved spatial resolution for the z ∼ 2 and above universe. Among the questions to be solved, we do not resist to the temptation of listing some of our favorite ones: what is the connection of large-scale structure formation with LIRG phases in galaxies? Can we probe the formation of distant clusters by the detection of the epoch when galaxies formed stars in such intense starbursts, producing galactic winds and enriching the intra-cluster medium with metals? Have we really resolved the bulk of the hard X-ray background, which peaks around 30 keV, and not left unknown some deeply buried AGNs which could make a larger than 20% fraction of the mid-IR light from distant LIRGs? What is the future of a distant LIRG, is it producing stars in a future bulge or disk? How uncertain is the interpolation to lower luminosities than those observed done in the models used to derive cosmic star formation history scenarios? This is of course only an example of a vast series of questions which indicate that the field opened for a large part by ISO has a long life to come.

References

Adelberger, K. L. and Steidel, C. C.: 2000, ApJ 544, 218. Altieri et al.: 1999, A&A 343, L65. Appleton, P. N., Fadda, D. T., Marleau, F. R., et al.: 2004, ApJS 154, 147. Aussel, H., Cesarsky, C. J., Elbaz, D., and Starck, J. L.: 1999, A&A 342, 313. Balland, C., Devriendt, J., Silk, J.: 2003, MNRAS 343, 107. Barger, A. J., Cowie, L. L., Sanders, D. B., et al.: 1998, Nature 394, 248. UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 117

Bernstein, R. A., Freedman, W. L., and Madore, B. F.: 2002, ApJ 571, 56. Bruzual, G. and Charlot, S.: 2003, MNRAS 344, 1000. Butcher, H. and Oemler, A.: 1984, ApJ 285, 426. Cardiel, N., Elbaz, D., Schiavon, R. P., et al.: 2003, ApJ 584, 76. Cesarsky, C., Abergel, A., Agnèse,P., et al.: 1996, A&A 315, L32. Chapman, S. C., Smail, I., Ivison, R., et al.: 2002, ApJ 573, 66. Chapman, S. C., Blain, A. W., Ivison, R. J., and Smail, I.: 2003, Nature 422, 695. Chary, R. R. and Elbaz, D.: 2001, ApJ 556, 562. Chary, R. R., Casertano, S., Dickinson, M. E., et al.: 2004, ApJS 154, 80. Ciliegi, P., McMahon, R. G., Miley, G., et al.: 1999, MNRAS 302, 222. Clements, D., et al.: 1999, A&A 346, 383. Cohen, J. G., Hogg, D. W., Blandford, R., et al.: 2000, ApJ 538, 29. Cole, S., et al.: 2001, MNRAS 326, 255. Désert, F.-X., Puget, J.-L., Clements, D., et al.: 1999, A&A 342, 363. Devriendt, J. and Guiderdoni, B.: 2000, A&A 363, 851. Dickinson, M., Papovich, C., Ferguson, H. C., and Budaári,T.: 2003, ApJ 587, 25. Dole, H., Gispert, R., Lagache, G., et al.: 2000, Springer Lecture Notes in Ph., v 548, 54 (astro- ph/0002283). Dole, H., Gispert, R., Lagache, G., et al.: 2001, A&A 372, 364. Dole, H.: 2002, Proceedings of the ESA SP-511 conference in C. Gry, et al.: (eds.) “Exploiting the ISO Data Archive – Infrared Astronomy in the Internet Age”, (astro-ph/0211310). Dole, H., Le Floc’h, E., Perez-Gonzalez, P. G., et al.: 2004, ApJS 154, 87. Duc, P.-A., Poggianti, B., Fadda, D., et al.: 2002, A&A 382, 60. Duc, P.-A., Fadda, D., Poggianti, B., et al.: 2004, Proceedings of the IAU Colloquium 195, Outskirts of Galaxy Clusters: Intense life in the suburbs, A. Diaferio Ed. (astro-ph/0404183). Efstathiou, A., Oliver, S., Rowan-Robinson, M., et al.: 2000, MNRAS 319, 1169. Elbaz, D., Cesarsky, C. J., Fadda, D., et al.: 1999, A&A 351, L37. Elbaz, D., Cesarsky, C. J., Chanial, P., et al.: 2002, A&A 384, 848. Elbaz, D. and Cesarsky, C. J.: 2003, Science 300, 270. Elbaz, D. and Moy, E.: 2004, Proceedgins of the “Multiwavelength Cosmology” Conference, held on Mykonos Island, Greece, M. Plionis Ed., Kluwer Academic Publishers, p.173 (astro-ph/0401617). Elbaz, D., Dole, H., Le F’loch, E., and Marcillac, D.: 2004 (submitted). Fadda, D., Elbaz, D., Duc, P.-A., et al.: 2000, A&A 361, 827. Fadda, D., Flores, H., Hasinger, G., et al.: 2002, A&A 383, 838. Fadda, D., Lari, G., Rodighiero, G., et al.: 2004, A&A (in press) (astro-ph/0407649). Fang, F., Shupe, D., Xu, C., and Hacking, P.: 1998, ApJ 500, 693. Fixsen, D. J., Dwek, E., Mather, J. C., Bennett, C. L., and Shafer, R. A.: 1998, ApJ 508, 123. Flores, H., Hammer, F., Thuan, T. X., et al.: 1999, ApJ 517, 148. Flores, H., Hammer, F., Elbaz, D., et al.: 2004, A&A 415, 885. Franceschini, A., Aussel, H., Cesarsky, C. J., Elbaz, D., and Fadda, D.: 2001, A&A 378,1. Franceschini, A., Berta, S., Rigopoulou, D., et al.: 2003, A&A 403, 501. Garrett, M.: 2002, A&A 384, L19. Genzel, R. and Cesarsky, C.: 2000, ARAA 38, 761. Goldader, J. D., Meurer, G., Heckman, T. M., et al.: 2002, ApJ 568, 651. Goldschmidt, et al.: 1997, MNRAS 289, 465. Gonzalez-Solares, E. A., Oliver, S., Gruppioni, C., et al.: 2004, MNRAS 352, 44. Gruppioni, C., Pozzi, F., Zamorani, G., et al.: 2003, MNRAS 341, L1. Haarsma, D. B., Partridge, R. B., Windhorst, R. A., and Richards, E. A.: 2000, ApJ 544, 641. Hauser, M. and Dwek, E.: 2001, ARAA 37, 249. Herbstmeierer, U., Abraham, P., Lemke, D., et al.: 1998, A&A 332, 739. 118 D. ELBAZ

Hughes, D. H., Serjeant, S., Dunlop, J., et al.: 1998, Nature 394, 241. Juvela, M., Mattila, K., and Lemke, D.: 2000, A&A 360, 813. Kakazu, Y., Sanders, D., Jpseph, R., et al.: 2002, IAU184, astro-ph/0201326. Kawara, K., Sato, Y., Matsuhara, H., et al.: 1998, A&A 336, L9. Kawara, K., Matsuhara, H., Okuda, H., et al.: 2004, A&A 413, 843. Kessler, M., Steinz, J., Anderegg, M., et al.: 1996, A&A 315, L27. King, A. J. and Rowan-Robinson, M.: 2003, MNRAS 339, 260. Kim, D.-C. and Sanders, D. B.: 1998, ApJS 119, 41. Kiss, C., Abraham, P., Klaas, U., et al.: 2001, A&A 379, 1161. Lagache, G., Abergel, A., Boulanger, F., et al.: 1999, A&A 344, 322. Lagache, G. and Puget J.-L.: 2000 , A&A 355, 17. Lagache, G., Dole, H., and Puget J.-L.: 2003, A&A MNRAS 338, 555. Lagache, G., Dole, H., Puget J.-L., et al.: 2004, ApJS 154, 112. Lari, C., Pozzi, F., Gruppioni, C., et al.: 2001, A&A 325, 1173. Lemke, D., Klaas, U., Abolins, J., et al.: 1996, A&A 315, L64. Liang, Y., Hammer, F., Flores, H., Elbaz, D., and Cesarsky, C. J.: 2004, A&A 423, 867. Lilly, S. J., Le Fevre, O., Hammer, F., and Crampton, D.: 1996, ApJ 460, L1. Linden-Voernle, M. J. D., Norgaard-Nielsen, H. U., Jorgensenet, H. E., et al.: 2000, A&A 359, L51. Madau, P., Ferguson, H. C., Dickinson, M. E., et al.: 1996, MNRAS 283, 1388. Madau, P. and Pozzetti, L.: 2000, MNRAS 312, L9. Malkan, M. and Stecker, F.: 2001, ApJ 555, 641. Marcillac, D., Elbaz, D., Charlot, S., et al.: 2004, submitted A&A (submitted). Matsuhara, H., Kawara, K., Sato, Y., et al.: 2000, A&A 361, 407. Metcalfe, L., Kneib, J.-P., McBreen, B., et al.: 2003, A&A 407, 791. Meurer, G. R., Heckman, T. M., and Calzetti, D.: 1999, ApJ 521, 64. Miville-DeschˆenesM.-A., Lagache, G., and Puget, J.-L.: 2002, 393, 749. Okumura, K.: 1998, ESA ISOCAM PSF Report, www.iso.vilspa.esa.es/users/expl lib/CAM/ psf rep.ps.gz. Oliver, S., et al.: 1996, MNRAS 280, 673. Oliver, S., Rowan-Robinson, M., Alexander, D., et al.: 2000, MNRAS 316, 749. Oliver, S., Mann, R. G., Carballo, R., et al.: 2002, MNRAS 332, 536. Papovich, C., Dole, H., Egami, E., et al.: 2004, ApJS 154, 70. Patris, J., Dennefeld, M., Lagache, G., and Dole, H.: 2003, A&A 412, 349. Pearson, C.: 2001, MNRAS 325, 1511. Puget, J.-L., Abergel, A., Bernard, J.-P., et al.: 1996, A&A 308, L5. Puget, J.-L., Lagache, G., Clements, D., et al.: 1999, A&A 345, 29. Puget J.-L. and Lagache, G.: 2000, IAU204, astro-ph/0101105. Rigopoulou, D., Spoon, H. W. W., Genzel, R., et al.: 1999, AJ 118, 2625. Rigopoulou, D., Franceschini, A., Aussel, H., et al.: 2000, ApJ 537, L85. Roche, N. and Eales, S.: 1999, MNRAS 307, 111. Rodighiero, G. and Franceschini, A.: 2004, A&A 419, L55. Rodighiero, G., Lari, C., Fadda, D., et al.: 2004, A&A (in press) (astro-ph/0407639). Rowan-Robinson, M.: 2001, ApJ 549, 745. Rowan-Robinson, M., Lari, C., Perez-Fournon, I., et al.: 2004, MNRAS 351, 1290. Rush, B., Malkan, M. A., and Spinoglio, L.: 1993, ApJS 89,1. Sanders, D. B. and Mirabel, I. F.: 1996, ARA&A 34, 749. Sanders, D., Surace, J. A., Ishida, C. M.: 1999, in J. E. Barnes and D. B. Sanders (eds.) “Galaxy Interactions at Low and High Redshift” IAU Symposium 186, Kyoto, Japan, (astro-ph/9909114). Sato, Y., Kawara, K., Cowie, L. L., et al.: 2003, A&A 405, 833. Sato, Y., Cowie, L. L., Kawara, K., et al.: 2004, AJ 127, 1285. UNDERSTANDING GALAXY FORMATION WITH ISO DEEP SURVEYS 119

Serjeant, S., Oliver, S., Rowan-Robinson, M., et al.: 2000, MNRAS 316, 768. Serjeant, S., Efstathiou, A., Oliver, S., et al.: 2001, MNRAS 322, 262. Smail, I., Ivison, R. J., Blain, A. W., and Kneib, J.-P.: 2002, MNRAS 331, 495. Stanev, T. and Franceschini, A.: 1998, ApJ 494, L159. Starck, J.-L., Aussel, H., Elbaz, D. et al.: 1999, A&AS 138, 365. Steidel, C. C., Adelberger, K. L., Giavalisco, M., et al.: 1999, ApJ 519,1. Stickel, M., Bogun, S., Lemke, D., et al.: 1998, A&A 336, 116. Stickel, M., Lemke, D., Klaas, U., et al.: 2000, A&A 359, 865. Takeuchi, T., Ishii, T., Hirashita, H., et al.: 2001, PASJ 53, 37. Tan, J., Silk, J. and Balland, C.: 1999, ApJ 522, 579. Taniguchi, Y., Cowie, L. L., Sato, Y., et al.: 1997, A&A 328, L9. Totani, T. and Takeuchi, T.: 2002, ApJ 570, 470. Tran, Q. D., Lutz, D., Genzel, R., et al.: 2001, ApJ 552, 527. Wang, Y.: 2002, A&A 383, 755. Wirth, G. D., Willmer, C. N. A., Amico, P., et al.: 2004, AJ 127, 3121. Xu, C., Hacking, P., Fang, F., et al.: 1998, ApJ 508, 576. Xu, C., Lonsdale, C., Shupe, D., et al.: 2001, ApJ 562, 179. Xu, C. K., Lonsdale, C. J., Shupe, D. L., et al.: 2003, ApJ 587, 90. SOLAR SYSTEM THE PLANETS AND TITAN OBSERVED BY ISO

THIERRY FOUCHET∗, BRUNO BEZARD´ and THERESE ENCRENAZ LESIA, Observatoire de Paris, 5 Place Jules Janssen, 92195 Meudon Cedex, France (∗Author for correspondence: E-mail: [email protected])

(Received 30 June 2004; Accepted in ﬁnal form 13 December 2004)

Abstract. Infrared spectroscopic observations of planets and Saturn’s satellite Titan with the Infrared Space Observatory led to many significant discoveries that improved our understanding on the formation, physics and chemistry of these objects. The prime results achieved by ISO are: (1) a new and consistent determination of the D/H ratios on the giant planets and Titan; (2) the first precise measurement of the 15N/14N ratio in Jupiter, a valuable indicator of the protosolar nitrogen isotopic ratio; (3) the first detection of an external oxygen flux for all giant planets and Titan; (4) the first detection of some stratospheric hydrocarbons (CH3,C2H4,CH3C2H, C4H2,C6H6); (5) the first detection of tropospheric water in Saturn; (6) the tentative detection of carbonate minerals on Mars; (7) the first thermal lightcurve of Pluto.

Keywords: solar system

1. Introduction

At the time of ISO launch, the infrared spectrum of planets and Titan was not a virgin territory. Interplanetary spacecraft had carried infrared spectrometers in the vicinity of Mars, Jupiter, Saturn and Titan, Uranus, and Neptune. These instruments demonstrated the predominance of CO2 in the martian atmosphere, N2 on Titan, and H2 and He on the giant planets, and gave insights on the planetary composition in minor species. Dedicated planetary missions were also specifically designed to provide spatial and temporal resolution on the atmospheric structure, in order to constrain the planets’ dynamics and meteorology. However, instruments aboard planetary spacecraft generally lacked spectral resolution and sensitivity. In this respect, ISO instruments gave planetary scientists a significant improvement that allowed them to obtain important results on the formation, chemistry, and dynamics of planetary objects. Hereafter we detail the major ISO achievements with emphasis on the measurement of the planetary isotopic composition and its significance for our understanding of Solar System formation, on the oxygen exogenic flux in the outer Solar System, and on the hydrocarbon photochemistry. Earlier reviews of ISO results can be found in Lellouch (1999), Bézard (2000) and Encrenaz (2003). Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, The Netherlands, and the United Kingdom), and with the participation of ISAS and NASA. Space Science Reviews (2005) 119: 123–139 DOI: 10.1007/s11214-005-8061-2 C Springer 2005 124 T. FOUCHET ET AL.

2. Isotopic Composition

2.1. DEUTERIUM IN THE GIANT PLANETS

The D/H ratio is one of the most useful diagnostic of the formation and evolution of the giant planets. They formed in the protosolar nebula composed of gas, mainly H2 and He, and icy grains, enriched in deuterium. Their current D abundances thus reflect the relative fraction of presolar gas and grains that participated in their formation. In Jupiter and Saturn, the majority of the mass came from the gaseous component of the presolar nebula. Their D/H ratios are therefore regarded as an estimate for the presolar D/H value. In contrast, the masses of Uranus and Neptune are dominated by ices, which have enriched their atmospheres in deuterium. For this reason, measuring their D/H ratios could yield the deuterium abundance in the proto-uranus and proto-neptunian ices, and help to constrain the condensation processes in the presolar nebula. Prior to the ISO mission, the deuterium abundance had been measured in Jupiter and Saturn from near-infrared lines of HD (Smith et al., 1989). However, the interpretation of these lines was complicated by some blending with absorptions from CH4 and uncertainties on the cloud physical properties. Therefore, the most accurate value for Jupiter was obtained in-situ by the mass spectrometer aboard the Galileo probe (Niemann et al., 1998). For the three other giant planets, we had to rely on measurements of the CH3D/CH4 ratio, and on the knowledge of the fractionation coefficient f induced by the isotopic exchange reaction HD+CH4 H2 +CH3D. The latter coefficient is quite uncertain as the reaction equilibrium constant and kinetics, and the planetary convection timescales are not well known. As predicted by Bézard et al. (1986), ISO allowed the first detection of HD rotational lines (Encrenaz et al., 1996; Griffin et al., 1996; Feuchtgruber et al., 1999a; Lellouch et al., 2001). These lines enabled a simple and consistent determination of the deuterium abundance on the four giant planets. The LWS instrument detected the R(0) line on all giant planets. The SWS instrument detected the R(2) line (37.7 μm) on each planet (Figure 1), while the R(3) line was only detected on Saturn. For Jupiter and Saturn, Lellouch et al. (2001) inferred D/H ratios in H2 of / = . ± . × −5 / = . +0.85 × −5 (D H)H2 (2 4 0 4) 10 and (D H)H2 1 85−0.60 10 , respectively. Using CH3D and CH4 bands in the 7–9 μm range, Lellouch et al. (2001) also measured the / = . ± . × −5 / = . +1.4 × −5 D/H in methane, (D H)CH4 (2 2 0 7) 10 and (D H)CH4 2 0−0.7 10 , respectively for Jupiter and Saturn. Combining these two measurements yielded . ± . × −5 . +0.75 × −5 global D/H ratios of (2 25 0 35) 10 and 1 70−0.45 10 , respectively. Cor- recting for a small atmospheric enrichment in deuterium from ices that constituted the proto-cores of these planets, Lellouch et al. (2001) obtained a protosolar ratio −5 of (D/H)ps = (2.1 ± 0.4) × 10 . Other estimates of the (D/H)ps ratio had been obtained analysing the evolution of the 3He/4He ratio from its protosolar value to its current solar value (Geiss and Gloeckler, 1998). The current solar ratio was measured in the solar wind, while PLANETS AND TITAN 125

Figure 1. The HD R(2) rotational line observed on the four giant planets.

the protosolar 3He/4He ratio was assumed equal to the current Jupiter ratio measured in situ by Galileo (Niemann et al., 1998), or equal to the ratio measured in meteorites (Geiss and Gloeckler, 1998). These values led to protosolar deuterium −5 −5 composition, (D/H)ps = (1.94 ± 0.5) × 10 and (D/H)ps = (2.1 ± 0.5) × 10 , respectively, in excellent agreement with the ISO measurement. This confirms that Jupiter is a reliable indicator of the protosolar D/H. The protosolar value also indicates a weak decrease of the D/H ratio in the Local Interstellar Medium −5 (LISM) since the formation of the Solar System: (D/H)LISM = (1.5 ± 0.1) × 10 (Linsky, 1998). For Uranus and Neptune, the retrieved isotopic ratios were larger than those / = . +3.5× −5 / = . +2.5× measured in Jupiter and Saturn: (D H)H2 5 5−1.5 10 and (D H)H2 6 5−1.5 10−5, respectively (Feuchtgruber et al., 1999a). This confirms that at least a fraction of the ices that constituted the proto-cores of these planets was mixed in their atmospheres. Using the interior models of Podolak et al. (1995), Feuchtgruber et al. (1995a) were able to evaluate the deuterium composition of the proto-uranian and / = . +7.6 × −5 / = . +8.5 × −5 proto-neptunian ices: (D H)ices 9 4−4.2 10 and (D H)ices 10 8−4.7 10 , respectively. These values differ by about a factor of three from the D/H ratios measured in cometary ices, ∼30 × 10−5 (see Altwegg and Bockelée-Morvan, 2003 for a review). Drouart et al. (1993) and Hersant et al. (2001) used this D/H variability within Solar System objects as a chronometer of their formation. They argued that the various D/H ratios resulted from the different levels of isotopic exchange between 126 T. FOUCHET ET AL. water enriched in deuterium by interstellar chemistry and the hydrogen gas. The sooner ices condensed in the protosolar nebula, the higher were their D/H ratios. Using turbulent models of the nebula Drouart et al. (1999) and Hersant et al. (2001) calculated that comets formed in ∼105 years, while Uranus and Neptune cores took ∼106 years to agglomerate.

2.2. DEUTERIUM IN TITAN

In Titan, methane is the prime deuterium carrier. Using PHT-S and SWS data, Coustenis et al. (2003) observed both CH3D and CH4 between 7 and 9 μmto / = . +3.2 × −5 derive a D/H ratio of (D H)CH4 8 7−1.9 10 . The ISO result is smaller than . +1.4 × −4 the value inferred from Voyager 1/IRIS spectra, 1 5−0.5 10 (Coustenis et al., 1989), but agrees with the ground-based determination of (7.75 ± 2.25) × 10−5 of Orton et al. (1992). Two different scenarios have been proposed to explain this deuterium enrichment in Titan’s atmosphere compared to that in the protosolar gas. Both scenarios account for the replenishment of Titan’s atmosphere in methane, which is effectively photolysed in heavier hydrocarbons by the solar ultraviolet ﬂux. In the ﬁrst scenario, Lunine et al. (1999) proposed that methane evaporates from a local, small reservoir, located near the surface. In this case, over 4.5 byr, solar photolysis may have enriched the D/H in atmospheric methane up to a factor of 4. In the second scenario, Mousis et al. (2002) argued that the atmosphere could be replenished by cryovolcanism. In this case Mousis et al. (2002) advocated that solar photolysis did not induce any fractionation, and that the observed deuterium enrichment is primordial. In this case, the D/H ratio in Titan could also be used as a chronometer for the condensation of ices at Saturn’s heliocentric distance in an approach similar to that used by Drouart et al. (1999) and Hersant et al. (2001).

2.3. 15N/14NRATIO IN JUPITER

The value of the 15N/14N in the protosolar nebula had long been an unsolved mystery. Measurements in comets (Altwegg and Bockel´ee-Morvan, 2003), in the solar wind trapped in the lunar regolith (Kerridge, 1993), in situ in the solar wind (Kallenbach et al., 1998), and in meteorites (Pillinger, 1984) gave values ranging between 2.5 × 10−3 and 6.9 × 10−3 that were impossible to explain within the framework of a consistent, global scenario of the nebula. In addition, measurements in Jupiter’s atmosphere yielded a ratio comparable to the terrestrial 15 14 −3 ratio (( N/ N)⊕ = 3.68 × 10 ), but with uncertainties as large as a factor of 2 (Encrenaz et al., 1978; Tokunaga et al., 1980). Using ISO-SWS observations at 10 μm, Fouchet et al. (2000a) reﬁned the measurement in Jupiter, yielding a 15 14 . +0.9 × −3 N/ N ratio of (1 9−1.0) 10 , much smaller than the terrestrial ratio (Figure 2). A similar result was found by the in situ measurements of Galileo (Owen et al., PLANETS AND TITAN 127

Figure 2. Comparison between Jupiter’s ISO spectrum (solid line) and two synhtetic models with terrestrial (dashed line) and half-terrestrial (dotted line) 15N/14N.

2001), and latter by the Cassini/CIRS observations (Abbas et al., 2004; Fouchet et al., 2004). Formation models of Jupiter (Owen et al., 2001; Hersant et al., 2004) suggest that the planet has retained most of the nitrogen present in its feeding zone. Therefore, the jovian nitrogen isotopic ratio can be seen as a tracer of the protosolar ratio. This view is supported by a recent reanalysis of lunar samples that gave an upper limit of 2.8 × 10−3 for the trapped solar wind (Hashizume et al., 2000). Such a difference between the protosolar and the terrestrial 15N/14N ratios bears important consequences for the origin of the Earth atmosphere. Terrestrial N2 must originate from a minor nitrogen reservoir in the protosolar nebula, highly fractionated (a factor of 2) with respect to the largest reservoir. If interstellar chemistry seems the most probable culprit (Charnley and Rodgers, 2002; Al´eonand Robert, 2004), the exact fractionation mechanism still remains to be identiﬁed, along with the main nitrogen carrier to the Earth. 128 T. FOUCHET ET AL.

3. Exogenic Oxygen to the Giant Planets

One of the most striking results obtained from ISO observations is the discovery of H2O in the stratospheres of the four giant planets (Feuchtgruber et al., 1997, 1999b; Lellouch et al., 2002), and Titan (Coustenis et al., 1998), as well as the detection of CO2 in the stratospheres of Jupiter, Saturn and Neptune (Feuchtgruber et al., 1997, 1999b; Lellouch et al., 2002). The presence of water above the tropopause cold trap can only be explained by an external influx. CO2 also condenses at the tropopause of Saturn and Neptune, while internal sources on Jupiter are believed to be weak, since carbon–oxygen bonded species present in the planetary interior are effectively converted to CH4 at shallow atmospheric levels. Therefore, the H2O and CO2 detections from ISO provided evidences for an exogenic supply of oxygen in the atmospheres of outer Solar System objects. Three different possible sources are possible: (i) infall of interplanetary dust particles (IDPs), (ii) sputtering from icy rings and satellites, and (iii) rare impacts from kilometre-sized bodies. In Jupiter, CO2 may orginates from the Shoemaker-Levy 9 (SL9) impacts that oc- cured in 1994. The shock chemistry induced by the impacts produced CO and H2O that were deposited by the plumes at atmospheric levels above the 100-μbar level. CO and H2O can subsequently react to form CO2, while all the chemical products are slowly transported, both horizontally from the location of the impacts (44◦S), and vertically. The ISO-SWS 14 arcsec×27 arcsec aperture allowed Lellouch et al. (2002) to crudely resolve the CO2 latitudinal variation. They derived column densities, (6.3±1.5)×1014 cm−2 in the southern hemisphere, (3.4±0.7)×1014 cm−2 in the Equatorial Region, and <7 × 1013 cm−2 in the northern hemisphere, reflecting the partial horizontal diffusion of SL9 products. Moreover, the ISO CO2 abundance agrees with photochemical predictions computed using the CO latitudinal distribution observed in the millimeter range (Moreno et al., 2003). Since the ISO spectrometers apertures were similar to the Jupiter angular size at wavelengths longer than 20 μm, it was not possible to observe such a latitudinal gradient for H2O. However, Lellouch et al. (2002) combined the SWS and LWS observations to demonstrate that Jupiter stratospheric H2O most likely originates from the SL9 impacts. As shown in Figure 3, the use of water lines with different intrinsic strengths in the two instrumental ranges led to the conlusion that water is restricted to pressures less than 0.5±0.2 mbar. If H2O originated from a continuous source like infall of IDPs, it would be present in the whole stratosphere. Rather, the water must originate from an event recent enough for H2O not to have diffused downwards, i.e. most likely from the Shoemaker-Levy 9 impacts. The disk-averaged column density of (2 ± 0.5) × 1015 cm−2 also agrees with predictions from shock chemistry models. In addition, Lellouch et al. (2002) put an upper limit of 8 × 4 −2 −1 10 cm s to the continuous exogenic flux of H2O. In Uranus and Neptune, Feuchtgruber et al. (1997) derived water column densities of 5 × 1013–12 × 1013 cm−2 and 3 × 1014–6 × 1014 cm−2, respectively. Using transport models for the upper atmospheres, these figures yielded PLANETS AND TITAN 129

Figure 3. The Jupiter H2O 39.38-μm and 66.44-μm line for different water vertical models. This shows that Jupiter water is restricted to pressures less than 0.5 ± 0.2 mbar.

estimates for continuous external fluxes of 0.6 × 105–1.6 × 105 cm−2 s−1 and 1.2×105–150×105 cm−2 s−1. Such fluxes can be accounted for by IDPs if the flux measured at 1 AU is extrapolated to 30 AU. In Neptune, the CO2 column density was estimated to 8 × 1014 cm−2. Such an abundance can be accounted for by two different sources. IDPs constitute the first possible origin, since the CO2/H2O ratio in Neptune is consistent with the CO2/H2O ratio measured in comets. A second possible source is the reaction of H2O with the neptunian stratospheric CO – whose origin remains uncertain. On Saturn, Moses et al. (2000a) estimated column densities of (6.3 ± 1) × 14 −2 15 −2 10 cm and (1.4 ± 0.4) × 10 cm , for CO2 and H2O, respectively. Using a photochemical model, Moses et al. (2000a) matched the observed abundances with an exogenic O flux of (4±2)×106 cm−2 s−1. Moreover, Moses et al. (2000a) found 130 T. FOUCHET ET AL. that, to reproduce the CO2/H2O ratio observed in Saturn’s stratosphere, most of the exogenic oxygen must be delivered in the form of CO and/or CO2, with a CO/H2O ratio of ∼3. Cometary ices do not match such a high ratio. On Titan, Coustenis . +1.9 × 14 −2 et al. (1998) detected water with a column density of 2 6−1.6 10 cm . Such an abundance could be achieved by an exogenic flux of 0.8 × 106–2.8 × 106 cm−2 s−1. The comparison between the exogenic oxygen fluxes derived for Jupiter, Saturn and Titan points towards a local source in the Saturn system. Indeed, the upper limit derived for the permanent flux at Jupiter, 8 × 104 cm−2 s−1, is much lower than the flux needed to account for the Saturn and Titan water. The variation of the IDPs flux with heliocentric distance is not well known, but it is expected to be lower at Saturn’s distance, where cometary activity is reduced. A SL9-like event can also be rejected for Saturn, as on the long term, most of the O deposited by the impacts is converted to CO, with a CO/H2O ratio of ∼100 (Bézard et al., 2002) higher than the ratio of ∼3 derived for Saturn by Moses et al. (2000a). Instead, it is sugested that the rings could be an efficient source of oxygen for Saturn, as well as Hyperion for Titan. Assessing the external sources of oxygen in Saturn and Titan constitutes one of the goals of the Cassini mission.

4. Upper Atmosphere

4.1. HYDROCARBON PHOTOCHEMISTRY

On the giant planets and Titan, the ultraviolet solar flux photolyses CH4 to produce CH, CH2 and CH3 radicals in the upper stratosphere (∼1 μbar). These radicals recombine in heavier hydrocarbons, such as C2H6,C2H4,C2H2,C3- and C4-species. Eddy diffusion then transports the newly formed molecules to the whole stratosphere. This hydrocarbon photochemistry still holds mysteries for theoretical mod- elers. In particular, some free parameters, such as the relative chemical production rates and the intensity of the eddy diffusion, require to be tuned. This can be achieved by measuring the relative abundance between different molecules, and observing how the abundance of an individual molecule evolves with altitude. ISO-SWS allowed Drossart et al. (1999) to observe for the first time the CH4 fluorescence at 3.3 μm on Jupiter and Saturn. The fluorescence probes low pressures, ∼1 μbar, where the solar energy absorbed by methane is re-emitted before being thermalized by collisions. By chance, this pressure level coincides with the homopause, the level where molecular diffusion becomes as efficient as eddy diffusion, resulting, at higher altitudes, in the diffusive separation of the elements according to their mass. The intensity of CH4 fluorescence constrains the CH4 column density, and is therefore a diagnostic of the homopause level and eddy 6 6 diffusion coefficient. Drossart et al. (1999) found KH = 6 × 10 –8 × 10 and 7 7 2 −1 KH = 3 × 10 –5 × 10 cm s for Jupiter and Saturn, respectively. These values agree with previous determinations made from Voyager observations. CH4 PLANETS AND TITAN 131

Figure 4. The ﬁrst detection of CH3 on Neptune compared with synthetic spectra.

fluorescence thus provides a new method to determine eddy mixing coefficients in the outer Solar System. The chain of reactions that produces hydrocarbons in the outer Solar System starts from radicals. However, these primitive blocks had never been observed prior to the ISO mission. For the first time, the SWS instrument allowed Bézard et al. (1998, 1999) to detect a radical, CH3, on Saturn and Neptune, from its ν2 band at 16.5 μm (Figure 4). The derived column densities are respectively 1.5 × 13 . × 13 . +1.2 × 13 −2 10 –7 5 10 and 1 6−0.9 10 cm . These values conflict with predictions from photochemical models, which overestimate the CH3 abundance by a factor of ∼5. Methyl is mostly sensitive to the eddy mixing coefficient at the homopause, and to the poorly known methyl recombination rate. Since eddy mixing coefficients at the homopause have been established consistently through various different means, Bézard et al. (1998, 1999) stressed the need for new laboratory data on the three- M body self-recombination of the CH3 radical 2CH3 −→ C2H6. The SWS and PHT-S instruments also led to the first detection of many products of the hydrocarbon photochemistry in the outer Solar System: C2H4 on Neptune 132 T. FOUCHET ET AL.

(Schulz et al., 1999), CH3C2H on Jupiter (Fouchet et al., 2000b), CH3C2H and C4H2 on Saturn (de Graauw et al., 1997), and benzene (C6H6) on Jupiter, Saturn and Titan (Bézard et al., 2001; Coustenis et al., 2003). (Benzene had been previously detected on Jupiter, but only over the polar regions (Kim et al., 1985), whereas the ISO detection pertains to the whole planet.) ISO-SWS also allowed Fouchet et al. (2000b), Moses et al. (2000b) and Coustenis et al. (2003), to determine the vertical distributions of C2H6 and C2H2 on Jupiter, and C2H2 on Saturn and Titan. All these observations made possible the development of refined photochemical models, in particular for Jupiter and Saturn (Moses et al., 2000b, 2001). Models generally fit the abundances and vertical distributions observed by ISO, but with a few exceptions, such as benzene on Jupiter and Saturn, and ethylene on Neptune. In the case of +4.5 × 14 . +2.1 × 13 −2 benzene, the observed column densities, 9−7.5 10 and 4 7−1.1 10 cm , respectively, are larger than that predicted by photochemistry. Wong et al. (2000) proposed that benzene production could be induced by auroral particle precipitation, but their model also fails to predict the observed global abundance on Jupiter. Therefore, new laboratory studies for benzene production in the conditions relevant to the outer Solar System are needed. In the case of ethylene on Neptune, the column densities predicted by photochemical models strongly depend on the CH4 branching ratios at Lyman α. The most recent laboratory measurements reproduce the observed column density, 1.1×1014–3×1014, but the need for more laboratory work is still present. On Uranus, Encrenaz et al. (1998) also improved the current photochemical models. From the non-detection of CH4 and the observed signature of C2H2, Encrenaz et al. (1998) were able to constrain the eddy mixing coefficient at Uranus tropopause in the range 0.5 × 104–1 × 104 cm2 s−1. This confirms that turbulent mixing is much less vigorous on Uranus than it is on Jupiter, Saturn and Neptune.

4.2. ATMOSPHERIC STRUCTURE

ISO-SWS observed the H2 S(0) and S(1) quadrupole lines in the four giant planets. These lines probe pressures levels located between 10 and 1 mbar and allowed Fouchet et al. (2003) to determine the stratospheric hydrogen para fraction for the ﬁrst time. In Jupiter, Saturn and Neptune, the measured para fraction present a signiﬁcant departure from thermodynamic equilibrium. Fouchet et al. (2003) inter- pretated this situation as a result from the lagged conversion between the ortho and para states as H2 is transported upward by eddy diffusion. In constrast, the uranian stratosphere lies close to thermodynamic equilibrium. Using a one-dimentional diffusion model, Fouchet et al. (2003) showed that the departure from the thermodynamic para fraction appeared anti-correlated with the density of stratospheric aerosols. This provided quantitative evidence for an equilibration mechanism dominated by conversion on aerosols rather than by conversion in the gas phase. PLANETS AND TITAN 133

5. Tropospheric Composition

With H2 and He the two main constituents (over 98% by volume), the atmosphere of the giant planets constitutes a highly reducing environment, where elements are mostly present in reduced forms, i.e. CH4,NH3,H2O, H2S, PH3. In the troposphere, the temperature decreases with increasing altitude to a minimum temperature of ∼110 K at Jupiter and down to ∼50 K at Uranus and Neptune. These low temperatures induce the condensation of gases: on Jupiter and Saturn, H2O, H2S, NH3 do condense between 10 and 1 bar, while on Uranus and Neptune CH4 also condenses out. This situation makes difficult the determination of the global composition, hidden below a thick cloud cover. Using ISO-SWS observations at 5 μm and 10 μm, and ISO-LWS observations in the range 43–197 μm, Fouchet et al. (2000a) and Burgdorf et al. (2004) derived two sharply different ammonia vertical profiles. At 10 μm and in the ammonia rotational lines, the atmosphere appears slighlty under-saturated, with a NH3 mixing ratio reaching the solar abundance at about 1 bar. In contrast, at 5 μm, Fouchet et al. (2000a) found a highly dessicated atmosphere, with a NH3 mixing ratio increasing with pressure to a solar abundance at ∼4 bar. This later vertical profile is quite similar to that measured by the Galileo Probe (Folkner et al., 1998) on its specific entry point. These two independent measurements confirm the spatial heterogeneity of Jupiter, covered by a majority of wet, cloudy, regions, with some spots, highly dry and almost cloud-free, mostly located in the North Equatorial Belt. If ISO-SWS could not resolve these spots, also called the 5-μm hot spots, they still dominate Jupiter spectrum at this wavelength. Dynamical models experience some difficulties in reproducing this structure. Hueso et al. (1999) showed that a traditional convection pattern, with ascent of wet air and descent of dry air, is not able to generate the high level of dessi- cation observed. Showman and Dowling (2000) suggested an alternative model, where the atmosphere and the gas vertical distributions are stretched by planeraty waves. On Saturn, the ISO-SWS observations at 5 μm led to the first detection of tropospheric water (de Graauw et al., 1997). The H2O mixing ratio was found strongly undersaturated, similarly to NH3 on Jupiter. This suggests that Saturn’s dynamics generates the same hot spot features as observed on Jupiter. However, Saturn images at 5 μm do not reveal the high radiance contrast observed on Jupiter. The Cassini mission, especially the VIMS instrument, will provide new insights on this issue. Phosphine was observed through its rotational lines both on Jupiter and Saturn (Davis et al., 1996; Burgdorf et al., 2004) using ISO-LWS. The retrieved vertical profiles (Burgdorf et al., 2004) agree well with those derived from 10 μm observations (Fouchet et al., 2000a; de Graauw et al., 1997) and clearly show a cut-off at pressures below 300 mbar due to the ultraviolet photolysis as predicted by photochemical models. 134 T. FOUCHET ET AL.

From both the SWS and LWS spectrometers, Burgdorf et al. (2003) used the collision-induced H2-He continuum to obtain a new determination of the He/H2 ra- . +1.7 tio in Neptune. The derived helium abundance, 14 9−2.2%, conﬁrms, with smaller error bars, previous determination obtained from the combined analysis of Voyager- 2 IRIS and radio occultation experiment (Conrath et al., 1991). This result is consistent with the protoplanetary estimate of the helium abundance, as expected from the internal models of Neptune which predict no differenciation (Marley, 1999).

6. Cloud Structure

6.1. MARS

Dust plays an important role in the geologic and climatic history of Mars. As an efficient absorber, it strongly affects the atmospheric thermal balance, and hence the atmospheric dynamics, while it erodes and coats geologic and mineralogic features that hold insights into the planet’s history. Therefore, the continuous monitoring of the martian dust, in particular of its spatial and temporal variations, constitutes an essential objective of the martian exploration. Using the ISO-SWS spectrum of Mars, Fedorova et al. (2002) applied a new method to sound the atmospheric dust. They used the optically thick 2.7-μmCO2 band, where solar photons are absorbed before reaching the ground. In clear atmospheric conditions, the flux in the core of the CO2 band is thus zero, whereas on dusty conditions a small fraction of the solar flux is reflected by aerosols at high altitudes before being absorbed. Indeed, ISO-SWS recorded a flux of ∼20 Jy at 2.68–2.70 μm and at 2.76– 2.78 μm (Figure 5), from which Fedorova et al. (2002) derived a dust optical depth of τ = 0.33±0.13. Their analysis also revealed the existence of a sharp absorption

Figure 5. The spectrum of Mars at 2.7 μm observed by ISO (solid line). Synthetic spectrum for grey aerosols model (dashed line). Synthetic spectrum for dust with phyllosilicates (dotted line). PLANETS AND TITAN 135 band in the dust spectrum between 2.7 and 2.8 μm, that probably indicates the presence of hydrated minerals, most likely phyllosilicates, in the dust. Therefore, ISO observations led to the development of a new dust sounding method that will be used by current and future missions to Mars. As a by-product, ISO observations of Mars were used to infer the mean water vapour vertical distribution, using both the SWS (Lellouch et al., 2000a) and LWS (Burgdorf et al., 2000; Sidher et al., 2000).

6.2. GIANT PLANETS

As stated earlier, several chemical species (H2O, H2S, NH3,CH4) are expected to condense in the tropospheres of the giant planets. This condensation is well established from the measurement of subsaturated abundances above the theoretical condensation levels, and from the obvious thick cloud covers of these planets. How- ever, ice spectral signatures had remained elusive. From the ISO-SWS spectrum at 3 μm, Encrenaz et al. (1996) and Brooke et al. (1998) first identified NH3 ice on Jupiter. The best-fit model included an upper cloud located at 0.55 bar, with 10-μm NH3 ice particles, above a grey thick cloud located between 1 and 2 bars. Since, the detection has been confirmed by Baines et al. (2002) from Galileo/NIMS observations, also at 3 μm. The spatial resolution provided by Galileo further demonstrated that pure NH3 ice is restricted to some specific regions on the planet, associated with strong convective updrafts. The reason why particles are featureless on the rest of the planet remains to be explained. In the case of Uranus and Neptune, the PHT-S spectra between 2.5 and 4 μm were used to constrain the cloud level and albedo (Encrenaz et al., 2000). For both planets, the albedo was found to be very low <1%, while the cloud level was found to be at about 3 bars. Encrenaz et al. (2000) also found evidence for cirrus cloud above ∼0.3 bar, presumably due to CH4 particles. Burgdorf et al. (2003) also found that their 46–185 μm spectrum is best reproduced by invoking a CH4 cloud in the upper troposphere with particle sizes lying between 15 and 40 μm.

7. Planetary Surfaces

7.1. MARS

The Martian atmosphere is transparent at nearly all wavelengths, allowing the sounding of the martian surface from the near infrared to the thermal infrared. ISO-SWS indeed measured the albedo and emissivity of Mars from 2.4 to 45 μm (Lellouch et al., 2000a). The spectrum shows the usual features due to hydrated minerals and silicates, but Lellouch et al. (2000a) also tentatively identified carbonate signatures at 5.3, 6.3, 7.2, 11.1, 26.5, 31, and 43.5 μm. The four signatures between 5 and 12 μm had been already identified in Mariner 9/IRIS spectra, but 136 T. FOUCHET ET AL. were not attributed to carbonates. The ISO detection remains tentative as no single carbonate species can be uniquely identified from this association of features. Since then, the detection of carbonates has been confirmed from MGS/TES observations by Bandfield et al. (2003). However, carbonate signatures were present only in a small fraction of the MGS/TES spectra. This leads Bandfield et al. (2003) to favour the presence of carbonates in the martian dust rather than on the surface itself. Such an origin was also supported from ISO observations (Lellouch et al., 2000a). This issue is important for our views on the evolution of the Martian atmosphere, since carbonates could represent a significant sink for the initial atmosphere of CO2. However, since carbonate minerals are currently only identified in the dust, they cannot yet be interpreted as the signature of a primordial thick atmosphere. More work remains to be done for current and future Martian orbiters.

7.2. PLUTO

Using ISO-LWS and PHT-Sobservations, Lellouch et al. (2000b) performed the first unambiguous detection of Pluto’s 60- and 100-μm lightcurve, and the first detection of its 150- and 200-μm emission. These observations definitely demonstrated that Pluto’s surface is non-isothermal. The infrared lightcurve is globally anti-correlated with the visible lightcurve, but not exactly. This lag was attributed to thermal inertia effects. Lellouch et al. (2000b) developed a model with three different geographical units that fits the complete set of visible and infrared observations. The three units are constituted of N2-ice, CH4-ice, and tholins. The latter two units show a small thermal inertia, 1.5 × 104–10 × 104 erg cm−2 s−1/2 K−1, and a high emissivity. This suggests that the shallow layer is made of porous ice. In this respect, Pluto seems similar to the icy satellites in the outer Solar System.

8. Conclusions

ISO observations of planetary objects proved extremely fruitful. ISO sensitivity and spectral resolutions led to the first detection of several chemical species, and to the development of new methods to probe planetary atmospheres. This brought new views on the formation, dynamics and chemistry of planets, and raised new questions. The work undertaken with ISO observations will be carried on by several infrared instruments aboard Earth orbiting or interplanetary missions: the Spitzer, Hershell and Cassini missions, several martian missions. In particular, Hershell will continue the ISO legacy by measuring more precisely the giant planets composition in deuterium and helium through the observations of the HD rotational lines and the H2 collision-induced continuum. Hershell, and Cassini specifically on the Saturn System, will bring more information on exogenic water in the outer Solar System, in order to sort out the various possible sources. Laboratory work PLANETS AND TITAN 137 will be done to better constrain chemical reactions in the conditions relevant to planetary atmospheres, especially the methyl recombination reaction and the benzene production. In association with the wealth of infrared observations provided by Cassini, we can expect significant progresses in our understanding of the hydrocarbon photochemistry in the giant planets and Titan. On Mars, Mars Express will look for carbonates and monitor the dust opacity using new ways opened by ISO.

References

Abbas, M. M., LeClair, A., Owen, T., Conrath, B. J., Flasar, F. M., Kunde, V.G., et al.: 2004, Astrophys. J. 602, 1063. Aléon,J. and Robert, F.: 2004, Icarus 167, 424. Altwegg, K. and Bockelée-Morvan, D.: 2003, Space Sci. Rev. 106, 139. Baines, K. H., Carlson, R. W., and Kamp, L. W.: 2002, Icarus 159, 74. Bandfield, J. L., Glotch, T. D., and Christensen, P. R.: 2003, Science 301, 1084. Bézard, B.: 2000, in Salama, A., Kessler, M. F., Leech, K., and Schulz, B. (eds.), ISO Beyond the Peaks, ESA Publication Division (ESA-SP456), Noordwijk, p. 3. Bézard, B., Gautier, D., and Marten, A.: 1986, Astron. Astrophys. 161, 387. Bézard, B., Feuchtgruber, H., Moses, J. I., and Encrenaz, T.: 1998, Astron. Astrophys. 334, L41. Bézard, B., Romani, P. N., Feuchtgruber, H., and Encrenaz, T.: 1999, Astrophys. J. 515, 868. Bézard, B., Drossart, P., Encrenaz, T., and Feuchgruber, H.: 2001, Icarus 154, 495. Bézard, B., Lellouch, E., Strobel, D., Maillard, J.-P., and Drossart, P.: 2002, Icarus 159, 95. Brooke, T. Y., Knacke, R. F., Encrenaz, T., Drossart, P., Crisp, D., and Feuchtgruber, H.: 1998, Icarus 136,1. Burgdorf, M. J., Encrenaz, T., Lellouch, E., Feuchtgruber, H., Davis, G. R., Swinyard, B. M., et al.: 2000, Icarus 145, 79. Burgdorf, M. J., Orton, G. S., Davis, G. R., Sidher, S. D., Feuchtgruber, H., Griffin, M. J., et al.: 2003, Icarus 164, 244. Burgdorf, M. J., Orton, G. S., Encrenaz, T., Davis, G. R., Sidher, S. D., Lellouch, E., et al.: 2004, Planet. Space Sci. 52, 379. Charnley, S. B. and Rodgers, S. D.: 2002, Astrophys. J. 569, L133. Conrath, B. J., Gautier, D., Lindal, G. F., Samuelson, R. E., and Shaffer, W. A.: 1991, J. Geophys. Res. 96, 18907. Coustenis, A., Bézard,B., and Gautier, D.: 1989, Icarus 82, 67. Coustenis, A., Salama, A., Lellouch, E., Encrenaz, T., Bjoraker, G. L., Samuelson, R. E., et al.: 1998, Astron. Astrophys. 336, L85. Coustenis, A., Salama, A., Schulz, B., Ott, S., Lellouch, E., Encrenaz, T., et al.: 2003, Icarus 161, 383. Davis, G. R., Griffin, M. J., Naylor, D. A., Oldham, P. G., Swinyard, B. M., Ade, P. A. R., et al.: 1996, Astron. Astrophys. 315, L393. de Graauw, T., Feuchtgruber, H., Bézard,B., Drossart, P., Encrenaz, T., Beintema, D. A., et al.: 1997, Astron. Astrophys. 321, L13. Drossart, P., Fouchet, T., Crovisier, J., Lellouch, E., Encrenaz, T., Feuchtgruber, H., et al.: 1999, in Cox, P., and Kessler, M. F. (eds.), The Universe as Seen by ISO, ESA Publication Division, Noordwijk, p. 169. Drouart, A., Dubrulle, B., Gautier, D., and Robert, F.: 1999, Icarus 140, 129. Encrenaz, T.: 2003, Planet. Space Sci. 51, 89. 138 T. FOUCHET ET AL.

Encrenaz, T., Combes, M., and Zéau,Y.: 1978, Astron. Astrophys. 70, 29. Encrenaz, T., de Graauw, T., Schaeidt, S., Lellouch, E., Feuchtgruber, H., Beintema, D. A., et al.: 1996, Astron. Astrophys. 315, L397. Encrenaz, T., Feuchtgruber, H., Atreya, S. K., Bézard,B., Lellouch, E., Bishop, J., et al.: 1998, Astron. Astrophys. 333, L43. Encrenaz, T., Schulz, B., Drossart, P., Lellouch, E., Feuchtgruber, H., and Atreya, S. K.: 2000, Astron. Astrophys. 358, L83–L87; erratum, 362, 1174. Fedorova, A. A., Lellouch, E., Titov, D. V., de Graauw, T., and Feuchtgruber, H.: 2002, Planet. Space Sci. 50,3. Feuchtgruber, H., Lellouch, E., de Graauw, T., Bézard,B., Encrenaz, T., and Griffin, M.: 1997, Nature 389, 159. Feuchtgruber, H., Lellouch, E., Bézard,B., Encrenaz, T., de Graauw, T., and Davis, G. R.: 1999a, Astron. Astrophys. 341, L17. Feuchtgruber, H., Lellouch, E., Encrenaz, T., Bézard,B., Coustenis, A., Drossart, P., et al.: 1999b, in Cox, P., and Kessler, M. F. (eds.), The Universe as Seen by ISO, ESA Publication Division, Noordwijk, p. 133. Folkner, W. M., Woo, R., and Nandi, S.: 1998, J. Geophys. Res. 103, 22847. Fouchet, T., Lellouch, E., Bézard,B., Encrenaz, T., Drossart, P., Feuchtguber, H., et al.: 2000a, Icarus 143, 223. Fouchet, T., Lellouch, E., Bézard,B., Feuchtgruber, H., Drossart, P., and Encrenaz, T.: 2000b, Astron. Astrophys. 355, L13. Fouchet, T., Lellouch, E., and Feuchtgruber, H.: 2003, Icarus 161, 127. Fouchet, T., Irwin, P. G. J., Parrish, P., Calcutt, S. B., Taylor, F. W., Nixon, C. A., et al.: 2004, Icarus 172, 50. Geiss, J. and Gloeckler, G.: 1998, Space Sci. Rev. 84, 239. Griffin, M. J., Naylor, D. A., Davis, G. R., Ade, P. A. R., Oldham, P. G., Swinyard, B. M., et al.: 1996, Astron. Astrophys. 315, L389. Hashizume, K., Chaussidon, M., Marty, B., and Robert, F.: 2000, Science 290, 1142. Hersant, F., Gautier, D., and Huré,J.-M.: 2001, Astrophys. J. 554, 391. Hersant, F., Gautier, D., and Lunine, J. I.: 2004, Planet. Space Sci. 52, 623. Hueso, R., Sánchez-Lavega, A., and Drossart, P.: 1999, Planet. Space Sci. 47, 1263. Kallenbach, R., Geiss, J., Ipavich, F. M., Gloeckler, G., Bochsler, P., Gliem, F., et al.: 1998, Astrophys. J. Lett. 507, L185. Kerridge, J. F.: 1993, Rev. Geophys. 31, 423. Kim, S. J., Caldwell, J., Rivolo, A. R., Wagener, R., and Orton, G. S.: 1985, Icarus 64, 233. Lellouch, E.: 1999, in Cox, P. and Kessler, M. F. (eds.), The Universe as Seen by ISO, ESA Publication Division, Noordwijk, p. 125. Lellouch, E., Encrenaz, T., de Graauw, T., Erard, S., Morris, P., Crovisier, J., et al.: 2000a, Planet. Space Sci. 48, 1393. Lellouch, E., Laureijs, R., Schmitt, B., Quirico, E., de Bergh, C., Crovisier, J., et al.: 2000b, Icarus 147, 220. Lellouch, E., Bézard,B., Fouchet, T., Encrenaz, T., Feuchtgruber, H., and de Graauw, T.: 2001, Astron. Astrophys 370, 610. Lellouch, E., Bézard,B., Moses, J. I., Davis, G. R., Drossart, P., Feuchtgruber, H., et al.: 2002, Icarus 159, 112. Linsky, J. L.: 1998, Space Sci. Rev. 84, 285. Lunine, J. I., Yung, Y. L., and Lorenz, R. D.: 1999, Planet. Space Sci. 47, 1291. Marley, M. S.: 1999, in Weissman, P. R., McFadden, L.-A., and Jonhson, T. V. (eds.), Encyclopedia of the Solar System, Acadamic Press, New York. Moreno, R., Marten, A., Matthews, H. E., and Biraud, Y.: 2003, Planet. Space Sci. 51, 591. PLANETS AND TITAN 139

Moses, J. I., Lellouch, E., Bézard,B., Gladstone, G. R., Feuchtgruber, H., and Allen, M.: 2000a, Icarus 145, 166. Moses, J. I., Bézard,B., Lellouch, E., Gladstone, G. R., Feuchtgruber, H., and Allen, M.: 2000b, Icarus 143, 244. Moses, J. I., Fouchet, T., Bézard,B., Lellouch, E., Gladstone, G. R., Feuchtgruber, H., et al.: 2001, Bull. Am. Astron. Soc. 33, 1044. Mousis, O., Gautier, D., and Coustenis, A.: 2002, Icarus 159, 156. Niemann, H. B., Atreya, S. K., Carignan, G. R., Donahue, T. M., Haberman, J. A., Harpold, D. N., et al.: 1998, J. Geophys. Res. 103, 22831. Orton, G. S., Lacy, J. H., Achtermann, J. M., Parmar, P., and Blass, W. E.: 1992, Icarus 100, 541. Owen, T., Mahaffy, P. R., Niemann, H. B., Atreya, S., and Wong, M. : 2001, Astrophys. J. 553, L77. Pillinger, C. T.: 1984, Geochim. Cosmochim. Acta 48, 2739. Podolak, M., Weizman, A., and Marley, M.: 1995. Planet. Space Sci. 43, 1517. Schulz, B., Encrenaz, T., Bézard,B., Romani, P. N., Lellouch, E., and Atreya, S. K.: 1999, Astron. Astrophys. 350, L13. Sidher, S. D., Griffin, M. J., Davis, G. R., Swinyard, B. M., Burgdorf, M., Orton, G. S., et al.: 2000, Icarus 147, 35. Showman, A. P. and Dowling, T. E.: 2000, Science 289, 1737. Smith, W. H., Schempp, W. V., Simon, J., and Baines, K. H.: 1989, Astrophys. J. 336, 962. Tokunaga, A. T., Knacke, R. F., and Ridgway, S. T.: 1980, Icarus 44, 93. Wong, A.-S., Lee, A. Y. T., Yung, Y. L., and Ajello, J. M.: 2000, Astrophys. J. Lett. 534, L215. COMETS, ASTEROIDS AND ZODIACAL LIGHT AS SEEN BY ISO

THOMAS G. MULLER¨ 1,∗,PETER´ ABRAH´ AM´ 2 and JACQUES CROVISIER3 1Max-Planck-Institut fur¨ extraterrestrische Physik, Giessenbachstraße, 85748 Garching, Germany 2Konkoly Observatory, H-1525 P.O. Box 67, Budapest, Hungary 3LESIA, Observatoire de Paris, 5 place Jules Janssen, F-92195 Meudon, France (∗Author for correspondence: E-mail: [email protected])

(Received 13 July 2004; Accepted in ﬁnal form 15 December 2004)

Abstract. ISO performed a large variety of observing programmes on comets, asteroids and zodiacal light – covering about 1% of the archived observations – with a surprisingly rewarding scientiﬁc return. Outstanding results were related to the exceptionally bright comet Hale–Bopp and to ISO’s capability to study in detail the water spectrum in a direct way. But many other results were broadly recognised: Discovery of new molecules in comets, the studies of crystalline silicates, the work on asteroid surface mineralogy, results from thermophysical studies of asteroids, a new determination of the asteroid number density in the main-belt and last but not least, the investigations on the spatial and spectral features of the zodiacal light.

Keywords: ISO, asteroids, comets, SSO (solar system objects), zodiacal light, dust

1. Introduction

The comet, asteroid and zodiacal light programme of ISO – without the planets – covered approximately 325 h, less than 1% of all ISO observations. Thus, the solar system community did not expect much from this small number of observations with a 60-cm telescope circulating just outside the Earth’s atmosphere. Yet, the ISO results made a big impact on our knowledge of objects on our celestial “doorstep”. Eleven Guaranteed Time projects and 31 Open Time projects were either dedicated to comets, asteroids and zodiacal light studies or included solar system science aspects. Surveys with observations close to the ecliptic plane, performed by ISOCAM and ISOPHOT, include many asteroids and several comets. In addition, the parallel and serendipity data – taken in non-prime instrument modes or while the satellite was slewing from one target to the next – contain by chance many interesting solar system sources (Müller, 2001). Even standard background observations can be considered as absolute measurements of the zodiacal light and have been added to the database of dedicated studies of the dust in the solar neighbourhood. First inventories of the ISO solar system contributions have been made by Encrenaz (2000), Fouchet (2003) and Müller(2003a). The latter two authors went through all solar system ISO projects and presented the proposed original ideas and comprared them with the published scientific results (as of June 2002). They also identified promising future fields on solar system science with ISO. There, one can Space Science Reviews (2005) 119: 141–155 DOI: 10.1007/s11214-005-8067-9 C Springer 2005 142 T. G. MULLER¨ ET AL. also find details on the proposals, target lists and instructions on how to search for specific moving objects in the ISO Data Archive. The results on the “life-element” water in the universe – not least its unexpected discovery in large amounts in the higher atmospheres of all giant planets – together with other highlights on planet science are covered by the contributions by Cernicharo and Crovisier (2005) and Fouchet (2005) in this volume. Different states of silicate dust under different conditions are the basis for the new research field on astro-mineralogy. Crystalline silicates are found around young and old stars, in proto-planetary disks and are also very common in our own solar system (e.g., in Hale–Bopp). The chapter on crystalline silicates (Molster, 2005, this volume) touches therefore solar system science, but from a different point of view. Here, the remaining part of the solar system field is followed up and ISO’s results on asteroids, comets and extended solar system structures are presented.

2. Comets

A total of 20 comets were observed with ISO (some of them with its four instruments): 2P/Encke, 7P/Pons-Winnecke, 22P/Kopff, 29P/Schwass- mann-Wachmann 1, 30P/Reinmuth 1, 32P/Comas Sola, 43P/Wolf-Harrington, 45P/Honda–Mrkos-Pajduˇsáková, 46P/Wirtanen, 55P/Tempel-Tuttle, 65P/Gunn, 81P/Wild 2, 95P/Chiron, 103P/Hartley 2, 117P/Helin-Roman-Alu 1, 126P/IRAS, 128P/Shoemaker-Holt 1B, 133P/Elst-Pizarro, and the long-period comets C/1995 O1 (Hale–Bopp) and C/1996 B2 (Hyakutake). A large part of the cometary programme was devoted to the exceptionally bright comet C/1995 O1 (Hale–Bopp), as a target of opportunity, resulting in a very rewarding scientific return. Unfortunately, ISO could not observe this comet at heliocentric distances smaller than 2.8 AU due to solar elongation constraints. Several observations of the fainter comets led to limited scientific results, which are still unpublished.

2.1. SPECTROSCOPY OF GAS SPECIES

The fundamental bands of vibration of volatile species released from the ices of comet nuclei are emitting in the mid-IR through ﬂuorescence excited by the Sun. High-resolution spectroscopy of comet Hale–Bopp (in September–October 1996, when it was at r = 2.8 AU from the Sun) with the SWS and LWS resulted in a detailed study of the water spectrum. The 6 μm band and several rotational lines (100–180 μm) of H2O were observed for the ﬁrst time in a comet. The bands around 2.7 μm could be studied in detail, resulting in estimates of the rotational and spin temperatures of water (Crovisier et al., 1997a,b, 1999b; Figure 1; see also Cernicharo and Crovisier, 2005). COMETS, ASTEROIDS AND ZODIACAL LIGHT AS SEEN BY ISO 143

Figure 1. The SWS spectrum of water in the 3 μm region observed in comets C/1995 O1 (Hale–Bopp) at 2.8 AU from the Sun, 103P/Hartley 2 at 1.0 AU, and 22P/Kopff at 1.9 AU (scale expanded by ×5). Adapted from Crovisier et al. (1997b, 1999a).

Carbon monoxide and dioxide were also easily detected. These very volatile species are the only ones which are observed far from the Sun (r = 4.6 AU). CO2 was very conspicuous from its strong ν3 band at 4.26 μm in the ISOPHOT-S spectra. This detection of CO2 by ISO (which was repeated in 103P/Hartley 2) is, besides the in situ observation in comet Halley, the only direct observation of this molecule in a comet. The Q-branch of the ν3 band of CH4 was also detected by the SWS at 3.31 μm (Crovisier, 2000a; Crovisier et al., in preparation). Evolution of the outgassing of H2O, CO and CO2 in comet Hale–Bopp could be investigated between 4.9 and 2.8 AU, pre- and post-perihelion, showing the transition from the CO-dominated to the H2O-dominated sublimation regimes (Crovisier et al., 1996, 1997b, 1999b; Crovisier, 2000b; Leech et al., 1999). H2O, CO and CO2 appear to be the most abundant volatile molecules, with relative production rates H2O:CO:CO2 equal to 100:70:20 at r = 2.9 AU in comet Hale–Bopp. The spectrum of the Jupiter-family comet 103P/Hartley 2 was investigated by ISOPHOT, SWS and ISOCAM, allowing detection of H2O, CO2 (but not CO) and crystalline silicates (Colangeli et al., 1999; Crovisier et al., 1999a, 2000; Figure 1). Long integration times were also spent on 22P/Kopff, but this comet was much fainter and only the water lines could be detected by the SWS (Crovisier et al., 1999a; Figure 1). 144 T. G. MULLER¨ ET AL.

The rotational temperatures retrieved for water in comets Hale–Bopp and Hart- ley 2 are 28 and 20 K, respectively. These cold temperatures agree with models of coma hydrodynamics and water excitation, which predict a rotational relaxation of water. The intensities of the ortho and para transitions in the ν3 band of water also permit an evaluation of the ortho-to-para ratio and of the water spin temperature. Tspin = 28 and 35 K were determined for the two comets, respectively. The mean- ing of the water Tspin is still to be understood (Crovisier et al., 1997b, 1999a,b; Crovisier et al., 2000; Cernicharo and Crovisier, 2005).

2.2. SPECTROSCOPY OF MINERALS

The 6–45 μm SWS spectrum of comet Hale–Bopp revealed emission features attributed to minerals in cometary dust (Figure 2). A ﬁrst look lead to the identiﬁcation of crystalline, Mg-rich olivine (forsterite) (Crovisier et al., 1997b). A more thorough investigation showed that crystalline pyroxenes and amorphous silicates are also present (Crovisier et al., 2000). No PAH signature could be found. This ISO spectrum of comet Hale–Bopp was the subject of many studies, in comparison with complementary spectra observed from the ground at other heliocentric distances

Figure 2. The 6–45 μm SWS spectrum of comet C/1995 O1 (Hale–Bopp) observed at 2.8 AU from the Sun, decomposed into several emission components: 22% of forsterite (Cry Ol), 8% of ortho- pyroxene (Cry o-Pyr), 70% amorphous silicates (Am Pyr), and two blackbody components (280 and 165 K). From Crovisier et al. (2000). COMETS, ASTEROIDS AND ZODIACAL LIGHT AS SEEN BY ISO 145 and comparison with laboratory spectra of various minerals (e.g., Wooden et al., 1999; Brucato et al., 1999; Harker et al., 2002). By probing the mid-IR spectroscopy of cosmic dust, ISO opened the new ﬁeld of astro-mineralogy, allowing us to study the cycle of dust in the Universe and in planetary systems. The resemblance between the spectra of cometary dust and of debris discs of young stars or of some Herbig Ae/Be stars observed by ISO was noted in many publications (e.g., Malfait et al., 1998; Molster et al., 1999; Waelkens et al., 1999; Crovisier, 2000b; Bouwman et al., 2003; Molster, 2005). However, the origin of crystalline silicates in comets is still an open issue. Silicates observed in interstellar clouds are amorphous. Annealing could occur in the hot inner Solar Nebula, in regions where cometary ices could not be present. Turbulent diffusion within the primitive Solar Nebula might be an explanation (Bockel´ee-Morvan et al., 2002). Crystalline water ice was shown to be present in the dust of comet Hale–Bopp at 2.9 AU from the Sun, from its signatures at 44 and 65 μm and possibly at 3.1 μm, observed by the LWS and ISOPHOT-S (Lellouch et al., 1998).

2.3. PHOTOMETRY AND IMAGING

The spectral energy distribution of several comets was observed by broad band photometry from 3.6 to 170 μm with ISOPHOT. For comet Hale–Bopp, which was the subject of five observations between r = 4.6 and 2.8 AU, an extensive analysis was possible (Grün et al., 2001). The colour temperature, dust production, dust-to- gas ratio and their evolution with heliocentric distance could be studied. The 7.3– 12.8 μm colour temperature is ≈1.5 times the expected blackbody temperature, pointing to the presence of small, superheated dust grains. Silicate features are always present and influence the emissions in the 10 and 25 μm region. Water ice grains are suggested at r = 4.6 AU from enhanced emission at 60 μm. Cometary dust and its distribution was studied by ISOCAM imaging of comets 2P/Encke, 46P/Wirtanen, 65P/Gunn and 103P/Hartley 2 (Colangeli et al., 1998; Epifani et al., 2001), allowing the authors to model the comet tail, the evolution of dust coma, grain sizes and velocities and to infer implications for meteoroid streams. Observations of comets in the thermal infrared with ISOCAM also allow one to separate the emission of the nucleus from that of dust. This was applied to C/1995 O1 (Hale–Bopp), 2P/Encke, 22P/Kopff, 46P/Wirtanen, 55P/Tempel-Tuttle, 95P/Chiron, 103P/Hartley 2 and 126P/IRAS (Fernández et al., 2000; Groussin et al., 2004a,b; Jorda et al., 2000; Lamy et al., 2002). When it is possible to combine such data with HST observations in the visible, the albedo and size of the nucleus can be securely determined (this was possible for 2P/Encke, 22P/Kopff, 55P/Tempel- Tuttle and C/1995 O1 (Hale–Bopp)). Thus ISO significantly contributed to the database of cometary nuclei physical properties (Lamy et al., 2004). Cometary trails were first discovered by IRAS. Their analysis in terms of mass loss rates and of dynamical evolution is crucial for our understanding of 146 T. G. MULLER¨ ET AL.

Figure 3. A mosaic image of comet 2P/Encke from ISOCAM. The comet tail shows up as a broad feature in the centre. The comet trail is the thin feature stretching diagonally from the lower left to the upper right. From Reach et al. (2000). the origins of insterstellar dust particles and of meteoroids. ISO investigated the trails of 22P/Kopff (Davies et al., 1997) and 2P/Encke (Reach et al., 2000; Figure 3). The trail of comet Hale–Bopp was even observed in the course of the ISOPHOT Serendipity Survey at 170 μm(M¨uller et al., 2002).

3. Asteroids

The ISO programmes included 40 different asteroids, among them 9 targets which are considered as transition objects between comets and asteroids. The whole asteroid programme was performed in less than 100 h, using all four instruments of ISO. Details on the proposals, the targets and the observing strategies and goals can be found in Müller(2003a). The main goals of the asteroid observations were the identification of surface minerals, composition, the connection to meteorites and comets, surface alteration processes and the interpretation of taxonomic classes through the identification of mid-infrared features of well-known minerals and COMETS, ASTEROIDS AND ZODIACAL LIGHT AS SEEN BY ISO 147 meteorites. The predominant observing technique in the asteroid programme was therefore spectroscopy.

3.1. SURFACE MINERALOGY

First results on ISOPHOT-S observations of bright main-belt asteroids (Dotto et al., 1999; Dotto et al., 2000; M¨uller et al., 2000) suggested the presence of silicates on the surface of all asteroids. Barucci et al. (2002) extensively studied 10 Hygiea and demonstrated the possible spectral similarity with CO carbonaceous chondrites at small grain sizes (Figure 4). Dotto et al. (2002) described the similarities between low-albedo asteroids and different meteorites which have been analysed in new laboratory experiments. The ISOPHOT-Sspectra seem to conﬁrm the theories about the thermal history of C-type

1.6

Hygiea

1.4

Warrenton 1.2 by ASTER Emissivity

Ornans 1.0 this paper

0.8

5 10 15 20 25 wavelength [μm]

Figure 4. Comparison between the ISO spectrum of 10 Hygiea and different meteorite samples. Adapted from Barucci et al. (2002). 148 T. G. MULLER¨ ET AL. asteroids by the match between the spectrum of 511 Davida and the emissivity of the CM-type carbonaceous chondrite meteorite Murchison. Müllerand Blommaert (2004) analysed the emissivity of 65 Cybele, which is a dark, very low albedo asteroid. Cybele shows an emissivity increase between 8.0 and 9.5 μm, which is close in wavelength to the Christiansen feature of CO and CO3 types of the carbonaceous chondrites. Very recent studies of the dark asteroid Polyxo, which showed in earlier investigations similarities with the Tagish lake meteorite (Hiroi and Hasegawa, 2003), were carried out by Dotto et al. (2004). They suggest a tentative spectral similarity with the Ornans meteorite and cannot confirm the analogy with the Tagish Lake meteorite in the measured wavelength range between 7 and 26 μm. All mid-infrared spectroscopic studies clearly reveal the difficulties of identifying mineralogic features based on disk-integrated, mid-IR spectroscopy. A wide wavelength range with uniform data quality covering different rotation phases, a clean calibration and a careful modelling of the underlying thermal emission are necessary to recognise signatures of the surface material. In addition, laboratory spectra of different meteorites at various grain sizes are needed to identify the low level, broad band features. But even if all prerequisites are fulfilled, the spectroscopic interpretation still remains difficult.

3.2. THERMOPHYSICAL STUDIES

A new thermo-physical model for asteroids (TPM, Lagerros, 1996, 1997, 1998) was developed during the ISO mission. This new TPM allows the detailed characterisation of the thermal behaviour of the surface and the derivation of physical parameters, like surface roughness or porosity (e.g., Müller et al., 1999a; Müller, 2002) from accurate thermal observations. Through large samples of photometric measurements at different wavelengths and observing geometries it was also possible to determine the thermal inertia and thermal conductivity for a few large main-belt asteroids (Müllerand Lagerros, 1998; Müller et al., 1999a): The thermal behaviour of many large asteroids can be explained by a regolith covered surface with a very low thermal inertia of about = 15Jm−2 s−0.5 K−1.Müller and Lager- ros (2002) demonstrated that the TPM matches all kind of thermal IR measurements from different observing geometries with a high accuracy. This quality proof for a sample of well-known asteroids was the prerequisite for applying the model to arbitrary thermal measurements. Depending on the a priory knowledge on a specific asteroids, the TPM allows to derive a variety of properties: From a simple diameter and albedo determination up to a full characterisation of the surface regolith (Müller, 2003b). For example, Müller et al. (2002) used 170 μm flux densities to derive simple size-albedo solutions for serendipitously seen asteroids which, apart from orbital elements and visual brightness, had no existing physical characterisation before. For 65 Cybele, on the other hand, a large database with visual and thermal observations was the basis for a sophisticated thermophysical analysis of the surface regolith (Müllerand Blommaert, 2004). The authors derived a size of COMETS, ASTEROIDS AND ZODIACAL LIGHT AS SEEN BY ISO 149

Figure 5. The radiometric diameter and albedo values for 65 Cybele calculated through the STM and the TPM. The observations are well-calibrated ISOCAM ﬁlter measurements. The STM results show clear trends with wavelength while the TPM produced a unique diameter/albedo solution. Adapted from M¨ullerand Blommaert (2004).

302×290×232 km (±4%) and a geometric visible albedo of 0.050±0.005 through a careful TPM analysis of the existing thermal observations. They also described the thermal behaviour in detail with respect to phase angle dependence, before/after opposition effects and wavelength dependencies. Figure 5 shows a comparison between the radiometric diameter-albedo determination using the TPM and the widely used Standard Thermal Model (STM; Lebofsky et al., 1986). The STM clearly shows unphysical, wavelength-dependent diameter and albedo values. The TPM efforts resulted in a convincing concept to use a few, well-known, large main-belt asteroids as new far-infrared photometric calibrators for ISO (M¨uller and Lagerros, 1998, 2002, 2003). The TPM predictions for 5–10 asteroids, which have large sets of independently calibrated infrared, submillimetre and millimetre observations, are now widely used to calibrate and characterise far-infrared and submillimetre instruments. Work is ongoing to extend this list to fainter targets in preparation for the ASTRO-F and HERSCHEL missions.

3.3. MISCELLANEOUS PROGRAMMES

3.3.1. Asteroid Number Density The superb mid-infrared sensitivity of ISO was used to determine the number of main-belt asteroids for the first time in a direct way (Tedesco and Desert, 2002). Twenty-two moving sources appeared in a 15 square field of repeated, 150 T. G. MULLER¨ ET AL. deep ISOCAM observations at 12 μm in the ecliptic plane. Most of the asteroids detected have flux densities less than 1 mJy, about 150–350 times fainter than any of the asteroids observed by IRAS. The study revealed 160 ± 20 asteroids per square degree at a 0.6 mJy detection limit. Based on a statistical asteroid model, the authors concluded that there are about 1.2 ± 0.5 × 106 asteroids (≥1 km in diameter) in our solar system, twice as many as previously believed.

3.3.2. Surface Properties through Thermal Polarimetry In a small programme Lagerros et al. (1999) and M¨uller et al. (1999b) investigated the surface properties of the asteroids 6 Hebe and 9 Metis through thermal polarimetry at 25 μm. Solid upper limits were interpreted through an extended version of the TPM to cover also polarimetric aspects. The Metis observations favored a low refractive index and high surface roughness, but the Hebe observations were inconclusive since they coincided with a minimum in the polarization curve.

3.3.3. NEOs and TNOs Two asteroid programmes performed observations on asteroids outside the main- belt-Jupiter region. Harris and Davies (1999) combined ISOPHOT observations with mid-infrared ground-based observations to investigate radiometric diameter and albedo values for three near-Earth asteroids. The original idea to search for cometary activity in near-Earth asteroids with ISOCAM turned out to be too am- bitious. The best achieved sensitivity corresponds to a mass-loss rate of above 0.1 kg s−1, which is the expected upper limit for near-Earth cometary asteroids. The second, very large programme with almost 28 h of ISO time was dedicated to 7 Kuiper-belt objects (Thomas et al., 2000). The authors reported marginal detections of the two trans-Neptunian objects (TNO) 1993 SC and 1996 TL66. Their radiometric modelling attempts indicated that the TNOs are large, spherical and very dark objects. The success of this programme was clearly limited by the ISOCAM and ISOPHOT sensitivities.

3.3.4. Serendipitous Observations Additionally to the dedicated asteroid programmes, many comets and asteroids have been seen serendipitously in the ISOCAM Parallel Survey, in the ISOPHOT Serendipity Survey and in large survey programmes, like e.g. the galactic plane survey ISOGAL (Müller, 2001). The ISOPHOT Serendipity slews went over 16 different bright asteroids, 2 planets and 7 comets (Müller et al., 2002). The resulting 170 μm fluxes served two purposes: (i) improvement of the flux calibration through a few well-known asteroids and planets; (ii) scientific investigations in terms of diameter-albedo determination and comet modeling. The nine slews over Hale– Bopp showed that the large particles are concentrated in the nucleus-trail transition region. The extraction of solar system targets for several other surveys is ongoing and interesting results can be expected. COMETS, ASTEROIDS AND ZODIACAL LIGHT AS SEEN BY ISO 151

4. Zodiacal Light

During its 71.9 h of dedicated programmes, ISO observed different extended structures of the solar system: zodiacal light, dust debris in the outer solar system and the Earth’s dust ring, comet tails and trails and planetary rings (for an inventory see M¨uller, 2003a). Three main goals were set by the observers: (i) reﬁnement of the post-DIRBE picture on the global structure of the Interplanetary Dust Cloud; (ii) study of sub-structures in the Zodiacal light; and (iii) reliable determination of the mid-infrared spectrum of the Zodiacal emission.

4.1. SPATIAL STRUCTURES

Technically, point (ii) is the least challenging, since differential – rather than absolute – photometric measurements, as well as images, can be utilised. With the aim of exploring how comets release fresh dust into the interplanetary space, Reach et al. (2000) observed Comet 2P/Encke with ISOCAM (see Figure 3). They detected a long, straight dust trail composed of large cm-sized grains following the orbit of the comet. These grains will become part of the Encke meteoroid stream, and – on longer term – of the zodiacal cloud. Computing the total mass loss from the 1997 apparition of Encke, the authors predict a limited (several thousands of years) lifetime for the comet. Abrahám´ et al. (2002) observed a similar trail of 22P/Kopff at longer wavelengths with ISOPHOT, and detected its emission at 0.25◦ behind the comet at 12 and 25 μm but not at 60 μm, consistent with a temperature of 215 K. The trail was only marginally seen at larger distances, showing that its intensity drops quickly behind the head. Abrahám´ et al. (2002) also studied the brightness profiles of the asteroidal bands parallel to the Ecliptic, and determined an upper limit for the intrinsic arcminute scale brightness variations of the bands. Searching for similar variations in the general zodiacal light distribution may also be interesting, since if the interplanetary dust was replenished from discrete sources – asteroids, comets – a structureless zodiacal light is difficult to explain. From ISOPHOT measurements at 25 μm, however, Abrahám´ et al. (1997b) deduced an upper limit for the underlying r.m.s. brightness fluctuations of ±0.2%, which corresponds at high ecliptic latitudes to ±0.04 MJy sr−1.

4.2. SPECTRAL CHARACTERISATION

The mid-infrared spectrum of the Zodiacal light, not known before ISO, was mea- surable with both ISOCAM-CVF and ISOPHOT-S. However, the technical difficulties involved are indicated by the fact that both instrument groups needed to revise their first results due to straylight contamination (ISOCAM) and to incorrect beam profile (ISOPHOT). The spectra derived from 29 ISOPHOT-S observations (Abrahám´ et al., 1997a; Leinert et al., 2002) could be well represented by black- body radiation in the 255–300 K range (see Figure 6). No spectral features in the 152 T. G. MULLER¨ ET AL.

Figure 6. Upper panel: Colour temperatures derived from the ISOPHOT-S templates using single temperature blackbody ﬁts. Lower pannel: Colour temperatures predicted by a model in the spectral range of ISOPHOT-S. Adapted from Leinert et al. (2002).

6−12 μm range were visible at the 10–20% level. With the ISOCAM-CVF, Reach et al. (1996, 2003) and Reach (1997) also found blackbody-shaped spectra – though their temperatures are somewhat lower than those from ISOPHOT-S – indicating large (>10 μm radius), low-albedo, rapidly rotating grey particles. Comparison to theoretical models provided acceptable ﬁts only for ‘astronomical silicates’, ruling out many other potential components of the zodiacal light. The ISOCAM spectra give also a hint for a 9–11 μm emission feature with an amplitude of 6% of the continuum. This feature might indicate the presence of a mixture of small (∼1 μm) amorphous, crystalline and hydrous silicates. The large scale temperature distribution on the sky was derived consistently from both instruments’ data sets: the temperature value increases with the ecliptic latitude and decreases with the distance from the sun (Figure 6). This result can be understood in terms of the geometry of the interplanetary dust cloud.

4.3. ABSOLUTE SURFACE BRIGHTNESS

The most challenging part of ISO’s Zodiacal light programme is no doubt the analysis of the absolute surface brightness measurements. This is still an on-going COMETS, ASTEROIDS AND ZODIACAL LIGHT AS SEEN BY ISO 153 project involving determination of the photometric zero points of the instruments, linearity and beam profile issues. Abrahám´ et al. (2002) discussed the possible advantages of ISOPHOT measurements of this type in refining the present picture on the large scale structure of the Zodiacal cloud be produced by DIRBE (e.g. small beam avoiding point sources; observations from September to December when DIRBE was not cooled; more far-infrared filters). They demonstrated that the ISOPHOT and DIRBE observations can be transformed to each other’s photometric systems (see also Holmes and Dermott, 2000). Thus once the absolute calibration of ISOPHOT is set, large samples of dedicated and serendipitous observations of the sky at different latitudes and at different solar elongations contained in the ISO Data Archive could be analysed. Improvements of the existing zodiacal light models might be possible and new structures, like asteroidal bands, could still be hidden in the observations. And it is also worth to mention that an accurate determination of the absolute level of the Zodiacal light is a strong prerequisite to extract the value of the infrared extragalactic background light, and that many of the absolute photometric data of ISO will not be superseeded by the coming new infrared satellites.

Acknowledgements

This work was partly supported by the grant OTKA 037508 of the Hungarian Re- search Fund. P. A.´ acknowledges the support of the Bolyai Fellowship. This work is based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, the Netherlands and the United Kingdom) and with the participation of ISAS and NASA.

References

Abrahám,P.,´ Acosta-Pulido, J. A., Klaas, U., Laureijs, R. J., Leinert, Ch., Lemke, D., et al.: 1997a, in A. M. Heras, K. Leech, N. R. Trams, and M. Perry (eds.), First Workshop on Analytical Spectroscopy, ESA SP-419, 119. Abrahám,P.,´ Leinert, C., and Lemke, D.: 1997b, A&A 328, 702. Abrahám,P.,´ Leinert, C., and Lemke, D.: 2002, in S. F. Green, I. P. Williams, J. A. M. McDonnell and N. McBride (eds.), Dust in the Solar System and Other Planetary Systems, COSPAR colloquia series, Vol. 15, p. 111. Barucci, M. A., Dotto, E., Brucato, J. R., et al.: 2002, Icarus 156, 202. Bockelée-Morvan, D., Gautier, D., Hersant, F., Huré,J.-M., and Robert, F.: 2002, A&A 384, 1107. Bouwman, J., de Koter, A., Dominik, C., and Waters, L. B. F. M.: 2003, A&A 401, 577. Brucato, J. R., Colangeli, L., Mennella, V., et al.: 1999, P&SS 47, 773. Cernicharo, J. and Crovisier, J.: 2005, Space Sci. Rev., this volume. Colangeli, L., Bussoletti, E., Pestellini, C. C., et al.: 1998, Icarus 134, 35. Colangeli, L., Epifani, E., Brucato, J. R., et al.: 1999, A&A 343, L87. Crovisier, J., Brooke, T. Y., Hanner, M. S., et al.: 1996, A&A 315, L385. Crovisier, J., Leech, K., Bockelee-Morvan, D., et al.: 1997a, Science 275, 1904. 154 T. G. MULLER¨ ET AL.

Crovisier, J., Leech, K., Bockelée-Morvan, D., Brooke, T. Y., et al.: 1997b, in A. M. Heras, K. Leech, N. R. Trams and M. Perry (eds.), First ISO Workshop on Analytical Spectroscopy, ESA SP-419, p. 137. Crovisier, J., Encrenaz, T., Lellouch, E., et al.: 1999a, in P. Cox and M. F. Kessler (eds.), The Universe As Seen by ISO, ESA SP-427, p. 161. Crovisier, J., Leech, K., Bockelée-Morvan, D., et al.: 1999b, in P. Cox and M. F. Kessler (eds.), The Universe As Seen by ISO, ESA SP-427, p. 137. Crovisier, J., Brooke, T. Y., Leech, K., et al.: 2000, in M. L. Sitko, A. L. Sprague and D. K. Lynch (eds.), Thermal Emission Spectroscopy and Analysis of Dust, Disks, and Regoliths, ASP Conf. Ser. 196, 109. Crovisier, J.: 2000a, in Y. C. Mihn and E. F. van Dishoeck (eds.), Astrochemistry: From Molecular Clouds to Planetary Systems (IAU Symp. 197), ASP, 461. Crovisier, J.: 2000b, in A. Salama, M. F. Kessler, K. Leech and B. Schulz (eds.), ISO Beyond the Peaks, ESA SP-456, 23. Davies, J. K., Sykes, M. V., Reach, W. T., et al.: 1997, Icarus 127, 251. Dotto, E., Barucci, M. A., Crovisier, J., et al.: 1999, in P. Cox and M. F. Kessler (eds.), The Universe As Seen by ISO, ESA SP-427, p. 165. Dotto, E., Müller, T. G., Barucci, M. A., et al.: 2000, A&A 358, 1133. Dotto, E., Barucci, M. A., Müller, T. G., et al.: 2002, A&A 393, 1065. Dotto, E., Barucci, M. A., Brucato, J. R., et al.: 2004, 308 Polyxo: ISO-SWS Spectrum up to 26 micron, A&A, accepted. Encrenaz, T.: 2000, in F. Casoli, J. Lequeux, and F. David (eds.), Infrared Space Astronomy, Today and Tomorrow (Ecole´ de Physique des Houches), EDP Sciences, Paris, p. 89. Epifani, E., Colangeli, L., Fulle, M., et al.: 2001, Icarus 149, 339. Fernández,Y. R., Lisse, C. M., Käufl,U. H., et al.: 2000, Icarus 147, 145. Fouchet, T.: 2003, in C. Gry, S. B. Peschke, J. Matagne et al. (eds), Exploiting the ISO Data Archive – Infrared Astronomy in the Internet Age, ESA SP-511, p. 41. Fouchet, T.: 2005, Space Sci. Rev., this volume. Groussin, O., Lamy, P., and Jorda, L.: 2004a, A&A, 413, 1163. Groussin, O., Lamy, P., Jorda, L., and Toth, I.: 2004b, A&A, 419, 375. Grün,E., Hanner, M. S., Peschke, S. B., et al.: 2001, A&A 377, 1098. Harker, D. E., Wooden, D. H., Woodward, C. E., and Lisse, C. M.: 2002, ApJ. 580, 579. Harris, A. W. and Davies, J. K.: 1999, Icarus 142, 464. Hiroi, T. and Hasegawa, S.: 2003, Antarct. Meteorite Res. 16, 176. Holmes, E. K. and Dermott, S. F.: 2000, in ISO Beyond Point Sources: Studies of Extended Infrared Emission, ESA-SP 455, p. 67. Jorda, L., Lamy, P., Groussin, O., et al.: 2000, in R.J. Laureijs, K. Leech and M. F. Kessler (eds.), ISO Beyond Point Sources, ESA-SP 455,p.61. Lagerros, J. S. V.: 1996, A&A 310, 1011. Lagerros, J. S. V.: 1997, A&A 325, 1226. Lagerros, J. S. V.: 1998, A&A 332, 1123. Lagerros, J. S. V., Müller, T. G., Klaas, U., and Erikson, A.: 1999, Icarus 142, 454. Lamy, P. L., Toth, I., Jorda, L., et al.: 2002, Icarus 156, 442. Lamy, P. L., Toth, I., Fernández,Y. R., and Weaver, H. A.: 2004, in M. Festoum, H. U. Keller, and H. A. Weaver (eds), Comets II, University Arizona Press, Tucson, in press. Lebofsky, L. A., Sykes, M. V., Tedesco, E. F., et al.: 1986, Icarus 68, 239. Leech, K., Crovisier, J., Bockelée-Morvan, D., et al.: 1999, Earth Moon and Planets 78, 81. Leinert, Ch., Abrahám,P.,´ Acosta-Pulido, J., Lemke, D., and Siebenmorgen, R.: 2002, A&A 393, 1073. Lellouch, E., Crovisier, J., Lim, T., et al.: 1998, A&A 339, L9. COMETS, ASTEROIDS AND ZODIACAL LIGHT AS SEEN BY ISO 155

Malfait, K., Waelkens, C., Waters, L. B. F. M., et al.: 1998, A&A 332, L25. Molster, F. J., Yamamura, I., Waters, L. B. F. M., et al.: 1999, Nature 401, 563. Molster, F. J.: 2005, Space Sci. Rev., this volume. Müller, T. G. and Lagerros, J. S. V.: 1998, A&A 338, 340. Müller, T. G., Lagerros, J. S. V., Burgdorf, M., et al.: 1999a, in P. Cox and M. F. Kessler (eds.), The Universe As Seen by ISO, ESA SP-427, p. 141. Müller, T. G., Lagerros, J. S. V., and Klaas, U.: 1999b, in R. J. Laureijs and R. Siebenmorgen (eds.), Workshop on ISO Polarisation Observations, ESA SP-435, p. 31. Müller, T. G., Dotto, E., and Barucci, A.: 2000, in A. Salama, M. F. Kessler, K. Leech and B. Schulz (eds.), ISO Beyond the Peaks: The 2nd ISO Workshop on Analytical Spectroscopy, ESA-SP 456, p. 33. Müller, T. G.: 2001, Planet. Space Sci. 49, 787. Müller, T. G., Hotzel, S., and Stickel, M.: 2002, A&A 369, 655. Müller, T. G. and Lagerros, J. S. V.: 2002, A&A 381, 324. Müller, T. G.: 2002, M&PS 37, 1919. Müller, T. G.: 2003a, in C. Gry, S. B. Peschke, J. Matagne et al. (eds), Exploiting the ISO Data Archive – Infrared Astronomy in the Internet Age, ESA SP-511, p. 47. Müller, T. G.: 2003b, ISO: Physical Properties and Morphology of Small Solar System Bodies, IAU JD 19, p. 29. Müller, T. G. and Lagerros, J. S. V.: 2003, in L. Metcalfe, A. Salama, S. B. Peschke and M. F. Kessler (eds.), The Calibration Legacy of the ISO Mission, ESA SP-481, p. 157. Müller, T. G. and Blommaert, J. A. D. L.: 2004, A&A 418, 347. Reach, W. T., Abergel, A., Boulanger, F., et al.: 1996, A&A 315, L381. Reach, W. T.: 1997, in H. Okuda, T. Matsumoto, and T. Rollig (eds.), Diffuse Infrared Radiation and the IRTS, ASP Conference Series 124, p. 33. Reach, W. T., Sykes, M. V., Lien, D., and Davies, J. K.: 2000, Icarus 148, 80. Reach, W. T., Morris, P., Boulanger, F., and Okumura, K.: 2003, Icarus 164, 384. Tedesco, E. F. and Desert, F. X.: 2002, AJ 123, 2070. Thomas, N., Eggers, S., Ip, W.-H., et al.: 2000, ApJ 534, 446. Waelkens, C., Malfait, K., and Waters, L. B. F. M.: 1999, Earth Moon and Planets 79, 265. Wooden, D. H., Harker, D. E., Woodward, C. E., et al.: 1999, ApJ 517, 1034. STARS and CIRCUMSTELLAR MATTER ISO OBSERVATIONS OF PRE-STELLAR CORES AND YOUNG STELLAR OBJECTS

BRUNELLA NISINI1,∗, ANLAUG AMANDA KAAS2, EWINE F. VAN DISHOECK3 and DEREK WARD-THOMPSON4 1INAF Osservatorio Astronomico di Roma, Via di Frascati 33, 00040 Monteporzio Catone, Italy 2Nordic Optical Telescope, Apdo 474, 38700 Santa Cruz de La Palma, Spain 3Leiden Observatory, P.O. Box 9513, 2300 RA Leiden, The Netherlands 4Department of Physics & Astronomy, Cardiff University, P.O. Box 913, Cardiff, CF24 3YB, U.K. (∗Author for correspondence: E-mail: [email protected])

(Received 9 July 2004; Accepted in ﬁnal form 13 December 2004)

Abstract. We summarize the observations of the Infrared Space Observatory (ISO) concerning the earliest stages of the stellar formation. The observations of samples of sources in different evolutionary stages are reviewed, addressing in particular how the physical and chemical properties of the protostellar environments change from the pre-stellar cores to the protostars at the end of their accretion phase. In addition, the mid-IR surveys in nearby star-forming regions are discussed, showing their implications for the understanding of the stellar initial mass function.

Keywords: star formation, young stellar objects, protostars, ISO, infrared observations

1. Introduction

The Infrared Space Observatory (ISO) has represented a fundamental step in enlarg- ing the observational data set needed to understand the still poorly known earliest stages of formation of a new star. Indeed, the spectral range covered by ISO (from 2 to 200 μm) is the most suited, both to identify new cold and heavily extincted protostars, and to probe the different processes characterizing the interaction between young stars and their parental cloud. In particular, observations at mid- and far-infrared wavelengths are able to probe the gas and dust cooling of young stellar objects (YSOs) close environments, which is in turn a function of the different heating mechanisms and of their time evolution. The formation of an isolated low-mass star begins with the collapse of a cold dense core (pre-stellar core) where no source of internal heating is present. As soon as a centrally condensed protostar develops, the main heating mechanisms of circumstellar gas and dust become shocks, either due to matter accretion onto the

Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, The Netherlands, and the United Kingdom), and with the participation of ISAS and NASA. Space Science Reviews (2005) 119: 159–179 DOI: 10.1007/s11214-005-8068-8 C Springer 2005 160 B. NISINI ET AL. disk and the protostellar surface, or developing along the energetic outflows (Class 0 sources, ages ∼104–105 year). Finally, when the protostellar phase ends and the star approaches its pre-main sequence tracks, the accretion/ejection mechanisms become less and less important up to the point where the stellar radiation field starts to represent the main source of heating (Class I/II, ∼105–106 year). The heating processes in high-mass protostars may follow a different time evolution. Indeed, due to their much shorter evolutionary time scales, high-mass sources develop large radiation-dominated regions while they are still accreting their mass, but the mechanisms at work are the same as in low-mass stars. ISO has gathered both photometric and spectroscopic observations of samples of sources with different mass over all these evolutionary stages. In addition, it has performed extended surveys in nearby molecular clouds, allowing the detection and classification of the young stellar and sub-stellar population of embedded clusters. This, in turn, has allowed us to extend the initial mass function (IMF) well into the brown dwarfs regime. In this article we will review the main results obtained by such observations. In particular, we will concentrate on the main mechanisms giving rise to the gas- and dust cooling in pre-stellar cores and young embedded sources (Class 0/I), with a particular emphasis on the identification of specific evolutionary trends to be used in the exploiting of present and future IR space missions. The observation of more evolved YSOs (pre-main sequence Class II sources) and their disks will be reviewed by Lorenzetti and Jourdain de Muizon in this volume.

2. Pre-Stellar Cores

The formation of YSOs is known to occur in dense cores within molecular clouds (e.g. Williams et al., 2000). Study of such regions has been hampered in the past by their very large optical depths at near-infrared and optical wavelengths. It is only since the opening up of the far-infrared to sub-millimetre regime that astronomers have been able to study molecular clouds in detail. ISO has played a key role in the progress that has been made. Theories of star formation have, until recently lacked a detailed observational determination of the initial conditions of the protostellar collapse phase (e.g. see: André et al., 2000 for a review). The pre-protostellar (or pre-stellar for short) core phase (Ward-Thompson et al., 1994) is the stage of star formation that precedes the formation of a protostar and hence should represent observationally the initial conditions of protostellar collapse. Some recent observations have indicated that the Initial Mass Function (IMF) of stars may be determined at the pre-stellar core stage (Motte et al., 1998, 2001). Furthermore, theory predicts that the core geometry prior to this stage is critical in determining the manner of protostellar collapse (e.g. Whitworth and Summers, 1985; Foster and Chevalier, 1993; Whitworth et al., 1996). Therefore, it is of vital ISO OBSERVATIONS OF YSOs 161 importance to star formation theories to determine observationally the physical parameters of pre-stellar cores. Molecular line surveys of dense cores by Myers and co-workers identified a significant number of isolated cores (Myers and Benson, 1983; Myers et al., 1988; Benson and Myers, 1989). Comparisons of these surveys with the IRAS point source catalogue identified those that already harboured embedded protostars, allowing those without to be classified as starless cores (Beichman et al., 1986). Subsequent millimetre and sub-millimetre surveys allowed those starless cores that were most centrally condensed – the pre-stellar cores – to be identified (Ward-Thompson et al., 1994, 1999; André et al., 1996). The millimetre and sub-millimetre work showed that the cores all follow a form of density profile that is relatively flat in the centre and steeper towards the edge (Ward-Thompson et al., 1994). This is similar to the profiles predicted by theory to produce a decreasing accretion rate with time (Whitworth and Summers, 1985; Foster and Chevalier, 1993). The profiles are not consistent with the scale-free power-law profile predicted by the singular isothermal sphere (Shu, 1977). ISO was used to image a number of pre-stellar cores, using the long-wavelength photopolarimeter ISOPHOT (Ward-Thompson et al., 2002). Figure 1 shows images of the pre-stellar core L1544 at each of three far-infrared wavelengths – 90, 170 and 200 μm. It can be seen that the core is very much less clearly defined at 90 μm than at the other two wavelengths. This is due to the temperature of the core being very low. Typically, pre-stellar cores have temperatures around 10 K (Ward-Thompson et al., 2002). The colour temperature variation can be calculated across each core by first subtracting the background level and then taking the ratio of the background-subtracted images at 170 and 200 μm. This can be converted to a colour temperature using the assumption of optically thin, modified black-body emission, namely: 3+β ν / ν [h 2 kT] − 1 e 1 (Fν /Fν ) = , 1 2 3+β ν / ν [h 1 kT] − 2 e 1 where the frequencies ν1 and ν2 correspond to wavelengths of 200 and 170 μm, respectively, and Fν1 and Fν2 are the flux densities at each of these frequencies. T is the dust temperature, β the dust emissivity index (set to 2), and h and k are the Planck and Boltzmann constants, respectively. Using these assumptions, the colour temperature map in Figure 1 was made. Examination of this map shows that there is a temperature gradient approximately centred on the core, with the edge warmer than the centre. This is consistent with external heating by the local inter-stellar radiation field (ISRF). ISOCAM was also used to observe pre-stellar cores. However, the cores are so dense that they do not emit at near- and mid-infrared wavelengths. Instead, the cores are seen in absorption by ISOCAM (Bacmann et al., 2000) against the general background emission of the mid-IR ISRF at a wavelength of 7 μm. 162 B. NISINI ET AL.

Figure 1. Grey-scale ﬂux density images, with isophotal contour maps superposed, of the pre-stellar core L1544.

Figure 2 shows the pre-stellar core L1696A (also known as Oph D) as seen in emission by SCUBA at 850 μm and in absorption at 7 μm by ISOCAM. The two maps show a very similar morphology. This indicates that the same dust that is absorbing in the near-IR is re-emitting in the sub-millimetre. Calculations of the amount of the absorbed energy and emitted energy show they are roughly equal, implying that pre-stellar cores are in approximate energy equilibrium with their surroundings and have no embedded energy sources (Ward-Thompson et al., 2002). The fact that the same dust is being traced at 7 μm allows one to probe the structure of a pre-stellar core at the higher angular resolution of ISOCAM – 6 arcsec. This type of study showed that in fact pre-stellar cores are very sharp edged. They can be modelled with proﬁles that reach density gradients steeper than ρ ∝ r −3 (Bacmann et al., 2000). These results are in direct contrast with singular isothermal sphere ISO OBSERVATIONS OF YSOs 163

Figure 2. The pre-stellar core L1696A (also known as Oph D) seen in sub-millimetre emission (left) and near-IR absorption (right). The similarity of the two maps shows that it is the same dust that is absorbing in the near-IR and re-emitting in the sub-millimetre (images from Kirk et al., 2005; Bacmann et al., 2000 respectively). models. However, some magnetically regulated models may be able to reproduce the observations.

3. Low-Mass YSOs

Systematic studies of solar mass embedded YSOs have become possible only after the IRAS survey, which has provided the discovery and census of large samples of protostellar candidates. The combination of ground-based and IRAS observations is still one of the main tool to define the relative evolutionary state of different YSOs based on the steepness of their spectral energy distribution (SED) from near to far-IR (e.g. Lada, 1999; André et al., 2000). ISO has provided two fundamental improvements with respect to the information given by IRAS. Firstly, it has given the possibility to extend the observations of protostellar SEDs up to 200 μm, which is where the youngest protostars (i.e. the Class 0 sources) have their emission peaks. Secondly, it has given for the first time, through its spectroscopic capabilities, the tools to probe the physical and chemical conditions of the warm circumstellar gas in large samples of sources, and to study their temporal evolution. These two aspects will be separately discussed in the following sections. Recent reviews on ISO observations of low-mass YSOs, which address these and other aspects of the subject can be found in van Dishoeck (2004), Nisini (2003) and Saraceno et al. (1999). 164 B. NISINI ET AL.

3.1. SPECTRAL ENERGY DISTRIBUTION (SED)

The shape of the YSOs SEDs in the mid- and far-infrared is a sensitive function of the temperature and density structure of their circumstellar envelopes and of the geometry of the dust distribution. Such a shape has been reconstructed with unprecedented details through both the photometric data by CAM and PHOT, and the spectro-photometric observations provided by SWS and LWS. Most of the analysis has been focussed on the young Class 0 protostars, which have their bulk of emission in the far-IR and sub-millimetre range. Adopting a simple model of grey-body emission from a spherical envelope, Froebrich et al. (2003) derived the Lbol and Tbol of seven young protostars using ISOPHOT data in conjunction with other sub-millimetre observations. The typical temperatures found range between 20 and 35 K. The location in a Lbol–Tbol diagram, compared with evolutionary tracks for young protostars, has been, in turn, used to infer their ages, which span from 15,000 to 30,000 years, thus confirming the youth of the objects. The SEDs of many IR/SMM sources in the Serpens cloud have been extracted from an LWS raster map of the region by Larsson et al. (2000). They show that most of the observed SEDs can be well represented as modified single-temperature black- bodies, as previously assumed, but that the derived temperatures are generally higher than what has been found by means of IRAS observations, which also implies higher Lbol values. For the individual Class 0 source SMM1/FIRS1 the model for an envelope with two escavated regions of lower densities produces a better fit also to the mid-IR ISOCAM data of the source (Larsson et al., 2002; Figure 3). Other young protostars studied with a similar analysis includes IRAS16293- 2422 in ρ Oph (Correia et al., 2004), and several sources of the Cederblad 110 nebula in the Cham I molecular cloud (Lehtinen et al., 2001). The SEDs of the more evolved Class I sources do appear much broader than those of the Class 0 sources due to their more tenuous envelopes, which allow to see their inner and warmer regions mainly emitting in the mid-IR. A detailed account for the SEDs from the near to the sub-millimetre have been given for two of these sources, namely L1551-IRS5 (White et al., 2000) and Elias 29 (Boogert et al., 2002). In both cases, and at variance with the Class 0 SEDs, good fits to the observed data require the inclusion of different emission components taking into account the presence of both a flaring disk and a more extended envelope. In Elias 29, in particular, a large and massive (up to 0.012M) disk has been suggested to explain the observed SED flattness between 10 and 200 μm (Figure 4).

3.2. GAS COOLING AND SPECTRAL EVOLUTION OF PROTOSTARS

ISO has gathered SWS, LWS and CAM-CVF spectra of samples of low-mass YSOs with different luminosities and evolutionary stages, providing a fundamental and ISO OBSERVATIONS OF YSOs 165

Frequencyν (Hz) 1013 1012 1011 1016 1016 N Q Small beams Medium beams Large beams 15 15 10 ISO CVF, LWS 10 Central source Equiv sphere Total fluxes 14 14 10 Beam-matched 10 ν 13 13

ν 10 10

1012 1012

1012 11 11 Energy distribution10 F (Hz Jy) 10

1010 1010

1011 10 109 109 10 100 1000 Wavelengthλμ ( m)

Figure 3. SED of the Class 0 source SMM1/FIRS1 from 10 μm to 3 mm, with a superimposed model ﬁt assuming an envelope emission where two bipolar cavities are escavated by the source outﬂow (solid line). The SED for the equivalent spherical envelope is also shown (from Larsson et al., 2002).

Figure 4. SED of the Class I source Elias 29 as observed by ISO SWS and LWS, combined with ground-based 1.3 mm flux (Boogert et al., 2002). The SED has been fitted by a model considering four different components, namely a black-body emission from the inner envelope, a flaring disk, an outer envelope and a foreground cloud. 166 B. NISINI ET AL. unique database to probe the gas physical conditions and their time evolution in the protostellar envelopes. The SWS and CAM-CVF spectra of low luminosity sources usually present only broad features due to abundant ices seen in absorption against the young source. These originate in the outer envelopes and, consequently, they are more sensitive to the amount of foreground dust and its chemical composition than to the conditions close to the source (Boogert et al., 2002; Alexander et al., 2003). The LWS spectra, on the other hand, are rich of lines in emission whose excitation mechanisms can be directly reconducted to intrinsic properties of the source itself, allowing us to discriminate among the different excitation mechanisms that dominate the gas heating of circumstellar envelopes during the protostellar evolution. The LWS spectral surveys of samples of low luminosity Class 0 and Class I sources have been discussed by Giannini et al. (2001) and Nisini et al. (2002), while the spectra of individual objects have been analysed by different authors (see van Dishoeck, 2004, for a complete list of references). The far-IR spectra of embedded YSOs indicate low excitation conditions in their envelopes, since only emission from the atomic lines [O I] 63 and 145 μm and the single ionised line [C II] 158 μm, as well as from molecular transitions of abundant species (CO, H2O, OH), are detected. Such lines probe gas temperatures and densities in the range ∼100–2000 K and 104–107 cm−3, respectively. The main observational evidence that appears from the analysis of the spectra of different sources is that, while the spectra from the Class 0 are dominated by both atomic and molecular emission, very few molecular lines, and only in form of CO and sporadically OH, are detected in Class I sources. Quantitatively, the total gas cooling due to molecular emission with respect to the bolometric luminosity decreases by about 1 order of magnitude going from Class 0 to Class I sources (see Figure 5). In Class 0 objects, water has been detected in 13 out of 17 observed sources, and its abundance is much higher (10−4–10−5) than in quiescent clouds with a trend to increase with the kinetic temperature. Several authors suggest that the main mechanism for gas excitation are the strong shocks expected to occur as the energetic protostellar jets impact towards the dense envelopes (see e.g. Giannini et al., 2001; Nisini et al., 2002; Larsson et al., 2002; Molinari et al., 2000; Froebrich et al., 2002). Different evidences are in favour of this interpretation: the physical conditions and relative abundances of the observed species, derived from the analysis of observed lines, are very similar to those inferred in shocked regions far from the source and can be explained in the framework of low velocity and dense shocks (see Section 5). In addition, the strongest molecular emission is observed towards sources where dense and warm molecular jets are present close to the object. Finally, there is a direct correlation between the total luminosity radiated by the traced gas component and the kinetic luminosity associated with the protostellar outflow. In the framework of shock excitation, the observed difference in the far-IR spectra of Class 0 and Class I objects can be explained by considering that the protostellar envelope with time ISO OBSERVATIONS OF YSOs 167

Figure 5. Continuum subtracted ISO-LWS spectrum towards the Class 0 source L1448-mm (upper) and the Class I source IRS 44 (bottom left). While the Class 0 source spectrum is very rich of molecular emission, especially H2O, the Class I spectrum has only few weak lines from CO and OH. (Bottom right) Histograms of the cooling due to molecular emission with respect to the bolometric luminosity for the Class 0 and Class I systems observed by ISO-LWS (Nisini et al., 1999, 2002). becomes more diffuse and escavated. Thus, the protostellar jet encounters a smaller ambient density, which gives rise to preferentially dissociative shocks, causing a significant decrease of the contribution of the molecular luminosity to the total gas cooling (Nisini et al., 2002). A different scenario is proposed by Ceccarelli et al. (2000) and Maret et al. (2002) which suggest that most of the water and [O I] emission observed in Class 0 sources is originated from the innermost regions of protostellar envelopes, where the relatively high dust temperatures can favour the evaporation of gaseous H2O from the grain mantles. Such an hypothesis is supported by a detailed model of the envelope thermal structure and radiative transfer which provide a good fit of the observed water line emission in NGC1333-IRAS4 and IRAS16293-2422, with −5 mass accretion rates of few times 10 M year. In this framework, the decline of water emission in Class I sources is explained by the water dissociation as soon as the source develops a significant FUV field. This hypothesis, however, requires 168 B. NISINI ET AL. a different mechanism to explain the strong CO emission always associated with the H2O lines, which cannot be reproduced in the envelope model. Ceccarelli et al. (2002) suggest that such CO emission in the Class I source Elias 29 could be originated in the circumstellar disk, an hypothesis which seems however difficult to explain the high CO temperatures (T = 300–1500 K) derived in most of the Class 0 sources.

4. High-Mass YSOs/UC H II Regions

In contrast with the case of low-mass YSOs, there is no well-established scenario for the evolution of high-mass YSOs. Many different observational phenomena are known to be associated with high-mass star formation – including radio emission from UltraCompact (UC) H II regions, H2O and CH3OH maser emission, millimetre lines of complex organic ‘hot core’ molecules, and bright mid-infrared continuum – but these have not yet been put in a clear evolutionary sequence. Moreover, most high-mass stars form in groups with the various YSOs often at slightly different evolutionary stages. For example, hot cores are often found within a few arcsec of UC H II regions (see Kurtz et al., 2000; Churchwell, 2002, for reviews). Within the ISO beams, these objects are usually blurred together, although often a single object dominates the mid- and far-infrared fluxes. A prototype example of a high-mass star-forming region is provided by Orion BN-KL. Within the large ISO beams, the Becklin-Neugebauer object dominates the flux at shorter wavelengths (<10 μm), whereas IRc2 is stronger at longer wavelengths. Due to the interaction of the massive young stars with their surroundings, a wealth of spectral features is revealed beautifully in SWS spectra by van Dishoeck et al. (1998), González-Alfonso et al. (1998, 2002), Wright et al. (2000), and Rosenthal et al. (2000), and in LWS data by Harwit et al. (1998) and Cernicharo et al. (1999). All phenomena associated with massive YSOs are reflected in the spectra, including foreground ionised gas, PDRs, shocks, and the warm gas associated with the deeply embedded YSOs themselves. In spite of these limitations, ISO observations have clearly contributed to our understanding of the evolution of high-mass YSOs, both in the deeply embedded early stages and in the more evolved H II region phase. The bulk of the new results are due to low- to medium-resolution spectroscopy with the SWS and LWS, which also provide a complete characterization of the SED.

4.1. DEEPLY EMBEDDED HIGH-MASS YSOS

The youngest high-mass pre-stellar cores and deeply embedded YSOs emit primarily at far-infrared wavelengths. The ISOPHOT Serendipity Survey has revealed a number of cold compact objects at 170 μm coinciding with sources detected at long wavelengths with other facilities (e.g. Krause et al., 2003). These objects have ISO OBSERVATIONS OF YSOs 169 low average dust temperatures (12–16 K) and masses up to a few hundred M, and are thus excellent candidates for massive protostars. About a dozen embedded high-mass YSOs have been observed with the SWS and LWS (e.g. Gerakines et al., 1999; Gibb et al., 2004). The combined SWS–LWS spectra show that the continuum flux peaks at 60–90 μm, characteristic of black- body emission of ∼30 K. Some sources are brighter at mid-infrared wavelengths around 12 μm, indicating the presence of warmer dust along the line of sight with temperatures up to a few hundred Kelvin. The SWS spectra are characterized by deep amorphous silicate and ice bands, which can absorb more than 50% of the continuum flux in the 2–20 μm region. Very few other features are seen in these deeply embedded stages: no PAH emission is apparent, and there are virtually no gas-phase molecular lines. Only a few atomic lines are detected in the LWS spectra, mostly the [C II] and [O I] fine-structure lines associated with nebulosity or PDRs. Thus, these objects do not resemble Orion BN-KL, even when scaled for the larger distance. In spite of their similar SEDs, the embedded objects do show evidence for progressive heating of the envelope. This is most clearly reflected in the shape of 13 a number of ice absorption bands, most notably the 15 μmCO2, 4.3 μm CO2 and 6.8 μm unidentified features (Gerakines et al., 1999; Boogert et al., 2000; Keane et al., 2001). In addition, the gas/solid abundance ratio of species like CO and H2O (e.g. Boonman and van Dishoeck, 2003), and the excitation temperatures of gaseous C2H2 and HCN (Lahuis and van Dishoeck, 2000) show a similar trend which correlates well with the far-infrared colour as defined by the 45/100 μm flux ratio. Thus, the objects can be put in a temperature sequence. Since these diagnostics involve absorption in a pencil beam along the line of sight as well as emission from the entire envelope, the heating must be ‘global’. The most likely scenario, put forward by van der Tak et al. (2000), is that the warmer sources have a lower ratio of envelope mass Menv compared with Lbol. If the lower Menv are due to gradual dispersion of the envelope rather than different initial conditions, the heating sequence would correspond to a time sequence. Thus, the mid-infrared spectral features provide a unique and powerful diagnostic of the first few ×104 years of protostellar evolution.

4.2. UC H II REGIONS

When the envelope is dispersed and the H II region developed, the spectral characteristics of massive YSOs change dramatically. An excellent example is provided by the comparison of Cep A with S 106 (see Figure 6; van den Ancker et al., 4 2000). Both sources have luminosities of 2–4 × 10 L, but the spectrum of Cep A is characteristic of the deeply embedded sources discussed earlier, whereas the S 106 spectrum is dominated by PAHs and by atomic lines from a wide range of ionisation stages characteristic of a young O8 star. The atomic and molecular H2 170 B. NISINI ET AL.

Figure 6. Combined SWS–LWS spectra of two massive YSOs at different evolutionary stages. Cep A(2.4 × 104 L) is in the deeply embedded phase, whereas S 106 (4.2 × 104 L) is in a more evolved stage where it has developed a compact H II region and a PDR (van den Ancker et al., 2000). and CO emission lines indicate that shock excitation dominates for Cep A, but PDR excitation for S 106. Peeters et al. (2002) present combined SWS–LWS spectra for a sample of 45 compact H II regions, and Giveon et al. (2002) analyse SWS spectra of 112 galactic H II regions. The SEDs of the compact H II regions peak at 40–60 μm, corresponding to Td ≈ 60–70 K, clearly warmer than those of sources in the deeply embedded phase. The spectra are dominated by PAHs and ionised atomic lines. Since the sources cover a wide range of galactocentric radii ranging from 0 to 18 kpc, the line ratios are used to study variations in the excitation, metallicity, and hardness of radiation ﬁeld on a galactic scale. No attempts have yet been made to use these features to study evolutionary effects within the H II region sample. ISO OBSERVATIONS OF YSOs 171

At an even later stage, the young stars have completely broken free from their parental clouds and the young O and B stars emit copious ultraviolet photons, which can create PDRs at the edges of neighbouring molecular clouds. Examples are given by M 17, Orion and Trumpler 14, with the latter source harbouring a cluster of at least 13 O stars (Brooks et al., 2003). They are readily recognized by strong atomic ﬁne-structure lines and H2 emission, and have been reviewed in detail by Hollenbach and Tielens (1997) prior to ISO. New insights due to ISO are summarized by van Dishoeck (2004).

5. Outﬂows and HH Objects

Shocks originated in outflows cools mainly radiatively through line emission in a wide range of wavelengths, from UV to the radio (Reipurth and Bally, 2001). In particular, shocks occurring in dense environments cool down preferentially through emission lines in the mid- and far-IR, mainly through the [O I]63μm transition and the pure rotational transitions of H2, CO and H2O (see e.g. Hollenbach, 1997). Due to the limited spectral and spatial resolution of the ISO observations, they do not allow in most cases to resolve the emission regions along the jets; nevertheless, they give unique information on global properties of the warm shocked gas, such as the total radiative cooling in the post-shocked layers and the relative abundances of species which in turn can be used to give tight constraints on the physics of the shocked region. Most of the well known molecular outflows from very young stars have been observed by ISO-LWS and SWS (see van Dishoeck, 2004 for a complete references list of the observed shocked regions). The physical conditions derived from the molecular emission are consistent with moderately warm and dense gas (n = 104– 106 cm−3, T = 300–1500 K) emitted by compact regions. Indeed, strong molecular emission is evidenced mainly in young outflows showing the presence of localized components of very high-velocity gas (v>50 km s−1). The far-infrared gas cooling can be a significant fraction of the total shock cooling in outflows from young embedded protostars, often representing more than 90% of the total radiated power −2 (e.g. Nisini et al., 2000). Such a shock-radiated power is ∼10 Lbol in outflows from Class 0 sources, while it is about 1 order of magnitude smaller for Class I outflows, indicating a clear decline of the outflow power with the evolution. The presence, in most of the observed spectra, of both strong [O I]63μm emission and strong molecular emission in form of CO and H2O supports the view where the interaction between the jet and the ambient medium occurs through a bow shock, where a high-velocity shock, mainly emitting in atomic transitions, is produced at the bow apex, while the molecular emission comes from lower velocity shocks in the bow wings. Such a bow structure is also able to produce the stratification of temperatures needed to take into account the H2 emission spectra, which show both low excitation pure rotational lines, as detected by SWS and CAM-CVF, 172 B. NISINI ET AL.

Figure 7. Left: LWS spectrum of the Cep E outflow (Giannini et al., 2001). Right: C-type bow shock model fit for Cep E, blue lobe, through ground-based roto-vibrational and ISO pure rotational H2 lines (upper panel) and LWS rotational CO lines (bottom panel). From Smith et al. (2003). and the roto-vibrational lines at higher excitation observed by ground (e.g. Wright et al., 1996; Moro-Mart´ın et al., 2001; Wilgenbus et al., 2002; Rosenthal et al., 2000). Detailed models for bow shocks have been successfully tested and applied to the ISO observations of several outflows (e.g. Froebrich et al., 2002; Smith et al., 2003; see Figure 7). Giving their low excitation temperature, the H2 0–0 lines can be also used to show deviations of the ortho/para ratio from the equilibrium value. In the Herbig- Haro object HH54, an ortho/para ratio ∼1 has been estimated which implies that the gas has not yet reached the equilibrium and it is therefore very young (Neufeld et al., 1998; Wilgenbus et al., 2000). In addition, the unique possibility offered by ISO to observe for the first time many transitions of thermally excited H2O, has allowed us to test predictions about the water formation in shocks and its importance in the gas cooling. Indeed, in magneto-hydrodynamic C-type shocks, water is expected to be one of the major coolant of the post-shocked gas, due to its rapid formation through chemical reactions which transform all the available neutral oxygen into H2O as soon as the gas temperature increases above ∼400 K (e.g. Kaufman and Neufeld, 1996). Pure rotational lines of H2O have been indeed observed in several outflows and high abundances (10−5–10−4) are often measured (Liseau et al., 1996; Harwit et al., 1998; Nisini et al., 2000; Giannini et al., 2001), which are much higher than the abundances measured in quiescent molecular clouds. The ISO spectra of many optically visible HH objects, show the presence of line emission also from the fundamental levels of ions, such as [Fe II], [Si II], [O III], ISO OBSERVATIONS OF YSOs 173

[N II] and [Ne II] which trace the presence of high excitation dissociative shocks (e.g. Molinari et al., 2001; White et al., 2000; Leﬂoch et al., 2003). Electronic densities derived through ratios among these lines are higher by at least 1 order of magnitude with respect to those inferred in the optical. Finally, from the luminosity of the [O I]63μm line observed in a sample of HH object, Liseau et al. (1997) have determined the mass loss rates of the HH exciting sources, which result to be in good agreement with those estimated in the associated CO outﬂow, implying that the two phenomena are strictly connected.

6. YSOs Census in Nearby Clouds

The origin of the stellar Initial Mass Function (IMF) is one of the most fundamental problems in astronomy. By studying statistically significant samples of YSOs in nearby clusters we hope to understand the relation between the final stellar IMF and the various phases of the star formation process. The observational challenge in obtaining a census of YSOs is (1) to detect the objects, as they are frequently embedded in dense clouds (AV ∼10–100 mag) and (2) to distinguish cluster members from field stars. The main objective of the ISOCAM surveys LNORDH.SURVEY 1 and GOLOFSSO.D SURMC within the ISO central programme was to make a deep search of YSOs in selected parts of nearby star formation regions in order to sample objects down to ∼0.03L. In terms of masses this corresponds to ∼0.05M for typical ages of ∼106 years of T Tauri type stars (Class II objects). In Table I we list the star-forming regions mapped in those surveys. The two broad band filters LW2 (5–8.5 μm) and LW3 (12–18 μm) had been designed to avoid the silicate features around 10 and 20 μm, and to trace well the mid-IR SED of YSOs. Be- fore ISOCAM our knowledge of the stellar content in these regions was in most cases limited to observations in the optical, the near-IR, and from the IRAS survey. Ground-based mid-IR surveys were few and limited to the LMNQ bands (see e.g. Lada and Wilking, 1984; Greene et al., 1994).

6.1. IMPROVED SAMPLES OF IR-EXCESS SOURCES

The 10 times higher spatial resolution of ISOCAM (3–6 per pixel) resolved much of the source confusion for clustered regions experienced with IRAS. The results of the photometry in the two broad bands LW2 (6.7 μm) and LW3 (14.3 μm) showed that the nominal integration times of 15–60 s gave a sensitivity of ∼2 mJy. In left part of Figure 8 we have shown the [14.3/6.7] versus Fν(14.3) diagram for several hundred sources in some of the star formation regions explored. A general result for all regions is the clear bimodal distribution in terms of colour – a blue (open triangles) and a red (ﬁlled circles) group. The blue colour is as expected for normal 174 B. NISINI ET AL.

TABLE I Overview of the surveyed regions.

Region D (pc) A (square degree) References

NGC1333 300 0.19 L1551 140 0.12 G˚alfalk et al. (2004) L1527 140 0.26 L1641 480 0.02 NGC2023 480 0.43 Abergel et al. (2002) LBS23 480 0.14 NGC2071 480 0.28 NGC2068 480 0.15 Cha I 160 0.59 Nordh et al. (1996), Persi et al. (2000) Cha II 178 0.20 Nordh et al. (1996), Persi et al. (2003) Cha III 150 0.10 Nordh et al. (1996) ρ Oph L1688 140 0.56 Bontemps et al. (2001) ρ Oph L1689N 140 0.07 Bontemps et al. (2001) ρ Oph L1689S 140 0.07 Bontemps et al. (2001) Serpens Core 260 0.13 Kaas et al. (2004) Serpens NH3 260 0.09 Kaas et al. (1999) RCrA 130 0.45 Olofsson et al. (1999) Cep A 730 0.09 G˚alfalk (1999)

Figure 8. Left: The clear separation between IR-excess sources (filled circles) and sources without IR- excess (open triangles) is shown in the ISOCAM colour–magnitude diagram [14.3/6.7] vs. Fν (14.3). The colour index [14.3/6.7] is defined as the logarithm of the flux ratio, and as an alternative the SED index α = d log(λFλ)/d log λ calculated between 6.7 and 14.3 μm is shown on the right hand y-axis. The top x-axis indicates the 14.3 μm band magnitude. The effect of 30 mag of visual extinction is indicated with an arrow. Right: The same sources shown in the J–H/H–K diagram (from Kaas and Bontemps, 2000). ISO OBSERVATIONS OF YSOs 175 photospheres, i.e. the slope in the Rayleigh–Jeans tail. This group thus contains a mixture of field stars and cluster members without IR excesses. The red sources have colours typical of classical T Tauri stars (CTTS) with mid-IR excesses arising from optically thick circumstellar disk emission. Comparing the same sources in the near-IR J–H/H–K diagram demonstrates a merging of these two groups around the reddening line (Figure 8 right). On average, only 50% of the IR-excess population found with ISOCAM can be extracted by using the J–H/H–K diagram. Thus, the ISOCAM surveys have revealed IR-excess populations about twice as large as previously known, thereby, providing the to date most complete samples of the youngest YSOs. In RCrA only 1/3 of the ISOCAM selected IR-excess sources were previously known YSOs, and in L1551 only 12%. Note that while most ISOCAM sources are detected in deep near-IR surveys, it is the clear separation of IR-excess from cloud reddening which is the strength of the mid-IR. YSOs without IR-excesses can not be distinguished from field stars by the use of broad band photometry alone. Therefore, the ISOCAM surveys do not sample these sources and can not be used to directly evaluate the number fraction of disks among YSOs. Other search tools such as e.g. X-ray surveys, optical spectroscopy or proper motion studies must be applied to make a census of this part of the YSO population. Only approximate estimates of disk fractions based on previous studies and galactic models of the mid-IR point source sky can be given, and these estimates give disk fractions from 65 to 100%, the largest value being for the central Serpens Cloud Core (see later).

6.2. NATURE AND SPATIAL DISTRIBUTION OF IR-EXCESS SOURCES

As emphasized by Nordh et al. (1996) the [14.3/6.7] colour for the IR-excess sources is essentially constant – although with a spread – over a 14.3 μm flux interval of almost 4 orders of magnitude, showing that the source characteristics are similar over a very wide range of luminosities. Also, the different types of IR excess YSOs: Class I and Class II sources, and the so-called flat-spectrum sources in a transition phase between them, were found to have about the same mid-IR colour on average (see also Alexander et al., 2003). To classify the IR-excess YSOs we used the SED index α = d log(λFλ)/d log λ calculated between 2.2 and 14.3 μm. While the Class I/Class II number ratio is 5/42 for Cha I and 16/123 for ρ Oph, as expected based on current ideas about the time scales for each evolutionary class, this ratio is 19/18 for the central Serpens Cloud Core. Together with the very strong sub-clustering of the Class I and flat-spectrum sources (0.12 pc diameter) this suggests a very recent burst of star formation in Serpens. The Class II sources also show sub-clustering, but on scales of 0.25 pc and with an additional scattered distribution, similar to what was found in ρ Oph. In Cha I there is also sign of sub-clustering of Class II sources, but the clusters being slightly more extended, about 0.5–0.7 pc. 176 B. NISINI ET AL.

Figure 9. (a) The Class II LF for ρ Ophiuchi. Dashed vertical line shows the completeness limit. Bold curve shows the best model LF, which is obtained with a constant star formation rate over 0.2–0.5 Myr and an underlying two-step power-law IMF with two free parameters: the slope of the low-mass −1.7 part α1, and the mass at which the slope changes Mﬂat. dN/d log M ∝ M for the higher mass part, in accordance with recent IMFs. (b) The mass function for the Class II sources in ρ Ophiuchi. The bold dashed vertical line shows the brown dwarf domain to the left. The completeness limit is at 0.055M. The best two-step power-law IMF is characterized by the two slopes α1 and α2, and the mass where the slope changes, Mﬂat. See Bontemps et al. (2001) for more information.

6.3. THE INITIAL MASS FUNCTION (IMF)

The improved Class II samples provide luminosity functions (LFs) with completeness limits at about 0.07L for RCrA, 0.06L for Cha I, 0.03L for ρ Ophiuchi, and 0.08L for the Serpens Core. The LFs reach down to below 0.01L, however, and it turns out that 20% of the Class II sources in Cha I, ρ Ophiuchi and Serpens are putative young brown dwarfs with disks. In RCrA about 30% of the ISOCAM selected YSOs are substellar candidates. As shown in Figure 9(a) the 123 Class II sources in ρ Ophiuchi produce a rather ﬂat LF below ∼2L. Furthermore, as shown by Bontemps et al. (2001) each of the four sub-clusters A, B, EF, and L1689S have individual LFs of this same shape. The Class II LF can be modelled with an IMF ﬂat below 0.55M (see Figure 9(b) down to 0.06M. There is no evidence for a turnover of the mass function at low masses. The Class II mass function is indistinguishable from the mass spectrum of the pre-stellar condensations (Motte et al., 1998), which suggests that the origin of the IMF is found at the early fragmentation stage of star formation.

7. Conclusions

The ISO observations have provided a fundamental database for studying the earliest stages of stellar evolution. From the ISO spectroscopic and photometric observations it has been possible to trace the evolution of the dust and gas emission properties during the earliest stages ISO OBSERVATIONS OF YSOs 177 of stellar formation. The appearance of a new protostar and the gradual dispersal of the initial dusty envelope is evidenced by the change in the shape of the mid- and far-IR SEDs of the YSOs. From these SEDs, average dust temperatures are derived, which increase from ∼10 K in the pre-stellar cores, to ∼20–40 K in the Class 0 phase, to temperatures of few hundreds of Kelvin in Class I sources, when the appearance of the inner and warmer regions produce significant temperature stratifications in the observed SEDs. For high-mass stars, a similar sequence of increasing temperatures can be seen passing from the deeply embedded phase to the UC H II regions. The ISO spectra, on the other hand, have evidenced for the first time how the prevailing gas excitation mechanisms change during protostellar evolution. In the dense pre-stellar cores, the gas is only externally heated by the local inter-stellar field, while, as soon as a protostellar object is formed, shock excitation, either due to accretion or outflows, dominates the gas heating producing a rich far-IR molecular spectrum. Finally, during the latest accretion phase, the stellar radiation field becomes the dominant source of heating causing, in high-mass stars in particular, the excitation of many atomic and ionic forbidden lines. In addition to the study of individual objects, the ISO photometric surveys have given a census of the IR-excess populations in nearby clouds which is about twice as large as previously known. These complete samples of young YSOs provide a unique database, which can be used for the selection of sources to be observed by new ground-based and space-born facilities. The present and future mid- and far-IR missions, such as Spitzer Space Tele- scope, Herschel, SOFIA and JWST, will certainly largely improve the observational scenario given by ISO, mainly by providing data acquired with a larger spatial and spectral resolution, and better sensitivity. All of these missions, however, will have in the ISO database a fundamental reference point for the definition and interpretation of their observations.

References

Abergel, A., Bernard, J. P., Boulanger, F., et al.: 2002, A&A 389, 239. Alexander, R. D., Casali, M. M., André,P., Persi, P., and Eiroa, C.: 2003, A&A 401, 613. André,P., Ward-Thompson, D., and Barsony, M.: 2000, in: Mannings, V., Boss, A. P., and Russell, S. S. (eds.), Protostars and Planets IV, University of Arizona Press, Tucson, p. 59. André,P., Ward-Thompson, D., and Motte, F.: 1996, A&A 314, 625. Bacmann, A., André,P., Puget, J.-L., Abergel, A., Bontemps, S., and Ward-Thompson, D.: 2000, A&A 361, 555. Beichman, C. A., Myers, P. C., Emerson, J. P., et al.: 1986, ApJ 307, 337. Benson, P. J. and Myers, P. C.: 1989, ApJS 71, 89. Bontemps, S., André,P., Kaas, A. A., et al.: 2001, A&A 372, 173. Boogert, A. C. A., Ehrenfreund, P., Gerakines, P. A., et al.: 2000, A&A 353, 349. Boogert, A. C. A., Hogerheijde, M. R., Ceccarelli, C., et al.: 2002, ApJ 570, 708. Boonman, A. M. S. and van Dishoeck, E. F.: 2003, A&A 403, 1003. 178 B. NISINI ET AL.

Brooks, K. J., Cox, P., Schneider, N., et al.: 2003, A&A 412, 751. Ceccarelli, C., Boogert, A. C. A., Tielend, A. G. G. M., Caux, E., Hogerheijde, M. R., and Parise, B.: 2002, A&A 395, 863. Ceccarelli, C., Castets, A., Caux, E., Loinard, L., Molinari, S., and Tielens, A. G. G. M.: 2000, A&A 355, 1129. Cernicharo, J., Perez-Martinez, S., González-Alfonso, et al.: 1999, in: P. Cox and M. F. Kessler (eds.), The Universe as Seen by ISO, ESA-SP 427, p. 651. Correia, J., Griffin, M., and Saraceno, P.: 2004, A&A 418, 607. Churchwell, E.: 2002, ARA&A 40, 27. Foster, P. N. and Chevalier, R. A.: 1993, ApJ 416, 303. Froebrich, D., Smith, M. D., and Eislöffel, J.: 2002, A&A 385, 239. Froebrich, D., Smith, M. D., Hodapp, K. W., and Eislöffel, J.: 2003, A&A 346, 163. G˚alfalk, M.: 1999, Master Thesis, Stockholm University, see www.astro.su.se/English/thesis/ MagnusG xjobb.pdf. G˚alfalk, M., Olofsson, G., Kaas, A. A., et al.: 2004, A&A 421, 623. Gerakines, P. A., Whittet, D. C. B., Ehrenfreund, P., et al.: 1999, ApJ 522, 357. Gibb, E. L., Whittet, D. C. B., Boogert, A. C. A., and Tielens, A. G. G. M.: 2004, ApJSS 151, 35. Giannini, T., Nisini, B., and Lorenzetti, D.: 2001, ApJ 555, 40. Giveon, U., Sternberg, A., Lutz, D., Feuchtgruber, H., and Pauldrach, A. W. A.: 2002, ApJ 566, 880. González-Alfonso,E., Cernicharo, J., van Dishoeck, E. F., Wright, C. M., and Heras, A.: 1998, ApJ 502, L169. Greene, T. P., Wilking, B. A., André,P., Young, E. T., and Lada, C. J.: 1994, ApJ 434, 614. Harwit, M., Neufeld, D. A., Melnick, G. J., and Kaufman, M. J.: 1998, ApJ 497, L105. Hollenbach, D.: 1997, in: Reipurth, B. and Bertout, C. (eds.), Herbig-Haro Flows and the Birth of Low Mass Stars, IAU Symposium 182, p. 181. Hollenbach, D. and Tielens, A. G. G. M.: 1997, ARA&A 35, 179. Kaas, A. A., Olofsson, G., Bontemps, S., et al.: 2004, A&A 421, 623. Kaas, A. A., Olofsson, G., Bontemps, S., et al.: 1999, The Universe as Seen by ISO, ESA-SP 427, p. 493. Kaas, A. A. and Bontemps, S.: 2000, in: Montmerle, T. and André,P. (eds.), From Darkness to Light: Origin and Evolution of Young Stellar Clusters, ASP Conference Proceedings 243, p. 367. Kaufman, M. J. and Neufeld, D. A.: 1996, ApJ 456, 611. Keane, J. V., Tielens, A. G. G. M., Boogert, A. C. A., Schutte, W. A., and Whittet, D. C. B.: 2001, A&A 376, 254. Kirk, J. M., Ward-Thompson, D., and André,P.: 2005, MNRAS 360, 1506 Krause, O., Lemke, D., Tóth,L. V., Klaas, U., Haas, M., Stickel, M., and Vavrek, R.: 2003, A&A 398, 1007. Kurtz, S., Cesaroni, R., Churchwell, E., Hofner, P., and Walmsley, C. M.: 2000, in: Mannings, V., et al. (eds.), Protostars & Planets IV, University of Arizona, Tucson, p. 299. Lada., C.: 1999, in: C. J. Lada and N. D. Kylafis (eds.), The Origin of Stars and Planetary Systems, Kluwer Academic Publishers, Dordrecht. Lada, C. J. and Wilking, B. A.: 1984, ApJ 287, 610. Lahuis, F. and van Dishoeck, E. F.: 2000, A&A 355, 699. Larsson, B., Liseau, R., and Men’shchikov, A. B., et al.: 2000, A&A 363, 253. Larsson, B., Liseau, R., and Men’shchikov, A. B.: 2002, A&A 386, 1055. Lefloch, B., Cernicharo, J., Cabrit, S., Noriega-Crespo, A., Moro-Martín, A., and Cesarsky, D.: 2003, ApJ 390, L41. Lehtinen, K., Haikala, L. K., Mattila, K., and Lemke, D.: 2001, A&A 367, 311. Liseau, R., Giannini, T., Nisini, B., et al.: 1997, in: Reipurth, B. and Bertout, C. (eds.), Herbig-Haro Flows and the Birth of Low Mass Stars, Vol. 111. ISO OBSERVATIONS OF YSOs 179

Liseau, R., Ceccarelli, C., Larsson, B., et al.: 1996, A&A 315, L181. Maret, S., Ceccarelli, C., Caux, E., Tielens, A. G. G. M., and Castets, A.: 2002, A&A 395, 573. Molinari, S., Noriega-Crespo, A., Ceccarelli, C., et al.: 2000, ApJ 538, 698. Molinari, S., Noriega-Crespo, A., and Spinoglio, L.: 2001, ApJ 547, 292. Moro-Mart´ın,A., Noriega-Crespo, A., Molinari, S., Testi, L., Cernicharo, J., and Sargent, A.: 2001, ApJ 155, 146. Motte, F., André,P., and Neri, R.: 1998, A&A 336, 150. Motte, F., AndréP., Ward-Thompson, D., and Bontemps, S.: 2001, A&A 372, L41. Myers, P. C. and Benson, P. J.: 1983, ApJ 266, 309. Myers, P. C., Heyer, M., Snell, R. L., and Goldsmith, P. F.: 1988, ApJ 324, 907. Neufeld, D. A., Melnick, G. J., and Harwit, M.: 1998, ApJ 506, L75. Nisini, B.: 2003, in: Gry, C., Peschke, S. B., Matagne, J., et al. (eds.), Proceedings of the Symposium ‘Exploiting the ISO Data Archive – Infrared Astronomy in the Internet Age’, ESA-SP 511, p. 111. Nisini, B., Benedettini, M., Giannini, T., et al.: 1999, A&A 350, 529. Nisini, B., Benedettini, M., Giannini, T., Codella, C., Lorenzetti, D., Di Giorgio, A., and Richer, J. S.: 2000, A&A 360, 297. Nisini, B., Giannini, T., and Lorenzetti, D.: 2002, ApJ 574, 246. Nordh, L., Olofsson, G., Abergel, A., et al.: 1996, A&A 315, L185. Olofsson, G., Huldtgren, M., Kaas, A. A., et al.: 1999, A&A 350, 883. Peeters, E., Mart´ın-Hernández,N. L., Dufour, F., et al.: 2002, A&A 381, 571. Persi, P., Marenzi, A. R., Olofsson, G., et al.: 2000, A&A 357, 219. Persi, P., Marenzi, A. R., Gómez,M., and Olofsson, G.: 2003, A&A 399, 995. Reipurth, B. and Bally, J.: 2001, ARA&A 39, 403. Rosenthal, D., Bertoldi, F., and Drapatz, S.: 2000, A&A 356, 705. Saraceno, P., Nisini, B., Benedettini, M., et al.: 1999, in: Cox, P. and Kessler, M. (eds.), Proceedings of the Conference ‘The Universe as Seen by ISO’, ESA-SP 427, p. 575. Shu, F. H.: 1977, ApJ 214, 488. Smith, M. D., Froebrich, D., and Eislöffel, J.: 2003, ApJ 592, 245. van den Ancker, M. E., Tielens, A. G. G. M., and Wesselius, P. R.: 2000, A&A 358, 1035. van der Tak, F. F. S., van Dishoeck, E. F., Evans, N. J., and Blake, G. A.: 2000, ApJ 537, 283. van Dishoeck, E. F.: 2004, ARA&A 42, 119. van Dishoeck, E. F., Wright, C. M., Cernicharo, J., González-Alfonso,E., de Graauw, Th., Helmich, F. P., and Vandenbussche, B.: 1998, ApJ 502, L173. Ward-Thompson, D., André,P., and Kirk, J. M.: 2002, MNRAS 329, 257. Ward-Thompson, D., Motte, F., and André,P.: 1999, MNRAS 305, 143. Ward-Thompson, D., Scott, P. F., Hills, R. E., and André,P.: 1994, MNRAS 268, 276. White, G. J., Liseau, R., Men’shchikov, A. B., et al.: 2000, A&A 364, 741. Whitworth, A. P., Bhattal, A. S., Francis, N., and Watkins, S. J.: 1996, MNRAS 283, 1061. Whitworth, A. and Summers, D.: 1985, MNRAS 214,1. Wilgenbus, D., Cabrit, S., Pineau Des Forêts,and Flower, D.: 2001, ASP Conference Proceedings, Vol. 243, p. 347. Williams, J. P., Blitz, L., and McKee, C. F.: 2000, in: Mannings, V., Boss, A. P., and Russell, S. S. (eds.), Protostars and Planets IV, University of Arizona Press, p. 97. Wright, C. M., Drapatz, S., Timmermann, R., et al.: 1996, A&A 315, L301. Wright, C. M., van Dishoeck, E. F., Black, J. H., Feuchtgruber, H., Cernicharo, J., González-Alfonso, E., et al.: 2000, A&A 358, 689. PRE-MAIN SEQUENCE STARS SEEN BY ISO

DARIO LORENZETTI INAF, Osservatorio Astronomico di Roma, Via Frascati 33-00040 Monte Porzio (Roma), Italy (E-mail: [email protected])

(Received 7 July 2004; Accepted in ﬁnal form 13 December 2004)

Abstract. Observations of pre-main sequence objects (T Tauri, Herbig Ae/Be and FU Orionis stars) obtained with the instrumentation on board the Infrared Space Observatory (ISO) are reviewed. All the observations have been mainly carried out by using the two spectrographs SWS and LWS, adopting their low resolution modes and such data have been used both for lines detection and to reconstruct the spectral energy distributions. Line emission and photometric behaviour of pre- main sequence objects have been analyzed in the framework of the current models, discussing the agreement (or disagreement) with them and trying to derive the questions which should be answered by the forthcoming FIR instrumentation.

Keywords: pre-main sequence, circumstellar matter, ISO, infrared spectroscopy

1. Introduction

The Infrared Space Observatory (ISO) has been the first facility outside the atmosphere able to provide, through its instrumental capabilities, a systematic diagnosis of the phenomena (photodissociation, shocks, PAH and dust emission) occurring in the close environment of the pre-main sequence stars. In this area of astrophysical interests, the importance of the spectroscopical approach has largely exceeded that of the photometric one: indeed the IRAS survey had already made possible the first evaluation of the amount of energy emitted by these systems, not allowing, however, any investigation of the mechanisms responsible for the gas excitation. Despite its peculiar nature of Observatory dedicated to selected targets, ISO has provided, in the specific topic of the pre-main sequence, also complete surveys of significant samples, whose unbiased nature has revealed as crucial, also representing a useful starting point for the FIR/sub-mm facilities of the next future. For reason of compactness, the present review will account only for published papers based on ISO data on this topic (37 during the period 1996 – beginning of 2004), although interesting information can be also retrieved from more than 50 conference proceedings and about 10 PhD theses dealing essentially with ISO studies on Herbig Ae/Be stars. The reader is also referred to the excellent review by van Dishoeck (2004) which gives a complete analysis of the ISO spectroscopic results. Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, The Netherlands, and the United Kingdom), and with the participation of ISAS and NASA. Space Science Reviews (2005) 119: 181–199 DOI: 10.1007/s11214-005-8064-z C Springer 2005 182 D. LORENZETTI

2. T Tauri and Herbig Ae/Be Stars

2.1. ISO OBSERVATIONS

Data from the spectrometers on board ISO (mainly SWS and LWS and PHOT-S to a lesser extent) allowed to obtain low resolution (R ∼ 200–1000) spectra between 3 and 200 μm for a significant samples of T Tauri (TTS) and Herbig AeBe (HAEBE) stars. Due to the intrinsic faintness of its members, the former class has been less investigated than the latter, and all the objects observed with ISO are listed in the first five columns of Table I, along with an indication of the papers referring to them. The binary system formed by T Tau (N and S) has been revealed as the pre-main sequence object richest of atomic and molecular lines (van den Ancker et al., 1999; Spinoglio et al., 2000). Two surveys, the one on pure rotational H2 lines (Thi et al., 1999, 2001), and the second on silicate features (Natta et al., 2000) have obtained significant results discussed in the following sections. Concerning the HAEBE, a complete compilation of SWS and LWS data has been recently presented by Elia et al. (2004) who searched the ISO mission data bank for the 108 bona fide HAEBE stars listed in the catalog of Thé et al. (1994, Table I) obtaining 36 objects with at least one SWS or LWS observation of good quality. More recently, Acke and van den Ancker (2004), in their study of disks around 46 HAEBE, have studied all the ISO spectra obtained with different facilities, namely

TABLE I ISO Observations of T Tauri and FU Orionis stars.

TTS SWS LWS TTS PHOT-S FU Ori SWS LWS

AA Tau 1 2, 4 CT Cha 7 RNO 1B 3, 9 DM Tau 1 Glass I 7 Z CMaa 10 3, 9 DR Tau 1 LkHα 332–20 7 V346 Nor 3, 9 GG Tau 1, 8 4 SX Cha 7 Re 13b 3, 9 GO Tau 1 VW Cha 7 V1057 Cyg 9 RY Tau 1 VZ Cha 7 V1331 Cyg 9 T Tau 5 3, 6 WX Cha 7 V1735 Cyg 3, 9 DG Tau 3 WW Cha 7 HL Tau 3 XX Cha 7 GM Aur 1 LkCa 15 1 2

aConsidered also as a HAEBE (see Table II). bIt is not a FU Ori. It lies in proximity of V346 Nor (less than 1 arcmin). References: 1. Thi et al. (2001); 2. Chiang et al. (2001); 3. Pezzuto et al. (2002); 4. Creech-Eakman et al. (2002); 5. van den Ancker et al. (1999); 6. Spinoglio et al.; (2000) 7. Natta et al. (2000); 8. Thi et al. (1999); 9. Lorenzetti et al. (2000); 10. Acke and van den Ancker (2004). PRE-MAIN SEQUENCE STARS SEEN BY ISO 183

SWS, ISOPHOT-S and ISOCAM CVF scans. Other ISOCAM CVF scans of 8 nearby HAEBE have been presented by Siebenmorgen et al. (2000). To have a complete census of the available ISO data concerning the HAEBE, spectrophotometric data have been complemented with far IR ISOPHOT photometry (obtained with 10 broad-band ﬁlters from 4.8 to 200 μm, partially overlapping the IRAS ones) of seven HAEBE (Abraham et al., 2000). All the investigated HAEBE are listed in Table II.

2.2. THE SPECTRAL ENERGY DISTRIBUTIONS

Low resolution ISO spectra (3–200 μm) have been extensively used, also in combination with literature data sets (UV, optical, near IR, sub-mm, radio) to construct accurate spectral energy distributions (SEDs). Only 3 TTS (namely AA Tau, GG Tau and LkCa 15) have been investigated in this sense, and the results are incorporated in the following discussion dealing with HAEBE. For these latter the dust continuum behaves very differently from source to source: some objects present their SED rising with the wavelength (up to about 100 μm), other rather flat or even declining (see Figure 1), confirming the fact that HAEBE do not constitute a homogeneous class of object, at variance with the T Tauri class. ISO data have played a major role in creating some convergence on the long- standing debate about the spatial distribution (spherical envelopes vs. flattened structures) of the circumstellar matter around HAEBE. The reason for that relies on the fact they represent a common starting point for different authors for evaluating the capabilities of the different models, namely a data-base which covers, at spectrophotometric resolution, a rather wide spectral range. Nevertheless we have to remember that the ISO spatial resolution is relatively poor, therefore ISO results have to be treated with a certain care and need sometime a further confirmation. The SWS data from 14 isolated Herbig stars with A–F spectral types, have been firstly presented by Meeus et al. (2001), who defined two key groups based upon

Figure 1. ISO spectra of two HAEBE, as obtained after accurate reduction of the spectrometer observations, showing two substantially different SEDs. 184 D. LORENZETTI

TABLE II ISO Observations of HAEBE.

Source SWS LWS CAM PHOT

LkHα198 1, 5, 27 3 V376 Cas 1 1, 5, 27 3 VX Cas 2 Elias 3-1 1, 2, 26, 30 AB Aur 1, 2, 8, 12, 13, 14, 15, 30 1, 9, 13 MWC 480 1, 2, 30 1, 9, 11, 28 UX Ori 30, 31 2 HD 34282 1, 2 1, 28 HD 34700a 22 HD 35187a 2 HD 36112 30 1, 9, 11, 28 CQ Tau 30 1, 9, 10, 11, 28 BF Ori 2 RR Tau 2 MWC 137 1 1, 28 Z CMab 1, 2 1, 27 He 3-554 2 2 HD 97048 1, 2, 26, 32 1, 5, 27 2, 4 2 HD 97300a 4, 20 7 HD 100453a 2, 8, 12, 15 2 HD 100546 1, 2, 8, 12, 15, 17, 29, 32 1, 17, 28 2 HD 104237 1, 2, 8, 12, 15, 31 1, 28 2 IRAS 12496-7650 1, 2 1, 5, 6, 18 2 HD 135344a 2, 8, 12, 15 2 HD 139614a 2, 8, 12, 15 2 HD 141569 1, 2 1, 28 2 HD 142527a 2, 8, 12, 15, 22 22 2 HD 142666 1, 2, 8, 12, 15, 31 1, 28 2 HD 144432 1, 2, 8, 12, 15, 31 1, 28 2 HR 5999 1, 2 2, 4 2 HR 6000a 4 He 3-1191 1, 2 HD 150193 1, 2, 8, 12, 15, 30, 31 2 AK Scoa 22 CoD−42◦11721 1, 2, 21, 32 1, 5, 27 2 KK Oph 32 HD 163296 1, 2, 8, 12, 13, 14, 15, 30, 31 2

(Continued on next page) PRE-MAIN SEQUENCE STARS SEEN BY ISO 185

TABLE II (Continued)

Source SWS LWS CAM PHOT

HD 169142a 2, 8, 12, 15 2 MWC 297 1, 2, 25 1, 5, 25, 27 VV Ser 2 MWC 300 1 TY CrA 1 1, 28 4 RCrA 1,2 1,5,6,27 2 T CrAa 227 2 HD 179218 1, 2, 8, 12, 15, 32 2, 4 HD 176386a 4 WWVul 1,2,30 2,10 BD+40◦4124 1, 2, 19, 24, 32 1, 9, 24, 28 2 LkHα224 1, 2, 24 1, 24 LkHα225a 2, 24 24 PV Cep 1, 2 1, 5, 27 2, 3 HD 200775 1, 2, 16, 23 1, 16, 28 2 V645 Cyg 1, 2 1, 5, 27 BD+65◦1637 3 BD+46◦3471 2 SV Cep 2 2 LkHα234 1, 5, 6, 27 3 LkHα233 1, 9, 28 3 MWC 1080 1, 2, 21 1, 5, 27 3

aNot listed in the Thé et al., Table I. bConsidered also as a FU Ori object (see Table I). References: 1. Elia et al. (2004); 2. Acke and van den Ancker (2004); 3. Abraham et al. (2000); 4. Siebenmorgen et al. (2000); 5. Lorenzetti et al. (1999); 6. Giannini et al. (1999); 7. Gürtler et al. (1999).; 8. Dominik et al. (2003); 9. Creech-Eakman et al. (2002); 10. Natta et al. (2001); 11. Chiang et al. (2001); 12. Meeus et al. (2001); 13. Bouwman et al. (2000); 14. van den Ancker et al. (2000a); 15. Bouwman et al. (2001); 16. Fuente et al. (2000); 17. Malfait et al. (1998); 18. Nisini et al. (1996); 19: Wesselius et al. (1996); 20. Siebenmorgen et al. (1998); 21. Benedettini et al. (1998); 22. Malfait et al. (1999); 23. Fuente et al. (1999); 24. van den Ancker et al. (2000b); 25. Benedettini et al. (2001); 26. van Kerckhoven et al. (2002); 27. Pezzuto et al. (2002); 28. Lorenzetti et al. (2002); 29. Bouwman et al. (2003); 30. Thi et al. (2001); 31. van Boekel et al. (2003); 32. Peeters et al. (2002). the spectral shape of the infrared region. The first group contains sources for which the continuum can be reconstructed by a power law (optically thick, geometrically thin disk) and by a black body at about 200 K (flare region); the second group sources only need a power law to fit their continuum. The presence or not of a solid state bands contribution further subdivides both groups. They explain the observed 186 D. LORENZETTI

SEDs of the second group by assuming that the inner part of the disk is partially optically thick, shielding the outer parts from direct stellar radiation, which prevents the disk from flaring, therefore the black body component is not observed. They also conclude that evidence exists for grain growth in the HAEBE disks: the first group sources may evolve into the second one which has evidence for larger grains and lacks of flaring. The same 14 HAe have been also considered by Dominik et al. (2003), who using the Dullemond et al. (2001) model, find that the different spectra can be interpreted as pure disk spectra without additional components. They have obtained satisfactory fits and deduced the disk parameters for a number of sources. They show how the first group sources can be fitted with disk with depleted inner region, while the second group requires more compact structures where the outer region are probably collapsed. Four different Herbig stars, always selected to show a late spectral type (A–F), and observed with ISOPHOT and LWS, have been studied by Natta et al. (2001). They, applying models of irradiated disk, propose how to solve the puzzle of the excess of near IR emission, with a peak at about 3 μm, as due to dust evaporation in disks where the gas component is optically thin to the stellar radiation, as expected if the accretion rate is very low. Creech-Eakman et al. (2002) and Chiang et al. (2001) have fitted the LWS spectra of 6 HAEBE (along with3TTauri stars) with a revised version of the Chiang and Goldreich (1997) model. Only a subset of HAe stars is adequately fitted, showing that the flux at wavelength longward of 55 μm arises from the cooler disk interior at large radius, which is heated diffusively by reprocessed radiation from the disk surface. In particular, Chiang et al. (2001) point out how the SED alone does not furnish sufficient information to uniquely constrain the geometry surrounding HAe stars. A different view is offered by Elia et al. (2004) who compare updated set of observational data spanning 5 orders of magnitude in wavelength and including SWS and LWS spectra, with the model spectra computed for different possible distribution of the circumstellar matter. In their previous model of a spherically symmetric distribution (Pezzuto et al., 1997), cases have been now considered with the circumstellar regions partially evacuated along the polar axis by the action of the stellar wind. The inclusion of the polar cavities indirectly allows geometries in which a small scale disklike flared structure around the central star is present. As a result on a sample of 36 HAEBE, they found that 17 objects are reasonably fitted, eight of which with a pure spherical model, and the remaining nine with the inclusion of the polar cavities. These latter objects are those in common with the Dominik et al. (2003) sample of isolated HAe stars. Such an agreement gives further support to the hypothesis that for low-mass, more evolved HAe stars, the circumstellar geometry can be similar to that found around the T Tauri stars. Moreover they conclude that a dust absorption coefficient typical of larger dust grains has to be associated to these eight cases. This is in agreement with the results by Meeus et al. PRE-MAIN SEQUENCE STARS SEEN BY ISO 187

(2001) and confirms the idea that the grain size is expected to increase with time in circumstellar envelopes as a result of the ongoing growth process activated in high density environments. Opening a new spectral window longward of 100 μm, ISO has allowed a classification of Young Stellar Objects (YSOs) more accurate than that provided by IRAS. A new tool has been introduced by Pezzuto et al. (2002) in form of two colours diagram [60–100] vs. [100–170]. The pre-main sequence stars cluster in a definite locus of the plot. The median colours are [60–100] = 0.11 ± 0.14 and [100–170] =−0.01 ± 0.12, corresponding to a colour temperature of about 50 K. This represents a powerful and simple tool to derive indications from the future photometric surveys. The only mid to far IR SEDs of HAEBE constructed with pure photometric ISO data have been presented by Abraham et al. (2000), who examined the ISOPHOT scans obtained with 10 large band filters between 4.8 and 200 μm) across 7 HAEBE (from B0 to A7). The observed SEDs can be described by a power law relationship, but their shapes are inconclusive for ruling out the presence of circumstellar disks or spherically symmetric envelopes. The emission at wavelengths larger than 100 μm is dominated by cold objects, unresolved by the 90 × 90 arcsec2 pixel of the 200 μm camera. These observations support the findings of previous sub-mm maps (Sandell and Weintraub, 1994; Henning et al., 1998) about the existence of dust cores (of ≤1 size) towards many HAEBE. Whether or not these cores coincide with huge massive envelopes (i.e. they are centrosymmetric structures), remains an open question.

2.3. DUST FEATURES AND PAH

Studies of the chemical and mineralogical composition of the dust have received a strong impulse by the ISO advent. It has allowed, for the first time, an unbiased spectroscopical approach to the complete wavelength regime (3–20 μm), revealing a large variety in dust properties, from small PAHs (features at 3.3, 6.2, 7.7, 8.6, 11.3, 12.7 μm – e.g. Légerand Puget, 1984), nanodiamond (NAN) features (3.4, 3.5 μm – van den Kerckhoven et al., 2002), to amorphous (9.7, 18 μm) and crystalline (11.3 μm) silicate dust. Disks around 9 TTS associated with the Cham I dark cloud have been studied by Natta et al. (2000) and Gürtler et al. (1999). All stars show evidence of amorphous silicate emission at 10 μm, from (sub)-micron sized grains, which is the only prominent feature in those spectra. They discuss how the emission originates in the disk atmosphere, namely the hot layer of the circumstellar disk, which, heated by the stellar radiation, is hotter than the disk midplane and therefore the emission feature can be seen. Chiang et al. (2001) investigate how the SED of the face-on disk around 2 TTS and 3 HAEBE depends on dust properties and grains distributions. A systematic investigation of ISO SWS and LWS spectra aimed to study the characteristics of the circumstellar dust around a significant sample of HAEBE has 188 D. LORENZETTI been conducted by Meeus et al. (2001). Beside their results on the dust distribution, discussed in the previous section, they also found that the PAH bands are different from the PAH band in the ISM and are only present in HAEBE with a large amount of warm dust. PAH cannot be correlated to any of the stellar parameters and are located in an extended region around the disk. Their presence is unrelated to the presence of the silicate bands and these latter are surprisingly laking in some sources. Acke and van den Ancker (2004) presented a complete atlas of all the features detected in 46 HAEBE retrieving the spectra (obtained with SWS, PHT-S, CAM-CVF) from the ISO archive. The emission detection rates of PAH, amorphous and crystalline silicates are 57, 52 and 24%, respectively, with some more embedded objects showing silicate in absorption. Following Meeus et al. (2002) and van Boekel et al. (2003), they also provide a classification of HAEBE in two groups based on the ratio between the NIR (1–5 μm) and mid-IR (20–100 μm) luminosity: the systems with stronger mid-IR (relative to their NIR), display significantly more PAH emission than those with weaker mid-IR. This circumstance has been interpreted as an observational support for the disk model by Dullemond et al. (2001) (see Section 2.2) in which flaring disks are associated with stronger mid-IR emission, while shielded systems with weaker contributions. Silicates behave differently: since they are expected to be produced mainly in the inner parts of the disk, no strong difference is expected between the two groups. The composition of the circumstellar dust around four individual HAe has been studied in detail through SWS and LWS spectra. Silicates in the crystalline form (similar to those found in comets), are found in HD100546 (Malfait et al., 1998); C-rich and O-rich components are respectively associated to the warm and cool dust component of HD142527 (Malfait et al., 1999), but no clear correlation exists between dust signatures and stellar parameters. The stars AB Aur and HD163296 (Bouwman et al., 2000; van den Ancker et al., 2000) also show the presence of crystalline silicates and larger (∼1 mm) grain sizes in HD 163296 compared to AB Aur suggests the former as a more evolved system. Bouwman et al. (2000) confirm the poor correlation between the degree of crystallization and the age of the central star and note that crystallization and grain growth are not necessarily coupled. Both these facts signal how the evolution of circumstellar dust in protoplanetary disks happens very rapidly or how such evolution is driven by other factors (than stellar mass and age). The good quality of the SWS data has allowed Peeters et al. (2002) to derive a striking aspect of the features in the 6–9 μm region, namely their variability: all the features shift in peak position from source to source, show different profile and each seems to be composed of several subfeatures. The HAEBE constitute a small sub-sample among the variety of the investigated objects: their observed variations are discussed in terms of the nature of the carrier and local physical condition, although some aspects remain still unknown. PRE-MAIN SEQUENCE STARS SEEN BY ISO 189

2.4. EMISSION LINES

Given the strong underlying continua, combined with the low spectral resolution of the ISO spectrometers (ISOCAM-CVF – ISOPHOT-S – SWS – LWS), rather strong emission lines have been detected, often for the ﬁrst time. The inventory of these lines is given in Table III. The following discussion will address separately the spectral surveys and the studies of individual regions.

2.4.1. Spectral Surveys Searches for the H2 pure rotational lines have been performed with SWS. A quite large sample of 7 TTS and 8 HAEBE with reasonable evidence for disks has been investigated by Thi et al. (1999, 2001). They report the detected fluxes (or, in few cases, the 3σ upper limits) of the S(0) and S(1) lines at 28.2 and 17.0 μm, respectively. Their conclusion that these lines arise in a warm gas disk component, which constitutes a significative fraction of the entire disk mass, has not been confirmed by the subsequent non-detection of the S(1) line in some of the same objects (Richter et al., 2002), which conversely suggests that H2 emission could be extended on scales of 5 or more. However, because of their sensitivity to small amount of warm gas, the mid-IR H2 lines remain crucial for more complete studies with future IR spectrometers. In the LWS range the ISO archive has been systematically searched to obtain a complete far IR spectrophotometric survey of HAEBE (Lorenzetti et al., 1999, 2002; Giannini et al., 1999). The investigated sample is constituted by 26 objects all presenting fine structure emission lines from different ions; 6 out of these also show molecular emission in form of CO and OH lines usually associated to local density peaks. The [O I] 63, 145 μm and [C II] 158 μm are the strongest features observed and have been used as a diagnostic of the excitation mechanism and then of the physical conditions. Line ratios and corresponding errors are indicated in Figure 2a. The observed line intensity ratios are superimposed on a grid of photodissociation models (Kaufman et al., 1999). The substantial difference between line intensity predictions of this model and those of other PDR models (e.g. Tielens and Hollenbach, 1985) is the inclusion of an additional heating: the ejection of photoelectrons from PAH and small grains. These latter are essential components of the HAEBE close environments, as demonstrated by many dedicated observations (see e.g. Meeus et al. (2001) and references therein). The extra heating source is essential because for a given value of density (n) and radiation field (Go), this model predicts increased O I and C II column densities, thus allowing to move at lower n and Go the transition point between CII dominated and OI dominated cooling. The location of the data points in Figure 2a, gives support to the interpretation in favour of phodissociation as dominating mechanism for far IR line excitation (see also Creech-Eakman et al., 2002). The existence, around the HAEBE stars, of a PDR originated by stellar rather than interstellar FUV photons has received further support by the correlation between [C II] 158 μm luminosity 190 D. LORENZETTI

TABLE III Emission lines observed in Pre-Main Sequence objects.

Wavelength Identiﬁcation Sources (μm)

2.4–7.5 H: Br, Pf, Hu CoD−42◦11721 (6), MWC 297 (7), BD+40◦4124 (9), Series HD200775 (10), MWC 1080 (6), ◦ 6.91 H2 S(5) BD+40 4124 (9), LkHα224 (9, 11), HD200775 (10) ◦ 8.02 H2 S(4) BD+40 4124 (9, 11), LkHα224 (9, 11), HD200775 (10) ◦ 9.66 H2 S(3) BD+40 4124 (9, 11), LkHα224 (9, 11), HD200775 (10) ◦ 12.28 H2 S(2) BD+40 4124 (9, 11), LkHα224 (9, 11), HD200775 (10)

17.03 H2 S(1) GG Tau (13), AB Aur (1), MWC 480 (1), HD36112 (1), CQ Tau (1), HD 150193 (1), HD163296 (1), BD+40◦4124 (9, 11), LkHα224 (9, 11), HD200775 (10) 3 3 ◦ 18.71 [S III] P2- P1 CoD−42 11721 (2) 3 3 25.25 [S I] P1- P2 LkHα224 (9) 6 6 ◦ 25.99 [Fe II] D7/2- D9/2 BD+40 4124 (9)

28.22 H2 S(0) GG Tau (13), AB Aur (1), UX Ori (1), CQ Tau (1), HD 150193 (1), HD163296 (1) 3 3 ◦ 33.48 [S III] P1- P0 CoD−42 11721 (2) 2 2 ◦ ◦ 34.81 [Si II] P3/2- P1/2 CoD−42 11721 (2), BD+40 4124 (9), LkHα224 (9), HD200775 (10) 3 3 ◦ 51.81 [O III] P2- P1 V1057 Cyg (12),BD+40 4124 (4, 8, 9) 2 2 57.33 [N III] P3/2- P1/2 V1057 Cyg (12) 3 3 63.18 [O I] P1- P2 RNO1B (12), ZCMa (12), V346 Nor (12), Re13 (12), V1057 Cyg (12), V1331 Cyg (12), V1735 Cyg (12), LkHα198 (2), V376 Cas (2), AB Aur (4, 8), CQ Tau (8), MWC 137 (4), HD97048 (2), HD100546 (4), HD104237 (4), IRAS 12496 (2, 5), CoD−42◦11721 (2), MWC 297 (2), TY CrA (4), R CrA (2), BD+40◦4124 (4, 8, 9), PV Cep (2), HD200775 (4, 10), V645 Cyg (2), LkHα234 (2), LkHα233 (4, 8), MWC 1080 (2) 65.69 CO 40-39 R CrA (2) 2 2 71.19 OH 1/2,7/2- 1/2,5/2 R CrA (3) 2 2 79.15 OH 1/2,1/2- 3/2,3/2 Z CMa (12), R CrA (2, 3), LkHα234 (3) 2 2 84.51 OH 3/2,7/2- 3/2,5/2 V1057 Cyg (12), R CrA (2, 3), LkHα234 (3) 3 3 88.36 [O III] P1- P0 V1057 Cyg (12), V1331 Cyg (12) CoD−42◦11721 (2), BD+40◦4124 (4, 8, 9), V645 Cyg (2) 3 3 121.80 [N II] P2- P1 V346 Nor (12), V1057 Cyg (12), V1331 Cyg (12), V1735 Cyg (12), CoD−42◦11721 (2), BD+40◦4124 (4, 8) 124.19 CO 21-20 R CrA (2) 130.37 CO 20-19 IRAS 12496 (2)

(Continued on next page) PRE-MAIN SEQUENCE STARS SEEN BY ISO 191

TABLE III (Continued)

Wavelength Identiﬁcation Sources (μm)

137.20 CO 19-18 IRAS 12496 (2, 3), R CrA (2, 3), LkHα234 (3) 144.78 CO 18-17 IRAS 12496 (2, 3), LkHα234 (3) 3 3 145.53 [O I] P0- P1 RNO1B (12), Z CMa (12), V346 Nor (12), V1331 Cyg (12), V1735 Cyg (12), MWC 137 (4), HD100546 (4), IRAS 12496 (2, 5), CoD−42◦11721 (2), MWC 297 (2), TY CrA (4), BD+40◦4124 (4, 8, 9), HD200775 (4, 10), LkHα234 (2), LkHα233 (4, 8), MWC 1080 (2) 153.27 CO 17-16 V1057 Cyg (12), IRAS 12496 (2, 3, 5), TY CrA (4), R CrA (2, 3), BD+40◦4124 (4, 9), LkHα234 (2, 3) 2 2 157.68 [C II] P3/2- P1/2 AA Tau (8), GG Tau (8), RNO1B (12), Z CMa (12), V346 Nor (12), Re13 (12), V1057 Cyg (12), V1331 Cyg (12), V1735 Cyg (12), LkHα198 (2), V376 Cas (2), AB Aur (4, 8), MWC 480 (4, 8), HD34282 (4), HD36112 (4, 8), CQ Tau (4, 8), MWC 137 (4), HD97048 (2), HD100546 (4), HD104237 (4), IRAS 12496 (2, 5), HD141569 (4), HD142666 (4), HD144432 (4), CoD−42◦11721 (2), MWC 297 (2), TY CrA (4), R CrA (2), BD+40◦4124 (4, 8, 9), PV Cep (2), HD200775 (4, 10), V645 Cyg (2), LkHα234 (2), LkHα233 (4, 8), MWC 1080 (2) 162.81 CO 16-15 Re13 (12), V1057 Cyg (12), IRAS 12496 (3, 5), TY CrA (4), R CrA (2, 3), BD+40◦4124 (4, 9), HD200775 (4), LkHα234 (2, 3) 173.63 CO 15-14 IRAS 12496 (3, 5), R CrA (2, 3), BD+40◦4124 (4, 9), LkHα234 (2, 3) ◦ 179.53 o-H2O212-101 BD+40 4124 (4) 186.00 CO 14-13 IRAS 12496 (3), R CrA (2, 3), BD+40◦4124 (4, 9), LkHα234 (2, 3)

T Tauri and FU Orionis stars are given in boldface. The remaining refer to HAEBE stars. An exception is done for T Tauri itself, whose spectrum presents a plenty of CO, OH, H2O and atomic lines which are not indicated in the present Table (see Figure 4). References: 1. Thi et al. (2001); 2. Lorenzetti et al. (1999); 3. Giannini et al. (1999); 4. Lorenzetti et al. (2002); 5. Nisini et al. (1996); 6. Benedettini et al. (1998); 7. Benedettini et al. (2001); 8. Creech- Eakman et al. (2002); 9. van den Ancker et al. (2000b); 10. Fuente et al. (2000); 11. Wesselius et al. (1996); 12. Lorenzetti et al. (2000); 13. Thi et al. (1999). vs. the bolometric luminosity of the central source (Lorenzetti et al. (1999), their Figure 5). Alternative models do not seem so much promising: J-shock models (e.g. Hollenbach and McKee, 1989) can be ruled out, since they predict substantially different line ratios whose intersection with the PDR plane is indicated in Figure 2a as the hatched area at the top right corner. Contributions from non dissociative 192 D. LORENZETTI

2 62.5 7 3 IRAS12496 3.54-5 1.5 5 1 RCrA 3 4 2 6 MWC361 1 3 0 4 TYCrA 5 Log n 0.5 5 Log n -1 -2 -1 0 1 2 -1 0 1 2

Figure 2. (a) Observed line ratios superimposed to PDR models by Kaufman et al. (1999). The hatched area identiﬁes the locus pertaining to the J-shock model predictions. (b) Observed ratios between molecular CO lines and ﬁne structure [O I] 145 μm and [C II] 158 μm superimposed on clumpy PDR models of Burton et al. (1990).

C-shocks (e.g. Draine et al., 1983) cannot be ruled out, but such mechanism can be considered as the main responsible for the gas excitation just in a narrow region of the parameter space (shock speed, magnetic field, pre-shock density). The molecular lines are usually weak (Giannini et al., 1999), but the authors do not believe the lack of detection is due only to an instrumental sensitivity limit, but it is also related to an intrinsic property of the circumstellar environment, namely the existence or not of some compact density peaks near the source where the column density is expected to substantially increase. For all sources the molecular emission has been consistently fitted with a Large Velocity Gradient (LVG) model and it results originated in a warm (T ≥ 200 K) and dense (n ∼ 105 cm−3)gas located in very compact regions having diameters of few hundreds of AU. By comparing the observed cooling ratios (L(CO)/L([OI]), L(CO)/L([CII])) with model predictions (Burton et al., 1990, see Figure 2b) the photondominated mechanism in high density or clumpy environments appears to be the likely source of excitation, while shocks can reasonably ruled out because of both the absence of any water vapour lines and the presence of OH lines, in contrast with predictions of molecular emission. Fine structure lines from ionized species [O III], [N II] are detected in those HAEBE having a spectral type B0 or earlier (see Table III), and these individual cases will be discussed in the next section. Finally, studies have been presented (Lorenzetti et al., 2002) to derive some general trend to be compared with similar behaviours from classes of younger objects. The far IR luminosity (LFIR) as due to all line emission contributions (fine structure line emission and total molecular cooling) is plotted in Figure 3 PRE-MAIN SEQUENCE STARS SEEN BY ISO 193

Figure 3. Far IR lines luminosity vs. bolometric luminosity of the central star. The equation representing the best linear fit to the points (dashed line) is given in the upper part. HAEBE are represented as filled circles, while younger Class 0 and Class I objects (see text) are indicated with triangles and open circles, respectively. as a function of the bolometric luminosity of the central object. The straight line, whose equation is given in the upper part of the Figure, represents the best linear fit to the points: the intercept on the y-axis is located at a LFIR value of about 10−4. A similar plot for the younger Class 0 objects (Nisini et al., 2002, here de- −2 picted as triangles) indicates LFIR 10 as the intercept value. An intermediate value of 10−3 is attributable to the Class I objects (open circles), which are expected to be somehow younger than HAEBE and approaching to their evolutionary stage. Although completely different mechanisms are invoked to account for the behaviour of the HAEBE and of the other objects (photodissociation vs. shock excitation), it is worthwhile noting how LFIR is a decreasing fraction of Lbol while the evolution goes on. The conversion of the bolometric luminosity of the central object into far IR line cooling in the circumstellar envelope occurs at progressively lesser efficiency while time elapses. Therefore ISO has allowed to extend to the far IR gas cooling the well consolidated statement according to which the rele- vance of the circumstellar vs. stellar properties diminishes during the pre-stellar evolution.

2.4.2. Emission Lines of Individual Objects Here, account is given for those TTS, HAEBE or nearby regions that have been individually investigated with the ISO spectrometers, through the analysis of emission lines different from those presented in the previous section, or through different techniques (e.g. spectroscopic mapping). T Tauri: This binary system presents the richest SWS (van den Ancker et al., 1999) and LWS (Spinoglio et al., 2000) spectrum (see Figure 4), which signals the presence of different excitation mechanisms: a partially ionized wind is likely 194 D. LORENZETTI

Figure 4. The ISO-LWS emission line spectrum of T Tau, from which the continuum has been subtracted. responsible for both the H I recombination lines and FIR fine structure lines through the onset of shocks. These latter are also traced by a low temperature (T ≈ 400 K) component of the pure rotational and rovibrational emission from H2. The role of the UV field is evidenced by both a higher temperature component (T ≈ 1500 K) of H2 and the overabundance of OH (more than a factor of 10). Most of the observed FIR molecular lines (high J (J = 14–25) CO, para- and ortho-H2O, OH) can be ∼ ∼ attribute to a very compact region (300–400 AU) with T 300–900 K and nH2 105–6 cm−3. CoD−42◦11721 and MWC1080: These two HAEBE (with Sp.T B0) have been studied by Benedettini et al. (1998). Their SWS spectra show hydrogen recombination lines of the Brackett, Pfund and Humphreys series (see Table III). The observed line decrements in each spectral series are consistent with emission from ionized winds, as expected for early type stars. Mass loss rates and reconsidered distance values have been derived along with the presence of ionization bounded compact regions of a few tens of stellar radii. In the case of CoD−42◦11721, a 2 −3 more diluted ionized region (with electron density ne ≈ 510 cm ), external to the optically thick H II region, is traced by the fine structure line ratios [S III] 33/19 and [O III] 88/52 (Lorenzetti et al., 1999). ISO observations indicate that density gradients, more than abundance gradients, are responsible for the observed line PRE-MAIN SEQUENCE STARS SEEN BY ISO 195 ratios, thus contrasting a previous interpretation in favour of an evolved nature of this star (Hutseméker and van Drom, 1990). MWC297: The same model of ionized wind has been applied to the plenty (more than 20) of recombination lines in the SWS range (Benedettini et al., 2001), providing consistent determinations of the distance (280 pc), the visual extinction 7 −1 (7.5 mag), the mass loss rate (9 × 10 M yr ) and the size of the emitting ionized region (30 stellar radii). BD+40◦4124: It belongs to a group together with the HAe LkHα 224 and the embedded source LkHα 225. In the first SWS study of this complex (Wesselius et al., 1996) the detection of H2 pure rotational lines is reported. The analysis indicates excitation temperatures of about 500 K and a very weak C-type shock wave driven into the cloud by stellar outflows as the most likely mechanism for the H2 excitation. More recently, van den Ancker et al. (2000) presented the spectra of both SWS and LWS of the three member of the group, where several contributions to the emission (pure rotational lines of H2, ro-vibrational transition of CO, PAHs,HI recombination lines and IR fine structure lines) were detected. Again, the data suggest that the HI lines of BD+40◦4124 arise in either a dense wind or a circumstellar disk, while a low density (≈102 cm−3) HII region combined with a clumpy PDR is responsible for the FIR fine structure lines. Lines from non-dissociative shock are observed in the lines of sight of LkHα 224 and LkHα 225 produced by a slow outflow driven by the latter, toward which also a dissociative shock is recognized, presumably located closer to this exciting source. NGC7023: It is a reflection nebulosity hosting the HAEBE star HD200775 and a very close (50) PDR that has been studied by Fuente et al. (1999), who provide a possible explanation (the advancing photodissociation front) for the nonequilibrium ortho-to-para ratio of 1.5–2 they found at variance with other ISO H2 observations toward several Galactic regions (see Habart, this issue). A larger portion of the same region, including also HD200775, has been investigated later (Fuente et al., 2000) and strong emission of the [Si II] 34.8 μm line has been detected towards the star, filling a ring constituted by the PDR tracers ([O I[, [C II]).This morphology can be only explained as arising in the diffuse gas filling the cavity around the star and its HII region. Indeed only up to AV ∼2 a significant fraction of Si remains in the gas phase, while for larger AV values Si is strongly depleted. NGC7129: LWS has also been used to provide a FIR spectroscopic map of this complex star forming region (Lorenzetti et al., 1999) where two HAEBE are located, ◦ namely BD+65 1637 and LkHα 234. [O I] 63,145 μm and [C II] 158 μm lines are detected everywhere in the mapped region and their intensity ratios are consistent with the model prediction of a PDR. The spatial distribution of the emission features indicates the presence of a peak which not coincide with the position of LkHα 234, but it seems more likely associated with the other HAEBE. This view is supported by the distribution of the 13CO emission obtained, at a comparable scale, by Fuente et al. (1998), where the cavity excavated by BD+65◦1637 is clearly evident. 196 D. LORENZETTI

3. FU Ori Systems

An other sample, interesting for having a complete picture of the pre-main sequence evolution, is represented by the FU Orionis stars: low luminosity objects which undergo a sudden increase in brightness. The widely accepted explanation of the FU Orionis phenomenon is a large increase in the accretion rate through a circumstellar disk surrounding a low luminosity young star (Hartmann and Kenyon, 1985). These events make the disk self luminous, with a luminosity which largely exceeds that of the central star (for more details see the review papers by Hartmann et al., 1993; Hartmann and Kenyon; and references therein). The observations obtained with the ISO instrumentation are reported in the last three columns of Table I. These data essentially consist of the LWS spectra presented by Lorenzetti et al. (2002). In addition to [O I]63μm and the [C II] 158 μm lines, that are commonly observed in all spectra, the observational novelty is the presence in most of the sources of the transition of ionized nitrogen [N II] 122 μm, which is not detected in other objects in a similar evolutionary phase. This line probes low ionization and low density 2 −3 material (ne < 5 × 10 cm ) not easily traced by other lines. No higher ionization stages of other elements are detected. It is worthwhile to mention that the [N II] detections in FU Ori are close to the limit of the observing capabilities and, as such, they could be arguable; however the circumstance to find this line practically at all the pointings on FU Ori objects, surely increases the statistical confidence in ascribing this property to this specific class. To emphasize this point in Figure 5 a sort of mean spectrum is reported: the result demonstrates that [N II] line is indeed real.

Figure 5. Mean spectrum in the range 119–124 μm obtained by averaging all the individual FU Ori spectra. The superposed formal ﬁt provides a detection at S/N > 8, with a width equal to the instrumental one, and centered at 121.75 μm. PRE-MAIN SEQUENCE STARS SEEN BY ISO 197

A possible interpretation of the observed FIR spectra is based on the simultaneous presence of two components: (i) a well localized J-shock, responsible for the [O I] emission, and (ii) an extended low density ionized medium, produced by 4 the UV photons from the disk boundary layer (Teff ∼ 3 × 10 K), responsible for the [N II] and [C II] emission. Infact, in FU Ori systems, the central star (F–G type) cannot provide the requisite ionising photons (Go ≈ 0.1 at 0.1 pc).

4. Future Perspectives

The FIR spectroscopical surveys presented here have clariﬁed many aspects of the interaction between pre-main sequence objects and their close environment, and have originated new questions which can be hopefully answered by using the forthcoming facilities. Some relevant aspects, addressed in this review, are summarized in the following:

– ISO instrumentation has evidenced for the first time the multiplicity of phenomena (outflows, shocks, disk emission, fluorescence) occurring in complex systems or in young binary objects (such as T Tau), but, due to the large beam size, one is often unable to specify which is the origin of the different manifestations. – SurveysofH2 pure rotational lines (in the 7–28 μm range) revealed themselves as crucial to probe small amounts of warm gas (maybe disks/planets) around young embedded objects, but enough spatial resolution (better than 1 arcsec) is required to impose compelling constrains. – The ISO-LWS mapping capability is enough to derive (from [O I] and [C II] line ratios) the gross distribution of Go and density. A better spatial resolution of ∼10 arcsec is in order to investigate the internal structure of the detected PDR’s by tracing the transition zones C+/C/CO, thus clarifying whether or not clumpiness is an important ingredient. – The observed ratio [OI]63μm/[OI] 145 μm is always lesser than predicted from both standard PDR and J-C-shocks models. To find a satisfactory agreement with the observed data, we pushed the PDR models toward the lowest [OI] 63/145 values. Studying whether this large scale behaviour stems from properties intrinsic to the emitting gas or results from averaging different emission vs. absorption components, deserves a spatial resolution much better (by ∼ an order of magnitude) than that of ISO-LWS. – Molecular emission in HAEBE has been detected on a few sources, coming from very compact (≤10 arcsec) regions. A systematic study carried out with a better sensitivity and spatial resolution could both give a definitive answer to the size of the emitting regions and significantly increase the detection rate of molecular emission (which is, at the moment, quite low). – ISO missed the fine structure transition of [N II]at205μm. The diagnostic diagram which uses the [N II] line pair (122/205 μm, see e.g. Rubin et al., 1994), 198 D. LORENZETTI

would provide the electron density in the N+ volume and the fractional ionization in the environment around FU Ori objects. This will allow to verify the hypothesis about the [N II] excitation mechanism in that class. Moreover, at higher spectral resolution, the [N II] 205 μm line (expected to be brighter than the 122 μm one) can be resolved, allowing a better constrain on the origin of this line.

References

Acke, B. and van den Ancker, M.: in press, Astron. Astrophys. 426, 151. Abraham, P., Leinert, C., Burkert, A., Henning, T., and Lemke, D.: 2000, Astron. Astrophys. 354, 965. Benedettini, M., Nisini, B., Giannini, T., et al.: 1998, Astron. Astrophys. 339, 159. Benedettini, M., Pezzuto, S., Giannini, T., Lorenzetti, D., and Nisini, B.: 2001, Astron. Astrophys. 379, 557. Bouwman, J., de Koeter, A., Dominik, C., and Waters, L. B. F. M.: 2003, Astron. Astrophys. 401, 577. Bouwman, J., de Koeter, A., van den Ancker, M. E., and Waters, L. B. F. M.: 2000, Astron. Astrophys. 360, 213. Bouwman, J., Meeus, G., de Koeter, A., Hony, S., Dominik, C., and Waters, L. B. F. M.: 2001, Astron. Astrophys. 375, 950. Burton, M. G., Hollenbach, D., and Tielens, A. G. G.: 1990, Astrophys. J. 365, 620. Chiang, E. I. and Goldreich, P.: 1997, Astrophys. J. 490, 368. Chiang, E. I., Joung, M. K., Creech-Eakman, M. J., Qi, C., and Kessler, J. E.: 2001, Astrophys. J. 547, 1077. Creech-Eakman, M. J., Chiang, E. I., Joung, R. M. K., Blake, G. A., and van Dishoeck, E. F.: 2002, Astron. Astrophys. 385, 546. Dominik, C., Dullemond, C. P., Waters, L. B. F. M., and Walch, S.: 2003, Astron. Astrophys. 398, 607. Draine, B. T., Roberge, W. G., and Dalgarno, A.: 1983, Astrophys. J. 264, 485. Dullemond, C. P., Dominik, C., and Natta, A.: 2001, Astrophys. J. 560, 957. Elia, D., Strafella, F., Campeggio, L., Giannini, T., Lorenzetti, D., and Nisini, B.: 2004, Astrophys. J. 601, 1000. Fuente, A., Marti´ın-Pintado,J., Bachiller, R., Neri, R., and Palla, F.: 1998, Astron. Astrophys. 334, 253. Fuente, A., Marti´ın-Pintado,J., Rodr´ıguez-Fernández,N. J., Chernicharo, J., and Gerin, M.: 2000, Astron. Astrophys. 354, 1053. Fuente, A., Marti´ın-Pintado,J., Rodr´ıguez-Fernández,N. J., Rodr´ıguez-Franco,A., de Vicente, P., and Kunze, D.: 1999, Astrophys. J. 518, L45. Giannini, T., Lorenzetti, D., Tommasi, E., et al.: 1999, Astron. Astrophys. 346, 617. Gürtler, J., Schreyer, K., Henning, T., Lemke, D., and Pfau, W.: 1999, Astron. Astrophys. 346, 205. Hartmann, L. W. and Kenyon, S. J.: 1985, Astrophys. J. 299, 462. Hartmann, L. W. and Kenyon, S. J.: 1996, Ann. Rev. Astron. Astrophys. 34, 207. Hartmann, L. W., Kenyon, S. J., and Hartigan, P.: 1993, in Levy, E. H. and Lunine, J. I. (eds.), Protostars and Planets, Vol. III, University of Arizona Press, Tucson, p. 497. Henning, T., Burkert, A., Launhardt, R., Leinert, C., and Stecklum, B.: 1998, Astron. Astrophys. 291, 546. Hollenbach, D. and McKee, C. F.: 1989, Astrophys. J. 342, 306. Hutseméker, D. and van Drom, E.: 1990, Astron. Astrophys. 238, 134. Kaufman, M. J., Wolfire, M. G., Hollenbach, D. J., and Luhman, M. L.: 1999, Astrophys. J. 527, 795. Léger, A. and Puget, J. L.: 1984, Astron. Astrophys. 137, L5. Lorenzetti, D., Giannini, T., Nisini, B., et al.: 2000, Astron. Astrophys. 357, 1035. PRE-MAIN SEQUENCE STARS SEEN BY ISO 199

Lorenzetti, D., Giannini, T., Nisini, B., et al.: 2002, Astron. Astrophys. 395, 637. Lorenzetti, D., Tommasi, E., Giannini, T., et al.: 1999, Astron. Astrophys. 346, 604. Malfait, K., Waelkens, C., Bouwman, J., de Koter, A., and Waters, L. B. F. M.: 1999, Astron. Astrophys. 345, 181. Malfait, K., Waelkens, C., Waters, L. B. F. M., Vandenbussche, B., Huygen, E., and de Graauw, M. S.: 1998, Astron. Astrophys. 332, L25. Meeus, G., Waters, L. B. F. M., Bouwman, J., van den Ancker, M. E., Waelkens, C., and Malfait, K.: 2001, Astron. Astrophys. 365, 476. Natta, A., Meyer, M. R., and Beckwith, S. V. W.: 2000, Astrophys. J. 534, 838. Natta, A., Pusti, T., Neri, R., Wooden, D., Grinin, V. P., and Mannings, V.: 2001, Astron. Astrophys. 371, 186. Nisini, B., Giannini, T., and Lorenzetti, D.: 2002, Astrophys. J. 574, 246. Nisini, B., Lorenzetti, D., Cohen, M., et al.: 1996, Astron. Astrophys. 315, L321. Peeters, E., Hony, S., van Kerckhoven, C., et al.: 2002, Astron. Astrophys. 390, 1089. Pezzuto, S., Grillo, F., Benedettini, M., et al.: 2002, MNRAS 330, 1034. Pezzuto, S., Strafella, F., and Lorenzetti, D.: 1997, Astrophys. J. 485, 290. Richter, M. J., Jaffe, D. T., Blake, G. A., and Lacy, J. H.: 2002, Astrophys. J. 572, L161. Rubin, R. H., Simpson, J. P., and Lord, S. D.: 1994, Astrophys. J. 420, 772. Sandell, G. and Weintraub, D. A.: 1994, Astron. Astrophys. 292, L1. Siebenmorgen, R., Natta, A., Kruegel, E., and Prusti, T.: 1998, Astron. Astrophys. 339, 164. Siebenmorgen, R., Prusti, T., Natta, A., and Müller, T. G.: 2000, Astrophys. J. 561, 1074. Spinoglio, L., Giannini, T., Nisini, B., et al.: 2000, Astron. Astrophys. 353, 1055. Thé.P.S., de Winter, D., Pérez,M.R.: 1994, Astrophys. Supp. 104, 315. Thi, W.F.,van Dishoeck, E. F.,Blake, G. A., van Zadelhoff, G.-J., Hogerheijde, M. R.: 1999, Astrophys. J. 521, L63. Thi, W. F., van Dishoeck, E. F., Blake, G. A., et al.: 2001, Astron. Astrophys. 385, 546. Tielens, A. G. G. M. and Hollenbach, D.: 1985, Astrophys. J. 291, 722. van Boekel, R., Waters, L. B. F. M., Dominik, C., et al.: 2003, Astron. Astrophys. 400, L21. van den Ancker, M. E., Bouwman, J., Wesselius, P. R., Waters, L. B. F. M., Dougherty, S. M., and van Dishoeck, E. F.: 2000a, Astron. Astrophys. 357, 325. van den Ancker, M. E., Wesselius, P. R., and Tielens, A. G. G. M.: 2000b, Astron. Astrophys. 355, 194. van den Ancker, M. E., Wesselius, P. R., Tielens, A. G. G. M., van Dishoeck, E. F., and Spinoglio, L.: 1999, Astron. Astrophys. 348, 877. van den Kerckhoven, C., Tielens, A. G. G. M., and Waelkens, C.: 2002, Astron. Astrophys. 384, 568. van Dishoeck, E. F.: 2004, Ann. Rev. Astron. Astrophys. 42, 119. Wesselius, P. R., van den Ancker, M. E., Young, E. T., et al.: 1996, Astron. Astrophys. 315, L197. DEBRIS DISCS AROUND STARS: THE 2004 ISO LEGACY

MARIE JOURDAIN DE MUIZON Leiden Observatory, PO Box 9513, 2300 RA Leiden, The Netherlands; LESIA, Paris Observatory, 92190 Meudon, France (E-mail: [email protected])

(Received 3 August 2004; Accepted in ﬁnal form 2 November 2004)

Abstract. Debris discs around stars were ﬁrst discovered by the Infrared Astronomical Satellite (IRAS) in 1983. For the ﬁrst time material orbiting another star than the Sun, but distinct from a circumstellar envelope, was observed through its far infrared emission. This major discovery motivated astronomers to investigate those discs by further analyzing the IRAS data, using ground-based telescopes for the hunting of exoplanets, developing several projects using the Infrared Space Observatory (ISO), and now exploiting the ISO Data Archive (IDA). This review presents the main ISO results, statistical as well as individual, on debris discs in orbit around pre-main-sequence and main-sequence stars.

Keywords: stars, debris discs, debris disks, planetary systems, infrared astronomy, Vega, β Pic, HD 207129, α PsA, ρ1 CnC

1. Introduction

Discs around stars appear during the early stages of stellar evolution. About 4.6 Gyr ago, the Sun – like any other star – formed in a local contraction of an interstellar cloud of molecular gas and dust. During its ﬁrst few million years the Sun was surrounded by a warm rotating disc of gas and dust that on one hand fed material onto the forming star and on the other hand led to the formation of comets, planets and planetesimals, asteroids and other objects. This kind of warm disc is also found around other pre-main-sequence stars, such as T Tauri stars, Herbig Ae/Be stars and ZAMS stars (Robberto et al., 1999). They contain some original interstellar dust, are optically thick and extend to a few AU. After a few 107 years, the warm inner part of the solar disc was dissipated while a cooler debris disc remained in the outer part of the newly formed solar system. Such a disc contains mainly interplanetary dust resulting from collisions, is optically thin and extends up to a few hundred AU. It is such a debris disc that was ﬁrst discovered by Infrared Astronomical Satellite (IRAS; Aumann et al., 1984) around Vega, α Lyr, one of the best calibrated and

Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, The Netherlands, and the United Kingdom), and with the participation of ISAS and NASA. Space Science Reviews (2005) 119: 201–214 DOI: 10.1007/s11214-005-8063-0 C Springer 2005 202 M. JOURDAIN DE MUIZON extensively used photometric standard in the visual range, then found to radiate more mid- and far-infrared emission than its photosphere can produce, and also around α PsA, ε Eri and β Pic. A systematic search for Vega-like discs in the IRAS survey led to several other identiﬁcations but was not very conclusive due to the limited sensitivity (Backman and Paresce, 1993; Plets and Vynckier, 1999). As follow-up of IRAS, several surveys were undertaken with Infrared Space Obser- vatory (ISO; Kessler et al., 1996). The most unbiased was obtained by Habing et al. (1999, 2001) but others concentrated on stars in open clusters (Spangler et al., 2001), G dwarfs (Decin et al., 2000), or selected MS stars (Fajardo-Acosta et al., 1999). Several case studies on the most ISO observed discs are also presented: Vega and β Pic (Heinrichsen et al., 1998, 1999), ρ1 Cnc (Dominik et al., 1998), HD 207129 (Jourdain de Muizon et al., 1999), ﬁve Vega-like stars including α PsA (Fajardo-Acosta et al., 1997). Finally ISO spectroscopic observations evidenced the presence of amorphous and crystalline silicates, forsterite, PAHs (Meeus et al.; 2001, Bouwman et al., 2001) and molecular hydrogen (Thi et al., 1999, 2001a,b) in the discs of several pre-main-sequence stars. See also two more general reviews on the subject by Lagrange et al. (2000) and Zuckerman (2001).

2. ISO Surveys

2.1. GENERAL STATISTICS

Habing et al. (1999, 2001) proposed to determine the incidence of Vega-like debris discs in a distance limited sample of main-sequence stars. The stars were carefully selected from Johnson and Wright (1983) who computed far-infrared ﬂuxes for 93% of the 2,150 stars in the “Woolley catalog of stars within 25 pc from the Sun.” After rejecting all stars either too faint for ISO at 60 μm or for which an infrared excess would be ambiguous to interpret as a disc, their ﬁnal selection consisted of 84 main-sequence stars of spectral types from A to K. Within this range, no spectral type was privileged. The stars were measured with ISO at 25, 60, 90, and 170 μm in a total observing time of 65 h. Based on ISOPHOT (Lemke et al., 1996) C100 3 × 3 minimaps at 60 μm, Habing et al. (1999, 2001) found that 17% of the stars in their sample do have a disc. Lachaume et al. (1999) determined the age of the 84 stars. It appeared that all stars younger than 300 Myr have a disc, 70% of the stars younger than 400 Myr still have a disc but this is the case for only 8% of the stars older than 1 Gyr. Thus, it seems that most stars arrive on the main-sequence surrounded by a disc, and for the majority of them the disc is dissipated within a few hundred Myr (Figure 1). The decay of the discs is attributed to the destruction and escape of planetesimals; indeed the collision of planetesimals is a good source of dust, necessary to replenish the disc because dust particles disappear on a much shorter timescale relative to the lifetime of the discs. The timescale of the dissipation of the disc corresponds to the DEBRIS DISCS AROUND STARS: THE 2004 ISO LEGACY 203

Figure 1. Cumulative distribution of stars with an infrared excess, as a function of the index, sorted by age. X-axis is rank at the bottom and age at the top of the figure. Y-axis is the number of stars up to that rank which have a disc. Each dot on the figure represents a star of the sample. For each star with an infrared excess, Nex increments by one. The two segments of a continuous straight line correspond to the perfect situation in which most debris discs are no longer detectable after 400 Myr (Habing et al., 2001). end of the heavy bombardment phase in the Solar System, which is identified by dating the cratering on the Moon, Mercury, Ganymede, and Callisto. Craters are due to impacts from planetesimals and this bombardment appears to have stopped about 4 Gyr ago (Shoemaker and Shoemaker, 1999; Morbidelli et al., 2001). Thus, the timescale of 400 Myr would not trace, properly speaking, the disappearance of the dust discs, but the termination of the production of dust by collisions of planetesimals, hence the drastic reduction in the occurrence of these larger bodies (Jourdain de Muizon et al., 2001). Using 25 μm ISOPHOT data on 81 stars from the Habing et al. (2001) sample, Laureijs et al. (2002) found that 5 stars (6%) have an infrared excess that can be attributed to a disc of dust temperature between 50 and 120 K. The 5 stars are younger than 400 Myr, thus indicating that warm debris discs are relatively rare and concern the younger stars only. The survey confirms that there seems to be an absence of detectable amounts of dust close to the stars (D ≤ 20 AU). From a survey of 38 main-sequence stars using IRAS and ISOPHOT data, Fajardo-Acosta et al. (1999) found no star with a significant excess at 12 μm, and a fraction of ∼14% excess stars at 20 μm. However, this fraction is difficult to interpret since the ISOPHOT data in their study were inconclusive and the detections needed confirmation. The absence of 12 μm excess indicates that the discs are not warmer than 200 K. Finally, Jourdain de Muizon and Laureijs (2005) found three more excess stars at 60 μm and one more at 90 μm by exploring the ISOPHOT chopped data, which had been obtained at an early stage in the ISO mission but were later replaced by minimaps (see Habing et al., 2001). 204 M. JOURDAIN DE MUIZON

2.2. G STARS

Decin et al. (2000) studied the incidence of the Vega phenomenon around G dwarfs. The stars were selected from the CORALIE sample (planet-search programme around stars closer than 50 pc and of spectral type from F8 to M1, excluding giants and faint cool dwarfs). Only southern stars were considered and the main selection criteria were observability and detectability by ISO at 60 μm. Confused systems such as multiple stars and those against a high cirrus background were rejected. Finally, this ISO programme consisted of 69 stars, 34 of which were effectively observed and results are given for 30 of them (ISOPHOT minimaps at 60 μm). Of the 30 G dwarfs in Decin et al. (2000), five (17%) have a debris disc and four out of these five are older than 3 Gyr. This is in good agreement with Habing et al. (2001) who have 21 G dwarfs in their sample and find that four (19%) of them have a debris disc; three out of these four are older than 5 Gyr (they are part of their 8% stars older than 1 Gyr and which still have their disc, see Section 2.1). However, two of the excess stars in Decin et al. (2000) have fractional luminosities of their disc comparable to the disc of β Pic which was a unique case so far, whereas the others and all the discs around G stars in Habing et al. (2001) are between one and two orders of magnitude fainter. Also no correlation was found between the existence of a disc and a planet around the stars of the Decin et al. (2000) sample. Thus, for both groups of authors, about 18% of the G stars do have a disc and about 80% of these discs are around G stars older than 3 Gyr. Why do G stars appear to keep their discs longer than A, F, or K stars is not yet understood. Around the Sun, the zodiacal light and the Kuiper Belt are probably some remnants of the earlier disc.

2.3. STARS IN OPEN CLUSTERS

Spangler et al. (2001) used ISOPHOT to observe a total of 148 stars, of which 87 young (50–700 Myr), nearby (d < 120 pc) main-sequence stars in open clusters (α Persei, Coma Berenice, Hyades, Pleiades, UMa) of spectral type A to K, 41 T Tauri stars in the clouds Chamaeleon I, Scorpius and Taurus (d ≈ 140–150 pc), and a sample of 19 isolated young nearby field stars (d < 60 pc) and another field star. They obtained ISOPHOT chopped observations with the C100 detector at 60 and 100 μm or raster at 60 and 90 μm. The goal was to determine an evolutionary sequence for circumstellar disc characteristics, hence their choice of well-studied clusters in which stellar ages are fairly well defined. Although spectral types span the range A to K the authors gave a preference to solar-type stars and privileged spectral types F and G in their star selection. They detected 36 stars, of which 33 show evidence for a far-infrared excess, i.e., 22% of their hole sample. More than one third were already thought to have an IR excess in their IRAS data, the rest are new ISO detections. These latter consist of 13 cluster stars, 5 young field stars and 1 other field star, with an ISO excess emission at 60, 90 or 100 μm. Among the DEBRIS DISCS AROUND STARS: THE 2004 ISO LEGACY 205 main-sequence excess stars, spectral type F clearly dominates. Given the sample size of the statistics, the 22% excess stars in Spangler et al. (2001) is not too far from the 17% in Habing et al. (2001) or in Decin et al. (2000). The difference is most likely due to the selection criteria, the sample of Habing et al. (2001) being the most unbiased and homogeneous. Discs around T Tauri stars are closer to accretion discs (optically thick) and are not properly called debris discs (optically thin) like those around main-sequence stars. We only mention them in this ISO review because they are precursors of the debris discs but they cannot be treated equally when defining the incidence of debris discs. A convenient parameter to describe systems exhibiting excess emission from circumstellar discs is the fractional excess luminosity, fd = Lex/L , where Lex is the luminosity of dust and L is the stellar bolometric luminosity. Although Lex is not easy to estimate as it requires to know the complete infrared spectrum of the excess, it can be done approximately using both the IRAS and ISO data. A very interesting result of Spangler et al. (2001) is how the IR excess evolves with time. They established that the fractional excess luminosity fd decreases with stellar age −1.76 according to the power law fd ∝ (age) . This is compatible with a collisionally replenished disc as suggested in Habing et al. (2001). Spangler et al. (2001) claim they do not see evidence for an abrupt cessation of the debris disc phenomenon as reported in Habing et al. (2001). It is indeed not obvious from their Figure 1, but in fact 27 out of their 33 excess stars (82%) are younger than 400 Myr, and the other 6 are between 400 and 625 Myr. This is in perfect agreement with Habing et al. (2001). It is clear that debris discs are mostly found around young stars. Spangler et al.’s sample is biased towards young stars anyway. All excess stars in their sample are younger than the end of the late heavy bombardment on the Moon. Robberto et al. (1999) present preliminary results on 97 very young stars in 5 open clusters, 3 of which are in common with those in Spangler et al. (2001), namely Chamaeleon, α Per and Pleiades, and they are all younger than 300 Myr. Using ISOPHOT with the C100 detector at 60 μm and P2 detector at 25 μm they detected only 4 stars, i.e., 4.1% of their hole sample. The four stars detected are all T Tauri stars in the Chamaeleon I cluster; three of them are classical T Tauri stars which are still in the accretion phase (there are only six such stars known in this cluster) and the fourth one is probably in the transition phase. Their results show that the transition from an optically thick disc (or accretion disc) to an optically thin one (or debris disc) occurs on a timescale of ∼10 Myr, with a transition phase lasting less than ∼0.3 Myr.

2.4. RE-ANALYSIS OF THE ABOVE SURVEYS

Greaves and Wyatt (2003) have addressed the issue of the incidence of debris discs among A- and G-type main-sequence stars from the various published ISO surveys and some IRAS discs (about 200 stars). Their analysis shows a much higher detection rate of debris disks towards A stars (9/20 hence 45%) than G stars (11/180 206 M. JOURDAIN DE MUIZON hence 6%), and this big difference in disc detection rate applies even when stars of similar age are considered. From the observations, disc lifetime is estimated to about 0.5 Gyr and may occur at any time during the main-sequence life of the star. Disc lifetime and high occurrence of discs among A stars were two conclusions already reached by Habing et al. (2001), but they are now also verified on a much larger sample. The fact that A stars have a much shorter lifetime and a larger disc mass than G-stars, implies the disc lifetime is a much higher fraction of the star lifetime in A stars than in G-stars and thus it is not surprising that A stars discs are more often seen than in G stars and also often seen in the first half of the A star lifetime (i.e., younger than 500 Myr). The disc mass scans a range of about 100, decline with time is slow and follows a power law not steeper than t−1/2. The conclusion that the debris disc lifetime can take place at any time during the main- sequence, and not necessarily at the beginning of the stellar life, would account for the existence of a few detected debris disks around some old G stars. According to the models (Kenyon and Bromley, 2002a,b), perturbing planetesimals can form slowly at large orbital radii as late as a few Gyr, thus producing dust by collisions and replenishing a disc, but the same models cannot really explain the more or less constant disc duration of about 0.5 Gyr. Also based on the various ISO surveys, the age dependence of the Vega phenomenon has been studied observationally and theoretically, respectively by Decin et al. (2003) and Dominik and Decin (2003). Decin et al. (2003) have critically re- examined the stellar age estimates from the existing ISO surveys. Two aspects are to be considered: (i) the time dependence of the disc dust mass, measured by fd the fractional (i.e., disc/star) luminosity and (ii) the incidence of debris discs versus stellar age. Decin et al. (2003) came to the conclusion that there is no clear trend in the time dependence of fd, in particular no power-law with an index of ≈−2 (e.g., Spangler et al., 2001), but rather a large spread irrespective of stellar age. However, −3 for a given stellar age, the maximum fd is about 10 (upper cut-off). They also identify a few cases of very young stars with intermediate or small excess (i.e., −5 below the lower cut-off for fd which is about 10 ), and they conclude that Vega- like stars are more common in young stars that in old stars, a conclusion which had already been reached by Habing et al. (1999). The results of Decin et al. (2003) are in agreement with Habing et al. (1999, 2001) who had the most unbiased sample of the ISO Vega-like studies (see Section 2.1). They differ significantly from those of Spangler et al. (2001), but that could be explained because the latter studied mainly stars in clusters, thus of similar age and evolution pattern, hence the rather steep power law they could derive for fd versus time. Dominik and Decin (2003) have tried to provide a coherent theoretical picture of the various ISO observations of debris discs. On one hand, without replenishment, the dust in a debris disc should disappear in few thousands years under the Poynting- Robertson effect, i.e., a drag on interplanetary particles caused by their interaction with solar (respectively stellar) radiation, which causes the particles to lose orbital DEBRIS DISCS AROUND STARS: THE 2004 ISO LEGACY 207 momentum and spiral into the sun (respectively star). On the other hand the discs last much longer than a few thousands years and some of them are seen in old stars, thus the dust must be replenished. The model must account for both the existence of the disc at a given stage of stellar evolution and also the disc evolution (from its formation to its dissipation). The former needs an accurate knowledge of stellar ages, the latter an accurate estimate of fd. According to Dominik and Decin (2003) the ISO observations can be explained by a simple model of collisional cascades, completed by some stirring in the case of high mass discs (cf. Kenyon and Bromley, 2002b). In a collisional cascade scenario where collisions occur with constant velocities, the decrease of the amount of dust can be described by a power law of index − 1 and this applies to all observed debris discs. In the case of very −3 low mass discs (about 10 M⊕, practically unobserved so far), the slope of this power law can be of the order of −2. In a scenario where the collisional cascade is continuously stirred (because the collision velocities increase with time), the power law slope can be steeper than −1.

3. Case Studies

A few stars were observed rather extensively and with several ISO observing modes or instruments because they were particularly interesting cases. The best examples are α Lyr (Vega), β Pic, α PsA (Fomalhaut), ρ1 CnC and HD 207129.

3.1. VEGA AND β PIC

The most detailed ISO studies of the discs around these two prototypes of debris discs are found in Heinrichsen et al. (1998, 1999). They are summarized in Table I. Based on the 25 μm ISO data, the authors argue that the disc around β Pic is in fact much more massive than the cool dust derived from the ISO 60 μm emission. They suggest that there is some warm dust (300 to 500 K), extension of the inner disc seen in the optical, and a signiﬁcant amount of cool dust in addition to that seen by ISO. The star could be surrounded by a large “Oort” cloud of comets. Walker and Heinrichsen (2000) present ISOPHOT-S spectra (from 6 to 12 μm) of 12 Vega-like stars including the four “prototypes” (Vega, β Pic, Fomalhaut and ε Eri) in search for silicates and PAHs features in debris discs. They found silicate dust emission towards two stars (HD 144432 and HD 139614), emission from carbon-rich molecules towards two others (HD 169142 and HD 34700) and emission from both towards HD 142666. For all their other stars, including the four “prototypes,” either only the photosphere is seen at these mid-infrared wavelengths, or some thermal featureless excess. The authors also present preliminary ISOPHOT maps at 60 and/or 90 μm, thus giving some insight in the extent and structure of the discs. The authors argue that the discs could be in the early stage of planet formation. However, they give no accurate information on stellar age, hence 208 M. JOURDAIN DE MUIZON

TABLE I ISOPHOT results on Vega and β Pic

ISOPHOT Vega β Pic

High-resolution scans 60 μm (P32/C100) 25, 60 μm Disc resolved Yes: face-on Yes: edge-on Distance 7.8 pc 19.3 pc Disc radius 86 AU at 60 μm 84AUat25μm 140 AU at 90 μm 140 AU at 60 μm Multiﬁlter-photometry 25, 60, 80, 100, 120, 150, 170 4.85, 7.3, 11.3, 12.8, 16, 25, 60, and 200 μm 80, 120, 150 and 170 μm Adopted dust emissivity Q(λ) ∝ 1/λ1.1 Q(λ) ∝ 1/λ − − Dust mass in disc (1–5) × 10 3 M⊕ (1.0–3.3) × 10 2 M⊕ −4 a −2 a (1.3–13) × 10 M⊕ (1.2–12) × 10 M⊕ Reference Heinrichsen et al. (1998) Heinrichsen et al. (1999) Proposal: Walker Proposal: Heinrichsen aHabing et al. (2001). the question: “Are the above features emitted by dust in a debris disc or in a much younger and thicker envelope or disc?”

3.2. α PSA

Fajardo-Acosta et al. (1997) have obtained P32 maps with ISOPHOT-C100 at 60 μm, with a spatial resolution of about 30 for six Vega-type systems. At 60 μm there is no excess emission detected towards α CrB, σ Her or α Cen. There is a marginal detection of extended emission, out to ∼800 AU or 30 from the star γ Oph. Only α PsA in their sample shows a convincing excess emission in the range ∼30 to 80, i.e., ∼210 to 560 AU, conﬁrming the IRAS temperature of 58–75 K and suggesting grains up to ∼10 μm in size. They estimate a dust mass −3 of ∼(2–6) × 10 M⊕. A comparison with β Pic shows that both discs are similarly extended at 60 μm (400 AU for β Pic vs. 560 AU for α PsA).

3.3. ρ1 CNC

The star ρ1 CnC is among the ﬁrst 10 stars found to host a planet. In 1996, following the news of a planet around ρ1 CnC, a special set of ISO observations was requested and carried out to search for a disc. The star ρ1 CnC was observed with ISOPHOT at 25 μm (PHT03), and at 60, 90, 135 and 170 μm (PHT22). Dominik et al. (1998) found an excess of 170 ± 30 mJy at 60 μm. They interpreted it as a debris disc −5 located at about 60 AU from the star and containing at least 4 × 10 M⊕ of dusty material. The star is a G8V, slightly metal-rich, located at a distance of 12.5 pc DEBRIS DISCS AROUND STARS: THE 2004 ISO LEGACY 209

Figure 2. The infrared excess toward HD 207129. Comparison between two different dust compositions. The dashed line indicates a Draine and Lee (1984) interstellar grain model, the solid line aLi and Greenberg (1997) cometary dust model. The two cases differ most around 30–40 μm, but they both fit the remaining far-IR excess equally well (Jourdain de Muizon et al., 1999). from the Sun. From CaII H and K lines, its age was estimated to be 5 Gyr. Butler et al. (1997) detected the presence of a planet of period 14,65 days, implying a semi-major axis of 0.11 AU. The inferred mass is M2 sin i = 0.84MJup. The far- infrared ISO excess was interpreted as a disc after any other possible origin had been examined and eliminated (companion M5, planet, cirrus knot in the background). The far-infrared spectrum of the excess is best fitted using cometary icy dust grains from the dust model by Li and Greenberg (1997) as seen in Figure 2 of Dominik et al. (1998). These observations brought the first evidence of the coexistence of a disc and a planet around a star other than the Sun.

3.4. HD 207129

This solar-type (G2V) star is particularly interesting because of its cold debris disc, may be one of the coldest observed so far. It caught the attention of Jourdain de Muizon et al. (1999) when they found a clear excess not only at 60 μm, but also at 170 μm in their ISOPHOT-C200 data. Additional discretionary ISO observations were requested and obtained to get a more complete infrared spectrum of the disc. Based on the independent measurements of the CaII K line made by two groups of authors (Pasquini, 1992; Henry et al., 1996), Lachaume et al. (1999) determined a stellar age of 4.7 Gyr similar to that of the Sun, adopted by Jourdain de Muizon et al. (1999) and Habing et al. (2001). Zuckermann and Webb (2000) argue that the star is a member of the Tucanae stream and must be as young as the Tucanae association (10–50 Myr) which is located at ∼45 pc. That is three times further than HD 207129 which is located at 15.6 pc from the Sun, according to the Hipparcos Catalogue (Perryman et al., 1997). This signiﬁcantly weakens the possibility of its 210 M. JOURDAIN DE MUIZON belonging to the Tucanae stars. In any case the age of the star is less of an issue if it is young because it would then strengthen the case of Habing et al. (1999, 2001) that debris disc is the privilege of young stars (less than 400 Myr). HD 207129 is indeed one of the few puzzling exceptions in their study. The infrared excess around HD 207129 is shown in Figure 2. The star emits approximately 1.1×10−4 of its luminosity longwards of 25 μm. The excess emission −2 is explained by assuming a disc of dust-like material equivalent to 2 × 10 M⊕. The dust temperature ranges from 10 to 50 K, which makes it the coldest debris disc around a Vega-like star known to date. The dust distribution in the disc presents a circular hole at a distance of ∼400 AU from the star. This hole should be ﬁlled within 10−3 of the stellar age by particles spiraling inward because of the Poynting- Robertson effect unless an agent sweeps it clean. It could thus be explained by the presence of one or more planets.

4. Dust and Gas Composition of the Discs

Debris discs are essentially made of dust particles and larger bodies; however, these latter cannot be detected by ISO. In addition, the younger discs may still have a signiﬁcant gas component. Before ISO, any attempt to trace the gas via CO molecular bands was unsuccessful. Direct measurement of H2 with ISO–SWS (de Graauw et al., 1996) has made a real break-through toward understanding the gas composition of very young discs.

4.1. DUST COMPOSITION

Amorphous silicates (r < 1 μm) were detected in the ISOPHOT-S spectra of nine classical T Tauri stars in the Chamaleon I dark cloud (Natta et al., 2000). They discuss a model that explains the origin of this material in a hot, optically thin surface layer of a disc around the star, i.e., in the disc atmosphere. Crystalline silicates were detected by van den Ancker et al. (2001) toward the pre-main-sequence (B9.5Ve, A0II–IIIe or AOV?) star 51 Oph, using SWS01 and LWS01 full scans. The solid-state bands and energy distribution indicate that the dust has formed recently; it is a very young disc, not a debris disc. Other emission bands from hot gas (∼350 K) such as CO, CO2,H2O and NO dominate the 4– 8 μm spectrum. Both these gas and dust bands are unusual for a young star and are more typical of evolved AGB stars, although 51 Oph does not seem at all to belong to that class. The authors explore various possibilities for the nature of 51 Oph, among them a recent episode of mass loss from a Be star, the collision of two gas-rich planets and the accretion of a solid body as the star expands at the end of its short main-sequence life. A variety of crystalline silicates, forsterite (Mg2SiO4), has been found toward the star HD 100546. Malfait et al. (1998) present SWS01 and LWS01 full scans of this isolated Herbig Ae/Be star. Forsterite is present in the micrometeorites DEBRIS DISCS AROUND STARS: THE 2004 ISO LEGACY 211 and interplanetary dust of our Solar System. The ISO spectrum of HD 100546 is very similar to that of comet Hale–Bopp published by Crovisier et al. (1997). The amount of forsterite in the disc of HD 100546 is equivalent to 1013 comets Hale– Bopp, strengthening the hypothesis by Grady et al. (1997) that the disc around HD 100546 contains a huge swarm of comets similar to the Oort cloud in the Solar System. Crystalline silicates have also been found in several other comets such as P/Halley. Malfait et al. (1998) argue that the crystallisation process occurs during the early phases of disc evolution around young stars. The ISO-SWS and -LWS spectra of an additional 14 isolated Herbig Ae/Be stars are presented in Meeus et al. (2001). These stars are believed to be the more massive analogs of T Tauri stars; they are seen as the progenitors of Vega-like stars (Waters and Waelkens, 1998). Meeus et al. (2001) obtained a variety of mid-infrared spectra ranging from the amorphous silicates (as in the M supergiant μ Cep) to the crystalline silicates (as in comet Hale–Bopp). The variations in the shape of the 8–14 μm part of the spectra indicate the prominence of one or the other form of silicates. Four out of the fourteen stars have no silicates. In most silicate stars in their sample, crystalline silicates are present; this is confirmed by the shape of the 15–28 μm ISO–SWS spectra of the stars. PAH bands are also identified toward half of the stars, and all features are superposed on a near-IR to mid-IR continuum excess. The authors interpret the continuum in term of disc geometry: an optically thick, geometrically thin, power-law component and an optically thin flare region (black body component). Bouwman et al. (2001) got a further detailed insight into the 6–14 μm spectrum of these stars in order to study the silicate grain processing. The 10 μm silicate profile is modelled using three components: silica (SiO2) responsible for the 8–9 μm blue shoulder in the silicate band, forsterite contributing to the 11.3 μm feature, and amorphous olivine with two typical grain sizes of 0.1 and 2.0 μm. They identify two causes for the observed shift in peak position of the silicate band: (i) a change in average grain size from small (0.1 μm) to big (2 μm), as a result of the depletion of small grains in the inner region of the disc, due to coagulation and other effects, (ii) a change in grain composition from amorphous silicate to a mixture of amorphous and Mg-rich crystalline silicate (forsterite) which could be the result of thermal annealing in the inner regions of the disc. However crystallization occurs on a longer timescale than coagulation. There is no correlation between dust composition and disc geometry.

4.2. GAS COMPOSITION:H2

Before ISO, any attempt to detect a gas component in Vega-like discs had been unsuccessful or inconclusive (see e.g., Liseau, 1999). The classical method of tracing molecular hydrogen by observing CO, widely applied in the interstellar medium, did not seem to help. Thi et al. (1999) discovered molecular H2 in the ISO–SWS spectra of the T Tauri star GG Tau. The detection of two pure rotational lines at 17.035 and 28.218 μm provided a direct measure of the total amount of warm molecular 212 M. JOURDAIN DE MUIZON gas in the disc around the star. The same group (Thi et al., 2001a,b) found H2 lines in the SWS spectra of4TTauri stars (spectral type K to M), 7 Herbig Ae stars (sp. type A to F), and 3 main-sequence stars including β Pic (spectral type A to F). Their data suggest the presence of warm gas (T ≈ 100–200 K), which mass ranges −4 −3 from ∼10 M up to 8 × 10 M, around all the stars they observed. This mass corresponds to 1–10% of the total disc mass inferred from millimeter continuum observations, and a much higher fraction in the case of debris discs. Additional CO observations show that CO is not a good tracer of the gas in circumstellar discs. The amount of CO gas is likely strongly affected by photodissociation due to the stellar and interstellar ultraviolet radiation in the surface layers of the disc and freeze-out onto grain surface in the midplane. Direct measurement of H2 leads to a gas-to-dust ratio of ∼100 in the debris discs of β Pic and HD 135344, similar to the interstellar medium ratio. The bulk of the gas around pre-main-sequence stars is cool (T ≈ 20–80 K), while the warm gas (T ≈ 100–200 K) may constitute the major gaseous component of debris discs around main-sequence stars, thus providing a reservoir for the formation of Jovian planets.

5. Conclusion and Open Questions

The ISO Data Archive (IDA) has not yet been fully exploited as far as debris discs and their precursors are concerned. Several aspects of disc evolution are not well understood. Why and how do most discs disappear after a lifetime of about 400 Myr? Why do some stars keep a disc for much longer, up to several Gyr? Do discs coexist with planets? Is it systematic, and if not, what conditions lead to a disc-planets or disc-only system? Does it depend on the stellar spectral type? Which physical and chemical processes occur during the evolution of a disc? Although the ISO data has already provided partial answers to some of these questions, it is clear that our picture is still incomplete. Beyond the IDA, space projects such as Spitzer and the Herschel Space Observatory do and will follow up ISO observations of debris discs. Spitzer has a similar wavelength range as ISO but a higher sensitivity; it is able to probe deeper and, hence, perform similar kind of statistics as ISO but on a larger volume around the Sun, thus expected to detect photometrically hundreds of debris disks. Some preliminary Spitzer results on the few brightest disks also address the structure of the disks, dust grain composition, and search for substellar companions. Herschel will open a new window since it extends to the submillimeter range. This will allow to trace the cold component of the discs such as the dust located in the Kuiper Belt, the Oort cloud and still further from the star.

References

Aumann, H. H., Gillett, F. C., Beichman, C. A., de Jong, T., Houck, J. R., Low, F. J., et al.: 1984, ApJ 278, L23. DEBRIS DISCS AROUND STARS: THE 2004 ISO LEGACY 213

Backman, D. and Paresce, F.: 1993, in E. Levy and J. Lunine (eds.), Protostars and Planets III, University of Arizona press, Tucson, p. 1253. Bouwman, J., Meeus, G., de Koter, A., Hony, S., Dominik, C., and Waters, L. B. F. M.: 2001, A&A 375, 950. Butler, R. P., Marcy, G. W., Williams, E., Hause, H., and Shirts, P.: 1997, ApJ 474, L115. Crovisier, J., Leech, K., Bockelee-Morvan, D., Brooke, T. Y., Hanner, M. S., Altieri, B., et al.: 1997, Science 275, 1904. Decin, G., Dominik, C., Malfait, K., Mayor, M., and Waelkens, C.: 2000, A&A 357, 533. Decin, G., Dominik, C., Waters, L. B. F. M., and Waelkens, C.: 2003, A&A 598, 636. de Graauw, Th., Haser, L. N., Beintema, D. A., Roelfsema, P. R., van Agthover, H., Barl, L., et al.: 1996, A&A 315, L49. Dominik, C., Laureijs, R. J., Jourdain de Muizon, M., and Habing, H. J.: 1998, A&A 329, L53. Dominik, C. and Decin, G.: 2003, A&A 598, 626. Draine, B. T. and Lee, H. M.: 1984, ApJ 285, 89. Fajardo-Acosta, S. B., Stencel, R. E., and Backman, D. E.: 1997, ApJ 487, L151. Fajardo-Acosta, S. B., Stencel, R. E., Backman, D. E., and Thakur, N.: 1999, ApJ 520, 215. Grady, C. A., Sitko, M. L., Bjorkman, K. S., Perez, M. R., Lynch, D. K., Russell, R. W., et al.: 1997, ApJ 483, 449. Greaves, J. S. and Wyatt, M. C.: 2003, MNRAS 345, 1212. Habing, H. J., Dominik, C., Jourdain de Muizon, M., Kessler, M. F., Laureijs, R. J., Leech, K., et al.: 1999, Nature 401, 456. Habing, H. J., Dominik, C., Jourdain de Muizon, M., Laureijs, R. L., Kessler, M. F., Leech, K., et al.: 2001, A&A 365, 545. Heinrichsen, I., Walker, H. J., and Klaas, U.: 1998, MNRAS 293, L78. Heinrichsen, I., Walker, H. J., Klaas, U., Sylvester, R. J., and Lemke, D.: 1999, MNRAS 304, 589. Henry, T. J., Soderblom, D. R., Donahue, R. A., and Baliunas, S. L.: 1996, AJ 111, 439. Johnson, H. M. and Wright, C. D.: 1983, ApJS 53, 643. Jourdain de Muizon, M., Laureijs, R. J., Dominik, C., Habing, H. J., Metcalfe, L., Siebenmorgen, R., et al.: 1999, A&A 350, 875. Jourdain de Muizon, M., Laureijs, R. J., Habing, H. J., Leech, K., Kessler, M. F., Metcalfe, L., et al.: 2001, Earth Moon Planets 85–86, 201. Jourdain de Muizon, M., Laureijs, R. J., and Augereau, J. C.: 2005, A&A in preparation. Kenyon, S. J. and Bromley, B. C.: 2002a, AJ 123, 1757. Kenyon, S. J. and Bromley, B. C.: 2002b, ApJ 577, L35. Kessler, M. F., Steinz, J. A., Anderegg, M. E., Clavel, J., Drechsel, G., Estaria, P., et al.: 1996, A&A 315, L27. Lachaume, R., Dominik, C., Lanz, T., and Habing, H. J.: 1999, A&A 348, 897. Lagrange, A.-M., Backman, D., and Artymowicz, P.: 2000, in V. Mannings, A. P. Boss, and S. S. Russell (eds.), Protostars and Planets IV, Tucson: University of Arizona Press, p. 639. Laureijs, R. J., Jourdain de Muizon, M., Leech, K., Siebenmorgen, R., Dominik, C., Habing, H. J., et al.: 2002, A&A 387, 285. Lemke, D., Klaas, U., Abolins, J., Abrah´am,P.,´ et al.: 1996, A&A 315, L64. Li, A. and Greenberg, J. M.: 1997, A&A 323, 566. Liseau, R.: 1999, A&A 348, 133. Malfait, K., Waelkens, C., Waters, L. B. F. M., Vandenbussche, B., Huygen, E., and de Graauw, M. S.: 1998, A&A 332, L25. Meeus, G., Waters, L. B. F. M., Bouwman, J., van den Ancker, M. E., Waelkens, C., and Malfait, K.: 2001, A&A 365, 476. Morbidelli, A., Petit, J.-M., Gladman, B., and Chambers, J.: 2001, Meteorit. Planet. Sci. 36, 371. Natta, A., Meyer, M. R., Beckwith, S. V. W.: 2000, ApJ 534, 838. 214 M. JOURDAIN DE MUIZON

Pasquini, L.: 1992, A&A 266, 347. Perryman, M. A. C., Lindegren, L., Kovalevsky, J., Hoeg, E., Bastian, U., Bernacca, P. L., et al.: 1997, A&A 323, 49. Perryman, M. A. C., ESA.: 1997, The Hipparcos and Tycho Catalogues. Astrometric and Photometric Star Catalogues Derived from the ESA Hipparcos Space Astrometry Mission, ESA SP Series Vol No. 1200, ISBN: 9290923997 (set), ESA Publications Division, Noordwijk, Netherlands. Plets, H. and Vynckier, C.: 1999, A&A 343, 496. Robberto, M., Meyer, M. R., Natta, A., and Beckwith, S. V. W.: 1999, in Cox, P., and Kessler, M. F. (eds.), The Universe as seen by ISO, ESA SP-427, p. 195. Shoemaker, E. and Shoemaker, C.: 1999, in Beatty, J. K., Petersen, C. and Chaikin, A. (eds.), The New Solar System 4th edn, Sky Publishing Corporation and Cambridge University Press, p. 69. Spangler, C., Sargent, A. I., Silverstone, M. D., Becklin, E. E., and Zuckerman, B.: 2001, ApJ 555, 932. Thi, W. F., van Dishoeck, E. F., Blake, G. A., van Zadelhoff, G. J., and Hogerheijde, M. R.: 1999, ApJ 521, L63. Thi, W. F., Blake, G. A., van Dishoeck, E. F., van Zadelhoff, G. J., Horn, J. M. M., Becklin, E. E., et al.: 2001a, Nature 409, 60. Thi, W. F., van Dishoeck, E. F., Blake, G. A., van Zadelhoff, G. J., Horn, J. M. M., Becklin, E. E., et al.: 2001b, ApJ 561, 1074. van den Ancker, M. E., Meeus, G., Cami, J., Waters, L. B. F. M., and Waelkens, C.: 2001, A&A 369, L17. Walker, H. J. and Heinrichsen, I.: 2000, Icarus 143, 147. Waters, L. B. F. M. and Waelkens, C.: 1998, ARA&A 36, 233. Zuckerman, B. and Webb, R. A.: 2000, ApJ 535, 959. Zuckerman, B.: 2001, ARA&A 39, 549. LATE STAGES OF STELLAR EVOLUTION

JORIS A. D. L. BLOMMAERT1,∗, JAN CAMI2, RYSZARD SZCZERBA3 and MICHAEL J. BARLOW4 1Instituut voor Sterrenkunde, K.U. Leuven, Celestijnenlaan 200B, B-3001 Leuven, Belgium 2NASA Ames Research Center, MS 245-6, Moffett Field, CA 94035, U.S.A. 3N. Copernicus Astronomical Center, Rabianska´ 8, 87-100 Torun,´ Poland 4Department of Physics & Astronomy, University College London, Gower Street, London, WC1E 6BT, U.K. (∗Author for correspondence: E-mail: [email protected])

(Received 18 October 2004; Accepted in ﬁnal form 2 November 2004)

Abstract. A large fraction of ISO observing time was used to study the late stages of stellar evolution. Many molecular and solid state features, including crystalline silicates and the rotational lines of water vapour, were detected for the ﬁrst time in the spectra of (post-)Asymptotic Giant Branch (AGB) stars. Their analysis has greatly improved our knowledge of stellar atmospheres and circumstellar environments. A surprising number of objects, particularly young planetary nebulae with Wolf-Rayet (WR) central stars, were found to exhibit emission features in their ISO spectra that are characteristic of both oxygen-rich and carbon-rich dust species, while far-IR observations of the PDR around NGC 7027 led to the ﬁrst detections of the rotational line spectra of CH and CH+.

Keywords: Infrared: stars, Stars: AGB and post-AGB, Stars: atmospheres, Stars: late-type, Stars: mass-loss, Stars: symbiotic, Stars: carbon, Stars: supergiants, Stars: Wolf-Rayet, Planetary Nebulae, Novae, dust, molecules

1. Introduction

ISO has been tremendously important in the study of the ﬁnal stages of stellar evolution. A substantial fraction of ISO observing time was used to observe different classes of evolved stars. IRAS had already shown the strong potential to discover many evolved stars with circumstellar shells in the infrared wavelength range. ISO provided the opportunity to extend this survey with higher sensitivity and spatial resolution, allowing many more sources to be detected in the inner regions of our Galaxy and in the Magellanic Clouds. But especially ISO’s spectroscopic capabilities in the 2–200 μm wavelength range allowed a detailed study of individual objects through observations of gas phase molecular bands, ionic forbidden lines and solid state bands. These studies are important not only to better understand the

2. The Asymptotic Giant Branch Stars

When low- and intermediate-mass stars have exhausted hydrogen in their cores, they ascend the HR diagram to become cool giants, ﬁrst on the Red Giant Branch (RGB) and later on the AGB. These phases are characterized by substantial mass loss; in the AGB phase, the mass loss effectively drives stellar evolution. The low effective stellar temperatures as well as the copious amounts of circumstellar material that form as a consequence of this mass loss make these objects ideal targets in the infrared. Asymptotic Giant Branch stars have been extensively studied with the various instruments onboard ISO, and have greatly enhanced the understanding of the composition, dynamics and interaction of the stellar atmosphere and the surrounding circumstellar material. In particular, SWS and LWS have greatly enhanced our knowledge of the composition of the molecular envelope and the sometimes thick dust shells surrounding these stars, while the sensitivity of ISOCAM greatly extended our knowledge of the fundamental properties of tens of thousands giant stars.

2.1. MOLECULES IN THE STELLAR ATMOSPHERES ...

Asymptotic Giant Branch stars are generally divided into oxygen-rich and carbon- rich stars. All stars start of with a photospheric C/O ratio which is lower than unity; during the evolution on the AGB however, dredge-ups bring freshly processed carbon to the surface, and gradually enrich the atmosphere in carbon. Depending (mainly) on the initial mass, some stars eventually experience a carbon preponder- ance in their atmosphere, and then they are called carbon stars. Asymptotic Giant Branch stars have very extended and relatively cool (typically 3000 K) atmospheres, allowing the formation of molecules in the stellar atmospheres. As the stellar envelope material cools down, CO is about the ﬁrst molecule to be formed. This process is so efﬁcient (and CO so stable) that for O-rich stars virtually all carbon is locked up in CO, leaving only oxygen for further chemistry. For carbon stars, all oxygen is locked up in CO and the chemistry is carbon driven. The difference in spectral appearance is enormous, both in the molecular and dust contents (see Tables I and II). LATE STAGES OF STELLAR EVOLUTION 217

TABLE I Overview of molecular bands and dust features reported in ISO observations of O-rich AGB stars

Component λ (μm) References

Molecular features H2O 2.3–197 [1–20, 29] OH 2.80–4.00 [2, 4, 13–18] CO Fundamental 4.6 [2, 4, 14–19] First overtone 2.3 [2, 4, 14–19] SiO Fundamental 8.0 [5, 14–19] First overtone 4.0 [2, 4, 13–19] CO2 ν3 Asymmetric stretch 4.2 [2, 4, 6, 14, 15, 17, 18] ν2 Bending 14.96 [3, 6, 9, 14, 15, 17, 18, 24, 28] Various combination bands 12–18 [3, 6, 9, 14, 15, 17, 18, 24] SO2 ν3 Asymmetric stretch 7.55 [5, 15, 17, 18]

Dust features Magnesiow¨ustite Mg(1−x)Fex O 19.5 [15, 21, 24] Alumina Al2O3 11.0 [15, 21] Spinel MgAl2O4 13.0 [15, 22–24] Crystalline water-ice 43,60 [25, 26] Metallic Fe 2–10 [27]

For an overview of the silicate dust we refer to (Molster and Kemper, this volume). [1] Neufeld et al. (1996); [2] Tsuji et al. (1997); [3] Justtanont et al. (1998); [4] Yamamura et al. (1999a); [5] Yamamura et al. (1999b); [6] Cami et al. (2000); [7] Tsuji (2000); [8] Zubko and Elitzur (2000); [9] Markwick and Millar (2000); [10] Tsuji (2001); [11] Jørgensen et al. (2001); [12] Jones et al. (2002); [13] Matsuura et al. (2002a); [14] Matsuura et al. (2002b); [15] Cami (2002); [16] Decin et al. (2003); [17] Van Malderen (2003); [18] Justtanont et al. (2004); [19] Van Malderen et al. (2004); [20] Truong-Bach et al. (1999); [21] Posch et al. (2002); [22] Posch et al. (1999); [23] Fabian et al. (2001); [24] Sloan et al. (2003); [25] Sylvester et al. (1999); [26] Dijkstra et al. (2003a); [27] Kemper et al. (2002a); [28] Ryde et al. (1999); [29] Barlow et al. (1996).

The presence of molecules can easily be inferred through their ro-vibrational bands in the infrared. Such molecular bands have been used – amongst others – with great success to extend the MK classiﬁcation to the infrared (Vandenbussche et al., 2002) and to improve the calibration of SWS (see Decin et al., 2003) using detailed models for A0–M2 stars; in turn, the SWS observations yielded many improvements in the models. Different molecules turn out to be sensitive to different stellar parameters, thereby, allowing to determine these stellar parameters from the ISO observations. The ISO observations of molecular bands have revealed the importance of fundamentally new ideas about modelling of such stars; one example is the development of fully consistent hydrodynamical codes, which include a frequency dependent radiative transfer (see e.g. H¨ofner, 1999; Gautschy-Loidl et al., 2004). 218 J. A. D. L. BLOMMAERT ET AL.

TABLE II Overview of molecular bands and dust features reported in ISO observations of C-rich AGB stars

Component λ (μm) References

Molecular features CO Fundamental 4.6 [1, 2, 4, 8] CH Fundamental 3.3–4.1 [4, 8] CS 1st overtone 3.9–4.2 [1, 4, 8] Fundamental 7–8 [1, 4, 5, 8] SiS First overtone 6.6–7 [4] HCN Various bands 2.58–3.86 [1, 2, 4, 7, 8] 2ν2 7.0 [1, 4, 8] ν2 14.04 [1, 5, 6, 8] ν0 − ν1 2 2 2 14.30 [5, 6, 8] C3 5.2 [1–3, 8] C2H2 ν3 3.05 [1, 2, 8] Various bands 3.8 [1, 2, 8] ν4 + ν5 7–8 [1, 2, 5, 8] ν5 13.7 [1, 2, 6, 8]

Dust features SiC 11.3 [2] MgS 26–35 [2, 9]

[1] Jørgensen et al. (2000); [2] Yang et al. (2004); [3] Hron et al. (1999); [4] Aoki et al. (1998); [5] Aoki et al. (1999); [6] Yamamura et al. (1999c); [7] Harris et al. (2003); [8] Gautschy-Loidl et al. (2004); [9] Hony et al. (2002a).

2.2. ...AND BEYOND:THE “EXTENDED”ATMOSPHERE OF AGB STARS

However, early in the ISO lifetime, it became clear that the molecular absorption seen in the SWS spectra of several O-rich AGB stars is quite different than was expected from detailed hydrostatic model calculations. Tsuji et al. (1997) noticed that CO and SiO absorption in early M giants was weaker than expected, while water vapour and CO2 bands were much stronger than predicted; in fact, the latter molecule was not predicted to be present in the stellar atmosphere at all. They found that the molecular features due to CO, CO2 and SiO correspond to excitation temperatures in the range 750–1250 K, indicating that they originate outside the stellar photosphere. This led to the terminology “extended atmosphere”, “warm molecular layer” or simply MOLsphere to denote that region close to the stellar photosphere (typically up to a few R ) where this excess molecular absorption or emission originates from and corresponds to a quasi-static molecular zone that was already hypothesized to exist based on ground-based data. LATE STAGES OF STELLAR EVOLUTION 219

Soon, a wealth of molecular bands from this extended atmosphere were discovered and/or analysed (see Table I). These molecular bands are now used as probes to map the physical conditions (such as temperature and density) in the direct surroundings of the AGB stars (e.g. Cami et al., 2000; Justtanont et al., 2004). Of particular importance is water vapour, which is detected in both SWS and LWS spectra of O-rich AGB stars (and many others; see Cernicharo and Crovisier, this volume) through literally millions of spectral lines, and is the source for large differences in spectral appearance between Semi-Regulars and Miras. For Semi- Regulars, water vapour is just a component in the spectra up to 10 μm, and the overall continuum shape is still determined by the stellar photosphere. In Mira variables on the other hand, the water vapour layer is dense, and due to the large opacity over a large wavelength range, this layer determines the shape of the IR continuum up to at least 10 μm in these objects. Analyses based on SWS observations suggested that this semiopaque layer typically extends to ∼2R ; this turns out to be in good agreement with interferometric observations in the infrared (Cami, 2003). Also, the pure rotational H2O lines were ﬁrst detected in the SWS (Neufeld et al., 1996) and LWS (Barlow et al., 1996) spectra of AGB stars. Molecular bands are also detected in the spectra of carbon stars. Absorption features due to C2, HCN (and HNC), C2H2 and C3 are commonly detected (see Table II); moreover, some (likely molecular) features detected in the SWS spectra are still unidentiﬁed (Yang et al., 2004). While the excitation temperatures of these molecular bands are in some cases comparable to the molecular bands in the O-rich stars, there is still debate about the exact origin of these molecular bands (see e.g. (Jørgensen et al., 2000)). It is clear, however, that for the carbon stars also, the ISO observations have greatly expanded the possibilities to study the stellar and (possibly) circumstellar envelope.

2.3. DUST

The importance of ISO observations for the study of circumstellar dust can hardly be overestimated. The discovery of crystalline silicates in a wide variety of objects (see Molster and Kemper, this volume) and its consequences for the study of circumstellar and interstellar dust opened up the ﬁeld of astromineralogy. Observa- tions with SWS and LWS have resulted in the detection and identiﬁcation of many more dust components in the circumstellar dust shells of AGB stars, and shed new light on the pathways to dust formation. Moreover, ISO observations have allowed to directly link the molecular features to the dust content in these stars (see e.g. Cami, 2002; Sloan et al., 2003). In the spectra of O-rich stars with low mass-loss rates, the typical silicate bands commonly found in Miras are often at most a minor component in the dust spectra. In these stars, prominent features in the ISO data have provided convincing evidence for a prominent role played by simple oxides such as Al2O3 (11 μm) or MgFeO (19.5 μm). The famous 13 μm feature (already detected in IRAS-LRS spectra; see 220 J. A. D. L. BLOMMAERT ET AL.

(Vardya et al., 1986) is attributed to spinel (MgAl2O4, Posch et al., 1999; Fabian et al., 2001), although there is still a lively debate about this identiﬁcation. For instance, Sloan et al. (2003) ﬁnd no correlation between the 13 μm feature and an observed band at 32 μm which was also presumably due to spinel. For stars with higher mass-loss rates, silicates (both amorphous and crystalline) are the dominant dust component (see Molster and Kemper, this volume). Also the presence of crystalline water-ice and metallic Fe has been demonstrated in these stars. Many more components (especially at longer wavelengths) can be observed in the more evolved stars (post-AGB and PNe, see later sections). For the carbon stars, the situation is quite different. The bulk of the dust is in the form of amorphous carbon, which determines the energy distribution but has no resonances in IR. Most energy distributions of the carbon stars seem to follow a sequence of decreasing TNIR which is interpreted as a cooling and thickening of the circumstellar dust shell. The infrared spectra of carbon stars furthermore show only a small number of solid state bands. Commonly detected dust features include the 11.3 μm feature due to SiC (Gilra, 1973) and the “30 μm feature”, which is believed to be due to MgS (see e.g. Hony et al., 2002a). A particular class of objects are the “mixed chemistry” objects. These are stars that show the presence of typically O-rich and C-rich dust at the same time. In all known cases, the O-rich material is cool and often highly crystalline, while the star itself and/or a warm dust component are typically rich in C. While such objects might represent stars that have only recently become a carbon star, it turns out that most of the well-studied objects in this class belong to a binary system. In such cases, the O-rich material is stored in a disk around the companion star (see e.g. Yamamura et al. (2000) or in a circumbinary disk (such as in the Red Rectangle, see Waters et al. (1998a)).

2.4. STUDIES OF AGB STARS IN DIFFERENT ENVIRONMENTS

Mid-infrared studies before ISO were mostly limited to relatively nearby objects; this was especially true for spectroscopical studies. ISO provided the sensitivity and the spatial resolution to investigate AGB stars in the inner regions of the Milky Way galaxy and in the Magellanic Clouds. Studies in these regions offer several advantages, ﬁrst it allows to study different populations (different metallicities) and second, often the distance is known and thus the luminosities.

2.4.1. The Inner Milky Way The second largest program performed with the ISO satellite, called ISOGAL, is a mid-infrared survey along the galactic plane, performed with ISOCAM (Omont et al., 2003). About 16 square degrees were observed mostly at two wavelengths: 7 and 15 μm. In comparison with IRAS, this survey provides data that are 2 orders of magnitude more sensitive and have 1 order of magnitude higher spatial LATE STAGES OF STELLAR EVOLUTION 221

Figure 1. ISOGAL [15] vs. (Ks − [15])0 diagram of galactic bulge sources, taken from Ojha et al. (2003). The different symbols are used to indicate sources for which different sorts of associations or data exist. The curves represent sequences of luminosities ranging from 500 to 20,000L, with increasing mass-loss rates. The approximate scales of mass-loss rates are displayed at the top. For further details on how these were derived (see Ojha et al., 2003).

resolution. A systematic cross-identification with near-infrared (I, J, Ks) DENIS (Epchtein et al., 1997) sources was performed for the ISOGAL detected sources. This has led to a five-wavelength ISOGAL catalogue, which can be obtained through VizieR. Several papers describe results on AGB stars obtained on fields towards the galactic bulge where the extinction is less strong than along the plane. Figure 1 shows a [15] versus (Ks − [15])0 diagram of bulge ISOGAL sources, taken from Ojha et al. (2003). A clear linear sequence of increasing (Ks − [15])0 colour for brighter 15 μm fluxes can be seen, as was also originally presented in the earlier Omont et al. (1999) paper for the [15] versus [7]–[15] diagram. The increasing red colour of (Ks − [15])0 as function of [15], can be explained by an increasing mass-loss rate. The exact amount of the mass-loss rate is model dependent, but is −9 −7 in the order of 10 toafew10 M per year (Ojha et al., 2003) for the sources in the sequence and higher for the redder sources ((Ks − [15])0 > 2.5) which are identified as Mira variables (see the following paragraph). An important point to mention is that the relatively low mass-loss rates observed by the 15 μm excess cannot be observed through the near-infrared colours. Omont et al. (1999) and Glass et al. (1999) demonstrated that the observed sources are red giants above the RGB limit. A comparison by Glass et al. (1999), 222 J. A. D. L. BLOMMAERT ET AL. and Glass and Schultheis (2002) of ISOGAL data with a spectroscopic survey of Baade’s window NGC 6522 shows that M giants as early as M2 are detected by ISOGAL, but that it is complete from M5 onwards. A further investigation of the circumstellar dust of the mass losing AGB stars comes from ISOCAM Circular Variable Filters measurements. About 20 sources were observed in three bulge fields (5–17 μm). As was expected from the colours, several show clear evidence of a mid-IR excess related to the circumstellar dust (Blommaert et al., 2000, manuscript in preparation). Most sources do not show the typical relatively peaked silicate 9.7 μm feature, but rather a broad feature with a maximum peaking at longer (∼12 μm) wavelengths. Such features are known from previous work on IRAS data (Sloan and Price, 1995, and references therein). They have been associated with amorphous aluminium oxide grains (Onaka et al., 1989; Speck et al., 2000a). Such features are mostly seen in low mass-loss rate sources, confirming the interpretation that ISOGAL is mainly detecting the onset of mass loss on the AGB. Another important characteristic of AGB stars is their variability. Typical examples are Mira and Semi-Regular variables, with periods in the range of hundred to a few hundred days. Mass loss and variability are related in the sense that the longer period and larger amplitude variable stars also show higher mass-loss rates with the extreme examples being the OH/IR stars, which have the longest periods −5 (up to a thousand days) and the highest mass-loss rates (∼10 M per year). Alard et al. (2001) and later also Glass and Schultheis (2002) combined the ISOGAL data with data coming from the gravitational lensing experiment MACHO for the Baade windows. In comparison to earlier searches for variable stars (e.g. Lloyd Evans, 1976), the new surveys find variability down to much smaller amplitudes (0.5 mag compared to better than 0.1 mag). Alard et al. (2001) conclude that almost all sources detected at 7 and 15 μm, and, thus, above the RGB limit are variable. A period of 70 days or longer is a necessary but not a sufficient condition for mass loss to occur (Figure 2). ISOGAL also allows to study the AGB stars as tracers for the underlying stellar population. First results are given in van Loon et al. (2003) for the bulge. The ISOGAL survey was also used to find the most extreme mass-loss rate AGB stars near the Galactic Centre. Ortiz et al. (2002) found infrared counterparts for all the objects detected in the OH (1612 MHz) maser surveys of the Galactic Centre. Blommaert (2003) observed the Spectral Energy Distributions of a sample of 26 OH/IR stars in the wavelength range of 2–60 μm. The main results from these studies are that the peak of the bolometric magnitude distribution falls at −5 and that only a relatively small number of stars have high luminosities (and thus less than 1 Gyr old). The OH/IR stars fall predominantly below the extension of the Mira Period–Luminosity relation. Through the modelling of the SEDs, indi- −6 −4 vidual total mass-loss rates were determined, which range from 10 to 10 M per year. LATE STAGES OF STELLAR EVOLUTION 223

Figure 2. A15μm excess (indicative of mass loss) vs. log(Period (days)), taken from (Alard et al., 2001).

2.4.2. The Magellanic Clouds Studying the AGB stars in the Magellanic Clouds has the advantage that the distances and thus the luminosities are known. ISO photometry and spectroscopy of a sample of 57 bright AGB stars and Red SuperGiants in the LMC are presented by Trams et al. (1999b). The sources were selected on basis of IRAS colours indicative of high mass-loss rates. From ISO colour–colour diagrams and the spectra, the spectral type of the dust was determined. About half the sample are classified as carbon stars. No M stars with carbon-rich dust were found, suggesting that Hot Bottom Burning (HBB) cannot efficiently turn back carbon stars into oxygen-rich ones. Mass-loss rates and luminosities for this sample have been determined by van Loon et al. (1999), using a radiative transfer code. The dust-enshrouded carbon stars are generally brighter than the optically bright carbon stars but less bright than the high mass-loss M-type ones. The interpretation given is that only stars between 2 and 4M become carbon stars, whereas more massive stars remain oxygen rich because of HBB. The existence of a few luminous carbon stars may be explained by HBB switching off near the very tip of the AGB and the occurrence of a final thermal pulse thereafter. Connected to this is the first detection of silicate dust around a carbon star (Trams et al., 1999a). In contrast to galactic carbon stars with silicate dust, IRAS04496-6958 is very luminous (Mbol =−6.8). Possibly, the silicate dust is a trace of its oxygen-rich past and it became carbon rich after its last thermal pulse. Most of the supergiants have lower mass-loss rates than the mass consumption rate by nuclear burning. Only two have a much higher mass-loss rate, suggesting that the RSGs lose most of their mantles in very short phases. Also AGB stars have episodes of intensified mass-loss rates, but the mass-loss rate is always higher than the nuclear burning rate. 224 J. A. D. L. BLOMMAERT ET AL.

Blommaert et al. (1999) studied the IRAS sources towards the SMC and find similar results as for the LMC study. There are, however, no carbon stars with luminosities near the AGB limit. All stars suffer extremely high mass-loss rates −5 (M˙ > 10 M per year), even stars with relatively low masses (∼1M). The earlier studies are based on sources that were already known from the IRAS survey and are thus restricted to the brightest AGB stars. ISOCAM was also used to detect new and lower luminosity AGB stars. Tanabé et al. (1998) performed a survey of MC clusters and detected very red stars at Mbol =−4.5. High mass-loss rates can thus also be obtained by lower luminosity stars with likely lower initial masses. Another survey with ISOCAM was performed at 4.5, 6.7 and 12 μmon several regions of the Large and Small Magellanic Clouds (Loup et al., 1999). First results were discussed in Cioni et al. (2003), who combined the ISO data with 2MASS and DENIS near-infrared data and light curves extracted from the MACHO database. They confirmed the AGB or supergiant nature of the detected point sources and classified the AGB stars in Mira or Semi-Regular pulsators.

3. Post-AGB Stars

The AGB phase of stellar evolution ends in an episode of extremely high mass −5 loss (so-called superwind), in which the mass-loss rate exceeds 10 M per year. The high mass loss during the end of AGB evolution creates a circumstellar envelope, which may completely obscure the star. The superwind rapidly strips the star of virtually its entire envelope, thus, effectively terminating the AGB phase. The central star subsequently evolves to higher temperatures on a timescale that depends on its mass and, at the same time, the ejected envelope expands and cools. The star is now in the post-AGB phase and depending on the rate of its evolution, the star may start to photoionize the ejected material, forming a planetary nebula (PN). Post-AGB objects hold the key to understanding how largely spherical AGB circumstellar shells transform into diverse aspherical morphologies (see e.g. Ueta et al., 2000, and references therein). The post-AGB objects that will become planetary nebulae in the near future are called proto-planetary nebulae (PPNe). Since masses of the central stars and, therefore, timescales of their evolution are very uncertain, the term proto-PN usually expresses only our belief that a PN will be created. In this review, however, both terms: post-AGB object and proto-PN will be used equivalently unless there are strong arguments suggesting that a PN will not be formed. For a more detailed discussion of the post-AGB phase of stellar evolution the reader is referred to the Toru´nConference Proceedings: Post-AGB objects as a phase of stellar evolution (edited by Szczerba and G´orny, 2001) and to the recent reviews by van Winckel (2003) and by Waelkens and Waters (2004). Just like in the previous sections, ISO results related to post-AGB objects with O-rich and C-rich circumstellar shells are discussed separately. Special importance will be given to the signatures of mixed chemistry. LATE STAGES OF STELLAR EVOLUTION 225

3.1. POST-AGB OBJECTS WITH O-RICH CIRCUMSTELLAR SHELLS

As far as dust is concerned, the most important ISO discovery seems to be the detection of crystalline silicates from both young and old stars. Lack of crystalline silicate features in the interstellar medium probably means that after their formation during stars’ death they do not survive until new stars are born. In a series of papers (Molster et al., 2002a,b,c) crystalline silicate features are investigated in the ISO spectra of 17 O-rich circumstellar dust shells surrounding evolved stars. The considered sample contains six post-AGB objects and HD 179821, for which it is yet a matter of debate whether this object is a massive supergiant or a low-mass object in its post-AGB phase of evolution. Apart from the broad 10 and 18 μm bands that are due to amorphous silicates, Molster et al. found at least 49 narrow bands between 10 and 60 μm (divided into seven complexes), which can be attributed to crystalline silicates (olivines and pyroxenes). The richness of the crystalline silicate features in the ISO spectra contains important information related to the conditions in which the grains were formed and processed. It allows detailed studies of the mineralogy of these dust shells. Note, however, that abundances of crystalline silicates are rather low and do not exceed about 10–15% of the total amount of silicates (see Molster, 2000, and references therein). Molster et al. (2002a) found that the crystalline silicate band strengths correlate with the geometry of circumstellar environment: disk (strong) or outflow sources (weak crystalline silicate bands). It has been suggested that in disk sources low temperature crystallization may take place (Molster et al., 1999), while in outflow sources high temperature crystallization in layers that are close to the star is more probable. Among six analysed post-AGB stars, four possibly have disk and two are classified as outflow sources. It is worth to note that two of the disk sources (Red Rectangle and Roberts 22) show also carbon-rich dust as evidenced by the polycyclic aromatic hydrocarbon (PAH) features. The PAHs are predominantly present in the scattering lobes, while the crystalline silicates are expected to be present in the disks (Waters et al., 1998a). A full 2–200 μm ISO spectrum was used to constrain the dust properties in a post-AGB star HD 161796 (Hoogzaad et al., 2002). A good fit to the spectral energy distribution was achieved using four co-spatial but distinct dust components: amorphous silicates (63%), forsterite (4%), enstatite (6%) and crystalline water-ice (27%). The derived temperature of water-ice suggest that the ice must be formed as a mantle on top of an amorphous silicate core, which requires high mass-loss −5 rates, exceeding 5 × 10 M per year. The derived mass-loss rate is high enough to allow for this process to be efficient. An interesting example of an OH/IR post-AGB star, still obscured by the ejected material, is IRAS 16342–3814 (OH 344.1+5.8, also called “the water fountain nebula”). This source is classified as a proto-PN (Sahai et al., 1999) and is the only OH/IR star known to have crystalline silicate absorption features seen up to almost 45 μm in its SWS spectrum (Dijkstra et al., 2003b). This suggests that 226 J. A. D. L. BLOMMAERT ET AL.

IRAS 16342−3814 recently had an extremely high mass-loss rate and is perhaps the youngest PPN observed until now. Besides the crystalline silicates (forsterite, diopside and possibly clino-enstatite) identified in the SWS spectrum, crystalline water-ice is also detected. In the LWS spectrum an unidentified feature at 48 μmis seen. A deep and wide water-ice band at 3.05 μm is seen also in another young OH/IR PPN: IRAS 22036 + 5306 (Sahai et al., 2003), together with wide absorption by silicates around 10 and a 60 μm feature, perhaps due to crystalline H2O-ice (the absence of a corresponding 43 μmH2O-ice feature may be due to a decline in its strength at low temperatures; Dijkstra et al. (2003a). The freeze-out of water onto dust grains indicated by the ice features requires high densities, which are likely to be characteristic of the dusty disk seen in the Hubble Space Telescope image. In addition, in the SWS spectrum of this object, an unidentified feature at 3.83 μm (tentatively attributed to the H2S) is detected. A sample composed of O-rich PPNe has been observed with the ISO spectrometers to search for atomic fine structure (FS) lines (Castro-Carrizo et al., 2001). The low-excitation transitions of [O I], [C II], [N II], [Si I], [S I], [Fe I] and [Fe II] were observed. Taking into account the sample of C-rich PPNe presented by Fong et al. (2001) (see Section 3.2), they concluded that PPNe emit in these atomic transitions only when the central star is hotter than about 10 000 K. This result suggests that such lines predominantly arise from photo-dissociation regions (PDRs), and not from the shocked regions.

3.2. POST-AGB OBJECTS WITH C-RICH CIRCUMSTELLAR SHELLS

ISO SWS observations of C-rich PPNe have been discussed by Hrivnak et al. (2000) with emphasis given to wavelengths longer than about 20 μm. At these wavelengths, prominent dust features at 21 and around 30 μm are seen. The 21 μm band has been discovered by Kwok et al. (1989) in the IRAS low-resolution spectra of some C-rich PPNe, while the 30 μm band in this class of objects has been detected in the Kuiper Airborne Observatory spectra by Omont et al. (1995). The 21 μm feature is not seen in the spectra of either the precursors to PPNe or the successors of PPNe. Note, however, that Hony et al. (2001b) reported tentative detection of this feature in two planetary nebulae with [WR] central stars, while Pei and Volk (2003) reported its detection in PN IC 418, and Volk et al. (2000) suggested possible presence of this band in two extreme carbon stars. On the other hand, the 30 μm band has been detected first in some C-rich AGB stars and in C-rich planetary nebulae (see Omont et al., 1995, and references therein). Thanks to the ISO observations, Volk et al. (1999) were able to obtain an intrinsic profile of the 21 μm feature, in fact centered at 20.1 μm. This feature has been attributed to various molecular and solid-state species, none of which satisfy all the constraints, although titanium carbide, PAHs and even silicon carbide (see LATE STAGES OF STELLAR EVOLUTION 227 discussion in Speck and Hofmeister, 2004, and references therein) seem to be the most favoured candidates. On the other hand, in spite of some problems with the SWS data above 26 μm, it seems that the 30 μm feature is not so smooth as the 21 μm one and, instead, shows a kind of substructure (Szczerba et al., 1999). In fact, analysis of the ISO data for other post-AGB sources apparently resolves the 30 μm feature into two components at about 26 and 33 μm (Volk et al., 2002). The derived intrinsic profiles for these features opens a possibility to search for their carrier(s). Recently, Hony et al. (2002a) analysed ISO spectra of C-rich sources including low mass carbon stars, extreme carbon stars, post-AGB objects and planetary nebulae. They modelled the whole range of the extracted “30” μm band profiles by using magnesium sulphide. They argued that in some sources a residual emission at approximately 26 μm can also be fitted using MgS grains but with a different grain shape distribution. The further strengthening of the MgS identification as a carrier of the “30” μm feature comes from detailed modelling of HD 56126, one of the post-AGB “21 μm object” (see Hony et al., 2003). The ISO spectra of C-rich sources show many dust features including Unidenti- fied Infrared Bands (UIRs), which are attributed very often to PAHs (see a comprehensive review by Tielens et al., 1999; Hony et al., 2001a). Besides the well known UIR bands at 3.3, 6.2, 7.7, 8.6 and 11.3 μm, ISO has revealed a wealth variety of weaker features, satellite bands and sub-features. While the final identification of the carrier(s) responsible for these features is still under discussion, there is an agreement that the observed UIR bands are very characteristic for aromatic structures (the C atoms are arranged in planar hexagons – like in graphite – with attached hydrogen atoms at the edges). A detailed analysis of SWS spectra for different kind of objects (Peeters et al., 2002) demonstrated that the UIR emission features in the region 6–9 μm clearly show profile variations. The observed variations in the characteristics of the UIR emission bands are linked to the local physical conditions, but they do not reflect – at least at first glance – a clear link between PAH spectrum and the post-AGB source evolutionary state (see Peeters et al., 2002; Peeters et al., this volume, for details). ISO observations of post-AGB objects demonstrated that there is an evolution of carbonaceous material from aliphatic structures (the C atoms are arranged in a tetrahedral network – like in diamond – with attached H atoms at the edges) to the aromatic structures during the evolution from the proto-planetary to the planetary nebulae phase (Kwok et al., 1999, 2001). AFGL 618 and AFGL 2688 were the first stars (already 30 years ago) to be considered to be in transition from AGB to PN, and they are now the best investigated sources among PPNe. When ISO results are discussed for C-rich PPNe, very often these two objects are given as an example and in this review the situation will be similar. Let us start with polarimetric imaging of AFGL 2688 at 4.5 μm with ISOCAM which allowed to conclude that the polarization of starlight induced by dust grains is almost independent of wavelength between 2 and 4.5 μm (Kastner et al., 2000). This finding indicates that scattering dominates over thermal emission at these wavelengths, and that the dust grains have characteristic radii less than 1 μm. 228 J. A. D. L. BLOMMAERT ET AL.

ISOPHOT linear scans of AFGL 2688 and AFGL 618 demonstrated that both these objects have extremely extended dust shells (see Speck et al., 2000b. As- suming constant expansion velocities, ages for these dust shells are of the order of 105 years. The infrared intensities show “periodic” enhancements, timescales for which (few times 104 years) coincide with thermal pulses on the AGB. ISO spectroscopy allowed also to detect molecular bands in spectra of some C-rich post-AGB objects. For example, Cernicharo et al. (2002) reported the detection of a molecular band at 57.5 μm in PPNe AFGL 618 and AFGL 2688 that has been tentatively assigned to the ν5 bending mode of C4. Polyynes such as C2H2,C4H2 and C6H2, and single aromatic species such as benzene, have been detected in AFGL 618 (Cernicharo et al., 2001a,b). These discoveries support the idea that more complex C-rich molecules could be formed in space. However, the mechanisms allowing the growth of carbon-rich molecules are still poorly known, and the full set of molecules that could participate in the chemical reactions leading to the formation of large complex carbon-rich species has yet to be identified. Analysis of the LWS spectra for the two aforementioned PPNe allowed to detect rotational lines of 12CO (J = 14–13 to J = 41–40) and lines of 13CO (J = 14–13 to J = 19–18), as well as HCN and HNC (Justtanont et al., 2000; Herpin et al., 2002). In the early stages, represented by AFGL 2688, the longwave emission is dominated by CO lines. In the more advanced stage (AFGL 618), very fast outflows are present, which, together with the strong UV radiation field from the central star, dissociate CO. The released atomic oxygen allows formation of new O-bearing species, such as H2O and OH, in a C-rich environment. In case of AFGL 618, several lines of OH and H2O are detected (Herpin and Cernicharo, 2000; Herpin et al., 2002). In CRL 618 the abundance of HNC is enhanced with respect to HCN as a result of chemical processes occurring in the PDR. HR 4049 was first suggested to be an object in the transition phase from the AGB to PN stage by Lamers et al. (1986) who discovered both a large IR excess and severe UV deficiency, indicating the presence of circumstellar dust. Presently, the source is considered to be the prototype of a group of post-AGB stars in a binary system with extremely metal-depleted atmospheres (van Winckel et al., 1995). The ISO/SWS spectrum of this object shows the signatures of C-rich dust (PAHs; Beintema et al., 1996) and the presence of O-rich gas (isotopomers of CO2; Cami and Yamamura, 2001). Analysis of the ISO/SWS spectrum together with other literature data allowed to propose that a very optically thick circumbinary disk could be response for the observed IR emission from HR 4049 (Dominik et al., 2003). A sample composed of C-rich PPNe has been observed with the ISO spectrometers to search for atomic FS lines (Fong et al., 2001). The low-excitation transitions of [O I], [C II], [Si I], [Si II], [S I], [Fe I], [Fe II], [Ne II] and [N II] were observed. However, only few lines in few PPNe were detected. LATE STAGES OF STELLAR EVOLUTION 229

3.3. OTHER POST-AGB OBJECTS

There are two classes of objects that, on the basis of their photospheric abundances, pulsation properties and circumstellar material have been classified as post-AGB stars. They are RV Tau stars and R CrB stars. Both groups of objects are known to have circumstellar dust, giving support to their evolved nature. There are few RV Tau stars observed with ISO. The typical SWS spectrum shows emission bands of amorphous silicates at 10 and 18 μm (see e.g. Szczerba et al., 2003). However, as demonstrated by van Winckel et al. (1998) in case of AC Her, the unusual broad feature at 8–12 μm is due to a high crystallization fraction of the circumstellar silicates. The authors argued that the circumstellar material is trapped in a long- lived disk in the system similar to what is observed in the case of the Red Rectangle. Matsuura et al. (2002b) analysed SWS spectra of the RV Tau star, R Scuti and demonstrated that those spectra are dominated by H2O emission bands with CO, SiO and CO2 bands also present. They argued, however, that R Sct may be a thermally pulsing AGB star, observed in a helium burning phase. Lambert et al. (2001) discussed spectra of some R Coronae Borealis stars. They reported detection of the sharp emission features, which coincides with those of UIRs, only in one of the analysed objects, V 854 Cen. The features coincide with those from laboratory samples of hydrogenated amorphous carbon. Since V 854 Cen, is of order of 1000 more abundant in hydrogen than other typical R CrB stars, the emission features are probably from a carrier containing hydrogen. The extreme H deficiency of the R CrB stars suggests that during their evolution some mechanism removed the entire H-rich stellar envelope. This may be related to a thermal pulse after the star left the AGB. A similar scenario probably applies to “Sakurai’s object”, a star that has experienced dramatic changes during the last decade. The ISO data demonstrate the presence of hot circumstellar dust around this star (Eyres et al., 1998a; Kerber et al., 1999). There is a group of less massive post-AGB objects with central stars of B spectral type (so-called hot post-AGB stars) that have small or even do not have infrared excess at all. One should be aware, however, that a group of Herbig Ae/Be stars (young intermediate-mass pre-main-sequence stars) can be easily confused with such post-AGB candidates. The ISO observations allowed to detect C-rich and O- rich dust features in SWS spectra of some hot post-AGB candidates (Gauba and Parthasarathy, 2004), thus indicating the dominant chemistry (C- or O-based) in the circumstellar dust shells and their central stars. Let us finish this part on post-AGB stars with an attempt to classify SWS01 spectra for 61 PPNe by Szczerba et al. (2003). The main dust and/or molecular features were taken into account to propose a division of all spectra into seven classes (four for C-rich and three for O-rich sources). On the basis of their classification, they discussed the connection between post-AGB objects and planetary nebulae with emphasis on possible precursors of PNe with [WR] central stars. 230 J. A. D. L. BLOMMAERT ET AL.

4. Planetary Nebulae and Pre-Planetary Nebulae

We have seen earlier in this paper that the terms ‘proto-planetary nebulae’ and ‘post-AGB objects’ are often used interchangeably to denote objects that have left the AGB but have not yet developed signiﬁcant photo-ionised regions. We include ‘pre-planetary nebulae’ in this group too; the term is sometimes taken to mean that signs of the early development of a circumstellar photoionized region are present for objects in this category.

4.1. PN AND PRE-PN WITH DUAL DUST CHEMISTRIES

ISO targeted many O-rich and C-rich objects for spectroscopy using the SWS and LWS instruments; the enormous wavelength coverage (2.4–197 μm) that these two instruments delivered, together with their improved sensitivity relative to NASA’s pioneering Kuiper Airborne Observatory, led to the discovery of many new dust spectral features. One of the big surprises was the relatively high incidence of objects that simultaneously exhibit dust features from both O-rich and C-rich carriers. Ever since its discovery, the Red Rectangle, AFGL 915, has been considered to exhibit the archetypal 3–13 μm UIR-band spectrum (usually attributed to aromatic carbon species, such as PAHs), yet the SWS spectrum of this object shows many crystalline silicate emission features longwards of 20 μm (Waters et al., 1998a). Waters et al. (1998b) found that two planetary nebulae with strong UIR-band spectra, BD+30◦3639 and He 2–113, also exhibited a similar array of crystalline silicate emission features longwards of 20 μm. These two nebulae both have carbon-rich nebulae and cool late-type [WCL] WR central stars, as does CPD–56◦8032, which Barlow (1997) and Cohen et al. (1999) found to exhibit emission features due to crystalline water-ice, as well as crystalline silicates, in its ISO LWS and SWS spectrum. Cohen et al. (2002) presented the SWS and LWS spectra of more [WCL] PNe showing both strong UIR bands and strong crystalline silicate features, and dubbed them as having ‘dual dust chemistries’. A number of interpretations of the dual dust chemistry phenomenon have been considered by the aforementioned authors, including (a) a recent thermal pulse that has brought up and ejected C-rich material into an O-rich nebula (the phenomenon appears too frequent, however, to be consistent with such dredge-up events, which ought to be encountered only rarely); or (b) the presence of comet clouds at large radii, which have been dis- rupted by the impact of expanding C-rich nebulae, liberating O-rich silicate and ice particles. However, the most likely explanation now seems to be (c) the presence of pre-existing dust disks around binary systems in these objects, which contain O-rich particles that have been captured and stored from previous O-rich AGB evolutionary stages. Self-shielding disks would enable crystalline silicate and ice particles to remain relatively cool, while C-rich particles synthesised during more recent evolutionary events would be channelled towards polar directions. Direct evidence for the presence of optically thick dust disks around dual dust chemistry LATE STAGES OF STELLAR EVOLUTION 231 objects has come from the speckle interferometric detection by Osterbart et al. (1997) of an obscuring disk or torus in the Red Rectangle, and the detection of a similar obscuring disk or torus in CPD−56◦8032 by De Marco et al. (2002), from HST STIS optical and UV long-slit spectroscopy. As discussed in Section 3.2, a strong 21 μm emission feature is often detected around carbon-rich post-AGB objects. Its non-detection in the spectra of AGB stars and PNe had led to suspicions that the carrier was a transitory species. Hony et al. (2001) have, however, claimed the detection of a 21 μm emission feature in the SWS spectra of two PNe having H-deficient WR central stars, NGC 40 and NGC 6369. Hony et al. (2002b) also reported the detection of a 23 μm emission feature in the SWS spectra of the PNe M2-43 and K3-17, which they attributed to iron sulphide, with M2-43 also showing evidence for weaker FeS features at 34, 38 and 44 μm. M2-43 has a H-deficient WR central star; the nature of the central star in K3-17 is not currently known. Harrington et al. (1998) reported the detection of an unusual 6.4 μm emission feature in the SWS spectra of the PNe Abell 78 and IRAS 15154- 5258 (PM1-89), both of which have H-deficient WC central stars. They identified the feature as an aromatic C C stretch feature produced in carbonaceous grains with little or no hydrogen. It is striking that most of the discoveries discussed earlier have been associated with PNe that have H-deficient WC, WR central stars, a group that accounts for no more than 15% of all PNe. NGC 6302 is a much more evolved, high-excitation PN than the dual dust chemistry [WCL] PNe, with a highly bipolar geometry. Figure 3 shows its continuum- subtracted 2.4–197 μm spectrum (Molster et al., 2001). As well as showing a UIR-band spectrum in the 3–13 μm region (attributed to C-rich particles), longwards of 17 μm the spectrum is seen to be dominated by strong emission features, identified with crystalline silicates and crystalline water-ice (the latter at 44 and 62 μm). There is a very broad feature peaking at about 90 μm in Figure 3, which was not identified by Molster et al., although they suggested that it might be due to hydrous silicates. Kemper et al. (2002b) identified this feature as due to a 92.6 μm band of calcite, CaCO3. They also suggested that dolomite, CaMg(CO3)2, contributes to the observed 62 μm feature. The interest in the calcite identification is that this material normally requires the presence of liquid water for its formation. Water-ice features are identified in the spectrum of NGC 6302, but a non-aqueous formation route would seem to be required for the suggested calcite component in NGC 6302. The relatively narrow emission feature at 69 μm in Figure 3 has been identified with the crystalline silicate material forsterite, Mg2SiO4.Boweyet al. (2001) have shown that the peak wavelength of this feature is a sensitive function of temperature for laboratory particle samples, shifting by nearly 1 μm as grains are cooled from 300 to 4 K. Bowey et al. (2002) used these results to derive dust temperatures of 30–140 K from the observed peak wavelengths of the 69 μm feature in the ISO LWS spectra of a range of PNe and post-AGB objects. 232 J. A. D. L. BLOMMAERT ET AL.

Figure 3. The continuum-subtracted spectrum of NGC 6302 from 2.4 to 197 μm. As well as many ionic emission lines, the 3–13 μm region is characterised by the presence of C-rich emission features, while from 17 to 80 μm the spectrum is dominated by emission bands attributable to crystalline silicates and crystalline water-ice (from Molster et al., 2001).

4.2. THE NEUTRAL ZONES AROUND PN

Many planetary nebulae have surrounding neutral atomic and molecular zones that emit strongly in the ISO spectral range. Their PDRs are sites of active chemistry, allowing a number of transient species to be produced and detected, in addition to the well known IR lines of CO and H2. In addition to the 11 rotational lines of CO, Liu et al. (1996) reported the detection of the 119.3 μm fundamental rotational line of OH from the PDR around NGC 7027, in a high signal to noise LWS spectrum. They also detected an emission line at 179.6 μm, which they identified with the fundamental rotational line of o-H2O. However, a re-analysis of the data by Cernicharo et al. (1997) led to the re-allocation of this line to the J = 2–1 rotational transition of CH+, supported by the detection in the same LWS spectrum of the 3–2, 4–3, 5–4 and 6–5 transitions of CH+, at 119.90, 90.03, 72.140 and 60.247 μm, the first detection of the rotational spectrum of CH+. Liu et al. (1997) reported the detection in the NGC 7027 LWS spectrum of emission lines at 149.18 and 180.7 μm which matched the wavelengths of CH rotational lines and represented the first detection LATE STAGES OF STELLAR EVOLUTION 233 of the far-IR lines of this radical. Detailed models for the PDR around NGC 7027 have been constructed by Yan et al. (1999) and by Hasegawa et al. (2000), and have provided good matches to the observed LWS rotational line fluxes from CO, OH, CH and CH+. Liu et al. (2001) determined the temperatures and densities of the PDRs around 24 PNe, from the observed line intensity ratios of [C II] 158 μm and [O I]63and 146 μm, with the deduced temperatures typically falling between 200 and 500 K and the densities ranging from 104 to 3×105 cm−3. With a temperature of 1600 K, NGC 7027 had one of the warmest, as well as one of the densest, PDRs studied. NGC 7027 is an example of a young, high-density and incompletely ionised PN. Turning to the other extreme of PN evolution, the old, low-density PN NGC 7293 (the Helix Nebula) was imaged by Cox et al. (1998) in the ISOCAM 6.9 and 15 μm filters. Their complementary ISOCAM CVF spectra revealed that the emission in the 6.9 μm image was completely dominated by the S(5) line of H2 and that no dust continuum or PAH UIR-band emission was detectable. The H2 line intensity distribution could be fitted by a rotational temperature of 900 K and column densities 18 −2 of ∼3×10 cm , while the total luminosity in H2 lines amounted to 6% of the central star luminosity, much higher than predicted for PDRs. The H2 emission appears to originate from the many neutral globules in the Helix Nebula that have been imaged at optical wavelengths and in mm-wave CO lines. Persi et al. (1999) have presented ISOCAM images and CVF spectrophotometry of a further six PNe, ranging from high-density to low-density objects.

4.3. THE IONISED REGIONS OF PN

Liu et al. (2001) presented ISO LWS [O III] 52 and 88 μm, [N III]57μm and [N II] 122 μm fluxes for a sample of 51 PNe. The electron densities derived from the [O III] 52/88 μm flux ratios were found to be systematically lower than those obtained for the same nebulae from higher critical density optical diagnostic line ratios, consistent with the presence of significant density variations within the nebulae, − with the [O III] 52 and 88 μm lines (critical densities of 3500 and 1500 cm 3, respectively) being quenched in the higher density nebular zones. However, Tsamis et al. (2004) have shown that for those nebulae for which the IR and optical density diagnostic ratios all indicate electron densities below the critical densities of the [O III] IR lines, the electron temperatures derived from the [O III] λ5007/λ4363 ratio and from the λ5007/(52 μm + 88 μm) ratio agree well with each other, ruling out the presence of significant temperature fluctuations in these nebulae, since the IR lines of [O III] have much lower excitation energies than do the optical lines such as λ4363. The SWS spectral region contains many ionic FS lines and a number of papers have been published presenting FS line fluxes and nebular abundance analyses for 18 different PNe (Beintema and Pottasch, 1999; Pottasch and Beintema, 1999; Pottasch 234 J. A. D. L. BLOMMAERT ET AL. et al., 2000, 2001, 2002, 2003a,b, 2004; Bernard-Salas et al., 2001, 2002, 2003; Ercolano et al., 2003; Surendiranath et al., 2004). These analyses take advantage of the low excitation energies of IR FS lines and thus the relative insensitivity of the derived ionic abundances to the adopted nebular electron temperature. If the particular FS line has a critical density below a few ×104 cm−3, i.e. within the typical electron density range for PNe, then the derived ionic abundance is quite sensitive to the adopted electron density, but many FS lines in the SWS domain have critical densities above this range and so make good abundance diagnostics. Marigo et al. (2003) have compared the abundances derived in many of the earlier- referenced papers to synthetic evolutionary models for the AGB phase, confirming the presence of ‘two groups of PNe, one indicating the occurrence of only the third dredge-up during the TP-AGB phase, and the other showing also the chemical signature of hot-bottom burning’. van Hoof et al. (2000) have used ISO SWS measurements of [Ne V] infrared FS lines to re-assess the accuracy of published collision strengths for the Ne4+ ion, concluding that the R-matrix calculations of Lennon and Burke (1994), which had appeared to conflict with some earlier observations, were accurate to at least 30%. Rubin et al. (2002) discussed the observed ratios of a number of density sensitive FS line ratios falling in the SWS domain and showed that several of the measured ratios for PNe were outside the range predicted using current atomic data (including the [Ne V] infrared FS lines discussed by van Hoof et al.), implying the need for improved atomic collision strengths. Feuchtgruber et al. (1997) used ISO SWS grating and Fabry–Pérotobservations of several PNe to measure improved rest wavelengths for 29 infrared FS lines, while Feuchtgruber et al. (2001) provided similar data for a further six FS lines.

5. Observations of Novae

Many of the same ionic FS lines observed in planetary nebulae can also be observed in the spectra of novae, and ISO spectroscopy of these lines has been able to reveal a great deal of information about the ionised gas component of nova shells. Salama et al. (1996) reported SWS and LWS observations of Nova V1974 Cyg 1992, obtained 1494 days after the outburst. From measured FS line flux ratios, they estimated electron temperatures in the ejecta and derived an enhanced Ne/O ratio of ∼4. Salama et al. (1997) presented further ISO SWS observations of Nova V1974 Cyg 1992, including detailed line profiles of the [Ne V] 14.32 and 24.32 μm lines at four different epochs. The density sensitive [Ne V] line ratio was found to decrease with time t as t−3, in accord with ejecta expansion expectations. Salama et al. (1997) also presented ISO spectra of Nova HR Del 1967 and Nova V705 Cas 1993. Further ISO observations of the latter nova, obtained between 950 and 1455 days after outburst, were presented by Salama et al. (1999). The [O IV] line at 25.89 μm was found to be strongly in emission. Upper limits to the fluxes of neon LATE STAGES OF STELLAR EVOLUTION 235

FS lines led to an upper limit of 0.5 solar for the Ne/O number ratio in the nova shell. As well as obtaining observations of older novae, a nova Target of Opportunity programme was carried out during the 1995–1998 duration of the ISO mission, which was combined with co-ordinated complementary ground-based optical and near-IR spectroscopic observations. The ToO programme was unlucky in that no bright dust-forming novae occurred during the mission. However, the novae that were observed produced a rich harvest of information on second and third row elemental abundance distributions with which to compare with explosive nucleosynthesis predictions. Lyke et al. (2001) presented ISO, IUE and ground-based data on Nova V1425 Aql 1995 and combined these with a detailed photoionization model for the evolving nova shell. The mass of the shell was derived to be −5 2.5–4.3×10 M, with the modelling indicating that C and O were enhanced by a factor of 9 with respect to solar, while N was enhanced by a factor of 100 and Ne was only slightly enhanced. The presence of a CO white dwarf in the system was confirmed and a distance estimate of 3.0 ± 0.4 kpc to the nova was obtained. From their SWS IR and AAT optical and near-IR spectra of the heavily reddened Nova CP Cru 1996, Lyke et al. (2003) derived a distance of 2.6 ± 0.5 kpc to the system. Relative to solar, they derived abundance enhancements for N, O and Ne of 75, 17 and 27, with Mg showing an approximately solar abundance. The enhanced neon abundance but relatively low Mg abundance led them to interpret this object as a ‘missing link’ between CO novae and ONeMg novae. Nova V723 Cas 1995 was observed at six different epochs by ISO, from 217 to 805 days after the outburst, showing prominent highly ionised coronal FS line emission (Evans et al., 2003). From the flux ratios of hydrogen recombination lines having similar wavelengths, they derived reddening insensitive electron temperature −3 8 and density estimates, finding Ne (cm ) ∼ 2 × 10 [250/t (days)], where t was the time since outburst. For a distance of 4 kpc, they estimated an ejected shell mass of −5 −4 between 2.6×10 and 4.3×10 M. No lines of neon were detected, excluding V723 Cas from the class of ONeMg novae and implying that it has a CO white dwarf core.

6. Symbiotic Stars

Symbiotic stars are long-period binaries that generically contain a cool red giant star and a hot, evolved object that produces ionisation in the circumstellar outflow. The hot component of the vast majority of symbiotic systems is typically a hot white dwarf orbiting close enough to the red giant that it can accrete material from its wind. There are two distinct classes of symbiotic systems: the S-type (stellar) with normal red giants and orbital periods of about 1–15 years, and D-type (dusty) with Mira primaries usually surrounded by a dust shell, and orbital periods longer than about 10 years (Nussbaumer and Stencel, 1988). Symbiotic stars are interacting 236 J. A. D. L. BLOMMAERT ET AL. binaries with the longest orbital periods, and their study is essential to understand the evolution and interaction of binary systems. IRAS measurements of the brighter symbiotic stars (most symbiotics were only detected with IRAS as mili-Jy sources), indicated the presence of dust shells with characteristic temperatures of 250–800 K and dimensions of tens of AU (Anandarao et al., 1988). These warm shells are surrounded by more extended ones that were detected recently in sub-mm wavelengths ((Mikolajewska et al., 2002), and references therein). Some symbiotic systems were observed with ISO and the results from the published observations are discussed later. The majority of symbiotics were however too faint for ISO and are in fact ideally suited for the Spitzer Space Telescope. Preliminary results from ISO observations of symbiotic stars were presented by Eyres et al. (1998b). ISO SWS observations of the symbiotic novae (or very slow novae) RR Tel and V1016 Cyg reveal prominent, broad 10 and 18 μm silicate dust features. In both cases, the 10 μm feature is well fitted by the crystalline silicate feature found in novae Her 91. There is some observational evidence that the silicate feature at 10 μm show time-dependent variations. In addition, SWS spectra show a number of narrow emission lines, which can be used to constrain abundances. ISOPHOT observations were made for the S-type symbiotics AG Dra and AG Peg, providing multi-filter photometry and spectrophotometry. The observed emission is dominated by the giant-component continuum. There are some additional features: a broad plateau between 2.47 and 3.0 μm; a broad feature around 3.2 μm (AG Dra) and at 3.8 μm (AG Peg); a clear excess between 10 and 15 μm in both cases. Note, however, that some of these features could be of instrumental origin. Schild et al. (1998) discussed spectral features in SWS and LWS spectra of CH Cyg. Much of the ISO wavelength range is dominated by the emission of silicate dust with almost no nebular emission lines (at the time of observations the symbiotic activity of CH Cygni was small). In addition, they find strong OH and weak H2O emission between 60 and 130 μm. Note that OH molecules are absent in spectra of single giants of similar spectral type. Apart from the well known photospheric absorptions of CO (the band heads of 12CO and 13CO are well discernible), OH and SiO, they also tentatively detected traces of HCl and relatively weak PAH features at 6.3 and 11.3 μm. The authors suggested that carbon-rich material may be ejected from the outbursting companion. Schild et al. (2001) presented and discussed SWS and LWS observations of the symbiotic star HM Sge. The emission from this star is dominated by silicate dust with a number of nebular emission lines superimposed on the dust continuum. In the short wavelengths, there are no molecular absorption features that suggest the presence of warm dust. Analysis of the overall spectral energy distribution of HM Sge allowed to suggest that in addition to a Mira dust shell there is a second dust component associated with the hot source. This second shell is located at the interface between the Mira dust shell and the dust-free region carved out by the hot companion. The dust envelope associated with the Mira is very extended and LATE STAGES OF STELLAR EVOLUTION 237 optically thin, whereas the second dust component is geometrically rather thin but dust rich. The general agreement of the presented model with ISO observations suggest that radiation (not collisions) is the main dust heating mechanism. There are still ISO spectra available (but probably of lower quality than the one discussed earlier) for some symbiotic stars which show many nebular emission lines and can be used together with optical data to better constrain chemical abundances in ionised parts of these systems.

7. Wolf-Rayet Stars

Wolf-Rayet stars represent one of the final stages of stellar evolution for massive stars. In general, they are thought to follow an evolutionary sequence divided into three phases: WN, WC and WO, which correspond to dominant emission lines in their spectra (see van der Hucht, 2001 for a review). These emission lines are formed −6 −4 in the stars’ hot winds (typical mass-loss rates are between 10 and 10 M per year) and have velocities in the range between 1000 and 2500 km/s. The WR phase is predicted to be brief, lasting only a few 105 years before ending in Type II supernova explosions. Despite their harsh environment, some WC-type WR stars are known to be surrounded by circumstellar dust (e.g. Williams, 1995). Therefore, ISO spectroscopic observations were suitable for both: to study chemical composition of ejected and ionised gas, as well as dust spectral features. Preliminary results from ISO SWS observations of seven WR stars were presented by van der Hucht et al. (1996). They discussed Ne and He abundances in WR 11 on the basis of SWS 04 spectra, and dust features in WC 8–9 stars: WR 48a, WR 98a, WR 104, WR 112 and WR 118 from medium resolution SWS 02 observations. Their target stars, except for WR 11, are heavily reddened and SWS spectra show absorptions by silicates and hydrocarbons, possibly of interstellar origin. Still preliminary results were discussed in more detail during “ISO’s view on stellar evolution” conference held in Noordwijkerhout in July, 1997 (Morris et al., 1998; Willis et al., 1998; Williams et al., 1998). One of the problems considered there was the neon abundance. Stellar evolution models predict an enormous enhancement in surface Ne abundances in WC stars, making it the fourth most abundant behind He, C and O in the WC stars. Before the launch of ISO, observations suggested that Ne/H is enhanced by a factor of 2 rather than >10 as models predicted (Barlow et al., 1988). This conflict was a major puzzle in WR research. SWS observations of several galactic WR stars have been obtained to provide tests of this controversy. Analysis of the SWS spectra showed that the neon abundance is rather close to that predicted by stellar evolution models in the cases of WR 146 and γ 2 Vel (WR 11), but rather lower than predicted in the case of WR 135. Surpris- ingly, an overabundance of Ne was found in a WN8 star, WR 147, and rotational mixing has been proposed as a plausible explanation. The WN8(h) + B0.5 V binary system WR 147 has been further analysed by Morris et al. (2000). Using SWS data, 238 J. A. D. L. BLOMMAERT ET AL. they were able to put constraints on the surface abundances of Ca, S and Ne which were in a good agreement with the solar values of these elements for the H-depleted environment of WR 147. A detailed analysis of the Ne abundance problem in WC stars, based on the SWS and ground-based data, was presented further by Willis et al. (1997) for WR 146 and improved by Dessart et al. (2000), who considered also other WC stars: WR 11, WR 90 and WR 135. They estimated that Ne/He = 3–4×10−3 in the case of WR 11 and put an upper limit of the same order on the three remaining WC stars. Neon is highly enriched, but the observed Ne/He ratios are still a factor of 2 lower than the predictions of current evolutionary models for massive stars (see e.g. Maeder, 1991; Meynet, 1999). They also suggested that an imprecise mass loss and distance were responsible for the much greater discrepancy in neon content identified by Barlow et al. (1988). Note that Smith and Houck (2001) have obtained ground- based mid-infrared (8–13 μm) spectra for a sample of galactic WR stars which, except for an unresolved discrepancy for WR 146, show agreement with the SWS observations. It is worth mentioning also the possibility to determine terminal speeds of winds in three single WN stars using [Ca IV] 3.207 μm line widths from SWS data for these stars (Ignace et al., 2001). A detailed analysis of these line profiles allowed them also to conclude that the effect of turbulence in the winds of single WN stars is significant. They also showed that the line profile shapes in WR binaries appear asymmetric. In addition, using SWS data, Ignace et al. (2003) were able to put constraints on the velocity law and clumpiness of the wind from the WN-type star WR 136. As far as dust in H-poor WC stars is concerned, there are two principal questions: how dust is formed near such hot stars and the apparent failure to observe in WR stars any of the expected precursors (such as C-chain molecules) of the C-rich dust. SWS observations of circumstellar dust around WR stars reveal a variety of spectral energy distributions (SEDs). The mid-infrared SEDs of circumstellar shells around late-WC stars are approximately Planckian, with strong interstellar absorption features (amorphous silicates, CO2-ice at 4.27 μm, and the 6.2 μm feature to which molecules with the C double bond may contribute). The results of SED modelling by Williams et al. (1998) allowed them to examine the contributions made by different dust grain types to the observed emission, as well as the relative contributions made by circumstellar and interstellar components to the observed reddening. The authors demonstrated that AC-type amorphous carbon (Colangeli et al., 1995) gave fairly good fits to the observed SEDs. From pre-ISO observations with the KAO of two dusty WC9 stars, Cohen et al. (1989) had reported the presence of the 7.7 μm but the absence of the 6.2 μm unidentified infrared features (frequently attributed to C C stretch modes of PAHs), while Chiar et al. (2002) reported the detection of features at ∼6.4 and 7.9 μmin the ISO SWS spectrum of the dusty WC8 WR star WR 48a. Since the atmospheres of these stars are H poor, the detected features are believed to be due to large LATE STAGES OF STELLAR EVOLUTION 239 hydrogen-poor carbonaceous molecules or amorphous carbon dust grains and not PAHs. Interestingly, the features in the spectrum of WR 48a resemble the features observed in the spectra of H-deficient planetary nebulae. These similarities point towards a similar origin for the dust in these H-deficient environments and highlight the apparent sensitivity of the bands to local physical conditions.

References

Alard, C., Blommaert, J. A. D. L., Cesarsky, C., et al. (ISOGAL and MACHO Collaborations), 2001, ApJ 552, 289. Anandarao, B. G., Taylor, A. R., and Pottasch, S. R.: 1988, A&A 203, 361. Aoki, W., Tsuji, T., and Ohnaka, K.: 1998, A&A 340, 222. Aoki, W., Tsuji, T., and Ohnaka, K.: 1999, A&A 350, 945. Barlow, M. J., Nguyen-Q-Rieu, Truong-Bach, Cernicharo, J., Gonzalez-Alfonso, E., et al.: 1996, A&A 315, L241. Barlow, M. J., Roche, P., and Aitken, D. K.: 1988, MNRAS 232, 821. Barlow, M. J.: 1997, ApSS 251, 15. Beintema, D. A., van den Ancker, M. E., Molster, F. J., et al.: 1996, A&A 315, L369. Beintema, D. A. and Pottasch, S. R.: 1999, A&A 347, 942. Bernard Salas, J., Pottasch, S. R., Beintema, D. A., and Wesselius, P. R.: 2001, A&A 367, 949. Bernard Salas, J., Pottasch, S. R., Feibelman, W. A., and Wesselius, P. R.: 2002, A&A 387, 301. Bernard Salas, J., Pottasch, S. R., Wesselius, P. R., and Feibelman, W. A.: 2003, A&A 406, 165. Blommaert, J. A. D. L., Trams, N. R., Groenewegen, M. A. T., et al.: 1999, in: Cox, P. and Kessler, M. F. (eds.), The Universe as Seen by ISO. ESA-SP 427, p. 273. Blommaert, J. A. D. L., Omont, A., and The ISOGAL Collaboration: 2000, Mem. Soc. Astron. Ital. 71, 623. Blommaert, J. A. D. L.: 2003, ASSL Vol. 283: Mass-Losing Pulsating Stars and their Circumstellar Matter, p. 355. Bowey, J. E., Lee, C., Tucker, C., Hofmeister, A. M., Ade, P. A. R., and Barlow, M. J.: 2001, MNRAS 325, 886. Bowey, J. E., Barlow, M. J., Molster, F. J., et al.: 2002, MNRAS 331, L1. Cami, J. and Yamamura, I.: 2001, A&A 367, L1. Cami, J.: 2002, PhD Thesis, University of Amsterdam. Cami, J.: 2003, in: ASSL Vol. 283: Mass-Losing Pulsating Stars and their Circumstellar Matter,p. 209. Cami, J., Yamamura, I., de Jong, T., Tielens, A. G. G. M., Justtanont, K., and Waters, L. B. F. M.: 2000, A&A 360, 562. Castro-Carrizo, A., Bujarrabal, V., Fong, D., et al.: 2001, A&A 367, 674. Cernicharo, J., Liu, X.-W., Gonzalez-Alfonso, E., et al.: 1997, ApJ 483, L65. Cernicharo, J., Heras, A. M., Pardo, J. R., et al.: 2001a, ApJ 546, L127. Cernicharo, J., Heras, A. M., Tielens, A. G. G. M., et al.: 2001b, ApJ 546, L123. Cernicharo, J., Goicoechea, J. R., and Benilan, Y.: 2002, ApJ 580, L157. Chiar, J. E., Peeters, E., and Tielens, A. G. G. M.: 2002, ApJ 579, L91. Cioni, M.-R. L., Blommaert, J. A. D. L., Groenewegen, M. A. T., et al.: 2003, A&A 406, 51. Cohen, M., Tielens, A. G. G. M., and Bregman, J. D.: 1989, ApJ 344, L13. Cohen, M., Barlow, M. J., Sylvester, R. J., et al.: 1999, ApJ 513, L135. Cohen, M., Barlow, M. J., Liu, X.-W., and Jones, A.: 2002, MNRAS 332, 879. Cox, P., Boulanger, F., Huggins, P. J., et al.: 1998, ApJ 495, L23. 240 J. A. D. L. BLOMMAERT ET AL.

Colangeli, L., Menella, V., Palumbo, P., et al.: 1995, A&AS 113, 561. Decin, L., Vandenbussche, B., Waelkens, C., et al.: 2003, A&A 400, 709. De Marco, O., Barlow, M. J., and Cohen, M.: 2002, ApJ 574, L83. Dessart, Luc, Crowther, P. A., Hillier, D. J., et al.: 2000, MNRAS 315, 407. Dijkstra, C., Dominik, C., Hoogzaad, S. N., et al.: 2003a, A&A 401, 599. Dijkstra, C., Waters, L. B. F. M., Kemper, F., et al.: 2003b, A&A 399, 1037. Dominik, C., Dullemond, C. P., Cami, J., and van Winckel, H.: 2003, A&A 397, 595. Epchtein, N., de Batz, B., Capoani, L., et al.: 1997, Messenger 87, 27. Ercolano, B., Morisset, C., Barlow, M. J., Storey, P. J., and Liu, X.-W.: 2003, MNRAS 340, 1153. Evans, A., Gehrz, R. D., Geballe, T. R., et al.: 2003, AJ 126, 1981. Eyres, S. P. S., Evans, A., Geballe, T. R., Salama, A., and Smalley, B.: 1998a, MNRAS 298, L37. Eyres, S. P. S., Evans, A., Salama, A., et al.: 1998b, Astrophys. Space Sci. 255, 361. Fabian, D., Posch, T., Mutschke, H., Kerschbaum, F., and Dorschner, J.: 2001, A&A 373, 1125. Feuchtgruber, H., Lutz, D., Beintema, D. A., et al.: 1997, ApJ 487, 962. Feuchtgruber, H., Lutz, D., and Beintema, D. A.: 2001, ApJS 136, 221. Fong, D., Meixner, M., Castro-Carrizo, A., et al.: 2001, A&A 367, 674. Gauba, G. and Parthasarathy, M.: 2004, A&A 417, 201. Gautschy-Loidl, R., Höfner, S., Jørgensen, U. G., and Hron, J.: 2004, A&A 422, 289. Gilra, D. P.: 1973, in: IAU Symposium 52: Interstellar Dust and Related Topics, p. 517. Glass, I. S., Ganesh, S., Alard, C., et al.: 1999, MNRAS 308, 127. Glass, I. S. and Schultheis, M.: 2002, MNRAS 337, 519. Harrington, J. P., Lame, N. J., Borkowski, K. J., Bregman, J. D., and Tsvetanov, Z. I.: 1998, ApJ 501, L123. Harris, G. J., Pavlenko, Y. V., Jones, H. R. A., and Tennyson, J.: 2003, MNRAS 344, 1107. Hasegawa, T., Volk, K., and Kwok, S.: 2000, ApJ 532, 994. Herpin, F. and Cernicharo, J.: 2000, ApJ 530, L129. Herpin, F., Goicoechea, J. R., Pardo, J.R., and Cernicharo, J.: 2002, ApJ 577, 961. Höfner, S.: 1999, A&A 346, L9. Hony, S., Bouwman, J., Keller, L. P., and Waters, L. B. F. M.: 2002b, A&A 393, L103. Hony, S., Van Kerckhoven, C., and Tielens, A. G. G. M., et al.: 2001a, A&A 378, L41. Hony, S., Waters, L. B. F. M., and Tielens, A. G. G. M.: 2001b, A&A 378, L41. Hony, S., Waters, L. B. F. M., and Tielens, A. G. G. M.: 2002a, A&A 390, 533. Hony, S., Tielens, A. G. G. M., Waters, L. B. F. M., and de Koter, A.: 2003, A&A 402, 211. Hoogzaad, S. N., Molster, F. J., Dominik, C., et al.: 2002, A&A 389, 547. Hrivnak, B., Volk, K., and Kwok, S.: 2000, ApJ 535, 275. Hron, J., Loidl, R., Höfner, S., Jørgensen, U. G., Aringer, B., Kerschbaum, F., et al.: 1999, in: IAU Symposium 191: Asymptotic Giant Branch Stars, p. 181. Ignace, R., Cassinelli, J. P., Quigley, M., and Babler, B.: 2001, ApJ 558, 771. Ignace, R., Quigley, M., and Cassinelli, J. P.: 2003, ApJ 596, 538. Jones, H. R. A., Pavlenko, Y., Viti, S., and Tennyson, J.: 2002, MNRAS 330, 675. Jørgensen, U. G., Hron, J., and Loidl, R.: 2000, A&A 356, 253. Jørgensen, U. G., Jensen, P., Sørensen, G. O., and Aringer, B.: 2001, A&A 372, 249. Justtanont, K., Barlow, M., Tielens, A. G. G. M., et al.: 2000, A&A 360, 1117. Justtanont, K., de Jong, T., Tielens, A. G. G. M., Feuchtgruber, H., and Waters, L. B. F. M.: 2004, A&A 417, 625. Justtanont, K., Feuchtgruber, H., de Jong, T., Cami, J., Waters, L. B. F. M., Yamamura, I., et al.: 1998, A&A 330, L17. Kastner, J. H., Li, J.-Q., Siebenmorgen, R., and Weintraub, D. A.: 2000, AJ 123, 2658. Kemper, F., de Koter, A., Waters, L. B. F. M., Bouwman, J., and Tielens, A. G. G. M.: 2002a, A&A 384, 585. LATE STAGES OF STELLAR EVOLUTION 241

Kemper, F., Jager, C., Waters, L. B. F. M., et al.: 2002b, Nature 415, 295. Kerber, F., Blommaert, J. A. D. L., Groenewegen, M. A. T., Kimeswenger, S., Käuﬂ, H. U., and Asplund, M.: 1999, A&A 350, L27. Kwok, S., Volk, K., and Hrivnak, B.: 1989, ApJ 345, L51. Kwok, S., Volk, K., and Hrivnak, B.: 1999, A&A 350, L35. Kwok, S., Volk, S., and Bernath, P.: 2001, ApJ 554, L87. Lambert, D. L., Rao, K. N., Pandey, G., and Ivans, I. I.: 2001, ApJ 555, 925. Lamers, H., Waters, L. B. F. M., Garmany, C., et al.: 1986, A&A 154, L20. Lennon, D. J. and Burke, V. M.: 1994, A&AS 103, 273. Liu, X.-W., Barlow, M. J., Nguyen-Q-Rieu, et al.: 1996, A&A 315, L257. Liu, X.-W., Barlow, M. J., Dalgarno, A., et al.: 1997, MNRAS 290, L71. Liu, X.-W., Barlow, M. J., Cohen, M., et al.: 2001, MNRAS 323, 343. Lloyd Evans, T.: 1976, MNRAS 174, 169. Loup, C., Cioni, M. R., Blommaert, J. A. D. L., et al.: 1999, in: Cox, P. and Kessler, M. F. (eds.), The Universe as Seen by ISO. ESA-SP 427, p. 369. Lyke, J. E., Gehrz, R. D., Woodward, C. E., et al.: 2001, AJ 122, 3305. Lyke, J. E., Koenig, X. P., Barlow, M. J., et al.: 2003, AJ 126, 993. Maeder, A.: 1991, A&A 242, 93; A&A 315, L241. Marigo, P., Bernard-Salas, J., Pottasch, S. R., Tielens, A. G. G. M., and Wesselius, P. R.: 2003, A&A 409, 619. Markwick, A. and Millar, T.: 2000, A&A 359, 1162. Matsuura, M., Yamamura, I., Cami, J., Onaka, T., and Murakami, H.: 2002a, A&A 383, 972. Matsuura, M., Yamamura, I., Zijlstra, A. A., and Bedding, T.R.: 2002b, A&A 387, 1022. Meynet, G.: 1999, in: van der Hucht, K. A., Koenigsberger, G., and Eenens, P. R. J. (eds.), Wolf-Rayet Phenomena in Massive Stars and Starburst Galaxies, IAU Symposium 193, ASP, San Francisco, p. 218. Mikolajewska, J., Ivison, R. J., and Omont, A.: 2002, Adv. Space Res. 30, 2045. Molster, F. J.: 2000, ‘Crystalline Silicates in Circumstellar Dust Shells’, PhD Thesis, University of Amsterdam. Molster, F. J., Lim, T. L., Sylvester, R. J., et al.: 2001, A&A 372, 165. Molster, F. J., Waters, L. B. F. M., Tielens, A. G. G. M., and Barlow, M. J.: 2002a, A&A 382, 184. Molster, F. J., Waters, L. B. F. M., and Tielens, A. G. G. M.: 2002b, A&A 382, 222. Molster, F. J., Waters, L. B. F. M., Tielens, A. G. G. M. et al.: 2002c, A&A 382, 241. Molster, F. J., Yamamura, I., Waters, L. B. F. M., et al.: 1999, Nature 401, 563. Morris, P. W., van der Hucht, K. A., Crowther, P. A., et al.: 2000, A&A 353, 624. Morris, P. W., van der Hucht, K. A., Willis, A. J., and Williams, P. M.: 1998, in: Waters, L. B. F. M., Waelkens, C., van der Hucht, K. A., and Zaal, P. A. (eds.), ISO’s View on Stellar Evolution, Astrophysics and Space Science 255, Kluwer Academic Publishers, Dordrecht/Boston/Leiden, p. 157. Neufeld, D. A., Chen, W., Melnick, G. J., de Graauw, T., Feuchtgruber, H., Haser, L., et al.: 1996, A&A 315, L237. Nussbaumer, H. and Stencel, R.: 1987, in: Exploring the Universe with the IUE Satellite (A88-22626 07-90), D. Reidel Publishing Co., Dordrecht, p. 203. Ojha, D. K., Omont, A, Schuller, F., et al.: 2003, A&A 403, p. 141. Omont, A., Moseley, S. H., Cox, P., et al.: 1995, ApJ 454, 819. Omont, A., Ganesh, S., Alard, C., Blommaert, J. A. D. L., Caillaud, B., et al.: 1999, A&A 348, 755. Omont, A., Gilmore, G. F., Alard, C., et al.: 2003, A&A 403, 975. Onaka, T., de Jong, T., and Willems, F. J.: 1989, A&A 218 169. Ortiz, R., Blommaert, J. A. D. L., Copet, E., et al.: 2002, A&A 388, 279. Osterbart, R., Langer, N., and Weigelt, G.: 1997, A&A 325, 609. 242 J. A. D. L. BLOMMAERT ET AL.

Peeters, E., Hony, S., Van Kerckhoven, C., et al.: 2002, A&A 390, 1089. Pei, A. and Volk, K.: 2003, in: Kwok, S., Dopita, M., and Sutherland, R. (eds.), Planetary Nebulae: Their Evolution and Role in the Universe, Astronomical Society of the Pacific. Persi, P., Cesarsky, D., Marenzi, A. R., et al.: 1999, A&A 351, 201. Posch, T., Kerschbaum, F., Mutschke, H., Dorschner, J., and Jäger,C.: 2002, A&A 393, L7. Posch, T., Kerschbaum, F., Mutschke, H., Fabian, D., Dorschner, J., and Hron, J.: 1999, A&A 352, 609. Pottasch, S. R. and Beintema, D. A.: 1999, A&A 347, 975. Pottasch, S. R., Beintema, D. A., and Feibelman, W. A.: 2000, A&A 363, 767. Pottasch, S. R., Beintema, D. A., Bernard Salas, J., and Feibelman, W. A.: 2001, A&A 380, 684. Pottasch, S. R., Beintema, D. A., Bernard Salas, J., Koornneef, J., and Feibelman, W. A.: 2002, A&A 393, 285. Pottasch, S. R., Hyung, S., Aller, L. H., et al.: 2003a, A&A 401, 205. Pottasch, S. R., Bernard-Salas, J., Beintema, D. A., and Feibelman, W. A.: 2003b, A&A 409, 599. Pottasch, S. R., Bernard-Salas, J., Beintema, D. A., and Feibelman, W. A.: 2004, A&A 423, 593. Rubin, R. H., Dufour, R. J., Colgan, S. W. J., et al.: 2002, RMexAC 12, 106. Ryde, N., Eriksson, K., and Gustafsson, B.: 1999, A&A 341, 579. Sahai, R., Te Lintel Hekkert, P., Morris, M., et al.: 1999, ApJ 514, L115. Sahai, R., Zijlstra, A., SánczezContreras, C., and Morris, M.: 2003, ApJ 586, L81. Salama, A., Evans, A., Eyres, S. P. S., Leech, K., Barr, P., and Kessler, M. F.: 1996, A&A 315, L209. Salama, A., Barr, P., Clavel, J., et al.: 1997, ApSS 255, 227. Salama, A., Eyres, S. P. S., Evans, A., Geballe, T. R., and Rawlings, J. M. C.: 1999, MNRAS 304, L20. Schild, H., Dumm, T., Folini, D., et al.: 1998, in: Kesler, M. and Cox, P. (eds.), The Universe as Seen by ISO, ESA SP-427, p. 397. Schild, H., Eyres, A. P. S., Salama, A., and Evans, A.: 2001, A&A 378, 146. Sloan, G. C. and Price, S. D.: 1995, ApJ 451, 758. Sloan, G. C., Kraemer, K. E., Goebel, J. H., and Price, S. D.: 2003, ApJ 594, 483. Smith, J. D. T. and Houck, J. R.: 2001, ApJ 121, 2115. Speck, A., Barlow, M. J., Sylvester, R. J., and Hofmeister, A. M.: 2000a, A&AS 146, 437. Speck, A., Meixner, M., and Knapp, G. R.: 2000b, ApJ 545, L145. Speck, A. and Hofmeister, A. M.: 2004, ApJ 600, 986. Surendiranath, R., Pottasch, S. R., and Garc´ıa-Lario,P.: 2004, A&A 421, 1051. Sylvester, R. J., Kemper, F., Barlow, M. J., de Jong, T., Waters, L. B. F. M., Tielens, A. G. G. M., et al.: 1999, A&A 352, 587. Szczerba, R. and Górny, S. K. (eds.): 2001, Post-AGB Objects as a Phase of Stellar Evolution, Kluwer Academic Publishers, Dordrecht/Boston/London. Szczerba, R., Stasińska,G., Siódmiak,N., and Górny, S. K.: 2003, in: Gry, C., Peschke, S., Matagne, J., Garcia-Lario, P., Lorente, R., and Salama, A. (eds.), Exploiting the ISO Data Archive. Infrared Astronomy in the Internet Age, ESA-SP, Vol. 511, p. 149. Szczerba, R., Henning, Th., Volk, K., et al.: 1999, A&A 345, L39. Tanabé,T., Nishida, S., Nakada, Y., et al.: 1998, Astrophys. Space Sci. 255, 515. Tielens, A. G. G. M., Hony, S., Van Kerckhoven, C., and Peeters, E.: 1999, in: Cox, P. and Kessler, M. F. (eds.), The Universe as Seen by ISO, ESA-SP, Vol. 427, p. 579. Trams, N. R., van Loon, J. Th., Zijlstra, A. A., et al.: 1999a, A&A 344, L17. Trams, N. R., van Loon, J. Th., Waters, L. B. F. M., et al.: 1999b, A&A 346, 843. Truong-Bach, Sylvester, R. J., Barlow, M. J., Nguyen-Q-Rieu, Lim, T., Liu, X. W., et al.: 1999, A&A 345, 925. Tsamis, Y. G., Barlow, M. J., Liu, X.-W., Storey, P. J., and Danziger, I. J.: 2004, MNRAS 353, 953. Tsuji, T.: 2000, ApJ Lett. 540, L99. LATE STAGES OF STELLAR EVOLUTION 243

Tsuji, T.: 2001, A&A 376, L1. Tsuji, T., Ohnaka, K., Aoki, W., and Yamamura, I.: 1997, A&A 320, L1. Ueta, T., Meixner, M., and Bobrowsky, M.: 2000, ApJ 528, 861. Vandenbussche, B., Beintema, D., de Graauw, T., et al.: 2002, A&A 390, 1033. van der Hucht, K. A.: 2001, New Astron. Rev. 45, 135. van der Hucht, K. A., Morris, P. W., Williams, P. M., et al.: 1996, A&A 315, L193. van Hoof, P. A. M., Beintema, D. A., Verner, D. A., and Ferland, G. J.: 2000, A&A 354, L41. van Loon, J. Th., Groenewegen, M. A. T., de Koter, A., et al.: 1999, A&A 351, 559. van Loon, J. Th., Gilmore, G. F., Omont, A., Blommaert, J. A. D. L., Glass, I. S., et al.: 2003, MNRAS 338, 857. Van Malderen, R.: 2003, PhD Thesis, K.U. Leuven, Leuven, Belgium. Van Malderen, R., Decin, L., Kester, D., Vandenbussche, B., Waelkens, C., Cami, J., and Shipman, R. F.: 2004, A&A 414, 677. van Winckel, H.: 2003, ARAA 41, 391. van Winckel, H., Waelkens, C., and Waters, L. B. F. M.: 1995, A&A 293, L25. van Winckel, H., Waelkens, C., and Waters, L.B.F.M., et al.: 1998, A&A 336, L17. Vardya, M. S., de Jong, T., and Willems, F. J.: 1986, ApJ Lett. 304, L29. Volk, K., Kwok, S., and Hrivnak, B.: 1999, ApJ 516, L99. Volk, K., Kwok, S., Hrivnak, B., and Szczerba, R.: 2002, ApJ 567, 412. Volk, K., Xiong, G.-Z., and Kwok, S.: 2000, ApJ 530, 408. Waelkens, C. and Waters, L. B. F. M., 2004, in: Habing, H. J. and Olofsson, H. (eds.), Asymptotic Giant Branch Stars, Springer-Verlag, New York. Waters, L. B. F. M., Cami, J., de Jong, T., Molster, F. J., van Loon, J. T., Bouwman, J., et al.: 1998a, Nature 391, 868. Waters, L. B. F. M., Beintema, D. A., Zijlstra, A. A., et al.: 1998b, A&A 331, L61. Williams, P. M.: 1995, in: van der Hucht, P. A. and Williams, P. M. (eds.), Wolf-Rayet Stars: Binaries; Colliding Winds, IAU Symposium 163, Kluwer Academic Publishers, Dordrecht, p. 335. Williams, P. M., van der Hucht, K. A., and Morris, P. W.: 1998, in: Waters, L. B. F. M., Waelkens, C., van der Hucht, K. A., and Zaal, P.A. (eds.), ISO’s View on Stellar Evolution, Astrophys. Space Sci. 255, p. 169. Willis, P. M., Dessart, Luc, Crowther, P. A., et al.: 1997, MNRAS 290, 371. Willis, P. M., Dessart, Luc, Crowther, P. A., et al.: 1998, in: Waters, L. B. F. M., Waelkens, C., van der Hucht, K. A., and Zaal, P. A. (eds.), ISO’s View on Stellar Evolution, Astrophys. Space Sci. 255, p. 167. Yamamura, I., de Jong, T., and Cami, J.: 1999a, A&A 348, L55. Yamamura, I., de Jong, T., Onaka, T., Cami, J., and Waters, L. B. F. M.: 1999b, A&A 341, L9. Yamamura, I., de Jong, T., Waters, L. B. F. M., Cami, J., and Justtanont, K.: 1999c, in: IAU Symposium 191: Asymptotic Giant Branch Stars, p. 267. Yamamura, I., Dominik, C., de Jong, T., Waters, L. B. F. M., and Molster, F. J.: 2000, A&A 363, 629. Yan, M., Federman, S. R., Dalgarno, A., and Bjorkman, J. E.: 1999, ApJ 515, 640. Yang, X., Chen, P., and He, J.: 2004, A&A 414, 1049. Zubko, V. and Elitzur, M.: 2000, ApJ Lett. 544, L137. INTERSTELLAR MEDIUM THE COOL INTERSTELLAR MEDIUM

ALAIN ABERGEL1,∗, LAURENT VERSTRAETE1, CHRISTINE JOBLIN2, RENE´ LAUREIJS3 and MARC-ANTOINE MIVILLE-DESCHENESˆ 1,4 1IAS, Universite´ Paris-Sud, Bat.ˆ 121, 91405 Orsay, France 2CESR, CNRS-Universite´ Paul Sabatier, 9 avenue du Colonel Roche, 31028 Toulouse, France 3Astrophysics Division, Research and Scientiﬁc Support Department of ESA, ESTEC, PO Box 299, 2200 AG Noordwijk, The Netherlands 4CITA, 50 St-Georges street, Toronto, Ontario, M5S3H8, Canada (∗Author for correspondence: E-mail: [email protected])

(Received 21 September 2004; Accepted in ﬁnal form 8 November 2004)

Abstract. Infrared spectroscopy and photometry with ISO covering most of the emission range of the interstellar medium has led to important progress in the understanding of the physics and chemistry of the gas, the nature and evolution of the dust grains and also the coupling between the gas and the grains. We review here the ISO results on the cool and low-excitation regions of the interstellar 5 −3 −4 medium, where Tgas 500 K, nH ∼ 100–10 cm and the electron density is a few 10 .

Keywords: infrared: spectroscopy, photometry; ISM: gas, dust; ISM: PDR, cirrus, cold clouds; ISM: physical processes

JEL codes: D24, L60, 047

1. The Cool Interstellar Medium

The interstellar medium (ISM) ﬁlls the volume of the Galaxy and its evolution is mostly driven by the stellar activity. It is composed primarily of gas and of a small amount (1% in mass) of submicronic dust grains. In spite of their low abundance, dust particles play a key role for the temperature and chemistry of the gas. Thus, the evolution of the interstellar gas through its different phases (from proto-stellar cores to the diffuse medium and vice-versa) and of dust grains are closely related. Although shocks and winds play an important role in the vicinity of stars (Blommaert et al.; Nisini et al.; Peeters et al., this volume), the ISM is primarily heated by stellar radiation. It then cools off by continuum and line emission in the infrared (IR) and submillimetre (submm) range. Observations of the IR emission from the ISM hence provide important clues on the gas state, the nature of the dust particles, the interactions between the gas and the dust as well as the dynamical processes affecting the gas. Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, The Netherlands, and the United Kingdom), and with the participation of ISAS and NASA. Space Science Reviews (2005) 119: 247–271 DOI: 10.1007/s11214-005-8056-z C Springer 2005 248 A. ABERGEL ET AL.

In the course of its lifecycle, the ISM assumes a number of phases described by the temperature Tgas and proton density nH of the gas as well as the electron abundance xe. Most of the ISM mass resides in cool, mostly neutral phases – the cool 5 −3 or low-excitation ISM – where Tgas 500 K, nH ∼ 100–10 cm and xe isafew 10−4. The cool ISM encompasses photon-dominated regions or PDRs generated by the far-ultraviolet (FUV) stellar radiation ﬁeld, shielded regions that are cold and molecular (the cold clouds), low-mass star forming clouds, the outskirts of massive star forming regions and the general ISM (the diffuse ISM seen in cirrus clouds). This medium is therefore widespread in the Galaxy (some 10% in volume), and observed to be structured on all scales (a few 0.01 to a few 10 pc, e.g., Falgarone et al., 1998; Miville-Deschˆenes et al., 2003). The cool ISM is the subject of this review while the more highly excited phases of the ISM, the regions around YSOs and late-type stars are discussed in this volume by Peeters et al.; Nisini et al.; Lorenzetti; Blommaert et al., respectively. The ISO mission was particularly well suited for ISM studies because of the complementarity of its four instruments. The CAM and PHOT instruments provided sensitive broad-band IR images at angular resolution of a few arcseconds for ISOCAM to 1 arcmin for ISOPHOT. The former also had a spectro-imaging facility, the CAM-CVF mode operating between 6 and 16 μm with a resolving power R ∼ 40 while the latter was equipped with a grating spectrometer used in the PHOT-S mode (two channels: 2.5–4.9 μm and 5.8–11.6 μm with R ∼ 90 and a 24 × 24 arcsec aperture). On the other hand, the SWS and LWS spectrometers combined relatively higher spectral resolution (R ∼ 200–2000) and full spectral coverage (2.4–45.2 μm for the SWS and 43–197 μm for the LWS). The ISO instruments have detected a great wealth of gas lines and dust spectral features in emission or absorption. High spectral resolution measurements with SWS and LWS have been performed mostly towards moderately to highly excited regions (H II regions, planetary nebulae, PDRs and active galaxies), while less excited regions could only be observed by CAM and PHOT at low spectral resolution or in broad bands. Thus, ISO has provided a complete and unbiased IR spectroscopic census of interstellar matter highlighting the nature and properties of dust, the physical conditions of the gas as well as the major processes (UV stellar irradiation, shock waves) driving the evolution of these species. While the results on the H2 and H2O molecules (Habart et al.; Cernicharo and Crovisier) and the dust features of silicates and ices (Molster and Kemper; Dartois) are discussed in dedicated chapters of this volume, we focus here on the properties of gas and dust in the cool ISM. We also refer the reader to the recent review by van Dishoeck (2004).

2. The Physics and Chemistry of the Gas

Combined observations of gas lines and dust emission proved to be a unique tool to constrain the physics and chemistry of the gas in the cool ISM (Section 2.1) and to quantify the importance of the interactions between gas and dust. These latter THE COOL ISM 249 are most directly studied in low-excitation PDRs (Section 2.2) while gas dynamical processes are best traced in the less excited, diffuse medium (Section 2.3).

2.1. PROBING INTERSTELLAR CHEMISTRY WITH SPECTROSCOPY OF GAS-PHASE MOLECULES

The wavelength range of the ISO spectrometers includes rotational and ro- vibrational transitions of many light interstellar molecules. The ISO spectroscopy data thus provided an important benchmark for models of chemistry. Furthermore, analysis of ISO spectra with models pointed at the dominant excitation processes and highlighted the structure of the observed regions. A wealth of molecular lines has been found towards the galactic centre. As detailed by Ceccarelli et al. (2002) the Sgr B2 region is a massive molecular cloud containing numerous compact and ultra-compact H II regions and hot cores embedded in a relatively dense and warm envelope (∼100 K), while the entire com- 3 −3 plex is surrounded by a hot (≥300 K) envelope of diffuse gas (nH ∼ 10 cm ). Sgr B2 is therefore an ideal target to detect molecular absorption. The detection of new absorption features has strongly constrained the abundances and the physical conditions (gas temperature and density) of the absorbing regions. Numerous + rotational lines of hydrides (OH, CH, HF, H2O, H3O ,NH3,NH2, NH) and hydrocarbons (C3,C4 or C4H) have been detected in absorption by LWS (see Tables 4 and 5 of Goicoechea et al., 2004 for the complete list of lines and references). Light hydrides such as OH, CH and H2O were found at all positions, indicating the 0 widespread presence of these species in molecular clouds. Lines of the H2O/OH/O and NH3/NH2/NH sequences were found to trace the importance of photodissociation processes and shock respectively (Goicoechea et al., 2004). The HD molecule has been detected by the LWS (Wright et al., 1999; Caux et al., 2002; Polehampton et al., 2002). A deuterium abundance of D/H ∼0.2–11×10−6 has been derived for molecular clouds in agreement with other determinations in the Solar neighbourhood. The hydrogen fluoride HF has been discovered through its J = 2 − 1 band at 121.7 μm (Neufeld et al., 1997). HF is the dominant reservoir of gas-phase fluorine and, assuming that the absorption arises in the warm envelope, the derived abundance (∼3 × 10−10) indicates a strong depletion of fluorine onto dust grains. However, Ceccarelli et al. (2002) mentioned that this absorption could also be due to the hot absorbing layer, in which case the abundance would be an order of magnitude larger. This illustrates that the interpretation of such data is not straightforward, since several components with different physical properties can participate to the detected absorption, as examplified in the work of Vastel et al. (2002) for the C+ and O0 fine-structure lines. Original results have also been obtained on hydrocarbons and carbon chains. Cernicharo et al. (2002) have found a new line at 57.5 μm which may be the first detection of C4 (another possible carrier is C4H). The required abundance of C4 would make it the most abundant carbon chain in interstellar and circumstellar 250 A. ABERGEL ET AL. media. Other polar carbon chains such as cyanopolynes (HC2n+1N with n = 1 − 5) were already detected in the millimeter, but ISO has allowed the detection of pure carbon chains through their low-energy bending modes which fall in spectral regions totally absorbed by the Earth’s atmosphere. On the other hand, ro-vibrational emission lines of HCN and C2H2 have been found in warm (a few 100 K) molecular cores with abundances ∼10−7 and a likely scenario is that these lines are pumped by warm dust IR emission (Boonman et al., 2003). The methyl radical CH3 has been discovered in absorption toward Sgr A with SWS through its ν2 Q-branch at 16.5 μm and the R(0) line at 16.0 μm (Feuchtgruber et al., 2000). The relatively −8 high abundance of CH3 (∼10 ) together with the low abundances of other hydrocarbon molecules (CH, CH4 and C2H2) is inconsistent with the results of published pure gas-phase models of dense clouds. The authors suggest that models including diffuse and translucent clouds or translucent clouds with gas–grain chemistry may resolve this discrepancy.

2.2. THE COUPLING BETWEEN GAS AND DUST IN LOW-EXCITATION PDRS

Photo Dissociation Regions (PDRs) are ubiquitous in the Galaxy since any interstellar cloud exposed to the ambient radiation field will develop such an interface. More generally, PDRs dominate the IR emission of galaxies and are used to trace their star formation activity (e.g., Förster-Schreiber et al., 2004). Early observations of bright, highly excited PDRs (e.g., Harper et al., 1976; Melnick et al., 1979) have led to detailed modelling where the thermal and chemical balance of the gas is fully described. Physical conditions (radiation field intensity and gas density, see below) in PDRs were then obtained from the comparison between model predictions and observations (see the review in Hollenbach and Tielens, 1999). These studies have pointed out the importance of gas–dust interactions and the need to study them further. Among these interactions, the most important are the photoelectric effect on dust which heats the gas and the formation of the H2 molecule on the surface of grains which triggers the chemistry.1 The gas state thus results of a balance between photon-driven processes (photoelectric effect on dust, photodissociation of H2, ...) and density dependent reactions (electron recombination, H2 formation on grains, ...). Hence, the excitation of PDRs is usually described with χ, a scaling factor for the intensity of the radiation field in the FUV,2 and with the local gas proton density nH. Observations – including the ISO mission – have been primarily performed on nearby PDRs which can be spatially resolved. Among these regions those with an edge-on geometry represent ideal targets for the study of gas–grain interactions, dust evolution and chemical stratification as a function of depth (or UV field intensity) within the cloud.

1Gas–grain and gain-grain collisions are also important. These processes, however, primarily affect the dust (size distribution, structure and composition) and they are discussed in Section 3.4. − 2For χ = 1 the stellar intensity is 4πνIν = 1.2 × 10 6 W/m2 at λ ∼ 100 nm, i.e. the average interstellar value in the solar neighbourhood (Mathis et al., 1983). THE COOL ISM 251

PDRs have been extensively observed with the four instruments of ISO, spanning a wide range of physical conditions. Thanks to the sensitivity of the ISO instruments, the sample of observed PDRs has been extended towards low-excitation 4 −3 regimes (χ a few 1000 and nH 10 cm ) allowing a detailed study of important microscopical processes. In low-excitation PDRs, the gas thermal budget is amenable to a few dominant processes (Kemper et al., 1999; Habart et al., 2001): the heating is due to the photoelectric effect on dust and the cooling occurs through the fine-structure line emission of C+ (158 μm) and O0 (63 and 145 μm). These lines measured by ISO-LWS have been used to derive the efficiency of the photoelectric heating, assuming thermal balance (heating = cooling). This efficiency noted represents the fraction of the stellar energy absorbed by a given dust population that goes into gas heating. In practice, can be determined observationally by cor- relating the gas cooling line to the dust emission. In such a detailed study towards −3 L1721, an isolated cloud heated by a single B-star (χ ∼ 4 and nH ∼ 3000 cm ), Habart et al. (2001) have shown that the photoelectric heating is dominated by the smallest grains (of radii 1 nm called PAHs, see Section 3) with PAH ∼ 3%. These results have confirmed the theoretical treatment of the gas photoelectric heating (Bakes and Tielens, 1994; Weingartner and Draine, 2001a) used in current PDR −3 models. Towards high-latitude, diffuse clouds (χ ∼ 1 and nH ∼ 100 cm ), Ingalls et al. (2002) have found a statistical correlation between the C+ line and the far-IR emission of big grains (radii 10 nm) as measured with IRAS at 60 and 100 μm which corresponds to an efficiency of ∼4.3%. This latter value, about 1.5 times larger than the theoretical predictions of Weingartner and Draine (2001a), suggests that small grains (PAHs and VSGs, see Section 3.1) undergoing temperature fluctuations contribute significantly to the IR emission. Other studies made use of PDR models treating in detail the gas thermal balance to analyse observed cooling lines. The [C+] 158 μm and [O0]63μm dominant lines were successfully explained with the current treatment of photoelectric heating (Timmermann et al., 1998; Kemper et al., 1999; Liseau et al., 1999; Habart et al., 2001, 2003; Li et al., 2002; Okada et al., 2003; Schneider et al., 2003) confirming independently the photoelectric heating rate used in models. These studies have also shown that the major cooling lines are optically thick, even in low-density 3 −3 PDRs (nH 10 cm ). In particular, the source geometry and line transfer proved to be important for the O0 transitions (Li et al., 2002; Okada et al., 2003). In most cases, the O0 emission line ratio 63 μm/145 μm cannot be readily explained (Liseau et al., 1999) and recent work (Mizutani et al., 2004) has suggested absorption in the [C+] 158 μm and [O0]63μm lines and/or a quenched photoelectric heating as the possible cause of this discrepancy. Finally, comparison of spatial profiles of the C+ and O0 lines to model results has shown that the abundance of PAHs is strongly depressed (by up to a factor 5) in shielded regions (Habart et al., 2001, 2003). A comparable result has been found in the diffuse medium, traced by the emission of small grains as seen by CAM (Section 3.4). 252 A. ABERGEL ET AL.

More indirect analysis of PDRs has also been conducted using the [O0] 63 and 145 μm and [C+] 158 μm gas cooling lines detected in absorption and/or emission with LWS towards the Galatic centre, in Sgr B2 (Vastel et al., 2002; Goicoechea et al., 2004). These observations have allowed to disentangle the contribution of the three layers (C+/C+-C0/CO) predicted by PDR models. The intensities of the three 3 4 3 4 −3 lines indicate that χ ∼ 10 –10 and nH ∼ 10 –10 cm in the PDRs (Goicoechea et al., 2004). The [C+] 158 μm line is seen in absorption, conﬁrming that this main cooling line of the ISM can be optically thick. Moreover, Vastel et al. (2002) show that atomic oxygen is associated with the three layers. Less than 30% of the total O0 column density (∼ 3 × 1019 cm−2) comes from the external (C+) layers. On the other hand, ∼70% of gaseous oxygen is in the atomic form and not locked up in CO in the molecular clouds lying along the line of sight. This latter result is not explained by current PDR models. This may be related to the low H2O abundance detected by the SWAS satellite and which is not explained by current chemical 0 networks (Roberts and Herbst, 2002). More observations of O ,O2 and H2Oin molecular clouds are necessary to solve this issue. For the ﬁrst time, the pure rotational bands of H2 have been detected with SWS towards many PDRs providing a direct probe of the gas temperature. In low- excitation PDRs, current PDR models failed to account for this emission suggesting an enhanced H2 formation rate. Observations also suggest that small grains (radii 10 nm) play an important role in the formation mechanism. This issue is discussed in the review on molecular hydrogen by Habart et al. (this volume).

2.3. NON-THERMAL GAS MOTIONS

Non-thermal motions can play an important role in exciting and mixing the interstellar gas. Some of the underlying processes are directly related to stars (outflows, winds) and are discussed elsewhere in this volume (Nisini et al.; Peeters et al.). We discuss here processes that pervade most of the ISM, namely, motions resulting from magneto-hydrodynamical (MHD) turbulence and motions driven by the ambient radiation field. MHD turbulence has been studied by Falgarone et al. (1999, 2005) from SWS observations of the pure rotational lines of H2 in emission along a deep sightline in the galactic plane sampling the cold diffuse ISM (i.e., avoiding star forming region and molecular clouds). The detected H2 excitation has shown the presence of a small fraction (a few percents) of warm gas (a few 100 K). This confirms more directly earlier absorption measurements of molecules (CH+, OH, HCO+) that can only form in warm gas through endothermic reactions. As illustrated in Figure 1, collisional excitation by MHD turbulence (in shocks or vortices) can successfully account for the excited H2. In addition, the required number of turbulent structures can also explain the observed CH+ abundances. More recent studies with FUSE towards stars in the Solar neighbourhood find similar fractions of warm gas (Gry et al., 2002). This suggests that turbulent motions occur with similar rates THE COOL ISM 253

Figure 1. SWS Intensities (solid squares) of the pure H2 rotational lines in the cold diffuse medium. Ju is the upper rotational number of the transition. Left: The contribution of background PDRs on the line of sight (upper solid line) consists of: (1) 19 diffuse PDRs in the Solar neighbourhood (χ = 1, −3 nH = 30 cm and AV = 0.3, lower solid line), (2) 5 diffuse PDRs in the molecular ring (χ = 10, −3 nH = 100 cm and AV = 2, dashed line) and (3) 7 dense PDRs in the molecular ring (χ = 10, 4 −3 nH = 10 cm and AV = 2, dotted line). The fluorescence contribution of low-density PDRs is too low to account for the fluxes of the excited H2 rotational lines. Center: Residual H2 line intensities once the PDR-type emission (see left panel) has been removed (solid squares). Models of H2 lines emission from a few tens of MHD shocks (dotted lines) at 8 (lower), 10 and 12 km s−1 (upper). Right: Same residuals compared to the emission of about 104 magnetized coherent vortices of rotational velocity 4 km s−1(dot-dashed) and 3.5 km s−1(dashed). The dotted line releases the assumption of statistical balance of the H2 level populations. Taken from Falgarone et al. (2005). throughout the inner parts of the Galaxy. In the near future, the spectrometer of the Spitzer Space Telescope (SST) will look at the H2 excitation towards a wider variety of galactic sightlines as well as in other galaxies. Turbulent motions have also been traced indirectly in cirrus clouds through their impact on the dust size distribution (Section 3.4). In the warmer environment of PDRs, non-thermal motions are dwarfed by thermal processes and their signature is more allusive. For instance, a possible explanation of the low ortho-to-para ratio of H2 observed in the warm layers (a few 100 K) of several PDRs are advection motions due to the propagation of photodissociation fronts (Habart et al., this volume). However, a quantitative estimate of the efficiency of this process awaits a dedicated modelling effort.

3. The Nature and Evolution of Interstellar Dust Grains

Observations of IR features in emission or absorption provide important constraints on the nature, composition and properties (size distribution, structure) of the dust 254 A. ABERGEL ET AL. particles. Work prior to the ISO mission (e.g., Sellgren, 1984; Désert et al., 1990; Dwek et al., 1997) has shown that interstellar dust comprises a population of small grains (radii a 10 nm) that absorbs about half the energy radiated away by stars and re-emit this energy during temperature fluctuations. As discussed in Section 2.2, small grains play an important role in the physics and chemistry of the interstellar gas. These small grains are believed to be predominantly carbonaceous (e.g., Draine, 2003) and we will discuss their size in terms of the number of carbon atoms they contain, NC. The smallest of them (a 1nmandNC a few 100) are the carriers of a family of IR emission bands between 3 and 13 μm characteristic of aromatic rings (Duley and Williams, 1982) and we call these features the aromatic IR bands (AIBs). To explain these bands, a family of large aromatic molecules, the Polycyclic Aromatic Hydrocarbons (PAHs) has been proposed (Légerand Puget, 1984; Allamandola et al., 1985). As discussed below (Sections 3.1 and 3.2), the PAH model explains the main properties of the AIB observations but still lacks a precise spectroscopic identification with terrestrial analogues. Dust emission at mid-IR wavelengths (e.g., in the IRAS 25 and 60 μm bands) has been ascribed to somewhat larger carbonaceous grains (1 a 10 nm) called Very Small Grains (VSGs) (Désert et al., 1990). Finally, the far-IR dust emission (λ 30 μm) is ascribed to a mixture of large graphite and silicated grains or big grains (a ∼ afew10 to a few 100 nm, e.g., Li and Draine, 2001; Weingartner and Draine, 2001b) which emit while being in thermal equilibrium with the radiation field. The ISO instruments have measured the IR emission of dust particles of all sizes in a wide variety of astrophysical environments. The analysis, still ongoing, of this large database has already lead to important steps forward in our understanding of interstellar dust as examplified here and in other reviews of this volume. ISO observations of the mid-IR dust emission spectrum have highlighted the properties of PAHs and VSGs as well as their evolution in the interstellar lifecycle. For the first time, direct evidence of important variations of the abundance of small grains in diffuse and cold clouds has been found in deep photometric ISO data. These latter results point at the importance of dust processing in the interstellar medium (see Section 3.4).

3.1. THE EMISSION SPECTRUM OF SMALL GRAINS AND ITS IMPLICATIONS

One of the prominent results of ISO was to show the ubiquitous presence of the AIBs (at 3.3, 6.2, 7.7, 8.6, 11.3 and 12.7 μm, see Figure 2) in the interstellar medium of the Galaxy (e.g., H II regions: Peeters et al., 2002; planetary nebulae: van Diedenhoven et al., 2004; Uchida et al., 2000; PDRs: Cesarsky et al., 2000b; Verstraete et al., 2001; cirrus and diffuse ISM: Boulanger et al. 1996; Mattila et al., 1996; Chan et al., 2001; Kahanp¨a¨a et al., 2003), but also in external galaxies (see the reviews by Sauvage et al., and Verma et al., in this volume). Furthermore, the ratio of the 7.7 μm band to the local continuum has now become a tool to measure the starburst activity in galaxies (Lutz et al., 1998). THE COOL ISM 255

Figure 2. Typical CAM (spectral resolution R ∼ 40, top and middle panels) and SWS (R ∼ 200, bottom panel, from Verstraete et al., 2001) spectra in NGC 2023 (from Abergel et al., 2002). For the CAM spectra, the aromatic lines are fitted with Lorentzian profiles, while the S(1), S(2) and S(3) pure rotational lines of H2 are adjusted with gaussian profiles. The continuum arises further inside the molecular cloud than the AIBs (top panel).

In the cool ISM, away from hot stars (of spectral type earlier than B2), the 6–13 μm AIB spectrum has been found to be strikingly similar considering the large range of excitation conditions spanned by the observations (χ ∼ 1–104, Boulanger et al., 2000; Uchida et al., 2000; Chan et al., 2001; Verstraete et al., 2001). Conversely, the AIB spectrum towards compact H II regions and circumstellar environments changes significantly as shown in Hony et al. (2001), Peeters et al. (2002) and van Diedenhoven et al. (2004). The authors argued that these variations are due to composition changes in a population of PAHs including substituted and metal complexed species. In fact, newly formed PAHs are expected to be found in the ejecta of evolved stars while in the general interstellar medium (which includes the cool ISM), the population of PAHs has already been processed by the ambient UV radiation field and shocks over long timescales (∼108yrs). In the following we call the AIBs observed in the interstellar medium, the IS-AIBs (interstellar AIBs). We discuss below the IS-AIBs while the AIBs around stars and H II regions are discussed in Peeters et al. (this volume). For the first time, most AIBs were spectrally resolved by the ISO spectrometers. Individual AIBs have been found to be broad (λ/FWHM ∼30–80), smooth and in some cases clearly asymmetrical (the 3.3, 6.2 and 11.3 μm bands all show a prominent red wing). As shown in Boulanger et al. (1998b) and Verstraete et al. (2001) the AIBs are well represented with Lorentzian profiles possibly indicating 256 A. ABERGEL ET AL. that the AIB carriers are large molecules. Other methods of decomposition have been applied to quantify the AIB intensities (Peeters et al., 2002; Uchida et al., 2000) illustrating the importance of the AIB definition for the extraction of the underlying continuum. SWS data has clearly shown that the 7.7 μm-feature consists of several components (Joblin et al., 2000; Peeters et al., 2002), including the two main features at ∼7.6 and 7.8 μm first reported by Bregman (1989). New features have been added to the PAH family, like the 16.4 μm band (van Kerckhoven et al., 2000; Moutou et al., 2000). This band has been assigned by the latter authors to PAHs containing pentagonal rings, species which are important in chemical networks of PAH and soot formation (Moutou et al., 2000). The 16.4 μm band is also present in CAM-CVF data but falls at the detection edge of the detectors (Figure 2). Some AIBs have also been detected in absorption (see Section 3.3). More recently an emission band at 4.65 μm has been assigned to deuterated PAHs (Peeters et al., 2004). More features are expected to be found in the ISO archive but also thanks to the Infrared Spectrograph (IRS) now operating on the SST. A recent study on the PDR NGC 7023 reports new features at 6.7, 10.1, 15.8, 17.4 and 19.0 μm (Werner et al., 2004). Features in the far-IR and submm range could also be detected in the future by the Herschel Space Observatory (Joblin et al., 2002). In the IS-AIB spectrum, the positions and widths of the major bands (3.3, 6.2, 8.6 and 11.3 μm) are stable (within a few cm−1, Verstraete et al., 2001), conversely to spectra towards H II regions and circumstellar environments where significant shifts of the 6–9 μm bands are observed (Peeters et al., 2002). To constrain the nature of PAHs, detailed comparisons between the observed AIBs and the laboratory and theoretical data on small (NC ≤ 50), planar, aromatic molecules have been carried out. Pech et al. (2002) have shown that the profiles of the 6.2 and 11.3 μm bands observed in the planetary nebula IRAS 21282+5050 can be explained with the emission of a mixture a medium-sized PAHs (30 < NC < 80) having similar spectroscopic properties. The band positions and widths arise then naturally from the molecular anharmonicity during temperature fluctuations (see Section 3.2). Thus, the strong variability of the vibrational spectrum of terrestrial PAHs as a function of charge state, hydrogen coverage and chemical structure (e.g., Hudgins et al., 2000; Pauzat et al. 1997; Langhoff 1996) is not reflected in the IS-AIBs of the cool ISM where PAHs are expected to be for instance neutral (χ ∼ 1 to a few 100) or ionized (χ ∼ afew103, Dartois and d’Hendecourt 1997). Furthermore, a precise, quantitative assignation of major AIBs is still not achieved: the components of the 7.7 μm and the 8.6 μm bands are not explained by the properties of small terrestrial PAHs (Verstraete et al. 2001; Peeters et al., 2002). Spectroscopic evidence (van Kerckhoven et al., 2000; Hony et al., 2001) and modelling (Schutte et al., 1993; Li and Draine, 2002) has pointed out the importance of large PAHs (NC ≥ 50) for the 6–13 μm AIBs. The vibrational bands of such large systems may be less sensitive to the structure, shape, charge and surface state and provide a better understanding of the observed AIBs. However, recent calculations by Bauschlicher (2002) on the THE COOL ISM 257 large PAH C96H24 seem to challenge this picture while laboratory measurements on large systems are still awaited. Another conspicuous aspect of the mid-IR spectrum of the ISM is the presence of broad features (λ/FWHM ∼10) and continua underneath the AIBs. These spectral components are seen in the very excited environments studied by Peeters et al. (2002) such as planetary nebulae but also in the cool ISM (see Figure 2). The intensity of the broad features around ∼7, 10 and 12 μm are best delineated in SWS data where all AIBs are spectrally resolved. Their detailed shape however is poorly constrained (Verstraete et al., 2001). This emission has been discussed in the past and proposed to be associated to the AIB spectrum but not directly correlated with the bands3 (Cohen et al., 1985, 1986, 1989; Roche et al., 1989; Bregman 1989). These broad features and continua may arise from larger species (NC a few 100) as suggested by several CAM studies of the spatial distribution of this broad-band emission relative to the AIBs in PDRs (Section 3.4). Here also laboratory and/or theoretical spectroscopy of carbon clusters is needed to quantitatively interpret these broad bands and continua.

3.2. EMISSION MECHANISM OF THE AROMATIC INFRARED BANDS AND MODELLING

The high sensitivity and dynamical range of the ISO instruments has provided detection of the AIBs from diffuse to bright regions. In particular, it has been found that the integrated flux of the AIBs globally scales with the intensity of the FUV radiation field as long as χ ≤ 104 (Boulanger et al., 1998a; Uchida et al., 2000). This strongly supports a picture in which small grains are transiently heated to high temperature each time they absorb a UV photon (Puget et al., 1985; Sellgren et al., 1985). It has also been shown that the AIBs can be excited by visible photons (Uchida et al., 1998) and this is well explained by AIB emission models, provided that PAHs are ionized or large (Li and Draine, 2002). Studies of the visible, near- IR absorption cross-sections of PAHs with different sizes and in different charge and hydrogenation states are important since, for instance, in the standard galactic radiation field (Mathis et al., 1983) PAHs absorb about half of their energy in this wavelength range. On the other hand, the high spectral resolution of the SWS has stimulated much work on the AIB profiles and on their formation mechanism. Indeed, this dedicated spectroscopy has provided a unique opportunity to compare to laboratory results. Specifically, detailed studies of the IR emission spectra from small, gas-phase PAHs excited by UV photons or thermally heated have shown that the spectral profile (position and width) of the bands depends on the internal vibrational energy (or temperature) of the molecule (Kim et al., 2001; Cook et al., 1998 and references

3The laboratory spectra of individual PAH molecules do not seem to present any continuum (e.g., Joblin et al., 1995). 258 A. ABERGEL ET AL. therein; Joblin et al., 1995). The observed emission profiles hence result from the sum of the individual lorentzian profiles formed during the cooling of the molecule after UV excitation. The asymmetrical profiles of the 3.3, 6.2 and 11.2 μm bands can thus be explained with an intramolecular broadening due to coupling between vibrational modes in PAHs (Pech et al., 2002). Assuming similar spectroscopic properties for a broad range of PAH size, the stability of the 6–13 μm IS-AIBs is naturally explained by an emission during thermal fluctuations which span a large range of temperatures. In this framework and from detailed profile modelling, one can also infer a photodissociation threshold (in terms of internal energy per carbon atom, a few 0.1 eV/C) beyond which small PAHs are destroyed. The smallest PAHs (NC 40) carrying the 3.3 μm band would then be “selected” by photodissociation leading to similar distribution of internal energy. This would explain why the profile of this AIB is so stable in increasingly excited PDRs (Verstraete et al., 2001; van Diedenhhoven et al., 2004). Finally, the mid-IR spectra observed by the ISO has allowed an empirical definition of the size-dependent IR cross-sections of PAHs and VSGs which have been used in recent quantitative modelling of the dust emission (Li and Draine, 2001).

3.3. MID-INFRARED EXTINCTION

The extinction curve is one of the main tool to study interstellar dust, to correct the photometry of embedded sources and to convert the measured extinctions into column densities. Before ISO, the mid-IR extinction curve was thought to be rel- −1.7 atively well understood, and described with a power law Aλ ∼ λ from ∼1μm down to a pronounced minimum around 7 μm due to the increasing contribution of the silicate absorption feature at 9.7 μm. It was explained by the absorption of a mixture of graphite and silicate grains (Draine, 1989). The situation is now more controversial, as illustrated in the Figure 4 of the recent review of Draine (2003). Using the measurements of hydrogen recombination lines with the SWS, Lutz et al. (1996) and Lutz (1999) have shown that the IR extinction towards Sgr A does not decrease with increasing wavelengths from 4 to 8 μm, and as a consequence does not present any minimum near 7 μm. On the other hand, Rosenthal et al. (2000) using rovibrational lines of H2 detected by the SWS in Orion have found that the mid-IR extinction decreases with increasing wavelengths and presents a minimum at ∼6.5 μm, as expected. The IR extinction curve depends on the composition of dust grains and to a lesser extent on their size distribution (for λ 2 μm) and may vary from place to place. New observations are therefore necessary to clarify this controversy. The IR extinction curve also presents a number of faint features detected with the SWS towards the galactic centre and other sources with large extinction (Schutte et al., 1998; Chiar et al., 2000; Bregman and Temi 2001). The ubiquitous 3.4 μm band is generally attributed to CH stretching in aliphatic hydrocarbons. The asymmetrical deformation modes of CH at ∼6.9 μm have been detected before ISO THE COOL ISM 259 with the KAO (Tielens et al., 1996), but the symmetrical deformation mode at 7.2 μm has been found for the ﬁrst time by Chiar et al. (2000) with SWS data. The strengths of the 3.4, 6.9 and 7.2 μm features appear compatible with the hypothesis that Hydrogenated Aliphatic Carbon (HAC) is the main carrier. Finally, the 3.3 and 6.2 μm bands are detected in absorption towards several lines of sight, and are attributed to PAHs. At this time the absorption of other aromatic bands have not been evidenced in the SWS data. As mentionned by Bregman and Temi (2001), a calibration problem in the SWS data at 11.1 μm may preclude the detection of the 11.2 μm band in absorption. These authors have discovered an absorption band centered at 11.2 μm from ground-based spectra of the molecular cloud surrounding Monoceros R2 IRS 3, and attributed this band to the C H out-of-plane vibrational mode of PAHs.

3.4. DUST EVOLUTION AND PROCESSING

The dense shells around evolved stars are known to form dust grains as demonstrated by observations of the dust IR emission and absorption features in these objects (e.g., Blommaert et al., Molster and Kemper and Peeters et al., in this volume). The timescale of this stardust injection is ∼109 yrs. On the other hand, theoretical studies show that dust grains are efficiently destroyed in supernova shocks (Jones et al., 1994) with a lifetime ∼108 yrs. Efficient mechanisms for grain growth must therefore exist in the other phases of the interstellar medium, most probably in dense, cold clouds where the rates of grain coagulation and accretion of gas phase species are expected to be the largest. High energy gas–grain collisions lead to erosion by vaporization (sputtering), while low energy collisions lead to the reverse process of gas accretion onto dust (Jones et al., 1994). Collisions between large grains (radii 100 nm) at velocities above a few km/s lead to fragmentation into small grains (radii 1 nm) or even grain destruction. By contrast, coagulation occurs through low energy collisions (Jones et al., 1996). UV irradiation can also play a major role by destroying dust grains in high excitation regions (with typically χ>104) and also photo-evaporating condensed species. Evidence for photoprocessing of PAHs and VSGs has been found towards H II regions where the 6–13 μm emission spectrum is dramatically different from the IS-AIBs (Roelfsema et al., 1996; Verstraete et al., 1996). Such processing must leave specific signatures in the dust size distribution and optical properties which in turn affect the gas thermal state and chemistry. As we discuss below, ISO has brought important observational evidences of the evolution of the dust size distribution in the ISM, from diffuse cirrus clouds to dense cores.

3.4.1. PDRs Grain dynamics driven by the radiation field could play a major role in the observed evolution of the mid-infrared spectrum across PDRs. Size segregation in regions anisotropically irradiated by stars may arise from radiation pressure and 260 A. ABERGEL ET AL. recoil forces (due to the photo-ejection of electrons, the photodesorption of weakly- bound surface species and the formation of molecular hydrogen, see Weingartner and Draine, 2001c). This process could strongly modify the size distribution at the illuminated surface of PDRs. Furthermore, the gas in PDRs is suddenly heated and ionized, and experiences a strong pressure gradient. It therefore flows supersoni- cally from the molecular cloud surface to the tenuous intercloud medium, carrying away dust grains. In this highly turbulent flow, grain-grain collision might be strong enough to fragment big grains. This process could explain spectro-imaging observations with CAM across several low-excitation PDRs (NGC 2023, 2068, 2071 and 7023; ρ Oph, see Abergel et al., 2002, 2003b; Rapacioli et al., 2005) showing that the ratio of the AIBs (especially around 7.7μm) to the underlying continuum is higher in the photodissociated interface than in more shielded regions. Variations related to a transition in the charge state of PAHs are also expected. Indeed, PAHs are predicted to be neutral inside dense clouds, and positively ionised in the low-density regions facing the illuminated surface of PDRs (e.g., Dartois and d’Hendecourt 1997). Laboratory data and theoretical calculations show that the single ionization of PAHs strongly enhances (by a factor of ∼10) the emissivity of the C C stretching (6.2 and 7.7 μm) and the C H in plane bending (8.6 μm) modes, with respect to the intensity of the C H out-of-plane bending features at 11.3 and 12.7 μm (Hudgins et al., 2000; Langhoff, 1996; Szczepanski and Vala, 1993). Yet, as already discussed in Section 3.1 a direct decomposition of the mid-IR dust emission spectrum shows only limited variations of the AIB ratios in the cool ISM (factor 2 or less, e.g. Uchida et al., 2000). Rapacioli et al. (2005) went one step further by applying the singular value decomposition method to CAM-CVF data in order to extract typical emission spectra for the different population of small grains in PDRs. Three elementary spectra could be extracted (Figure 3):

– The ﬁrst two spectra are dominated by the AIBs and attributed to neutral PAH and PAH cations. The differences between these two spectra follow the expected changes due to ionisation (increase of the (6.2–7.7–8.6 μm)/(11.3–12.7 μm) band ratio and blueshift of the 11.3 μm band in cations according to Hudgins and Allamandola (1999). The PAH cation spectrum is dominant in the low-density gas surrounding the star whereas the PAH neutral spectrum dominates at the interface with the molecular cloud. – The third spectrum contains a strong continuum and dominates behind the illuminated edge of the PDR (far from the illuminating star), and is attributed to very small carbonaceous grains. The fact that a similar spectrum has been seen on a few pixels of the CAM map of the illuminated edge of Ced 201 (Cesarsky et al., 2000a) lends conﬁdence to the analysis performed by Rapacioli et al. (in press). Several features visible in the spectrum indicate a possible aromatic nature of the carriers. Interestingly, the 7.8 μm component of the broad 7.7 μm feature appears to be carried by this population. The authors also argued that these carbonaceous VSGs could be the progenitors of free PAHs. Indeed, the THE COOL ISM 261

Figure 3. Elementary spectra extracted from spectro-imaging CAM observations of NGC 7023 (Rapacioli et al., in press). These spectra are attributed to PAH cations (upper panel), neutral PAHs (middle panel) and carbonaceous clusters (lower panel).

drop of the VSG emission at the cloud edge appears to be correlated with the increase of the neutral PAH emission. These VSGs may be clusters of PAHs which photoevaporate at the surface of clouds. This idea is consistent with the strong spatial variations observed in the IRAS 12/100 μm ratio at the surface of molecular clouds (Boulanger et al., 1990; Bernard et al., 1993). SST observations are underway to confirm these results on a larger sample of objects (the CAM- CVF data suffer from stray light artefacts which limit their useful dynamics, see Boulanger et al., in press), and search for new features at longer wavelengths which can be attributed to these VSGs. It should be emphasized that the question of the formation of PAHs in the ISM is still open. Some PAHs could be formed in the outflows of evolved carbon-rich stars in a chemical kinetic scheme based on soot formation in hydrocarbon flames (Frenklach and Feigelson, 1989; Cherchneff et al., 1992). Cernicharo et al. (2001) detected C4H2,C6H2 and C6H6 in the proto-planetary nebula CRL 618, showing that aromatic molecules can indeed be formed in such objects and that energetic processes such as UV irradiation or shocks play an important role in the chemistry of these species.

3.4.2. Cirrus Clouds Important variations of the PAH abundance have been observed in diffuse regions like cirrus clouds. Based on a comparison between CAM-LW2 (6.7 μm broad band) and radio data, Miville-Deschˆenes et al. (2002) have shown variations of 262 A. ABERGEL ET AL.

Figure 4. The L1780 cloud (from Miville-Deschênes et al., 2003). Left: IRAS 100 μm (top) and 60 μm (middle), and CAM 6.7 μm (bottom) maps across L1780. Right: Average emission as a function of right ascension. The 6.7 and 60 μm were scaled to be compared to the 100 μm emission. The spectacular decrease of the 6.7 μm / 100 μm ratio cannot be explained by the attenuation of the radiation field. The abundance of PAHs likely decreases, a trend which is spatially correlated with the apparition of CO emission. the abundance of PAHs by up to a factor 10 in a high latitude cirrus cloud. High PAH abundances were found in regions with strong turbulent motions while low abundances were associated with regions of shallower turbulent motions where CO emission is also detected (see Figure 4). This study highlights for the first time the important role of the gas turbulent motions on the evolution of the grain size distribution. The depletion of small grains observed in the denser parts of cirrus clouds is the first step of the dust coagulation seen in cold clouds.

3.4.3. Cold Clouds Depletion of PAHs and VSGs in the dense regions of molecular clouds has also been shown by ISO data suggesting the importance of dust coagulation processes in the ISM. IRAS images first showed that the emission from small grains at 12, 25 and 60 μm, relative to that from big grains (in thermal equilibrium with the radiation field) at 100 μm, varies by one order of magnitude among and within translucent clouds in the nearby ISM (Boulanger et al., 1990). Standard model calculations predict that the FIR emission of dust in thermal equilibrium with the radiation field follows a modified black body with a spectral index β between 1 and 2 (e.g., Hildebrandt, 1983; Draine and Lee, 1984). The brightness per hydrogen λ −β τλ atom can be written: Iλ = λ × ( ) × Bλ(T ) = × Bλ(T ). The dust opacity 0 λ0 NH τλ can be derived if the brightness, the column density and the temperature of the emitting particles are known at a given wavelength. However, the main unknown in our understanding of the big dust grains is the FIR emissivity law, due to the observational difficulty to accurately determine the temperature and β independently, together with the column density. FIR observations of a variety of mostly THE COOL ISM 263 quiescent clouds (i.e. no high mass star formation) have been used to solve this problem. The longest wavelength that can be measured with PHOT and LWS is slightly more than 200 μm, which is not sufficient to allow an independent determination of β and T. On the other hand, the wavelength coverage of the balloon-borne experiment SPM/PRONAOS (λ = 200–580 microns, Lamarre et al., 1994) is extremely suitable for the determination of β, but PRONAOS is less sensitive to dust temperature (except in the coldest objects, T < 15 K, where the spectral energy distribution peaks around the shortest PRONAOS band), so IRAS or PHOT must also be used. For comparison with dust models, the FIR opacity needs to be gauged against the extinction curve in either the visible or near-infrared. This can be done by determining the optical extinction directly from colour excesses or from star counts. The former method is constrained to a limited amount of sightlines and a maximum extinction. The latter method is also limited to a maximum extinction (of AV < 8), and suffers from higher uncertainties in the denser regions due to a lack of statistical accuracy. A study by Arce and Goodman (1999) showed that for low and moderate density regions (AV < 4), the different extinction estimates are reliable and give similar results. With the advent of large stellar databases (e.g. USNO, 2MASS) it has become much easier to perform star counts over any area on the sky, providing valuable extinction information. Alternatively, one can estimate the visual extinction from the total gas column density and assuming a ratio AV/Ngas. This indirect method depends on the assumed ratio between gas and dust and the conversion between molecular line strength and column density. Recent observations by FUSE (Rachford et al., 2002) confirmed that ratio between the total hydrogen / − − ∼ column density and reddening NH+H2 E(B V ) is constant up to E(B V ) 1 and well described by the constant presented by Bohlin et al. (1978), obtained from the Copernicus satellite data. Very high extinctions can be measured from near- infrared extinction of background emission, first noticed in the ISOCAM images (Perault et al. 1996; Abergel et al., 1996) and used to study the very dense cores (Nisini et al., this volume). For the diffuse ISM associated with atomic hydrogen, Boulanger et al. (1996) found a value of β close to 2.0 and a mean temperature of 17.5 K by correlat- ing COBE data with Dwingeloo HI data. The derived FIR emissivity τ(λ)/NH = 1.0 × 10−25 (λ/250 μm)−2 cm2 is close to the value predicted by Draine and Lee (1984) based on a standard dust model. This result suggests that spherical grains are more likely to occur in the diffuse ISM than porous or fractal grains. Assuming RV = −4 −2 3.1, the result by Boulanger et al. implies τ(λ)/AV = 2.9×10 (λ/200 μm) .We will use this value as a reference in the discussion below. In case of grain processing towards fractal or porous grains, one would expect an increase of the FIR emissivity and a corresponding increase in the value of τ(λ)/AV (e.g., Ossenkopf and Henning, 1994). The first ISO studies of moderate density regions indicated dust temperatures of 13.5 ± 2 K assuming β = 2 (Laureijs et al., 1996; Lehtinen et al., 1997). The 264 A. ABERGEL ET AL. observed ratios τ(200)/AV had large uncertainties but were not inconsistent with the predictions for a standard dust model. For increasing densities (in the range 100 to 104 cm−3) the temperature is generally found to decrease from ∼17.5 K to as low as 12 K (Tóth et al., 2000; Juvela et al., 2002; del Burgo et al., 2003; Lehtinen et al., 2004). The first strong indication that the FIR properties of the dust in moderate density high latitude clouds might deviate from the diffuse ISM was presented by Bernard et al., (1999) based on SPM/PRONAOS and ISO observations of a cirrus cloud in Polaris. Bernard et al. find cold dust at 13.0 K with β = 12.2 ± 0.3. Since the cloud has a low visual extinction (AV < 1), only an anomalously low ambient radiation field could cause such a low temperature, which is ruled out. It was therefore concluded that the dust properties of the large grains have changed in −4 the cloud. The inferred ratio τ(200)/AV = 17 ± 3×10 is more than 5 times the value for the diffuse ISM. Other PRONAOS observations combined with IRAS data of a dense filament in the Taurus molecular cloud by Stepnik et al., (2003) has evidenced temperatures as low as of 12.1 ± 0.2 K in the densest regions (with β = 2), together with an increase of the FIR emissivity by a factor larger than 3 compared with the diffuse ISM. Assuming β = 2, del Burgo et al. (2003) have analysed a sample of moderate density regions (AV < 6) observed with ISO. The regions have dust temperatures in the range from 13.5–15.6 K and one exceptional region with T = 18.9 K. Com- bination with extinction data from star counts shows that the ratio τ(200)/AV is systematically higher than for the diffuse ISM. In fact a trend was found where τ(200)/AV increases with decreasing temperature. A similar trend, although less apparent, was found by Cambrésy et al. (2001) who correlated large scale extinction measurements from starcounts with far-infrared DIRBE data in the Polaris Flare region. They find τ(200)/AV in the coldest regions with T < 15.5 K to be several times higher than in the diffuse ISM. Lagache et al. (1998), using COBE data, and Tóth et al. (2000), using ISO serendipity mode data, found that the temperature of the dust can be significantly below 17.5 K over large areas. In these areas, the abundance of small grains is low as inferred from the absence of 60 μm emission. Earlier studies have revealed a strong anticorrelation between the 60 μm emission and dense gas traced by the millimetre transitions of 13CO (Laureijs et al., 1991; Abergel et al., 1994). In the dense molecular cloud TMC2, a large scale cold component with a uniform temperature of about 12.5 K is evidenced from PHOT data, while only the densest cores show lower temperatures (del Burgo and Laureijs, 2005). The dust spatially associated with the 60 μm emission is significantly warmer (∼19 K with β = 2). In addition the morphology of this “warm” component is different from the cold component suggesting that the cold and warm components are spatially separate. The evolution of dust properties with the gas density is believed to be due to the process of gas accretion onto grains and grain coagulation. Gas accretion forms grain mantles typically for AV 2 mainly composed of water ice and carbon oxides THE COOL ISM 265

(e.g., Teixeira and Emerson, 1999). However, Preibisch et al. (1993) has shown that the presence of such mantles on dust grains does not significantly increase the FIR emissivity. As the gas cools down and condenses, gas turbulent motions become shallow leading to grain-grain collisions with sufficiently low energy to enable grain sticking and coagulation, which can lead to the formation of clusters with an enhancement of the FIR emissivity (e.g., Ossenkopf and Henning, 1994; Stepnik, 2001; Stepnik et al., in press) compatible with the PHOT and SPM/PRONAOS observations. Such a scenario can also qualitatively explain the observed general trend in the extinction curve, i.e., an increase in RV with the density (e.g., Cardelli et al., 1989). At first order, the coagulation process appears to be compatible with the dynamical and cloud evolutionary time-scales. The formation of ice mantles is also believed to increase the sticking probability for coagulation. The physics of grain–grain collisions, coagulation and growth has been studied theoretically (Dominik and Tielens 1997) and in the laboratory (Blum, 2004) in the context of proto-planetary disks. However, the efficiency of these processes in the ISM still needs a dedicated modelling taking into account the velocity structure of the ISM (e.g., turbulent and intermittent processes). In the last decade laboratory measurements (Agladze et al., 1996; Mennella et al., 1998) have also indicated that, at low temperatures (T < 30 K), the dust optical properties can change significantly likely because of low-energy structural transformations. A systematic analysis of β from a large compilation of PRONAOS data by Dupac et al. (2003) has shown that β is not a constant but varies with dust temperature: high indices (1.6–2.4) being observed in cold regions (11–20 K) while low indices (0.8–1.6) are observed in warm regions (35–80 K). Although systematic variations in grains size and/or chemical composition as a function temperature cannot be ruled out, the most likely explanation is that the spectral index of dust can change with temperature due to a solid state quantum process. Such studies allow a quantitative understanding of the FIR emission of dust in our Galaxy and in outer galaxies (a cold dust component has also been detected with PHOT in non-active galaxies, see the chapter by Sauvage et al., this volume) and pave the way for the scientific analysis of the Herschel and Planck missions.

4. Structure of the Cool ISM

The structure of the ISM strongly affects the chemistry and the star formation activity. Observations of the IR dust emission is unique for tracing interstellar matter over a wide range of physical conditions and, in particular, across the atomic to molecular transition where neither H I nor molecular millimetric lines are robust tracers because of the local variations of the abundance and the excitation conditions. CAM broad-band mapping of molecular clouds has shown the spatial distribution of the emission due to small grains with an angular resolution two orders of magnitude better than IRAS at 12 μm. The images are, to ﬁrst order, anti-correlated with the visible emission (e.g., Figure 5), and contain small scale brightness ﬂuctuations 266 A. ABERGEL ET AL.

Figure 5. L1630 in the visible (35 × 40 field from the UK Schmidt Telescope and extracted from the DSS produced at the STSI) and with CAM (blue: 5–8.5 μm, red: 12–18 μm) from Abergel et al. (2002). The central region (NGC 2023) is saturated because of the dynamic range. THE COOL ISM 267 related to illuminated edges of dense regions. For nearby objects with edge-on geometry, IR filaments as sharp as ∼10 in. (or ∼0.02 pc at a distance of 400 pc) are detected due to the combined effects of steep density gradients and extinction in the dense regions (e.g., Abergel et al., 2003a in the illuminated edge of the Horsehead nebula). Associated measurements of the penetration length of the incident radiation field gives estimates of the density just behind the illuminated edges. More quantitative constraints on the density profiles can be derived from the comparison of PDR modelling with spatial profiles of the AIBs as well as of cooling lines or molecular emissions taken with ISO or from the ground (Habart et al., 2005; Pety et al., 2005; Habart et al., 2001). Yet, such analysis should be carried out bearing in mind the discrepancies between the line intensities predicted by different PDR codes (van Dishoeck, 2004). Moreover, caution should be also exercised in analysing the spatial profiles observed across PDRs because of the complex interface geometry seen in the high angular resolution data of the ISO cameras (e.g., Habart et al., 2005). This is also true of ground-based and of recent SST data. The spatial structure of the cool ISM appears self-similar in a wide range of scales. With PHOT images of cirrus emission at 90 μm, the power law relation established from IRAS data (Gautier et al., 1992) between the power of the fluctuations and the spatial scale has been extended to twice as high spatial fequencies, while at 170 μm fluctuations at arcmin scales are studied for the first time (Herb- stmeier et al., 1998; Kiss et al., 2003). The spectral index is found to vary from field to field (−5.3 <α<−2.1), depending on the absolute surface brightness. Higher values of α are found at 170 μm rather than at 90 μm and are attributed to temperature variations of dust. Compared to ISO, the SST has the sensitivity, the angular resolution and the mapping capabilities to extend significantly these studies (see the first results by Ingalls et al., 2005).

5. Conclusions

The ISO mission has provided an incredibly rich harvest of gas and dust features in the general interstellar medium away from hot stars, the cool ISM. These results have highlighted the physics and chemistry of interstellar matter and its evolution. Emission and absorption lines of gas phase species were used to constrain the excitation processes (shocks, stellar UV irradiation) of the ISM and represented a major benchmark for the modelling. Similarly, spectroscopy of dust features has allowed a detailed comparison to the properties of terrestrial analogues of interstellar grains. These studies have pointed out the importance of nanometric carbon clusters in space and the need to build a laboratory and/or theoretical database of spectra for a large size range of these species. The spatial distribution of the dust emission observed by the photometers has clearly shown and for the ﬁrst time, the importance of evolution processes (e.g., photodestruction, grain coagulation) affecting the dust 268 A. ABERGEL ET AL. size distribution. The joint observations of dust and gas features have highlighted the gas–grain couplings (photoelectric effect, H2 formation) and allowed quantitative studies of the underlying microphysics. Furthermore, the quantitative understanding of the gas-to-dust correlation for all dust components is a key issue for the analysis of the extended emission in the future surveys to be conducted with Planck and Herschel, allowing to disentangle brightness ﬂuctuations due to the ISM from those of the cosmic IR and microwave backgrounds.

References

Abergel, A., Boulanger, F., Mizuno, A., et al.: 1994, Astrophys. J. 423, 59. Abergel, A., Bernard, J.-P., Boulanger, F., et al.: 1996, Astron. Astrophys. 315, 329. Abergel, A., Bernard, J. P., Boulanger, F., et al.: 2002, Astron. Astrophys. 389, 239. Abergel, A., Teyssier, D., Bernard, J. P., et al., 2003a, Astron. Astrophys. 410, 577. Abergel, A., 2003b, in C. Gry, S. Peschke, J. Matagne, P. Garcia-Lario, R. Lorente, and A. Salama (eds.), Exploiting the ISO Data Archive, ESA SP-511, p. 177. Agladze, N. I., Sievers, A. J., Jones, S. A., et al.: 1996, Astrophys. J. 462, 1026. Allamandola, L. J., Tielens, A. G. G. M., and Barker, J. R.: 1985, Astrophys. J. 290, L25. Arce, H. G. and Goodman, A. A.: 1999, Astrophys. J. 517, 264. Bakes, E. L. O. and Tielens, A. G. G. M.: 1994, Astrophys. J. 427, 822. Bauschlicher, C. W.: 2002, Astrophys. J. 564, 782. Bernard, J. P., Boulanger, F., and Puget, J. L.: 1993, Astron. Astrophys. 277, 609. Bernard, J.-P., Abergel, A., Ristorcelli, I., et al.: 1999, Astron. Astrophys. 347, 640. Bohlin, R. C., Savage, B. D., and Drake, J. F.: 1978, Astrophys. J. 224, 132. Boonman, A. M. S., van Dishoeck, E. F., and Lahuis, F.: 2003, Astron. Astrophys. 399, 1047. Boulanger, F., Falgarone, E., Puget, J.-L., et al.: 1990, Astrophys. J. 364, 136. Boulanger, F., Abergel, A., Bernard, J. -P., et al.: 1996, Astron. Astrophys. 312, 256. Boulanger, F., Abergel, A., Bernard, J.-P., et al., 1998a, in J. L. Yun and R. Liseau (eds.), Star Formation with ISO, Astronomical Society of the Paciﬁc, San Francisco, p. 15. Boulanger, F., Boissel, P., Cesarsky, D., and Ryter, C., 1998b, Astron. Astrophys. 339, 194. Boulanger, F., Abergel, A., Cesarsky, D., et al.: 2000, in R. J. Laureijs, K. Leech, and M. F. Kessler (eds.), ISO Beyond Point Sources: Studies of Extended IR Emission, ESA SP-455. Boulanger, F., Lorente, R., Miville-Deschˆenes,M. A., et al., 2005, Astron. Astrophys. 436, 1151. Blum, J.: 2004, in A. N. Witt, G. C. Clayton, and B. T. Draine (eds.), Astrophysics of Dust, ASP Conference Series, vol. 309, p. 369. Bregman, J. D. and Temi, P.: 2001, Astrophys. J. 554, 126. Bregman, J. D.: 1989, in Interstellar dust, Proceedings of IAU Symposium, vol. 135, p. 109. Cambre´sy, L., Boulanger, F., Lagache, G., and Stepnik, B.: 2001, Astron. Astrophys. 375, 999. Cardelli, J., Clayton, G. C., and Mathis, J. S.: 1989, Astrophys. J. 345, 245. Caux, E., Ceccarelli, C., Pagani, L., et al.: 2002, Astron. Astrophys. 383, L9. Ceccarelli, C., Baluteau, J.-P., Walmsley, M., et al.: 2002, Astron. Astrophys. 383, 603. Cernicharo, J., Heras, A. M., Tielens, A. G. G. M., et al.: 2001, Astrophys. J. 546, L123. Cernicharo, J., Goicoechea, J. R., and Benilan, Y.: 2002, Astrophys. J. 580, 157. Cesarsky, D., Lequeux, J., Ryter, C., and Gerin, M., 2000a, Astron. Astrophys., 354, L87. Cesarsky, D., Jones, A. P., Lequeux, J., et al., 2000b, Astron. Astrophys. 358, 708. Chan, K.-W., Roellig, T. L., Onada, T., et al.: 2001, Astrophys. J. 546, 273. Cherchneff, I., Barker, J. R., and Tielens, A. G. G. M.: 1992, Astrophys. J. 401, 269. Chiar, J. E., Tielens, A. G. G. M., Whittet, D. C. B., et al.: 2000, Astrophys. J. 537, 749. THE COOL ISM 269

Cohen, M., Tielens, A. G. G. M., Allamandola, L. J., et al.: 1985, Astrophys. J. 299, 93. Cohen, M., Allamandola, L. J., Tielens, A. G. G. M., et al.: 1986, Astrophys. J. 302, 737. Cohen, M., Tielens, A. G. G. M., Bregman, J., et al.: 1989, Astrophys. J. 341, 246. Cook, D. J., Schlemmer, S., Balucani, N., et al.: 1998, J. Phys. Chem. A, 102, 1465. Dartois, E. and d’Hendecourt, L. B.: 1997, Astron. Astrophys. 323, 534. Désert, F.-X., Boulanger, F., and Puget, J.-L.: 1990, Astron. Astrophys. 237, 215. Dominik, C. and Tielens, A. G. G. M.: 1997, Astrophys. J. 480, 647. Draine, B. T.: 1989, in B. H. Kaldeich (ed.), Infrared Spectroscopy in Astronomy, ESA SP-290, p. 93. Draine, B. T.: 2003, Annu. Rev. Astron. Astrophys. 41, 241. Duley, W. W. and Williams, D. A.: 1981, Mon. Not. Roy. Astron. Soc 196, 269. Dupac, X., Bernard, J.-P., Boudet, N., et al.: 2003, Astron. Astrophys. 404, L11. Dwek, E., Arendt, R. G., and Fixsen, D. J.: 1997, Astrophys. J. 475, 565. del Burgo, C., Laureijs, R. J., Abraham, P., and Kiss, C., 2003, Mon. Not. Roy. Astron. Soc. 346, 403. del Burgo, C. and Laureijs, R. J.: 2004, MNRAS 360, 901. Falgarone, E., Panis, J. -F., Heithausen, A., et al.: 1998, Astron. Astrophys. 331, 669. Falgarone, E., Verstraete, L., Pineau des Forêts, G., Flower, D., and Puget, J.-L.: 1999, in F. Combes and G. Pineau des Forêts (eds.), H2 in Space, Cambridge University Press, Cambridge, p. 225. Falgarone, E., Verstraete, L., Pineau des Forêts, G., et al., in press, Astron. Astrophys. Feuchtgruber, H., Helmich, F. P., van Dishoeck, E. F., and Wright, C. M., 2000, Astrophys. J. 535, 111. Förster-Schreiber, N. M., Roussel, H., Sauvage, M., and Charmandaris, V.: 2004, Astron. Astrophys. 419, 501. Frenklach, M. and Feigelson, E. D.: 1989, Astrophys. J. 341, 372. Gautier, T. N. III, Boulanger, F., Perault, M., and Puget, J. L.: 1992, Astron. J. 103, 1313. Goicoechea, J. R., Rodriguez-Fernandez, N. J., and Cernicharo, J.: 2004, Astron. Astrophys. 600, 214. Gry, C., Boulanger, F., Nehmé,C., et al.: 2002, Astron. Astrophys. 391, 675. Habart, E., Verstraete, L., Boulanger, F., et al.: 2001, Astron. Astrophys. 373, 702. Habart, E., Boulanger, F., Verstraete, L., et al.: 2003, Astron. Astrophys. 397, 623. Habart, E., Abergel, A., Walmsley, M., Teyssier, D., and Pety, J.: 2005, Astron. Astrophys. 437, 177. Harper, D. A., Low, F. J., Rieke, G. H., and Thronson, H. A. Jr., 1976, Astrophys. J. 205, 136. Herbstmeier, U., Abraham, P., Lemke, D., et al.: 1998, Astron. Astrophys. 332, 739. Herlin, N., Bohn, I., Reynaud, C., et al.: 1998, Astron. Astrophys. 330, 1127. Hollenbach, D. J. and Tielens, A. G. G. M.: 1999, Rev. Mod. Phys. 71, 173. Hony, S., van Kerckhoven, C., Peeters, E., et al.: 2001, Astron. Astrophys. 370, 1030. Hudgins, D. M. and Allamandola, L. J.: 1999, Astrophys. J. 516, L41. Hudgins, D. M., Bauschlicher, C. W., Allamandola, L. J., and Fetzer, J. C., 2000, J. Phys. Chem. A, 104, 3655. Ingalls, J. G., Reach, W. T., and Bania, T. M.: 2002, Astrophys. J. 579, 289. Ingalls, J. G., Miville-Deschênes, M. A., Reach, W. T., Noriega-Crespo, A., Carey, S. J., Boulanger, F., et al.: 2004, Astrophys. J. SS 154, 281. Joblin, C., Abergel, A., Bregman, J., et al.: 2000, in A. Salama, M. F. Kessler, K. Leech, and B. Schulz (eds.), ISO Beyond the Peaks: The 2nd ISO Workshop on Analytical Spectroscopy, ESA-SP 456. Joblin, C., Boissel, P., and Léger, A.: 1995, Astron. Astrophys. 299, 835. Joblin, C., Toublanc, D., Boissel, P., and Tielens, A. G. G. M.: 2002, Mol. Phys. Jones, A. P., Tielens, A. G. G. M., Hollenbach, D. J., and McKee, C. F.: 1994, Astrophys. J. 433, 797. Jones, A. P., Tielens, A. G. G. M., and Hollenbach, D. J.: 1996, Astrophys. J. 469, 740. Jones, A. P. and d’Hendecourt, L. B.: 2000, Astron. Astrophys. 355, 1191. Juvela, M., Mattila, K., Lehtinen, K., et al.: 2002, Astron. Astrophys. 382, 583. Kahanpää, J., Mattila, K., Lehtinen, K., et al.: 2003, Astron. Astrophys. 405, 999. 270 A. ABERGEL ET AL.

Kemper, C., Spaans, M., Jansen, D. J., et al.: 1999, Astrophys. J. 515, 649. Kim, H. S., Wagner, D. R., and Saykally, R. J. 2001, Phys. Rev. Lett. 86, 5691. van Kerckhoven, C., Hony, S., Peeters, E., et al.: 2000, Astron. Astrophys. 357, 1013. Kiss, Cs., Abrahm, P., Klaas, U., et al.: 2003, Astron. Astrophys. 399, 177. Lagache, G., Abergel, A., Boulanger, F., et al.: 1998, Astron. Astrophys. 333, 709. Lamarre, J.-M., Pajot, F., Torre, J.-P., Guyot, G., Bernard, J. P., de Luca, A., et al.: 1994, Infrared Phys. Technol. 35, 277. Langhoff, S. R.: 1996, J. Phys. Chem. 100, 2819. Laureijs, R. J., Clark, F. O., and Prusti, T.: 1991, Astrophys. J. 371, 602. Laureijs, R. J., Haikala, L., Burgdorf, M., et al.: 1996, Astron. Astrophys. 315, 316. Le Bourlot, J., Pineau des Forêts,G., Roueff, E., and Flower, D., 1993, Astron. Astrophys. 267, 233. Léger, A. and Puget, J. L.: 1984, Astron. Astrophys. 137, L5. Lehtinen, K., Lemke, D., Mattila, K., and Haikala, L. K.: 1998, Astron. Astrophys. 333, 702. Li, A. and Draine, B. T.: 2001, Astrophys. J. 554, 778. Li, A. and Draine, B. T.: 2002, Astrophys. J. 572, 232. Liseau, R., White, G. J., Larsson, B., et al.: 1999, Astron. Astrophys. 344, 342. Lutz, D., Feuchtgruber, H., Genzel, R., et al.: 1996, Astron. Astrophys. 315, 269. Lutz, D., Spoon, H. W. W., Rigopoulou, D., et al.: 1998, Astrophys. J. 505, L103. Lutz, D.: 1999, in P. Cox and M. F. Kessler (eds.), The Universe as Seen by ISO, ESA-SP 427, vol. 2, p. 623. Mattila, K., Lemke, D., Haikala, L. K., et al.: 1996, Astron. Astrophys. 315, L353. Mathis, J. S., Mezger, P. G., and Panagia, N.: 1983, Astron. Astrophys. 128, 212. Melnick, G., Gull, G. E., and Harwit, M.: 1979, Astrophys. J. 227, 29. Mennella, V., Brucato, J. R., Colangeli, L., et al.: 1998, Astrophys. J. 496, 1058. Miville-Deschênes,M. A., Boulanger, F., Joncas, G., and Falgarone, E.: 2002, Astron. Astrophys. 381, 209. Miville-DeschênesM. A. 2003, in Chemistry as Diagnostic of Star Formation, NRC press, Canada, p. 363. Miville-Deschênes,M.-A., Joncas, G., Falgarone, E., and Boulanger, F.: 2003, Astron. Astrophys. 411, 109. Mizutani, M., Onaka, T., and Shibai, H.: 2004, Astron. Astrophys., 423, 579. Moutou, C., Verstraete, L., LégerA., et al.: 2000, Astron. Astrophys. 354, 17. Neufeld, D. A., Zmuidzinas, J., Schilke, P., and Phillips, T. G.: 1997, Astrophys. J. 488, 141. Ossenkopf, V. and Henning, T.: 1994, Astron. Astrophys. 291, 943. Pauzat, F., Talbi, D., and Ellinger, Y.: 1997, Astron. Astrophys. 319, 318. Pech, C., Joblin, C., and Boissel, P.: 2002, Astron. Astrophys. 388, 639. Peeters, E., Hony, S., van Kerckhoven, C., and Tielens, A. G. G. M: 2002, Astron. Astrophys. 390, 1089. Peeters, E., Allamandola, L. J., Bauschlicher, C. W. Jr., et al.: 2004, Astrophys. J. 604, 252. Perault, M., Omont, A., Simon, G., et al.: 1996, Astron. Astrophys. 315, 165. Pety, J., Teyssier, D., Fossé,D., Gerin, M., Roueff, E., Abergel, A., et al.: 2005, Astron. Astrophys. 435, 885. Polehampton, E. T., Baluteau, J.-P., and Ceccarelli, C.: 2002, Astron. Astrophys. 388, L44. Preibisch, T., Ossenkopf, V., Yorke, H.W., et al.: 1993, Astron. Astrophys. 279, 577. Puget, J.-L., Léger, A., and Boulanger, F.: 1985, Astron. Astrophys. 142, L19. Rachford, B. L., Snow, T. P., Tumlinson, J., et al.: 2002, Astrophys. J. 577, 221. Rapacioli, M., Joblin, C., and Boissel, P.: in press, Astron. Astrophys. 429, 193. Roberts, H. and Herbst, E.: 2002, Astron. Astrophys. 395, 233. Roche, P. F., Aitken, D. K., and Smith, C. H.: 1989, Mon. Not. Roy. Astron. Soc. 236, 485. Roelfsema, P., Cox, P., Tielens, A. G. G. M., et al.: 1996, Astron. Astrophys. 315, L289. THE COOL ISM 271

Rosenthal, D., Bertoldi, F., and Drapatz, S.: 2000, Astron. Astrophys. 356, 705. Schutte, W. A., Tielens, A. G. G. M., and Allamandola, L. J.: 1993, Astrophys. J. 415, 397. Schutte, W. A., van der Hucht, K. A., Whittet, D. C. B., et al.: 1998, Astron. Astrophys., 337, 261. Schneider, N., Simon, R., Kramer, C., et al.: 2003, Astron. Astrophys. 406, 915. Sellgren, K.: 1984, Astrophys. J. 277, 623. Sellgren, K., Allamandola, L. J., Bregman, J. D., et al.: 1985, Astrophys. J. 299, 416. Stepnik, B.: 2001, Thesis, Universitéof Paris. Stepnik, B., Abergel, A., Bernard, J.-P., et al.: 2003, Astron. Astrophys. 398, 551. Stepnik, B., Abergel, A., Jones, A. P., et al.: in press, Astron. Astrophys. Szczepanski, J. and Vala, M.: 1993, Astrophys. J. 414, 646. Teixeira, T. C. and Emerson, J. P.: 1999, Astron. Astrophys. 352, 517. Tielens, A. G. G. M., Wooden, D. H., Allamandola, L. J., et al., 1996, Astrophys. J. 461, 210. Timmermann, R., Köster, B., and Stutzki, J.: 1998, Astron. Astrophys. 336, L53. Tóth, L. V., Hotzel, S., Krause, O., et al.: 2000, Astron. Astrophys. 364, 769. Uchida, K. I., Sellgren, K., and Werner, M. W.: 1998, Astrophys. J. 493, L109. Uchida, K. I., Sellgren, K., Werner, M. W., et al.: 2000, Astrophys. J. 530, 817. van Diedenhoven, B., Peeters, E., van Kerckhoven, C., et al., 2004, Astrophys. J. 611, 928. van Dishoeck, E. F.: 2004, Annu. Rev. Astron. Astrophys. 42, 119. Vastel, C., Polehampton, E. T., Baluteau, J.-P., et al.: 2002, Astrophys. J. 581, 315. Verstraete, L., Puget, J. L., Falgarone, E., et al.: 1996, Astron. Astrophys. 315, L337. Verstraete, L., Pech, C., Moutou, C., et al.: 2001, Astron. Astrophys. 372, 981. Weingartner, J. C. and Draine, B. T., 2001a, Astrophys. J. Suppl. 134, 263. Weingartner, J. C. and Draine, B. T., 2001b, Astrophys. J. 548, 296. Weingartner, J. C. and Draine, B. T., 2001c, Astrophys. J. 553, 581. Werner, M. W., Uchida, K. I., Sellgren, K., et al.: 2004, Astrophys. J. Suppl. 154, 309. Wright, C. M., van Dishoeck, E. F., Cox, P., et al.: 1999, Astrophys. J. 515, L29. HIGH EXCITATION ISM AND GAS

ELS PEETERS1,∗, NIEVES LETICIA MARTIN-HERN´ ANDEZ´ 2, NEMESIO J. RODRIGUEZ-FERN´ ANDEZ´ 3 and XANDER TIELENS4 1NASA Ames Research Center, U.S.A. 2Observatoire de Geneve, Switzerland 3LUTH/LERMA – Observatoire de Paris, France 4Kapteyn Astronomical Institute, The Netherlands (∗Author for correspondence: E-mail: [email protected])

(Received 6 October 2004; Accepted in ﬁnal form 3 December 2004)

Abstract. An overview is given of ISO results on regions of high excitation ISM and gas, i.e. H II regions, the Galactic Centre and Supernova Remnants. IR emission due to ﬁne-structure lines, molecular hydrogen, silicates, polycyclic aromatic hydrocarbons and dust are summarised, their diagnostic capabilities illustrated and their implications highlighted.

Keywords: IR, H II regions, galactic centre, supernova remnants, LATEX

1. Introduction

Massive stars provide much of the radiative stellar energy of the Milky Way. Their copious amount of UV radiation has a great impact on the surroundings of the massive star. Indeed, their UV radiation dissociates molecules and dust in the interstellar medium (ISM) and ionises hydrogen, creating large H II regions. In addition, their powerful winds and supernova explosions provide most of the mechanical energy of the Galaxy which dominate the structure of the ISM. Furthermore, the nuclear reactions during their lifetime and death scene synthesis most of the intermediate mass elements and likely the r-process elements. Much of these nucleosynthetic products may condense in the form of small dust grains. Indeed, supernovae may dominate the dust mass budget of the ISM. Because massive stars evolve so fast, they are generally associated with the remnants of their “cradle” and are heavily enshrouded in dust and gas. This dust and gas absorbs most of the stellar luminosity which is then re-emitted in the infrared. Hence, due to the high degree of obscuration, massive stars can be best studied at wavelengths longer than 2 μm. Offering unique combinations of wavelength coverage, sensitivity, and spatial and spectral resolutions in the infrared spectral region, ISO opened the infrared Universe. Its spectral coverage (from 2.3–196 μm) gives for the ﬁrst time access to nearly all the atomic ﬁne-structure and hydrogen recombination lines in the infrared Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, The Netherlands, and the United Kingdom), and with the participation of ISAS and NASA. Space Science Reviews (2005) 119: 273–292 DOI: 10.1007/s11214-005-8070-1 C Springer 2005 274 E. PEETERS ET AL. range. In addition to the atomic lines, ISO revealed the shape and strength of the dust continuum and several emission features. This presented an unprecedented view of the luminous but dusty and obscured Universe, from the Galactic Centre to the most extincted regions of the Milky Way and other galaxies. This chapter reviews the major achievements of ISO on massive stars. Section 2 summarises the results on H II regions, the Galactic Centre is discussed in Section 3. Subsequently, Section 4 highlights the ISO results on supernova remnants.

2. H II Regions

The overall mid-IR (MIR) spectrum of H II regions is dominated by a dust continuum which rises strongly towards longer wavelengths (Figure 1). On top of this continuum, there are a multitude of ﬁne-structure lines and hydrogen recombination lines. The continuum emission is dominated by strong emission features at 3.3, 6.2, 7.7, 8.6 and 11.2 μm, generally attributed to Polycyclic Aromatic Hydro- carbons (PAHs). In some case, broad absorption features due to simple molecules (H2O, CO, CO2) in an icy mantle and/or narrow emission or absorption lines due to gaseous molecules (H2,H2O, CO2, OH, C2H2) are present.

2.1. THE LINE SPECTRUM

2.1.1. The Ionised Gas H II regions are prime targets to derive the present-day elemental abundances of the ISM. They are bright and are characterised by a large number of emission

Figure 1. The ISO-SWS/LWS spectra of a typical H II region (Peeters et al., 2002b). HIGH EXCITATION ISM AND GAS 275 lines. Optical studies of abundance gradients in the Galaxy, however, present the problem that many H II regions are highly obscured by dust lying close to the Galactic plane. The arrival of the Kuiper Airborne Observatory (e.g. Simpson et al., 1995; Afflerbach et al., 1997; Rudolph et al., 1997) and the Infrared Astronomical Satellite (e.g. Simpson and Rubin, 1990) permitted measurements of elemental abundances of embedded compact and ultra-compact H II regions, obscured at optical wavelengths but observable in the infrared, and gave access for the first time to the central regions of the Galaxy. Regarding the determination of elemental abundances, the use of infrared fine-structure lines present clear advantages with respect to the optical lines (e.g. Rubin et al., 1988; Simpson et al., 1995): (1) they are atten- uated much less due to the presence of dust; (2) they are practically insensitive to the precise temperature of the emitting gas since they are emitted from levels with very low excitation energies; and (3) the infrared range is the only wavelength regime + to measure the dominant form of nitrogen in highly excited H II regions (N2 ). ISO provided a unique opportunity to measure the full spectral range from 2.3 to 196 μm of a large number of (ultra)compact H II regions with relatively good spectral and spatial resolution (e.g. Peeters et al., 2002b). Figure 1 shows the combined SWS and LWS spectrum of a typical H II region. The spectrum is dominated by recombination lines of H and fine-structure lines of C, N, O, Ne, S, Ar and Si. The lines of [C II], [O I], and [Si II] are produced by ions with ionisation potentials lower than 13.6 eV and are thus expected to be mostly emitted in the PDR surrounding the H II region (see Section 2.1.2). Two lines are present for some of the ions (O2+,Ne2+ and S2+), providing a handle on the electron density. Typically, [O III] 52, 88 μm densities towards Galactic H II regions range between ∼100 and − 3000 cm 3 (cf. Mart´ın-Hernández et al., 2002). The use of the [S III] 18.7, 33.5 μm and [Ne III] 15.5, 36.0 μm line ratios requires careful aperture corrections. When such corrections are applied, the [S III] and [Ne III] densities agree well with the [O III] densities within the errors (cf. Mart´ın-Hernández et al., 2003). N, Ne, Ar and S are observed in two different ionisation stages, which enormously alleviates the problem of applying ionisation correction factors (cf. Mart´ın-Hernández et al., 2002). The ISO observations of H II regions covered the Galactic plane from the centre to a Galactocentric distance of about 15 kpc, giving thus the possibility of investi- gating trends of relative and absolute elemental abundances across a large part of the Galactic disk. The gradients resulting from the full samples of Mart´ın-Hernández et al. (2002) and Giveon et al. (2002) are log(Ne/H) =−0.039 ± 0.007, log(Ar/H) =−0.047 ± 0.007, log(S/H) =−0.023 ± 0.014 – re-computed by Giveon et al. (2002) applying the same extinction and Te corrections to both data sets – and log(N/O) =−0.056±0.009 dex kpc−1. Figure 2 shows the neon abundance as a function of the distance from the Galactic Centre. The direct observation of two different ionisation stages not only facilitates the determination of elemental abundances, but allows us to probe the ionisation structure of the H II regions and constrain the stellar energy distribution (SED) of the ionising stars. Line ratios such as [N III]/[N II] 57/122 μm, [Ne III]/[Ne II] 276 E. PEETERS ET AL.

Figure 2. Neon abundance as a function of the distance to the Galactic Centre (cf. Mart´ın-Hern´andez et al., 2002a).

15.5/12.8 μm, [S IV]/[S III] 10.5/18.7 μm and [Ar III]/[Ar II] 9.0/7.0 μm probe the ionising stellar spectrum between 27.6 and 41 eV and depend on the shape of the SED and the nebular geometry (see e.g. Morisset et al., 2002, who present a detailed model of the well studied H II region G29.96–0.02 based on their infrared lines and Morisset et al., 2004, who compared predicted ionising spectra against ISO observations of Galactic H II regions). These line ratios are found to correlate well with each other for the large sample of H II regions observed by ISO (cf. Figure 3) and to increase with Galactocentric distance (cf. Giveon et al., 2002; Mart´ın-Hernández et al., 2002). The observed [Ne III]/[Ne II] and [S IV]/[S III] line ratios have been compared with diagnostic diagrams built from extensive photoionisation model grids computed for single-star H II regions using stellar atmosphere models from the WM-Basic code (Pauldrach et al., 2001) where the metallicities of both the star and the nebula have been taking into account (Morisset, 2004). This comparison finds no evidence of a gradient of the effective temperature of the ionising stars with the Galactocentric distance, attributing the observed increase of excitation with distance mainly to the effect of the metallicity gradient on the SED (see also Mart´ın-Hernández et al., 2002 and Mokiem et al., 2004). The relation between the [Ne III]/[Ne II] line ratio and the Ne/H elemental abundance for a combined sample of Galactic and extra-galactic H II regions is shown in the right panel of Figure 3. As it is evident from this figure, a clear correlation between excitation and metallicity exists, albeit with a large scatter. Extended emission of highly ionised species (fine-structure lines of N+,N2+, O2+ and S2+) have been detected by LWS and SWS (Mizutani et al., 2002; Okada et al., 2003; Goicoechea et al., 2004). This reveals the presence of extended, highly ionised gas surrounding H II regions and probably ionised by the central O-star(s). The electron density has been derived from the [O III]52to88μm ratio, and values HIGH EXCITATION ISM AND GAS 277

Figure 3. Left: Relation between the [S IV]/[S III] 10.5/18.7 μm and [Ne III]/[Ne II] 15.5/12.8 μm line ratios for a sample of H II regions (cf. Mart´ın-Hernández et al., 2002b). Indicated by various symbols are Galactic regions at Galactocentric distance Rgal < 7 kpc (solid triangles), Galactic regions at Rgal > 7 kpc (solid circles), LMC regions (open triangles, except 30 Doradus, which is indicated by an open square) and SMC regions (open diamonds). The positions of Sgr A (Lutz et al., 1996) and the Orion nebula (Simpson et al., 1998) are indicated by a solid star and a solid square, respectively. The dotted line is a least squares fit to the data. A typical error bar is given in the upper left corner. The arrow in the lower right corner indicates the correction due to an extinction AK = 2 mag. Right: Relation between the [Ne III]/[Ne II] 15.5/12.8 μm line ratio and the Ne/H elemental abundance for a combined sample of H II regions (Mart´ın-Hernández et al., 2003). Indicated by various symbols are

Galactic H II regions (solid circles), Sgr A (solid star), the orior nebula (star), the Pistol and the Sickle (open and solid diamonds, respectively; Rodr´ıguez-Fern´andez et al., 2001a), a sample of H II regions in M33 (open squares; Willner and Nelson-Patel, 2002), LMC H II regions (reverse open triangles, except 30 Doradus, which is plotted as a plus sign) and 2 regions in the SMC (open triangles). of a few 10 to a few 100 cm−3 have been found. This extended, ionised gas – denser than the galactic warm ionised medium – would represent a new phase of the ISM as stressed by Mizutani et al. (2002). Extragalactic studies of H II regions have been performed in the Magellanic Clouds (cf. Vermeij et al., 2002a; Vermeij and van der Hulst, 2002b) and M33 (Willner and Nelson-Patel, 2002). H II regions in the Magellanic Clouds are characterised by low metallicities and high excitation (see e.g. Figure 3). Towards M33, the distribution of neon abundances as a function of Galactocentric radius is best described as a step gradient, with a slope of −0.15 dex from 0.7 to 4.0 kpc and −0.35 dex from 4.0 to 6.7 kpc.

2.1.2. The Photodissociation Region Photodissociation Regions (PDRs) are regions where FUV (6 < hν<13.6 eV) photons dissociate, ionise and heat neutral atomic/molecular gas surrounding H II regions (Tielens and Hollenbach, 1985; Hollenbach and Tielens, 1999). The gas reaches temperatures in the range 100–1000 K, much warmer than the dust (10– 100 K), and radiates its energy in low energy atomic fine-structure lines, particularly 278 E. PEETERS ET AL. the [O I]63μm, [C II] 157 μm, and [Si II]34μm lines, and in pure rotational molecular hydrogen lines (at wavelengths between 28 and 2 μm). In addition, the strong FUV flux pumps molecular hydrogen molecules into excited vibrational states and their cascade produces strong fluorescent ro-vibrational lines in the near-IR. The dust gives rise to a continuum at long wavelengths (λ>25 μm). In addition, the IR spectrum of PDRs shows broad emission features at 3.3, 6.2, 7.7, 8.6, 11.2, and 12.7 μm due to fluorescent emission of FUV pumped Polycyclic Aromatic Hydrocarbon (PAHs) molecules (see Section 2.2). Here, we will focus on PDRs associated to H II regions; other PDRs are extensively discussed elsewhere in this book. Because of their luke warm temperatures, PDRs were only ‘discovered’ when the IR window was opened by the Kuiper Airborne Observatory in the late 1970s and 1980s, but – because of limited sensitivities – these studies focused on the brightest objects in the sky (the Orion bar; Melnick et al., 1979; Storey et al., 1979; Russell et al., 1980, 1981). The increased sensitivities of the SWS and LWS and, in particular, their wide wavelength coverage allowed for the first time a systematic study of the properties of PDRs. Much of this data are however still resting in the archives. Except for the some of the ionic fine-structure lines and the H I recombination lines, much of the IR emission characteristics of H II regions (Figure 1), originate in the associated PDRs. The SWS is particularly particularly suited for observations of the pure rotational lines of molecular hydrogen. These lines provide a direct handle on the physical properties of the emitting gas. Because of the small Einstein A’s of the associated quadrupole transitions, the lowest levels of H2 have very low critical densities and hence their populations are in local thermodynamic equilibrium for typical PDR conditions (Burton et al., 1992). A rotational diagram provides then directly the temperature and column density of the emitting gas. This is illustrated in Figure 4, left panel, for the PDR associated with the blister H II region, S140 (Timmermann et al., 1996). The optically bright rim is produced by the ionisation of a dense clump in the molecular cloud, L1202/1204 by the nearby B0.5 star, HD 211880 (Blair et al., 1978; Evans et al., 1987). The SWS spectrum of this source shows pure rotational H2 lines up to 0 − 0 S(9). The column density and temperature derived 20 −2 from these observations are N 10 cm and Trot = 500 K (Timmermann et al., 1996). Similar studies of the pure rotational lines in the spectra of S106 and Cep A resulted in excitation temperatures of 500 and 700 K (van den Ancker et al., 2000; Wright et al., 1996). In addition, SWS has observed several ro-vibrational lines in S140 belonging to the v = 1–0, 2–1, and 3–2 series. These levels are populated by FUV pumping. Figure 4, left panel, compares the observations with a detailed model for the collisional and FUV-pumping excitation of H2 in this source (Timmermann et al., 1996). Good agreement is obtained for a density of 104 cm−3 in the PDR and an incident FUV flux of 400 Habing. These values are consistent with estimates based upon pressure equilibrium across the ionisation front coupled with the ionised gas density in the associated bright rim and the properties of the illuminating star. HIGH EXCITATION ISM AND GAS 279

Figure 4. Left: The rotational diagram of molecular lines observed towards the bright rim cloud, S140 (Timmermann et al., 1996). Observed (filled symbols) and model calculated (open symbols) column densities (ratioed to the statistical weight of the level involved) are plotted as a function of the energy of the upper level. A Boltzmann distribution will show a straight line on this plot. The low lying pure rotational lines are consistent with gas at a kinetic temperature of 500 K. The higher excitation temperature of the vibrationally excited levels at high energies indicate the increased importance of FUV pumping for the populations of these levels. The models results are calculated for an incident FUV field of 400 times the average interstellar radiation field and a gas density of 104 cm−3.A foreground extinction of 0.3 magnitudes in the K band has been adopted. To indicates the adopted peak temperature of the gas in the PDR. Right: A diagnostic diagram for PDRs based upon the intensity ratio of the [C II] 158 μm and the [O I]63μm lines and the overall cooling efficiency of the gas in the PDR. The lines present the results of detailed model calculations for a range in densities and incident FUV fields (Kaufman et al., 1999). The data are taken from LWS observations of a sample of H II regions (Peeters et al., 2002b; Peeters et al., 2005, in preparation; Vermeij et al., 2002a).

While this agreement between observations and models of H2 excitation is very comforting (Figure 4, left panel), PDR models have a problem in matching the observed kinetic temperatures and column densities derived from the low lying pure rotational lines. In particular, for such low FUV fields and densities as derived for S140, the calculated gas temperatures in the PDR are much lower than the observed 500 K (e.g. Draine and Bertoldi, 1999). This problem seems to be more general, although a detailed comparison of observations and models for regions spanning a wide range in physical properties is still wanting. Nevertheless, it seems that PDR models may be missing an essential part of the energetic coupling between the gas and the illuminating stellar FUV radiation field. The far-IR atomic fine-structure lines provide important probes of the physical conditions in PDRs. However, at present, little has been done to harvest the data provided by the LWS. The [O I]63 μm and [C II] 157 μm lines are the dominant coolants of PDRs (Tielens and Hollenbach, 1985). Hence, the summed flux of these lines measures the total heating of PDR gas by the FUV radiation field. Compar- ison with the total incident FUV flux, as measured by the far-IR dust continuum 280 E. PEETERS ET AL. provides a measure of the heating efficiency of PDRs. The levels involved in these two transitions have very different critical densities (∼105 and 3 × 103 cm−3) and excitation energies (92 and 228 K). Not coincidentally, these values are precisely the range expected for gas in PDRs and, hence, the ratio of these lines has evolved to a major diagnostic of the properties of PDRs. This is illustrated in Figure 4, right panel, where the results of PDR models for this line ratio are plotted against the calculated ratio of the cooling lines as a function of density and incident FUV field (Kaufman et al., 1999). Over much of the relevant range of these observables, the model result segregate out well and hence this diagram is very useful for the analysis of the conditions in PDRs. Observations towards a sample of H II regions are also shown in Figure 4, right panel. The conditions in these regions span a range in density, 30–3×103 cm−3, and incident FUV fields, 3 × 102–3 × 104 Habing. These values are within the range expected for PDRs associated with (evolved) compact H II regions. In addition, there seems to be a trend in the distribution of the observed points with location in the galaxy or, equivalently, metallicity. That is regions with lower metallicity seem to be characterised with a higher heating efficiency and lower [C II]/[O I] ratio than regions with higher metallicity. If the [O I] line is the dominant cooling line, a decrease in the metallicity is expected to result in a decreased [C II]/[O I] ratio (Tielens and Hollenbach, 1985). However, the decrease in the heating rate with increasing metallicity is not expected. Further analysis of these types of observations is clearly warranted. The fine-structure lines observed towards several individual high mass star forming regions have been analysed in detail. Analysis of the observations of W49N 5 show that the PDR is illuminated by an intense FUV field (G0 = 3 × 10 ). The density and temperature are, however, quite moderate (n = 104 cm−3, T = 130 K). Reflecting this high FUV field and low density, the observed heating efficiency is comparatively low, ∼10−4 (Vastel et al., 2001). Towards S106IR, a hot (200–500 K) dense (n > 3 × 105 cm−3) gas component is present. Cooler gas is associated with the bulk of the emission of the molecular cloud and is characterised by two emission peaks which have densities of 105 cm−3 and are illuminated by a radiation field, 2 3.5 G0, ranging from 10 to 10 (Schneider et al., 2003). The H II region S 125 is modelled self-consistently with a two-dimensional geometrical blister model. In order to fit the spatial profile, a systematic increase of the gas temperature along the PDR boundary with decreasing distance from the ionising star is necessary, which is not readily understood within present-day PDR models (Aannestad and Emery, 2003). Despite the rich concentration of massive stars in the stellar cluster Trumpler 14, in the Carina nebula, the physical conditions in the associated PDR are much less extreme than for Orion or M17 (Brooks et al., 2003). The data for the PDR in NGC 2024, on the other hand, are consistent with emission from dense (n 106 cm−3), coolish (T 100 K) clumps (Giannini et al., 2000). Finally, the lowest pure rotational J = 1 − 0 transition of the HD molecule at 112 μm was detected towards the Orion Bar using the LWS on ISO (Wright HIGH EXCITATION ISM AND GAS 281 et al., 1999). The observations imply a total HD column density of (3 ± 0.8) × 1017 cm−2 for adopted temperatures in the range 85–300 K. This corresponds to an elemental D/H abundance ratio of (2 ± 0.6) × 10−5, comparable to that derived from D I and H I ultraviolet absorption measurements in the solar neighbourhood.

2.2. PAHS AND DUST

Besides the well-known UIR bands at 3.3, 6.2, 7.6/7.8, 8.6, 11.2 and 12.7 μm, and very broad structures present underneath these bands, many discrete weaker bands and subcomponents can be found at 3.4, 3.5, 5.2, 5.7, 6.0, 6.6, 7.2–7.4, 8.2, 10.8, 11.0, 12.0, 13.2 and 14.5 μm. ISO extended the observable wavelength range and found there new feature at 16.4 μm and 17.4 μm as well as a weak, variable, emission plateau between 15 and 20 μm (Moutou et al., 2000; Van Kerckhoven et al., 2000; Van Kerckhoven, 2002). It should be emphasized that not all sources show all these emission features at the same time. ISO also allowed, for the first time, a systematic analysis of the UIR bands in a wide variety of environments. It is now firmly established that the detailed characteristics (intensity, peak position, profile) of the UIR features vary from source to source and also spatially within extended sources (for a recent review, see Peeters et al., 2004b). ISO-SWS spectra indicated that the UIR band profiles and positions of all H II regions are equal to each other and to those observed in Reflection Nebulae (RNe), non-isolated Herbig AeBe stars and the ISM, while they are clearly distinct from those found around evolved stars and isolated Herbig AeBe stars (Figure 5, Peeters et al., 2002a; van Diedenhoven et al., 2004). However, spatial variations within RNe (Bregman and Temi, 2005) and evolved stars (Kerr et al., 1999; Song et al., 2004; Miyata et al., 2004) are observed suggesting that in all sources profile variations occur on a small spatial scale. In addition – despite of what the integrated spectra of sources might suggest – the profiles are not unique to certain object types. In contrast, their relative strength in H II regions varies both from source to source and within sources (Figure 5). Integrated spectra of sources show variations in the relative strength of the CC modes (6–9 μm range) relative to the CH modes (at 3.3 and 11.2 μm), as is the 11.2/12.7 band ratio while the 3.3 μm feature correlates with the 11.2 μm feature and the 6.2 μm feature with the 7.7 μm feature (Figure 5, Roelfsema et al., 1996; Verstraete et al., 1996; Hony et al., 2001; Vermeij et al., 2002c). These variations are directly related to object type (Hony et al., 2001) and to the local environment (Figure 5, Verstraete et al., 1996; Vermeij et al., 2002c). Within extended sources, additional variations are found, nicely demonstrated by that of the 6.2/7.7 band ratio in S106 (Joblin et al., 2000). Within a single object, variations in G0 are important in driving the UIR emission spectrum (Onaka, 2000; Bregman and Temi, 2005). However, the source to source variation does not seem to follow G0 (Verstraete et al., 2001; Hony et al., 282 E. PEETERS ET AL.

Figure 5. Left: Source to source variations for the 3–12 μm UIR bands. Profile A represent H II regions, Reflection Nebulae and the ISM, profile B isolated HAeBe stars and most PNe and profile C 2 post-AGB stars (Peeters et al., 2002a; van Diedenhoven et al., 2004). Right: PAH band strength ratios for H II regions. Galactic sources are shown as (), non-30 Dor sources as (), 30 Dor pointings as () and the SMC B11 molecular cloud (Reach et al., 2000a) by an asterisk (Vermeij et al., 2002).

2001; Peeters et al., 2002a, 2004c; Vermeij et al., 2002c; van Diedenhoven et al., 2004). The observed variations of the UIR bands provide direct clues to the physical and chemical properties of their carriers, the local physical conditions and/or the local history of the emitting population. These spectral variations have been attributed to a range of different physical or chemical characteristics of the carriers, including charge state, anharmonicity, different subcomponents with variable strength, family of related species with varying composition, substituted/complexed PAHs, isotope variations, clustering, molecular structure, molecular composition etc. (Verstraete et al., 1996, 2001; Joblin et al., 2000; Hony et al., 2001; Pech et al., 2002; Peeters et al., 2002a; Wada et al., 2003; Song et al., 2004; Bregman and Temi, in press; Hudgins and Allamandola, 2004; van Diedenhoven et al., 2004). In particular, the variations in the 11.2/6.2 and 3.3/6.2 μm ratios likely reflect variations in the average ionisation state of the emitting PAHs. In contrast, the variations in the peak position of the 6.2 μm band point towards variations in the molecular structure of the emitting species perhaps related to the incorporation of N into the ring structure of the PAHs. Since the UIR bands are omnipresent, they serve as important probes of the different emission zones. For example, the presence and strength of the UIR bands are generally thought to trace star formation and so they are used as qualitative and quantitative diagnostics of the physical processes powering Galactic nuclei (see this HIGH EXCITATION ISM AND GAS 283 issue). In addition, deuterated PAHs are tentatively detected at 4.4 and 4.65 μmat high abundances in the Orion Bar and M17 (Peeters et al., 2004a). If born out, they can serve as probes of PAH D/H ratio in various regions. Several studies were devoted to the spatial distribution of the different emission components (Cesarsky et al., 1996; Verstraete et al., 1996; Pilbratt et al., 1998; Créte et al., 1999; Klein et al., 1999; Zavagno and Ducci, 2001; Urquhart et al., 2003; Verma et al., 2003). It is overwhelmingly clear that the UIR bands peak in the PDR, though closer to the ionised star with respect to H2, while the strong dust continuum peaks towards the H II region. Inside the H II region, where UIR bands are weak or absent, remains a broad underlying emission band which is attributed to very small grains (VSG, Cesarsky et al., 1996; Jones et al., 1999). The spatial distribution of the UIR bands is also displaced from that of the ERE emission in Sh 152 (Darbon et al., 2000), similar to the Red Rectangle (Kerr et al., 1999). It is clear that the carrier of the UIR bands is not also responsible for the ERE. Moreover, because the ERE emission can sometimes be traced to inside the ionised gas and recalling that the PAHs do not survive in the H II region, the carrier of the ERE is unlikely to be molecular in origin. The dust/UIR band emission has been modeled for several sources. Créte et al. (1999) modeled M17 using a two-component model composed of PAHs as the carrier of the UIR bands and VSG responsible for the “warm” continuum and a two-component model composed of coal dust emission and big grain emission. In the frame of the coal dust model, the observations required nano-sized, transiently heated particles. The IR emission in the M17-SW H II region and Orion is modeled with a mixture of amorphous carbons, silicates and possibly crystalline silicates (see below, Jones et al., 1999; Cesarsky et al., 1996). PAHs seem to be depleted inside the H II region consistent with their destruction in the intense radiation field of M17 and Orion (Jones et al., 1999; Cesarsky et al., 2000; Créte et al., 1999). Modelling of the FIR emission of Sh 125 also indicates the dust to be severely depleted in the H II region while normal dust/gas ratio is observed in the PDR, but with a major fraction in the form of VSG (Aannestad and Emery, 2001). ISO also increased the number of H II regions showing silicate in emission (Cesarsky et al., 2000; Peeters et al., 2002b; Peeters, 2002) and detected several new dust features:

(1) A very strong, broad 8.6 μm band is detected towards the M17-SW H II region and three compact H II regions, clearly distinct from the “classical” 8.6 μmPAH band and originating inside the H II regions (Roelfsema et al., 1996; Cesarsky et al., 1996; Verstraete et al., 1996; Peeters et al., 2002a; Peeters et al., 2005). Detailed analysis reveals the feature having a proﬁle peaking at 8.9 μm with a FWHM ∼1 μm and likely not being related to PAH emission (Peeters et al., 2005). (2) A broad 22 μm feature is detected in the M17-SW H II region (Jones et al., 1999), the Carina Nebula and two starburst galaxies (Chan and Onaka, 2000). 284 E. PEETERS ET AL.

The feature shape is similar to that of the 22 μm band observed in Cassiopeia A indicating that supernovae are probably the dominant production sources (Chan and Onaka, 2000). (3) Several (possible) detections of crystalline silicates associated with H II regions are reported, i.e. (i) emission at 34 μm in the M17 H II region attributed to Mg-rich crystalline olivine (Jones et al., 1999); (ii) emission bands at 9.6 μm in θ 2 Ori A and at ∼15–20 μm, ∼20–28 μm and longward of 32 μm in the Orion H II region attributed to crystalline Mg-rich silicate fosterite (Cesarsky et al., 2000). The latter emission bands are however a combination of instrumental effects and the PAH 15–20 μm plateau (Kemper et al., in preparation); (iii) a band at 65 μm in the Carina Nebula and Sh 171 attributed to diopside (Ca- bearing crystalline silicate – Onaka and Okada, 2003). An upper-limit for the crystallinity of the silicates in the ISM is derived being a few percents (Jones et al., 1999), less than 5% (Li and Draine, 2001) and less than 0.4% (Kemper et al., 2004). (4) To end, a broad emission feature is found at 100 μm in the Carina Nebula and Sh 171, possibly due to carbon onion grains (Onaka and Okada, 2003). A similar broad feature around 90 μm is reported in evolved stars and a low mass protostar, attributed to calcite, a Ca-bearing carbonate mineral (Ceccarelli et al., 2002; Kemper et al., 2002).

3. The Galactic Centre

The 500 central pc of the Galaxy (hereafter Galactic center, GC) are an extended source of emission at all wavelengths from radio to γ -rays (see Mezger et al., 1996 for a review). The non-thermal radio emission as well as γ - and X-rays observations show the presence of a very hot plasma (107–108 K) and a recent episode of nucleosynthesis (∼106 year ago) that suggest an event of violent star formation in the recent past. On the other hand, the cold medium in the GC is mainly known by radio observations of the molecular gas. In between these hot and cold phases, there is a warm (few hundred K) neutral medium and a thermally ionised medium. ISO has made important contributions to our understanding of the heating of the warm neutral gas and the properties of the ionised gas in the GC. Before ISO, the ionised gas has been mainly studied by radio continuum and hydrogen radio recombination lines observations. They showed an extended diffuse ionised gas component and a number of discrete H II regions, most of them associated with the Sgr A, B, ... complexes. Only the most prominent H II regions (Sgr A) and ionised nebulae (the Sickle) had been observed in infrared ﬁne-structure lines. ISO observations of ﬁne-structure lines allowed studying the properties of the ionised gas and the ionising radiation over the whole GC region. In the vicinity of Sgr A∗, where the ionising source is the central stellar cluster, Lutz et al. HIGH EXCITATION ISM AND GAS 285

(1996) derived an effective temperature of the ionising radiation of ∼35000 K and Rodr´ıguez-Fernández et al. (2001a) derived a similar effective temperature in the Radio Arc region. They also showed unambiguously that the Quintuplet and the Arches clusters ionise the Sickle nebula and the thermal filaments that seem to connect Sgr A to the Radio Arc. Indeed, the combined effects of both clusters ionise a region of more that 40 × 40 pc2. Their photoionisation model simulations showed that, in order to explain the long-range effects of the radiation, the ISM that surrounds the clusters must be highly inhomogeneous. A similar scenario has been invoked by Goicoechea et al. (2004) to explain the extended ionised gas component in the Sgr B molecular complex. Indeed, the whole GC region is permeated by relatively hot (∼35000 K) but diluted ionising radiation field (Rodr´ıguez-Fernández and Mart´ın-Pintado,2005). The line ratios ([Ne III]/[Ne II] 15.5/12.8 μm, [N III]/[N II]57/122 μm ...) measured in the GC are similar to those found in low excitation starburst galaxies as M 83 or IC 342 but somewhat lower than those measured in other starburst galaxies as NGC 3256 or NGC 4945. However, the low ratios do not imply intrinsic differences in the star formation activity as a lower upper mass cutoff of the initial mass function. On the contrary, they are probably due to starburst of short duration that produce hot massive stars but whose age quickly softens the radiation (Thornley et al., 2000). The ISO observations of the GC are consistent with a burst of star formation less than 7 Myr ago. On the other hand, in the GC there is a widespread neutral gas component with temperatures of ∼200 K and without associated warm dust. Before ISO, this warm neutral gas had only been studied by radio observations of symmetric top molecules like NH3. How this gas is heated has been a long-standing problem. Radiative heating mechanisms are usually ruled out due to the apparent lack of continuum sources and the discrepancy of gas and dust temperatures. ISO was perfectly suited to study the heating of this gas. First, Rodr´ıguez-Fernández et al. (2001b) have measured for the first time the total amount of warm gas in the GC clouds by observing H2 pure rotational lines. These lines trace gas with temperature from 150 K to 500 K. The column density of gas with a temperature of 150 K is a few 1022 cm−2 and, on average, it is around 30% of the total gas column density. Second, ISO has observed the main coolants of the neutral gas (the [C II] 158 and [O I]63 μm lines) with temperatures of a few hundred K and the full continuum spectrum of the dust emission around the maximum. The dust emission can be characterized by two temperature components, one with a temperature of ∼15 K and a warmer one with a temperature that varies from source to source from ∼25 to ∼45 K (Lis et al., 2001; Rodr´ıguez-Fernández et al., 2004; Goicoechea et al., 2004). A detailed comparison of the lines and continuum emission with shocks and PDR models show that the fine structure lines and the H2 emission from excited levels arise in a PDR with a far-ultraviolet incident radiation field 103 times higher than the average local interstellar radiation field and a density of 103 cm−3 (Figure 6; Rodr´ıguez-Fernández et al., 2004). The warm dust emission and 286 E. PEETERS ET AL.

Figure 6. Intensities of two fine-structure lines and two H2 pure-rotational lines versus the far- ultraviolet incident flux derived for the sources observed by Rodr´ıguez-Fernández et al. (2004). For comparison, the PDR model predictions (Tielens and Hollenbach, 1985) for different densities are also shown. the fine-structure lines probably arise in PDRs located in the interface between the ionised gas discussed above and the cold molecular gas. Rodr´ıguez-Fernández et al. (2004) also showed that an important fraction (∼50%) of the [C II] 158 and [Si II]35 μm can arise from the diffuse ionised gas instead of the PDR itself. Regarding the warm H2, only 10–20% of the total column density of gas with temperature of 150 K can be accounted for in the widespread PDRs scenario. The most likely heating mechanism for the bulk of the warm neutral gas is the dissipation of magnetohydrodynamic turbulence or low velocity shocks. Spectroscopic imaging/mapping will be useful to investigate the interplay between the different ISM phases in the GC. This could be done with the IRS instrument onboard the Spitzer Space Telescope. In addition, SOFIA and Herschel will be able to map the ionised and the warm neutral gas with high spectral resolution, which is needed due to the crowded velocity fields in the GC region.

4. Supernova Remnants (SNRs)

As the prototype of O-rich young SNRs, Cassiopeia A is well-studied by ISO (Lagage et al., 1996; Tuffs et al., 1997; Arendt et al., 1999; Douvion et al., 1999; Douvion et al., 2001b). Several ﬁne-structure lines are observed at high velocity with [OIV] being the most prominent (e.g. Figure 7). However, no iron emission is observed (Arendt et al., 1999). These ﬁne-structure lines are spatially correlated HIGH EXCITATION ISM AND GAS 287

Figure 7. Contour maps of the [Ne II] 12.8 μm line emission (dotted contours) and of the 9.5 μm silicate dust emission (full contours) plotted onto an optical image of the Cassiopeia-A supernova remnant. Examples of the variety seen in the MIR spectra are shown on the right (Taken from Douvion et al., 1999). with the Fast Moving Knots (FMKs) known to be made of nuclear burning products from the progenitor star. While fainter diffuse dust emission is present, the brightest continuum emission is spatially correlated with the fine-structure lines indicating that dust has freshly condensed in the FMKs. The high dust temperature derived is consistent with emission from collisionally heated dust in the post-shock region of FMKs. A new 22 μm feature is found and identified with Mg protosilicates (Arendt et al., 1999). However, Mg protosilicates give rise to a 10 μm feature quite different from the observed one (Douvion et al., 2001b). These authors fitted a 6–30 μm spectrum of Cas A with SiO2, MgSiO3,Al2O3. In addition to dust continuum emission, silicate emission is present at several positions within Cas A (Figure 7). The presence of silicate and the presence of neon is anti correlated in many knots while argon and sulfur are present in most of the neon knots (Figure 7), suggesting a different degree of mixing occurred (Douvion et al., 1999). These results also indicate an incomplete condensation of silicate elements. If born out for other core collapse SNe, these SNe can no longer be considered the dominant source of silicates (Douvion et al., 1999). Similar conclusion was reached by Tuffs et al. (1997) based upon the derived dust mass being less than a few hundredths of a solar mass. In contrast to Cas A, the dust emission in the Kepler and Tycho SNRs is spatially correlated with Hα emission and therefore attributed to shocked circumstellar (CS) 288 E. PEETERS ET AL. material and CS or interstellar (IS) material respectively and not to SN condensates. Instead of dust signatures, the Crab SNR is dominated by synchrotron radiation with a spectral index of −0.3 to −0.65. Fine-structure lines of [Ne II] and [Ar II] are observed for the Kepler SNR and of [Ne II], [Ne III], [Ne V], [Ar II], [Ar III], [S IV], [Ni II] for the Crab SNR. The [Ne III] emission in the Crab SNR shows the same filamentary structure as seen in optical lines, known to be composed predominately of SN ejecta (Douvion et al., 2001a). The IR spectrum of RCW103, a young and fast SNR, is dominated by prominent lines from low-excitation species allowing estimates of electron density, temperature and abundances. The ionic lines are consistent with post-shock emission. The absence of significant emission from the pre-shock region suggests that shock models may overestimate the importance of the precursor. All – but one – abundances are solar and hence the emitting gas is ISM material in which the dust grains have been destroyed by the shock front (Oliva et al., 1999b). The interaction of SNRs with molecular clouds has been extensively studied with ISO. The IR spectra of the SNRs 3C 391, W44 and W28 exhibit fine-structure line emission, dust emission along the line of sight and shocked dust emission shortwards of 100 μm (Reach and Rho, 1996, 2000; Reach et al., 2002). In addition, shock-excited FIR emission of H2O, CO and OH emission is detected in 3C 391 (Reach and Rho, 1998). These IR observations provide evidence for multiple pre- and post-shock conditions which are found to be spatially separated. Indeed, ISOCAM-CVF observations clearly resolves the ionic emission from the molecular emission tracing moderate-density to high density regions in 3C 391. Abundance analysis reveal partial dust destruction consistent with theoretical models of grain destruction. Similar, spectral variations are observed spatially within IC443 (Cesarsky et al., 1999; Oliva et al., 1999a; Rho et al., 2001). The north- easthern filamentary structure is dominated by prominent [Ne II], [Fe II], [S III], [O IV] and [O I] line emission while the southern clumpy region is dominated by H2 emission indicating different types of shock occurring in different conditions. Both regions can alone account for virtually all the observed IRAS flux in the 12 and 25 μm bands – likely due to the limited ISO field-of-view and limited IRAS spatial resolution. Despite this, the unusually blue IRAS colors of IC443 reflect line (molecular and/or atomic) contamination instead of a large population of small grains. The late emission of SN1987A probes directly the elemental abundances deep in the stellar ejecta and as such constrains models of SN explosion and the explosive nucleosynthesis. Particularly important for this are the 56Ni, 57Ni and 44Ti masses. The former two are known, the latter is more uncertain and can be derived from the [Fe I]24.05 μm and [Fe II] 25.99 μm fine-structure lines. These lines are however not detected in the ISO spectra of SN1987A indicating – based upon time-dependent −4 models – an upper limit of 1.1 × 10 M (Lundqvist et al., 1999, 2001). The central region around the SN position is resolved in the MIR and consistent with collisionally heated grains in the shocked CS gas around SN1987A (Fischera et al., HIGH EXCITATION ISM AND GAS 289

2002a). The MIR spectrum can be modeled by a dust mixture of silicate–iron or silicate–graphite or pure graphite grains and indicates a low dust-to-gas ratio (Fischera et al., 2002a). Hence, observations of ionic lines allow to diagnose the physical conditions of the medium in which the SNR blast wave is propagating, the shock kinematics, the elemental abundances and hence also dust destruction. ISO provided evidence that dust can form in supernovae and dust can be (partially) destroyed in supernovae.

Acknowledgements

EP acknowledges the support of the National Research Council. NJR-F acknowledges support by a Marie Curie Fellowship of the European Community under contract number HPMF-CT-2002-01677.

References

Aannestad, P. A. and Emery, R. J.: 2001, Astron. Astrophys. 376, 1040. Aannestad, P. A. and Emery, R. J.: 2003, Astron. Astrophys. 406, 155. Afflerbach, A., Churchwell, E., and Werner, M. W.: 1997, Astrophys. J. 478, 190. Arendt, R. G., Dwek, E., and Moseley, S. H.: 1999, Astrophys. J. 521, 234. Blair, G. N., Evans, N. J. II, vanden Bout, P. A., and Peters, W. L., III: 1978, Astrophys. J. 219, 896. Bregman, J. and Temi, P.: 2005, Astrophys. J. 621, 831. Brooks, K. J., Cox, P., Schneider, N., Storey, J. W. V., Poglitsch, A., et al.: 2003, Astron. Astrophys. 412, 751. Burton, M. G., Hollenbach, D. J., and Tielens, A. G. G. M.: 1990, Astrophys. J. 365, 620. Ceccarelli, C., Caux, E., Tielens, A. G. G. M., Kemper, F., Waters, L. B. F. M., and Phillips, T.: 2002, Astron. Astrophys. 395, L29. Cesarsky, D., Lequex, J., Abergel, A., Perault, M., Palazzi, E., et al.: 1996, Astron. Astrophys. 315, L309. Cesarsky, D., Cox, P., Pineau des Forets, G., van Dishoeck, E. F., Boulanger, F., and Wright, C. M.: 1999, Astron. Astrophys. 348, 945. Cesarsky, D., Jones, A. P., Lequex, J., and Verstraete, L.: 2000, Astron. Astrophys. 358, 708. Chan, K.-W., and Onaka, T.: 2000, Astrophys. J. 533, L33. Créte,E., Giard, M., Joblin, C., Vauglin, I., Léger, A., and Rouan, D.: 1999, Astron. Astrophys. 352, 277. Darbon, S., Zavagno, A., Perrin, J.-M., Savine, C., Ducci, V.,and Sivan, J.-P.: 2000, Astron. Astrophys. 364, 723. Draine, B. T. and Bertoldi, F.: 1999, The Universe as Seen by ISO, ESA-SP 427, p. 553. Douvion, T., Lagage, P. O., and Cesarsky, C. J.: 1999, Astron. Astrophys. 352, L111. Douvion, T., Lagage, P. O., Cesarsky, C. J., and Dwek, E.: 2001a, Astron. Astrophys. 373, 281. Douvion, T., Lagage, P. O., and Pantin, E.: 2001b, Astron. Astrophys. 369, 589. Evans, Neal, J., II, Kutner, M. L., and Mundy, Lee, G.: 1987, Astrophys. J. 323, 145. Fischera, Jg., Tuffs, R. J., and Volk, H. J.: 2002a, Astron. Astrophys. 386, 517. Fischera, J.g., Tuffs, R. J., and Volk, H. J.: 2002b, Astron. Astrophys. 395, 189. 290 E. PEETERS ET AL.

Fuente, A., Mart´ın-Pintado,J., Rodr´ıguez-Fernández,N. J., Rodr´ıguez-Franco,A., de Vicente, P., and Kunze, D.: 1999, Astron. Astrophys. 518, L45–L48. Giannini, T., Nisini, B., Lorenzetti, D., Di Giorgio, A. M., Spinoglio, L., et al.: 2000, Astron. Astrophys. 358, 310. Giveon, U., Sternberg, A., Lutz, D., Feuchtgruber, H., and Pauldrach, A. W. A.: 2002, Astrophysical Journal 566, 880. Giveon, U., Morisset, C., and Sternberg, A.: 2002, Astron. Astrophys. 392, 501. Goicoechea, J. R., Rodr´ıguez-Fernández,N. J., and Cernicharo, J.: 2004, Astrophys. J. 600, 1214. Gruenwald, R. B. and Viegas, S. M.: 1992, Astrophys. J., Suppl. Ser. 78, 153–178. Hollenbach, D. J. and Tielens, A. G. G. M.: 1999, Rev. Mod. Phys. 71, 173. Hony, S., Van Kerckhoven, C., Peeters, E., Tielens, A. G. G. M., Hudgins, D. M., and Allamandola, L. J.: 2001, Astron. Astrophys. 370, 1030. Hudgins, D. M. and Allamandola, L. J.: 2004, Astrophys. of Dust/Astron. Soc. Pacif. Conference Series, vol. 309, 665. Joblin, C., Abergel, A., Bregman, J., d’Hendecourt, L., Heras, A. M., et al.: 2000, ISO Beyond the Peaks, ESA-SP 466. Jones, A. P., Frey, V., Verstraete, L., Cox, P., and Demyk, K.: 1999, The universe as seen by ISO, ESA-SP 427, p 679. Kaufman, M. J., Wolfire, M. G., Hollenbach, D. J., and Luhman, M. L.: 1999, Astrophys. J. 527, 795. Kemper, F., Molster, F. J., Jäger, C., and Waters, L. B. F. M.: 2002, Astron. Astrophys. 394, 679. Kemper, F., Vriend, W. J., and Tielens, A. G. G. M.: 2004, Astron. Astrophys. 609, 826. Kerr, T. H., Hurst, M. E., Miles, J. R., and Sarre, P. R.: 1999, Mon. Not. R. Astron. Soc. 303, 446. Klein, R., Henning, T., and Cesarsky, D.: 1999, Astron. Astrophys. 343, L53. Li, A. and Draine, B. T.: 2001, Astrophys. J. 550, L213. Lagage, P. O., Claret, A., Ballet, J., Boulanger, F., Cesarsky, C. J., et al.: 1996, Astron. Astrophys. 315, L274. Lis, D. C., Serabyn, E., Zylka, R., and Li, Y.: 2001, Astrophys. J. 550, 761. Lundqvist, P., Sollerman, J., Kozma, C., Larsson, B., Spyromilio, J., et al.: 1999, Astron. Astrophys. 347, 500. Lundqvist, P., Kozma, C., Sollerman, J., and Fransson, C.: 2001, Astron. Astrophys. 374, 629. Lutz, D., Feuchtgruber, H., Genzel, R., Kunze, D., Rigopoulou, D., et al.: 1996, Astron. Astrophys. 315, L269. Mart´ın-Hernández,N. L., van der Hulst, J. M., and Tielens, A. G. G. M.: 2003, Astron. Astrophys. 407, 957. Mart´ın-Hernández,N. L., Vermeij, R., Tielens, A. G. G. M., van der Hulst, J. M., and Peeters, E.: 2002, Astron. Astrophys. 389, 286. Mart´ın-Hernández,N. L., Peeters, E., Morisset, C., Tielens, A. G. G. M., Cox, P., et al.: 2002, Astron. Astrophys. 381, 606. Mathis, J. S. and Rosa, M. R.: 1991, Astron. Astrophys. 245, 625. Melnick, G., Gull, G. E., and Harwit, M.: 1979, Astrophys. J. 227, L29. Mezger, P. G., Duschl, W. J., and Zylka, R.: 1996, Astron. Astrophys. Rev. 7, 289. Mizutani, M., Onaka, T., and Shibai, H.: 2002, Astron. Astrophys. 382, 610. Miyata, T., Kataza, H., Okamoto, Y. K., Onaka, T., Sako, S., et al.: 2004, Astron. Astrophys. 415, 179. Mokiem, M. R., Mart´ın-Hernández,N. L., Lenorzer, A., de Koter, A., and Tielens, A. G. G. M.: 2004, Astron. Astrophys. 419, 319. Morisset, C., Schaerer, D., Mart´ın-Hernández,N. L., Peeters, E., Damour, F., et al.: 2002, Astron. Astrophys. 386, 558. Morisset, C.: 2004, Astrophys. J. 601, 858. Morisset, C., Schaerer, D., Bouret, J.-C., and Martins, F.: 2004, Astrophys. J. 415, 577. HIGH EXCITATION ISM AND GAS 291

Moutou, C., Verstraete, L., L’eger, A., Sellgren, K., and Schmidt, W.: 2000, Astron. Astrophys. 354, L17. Oliva, E., Moorwood, A. F. M., Drapatz, S., Lutz, D., and Sturm, E.: 1999b, Astron. Astrophys. 343, 943. Oliva, E., Lutz, D., Drapatz, S., and Moorwood, A. F. M.: 1999a, Astron. Astrophys. 341, L75. Okada, Y., Onaka, T., Shibai, H., and Doi, Y.: 2003, Astron. Astrophys. 412, 199. Onaka, T.: 2000, Adv. Space Res. 25, 2167. Onaka, T. and Okada, Y.: 2003, Astrophys. J. 585, 872. Pauldrach, A. W. A., Hoffmann, T. L., and Lennon, M.: 2001, Astron. Astrophys. 375, 161. Pech, C., Joblin, C., and Boissel, P.: 2002, Astron. Astrophys. 388, 639. Peeters, E.: 2002, PhD thesis, University of Groningen, The Netherlands. Peeters, E., Allamandola, L. J., Bauschlicher, C. W., Hudgins, D. M., Sandford, S. A., and Tielens, A. G. G. M.: 2004a, Astrophys. J. 604, 252. Peeters, E., Allamandola, L. J., Hudgins, D. M., Hony, S., and Tielens, A. G. G. M.: 2004b, Astrophys. of Dust/Astron. Soc. Pacif. Conference Series, vol. 309, 141. Peeters, E., Hony, S., Van Kerckhoven, C., Tielens, A. G. G. M., Hudgins, D. M., Allamandola, L. J., and Bauschlicher, C. W.: 2002a, Astron. Astrophys. 390, 1089. Peeters, E., Spoon, H. W. W., and Tielens, A. G. G. M.: 2004c, Astrophys. J., 613. Peeters, E., Tielens, A. G. G. M., Boogert, A. C. A., Hayward, T. L., and Allamandola, L. J.: 2005, Astrophys. J. 620, 774. Peeters, E., Mart´ın-Hernández,N. L., Damour, F., Cox, P., Roelfsema, P. R., et al.: 2002b, Astron. Astrophys. 381, 571. Peimbert, M. and Torres-Peimbert, S.: 1977, Mon. Not. R. Ast. Soc. 179, 217–234. Pilbratt, G. L., Altieri, B., Blommaert, J. A. D. L., Fridlund, C. V., Tauber, J. A., and Kessler, M. F.: 1998, Astron. Astrophys. 333, L9. Reach, W. T. and Tho, J.: 1996, Astron. Astrophys. 315, L277. Reach, W. T. and Tho, J.: 1998, Astrophys. J. 507, L93. Reach, W. T., Boulanger, F., Contursi, A., and Lequeux, J.: 2000a, Astron. Astrophys. 361, 895. Reach, W. T. and Tho, J.: 2000b, Astrophys. J. 544, 843. Reach, W. T., Tho, J., Jarrett, T. H., and Lagage, P.-O.: 2002, Astrophys. J. 564, 302. Rho, J., Jarrett, T. H., Cutri, R. M., and Reach, W. T.: 2001, Astrophys. J. 547, 885. Rodr´ıguez-Fernández,N. J., Mart´ın-Pintado,J., and de Vicente, P.: 2001a, Astron.y & Astrophysics 377, 631. Rodr´ıguez-Fernández, N. J., Mart´ın-Pintado, J., Fuente, A., de Vicente, P., Wilson, T. L., and Hüttemeister: 2001b, Astron. Astrophys. 377, 631. Rodr´ıguez-Fernández,N. J. and Mart´ın-Pintado:2005, Astron. Astrophys. 429, 923. Rodr´ıguez-Fernández,N. J., Mart´ın-Pintado,J., Fuente, A., and Wilson, T. L.: 2004, Astron. Astrophys. 427, 217. Roelfsema, P. R., Cox, P., Tielens, A. G. G. M., Allamandola, L. J., Baluteau, J.-P., et al.: 1996, Astron. Astrophys. 315, L289. Rubin, R. H., Simpson, J. P., Erickson, E. F., and Haas, M. R.: 1988, Astrophys. J. 327, 377. Rudolph, A. L., Simpson, J. P., Haas, M. R., Erickson, E. F., and Fich, M.: 1997, Astrophys. J. 489, 94. Russell, R. W., Melnick, G., Smyers, S. D., Kurtz, N. T., Gosnell, T. R., et al.: 1981, Astrophys. J. 250, L35. Russell, R. W., Melnick, G., Gull, G. E., and Harwit, M.: 1980. Astrophys. J. 240, L99. Schneider, N., Simon, R., Kramer, C., Kraemer, K., Stutzki, J., and Mookerjea, B.: 2003, Astron. Astrophys. 406, 915. Simpson, J. P. and Rubin, R. H.: 1990, Astrophys. J. 354, 165. 292 E. PEETERS ET AL.

Simpson, J. P., Colgan, S. W. J., Rubin, R. H., Erickson, E. F., and Haas, M. R.: 1995, Astrophys. J. 444, 721. Simpson, J. P., Witteborn, F. C., Price, S. D., and Cohen, M.: 1998, Astrophys. J. 508, 268. Sollerman, J.: 2002, New Astron. Rev. 46, 493. Song, I.-O., Kerr, T. H., McCombie, J. and Sarre, P. J.: 2004, Mon. Not. R. Astron. Soc. 346, L1. Stasi´nska,G.: 1990, Astron. Astrophys. Suppl. Ser. 83, 501. Storey, J. W. V., Watson, D. M., and Townes, C. H.: 1979, Astrophys. J. 233, 109. Tielens, A. G. M. M. and Hollenbach, D.: 1985, Astrophys. J. 291, 722. Timmermann, R., Bertoldi, F., Wright, C. M., Drapatz, S., Draine, B. T., et al.: 1996, Astron. Astrophys. 315, L281. Thornley, M. D., Schreiber, N. M. F., Lutz, D., Genzel, R., Spoon, H. W. W., et al.: 2000, Astrophysical Journal 539, 641–657. Torres-Peimbert, S. and Peimbert, M.: 1977, Rev. Mexicana de Astron. Astrof´ıs. 2, 181–207. Tuffs, R. J., Drury, L. O. ’C., Fischera, J., Heinrichsen, I., Rasmussens, I. et al.: 1997, First ISO Workshop on Analytical Spectroscopy, ESA-SP419, 177. Urquhart, J. S., White, G. J., Pilbratt, G. L., and Fridlund, C. V. M.: 2003, Astron. Astrophys. 409, 193. van den Ancker, M. E., Tielens, A. G. G. M., and Wesselius, P. R.: 2000, Astron. Astrophys. 358, 1035. van Diedenhoven, B., Peeters, E., Van Kerckhoven, C., Hony, S., Hudgins, D. M., et al.: 2004, Astrophysical Journal 611, 928. Van Kerckhoven, C., Hony, S., Peeters, E., Tielens, A. G. G. M., and Allamandola, L. J., et al.: 2000, Astron. Astrophys. 357, 1013. Van Kerckhoven, C.: 2002, PhD thesis, Catholic University of Leuven, Belgium. Vastel, C., Spaans, M., Ceccarelli, C., Tielens, A. G. G. M., and Caux, E.: 2001, Astron. Astrophys. 376, 1064. Verma, R. P., Ghosh, S. K., Mookerjea, B., and Rengarajan, T. N.: 2003, Astron. Astrophys. 398, 589. Vermeij, R., Damour, F., van der Hulst, J. M., and Baluteau, J.-P.: 2002a, Astron. Astrophys. 390, 649. Vermeij, R. and van der Hulst, J. M.: 2002b, Astron. Astrophys. 391, 1081. Vermeij, R., Peeters, E., Tielens, A. G. G. M., and van der Hulst, J. M.: 2002c, Astron. Astrophys. 382, 1042. Verstraete, L., Puget, F. L., Falgarone, E., Drapatz, S., Wright, C. M., and Timmermann, R.: 1996, Astron. Astrophys. 315, L337. Verstraete, L., Pech, C., Moutou, C., Sellgren, K., Wright, C. M., et al.: 2001, Astron. Astrophys. 372, 981. Wada, S., Onaka, T., Yamamura, I., Murata, Y., and Tokunaga, A. T.: 2003, Astron. Astrophys. 407, 551. Willner, S. P. and Nelson-Patel, K.: 2002, Astrophys. J. 568, 679. Wright, C. W., Drapatz, S., Timmermann, R., van der Werf, P. P., Katterloher, R., and de Graauw, Th.: 1996, Astron. Astrophys. 315, L301. Wright, C. W., van Dishoeck, E. F., Cox, P., Sidher, S. D., and Kessler, M. F.: 1999, Astrophys. J. 515, L29. Zavagno, A. and Ducci, V.: 2001, Astron. Astrophys. 371, 312. THE ICE SURVEY OPPORTUNITY OF ISO

EMMANUEL DARTOIS Institut d’Astrophysique Spatiale, Astrochimie Experimentale,´ UMR-8617 Universite´ Paris-Sud, batimentˆ 121, F-91405 Orsay, France (E-mail: [email protected])

(Received 30 July 2004; Accepted in ﬁnal form 27 October 2004)

Abstract. The instruments on board the Infrared Space Observatory have for the first time allowed a complete low (PHOT, CVF) to medium resolution (SWS) spectroscopic harvest, from 2.5 to 45 μm, of interstellar dust. Amongst the detected solids present in starless molecular clouds surrounding recently born stellar and still embedded objects or products of the chemistry in some mass loss envelopes, the so-called “ice mantles” are of specific interest. They represent an interface between the very refractory carbonaceous and silicates materials that built the first grains with the rich chemistry taking place in the gas phase. Molecules condense, react on ices, are subjected to UV and cosmic ray irradiation at low temperatures, participating efficiently to the evolution toward more complex molecules, being in constant interaction in an ice layer. They also play an important role in the radiative transfer of molecular clouds and strongly affect the gas phase chemistry. ISO results shed light on many other species than H2O ice. The detection of these van der Waal’s solids is mainly performed in absorption. Each ice feature observed by ISO spectrometer is an important species, with abundance −4 −7 in the 10 –10 range with respect to H2. Such high abundances represent a substantial reservoir of matter that, once released later on, replenishes the gas phase and feeds the ladder of molecular complexity. Medium resolution spectroscopy also offers the opportunity to look at individual line profiles of the ice features, and therefore to progressively reveal the interactions taking place in the mantles. This article will give a view on selected results to avoid to overlap with the numerous reviews the reader is invited to consult (e.g. van Dishoeck, in press; Gibb et al., 2004.).

Keywords: ices, dust, infrared absorption

1. Ices Everywhere

Interstellar ices have been observed with ISO-SWS in many different lines of sight among which the most common are: r OH-IR (and by extension, evolved stars circumstellar shells) which are oxygen rich post main sequence stars loosing mass at a high rate. The dust and molecules formed in this mass loss ejecta feed a circumstellar shell, opaque to visible light and therefore re-emitting in the infrared. In these lines of sight, H2O is likely to be the only ice present on dust grains, generally in crystalline

form (see Figure 1) as it was formed in the gas phase and then condensed at r high temperature. External galaxies (e.g. M82, Figure 2). In such lines of sight, the ocurrence of ices is strongly dependent on geometrical effects. Not surprisingly, galaxies do contain large amounts of ices inside molecular clouds.

Figure 1. Short Wavelength Spectrometer (AOT01) observations of two OH-IR circumstellar shells, displaying silicates features and water ice signatures. The insert contains the continuum divided spectra in the 3 μm region of the spectrum, in order to show the OH stretching mode of crystalline water ice. Extracted from the ISO database (http://isowww.estec.esa.nl/).

Figure 2. Short Wavelength Spectrometer (AOT01) observations of the Andromeda galaxy. The galactic continuum is absorbed by water ice present in the dense clouds of the galaxy. Extracted from the ISO database. ICE SURVEY OPPORTUNITY OF ISO 295

Figure 3. Short Wavelength Spectrometer (AOT06) observation of the Field star Elias 3–16, located behind the Taurus dark cloud. The infrared continuum of the star is absorbed selectively by the ice mantles located in the cloud. Extracted from the ISO database. r Field Stars are located behind molecular clouds (e.g. Elias 3–16, Figure 3) and thus allow one to use their infrared pencil to probe the foreground molecular cloud ice composition. They are crucial to understand ice evolution and chemistry during the earliest stage of starless clouds formation. The drawbacks are that they are statistically scarse and faint in the infrared. They are typically main sequence stars, often possessing intrinsic photospheric absorptions one r has to take into account to interpret the data. Embedded protostellar objects (e.g. Figure 4) which constitute by far the more abundant and richest database of ice species, with an observational bias toward the brightest high mass sources. The drawback of these sources lines of sight is that the central object has generally started the interaction with the ice containing parental cloud, stimulating new processes but also complicating the analysis. Such an extended panel of common line of sight, to which must be added the ones that to date escape detections such as optically thick disks, suggests ices are very common and abundant constituents in the life-cycle of dust in galaxies.

2. Inventory of the Detected Ice Features

ISO has brought the number of ice features to more than 20 for clear detections, to which must be added a few ones to conﬁrm. These features, the corresponding molecule, vibrational mode and integrated band strength are summarized in Table I. The large wavelength coverage of the ISO spectrometers allowed to observe in the same spectra several modes of the same ice constituent. 296 EMMANUEL DARTOIS

Figure 4. Short Wavelength Spectrometer (AOT06) observations of the rich ice spectrum of embedded intermediate to high mass protostars. From top to bottom: Elias 2–29, RCRA, GL2136, NGC7538 IRS9. Spectra have been shifted in ﬂux by arbitrary constants for better clarity. The more common identiﬁed features are labeled below the spectra. Extracted from the ISO database

3. Molecule by Molecule

3.1. H2O: THE LORD OF THE ICES

One of the ﬁrst but unsuccessful attempt to detect water ice was performed by Knacke et al. (1969) in the diffuse medium. Water ice was observed later on through ICE SURVEY OPPORTUNITY OF ISO 297

TABLE I Line list of ices absorptions observed by ISO.

Wavelength Vibrational A (μm) Molecule mode (10−18 cm/mol) Reference

2.70 CO2 Combination (ν1 + ν3) – 2.78 CO2 Combination (2ν2 + ν3) – 2.97 NH3 N H stretch (ν1) 10–11 (1,2) 3.05 H2OOH stretch (ν1 − ν3) 200 (1,3) 3.47 NH3 ···H2O Hydrate O H stretch ∼200 (4) 3.53 CH3OH CH3 s-stretch (ν3) 5.3–7.6 (1,5) 3.84–3.95 CH3OH Combination 2.6–3.2 (1,5) 4.27 CO2 Antisym. stretch (ν3) 76 (3) 13 4.38 CO2 Antisym. stretch (ν3) 78 (3) 4.55 H2O Libration overt. (3νL) ∼10 (3) 4.62 OCN− C N stretch (ν3) 40–80/(130 ?) (1,6,7,8)/(9) 4.67 CO CO stretch 11 (3) 4.78 13CO CO stretch 13 (3) 4.90 OCS CO stretch (ν1) 150–170 (5) 5.81 H2CO CO stretch (ν2) 9.6 (10) 5.83 HCOOH CO stretch (ν3) 67 (11) 5.85 CH3CHO CO stretch 13 (11) 6.02 H2OOH bend (ν2) 8.4–12 (1,3) 6.33 HCOO− CO stretch 100 (11) + ν 6.85 NH4 NH def ( 4) 40–44 (14) 7.25 HCOOH C H bend. (ν4) 2.6 (11) 7.25 HCOO− CO stretch 8 (11) 7.41 HCOO− 17 (11) 7.41 CH3CHO CO stretch 1.5 (11) 7.7 CH4 CH def. (ν4) 6.4–7.3 (1,12) 8.85 CH3OH CH3 rock (ν7) 1.3–1.8 (1,5) 9.35 NH3 Umbrella (ν2) 12–20 (2,13) 9.75 CH3OH CO stretch (ν8) 18 (1,5) 11–14 H2O Libration(νL) 28–31 (1,5) 15.2 CO2 CO2 bending (ν2) 11 (3) Note: The italicized detected molecules are either tentative, need confirmations or their abundance are still subject to strong debate. References: (1) Dhendecourt and Allamandola (1986); (2) Sandford and Allamandola (1993); (3) Gerakines et al. (1995); (4) Dartois and d’Hendecourt (2001); (5) Hudgins et al. (1993); (6) Grim and Greenberg (1987); (7) Schutte and Greenberg (1997); (8) Demyk et al. (1998); (9) van Broekhuizen et al. (2004) a value 3 σ above the other determinations; (10) Schutte et al. ((1996a)); (11) Schutte et al. (1999); (12) Boogert et al. (1997); (13) Kerkhof et al. (1999); (14) Schutte and Khanna (2003). 298 EMMANUEL DARTOIS its OH stretching mode at 3 μm, from the ground by Gillett and Forrest (1973). Since then, many line of sight were surveyed from the ground in this mode (e.g. Merrill et al., 1976; Léger et al., 1979; Willner et al., 1982; Whittet et al., 1983; Smith, Sellgren et al., 1989; Eiroa and Hodapp, 1989; Tanaka et al., 1990; Sato et al., 1990; Chen and Graham, 1993; Smith et al., 1993). After the discovery of the vibrational fundamental modes, phonon modes around 44 and 66 μm were found in emission by Omont et al. (1990) and followed by laboratory analysis (Moore and Hudson, 1992; Hudgins et al., 1993; Smith et al., 1994; Maldoni et al., 1999). So many others authors contributed to this search that a full list is out of scope. In this context, the insight in a better understanding of water ice performed by ISO has first been to substantially extend the number of astrophysical lines of sight where water ice is observed and, more interestingly, to offer a complete and simultaneous coverage of the six principal modes/combinations occurring in the infrared: the OH stretch (∼3.05 μm), the libration overtone (∼4.5 μm), the OH bend (∼6.0 μm), the libration (∼11–13 μm), the longitudinal and transverse optical mode at ∼44 μm and ∼66 μm, respectively. The two latter modes are detected either in absorption (Dartois et al., 1998) or emission (Malfait et al., 1999; Molinari et al., 1999; Hoogzaad et al., 2002; Molster et al., 2002; Maldoni et al., 2003) betraying the radiative transfer effects at long wavelength. For the water ice seen in emission toward some lines of sight, care must however be taken to analyse the spectra, as silicates far infrared features (especially enstatite) contribute to the lines. In addition to their detections, ISO has brought truly new spectroscopic constraints on the water ice modes. As an example, the OH bending mode profiles reveal the contribution of other features, unaccessible before (Keane et al., 2001), and whose identification will be discussed later on. Water ice is the dominant solid state frozen species and consequently the abundances of other ices are referred to this one on a relative scale (see Figure 6). On an absolute scale, in clouds with substantial visual extinction, water ice is the second −4 −5 most abundant species after H2 ([H2O]/[H2] ≈ 10 –10 ), comparable to or even more abundant than gas phase CO.

3.2. UNRAVELING THE UBIQUITOUS CO2

Before the launch of ISO, the carbon dioxide existence in ice mantles was founded on very few astrophysical data. Its inference was indirect and based on the possible influence of CO2 interacting with other molecules on their line profiles (Sandford et al., 1988; Kerr et al., 1991; Tielens et al., 1991). Indeed, besides an infrared oscillator strength amongst the strongest of ISM solid molecules, the atmospheric carbon dioxide prevents any ground based or airborne observations detection of its space parent. The very first direct detection was performed in 1989, on unresolved CO2 bending mode profiles, using IRAS Low Resolution Spectrometer spectra (d’Hendecourt and Jourdain de Muizon, 1989). The ubiquitous nature of carbon dioxide was revealed and firmly assessed with ISO spectrometers (Guertler et al., ICE SURVEY OPPORTUNITY OF ISO 299

1996; Gerakines et al., 1999; Alexander et al., 2003; de Graauw et al., 1996). CO2 then appears as the most abundant molecule after H2O in ice mantles, with a mean CO2/H2O ratio ranging generally from 10 to 50% in protostellar environments, with a mean around 15–20%. 13 The CO2 stretching mode was detected, and its line profile analysed by Boogert et al. (2000). This is a valuable information, because being about 60 times less abundant than the main molecule, this isotopomer is experiencing and therefore tracing the ice matrix environment field. It allows to a certain point to distinguish between different thermal processing having affected in the past the ice mantles. Some combination modes (symmetric+antisymmetric stretching and twice the bending+symmetric stretching modes) have also been detected in S140 IRS1 (Keane et al., 2001), being the first detection of solid state identified combinations in space ice mantles. The position and profiles of these combinations implies that the carbon dioxide is mixed in roughly equal proportions with methanol and water ice. The ability to observe combination modes, with low integrated absorption cross sections further demonstrate these ices are abundant components of molecular clouds. Another results of ISO medium resolution spectra is the first detection of interstellar molecular complexes implying the interaction of a carbon atom of the carbon dioxide with the lone electron pairs of the methanol oxygen (Ehrenfreund et al., 1999; Dartois et al., 1999; Klotz et al., 2004). This interaction deforms the CO2 molecule in such a way that the degenerescence of the bending mode is broken, giving rise to two individual absorption components around 15.2 μm (see Figure 5, right panel), whereas the stretching mode is only slightly affected. The combination of this complexation with the thermal evolution of ice mantles gives rise to characteristic substructures, possessing from one to three subpeaks (Figure 5, left panel, Dartois, et al., 1999) in the bending mode regions of carbon dioxide infrared

Figure 5. Left: SWS01 absorption spectrum of RAFGL7009 in the CO2 bending mode region displaying a triple substructure, compared with laboratory ice spectrum + gas phase model. Right: This structure partly arises due to the break of degeneracy of the in-plane and out-of-plane CO2 bending mode (see text for details). 300 EMMANUEL DARTOIS spectra. Such a spectroscopic constraint is of major importance for the understanding of the structure in which are organised here the carbon dioxide and methanol ices. These structures are not fortuitous, requires a speciﬁc stœchiometry, and the premice of better understanding of their formation history.

3.3. CO

Solid CO is a molecule observed almost routinely from ground based telescopes and recently its principal isotopomer 13CO was reported (Boogert et al., 2002). Ground- based observations allow to study the carbon monoxide line proﬁles (Pontoppidan et al., 2003) with unprecedented spectral resolution. The reader is referred to the huge literature on the subject, notably on the various proﬁles induced by the polar or apolar environment of this molecule. One of the results of ISO was to allow a comparison with the same infrared pencil beam of gas-to-solid state CO ratio (Dartois et al., 1998; van Dishoeck et al., 1996), showing that even if it is an abundant component of ice mantles (typically 3–20%), it resides principally in the gas phase in the general case.

3.4. OCS AND SULFUR-CONTAINING MOLECULES

Sulfur is an abundant element of interstellar medium and one expects therefore to observe it in ices. Geballe et al. (1985) detect toward W33 A two absorptions at 3.9 and 4.9 μm, respectively attributed by these authors to H2S and another sulfur- containing molecule produced in laboratory experiments. The 3.9 μm feature will be attributed later on to solid methanol combination modes, the only really containing sulfur molecule giving rise to the feature at 4.9 μm. Smith (1991) discuss the search for H2S in six objects and put an upper limit of about 1% on its presence. Palumbo et al. (1997) and Palumbo et al. (1995) observe the 4.9 μm feature, identify it with carbonyl sulﬁde and discuss its abundance toward various lines of sight. The line proﬁle is best reproduced when OCS is mixed with methanol. In fact there is an overlap of the second overtone 2ν8 of methanol with this line. With ISO, OCS has been evaluated in many lines of sights (see references in Gibb et al., 2004), but most of the time without evaluating the methanol overtone contribution, especially towards RAFGL 7009 and W33 A. The CO stretch of OCS is one of the more intense infrared bands and therefore OCS is the molecule detected with the lowest abundance with respect to H2O, of about 0.05–0.15%. This can be considered as the limit of detection of ices.

3.5. CH4

The ground based very difﬁcult detections of methane were started by Lacy et al. (1991). With ISO, the methane molecule was detected in several sources via its ICE SURVEY OPPORTUNITY OF ISO 301 deformation mode at 7.67 μm mode, with abundances up to 4% with respect to H2O (d’Hendecourt et al., 1996), but more generally at the 1–2% level (Boogert et al., 1996). The gas-to-solid methane ratio in such lines of sight is less than unity, favouring a formation of this hydride by hydrogenation of atomic carbon directly on the surface of grains. The actual abundance of methane does not reﬂect its initial abundance in the solid phase as it is one of the easiest molecule to dissociate upon UV secondary photons photolysis or ion irradiation, typical of dark clouds.

3.6. H2CO

Most of the molecules containing a carbonyl group strongly absorb in the 5.7– 6 μm region of the spectrum, the simplest of them being the formaldehyde (e.g. Schutte et al., 1993). The first attempt to detect it was made by Schutte et al., (1996a) in GL 2136 but the abundance controversed due to a mixing of formaldehyde modes with methanol ones, also present in large amount in the same line of sight. An estimate was also given for W33 A by Brooke et al. (1999). These “pre-ISO” estimates were based on C H stretching modes rather difficult to disentangle from the methanol CH stretches present in the spectra. Formaldehyde strongest modes fall at 5.8 and 6.68 μm, both inaccessible from the ground. ISO therefore allowed the search of these reliable modes for the first time, even if protostellar sources do possess strong absorptions in the same wavelength region due to the H2O bending mode and the so-called 6.85 μm feature. Formaldehyde abundance in ISO spectra was estimated in five high mass protostellar envelopes (W33 A, AFGL7009, GL2136, GL989, NGC7538 IRS9) to range between 1 and 3% (see Keane et al., 2001 (Table III); Dartois et al., 1999; Gibb et al., 2000; and references therein).

3.7. CH3OH

CH3OH is mainly detected from ground-based observations of its ν3 mode around 3.53 μm . It was early detected through this mode (Baas et al., 1988; Grim et al., 1991; Allamandola et al., 1992) in massive protostars lines of sight. The presence of the methanol–carbon dioxide complex, discussed above in the CO2 section, raised a new interest for its search (e.g. Dartois et al., 1999). Only recently and from the ground, it was shown by Pontoppidan et al. (2003) that methanol is also present around some low mass protostars. Methanol is a key molecule in ices, as it sometimes represent up to several tenth of the water ice abundance. Its formation route is under a controversial debate. Some authors think it is formed directly from grain surface hydrogenation of CO, but to date laboratory experiments performed give contradictory answers to the hypothesis (Hiraoka et al., 2002; Watanabe et al., 2003). 302 EMMANUEL DARTOIS

Others think that it can be created by energetic processes (UV or cosmic rays) (Hudson and Moore, 1999; Moore and Hudson, 1998), eventually followed by selective desorption enhancement effects, methanol being a rather refractive ice.

3.8. NITROGEN BEARING SPECIES

3.8.1. N2 Nitrogen being an abundant cosmic element and N2 a stable molecule, its presence would be logical. In practice, this homonuclear molecule is, due to its symmetry, inactive in the infrared. However, when embedded in an ice matrix, the crystal field breaks this symmetry and an infrared transition is activated around 4.295 μm. This weakly intense transition, with an integrated absorption cross section of 10−6– 10−4 the OH stretching one of water ice in the best activated mixtures (Sandford et al., 2001; Ehrenfreund and van Dishoeck, 1998 and references therein), is a challenge to detection limits. In addition, this feature falls in the red wing of the strong CO2 antistretch absorption, both in the solid and gas phase. To date, the only few upper limits derived for the presence of N2 when embedded into water ices, and using ISO SWS data (see above cited references), are not sufficiently low to provide strong constraints. Only very (too ?) specific ice mantles compositions (such as pure N2:CO2 ice mixtures) could lead to push down this limit by enhancing the activation of the otherwise inactive N2 stretching mode, which remains to be proved in an ice dominated by water ice.

3.8.2. NH3 After N2, the most abundant nitrogen bearing molecule is certainly the ammonia one. Knacke et al. (1982) were among the firsts to claim for an identification of NH3 in interstellar grains from an infrared absorption at 2.97 μm (NH stretching mode), a detection refuted later on Knacke and McCorkle (1987). Hagen et al. (1983) show that the H2O stretching mode profile is best reproduced with an ice containing a significant fraction of ammonia. Since then, many authors entered into the still opened debate of ammonia ice abundance in grain mantles. Ammonia possess three main modes at 2.97 μm (NH stretch), 6.2 μm (NH bend) and 9. μm(umbrella mode, which shifts to longer wavelength if pure ammonia). Importantly, and often underestimated, when mixed with H2O, the major ice mantle constituent, ammonia forms an hydrate that gives rise to a large and intense band around 3.47 μm (with an integrated absorption cross section similar to the water OH stretch). All modes, except for the hydrate one, fall in spectroscopic regions overlapping with the dominant water ice or silicate features, which renders the detection task difficult. In addition, toward lines of sight where methanol is present, the CH3 rock can contribute significantly to the ammonia umbrella mode. Detecting the relatively intense 9 μm umbrella mode in the wing of the strongly absorbed silicates features has always led to high overestimates of the true ammonia contents. ICE SURVEY OPPORTUNITY OF ISO 303

A school case to support this view comes from the well-studied W33 A line of sight. The methanol contribution at 8.9–9 μm, inferred from other modes, should provide an optical depth around ∼0.1. To show consistency in interpreting spectra of this source, and before any ammonia search, this CH3OH mode must be evidenced and subtracted to the data. Based on ISO PHOT40 and SWS spectra extracts around 9 μm, various authors derive quite different column densities of NH3, namely 1.7 × 1018 cm−2 (Gibb et al., 2001), 1.2–2.2 × 1018 cm−2 (G¨urtler et al., 2002), representing a high percentage (11–20%) as compared to water ice. These numbers were in fact all contradicted by the recent ground based upper limit toward this source by Taban et al. (2003), which lies about two to three times lower. In the analysis presented in Dartois and d’Hendecourt (2001) and Gibb et al. (2001), on 20 different sources lines of sight (mainly massive YSOs), only in three of the Gibb’s article are claimed NH3 detections. For these three sources : (i) W33 A detection was later on rejected by Taban et al. (2003); (ii) NGC7538 IRS9 possess a considerable uncertainty on ammonia abundance given the scatter in the derived NH3/H2O ratios of 10% (Lacy et al., 1998), 15% (Gibb et al., 2001), <5% with PHOT and ∼8% with SWS (G¨urtler et al., 2002); (iii) GL989 possess at most 5% of ammonia (Dartois and d’Hendecourt, 2001; Gibb et al., 2001). Based on the absence of the 2.97 μm NH stretching mode and the maximum contribution of the ammonia hydrate at 3.47 μm,Dartois and d’Hendecourt (2001) deduced that a conservative upper limit of 5% ammonia relative to water ice can be hidden in the ice mantles. ISO has allowed to cover the full range of main ammonia modes and this is a great advantage when searching for a molecule with several infrared active modes. A detection should therefore be compatible with the complete spectrum as discussed thoroughly in Dartois et al. (2002). The uncertainties and scatter in the determination of ammonia column densities using only its umbrella mode evidence this mode is not reliable when extracted on the deep silicate stretching mode absorption.

3.8.3. OCN− A special nitrogen molecule in interstellar ices is represented by the cyanate anion which possess its strongest infrared active mode around 4.62 μm, aside to the CO feature. Falling in an atmospheric window, it has been seen early by (McGregor et al., 1982; Lacy et al., 1984) and assigned then to a nitrile bond ( C≡N) although it was present but unresolved (i.e. mixed with solid CO) in the observations of W33 A by Soifer et al. (1979). It was then ﬁrst assigned by Grim and Greenberg (1987) to the − 13 15 OCN cyanate anion, by photolysis of CO/N3 ice mixtures and using C and N isotopic labelling. Such species was conﬁrmed by further experiments using 18O substitution (Schutte and Greenberg, 1997). Many others authors have pursued the “XC≡N” in the laboratory through isotopes substitutions, proton and UV irradia- tions of ices, etc., among which one can cite the non exhaustive list of Novozamsky et al. (2001), Hudson et al. (2001), Hudson and Moore (2000), Bernstein et al. (2000), Palumbo et al. (2000), Pendleton et al. (1999), Demyk et al. (1998). 304 EMMANUEL DARTOIS

The observation of OCN− was long though to be very important because it was the molecule giving evidence of the ultraviolet field present inside dense clouds. However, recently, the debate about its formation mechanism has been relaunched as several laboratory investigations gave an alternative route via a direct proton − + transfer from HOCN (or HNCO) to NH3, leading to the formation of OCN NH4 (Demyk et al., 1998; Raunier et al., 2004; van Broekhuizen et al., 2004; Raunier et al., 2004). Such a proton transfer could be so efficient, even at low temperature, that it would be impossible to see the HNCO precursor. The abundance of OCN− relies on the adopted integrated cross section, which varies from about A = 4 ×10−17 to 1.3 × 10−16 using different approaches to estimate it (Dhendecourt et al., 1986; Grim and Greenberg, 1987; Demyk et al., 1998; van Broekhuizen et al., 2004), but favouring the lower value for A. From the ISO observational point of view, the exact abundance or upper limits reported on various ISO lines of sights (e.g. Gibb et al. (2004)), completed by ground based recent VLT ones (von Broekhuizen, private communication) of OCN− represent at most a few percent of the water ice content, up to 4% for the highest column density found, for W33A. + 3.8.4. NH4 To preserve the global neutrality of the ices, and as a consequence of the processes − + proposed for the existence of the OCN anion, the counterion NH4 has been proposed years ago (Grim et al., 1989). Such an assignment was compatible with the early detection of a band at around 6.85 μm (Russell et al., 1977), therefore associated with the ν4 mode of this cation. The first medium resolution spectra of this band were obtained by ISO (e.g. Schutte et al., 1996b). A rather complete analysis of the various hypothesis to explain this 6.85 μm absorption was performed by Keane et al. (2001), on lines of sight with different environmental conditions. They conclude that this band is composed of two components, related through thermal processing. In the laboratory, after many years of ammonia containing ice mixtures or ther- + mal processing of NH3 with acids (formic acid, cyanic acid, isocyanic acid), NH4 sould explain a part if not all the 6.85 μm feature (Schutte and Khanna, 2003 and references therein for a view of laboratory work), and other counterions than OCN− would be present in the ices, forming an infrared absorption continuum that is hardly detectable.

3.9. O2

Like N2, the homonuclear oxygen molecule O2 is infrared inactive, except when embedded in an ice matrix, where it is slightly activated (Ehrenfreund, 1992; Ehrenfreund and van Dishoeck, 1998) and gives rise to an absorption around 6.45 μm. Considerable laboratory work was dedicated to infer the presence of O2 by studying its inﬂuence on other modes, among which Ehrenfreund et al. ICE SURVEY OPPORTUNITY OF ISO 305

(1997), Strazzulla et al. (1997), Elsila et al. (1997), Ehrenfreund et al. (1996). With ISO, the upper limits found by Vandenbussche et al. (1999) are not sufﬁciently low to give any constraint on the ice mantles oxygen content, but still tell us that no more than about 6% of the oxygen cosmic budget can be hidden in this particular component. Molecular oxygen seems not extremely abundant in ice mantles. Ad- ditional indirect evidence for that might come from the absence of detection of its photoproducts like CO3 or O3. A high abundance of molecular oxygen would also inﬂuence drastically the ice mantles spectra observed by creating a lot of new radicals that can be seen in all the relevant irradiated mixtures found in the references of this review. The fact that the SWAS satellite also demonstrated O2 is not abundant in the gas phase tends to demonstrate the oxygen atom assume other forms. In particular, in dense clouds, carbon monoxide and dioxide (gas and solid), water ice and silicates already consume a large part of the total oxygen budget.

3.10. RELATIVE ABUNDANCES

A schematic diagram showing the typical relative abundances of species with respect to the main H2O ice, based on ISO references of this review, is presented in Figure 6. Grey histogram implies the values can be lower than typical, especially for sources lines of sight where a warming of the mantle has begun (e.g. S140). The ammonia case is speciﬁc as no truly reliable column densities are in the general case available (see corresponding section), but its occurrence is however certain. On the right side, an approximate absolute abundance is given, assuming that about

Figure 6. Schematic diagram of solid species typical abundances, normalized to water ice. The grey line are here to represent the selective disappearance of species from the solid phase by e.g. selective thermal desorption. Approximate absolute abundances are given on the right scale. See text for explanation. 306 EMMANUEL DARTOIS

10–15% of total cosmic oxygen is locked into water ice. Such a reference is based on typical correlation found in dark clouds between visual extinction and H2O ice (e.g. Murakawa et al., 2000 and references in the H2O section).

4. Tentative Detections, Requested Conﬁrmations

Other species, with low fractional abundances have been proposed/suggested to explain some minor features in ISO ice mantles spectra. The presence of most of them requires however confirmations to assess their role in ices. In this list one can − find HCOO ,CH3CHO, HCONH2 (Schutte et al., 1999) and (NH2)2CO (Raunier et al., 2004). A feature still under debate is the 6.85 μm one. It falls in wavelength range which is an infrared spectroscopic accumulation point. In fact, many molecules absorb in this range and if it is due to the superposition of many distinct species, it will be difficult to sort out which ones. This feature can then also be considered as needing further confirmations for the contributing species.

5. ISO

ISO has opened a new era of ice mantles compositional determination, by providing for the first time a complete coverage of the vibrational spectrum of interstellar ices. The analysis is now based on a laboratory simultaneous match of all the modes of a species to be able to claim an identification, which from the ground is simply impossible. ISO showed that ices are a truly ubiquitous dense medium component, serving as interface between gas and solid state. It also provided insight into the initial conditions from which such dense media evolves to protoplanetary systems, and therefore the ice compositions to be explored by laboratory photo-thermo and/or cosmic rays processing of grain mantles. Typical abundances of ices summarized in Figure 6 tells us that the detection limit for ices by direct absorption spectroscopy against infrared sources continuum −7 −8 is about 10 –10 with respect to H2. This fact contrasts with radio astronomy detection of gas phase molecules that can push down this value more than five orders of magnitude below. One must keep in mind this observational bias and remember that it does show that every detected ice with ISO is a major element of the astrochemical evolution of clouds. In addition to a complete view on relative abundances of ice components, ISO showed that medium resolution spectra are the prerequisite to understand physically the behavior and interactions inside interstellar ice mantles. Without such a resolution, astrophysicists would have been unable to detect the CO2 combinations and overtones, to decipher the formation of CH3OH−CO2 intermolecular complexes, ICE SURVEY OPPORTUNITY OF ISO 307 the NH3 H2O hydrate, to understand the gas-to-solid state ratios probed in the same lines of sight, to cite few of them.

References

Allamandola, L. J., Sandford, S. A., Tielens, A. G. G. M., and Herbst, T. M.: 1992, Astrophys. J. 399, 134. Alexander, R. D., Casali, M. M., Andr´e,P., Persi, P., and Eiroa, C.: 2003, Astron. Astrophys. 401, 613. Baas, F., Grim, R. J. A., Schutte, W. A., Greenberg, J. M., and Geballe, T. R.: 1988, Dust in the Universe, Cambridge University Press, Cambridge, p. 55. Bernstein, M. P., Sandford, S. A., and Allamandola, L. J.: 2000, Astrophys. J. 542, 894. Boogert, A. C. A., Blake, G. A., and Tielens, A. G. G. M.: 2002, Astrophys. J. 577, 271. Boogert, A. C. A., Tielens, A. G. G. M., Ceccarelli, C., Boonman, A. M. S., van Dishoeck, E. F., Keane, J. V., Whittet, D. C. B., and de Graauw, T.: 2000, Astron. Astrophys. 360, 683. Boogert, A. C. A., et al.: 2000, Astron. Astrophys. 353, 349. Boogert, A. C. A., Schutte, W. A., Helmich, F. P., Tielens, A. G. G. M., and Wooden, D. H.: 1997, Astron. Astrophys. 317, 929. Boogert, A. C. A., et al.: 1996, Astron. Astrophys. 315, L377. Brooke, T. Y., Sellgren, K., and Geballe, T. R.: 1999, Astrophys. J. 517, 883. van Broekhuizen, F. A., Keane, J. V., and Schutte, W. A.: 2004, Astron. Astrophys. 415, 425. Chen, W. P. and Graham, J. A.: 1993, Astrophys. J. 409, 319. Dartois, E., d’Hendecourt, L., Thi, W., Pontoppidan, K. M. and van Dishoeck, E. F.: 2002, Astron. Astrophys. 394, 1057. Dartois, E. and d’Hendecourt, L.: 2001, Astron. Astrophys. 365, 144. Dartois, E., Demyk, K., d’Hendecourt, L., and Ehrenfreund, P.: 1999, Astron. Astrophys. 351, 1066. Dartois, E., Schutte, W., Geballe, T. R., Demyk, K., Ehrenfreund, P., and d’Hendecourt, L.: 1999, Astron. Astrophys. 342, L32. Dartois, E., et al.: 1998, Astron. Astrophys. 338, L21. Dartois, E., Demyk, K., Gerin, M., and d’Hendecourt, L.: 1999, Solid Interstellar Matter: The ISO Revolution, Les Houches Workshop, February 2–6, 1998, L. d’Hendecourt, C. Joblin, and A. Jones (eds.), EDP Sciences Les Ulis; Springer-Verlag, Berlin, Heidelberg, p. 162. Dartois, E., d’Hendecourt, L., Boulanger, F., Jourdain de Muizon, M., Breitfellner, M., Puget, J.-L., and Habing, H. J.: 1998, Astron. Astrophys. 331, 651. Demyk, K., Dartois, E., d’Hendecourt, L., Jourdain de Muizon, M., Heras, A. M., and Breitfellner, M.: 1998, Astron. Astrophys. 339, 553. d’Hendecourt, L., et al.: 1996, Astron. Astrophys. 315, L365. d’Hendecourt, L. B. and Jourdain de Muizon, M.: 1989, Astron. Astrophys. 223, L5. Dhendecourt, L. B., Allamandola, L. J., Grim, R. J. A., and Greenberg, J. M.: 1986, Astron. Astrophys. 158, 119. Dhendecourt, L. B. and Allamandola, L. J.: 1986, Astron. Astrophys. Suppl. 64, 453. Ehrenfreund, P., et al.: 1999, Astron. Astrophys. 350, 240. Ehrenfreund, P. and van Dishoeck, E. F.: 1998, Adv. Space Res. 21, 15. Ehrenfreund, P., Boogert, A. C. A., Gerakines, P. A., Tielens, A. G. G. M., and van Dishoeck, E. F.: 1997, Astron. Astrophys. 328, 649. Ehrenfreund, P.,Boogert, A. C. A., Gerakines, P.A., Jansen, D. J., Schutte, W. A., Tielens, A. G. G. M., and van Dishoeck, E. F.: 1996, Astron. Astrophys. 315, L341. Eiroa, C. and Hodapp, K.-W.: 1989, Astron. Astrophys. 210, 345. Elsila, J., Allamandola, L. J., and Sandford, S. A.: 1997, Astrophys. J. 479, 818. 308 EMMANUEL DARTOIS

Geballe, T. R., Baas, F., Greenberg, J. M., and Schutte, W.: 1985, Astron. Astrophys. 146, L6. Gerakines, P. A., et al.: 1999, Astrophys. J. 522, 357. Gerakines, P. A., Schutte, W. A., Greenberg, J. M., and van Dishoeck, E. F.: 1995, Astron. Astrophys. 296, 810. Gibb, E. L., Whittet, D. C. B., Boogert, A. C. A., and Tielens, A. G. G. M.: 2004, Astrophys. J. Suppl. 151, 35. Gibb, E. L., Whittet, D. C. B., and Chiar, J. E.: 2001, Astrophys. J. 558, 702. Gibb, E. L., et al.: 2000, Astrophys. J. 536, 347. Gillett, F. C. and Forrest, W. J.: 1973, Astrophys. J. 179, 483. Grim, R. J. A., Baas, F., Greenberg, J. M., Geballe, T. R., and Schutte, W.: 1991, Astron. Astrophys. 243, 473. Grim, R. J. A., Greenberg, J. M., Schutte, W. A., and Schmitt, B.: 1989, Astrophys. J. 341, L87. Grim, R. J. A., Greenberg, J. M., de Groot, M. S., Baas, F., Schutte, W. A., and Schmitt, B.: 1989, Astron. Astrophys. Suppl. 78, 161. Grim, R. J. A. and Greenberg, J. M.: 1987, Astrophys. J. 321, L91. G¨urtler, J., Klaas, U., Henning, T., Abrah´am,P.,´ Lemke, D., Schreyer, K., and Lehmann, K.: 2002, Astron. Astrophys. 390, 1075. Guertler, J., Henning, T., Koempe, C., Pfau, W., Kraetschmer, W., and Lemke, D.: 1996, Astron. Astrophys. 315, L189. Hagen, W., Greenberg, J. M., and Tielens, A. G. G. M.: 1983, Astron. Astrophys. 117, 132. Hiraoka, K., Sato, T., Sato, S., Sogoshi, N., Yokoyama, T., Takashima, H., and Kitagawa, S.: 2002, Astrophys. J. 577, 265. Hoogzaad, S. N., Molster, F. J., Dominik, C., Waters, L. B. F. M., Barlow, M. J., and de Koter, A.: 2002, Astron. Astrophys. 389, 547. Hudgins, D. M., Sandford, S. A., Allamandola, L. J. and Tielens, A. G. G. M.: 1993, Astrophys. J. Suppl. 86, 713. Hudson, R. L., Moore, M. H., and Gerakines, P. A.: 2001, Astrophys. J. 550, 1140. Hudson, R. L. and Moore, M. H.: 2000, Astron. Astrophys. 357, 787. Hudson, R. L. and Moore, M. H.: 1999, Icarus 140, 451. Keane, J. V., Tielens, A. G. G. M., Boogert, A. C. A., Schutte, W. A., and Whittet, D. C. B.: 2001, Astron. Astrophys. 376, 254. Keane, J. V., Boogert, A. C. A., Tielens, A. G. G. M., Ehrenfreund, P., and Schutte, W. A.: 2001, Astron. Astrophys. 375, L43. Kerkhof, O., Schutte, W. A., and Ehrenfreund, P.: 1999, Astron. Astrophys. 346, 990. Kerr, T. H., Adamson, A. J., and Whittet, D. C. B.: 1991, Mon. Not. Roy. Astron. Soc. 251, 60. Klotz, A., Ward, T., and Dartois, E.: 2004, Astron. Astrophys. 416, 801. Knacke, R. F. and McCorkle, S. M.: 1987, Astron. J. 94, 972. Knacke, R. F., McCorkle, S., Puetter, R. C., Erickson, E. F., and Kraetschmer, W.: 1982, Astrophys. J. 260, 141. Knacke, R. F., Cudaback, D. D., and Gaustad, J. E.: 1969, Astrophys. J. 158, 151. Lacy, J. H., Faraji, H., Sandford, S. A., and Allamandola, L. J.: 1998, Astrophys. J. Lett. 501, L105. Lacy, J. H., Carr, J. S., Evans, N. J., Baas, F., Achtermann, J. M., and Arens, J. F.: 1991, Astrophys. J. 376, 556. Lacy, J. H., Baas, F., Allamandola, L. J., van de Bult, C. E. P.,Persson, S. E., McGregor, P.J., Lonsdale, C. J., and Geballe, T. R.: 1984, Astrophys. J. 276, 533. Leger, A., Klein, J., de Cheveigne, S., Guinet, C., Defourneau, D., and Belin, M.: 1979, Astron. Astrophys. 79, 256. Maldoni, M. M., Egan, M. P., Smith, R. G., Robinson, G., and Wright, C. M.: 2003, Mon. Not. Roy. Astron. Soc. 345, 912. ICE SURVEY OPPORTUNITY OF ISO 309

Maldoni, M. M., Robinson, G., Smith, R. G., Duley, W. W., and Scott, A.: 1999, Mon. Not. Roy. Astron. Soc. 309, 325. Malfait, K., Waelkens, C., Bouwman, J., de Koter, A., and Waters, L. B. F. M.: 1999, Astron. Astrophys. 345, 181. McGregor, P. J., Persson, S. E., Lacy, J. H., Baas, F., Allamandola, L., Lonsdale, C. J., and Greballe, T. R.: 1982, ESA SP-192: Galactic and Extragalactic Infrared Spectroscopy, ESTEC Reproduction Services, Noordwijk, Holland, p. 51. Merrill, K. M., Russell, R. W., and Soifer, B. T.: 1976, Astrophys. J. 207, 763. Molinari, S., Ceccarelli, C., White, G. J., Saraceno, P., Nisini, B., Giannini, T., and Caux, E.: 1999, Astrophys. J. 521, L71. Molster, F. J., Waters, L. B. F. M., and Tielens, A. G. G. M.: 2002, Astron. Astrophys. 382, 222. Moore, M. H. and Hudson, R. L.: 1998, Icarus 135, 518. Moore, M. H. and Hudson, R. L.: 1992, Astrophys. J. 401, 353. Murakawa, K., Tamura, M., and Nagata, T.: 2000, Astrophys. J. Suppl. 128, 603. Novozamsky, J. H., Schutte, W. A., and Keane, J. V.: 2001, Astron. Astrophys. 379, 588. Omont, A., Forveille, T., Moseley, S. H., Glaccum, W. J., Harvey, P. M., Likkel, L., Loewenstein, R. F., and Lisse, C. M.: 1990, Astrophys. J. 355, L27. Palumbo, M. E., Pendleton, Y. J., and Strazzulla, G.: 2000, Astrophys. J. 542, 890. Palumbo, M. E., Geballe, T. R., and Tielens, A. G. G. M.: 1997, Astrophys. J. 479, 839. Palumbo, M. E., Tielens, A. G. G. M., and Tokunaga, A. T.: 1995, Astrophys. J. 449, 674. Pendleton, Y. J., Tielens, A. G. G. M., Tokunaga, A. T., and Bernstein, M. P.: 1999, Astrophys. J. 513, 294. Pontoppidan, K. M., et al. : 2003, Astron. Astrophys. 408, 981. Pontoppidan, K. M., Dartois, E., van Dishoeck, E. F., Thi, W.-F., and d’Hendecourt, L.: 2003, Astron. Astrophys. 404, L17. Raunier, S., Chiavassa, T., Marinelli, F., Aycard, J. P.: 2004, Chem. Phys. 302(1–3), 259. Raunier, S., Chiavassa, T., Duvernay, F., Borget, F., Aycard, J. P., Dartois, E., and d’Hendecourt, L.: 2004, Astron. Astrophys. 416, 165. Russell, R. W., Soifer, B. T., and Puetter, R. C.: 1977, Astron. Astrophys. 54, 959. Sandford, S. A., Bernstein, M. P., Allamandola, L. J., Goorvitch, D., and Teixeira, T. C. V. S.: 2001, Astrophys. J. 548, 836. Sandford, S. A. and Allamandola, L. J.: 1993, Astrophys. J. 417, 815. Sandford, S. A., Allamandola, L. J., Tielens, A. G. G. M., and Valero, G. J.: 1988, Astrophys. J. 329, 498. Sato, S., Nagata, T., Tanaka, M., and Yamamoto, T.: 1990, Astrophys. J. 359, 192. Schutte, W. A. and Khanna, R. K.: 2003, Astron. Astrophys. 398, 1049. Schutte, W. A., et al.: 1999, Astron. Astrophys. 343, 966. Schutte, W. A. and Greenberg, J. M.: 1997, Astron. Astrophys. 317, L43. Schutte, W. A., et al.: 1996, Astron. Astrophys. 315, L333. Schutte, W. A., Gerakines, P. A., Geballe, T. R., van Dishoeck, E. F., and Greenberg, J. M.: 1996, Astron. Astrophys. 309, 633. Schutte, W. A., Allamandola, L. J., and Sandford, S. A.: 1993, Icarus 104, 118. Smith, R. G., Robinson, G., Hyland, A. R., and Carpenter, G. L.: 1994, Mon. Not. Roy. Astron. Soc. 271, 481. Smith, R. G., Sellgren, K., and Brooke, T. Y.: 1993, Mon. Not. Roy. Astron. Soc. 263, 749. Smith, R. G.: 1991, Mon. Not. Roy. Astron. Soc. 249, 172. Smith, R. G., Sellgren, K., and Tokunaga, A. T.: 1989, Astrophys. J. 344, 413. Soifer, B. T., Puetter, R. C., Russell, R. W., Willner, S. P., Harvey, P. M., and Gillett, F. C.: 1979, Astrophys. J. 232, L53. Strazzulla, G., Brucato, J. R., Palumbo, M. E., and Satorre, M. A.: 1997, Astron. Astrophys. 321, 618. 310 EMMANUEL DARTOIS

Taban, I. M., Schutte, W. A., Pontoppidan, K. M., and van Dishoeck, E. F.: 2003, Astron. Astrophys. 399, 169. Tanaka, M., Sato, S., Nagata, T., and Yamamoto, T.: 1990, Astrophys. J. 352, 724. Tielens, A. G. G. M., Tokunaga, A. T., Geballe, T. R., and Baas, F.: 1991, Astrophys. J. 381, 181. van Dishoeck, E. F.: 2004, Annu. Rev. Astron. Astrophys. 42, 119. van Dishoeck, E. F., et al.: 1996, Astron. Astrophys. 315, L349. Vandenbussche, B., et al.: 1999, Astron. Astrophys. 346, L57. Watanabe, N., Shiraki, T., and Kouchi, A.: 2003, Astrophys. J. 588, L121. Willner, S. P., et al.: 1982, Astrophys. J. 253, 174. Whittet, D. C. B., Bode, M. F., Baines, D. W. T., Longmore, A. J., and Evans, A.: 1983, Nature 303, 218. OUR LOCAL UNIVERSE . . . NORMAL NEARBY GALAXIES

MARC SAUVAGE1,∗, RICHARD J. TUFFS2 and CRISTINA C. POPESCU2 1CEA/DSM/DAPNIA/Service d’Astrophysique, C.E. Saclay, 91191 Gif sur Yvette CEDEX, France 2Max Planck Institut fur¨ Kernphysik, Astrophysics Department, Saupfercheckweg 1, 69117 Heidelberg, Germany (∗Author for correspondence: E-mail: [email protected])

(Received 30 August 2004; Accepted in ﬁnal form 29 September 2004)

Abstract. Following on from IRAS, ISO has provided a huge advancement in our knowledge of the phenomenology of the infrared (IR) emission of normal galaxies and the underlying physical processes. Highlights include the discovery of an extended cold dust emission component, present in all types of gas-rich galaxies and carrying the bulk of the dust luminosity; the deﬁnitive characterisation of the spectral energy distribution in the IR, revealing the channels through which stars power the IR light; the derivation of realistic geometries for stars and dust from ISO imaging; the discovery of cold dust associated with H I extending beyond the optical body of galaxies; the remarkable similarity of the near-IR (NIR)/mid-IR (MIR) SEDs for spiral galaxies, revealing the importance of the photodissociation regions in the energy budget for that wavelength range; the importance of the emission from the central regions in shaping up the intensity and the colour of the global MIR luminosity; the discovery of the “hot” NIR continuum emission component of interstellar dust; the predominance of the diffuse cold neutral medium as the origin for the main interstellar cooling line, [C II] 158 μm, in normal galaxies.

Keywords: galaxies: spiral, galaxies: dwarf, galaxies: ellipticals, galaxies: ISM

1. Introduction

Although IRAS provided the first systematic survey of infrared (IR) emission from normal galaxies, it has been the photometric, imaging and spectroscopic capabilities of ISO which have unravelled the basic physical processes giving rise to this IR emission. Thanks to the broad spectral grasp of ISO (Kessler et al., 1996), the bulk of the emission from dust could be measured, providing the first quantitative assessment of the fraction of stellar light re-radiated by dust. The battery of filters has led to a definitive characterisation of the spectral energy distribution (SED) in the IR, revealing the contribution of the different stellar populations in powering the IR emission. The imaging capabilities have unveiled the complex morphology of galaxies in the IR, and their changing appearance with IR wavelength. They also allowed the exploration of hitherto undetected faint diffuse regions of galaxies. The contribution of different grain populations to the emission has been measured Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, The Netherlands, and the United Kingdom), and with the participation of ISAS and NASA. Space Science Reviews (2005) 119: 313–353 DOI: 10.1007/s11214-005-8071-0 C Springer 2005 314 M. SAUVAGE ET AL. through their characteristic spectroscopic signatures. Knowledge of the emission from cooling lines of the interstellar medium (ISM) has been extended to low luminosity quiescent spiral and dwarf galaxies. In this review we will concentrate on the mid-IR (MIR) to far-IR (FIR) properties of normal nearby galaxies. By normal we essentially mean that their SEDs are not powered by accretion. We will begin with spiral galaxies, since these have attracted most of the ISO observers’ attention. From these objects we will move to the other class of gas-rich galaxies, the dwarfs. However, as the nearby extragalactic population is not only made of spirals and dwarfs, we will conclude by the exploring the very varied IR properties of early-type galaxies.

2. Spiral Galaxies

2.1. SPATIAL DISTRIBUTIONS

2.1.1. FIR Morphologies ISOPHOT (Lemke et al., 1996), imaged three nearby galaxies (M 31: Haas et al., 1998; M 33: Hippelein et al., 2003 and M 101: Tuffs and Gabriel, 2003) in the 60– 200 μm range, with sufﬁcient linear resolution to easily distinguish between the main morphological components in the FIR – nucleus, spiral arms and underlying disk. The main discovery, made possible by the unprecedented surface brightness sensitivity longwards of 100 μm, was the existence of large amounts of cold dust associated both with the spiral arms and with the underlying disk. This dust was too cold to have been seen by IRAS. Furthermore, ground-based submillimeter (submm) facilities lacked the surface brightness sensitivity to map the diffuse component of the cold dust associated with the underlying disk, though they detected the component associated with the spiral arms (e.g. Bianchi et al., 2000b). In the case of the Sab galaxy M 31, most of the emission at 170 μm arises from the underlying disk, which has a completely diffuse appearance (Figure 1, left panel). This diffuse disk emission can be traced out to a radius of 22 kpc, so the galaxy has a similar overall size in the FIR as seen in the optical bands. Figure 1 (left panel) also shows that at 170 μm the spiral arm component is dominated by a ring of 10 kpc radius. In addition, there is a faint nuclear source, which is seen more prominently in HIRES IRAS 60 μm maps at similar resolution and in Hα. The overall SED (Figure 1, right panel) can be well described as a superposition of two modiﬁed (β = 2) Planck curves, with dust temperatures TD of 16 and 45 K. The cold dust component at 16 K arises from both the ring structure (30%) and the diffuse disk (70%; Haas, private communication), illustrating the importance of the diffuse emission at least for this example. The 45 K component matches up well with H II regions within the star-formation complexes in the ring structure. Associated with each star-formation complex are also compact, cold emission sources (see Figure 3 of Schmitobreick et al., 2000) with dust temperatures in the 15–20 K range. These NORMAL NEARBY GALAXIES 315

Figure 1. Left: ISOPHOT 170 μm map of M 31 (Haas et al., 1998), with an angular resolution of 1.3. North is towards the top, and East is towards the left. The field size is 2.9×2.9 degrees. Right: Infrared SED of M 31 (Haas et al., 1998). The data are shown by symbols (diamonds ISO, crosses IRAS, triangles DIRBE) with the size being larger than the errors. The blackbody curves with emissivity proportional to λ−2 are shown by lines. The dotted line with T = 22 K through the IRAS 60 and 100 μm data points indicate what one would extrapolate from this wavelength range alone without any further assumptions. could well represent the parent molecular clouds in the star-formation complexes which gave rise to the H II regions. Detailed examination of the morphology of the ring shows a smooth component of cold dust emission as well as the discrete cold dust sources. Finally, the nuclear emission was fitted by a 28-K dust component. The ISOPHOT maps of the Sc galaxies M 33 and M 101 show the same morphological components as seen in M 31, with the difference that the spiral arm structure can be better defined in these later-type spirals. Also the star-formation complexes in the spiral arms show similar SEDs to those seen in M 31. In conclusion, the characteristics of the FIR emission from the main morphological components of spiral galaxies are:

– nucleus: an unresolved warm source with TD ∼30 K – spiral arms: a superposition of: r localised warm emission with 40 ≤ TD ≤ 60 K from H II regions r ≤ ≤ r localised cold emission with 15 TD 20 K from parent molecular clouds diffuse emission running along the arms

– disk: an underlying diffuse (predominantly cold) emission with 12 ≤ TD ≤ 20 K

2.1.2. MIR Morphologies Generally, ISOCAM (Cesarsky et al., 1996), maps of nearby galaxies are difficult to compare with their FIR counterparts from ISOPHOT, for two reasons. On the one 316 M. SAUVAGE ET AL. hand ISOCAM’s far superior angular resolution enables the mapping of detailed structures even within the spiral arm. On the other hand ISOPHOT’s surface brightness sensitivity is far superior to that of ISOCAM, enabling faint diffuse emission on scales of galactic disks to be traced. ISOCAM’s smaller PSF also meant that different galaxies were generally mapped in the MIR than in the FIR. M 31 and M 33 were too large to be completely mapped with ISOCAM. Observations of a portion of the southern disk of M 31 were presented by Pagani et al. (1999). Due to the vastly different scales, these observations are particularly difficult to relate to those of Haas et al. (1998) but they nevertheless reveal that the detected MIR emission originates mostly in the regions of the spiral arms of Figure 1. In fact, M 31 was completely mapped by Kraemer et al. (2002) at a number of MIR wavelengths with the Midcourse Space Experiment and their 8.3 μm map (their Figure 1) is strikingly similar to that of Haas et al. (1998), with the difference that the diffuse emission detected in the FIR beyond the main spiral arms is not detected in the MIR. This difference could for instance be understood as being due to different parts of the grain size distribution having different spatial extents. However, generalising such a conclusion on the basis of the observation of M 31 alone could be dangerous, as it is known that some of the dust properties of this galaxy are very different from those of other galaxies, as revealed by its atypical MIR spectrum (Cesarsky et al., 1998). A better grasp of the MIR morphology of spiral galaxies can be obtained from studies that completely mapped their targets. Such studies were predominantly made in the 6.75 and 15 μm broad band filters of ISOCAM (occasionally the ISOCAM 12 μm filter was used). These filters were chosen to trace different components of the MIR emission in galaxies. The 6.7 μm filter includes the most prominent spectral features emitted by Polycyclic Aromatic Hydrocarbons (PAHs) powered by UV photons, as revealed by spectroscopic studies of galactic sources. The 15 μm filter was originally chosen to trace stochastic emission from very small grains, though it subsequently turned out that PAH emission from the 12.7 μm feature and an underlying PAH continuum can also contribute and even dominate the signal seen towards the spiral arms. Morphological studies made using the above mentioned filters were presented by Malhotra et al. (1996) for NGC 6946, by Smith (1998) for NGC 7331, by Sauvage et al. (1996) for M 51 as well as by Bendo et al. (2002a), Dale et al. (2000) and Roussel et al. (2001c) for larger samples of spiral galaxies. These papers show that, in the MIR, spiral galaxies present a morphology that is quite similar to that observed at other optical/NIR wavelengths. The spiral arms are very prominent, with the giant H II region complexes showing up as emission enhancements along the arms. A diffuse underlying MIR emitting disk is often detected in the interarm regions of the inner disk. Finally, a central region of varying importance is observed. A clear example of these morphological features is given in Figure 2 where the 6.75 μm map of M 101 from Roussel et al. (2001c) is presented. Although the central region of spiral galaxies is generally dominated optically by the bulge of old stars, one should not jump to the conclusion that the central MIR NORMAL NEARBY GALAXIES 317

Figure 2. The 6.75 μm map of M 101, from Roussel et al. (2001c). The different features of the MIR morphology, central region, H II regions, spiral arms, are clearly seen. Diffuse interarm emission is also present but it becomes rapidly too faint to be detected. component is dominated by the Rayleigh–Jeans emission of cold photospheres or circumstellar envelopes. The extent of the central MIR source is generally different from that of the stellar bulge and furthermore Roussel et al. (2001b) showed that its IR colours are generally not those expected from stars (i.e. the 15/6.75 μm flux ratio is generally larger than unity whereas the opposite applies in early-type galaxies where the stars provide most of the 6.75 μm emission, see Athey et al., 2002 or Xilouris et al., 2004, and Figure 9). In fact, most of the emission from the central regions of spiral galaxies can be attributed to dust emission powered by enhanced star formation (Roussel et al., 2001b). A further detailed study of the morphological features in the MIR is the multiwavelength study of M 83 by Vogler et al. (2005) who present the azimuthally averaged radial profiles of a number of emission components of the galaxy. These profiles show that the MIR emission of M 83 can be decomposed into a central component, associated with the central starburst, the prominent arm structure and a diffuse exponentially fading disk emission. The arm/interarm contrast was found to be larger in the MIR than it is in the optical, although it was found to be smaller than what is observed in Hα. This was interpreted as being again due to the important, if not dominant role of star formation in powering the MIR emission, through the generation of UV photons. The fact that particles responsible for the MIR emission can be excited by a broad range of photon energies (see e.g. Uchida et al., 2000; Li and Draine, 2002) explains why the influence of star-forming regions is seen on broader scales than the actual extent of H II regions (see also Section 2.2.2). An obvious question arising from these investigations pertains to the colour of the various MIR emission components (central regions, spiral arms, interarm regions), which could potentially be used to reveal the heating sources of MIR emission. Searches made in this direction came up with mixed results. For instance 318 M. SAUVAGE ET AL.

Sauvage et al. (1996) reported a systematic variation of the 15/6.75 μm flux ratio along some cuts through the arms of M 51, while Helou et al. (1996) showed that, on average, the 15/6.75 μm flux ratio is the same for the arm and interarm regions of NGC 6946, and Smith (1998) reported no systematic spectral variations in the disk of NGC 7331. The only systematic colour variations are observed in the central regions of spiral galaxies, where the 15/6.75 μm flux ratio can reach values observed in starburst galaxies (Roussel et al., 2001b). Another study of the physical interpretation of the MIR morphology of spiral galaxies was presented by Dale et al. (1999). This shows that although star-formation activity is the primary driver for surface brightness variations inside galaxies, it is only for regions of intense activity that a colour variation occurs. Column density also appears as a secondary parameter to explain the difference in disk surface brightness between galaxies. This relatively fast exploration of the IR morphology of spiral galaxies thus reveals how the same processes that affect their visible shapes, i.e. intensity and distribution of star-forming regions, play an important part in determining their IR aspect. It already suggests that the MIR and FIR part of the spectrum behave differently, in the sense that spectral modifications are more apparent in the FIR than in the MIR. This is not unexpected given the different thermodynamical states of the grains that produce these emissions, thermal equilibrium for most of the FIR, and stochastic heating for most of the MIR.

2.1.3. The Extent of Spiral Disks in the FIR Further information about the true distribution of dust in spiral disks is provided by FIR observations of galaxies more distant than the highly resolved local galaxies discussed in Section 2.1.1, but still close enough to resolve the diffuse disk at the longest FIR wavelengths accessible to ISO. In a study of eight spiral galaxies mapped by ISOPHOT at 200 μm, Alton et al. (1998a; see also Davies et al., 1999 for NGC 6946) showed that the observed scalelength of FIR emission at 200 μm is greater than that found by IRAS at 60 and 100 μm. Thus, the scalelength of the FIR emission increases with increasing FIR wavelength. This result was reinforced using LWS measurements of the dust continuum by Trewhella et al. (2000), and can also be inferred from Figure 2 of Hippelein et al. (2003) for M 33. We note here that this implies that the bulk of the 200 μm emission arises from grains heated by a radially decreasing radiation field, as would be expected for grains in the diffuse disk. If most of the 200 μm emission had arisen from localised sources associated with the parent molecular clouds within the spiral arms, there should be no FIR colour gradient in the galaxy, since the SEDs of the localised sources should not depend strongly on position. The second result to come out of the studies by Alton et al. (1998a) and Davies et al. (1999) is that the observed scalelength at 200 μm is comparable to or exceeds the scalelength of the optical emission (see also Tuffs et al., 1996). As noted by Alton et al., this result implies that the intrinsic scalelength of the dust in galaxies is greater than that of the stars. This is because the apparent scalelength of stars should increase NORMAL NEARBY GALAXIES 319 with increasing disk opacity (since the inner disk is expected to be more opaque than the outer disk) whereas the apparent scalelength of the dust emission will be less than the intrinsic scalelength (due to the decrease in grain temperature with increasing galactocentric radius). The extraction of the precise relation between the intrinsic scalelengths of stars and dust requires a self-consistent calculation of the transfer of radiation through the disk (see Section 2.3). The reason for the difference between the intrinsic scalelength of stars and dust in galaxies is not self-evident, since it is the stars themselves which are thought to be the sources of interstellar grains (produced either in the winds of evolved intermediate mass stars or perhaps in supernovae). One might speculate either that there is a mechanism to transport grains from the inner disk to the outer disk, or that the typical lifetimes of grains against destruction by shocks is longer in the outer disk than it is in the inner disk. While Alton et al. and Davies et al. showed that the scalelength of the 200 μm emission was comparable to or slightly larger than that of the optical emission, these studies did not actually detect grain emission beyond the edge of the optical disk. Since spiral galaxies in the local universe are commonly observed to be embedded in extended disks of neutral hydrogen – the so called “extended H I disks”, it is a natural question to ask whether these gaseous disks contain grains. This question was answered in the affirmative by Popescu and Tuffs (2003), through dedicated deep FIR maps of a large field encompassing the entire H I disk of the edge-on spiral galaxy NGC 891, made using ISOPHOT at 170 and 200 μm (see Figure 3).

Figure 3. The radial profiles of H I emission (from Swaters et al., 1997) convolved with the ISOPHOT PSF (solid line) and of 200 μm FIR emission (symbols) of NGC 891 (Popescu and Tuffs, 2003). Note that the extent and asymmetry of the 200 μm emission follow that of the H I emission. The profiles are sampled at intervals of 18. The negative radii correspond to the southern side of the galaxy and the galaxy was scanned at 60 degrees with respect to the major axis. The units of the FIR profile are W/Hz/pixel, multiplied by a factor of 2 × 10−22 and the error bars are smaller than the symbols. The horizontal bar delineates the FWHM of the ISOPHOT PSF of 93. The vertical arrows indicate the maximum extent of the optically emitting disk. The dotted line represents a modified H I profile obtained in the southern side from the original one by cutting off its emission at the edge of the optical disk and by convolving it with the ISOPHOT PSF. 320 M. SAUVAGE ET AL.

The large amounts of grains found in the extended H I disk (gas-to-dust ratio of ∼1%) clearly shows that this gaseous disk is not primordial, left over from the epoch of galaxy formation. It was suggested that the detected grains could have either been transported from the optical disk (via the halo, using mechanisms such as those proposed by Ferrara (1991), Davies et al. (1998), Popescu et al. (2000a) or through the action of macro turbulence) or that they could have been produced outside the galaxy (for example transferred in interactions with other galaxies). It is interesting to note that, although the dust emission is seen towards the H I component, the grains may not actually be embedded in the neutral ISM. Instead, this dust could trace an “unseen” molecular component, as proposed by Pfenniger and Combes (1994), Pfenniger et al. (1994), Gerhard and Silk (1996), and Valentijn and van der Werf (1999b). This cold molecular gas component has been invoked as a dark matter component to explain the ﬂat rotation curves of spiral galaxies. Its presence might also reconcile the apparent discrepancy between the very low metallicities measured in H II regions in the outer disk (Ferguson et al., 1998) and the high ratio of dust to gas (on the assumption that all gas is in form of H I) found by Popescu and Tuffs (2003) in the extended H I disk of NGC 891.

2.1.4. The Extent of Spiral Disks in the MIR As alluded to in the discussion of the MIR aspect of M 31 above, the MIR disks of spiral galaxies are generally quite short. At first glance they can even appear truncated since in the outer region the emission is generally dominated by the H II regions of the arms, superimposed on a very diffuse disk emission. This appearance is however mostly an “optical” illusion. The radial profiles of the exponentially declining MIR disk emission of M 83 Vogler et al. (2005) show a scalelength that is clearly shorter than that seen in any optical bands, but with no trunca- tion. The strong decrease of star-formation activity in the outer regions of galaxies induces a rapid disappearance of the MIR emission, although, as is definitely revealed by long-wavelength observations, dust is still present at large galactocentric distances. A more systematic study of the extent of MIR spiral disks was presented in Roussel et al. (2001c). These authors have defined a MIR disk diameter, DMIR,in a similar way as the definition used in the optical, by measuring the diameter of the 5 μJy/arcsec−2 isophote at 6.75 μm (this is roughly the detection limit of this band). This definition allows to compare optical and MIR diameters on a broad range of galaxies. It shows that the IR diameter (as defined above) is systematically smaller than the optical one (the largest value of DMIR/Dopt observed in this sample of about 70 galaxies is slightly less than 1, while the smallest value observed is ∼0.3). The variation of DMIR/Dopt can be related to both the morphological type of the galaxies or their H I deficiency in the sense that earlier type or more severely H I deficient galaxies have a smaller DMIR/Dopt. Roussel et al. (2001c) indicate that the H I deficiency is likely the dominant parameter given that H I-deficient NORMAL NEARBY GALAXIES 321 galaxies are often classified as early-type galaxies, due to the absence of star- forming regions at large galactocentric distances. The MIR diameter of a galaxy is thus essentially determined by the extent of the star-forming activity through its disk, demonstrating the crucial role of star formation in powering the MIR emission. The fact that the MIR disks appear so much smaller than the FIR disks immediately prompts the following question: does the dust responsible for the MIR emission disappear at large galactocentric distances, or phrased differently, should we observe a MIR counterpart to the FIR extended distributions? As mentioned before, ISOCAM was probably less sensitive than ISOPHOT to faint diffuse emission, which would bias the ISOCAM maps toward small-scale bright structures such as the H II regions in the disks. This would explain why MIR emission is not detected far out in the disk, but still leaves open the question of why the scalelength of the MIR emission is so much shorter than that of the FIR emission.

2.1.5. Comparison with Morphologies at other Wavelengths 2.1.5.1. Comparison with the Optical. Although one immediately associates IR light with dust, the fact that ISO could reach relatively short IR wavelengths prompts the question of the contribution of stars to the shortest ISO bands. And indeed, when studied quantitatively through the use of light concentration indices, as done for instance by Boselli et al. (2003b), the morphology of galaxies at MIR and NIR wavelengths show surprising similarities. This indicates at least a similar spatial distribution of dust and stars, which is not too surprising given the origin of the MIR emission (see later). However, determining the actual contribution of stellar photospheres and circumstellar shells to the ISO fluxes is harder, as it requires first the construction of a pure stellar image and second the correct extrapolation in wavelength of this image to the ISO wavelengths, taking into account the possible stellar population gradients in the galaxy. This type of work led Boselli et al. (2003a) to conclude that a large fraction of the 6.75 μm flux in galaxies could be due to stars, ranging from 80% in Sa to 20% in Sc galaxies. This conclusion is however not supported by the analysis of Lu et al. (2003) who showed that already at 6 μm, the mean flux of their galaxies is one order of magnitude larger than what is observed in galaxies devoid of interstellar medium (see Figure 9). Therefore it appears unlikely that stars make a large contribution to the IR emission of spiral galaxies longwards of 6 μm, even though the spatial distribution of the emission shares some morphological properties with that of the stars, a fact which is due to their common link with the star-formation process.

2.1.5.2. Comparison with Hα. For the highly resolved galaxy M 33, Hippelein et al. (2003) showed that there is a strong resemblance between the morphology μ l of the localised warm dust component at 60 m(F60; obtained at each direction by subtracting the 170 μm map scaled by the ratio of the 60/170 μm brightness 322 M. SAUVAGE ET AL.

μ l Figure 4. Left: Distribution of the localised warm dust component at 60 m, F60, in M 33 (Hippelein et al., 2003). This is the scaled difference map 2(F60 − 0.165 × F160), with the factor 0.165 given by the average ﬂux density ratio F60/F170 in the interarm regions. Right:Hα map of M 33 convolved to a resolution of 60.

μ l Figure 5. The ratio of FHα to the localised warm dust component at 60 m, F60, for the star-forming regions in M 33 (Hippelein et al., 2003) versus distance from the galaxy centre. Symbol sizes indicate the brightness of the sources. in the interarm regions) and the morphology of the Hα emission (see Figure 4), indicating that the 60 μm localised emission traces the star-formation complexes. / l The FHα F60 ratio (see Figure 5) for the star-formation complexes shows a clear systematic increase with increasing radial distance from the centre (allowing for the [N II] line contribution decreasing with distance, the slope would be even steeper). NORMAL NEARBY GALAXIES 323

Veryprobably this is due to the presence of a large scale gradient of opacity affecting the recombination line ﬂuxes. Similarly, in the MIR a number of authors have noted a striking resemblance between MIR and Hα maps of galaxies, which should not come as a surprise given the numerous links between MIR emission and star formation mentioned above. This is both the case for the well-studied galaxies M 51 (Sauvage et al., 1996), NGC 7331 (Smith, 1998), M 83 (Vogler et al., 2005) and NGC 6946 (Walsh et al., 2002) and for larger samples of more distant objects (see for instance Contursi et al., 2001 or Domingue et al., 2003). However, few authors have studied the relation between the two emissions within galaxies. Sauvage et al. (1996) presented a comparison of the MIR/Hα ratio for a number of H II region complexes in M 51. This reveals that extinction affecting the measure of Hα is responsible for most of the variation observed in the MIR/Hα ratio. One should note that the apparently strong spatial correlation between the MIR and Hα maps is slightly misleading, in the sense that the limited spatial resolution tends to make the maps more similar than they really are. Studies of resolved H II regions indeed show that the MIR emission tends to avoid the peaks of Hα emission and rather delineates the edges of molecular clouds exposed to the radiation produced by newly formed stars (see e.g. Contursi et al., 1998 or Contursi et al., 2000 and Figure 8). The complexity of the local relation between the MIR emission and Hα is revealed by the studies of Vogler et al. (2005) and Walsh et al. (2002). Using maps at the same resolution, and only taking independent points, these authors showed that the local MIR-Hα correlation is quite poor and highly non-linear. As we will see, this is in contrast to the global MIR-Hα correlation, which is linear and tighter. A plausible explanation of this, ﬁrst introduced in the context of the FIR emission by Popescu et al. (2002; see also Pierini et al., 2003b), would be the existence of a diffuse component of the MIR emission from the underlying disk, powered by the star-formation activity, but on larger scales than that used by Walsh et al. (2002) or Vogler et al. (2005).

2.1.5.3. Comparison with UV. A fundamental property of spiral galaxies is the fraction of light from young stars which is re-radiated by dust. This property can be investigated as a function of position in the galaxy by a direct comparison of ISOPHOT maps at 60, 100 and 170 μm with UV maps obtained with Galaxy Evolution Explorer (GALEX; Martin et al., 2004) in its near-UV (NUV; 2310 A)˚ and far-UV (FUV; 1530 A)˚ bands. Such a comparison was performed for M 101 by Popescu et al. (2004b). The top panels in Figure 6 display the 100 μm ISOPHOT image (left) together with the corresponding “total UV” (integrated from 1412 to 2718 A)˚ image (right). Comparison between the ratio image (100 μm/UV) (bottom left panel) and an image of the “spiral arm fraction” (the fraction of the UV emission from the spiral arm within an ISOPHOT beam; bottom right panel) shows that the high values of the 100 μm/UV ratio trace the interarm regions. In other words the 324 M. SAUVAGE ET AL.

Figure 6. FIR-UV comparison for M 101 (Popescu et al., 2004b). Top left panel: Filter-integrated 100 μm ISOPHOT image. Top right panel: “Total UV” image converted to the orientation, resolution and sampling of the 100 μm ISOPHOT image. Bottom left panel: The ratio image of the ﬁlter- integrated 100 μm ISOPHOT image divided by the corresponding “total UV” image. Bottom right panel: The image of the “spiral arm fraction” at the orientation, resolution and sampling of the 100 μm ISOPHOT image. All panels depict a ﬁeld of 27.7 ×27.1 centered at α2000 = 14h03m13.11s; δ2000 = 54◦2106.6. The pixel size is 15.33 × 23.00.

“spiral features” in the ratio image are in reality regions of diffuse emission which are interspaced with the real spiral features, as seen in the “spiral arm fraction” image. The trend for the FIR/UV ratio to be higher in the diffuse interarm regions than in the spiral-arms is seen in Figure 7 from the segregation of the blue dots and red diamonds at a given radius. This apparently surprising result was explained in terms of the escape probability of UV photons from spiral arms and their subsequent scattering in the interarm regions, and in terms of the larger relative contribution of optical photons to the heating of the dust in the interarm regions. The combined NORMAL NEARBY GALAXIES 325

Figure 7. The pixel values of the FIR/UV ratio map of M 101 (Popescu et al., 2004b) at the resolution of the 170 μm image versus angular radius. The dots are for lines of sight towards interarm regions and the diamonds towards the spiral arm regions. The solid line is an offset exponential ﬁt to the data.

effect of the optical heating and the scattering of the UV emission means that the FIR/UV ratio will not be a good indicator of extinction in the interarm region. Despite these local variations, the main result of Popescu et al. (2004b) is the discovery of a tight dependence of the FIR/UV ratio on radius, with values monotonically decreasing from ∼4 in the nuclear region to nearly zero towards the edge of the optical disk (see Figure 7). This was interpreted in terms of the presence of a large-scale distribution of diffuse dust having a face-on optical depth which decreases with radius and which dominates over the more localised variations in opacity between the arm and interarm regions. A comparison of the UV and MIR morphologies, analogous to the UV-FIR comparison for M 101, has been made for the south-west ring of M 31 by Pagani et al. (1999). Here the 20 resolution 200 nm map from the FOCA 1000 balloon experiment is compared to the ∼6 resolution 6.75 μm ISOCAM map (Figure 8). This study, and especially its Figure 8, shows that although the overall morphology of the two maps appears similar, following the ring-like structure observed in that region of the galaxy, the UV and MIR emission actually complement each other: UV-emitting regions ﬁll the holes in the distribution of the 6.75 μm emission. In fact, the same study shows a much tighter spatial correlation of the MIR emission with the gaseous components (H I or CO) than with the UV emission. Again, this is interpreted as showing that even if most of the energy that ultimately powers the MIR emission is generated in the giant H II region complexes that are seen in Hα or UV light, most of the MIR emission comes from the surfaces of the molecular clouds that surround these complexes. 326 M. SAUVAGE ET AL.

Figure 8. The SW ring of M 31 observed by Pagani et al. (1999). On the left panel, darker Hα contours are superimposed on a grey scale map at 6.75 μm. On the right panel, UV contours are superposed on the same map. In both maps, though the overall ring-like structure is preserved at all wavelengths, it is clear that the UV and Hα emissions tend to be located in regions of minimal 6.75 μm emission, and that reversely, the MIR emission is located on the edges of the Hα and UV regions.

These differences in the UV and MIR spatial distribution are conﬁrmed by the recent observations of the Spitzer satellite, which also reveal that at 24 μm the association between IR and UV or Hα emission becomes tighter, even at high spatial resolution (Helou et al., 2004). The reason for this will be elucidated in Section 2.2.2.

2.2. INTEGRATED PROPERTIES

The overriding result of all ISOPHOT studies of the integrated properties of normal galaxies in the FIR is that their SEDs in the 40–200 μm spectral range require both warm and cold dust emission components to be fitted. Although the concept of warm and cold emission components is as old as IRAS (de Jong et al., 1984), it only became possible to directly measure and spectrally separate these components using ISOPHOT’s multi-filter coverage of the FIR regime out to 200 μm. In the MIR the two most stunning results are the high similarity of the disk’s SED in spiral galaxies and the almost complete invariance of the global MIR colours to the actual star-formation activity of the galaxy as a whole. The former result shows that the MIR emission comes predominantly from PDRs, whereas the latter result demonstrates the importance of the emission from the central region to the MIR, a characteristic which could only be uncovered due to the high spatial resolution offered by ISOCAM. In order to investigate the integrated properties of local universe gas-rich galaxies, a number of statistical samples were constructed, which were primarily NORMAL NEARBY GALAXIES 327 observed by ISOPHOT and ISOCAM. All these projects were complementary in terms of selection and observational goals. In descending order of depth (measured in terms of a typical bolometric luminosity of the detected objects), the published surveys are as following. The ISO Virgo Cluster Survey is a combination of the ISOPHOT Virgo Cluster Deep Survey (IVCDS; Tuffs et al., 2002a,b; Popescu et al., 2002) and ISOCAM measurements (Boselli et al., 2003b, see also Boselli et al., 1997) of the same sample supplemented by additional targets on the eastern side of the cluster. The IVCDS represents the deepest survey (both in luminosity and surface brightness terms) of normal galaxies measured in the FIR with ISO. A complete volume- and luminosity- limited sample of 63 gas-rich Virgo cluster galaxies selected from the Virgo Cluster Catalogue (Binggeli et al., 1985; see also Binggeli et al., 1993) with Hubble types later than S0 and brighter than BT ≤ 16.8 were mapped with ISOPHOT at 60, 100 and 170 μm. A total of 99 galaxies were mapped with ISOCAM in the 6.7 and 15 μm filters. The IVCDS sample was (in part) also observed with the LWS (Leech et al., 1999). The ISO Virgo cluster survey provides a database for statistical investigations of the FIR and MIR SEDs of gas-rich galaxies in the local universe spanning a broad range in star-formation activity and morphological types, including dwarf systems and galaxies with rather quiescent star-formation activity. The Coma/A1367 Survey (Contursi et al., 2001) consists of 6 spiral and 12 irregular galaxies having IRAS detections at 60 μm. The galaxies were selected to be located within 2 or 1 degrees of the X-ray centres of Coma and A1367 clusters, respectively, with emphasis on peculiar optical morphologies. Each galaxy was observed in a single pointing with ISOPHOT, at 120, 170 and 200 μm, as well as mapped with ISOCAM in the 6.75 and 15 μm broadband filters. The sample provides a database of integrated flux densities for a pure cluster sample of high luminosity spiral and irregular galaxies. The ISO Bright Spiral Galaxies Survey (Bendo et al., 2002a,b; 2003) consists of 77 spiral and S0 galaxies chosen from the Revised Shapley-Ames Catalog (RSA), with BT ≤ 12.0. Almost all are IRAS sources. Mainly an ISOCAM mapping survey with the 12 μm filter, the project also used ISOPHOT to take 60, 100 and 170 μm short stares towards the nucleus of the galaxies and towards background fields. The sample provides a database of MIR morphologies and FIR surface brightnesses of the central regions of bright spiral galaxies, including S0s. The ISO Bright Spiral Galaxies Survey and the ISO Virgo cluster survey represent the principle investigations of optically selected samples of normal galaxies. It should be emphasised that the main difference between them is primarily one of shallow versus deep, rather than field versus cluster, since by design the Virgo sample predominantly consists of infalling galaxies from the field, and no cluster specific effects could be found (see also Contursi et al. for Coma/A1367 Sample). The ISO Key Project on Normal Galaxies Sample (Dale et al., 2000) consists of 69 galaxies selected to span the whole range of the classical IRAS colour-colour 328 M. SAUVAGE ET AL. diagram (Helou, 1986). Since IRAS detected a vast number of galaxies in its four bands, the selection was also made to span the Hubble sequence evenly and provide a broad range of IR luminosities, dust temperatures (as determined by IRAS) and star-formation activity. Galaxies in the sample were mapped with ISOCAM, their main cooling lines in the FIR were measured with ISOLWS (Malhotra et al., 2001) and their 2.5–12 μm spectra were measured with ISOPHOT-S (Lu et al., 2003). The ISOCAM Parallel Mode Survey (Ott et al., 2003) is drawn from observations made with ISOCAM, mostly around 6 μm, while another ISO instrument was prime. It detected around 16,000 distinct objects down to 0.5 mJy. Identifications of these objects by cross-correlation with other catalogues is on-going and ∼25% of the objects are expected to be galaxies. The ISOPHOT Serendipity Survey (Stickel et al., 2000) has initially catalogued 115 galaxies with Sν ≥ 2 Jy at 170 μm and with morphological types predominantly S0/a - Scd. This sample provides a database of integrated 170 μm flux densities for relatively high luminosity spiral galaxies, all detected by IRAS at 60 and 100 μm. Recently a catalogue of 1,900 galaxies was released (Stickel et al., 2004), of which a small fraction does not have IRAS detections. Most of the 1,900 galaxies are spirals. The measured 170 μm flux densities range from just below 0.5 Jy up to ∼ 600 Jy. Finally about 30 additional spiral galaxies were mapped in various ISO guaranteed and open time programs (see Roussel et al., 2001c, for a compilation)

2.2.1. The FIR Spectral Energy Distribution: Dust Temperatures, Masses and Luminosities The presence of a cold dust emission component peaking longwards of 120 μm was inferred from studies of the integrated SEDs of individual galaxies (see Section 2.1.1), from statistical studies of small samples (Krügel et al., 1998; Sieben- morgen et al., 1999) and was confirmed and generalised by studies of the larger statistical samples mentioned above. The latter studies also demonstrated the uni- versality of the cold dust component, showing it to be present within all types of spirals (Tuffs and Popescu, 2003). The cold emission component predominantly arises from dust heated by the general diffuse interstellar medium and the warm component from locally heated dust in H II regions, an interpretation consistent with what has been seen in the ISOPHOT maps of nearby galaxies (see Section 2.1.1) and with self-consistent modelling of the UV-FIR SEDs (see Section 2.3). The cold dust component is most prominent in the most “quiescent” galaxies, like those contained in the IVCD sample, where the cold dust temperatures were found to be broadly distributed, with a median of 18 K (Popescu et al., 2002), some 8–10 K lower than would have been predicted by IRAS. The corresponding dust masses were found to be increased by factors of typically 2–10 (Stickel et al., 2000) for the serendipity sample and by factors 6–13 (Popescu et al., 2002) for the IVCD sample, with respect to previous IRAS determinations. As a consequence, the derived gas-to-dust ratios are much closer to the canonical value of ∼160 for NORMAL NEARBY GALAXIES 329 the Milky Way (Stickel et al., 2000; Contursi et al., 2001; see also Haas et al., 1998 for M 31), but with a broad distribution of values (Popescu et al., 2002). It was found that the cold dust component provides not only the bulk of the dust masses, but even the bulk of the FIR luminosity, in particular for the case of the most quiescent spirals, like those in the IVCD sample. In contrast to the SEDs found by the other ISOPHOT studies, which typically peaked at around 170 μm, Bendo et al. (2003) derived spatially integrated SEDs typically peaking at around 100 μm. The result of Bendo et al. may reflect the fact that these observations were single pointings, made towards the nucleus of resolved galaxies extending (in the main) beyond the field of view of ISOPHOT, and were therefore biased towards nuclear emission, which is warmer than the extended cold dust emission missed (or only partially covered) by these measurements. Nevertheless, the measurements of Bendo et al. constitute a useful probe of the FIR emission of the inner disks. This emission (normalised to K-band emission) was found to increase along the Hubble sequence (Bendo et al., 2002b). Since the FIR carries most of the dust luminosity, it is interesting to re-evaluate the question of the fraction of stellar photons converted via grains into IR photons, taking into account the comprehensive measurements of the cold dust emission component made available by ISOPHOT. This was done by Popescu and Tuffs (2002a), who showed that the mean percentage of stellar light reradiated by dust is ∼30% for the Virgo cluster spirals contained in the IVCD sample. This study also included the dust emission radiated in the NIR-MIR range. The fact that the mean value of ∼30% found for the Virgo cluster spirals is the same as the canonical value obtained for the IRAS Bright Galaxy Sample (BGS; Soifer and Neugebauer, 1991) is at first sight strange, since IRAS was not sensitive to the cold dust component. However the BGS sample is an IR selected sample and biased towards galaxies with higher dust luminosities, while the Virgo sample is optically selected and contains a full representation of quiescent systems. So the deficit in FIR emission caused by sample selection criteria for the Virgo sample is compensated for by the inclusion of the cold dust component. Popescu and Tuffs (2002a) also found evidence for an increase of the ratio of the dust emission to the total stellar emitted output along the Hubble sequence, ranging from typical values of ∼15% for early spirals to up to ∼50% for some late spirals. This trend was confirmed by Boselli et al. (2003a) who further utilised the new ISO data on dust emission to constrain the corresponding absorption of starlight and thus improve extinction corrections (using the technique pioneered by Xu and Buat (1995) for the IRAS data).

2.2.2. The NIR to MIR Spectral Energy Distribution Using the large body of NIR-MIR data – both photometric and spectroscopic – from the statistical studies mentioned previously, one can also determine the general properties of the SED of galaxies from the NIR to the MIR. Concerning the colours of galaxies in the MIR, the main property that is evidenced is a remarkable uniformity of the 6.75/15 μm flux ratio. In contrast with the 330 M. SAUVAGE ET AL. well-known IRAS colour-colour diagram, the composite ISO-IRAS colour diagram (showing the 6.75/15 μm flux ratio versus the 60/100 μm flux ratio, see Dale et al., 2000) shows a plateau of the 6.75/15 μm flux ratio for a large range of 60/100 μm flux ratios. A decrease of the ISO colour is only observed for the largest values of the IRAS colour. A rapid interpretation of that behaviour is that for most galaxies, the ISOCAM filters collect emission from transiently heated PAHs, which shows no spectral dependence on the heating intensity. It is only for the most actively star-forming galaxies that a fraction of the 15 μm flux originates in hot dust located in the H II regions. This simple interpretation has to be taken with caution since it assumes that the integrated colour reflects a global property of the object (e.g. its star-forming activity). Indeed, following up on the fact that in the MIR the central regions of spiral galaxies appear to differ from the disk (see Section 2.1.2), Roussel et al. (2001b) inspected the behaviour of the 6.75/15 μm flux ratio separately for disks and central regions. This revealed that the disk colours are extremely similar from one galaxy to the other and that it is the fraction of the total flux emitted from the central region combined with the colour of that region that dominate the variation of the global 6.75/15 μm flux ratio, rather than the global level of star-formation activity (Figure 9). We remark that, in principle, the decrease of the 6.75/15 μm flux ratio in the central regions could correspond to a drop of the 6.75 μm flux rather than to an increase of the 15 μm flux. That this is not the case is clearly demonstrated by the ISOCAM spectra of these central regions (Roussel et al., 2001b) that show the appearance of the expected small grains continuum at the long wavelength edge of the bandpass.

Figure 9. On the left panel, we show the distribution of the 15/6.75 μm ﬂux ratio for the disks and central regions of spiral galaxies (taken from Roussel et al., 2001b). It is clear that the disks have a very uniform colour while the central regions explore a much larger range. On the left side, the global NIR-MIR SED (from Lu et al., 2003) reveals the dominance of the PAH features in the MIR, as well as the very small dispersion from one object or the other. Both PAH dominated spectra represent the average spectrum of the sample, with two different normalisation. The very small dispersion bars reveal the homogeneity of the sample. The thin solid and dotted lines are spectra from early-type galaxies, dominated by stars. NORMAL NEARBY GALAXIES 331

This uniformity of the disk emission is confirmed by the spectroscopic study of Lu et al. (2003). With the spectroscopic capabilities of ISOPHOT-S, they showed that the 2.5–12 μm spectrum of spiral galaxies is indeed extremely uniform. The NIR part is dominated by the tail of photospheric emission, while the MIR part presents the classical features of PAH emission, with extremely small spectral variations of these features from one galaxy to the other (of the order of 20%, see Figure 9). Furthermore, these variations are not related to the IRAS 60/100 μm colour or to the FIR-to-blue luminosity ratio. Thus, as revealed by the broadband colours, it appears that the detailed SEDs of spiral galaxies in the MIR show very little dependence on the actual level of star formation in the galaxy. Studies of the MIR SEDs of various regions in our own galaxies have shown that a spectrum such as that observed globally in spiral galaxies is prototypical of photo-dissociation regions (PDRs), e.g. the outer layers of molecular clouds exposed to moderately strong radiation fields (see e.g. van Dishoeck, 2004). Hence the spatial distribution of the MIR emission revealed in Section 2.1 is fully consistent with the observed SED. It should be noted, however, that the ISOPHOT-S spectral range is shorter than the one explored by the 15 μm filter of ISOCAM, and in fact covers a wavelength range that is not very sensitive spectrally to activity variations, as demonstrated by the study of starburst galaxies presented by Förster-Schreiber et al. (2003). Therefore it is from the combination of the nearly constant colours observed by Roussel et al. (2001b) and the constant spectrum observed by Lu et al. (2003) that one can conclude that the 5–18 μm spectrum of galactic disks is dominated by the emission from PDRs, rather than from the H II regions themselves. The systematic exploration of the NIR-MIR SED of spiral galaxies made possible by ISO offered a number of other important discoveries. First Lu et al. (2003) evidenced a new component of dust emission consisting of a “hot” NIR continuum emission. Its association with the ISM was made clear by the existence of a correlation between that NIR excess and the strength of the PAH features. However a physical explanation of this continuum is still lacking. Second, the strength of the PAH emission (which represents most, but not all of the MIR flux when defined as the 5–20 μm luminosity for instance), was found to be well correlated spatially with the 850 μm flux of bright regions inside a number of galaxies (Haas et al., 2002). This can be interpreted as showing a close physical association between cold dust clouds and the PAHs, but, given the limited spatial resolutions involved, more likely indicates that the localised 850 μm emission comes essentially from the inside of the molecular clouds whose surfaces produce the PAH features. Finally Helou et al. (2001) found a strong correlation between the 5–10 μm luminosity and the [C II] line luminosities of actively star-forming galaxies from the sample of Malhotra (2001). Since the [C II] emission from this sample mainly originates in PDRs (see Section 2.4.1), this correlation implies that PAHs, which dominate the 5–10 μm luminosities (Lu et al., 2003), are also responsible for most of the gas heating in strongly star-forming galaxies. 332 M. SAUVAGE ET AL.

2.2.3. Two Outstanding Questions of the IRAS Era 2.2.3.1. The Radio-FIR Correlation. One of the most surprising discoveries of the IRAS all-sky survey was the very tight and universal correlation between the spatially integrated FIR and radio continuum emissions (de Jong et al., 1985; Helou et al., 1985; Wunderlich et al., 1987; see Völkand Xu, 1994 for a review). However all the pre-ISO studies of the FIR/radio correlation were based on FIR luminosities derived from the IRAS 60 and 100 μm flux densities, and thus were missing the bulk of the cold dust luminosity. The ISOPHOT measurements at 60, 100 and 170 μm were used to redefined the FIR/radio correlation (Pierini et al., 2003b) for a statistical sample of spiral galaxies. The inclusion of the cold dust component was found to produce a tendency for the total FIR/radio correlation to become more non-linear than inferred from the IRAS 60 and 100 μm observations. The use of the three FIR wavelengths also meant that, for the first time, the correlation could be directly derived for the warm and cold dust emission components. The cold FIR/radio correlation was found to be slightly non-linear, whereas the warm FIR/radio correlation is linear. Because the effect of disk opacity in galaxies would introduce a non-linearity in the cold-FIR/radio correlation, in the opposite sense to that observed, it was argued that both the radio and the FIR emissions are likely to have a non-linear dependence on SFR. For the radio emission an enhancement of the small free-free component with SFR can account for this effect. For the cold FIR emission a detailed analysis of the dependence of local absorption and opacity of the diffuse medium on SFR is required to understand the non-linear trend of the correlation (see Pierini et al., 2003b). The improved angular resolution of ISO compared with IRAS also allowed a more detailed examination of the local FIR-radio correlation on sub-kpc size galactic substructures. Hippelein et al. (2003) established the correlation for the star-forming regions in M 33. This correlation is shown in Figure 10, overplotted on the correlation for integrated emission from galaxies. It is apparent that the local correlation has a shallower slope (of the order of 0.9) than for the global correlation. It was argued that the local correlation is attributable to the increase with SFR of dust absorption in increased dust densities, and to local synchrotron emission from within supernova remnants, still confining their accelerated electrons. Both emission components play only a minor role in the well-known global radio- FIR correlation, that depends on the dominant large-scale absorption/re-emission properties of galaxies. This local correlation has also been investigated in the MIR by Vogler et al. (2005). A good correlation was found between the two emissions on scales of ∼500 pc, which is comparable to the scales of the individual regions considered for the local FIR-radio correlation in M 33. The slope of the local MIR-radio correlation was comparable to that of the local FIR-radio correlation. This is consistent with the MIR and FIR emissions on scales of ∼ 500 pc being powered by the same UV photons. NORMAL NEARBY GALAXIES 333

Figure 10. Plot of the monochromatic radio luminosity versus the FIR luminosities for M 33 (Hippelein et al., 2003) and its star-forming regions (ﬁlled circles) together with the data for the Effelsberg 100-m galaxy sample (Wunderlich et al., 1987, open circle). The dotted line has a slope of 1.10 (Wunderlich and Klein, 1988).

2.2.3.2. Infrared Emission as a Star-Formation Tracer. Because star-forming regions and H II complexes shine brightly in the IR, and because starburst galaxies output most of their energy in the IR, this wavelength domain has always been associated with star formation. As we have seen earlier, this is supported by a qualitative analysis of the IR emission of galaxies, but a quantitative assessment of this association has proven difficult. With the enhanced capacities of ISO, this question has been addressed in much more detail, and this reveals that the different regions of the IR SED have a different link with star formation. In the FIR, the correlation with the most widely used star-formation tracer, Hα, is extremely non-linear, which has long led authors to suspect that the FIR emission is the result of more than one component (see e.g. Lonsdale-Persson and Helou, 1987). With ISOPHOT the correlation was separately established for the warm and cold dust emission components. A good linear correlation was found between the warm FIR luminosities (normalised to the K-band luminosity) derived for the IVCD sample and their Hα EW (Popescu et al., 2002). This is in agreement with the assumption that the warm dust component is mainly associated with dust locally heated within star-formation complexes. The scatter in the correlation was attributed to a small component of warm emission from the diffuse disk (produced either by transiently heated grains or by grains heated by old stellar population), as well as to the likely variation in H II region dust temperatures within and between galaxies. A good but non-linear correlation was found between the cold FIR luminosities of the galaxies from the IVCD sample and their Hα EW, in the sense that FIR increases more slowly than Hα. Since the bulk of the cold FIR emission arises from the diffuse disk, the existence of this correlation implies that 334 M. SAUVAGE ET AL. the grains in the diffuse disk are mainly powered by the UV photons (see also Section 2.3). The non-linearity of the correlation is consistent with there being a higher contribution from optical photons in heating the grains in more quiescent galaxies. For the late-type galaxies in the Coma and A1367 clusters, Contursi et al. (2001) derived the relationships between the IR flux densities at 200, 170, 120, 100 and 60 μm, normalised to the H-band flux, as a function of the Hα EW. It was found that the poorer correlation is in the 200 μm band and that the values of the fitted slopes decrease as the FIR wavelength increases. These results should be interpreted in terms of the increasing contribution of the diffuse component with increasing FIR wavelength. Using the ISOPHOT observations from Tuffs et al. (2002a,b) and Stickel et al. (2000) in combination with UV- and K-band photometry, Pierini and Möller(2003) have tried to quantify the effects of optical heating and disk opacity on the derivation of SFR from FIR luminosities. For this they investigated trends in the ratio of the far-IR luminosity to the intrinsic UV luminosity, Ldust/LUV, with both disk opacity and disk mass (as measured by the intrinsic K-band luminosity). Using a separate relation between disk opacity and K-band luminosity they were able to re-express Ldust/LUV in terms of a single variable, the galaxy mass. In this way they found evidence for the relative importance of optical photons in heating dust to increase with increasing galaxy mass. In the MIR, the correlation of global MIR and Hα luminosities is also non- linear (Roussel et al., 2001a), but since the MIR SED is so constant from galaxy to galaxy, an explanation along the lines developed for the FIR correlation is not possible. However, Roussel et al. (2001a) have shown that the non-linearity is due again to the fact that the global MIR flux mixes together the contributions from the central regions and the disk. Restricting their studies to the disk emission, they showed that both the 6.75 μm and the 15 μm luminosities are linearly correlated with the extinction-corrected Hα luminosities (see Figure 11). Therefore, although the MIR emission from the disk of spiral galaxies does not originate directly in the star-forming regions, but rather from the PDRs around them, it is the energy from the star-formation process that powers the MIR luminosity, and Roussel et al. provide the conversion factors that allow the derivation of the star-formation rate (SFR) from the MIR luminosities. In a follow-up study, Förster-Schreiber et al. (2004) demonstrated that the 6.75 μm luminosity is still as good a tracer of star formation in the central regions of spiral galaxies as it was in their disks. The non-linearity observed in the global 6.75-Hα diagram is attributed to the difference in extinction between disk and central star-forming regions. For the 15 μm luminosities, Förster- Schreiber et al. (2004) showed that the star-formation activity observed in the central regions of galaxies is such that it shifts the emission from very small grains (usually emitting in the 20–60 μm range) toward shorter wavelengths. Therefore, even when the Hα emission is corrected for the higher central extinction, the correlation with NORMAL NEARBY GALAXIES 335

Figure 11. The MIR-Hα correlation for the disks of spiral galaxies (adapted from Roussel et al., 2001a). Here only the disk fraction of the MIR flux is used, and the Hα fluxes are uniformly corrected for 1 magnitude of extinction. All fluxes are normalised by the galaxy size to remove size-related biases. For both fluxes the correlation is extremely linear. Lines on the graphs correspond to different adjustment methods, but all give a slope very close to unity. the 15 μm emission is not linear as grains that make the IR emission change their thermal regime.

2.3. QUANTITATIVE INTERPRETATION OF FIR SEDS

When considered in isolation, the FIR luminosity of normal galaxies is a poor estimator of the SFR. There are two reasons for this. First, these systems are only partially opaque to the UV light from young stars, exhibiting large variations in the escape probability of UV photons between different galaxies. Second, the optical luminosity from old stellar populations can be so large in comparison with the UV luminosity of young stellar populations that a significant fraction of the FIR luminosity can be powered by optical photons, despite the higher probability of absorption of UV photons compared to optical photons. One step towards a quantitative interpretation of FIR SEDs was achieved by the semi-empirical models of Dale et al. (2001) and Dale and Helou (2002). These authors have used ISO observations to develop a family of templates to fit the variety of the observed forms of the IR SEDs. This work assumes a power law distribution of dust masses with local radiation field intensity to provide a wide range of dust temperatures. It appears to indicate that the SEDs of star-forming galaxies can be fitted by a one-parameter family of curves (characterised by the 60/100 μm colour), determined essentially by the exponent, α, of the power law distribution of dust masses with the radiation field intensity, where the radiation field is assumed to have the colour of the local ISRF. This calibration method has been extensively used to predict FIR flux densities longwards of 100 μm for the galaxies not observed by ISOPHOT. A comparison of the FIR continuum levels 336 M. SAUVAGE ET AL.

Figure 12. Comparison of observed FIR continuum levels observed by ISO with the model predictions of Dale and Helou (2002). The circles derive from the ISO LWS templates and the triangles represent 170 μm data from the ISOPHOT serendipity survey (Stickel et al., 2000). observed by ISO with the Dale and Helou (2002) model predictions is shown in Figure 12. However, a quantitative interpretation of dust emission in terms of SFRs and star-formation histories requires a combined analysis of the UV-optical/FIR- submm SEDs, embracing a self-consistent model for the propagation of the photons. There are a few such models in use which incorporate various geometries for the stellar populations and dust (Silva et al., 1998; Bianchi et al., 2000a; Charlot and Fall, 2000; Popescu et al., 2000b). We will concentrate here on the model of Popescu et al., since this is the only model which has been used to make direct predictions for the spatial distribution of the FIR emission for comparison with the ISO images. This model also ensures that the geometry of the dust and stellar populations is consistent with optical images. Full details are given by Popescu et al. (2000b), Misiriotis et al. (2001), Popescu and Tuffs (2002b) and Tuffs et al. (2004). In brief, the model includes solving the radiative-transfer problem for a realistic distribution of absorbers and emitters, considering realistic models for dust, taking into account the grain-size distribution and stochastic heating of small grains and the contribution of H II regions. The FIR-submm SED is fully determined by just three parameters: the star-formation rate (SFR), the dust mass Mdust associated with the young stellar population, and a factor F, deﬁned as the fraction of non-ionising UV photons which are locally absorbed in H II regions around the massive stars. A self-consistent theoretical approach to the calculation of the F factor is given by Dopita et al. (2004) in the context of modelling SEDs of starburst galaxies. Popescu et al. (2000b) illustrated their model with the example of the edge-on galaxy NGC 891, which has been extensively observed at all wavelengths (including the complete submm range; Dupac et al., 2003), and also mapped with ISOPHOT at 170 and 200 μm (Popescu et al., 2004a). A particularly stringent test of the model was to compare its prediction for the radial proﬁle of the diffuse dust emission NORMAL NEARBY GALAXIES 337

Figure 13. Left: The radial profile of NGC 891 at 170 μm (Popescu et al., 2004a) produced by integrating the emission parallel to the minor axis of the galaxy for each bin along the major axis. Solid line: model prediction; diamonds: observed profile; dotted line: beam profile. Right: The fractional contribution of the three stellar components to the FIR emission of NGC 891 (Popescu et al., 2000b). component near the peak of the FIR SED with the observed radial profiles. This comparison (see Figure 13, left panel) done in Popescu et al. (2004a) showed a remarkable agreement. The excess emission in the observed profile with respect to the predicted one for the diffuse emission was explained in terms of two localised sources. At 170 and 200 μm the model for NGC 891 predicts that the bulk of the FIR dust emission is from the diffuse component. The close agreement between the data and the model predictions, both in integrated flux densities, but especially in terms of the spatial distribution, constitutes a strong evidence that the large-scale distribution of stellar emissivity and dust predicted by the model is in fact a good representation of NGC 891. In turn, this supports the prediction of the model that the dust emission in NGC 891 is predominantly powered by UV photons. Depending on the FIR/submm wavelength, the UV-powered dust emission arises in different proportions from within the localised component (H II regions) and from the diffuse component (Figure 13; right panel). For example at 60 μm, 61% of the FIR emission in NGC 891 was predicted to be powered by UV photons locally absorbed in star-forming complexes, 19% by diffuse UV photons in the weak radiation fields in the outer disk (where stochastic emission predominates), and 20% by diffuse optical photons in high energy densities in the inner part of the disk and bulge. At 100 μm the prediction was that there are approximately equal contributions from the diffuse UV, diffuse optical and locally absorbed UV photons. At 170, 200 μm and submm wavelengths, most of the dust emission in NGC 891 was predicted to be powered by the diffuse UV photons. The analysis described above does not support the preconception that the weakly heated cold dust (including the dust emitting near the peak of the SED sampled by the ISOPHOT 338 M. SAUVAGE ET AL. measurements presented here) should be predominantly powered by optical rather than UV photons. The reason is as follows: the coldest grains are those which are in weaker radiation fields, either in the outer optically thin regions of the disk, or because they are shielded from radiation by optical depth effects. In the first situation the absorption probabilities of photons are controlled by the optical properties of the grains, so the UV photons will dominate the heating. The second situation arises for dust associated with the young stellar population, where the UV emissivity far exceeds the optical emissivity.

2.4. THE GASEOUS ISM

2.4.1. The Atomic Component Most ISO studies of the ISM in normal galaxies focussed on observations with LWS of the [C II] 158 μm fine structure transition, which is the main channel for cooling of the diffuse ISM. Prior to ISO, knowledge of [C II] emission from external galaxies was limited to measurements of a total of 20 strongly star-forming galaxies sufficiently bright to be accessible to the Kuiper Airborne Observatory (Crawford et al., 1985; Stacey et al., 1991; Madden et al., 1993). The main conclusion of these early investigations was that the [C II] line emission arose from PDRs, where the UV light from young stars impinged on molecular clouds in star-forming regions. The only exception to this was Madden et al.’s resolved study of NGC 6946, who showed that the line could also be emitted from the diffuse disk. The principal contributions of ISO to the knowledge of the neutral gaseous ISM have been twofold: First, knowledge of the [C II] line emission has been extended to quiescent spiral galaxies and dwarf galaxies. These targets are difficult to access from beneath the atmosphere on account of their low luminosities and brightnesses. Second, ISO has tripled the number of higher luminosity systems with [C II] measurements, and has also observed a whole range of other fine structure lines in these systems. These lines have been used to probe in detail the identity of the emitting regions and their physical parameters. Fundamental to the studies of the quiescent spiral galaxies were the measurements by Leech et al. (1999), who detected 14 out of a sample of 19 Virgo cluster spirals (a subset of the IVCDS sample – see Section 2.2) in the [C II] line using LWS. These galaxies are the faintest normal galaxies ever measured in the [C II] line, and can be regarded as quiescent in star-forming activity (Hα EW is less than 10 A˚ for 8 galaxies). In a series of papers Pierini and collaborators (Pierini et al., 1999; Pierini et al., 2001, 2003a) used these LWS data, in combination with the existing KAO measurements of more actively star-forming galaxies, to investigate the physical origin of the line emission in normal galaxies as a function of star-formation activity. The main result was that the [C II] emission of quiescent systems predominantly arises from the diffuse ISM (mainly the cold neutral medium; CNM), as evidenced from the fact that these systems occupy a space in the L[C II]/LCO versus

L[C II]/LFIR diagram devoid of individual star-formation regions or giant molecular NORMAL NEARBY GALAXIES 339 clouds (Pierini et al., 1999). An origin of the [C II] emission from the CNM had earlier been proposed by Lord et al. (1996) as an alternative to an origin in localised PDRs, on the basis of a full LWS grating spectrum of the spiral galaxy NGC 5713.

The luminosity ratio of [C II] line emission to FIR dust emission, L[C II]/LFIR,was found to be in the range 0.1–0.8% for the quiescent spirals. This is consistent with the basic physical interpretation, originally established for the more active galaxies, of a balance between gas cooling (mainly through the [C II] line) and gas heating through photoelectric heating from grains.1 However, the range of a factor of almost an order of magnitude in the observed values of L[C II]/LFIR also suggests that the relation between the [C II] line emission and the SFR may be quite complex. This impression is reinforced by the large scatter and non-linearities found by Boselli et al. (2002) in the relations between the [C II] luminosity and SFRs derived from Hα measurements for a sample principally composed of the Virgo galaxies measured by Leech et al. (1999) with the LWS. One consequence of the physical association of the [C II] emission with the diffuse dusty ISM in quiescent galaxies is that the observed L[C II]/LFIR ratio will be reduced due to an optically (rather than UV-) heated component of the FIR dust emission. A model quantifying this effect was developed by Pierini et al. (2003a) in which L[C II]/LFIR depends on the fractional amount of the non-ionising UV light in the interstellar radiation field in normal galaxies. Overall, systematic variations in the L[C II]/LFIR ratio for star-forming galaxies can be summarised as follows: In progressing from low to high star-formation activities, L[C II]/LFIR first increases in systems in which the [C II] emission is mainly from the diffuse medium, due to a decrease in importance of optical photons in heating the diffuse dust. After reaching a maximum, L[C II]/LFIR then decreases as the star-formation activity is further increased to starburst levels, due to an increase in the fraction of [C II] emission arising from localised PDRs associated with star-forming regions, coupled with the quenching of the [C II] emission through the decreased efficiency of photoelectric heating of the gas in high radiation fields. Thus, the dominance of the [C II] emission from PDRs turned out to be the asymptotic limit for high SFR in gas-rich galaxies. The LWS observations of Virgo spirals also showed that the linear relation between L[C II] and LCO previously established for starburst systems extends into the domain of the quiescent spirals, though with an increased scatter (see Figure 14). α The tight relation between L[C II]/LCO and the H equivalent width (EW) found by Pierini et al. (1999) indicated that this scatter may be induced by different strengths of the far-UV radiation field in galaxies. Alternatively, as discussed by Smith and

Madden (1997) in their study of five Virgo spirals, the fluctuations in L[C II]/LCO from one galaxy to the next might reflect changes in LCO from PDRs caused by

1The exception is the Virgo cluster galaxy NGC 4522, which was found by Pierini et al. (1999) to have an abnormally high value for L[C II]/LFIR, possibly indicating mechanical heating of the interstellar gas as the galaxy interacts with the intracluster medium. 340 M. SAUVAGE ET AL.

Figure 14. The relationship between the observed central [C II] line intensity, I[C II], and the central CO line intensity, ICO from Pierini et al. (2001). Asterisks and filled triangles denote starburst galaxies and gas-rich galaxies of the KAO sample, while open circles identify spiral galaxies of the ISO sample. The long- and short-dashed lines show the average ratios of the two observables obtained by Stacey et al. (1991) for the starburst galaxies and the gas-rich galaxies, respectively. variations in metallicity (and dust abundance), as well as varying fractions of [C II] emission arising from PDRs and the cold neutral medium. The measurements of the [C II]/CO line ratios derived from ISO observations have also thrown new light on the much debated relation between the strength of the CO line emission and the mass of molecular hydrogen in galaxies. Bergvall et al. (2000) emphasise how the low metallicity, the intense radiation field and the low column density in the dwarf starburst galaxy Haro 11 can explain the extremely high observed [C II]/CO flux ratio, indicating that CO may be a poor indicator of the H2 mass in such systems. Similarly, the enhanced [C II]/CO ratios found in the Virgo spirals observed by Smith and Madden (1997) were attributed to the low metallicities in these galaxies, although, as also pointed out by these authors, an alternative explanation could be that radiation from diffuse H I may dominate the [C II] emission. Mapping observations of the [C II] line using the LWS provide a means to directly probe the relative amounts of [C II] arising from localised PDRs associated with star-forming regions in the spiral arms and the diffuse disk, as well to investigate the relation between the [C II] emission and the neutral and molecular components of the gas. They have also allowed the nuclear emission to be separately studied. Contursi et al. (2002) mapped the two nearby late-type galaxies NGC 1313 and NGC 6946 in the [C II] line, finding that the diffuse H I disk contributes ≤∼40% and ∼30% of NORMAL NEARBY GALAXIES 341 the integrated [C II] emission in NGC 6946 and NGC 1313, respectively. CO(1-0) and [C II] were also found to be well correlated in the spiral arms in NGC 6946, but less well so in NGC 1313. Stacey et al. (1999) mapped the barred spiral M83 in a variety of fine structure lines in addition to the [C II] line – the [O I] 63 and 145 μm lines, the [N II] 122 μm line and the [O III]88μm line, and obtained a full grating scan of the nucleus. At the nucleus, the line ratios indicate a strong starburst headed by O9 stars. Substantial [N II] emission from low density H II regions was found in addition to the anticipated component from PDRs. A further resolved galaxy studied with the LWS is M33, for which Higdon et al. (2003) made measurements of fine structure lines and dust continuum emission towards the nucleus and six giant H II regions. Overall, the picture presented by these investigations of resolved galaxies is broadly in accordance with the theory that the integrated [C II] emission from spirals galaxies is comprised of the sum of components from the diffuse disk and from localised PDRs and diffuse ionised gas associated with star-formation regions in the spiral arms. Fundamental to the statistical investigations of gas in more active galaxies are the full grating spectra of 60 normal galaxies from the ISO key-project sample of Helou et al. (1996), which were obtained by Malhotra et al. (1997, 2001) with the LWS. One of the main results was the discovery of a smooth but drastic decline in L[C II]/LFIR as a function of increasing star-formation activity (as traced by the L60/L100 IRAS colour ratio or LFIR/LB) for the most luminous third of the sample. This trend was accompanied by a increase in the ratio of the luminosity of the [O I] 63 μm line to that of the [C II] line. This is readily explained in PDR models in terms of a decreased efficiency of photoelectric heating of the gas as the intensity of the UV radiation field is increased, coupled with an increased importance of the [O I] cooling line as the gas temperature is increased. Both Malhotra et al. (2001) and Negishi et al. (2001) present a comprehensive analysis of the ratios of the most prominent lines ([C II] 158 μm; [O I] 63 and 145 μm; [N II] 122 μm; [O III] 52 and 88 μm and [N III]57μm) in terms of the densities of gas and radiation fields in PDRs, as well as of the fraction of emission that comes from the diffuse ionised medium (as traced by [N II] 122 μm). Malhotra et al. (2001) found that the α radiation intensities G0 increases with density n as G0 ∝ n , with α being ∼1.4. They interpret this result, together with the high PDR temperatures and pressures needed to fit the data, as being consistent with the hypothesis that most of the line and continuum luminosity of relatively active galaxies arises from the immediate proximity of the star-forming regions. By contrast, on the basis of a comparison of the [C II] line with the [N II] 122 μm line, Negishi et al. (2001) argue that a substantial amount (of order 50%) of the [C II] line emission might arise not from PDRs but instead from low density diffuse ionised gas. Negishi et al. (2001) also found a linear (rather than non-linear) relation between G0 and n, which would argue against a decreased photoelectric heating efficiency being responsible for the decline in L[C II]/LFIR with star-formation activity. Instead, Negishi et al. invoke either an increase in the collisional de-excitation of the line with increasing density, or a 342 M. SAUVAGE ET AL. decrease in the diffuse ionised component with increasing star-formation activity as being responsible for this effect.

An explanation for the decrease in L[C II]/LFIR in terms of a decreased efﬁciency of photoelectric heating of the gas (Malhotra et al., 2001) or due to a collisional de-excitation of the [C II] line (Negishi et al., 2001) would predict that the L[C II] will not be a good probe of high redshift starburst galaxies. However, an alternative

(or complementary) explanation for the decrease in L[C II]/LFIR with star-formation activity, proposed by Bergvall et al. (2000), is that the line becomes self-absorbed in local universe starbursts with high metallicity. This scenario would also explain why the metal poor luminous blue compact dwarf galaxy Haro 11 was found by

Bergvall et al. to have a high ratio of L[C II]/LFIR, despite the intense radiation ﬁelds present.

2.4.2. The Molecular Component The lowest energy transitions of the most abundant molecule in the Universe, H2, are located in the MIR, typically between 3 and 30 μm. These pure rotational lines potentially offer a new access to the molecular component of the ISM, to be compared with the less direct CO tracer. One caveat though is that the lines can be excited by different mechanisms, typically shocks or UV-pumping, which renders their interpretation difﬁcult. Furthermore, the emission is very rapidly dominated by that from the warmest fraction of the gas, and therefore this new tracer offers only a partial and complex access to the molecular phase. Most of the extragalactic H2 detections were made with ISOSWS, since high spectral resolution is necessary to isolate the very thin molecular lines from the broader features that abound in the MIR regime (although a tentative detection in NGC 7714 made with ISOCAM is presented by O’Halloran et al., 2000). For sensitivity reasons, the galaxies that were observed in these lines tended to be starburst or active galaxies and are discussed elsewhere in this volume. Nevertheless, the S(0) and S(1) transitions were observed in the central region and in the disk of NGC 6946 (Valentijn et al., 1996; Valentijn and van der Werf, 1999a) and all along the disk of NGC 891 (Valentijn and van der Werf, 1999b). In the central region of NGC 6946, the molecular component detected with ISOSWS is clearly the warm fraction (of order 5–10%) of the molecular gas present in the region and is raised to the observed temperature of 170 K by the enhanced star-formation activity in the center of the galaxy. H2 is also detected further out (4 kpc) in the disk, with a cooler temperature, though still relatively warm. The observations of H2 in the disk of NGC 891 reach much larger distances, extending about 10 kpc on both sides of the nucleus. Valentijn and van der Werf (1999b) argue that the observed line ratios indicate the presence of both warm (150– 230 K) molecular gas, identical as that observed in NGC 6946, and cool (80–90 K) molecular gas. This cooler gas needs to have a very large column density to explain the observed surface brightness and the masses derived by Valentijn and van der NORMAL NEARBY GALAXIES 343

Werf (1999b) are in fact so large that they would resolve the missing mass problem in the optical disk of NGC 891. This highly controversial conclusion however rests on a rather uncertain basis dealing with the actual spatial distribution of the H2- emitting clouds within the ISOSWS beam. Upcoming observations with the IRS on Spitzer should remove the remaining uncertainties regarding the presence of large amounts of cold H2 in the disks of galaxies.

3. Dwarf Galaxies

3.1. COLD DUST SURROUNDING DWARF GALAXIES

Gas-rich dwarf galaxies and in particular Blue Compact Dwarfs (BCDs) were originally expected to have their FIR emission dominated by dust heated locally in H II regions. Temperatures of 30 K or more were anticipated. This was the a priori expectation in particular for the BCDs, and became the standard interpretation for the IRAS results obtained for these systems. Hoffman et al. (1989), Helou et al. (1988) and Melisse and Israel (1994) each found that the 60/100 μm colours of BCDs were clearly warmer than those of spirals. The IVCD Survey (Tuffs et al., 2002a,b) changed that simple picture of the FIR emission from dwarf galaxies. The IVCDS included measurements at 60, 100, and (for the first time) at 170 μm of 25 optically selected gas-rich dwarf galaxies. The observations at 60 and 100 μm were consistent with the previous IRAS results on this class of galaxies, though extending knowledge of these systems to lower intrinsic luminosities. Unexpectedly however, high ratios of 170/100 μm luminosities were found in many of the surveyed systems. Such long-wavelength excesses were found both in relatively high-luminosity dwarfs, such as VCC655, as well as in fainter objects near the limiting sensitivity of the survey. These observations imply the presence of large amounts of cold dust. As shown by Popescu et al. (2002), it seems unlikely that the cold dust resides in the optically thick molecular component associated with star-formation regions, since the implied dust masses would be up to an order of magnitude greater than those typically found in giant spirals. Such large masses could not have been produced through star formation within the dwarfs over their lifetime. One alternative possibility is that the dust originates and still resides outside the optical extent of these galaxies. In fact, some evidence for this was provided by the ISOPHOT observations themselves, since they were made in the form of scan maps, from which estimates of source sizes could be determined. Even with the relatively coarse beam (1.6 FWHM at 170 μm), the extended nature of the sources could be clearly seen in a few cases for which the FWHM for the 170 μm emission exceeds the optical diameters of the galaxies (to 25.5 Bmag/arcsec2)byfac- tors of between 1.5 and 3.5. It is interesting to note that two of these galaxies (VCC 848 and VCC 81) have also been mapped in H I (Hoffman et al., 1996), 344 M. SAUVAGE ET AL. revealing neutral hydrogen sizes comparable to the 170 μm extent. This raises the possibility that the cold dust is embedded in the extended H I gas, external to the optical galaxy. This would be analogous to the case of the edge-on spiral NGC 891, where Popescu and Tuffs (2003) discovered a cold-dust counterpart to the extended H I disk (see Section 2.1.3). In this context the main observational difference between the giant spiral and the dwarfs may be that for the dwarfs the integrated 170 μm emission is dominated by the extended emission component external to the main optical body of the galaxy, whereas for the giant spirals the long-wavelength emission predominantly arises from within the confines of the optical disk of the galaxy. Apart from the SMC (which is discussed in Section 3.2 and is too extended for ISOPHOT to map beyond its optical extent), only three dwarf galaxies were observed by ISOPHOT in the field environment (all serendipity survey sources; Stickel et al., 2000). All three sources have comparable flux densities at 100 and 170 μm. But the small statistics mean that it is still an open question to which extent the cold dust emission associated with the extended H I component in dwarf galaxies is a cluster phenomenon or not. The existence of large quantities of dust surrounding gas-rich dwarf galaxies may have important implications for our understanding of the distant Universe. According to the hierarchical galaxy-formation scenarios, gas-rich dwarf galaxies should prevail at the earliest epochs. We would then expect these same galaxies to make a higher contribution to the total FIR output in the early Universe, certainly more than previously expected.

3.2. INFRARED EMISSION FROM WITHIN DWARF GALAXIES

The distribution of cold dust within dwarf galaxies could be studied in only one case, namely in the resolved (1.5 resolution) 170 μm ISOPHOT map of the Small Magellanic Cloud (Wilke et al., 2003, 2004). The 170 μm ISO map of the SMC reveals a wealth of structure, not only consisting of ﬁlamentary FIR emitting regions, but also of numerous (243 in total) bright sources which trace the bar along its major axis as well as the bridge which connects the SMC to the LMC. Most of the brighter sources have cold components, associated with molecular clouds. The discrete sources were found to contribute 28, 29 and 36% to the integrated ﬂux densities at 60, 100 and 170 μm, respectively. The SED was modelled by the superposition of 45, 20.5 and 10 K blackbody components with emissivity index β = 2. The average dust colour temperature (averaged over all pixels of the 170/100 colour map) was found to be TD = 20.3K. Bot et al. (2004) have compared the FIR map of the SMC with an H I map of similar resolution. This reveals a good spatial correlation of the two emissions in the diffuse regions of the maps (regions that fall outside of the correlation are either hot star-forming regions, or cold molecular clouds with no associated H I). Adding the IRAS data allows them to compute the FIR emissivity per unit H atom. Bot et al. NORMAL NEARBY GALAXIES 345

(2004) found that this emissivity is lower than in the Milky Way, and in fact it is even lower than the lower metallicity (Z/10) of the SMC would imply, suggesting that depletion mechanisms at work in the ISM have more than a linear dependence on metallicity. Although FIR emission from dwarf galaxies has been associated only with gas- rich dwarfs, in one particular case such emission has been detected in a dwarf elliptical galaxy, as well. This is the case of NGC 205, one of the companions of M 31, classified as a peculiar dE5. This galaxy shows signatures of recent star formation (Hodge, 1973) and of extended H I emission (Young and Lo, 1997), and was detected by IRAS (Rice et al., 1988; Knapp et al., 1989), with a SED steeply rising between 60 and 100 μm. Based on ISOPHOT observations, Haas (1998) showed that the FIR emission is resolved and similar to that seen in H I. He also presented evidence for a very cold dust component, of 10 K, coming from the center of the galaxy. The properties of gas-rich dwarf galaxies were found to markedly differ from those of spiral galaxies at shorter IR wavelengths. Thus, when comparing the MIR emission of five IBm galaxies to their Hα or FIR properties, Hunter et al. (2001) discovered a noticeable deficit of the 6.75 μm component in their galaxies. This deficit is not observed at 15 μm, i.e. the 6.75/15 μm flux ratio is very different in IBm galaxies than in spirals, which is interpreted as tracing both a deficit of PAH, possibly destroyed by the radiation field, and an increase of the mean temperature of the very small grains, again linked to the radiation field (see also O’Halloran et al., 2000). This interpretation of the observations is confirmed by detailed spectroscopic measurements of similar objects. In NGC 5253 the MIR SED is dominated by a “hot” continuum from which the PAH features are absent (Crowther et al., 1999). The blue compact galaxy II Zw 40 presents a similar spectrum, i.e. a strongly rising continuum with no bands (Madden, 2000). In both cases this is interpreted as being the result of very hard radiation fields that both destroy the PAH and shift the continuum emission from very small grains toward shorter wavelengths. These galaxies are essentially small starburst regions and their IR spectrum thus reflects the IR properties of starburst galaxies discussed elsewhere in this volume. That we are indeed seeing a gradual destruction of PAHs, and not an effect of metallicity that would prevent the formation of these species is demonstrated by the fact that we find galaxies of similar metallicity (NGC 5253 and NGC 1569 for instance at Z/4 − Z/5) without or with PAH emission (see Galliano et al., 2003 for NGC 1569). In fact, even the SMC at Z/10 has a MIR spectrum with pronounced PAH features (Reach et al., 2000). Many of the qualitative indications about grain and PAH abundance indicated above have been quantitatively estimated by Galliano et al. (2003) in their detailed modelling of the UV-optical/MIR-FIR-submm SED of the low metallicity nearby dwarf galaxy NGC1569 (Figure 15). This study is noteworthy in that it constrains the grain size distribution through the MIR-FIR ISOCAM and ISOPHOT observations and therefore gives more specific information about grain properties in dwarf 346 M. SAUVAGE ET AL.

Figure 15. NGC 1569 observations and modeled SED from Galliano et al. (2003). The data are indicated by crosses: vertical bars are the errors on the flux density values and the horizontal bars indicate the widths of the broadbands. The lines show the predictions of the dust model with its different components. Diamonds indicate the model predictions integrated over the observational broadbands and colour-corrected. galaxies. The results confirm the paucity of PAHs due to an enhanced destruction in the intense ambient UV radiation field, as well as an overabundance (compared to Milky Way type dust) of small grains of size ∼3 nm, possibly indicative of a redistribution of grain sizes through the effect of shocks.

4. Early-Type Galaxies

Here we use the term “early-type” to embrace both elliptical and lenticular galaxies. Their IR emission, when detected, has been interpreted very differently from the IR emission from spiral galaxies. Indeed, early-type galaxies belong to a subset of the Hubble sequence where star-formation activity is at its minimum, or has stopped, and the ISM is almost exhausted. Given what we have seen above, one does not expect these objects to be IR bright. This is indeed the case, judging from the small fraction of early-type objects in the IRAS catalogues (see e.g. Jura et al., 1987), and the large uncertainties associated with their ﬂuxes (Bregman et al., 1998). For those few galaxies where IR emission is detected, it is commonly attributed either to the Rayleigh–Jeans tail of the stellar emission, or to dust in mass-losing stars. These interpretations however rest on the four broadband low spatial resolution data of IRAS. ISO, with its much improved sensitivity to both spatial and spectral details, has shown a sharper light on the subject. NORMAL NEARBY GALAXIES 347

The FIR observations of early-type galaxies selected from the ISO archive have been analysed by Temi et al. (2004). They found that the FIR SED requires emission from dust with at least two different temperatures, with mean temperatures for the warm and cold components of 43 and 20 K, respectively. This result is quite similar to the results obtained for quiescent spiral galaxies, in particular to the result obtained by Popescu et al. (2002) for the Virgo cluster spiral galaxies. Since ellipticals are known to be ISM-poor and have reduced star-formation activity, the similarity between their FIR properties and those of spirals is surprising. One explanation could be the fact that the galaxies included in the ISO archive are not representative of typical ellipticals. As pointed out by Temi et al., many galaxies selected for ISO observations were chosen because they were known to have large IRAS fluxes or large masses of cold gas, and in this sense they have properties closer to those of spirals. For the same reason Temi et al. considered that these IR- selected galaxies are likely to have experienced unusual dust-rich mergers, and that the dust in these galaxies could have an external origin. The lack of a correlation found between the FIR luminosity (or dust mass) and the luminosity of the B-band emission was taken as evidence for this scenario. Another scenario proposed is one in which most of the FIR is emitted by central dust clouds, where the clouds could be disturbed at irregular intervals by low-level AGN activity in the galactic cores, creating the stochastic variation in the FIR luminosity. It is difficult to distinguish between these two scenarios, also because the sensitivity of ISO was not sifficient to detect FIR emission from many optically luminous elliptical galaxies. In order to distinguish between an external and an internal origin for the dust, deeper FIR observations of an optically selected sample of elliptical galaxies are needed. It is even possible that such surveys may occasionally detect bright FIR emission from galaxies traditionally classified as ellipticals, but which in fact have hidden starburst activity in their central regions. This is the case for the IRAS source F15080 + 7259, which was detected in the ISOPHOT serendipity survey, and was shown by Krause et al. (2003) to harbour large quantities of gas and cold dust. Another elliptical detected by ISOPHOT is the Virgo cluster galaxy M 86 (NGC 4406). The deep FIR imaging data obtained with ISOPHOT at 60, 90, 150 and 180 μm revealed a complex FIR morphology for this galaxy (Stickel et al., 2003). It was found that the FIR emission originates from both the centre of M 86 and from the optically discovered dust streamers, and has a dust temperature of 18 K. Again, this elliptical is not representative for this class of objects, since it was considered either to be experiencing ram-pressure stripping or to be gravitationally interacting with its neighbouring galaxies. In particular the ISOPHOT observations were interpreted by Stickel et al. as being consistent with M 86 having a gravitational interaction with its nearest spiral galaxy NGC 4402. In this case the detected dust would again have (at least in part) an external origin. The importance of gravitational interaction as a driver for FIR emission in early-type galaxies was demonstrated by Domingue et al. (2003), who showed these systems to have a 50% detection rate with ISOPHOT when they are paired with spirals. 348 M. SAUVAGE ET AL.

In the MIR ISOCAM was used to observe both elliptical and lenticular galaxies, though the objects selected differ from those observed by ISOPHOT. An unbiased sample of 16 IRAS-detected lenticular galaxies was observed using low resolution MIR spectroscopy by Sauvage et al. (2006). The vast majority of the detected galaxies present MIR SEDs extremely similar to those of later-type spiral galaxies (i.e. PAH bands). Thus, from their MIR spectral properties, lenticular galaxies appear to belong to the spiral galaxy group, although with smaller luminosities. Elliptical galaxies have drawn much more interest from ISOCAM observers, even though they are ISM poorer than the lenticulars. The main result from these studies is that elliptical galaxies are really a mixed bag, as far as their MIR properties are concerned. Quillen et al. (1999) mapped with ISOCAM a number of E+A galaxies2 in the Coma cluster, but detected only those which showed emission lines characteristic of an on-going star-formation episode. This might be thought to imply that star formation is required for ellipticals to emit in the MIR, but in fact more likely reflects the sensitivity reached by Quillen et al.. With deep ISOCAM observations of closer “normal” elliptical galaxies, Athey et al. (2002) showed that their MIR emission is dominated by the emission from K and M stars, with the addition of a spectral feature near 10 μm that they attribute to silicates in the circumstellar envelopes of AGB stars. Xilouris et al. (2004) brought balance to the star/ISM debate for the origin of MIR emission in early-type galaxies by showing that the complete range of phenomena occur in this class of galaxies. Systems vary from those that are dominated up to ∼20 μm by the emission from their old stellar populations to those in which a small star-forming episode, possibly fueled by recent accretion, dominates the MIR luminosities. In between one also finds galaxies where only the long-wavelength part of the MIR spectrum deviates from the extrapolation of the stellar light, which indicates the presence of small amounts of diffuse dust much hotter than what is commonly found in spirals. This is most likely to be due to the higher interstellar radiation field generated by the denser stellar populations. Although ISO has had some success in detecting and measuring the IR SED from early-type galaxies, a critical look at the observations done reveals that we still do not know the IR properties of typical elliptical galaxies. To achieve this, systematic deep observations of optically selected samples are required.

5. Conclusions and Outlook

ISO has not only advanced the knowledge of IR properties of normal galaxies but has also made unexpected discoveries. The bulk of the emission from dust has been measured, revealing cold dust in copious quantities. This dust is present in all types of normal galaxies and is predominantly distributed in a diffuse disk with an intrinsic scalelength exceeding that of the stars. Cold dust has been found beyond the

2Early-type galaxies showing evidence of a post-starburst population. NORMAL NEARBY GALAXIES 349 optical regions of isolated galaxies, associated with the extended H I disks of spiral galaxies or with the H I envelopes of dwarf galaxies. The fraction of the bolometric luminosity radiated by dust has been measured for the ﬁrst time. Realistic geometries for stars and dust have been derived from ISO imaging observations, enabling the contribution of the various stellar populations to the dust heating to be accurately derived. The NIR/MIR SEDs were shown to be remarkably similar for normal spiral galaxies and to originate mostly in the PDRs surrounding star-forming regions. A diffuse component of the MIR emission was found, though puzzlingly with a much smaller scalelength than its FIR counterpart. The importance of the central regions in shaping up the intensity and the colour of the MIR global emission was revealed thanks to the high spatial resolution offered by ISO. A new component of interstellar dust emission consisting of a “hot” NIR continuum emission was discovered. The relative contribution of photospheric and very small grains/ PAH emission has been established in the NIR/MIR spectral region. Spectroscopic observations have revealed that the main interstellar cooling line, [C II], is predominantly carried by the diffuse cold neutral medium in normal galaxies, with emission from localised PDRs only dominating for galaxies having high star-formation activity. This enormous advancement in the understanding of normal galaxies in the nearby Universe has laid the foundation for more detailed investigations with Spitzer and Herschel. A clear priority is to ﬁll the gap left by ISO in knowledge of the SEDs of normal galaxies between 20 and 60 μm. Another priority is to increase the number of galaxies with detailed imaging information and to provide better statistics on carefully selected samples, especially those selected in the optical/NIR bands. Ultimately, the improved sensitivity of the new infrared space observatories will allow knowledge of the dust emission from normal galaxies to be extended beyond the nearby universe.

Acknowledgements

The authors would like to take this opportunity to thank all the individuals that helped make the ISO mission a success. M. Sauvage acknowledges the Max Planck Institut fürKernphysik for its support during the final adjustments of the manuscript. R.J. Tuffs and C.C. Popescu would also like to thank Heinrich J. Völkfor enlight- ening discussions.

References

Alton, P. B., Bianchi, S., Rand, R. J., Xilouris, E. M., Davies, J. I., et al.: 1998b, Astrophys. J. 507, L125. Alton, P. B., Trewhella, M., Davies, J. I., Evans, R., Bianchi, S., et al.: 1998a, Astron. Astrophys. 335, 807. Athey, A., Bregman, J., Bregman, J., Temi, P., and Sauvage, M.: 2002, Astrophys. J. 571, 272. 350 M. SAUVAGE ET AL.

Bendo, G. J., Joseph, R. D., Wells, M., Gallais, P., Haas, M., et al.: 2002a, Astron. J. 123, 3067. Bendo, G. J., Joseph, R. D., Wells, M., Gallais, P., Haas, M., et al.: 2002b, Astron. J. 124, 1380. Bendo, G. J., Joseph, R. D., Wells, M., Gallais, P., Haas, M., et al.: 2003, Astron. J. 125, 2361. Bergvall, N., Masegosa, J., Ostlin,¨ G., and Cericharo, J.: 2000, Astron. Astrophys. 359, 41. Bianchi, S., Davies, J. I., and Alton, P. B.: 2000a, Astron. Astrophys. 359, 65. Bianchi, S., Davies, J. I., Alton, P. B., Gerin, M., and Casoli, F.: 2000b, Astron. Astrophys. 353, L13. Binggeli, B., Popescu, C. C., and Tammann, G. A.: 1993, Astron. Astrophys. Suppl. 98, 275. Binggeli, B., Sandage, A., and Tammann, G. A.: 1985, Astron. J. 90, 1681. Bot, C., Boulanger, F., Lagache, G., Cambr´esy, L., and Egret, D.: 2004, Astron. Astrophys. 423, 567. Boselli, A., Gavazzi, G., Lequeux, J., and Pierini, D.: 2002, Astron. Astrophys. 385, 454. Boselli, A., Gavazzi, G., and Sanvito, G.: 2003a, Astron. Astrophys. 402, 37. Boselli, A., Lequeux, J., Contursi, A., Gavazzi, G., Boulade, O., et al.: 1997, Astron. Astrophys. 324, L13. Boselli, A., Sauvage, M., Lequeux, J., Donati, A., and Gavazzi, G.: 2003b, Astron. Astrophys. 406, 867. Bregman, J. N., Snider, B. A., Grego, L., and Cox, C. V.: 1998, Astrophys. J. 499, 670. Cesarsky, D., Lequeux, J., Pagani, L., Ryter, C., Loinard, L., and Sauvage, M.: 1998, Astron. Astrophys. 337, L35. Cesarsky, C. J., et al.: 1996, A&A 315, L32. Charlot, S. and Fall, S. M.: 2000, Astrophys. J. 539, 718. Contoursi, A., Kaufman, M. J., Helou, G., et al.: 2002, Astron. J. 124, 751. Contursi, A., Boselli, A., Gavazzi, G., Bertagna, E., Tuffs, R., and Lequeux, J.: 2001, Astron. Astrophys. 365, 11. Contursi, A., Lequeux, J., Hanus, M., Heydari-Malayeri, M., Bonoli, C., Bosma, A., et al.: 1998, Astron. Astrophys. 336, 662. Contursi, A., Lequeux, J., Cesarsky, D., Boulanger, F., Rubio, M., Hanus, M., et al.: 2000, Astron. Astrophys. 362, 310. Crawford, M. K., Genzel, R., Townes, C. H., and Watson, D. M.: 1985, Astrophys. J. 291, 755. Crowther, P. A., Beck, S. C., Willis, A. J., Conti, P. S., Morris, P. W., and Sutherland, R. S.: 1999, Mon. Not. R. Astron. Soc. 304, 654. Dale, D., and Helou, G.: 2002, Astrophys. J. 576, 159. Dale, D. A., Helou, G., Contursi, A., Silbermann, N. A., and Kolhatkar, S.: 2001, Astrophys. J. 549, 215. Dale, D. A., Helou, G., Silbermann, N. A., Contursi, A., Malhotra, S., and Rubin, R. H.: 1999, Astron. J. 118, 2055. Dale, D. A., Silbermann, N. A., Helou, G., Valjavec, E., Malhotra, S., Beichmann, C. A., et al.: 2000, Astron. J. 120, 583. Davies, J. I., Alton, P., Bianchi, S., and Trewhella, M.: 1998, Mon. Not. R. Astron. Soc. 300, 1006. Davies, J. I., Alton, P., Trewhella, M., Evans, R., and Bianchi, S.: 1999, Mon. Not. R. Astron. Soc. 304, 495. de Jong, T., Clegg, P. E., Soifer, B. T., et al.: 1984, Astrophys. J. 278, L67. de Jong, T., Klein, U., Wielebinski, R., and Wunderlich, E.: 1985, Astron. Astrophys. 147, L6. Domingue, D. L., Sulentic, J. W., Xu, C., Mazzarella, J., Gao, Y., and Rampazzo, R.: 2003, Astron. J. 125, 555. Dopita, M. A., Groves, B. A., Fischera, J., Sutherland, R. S., Tuffs, R. J., et al.: 2005, Astrophys. J. 619, 755. Dupac, X., del Burgo, C., Bernard, J.-P., et al.: 2003, Mon. Not. R. Astron. Soc. 344, 105. Ferguson, A., Gallagher, J. S., and Wyse, R.: 1998, Astron. J. 116, 673. Ferrara, A., Ferrini, F., Barsella, B., and Franco, J.: 1991, Astrophys. J. 381, 137. NORMAL NEARBY GALAXIES 351

Förster-Schreiber, N., Roussel, H., Sauvage, M., and Charmandaris, V.: 2004, Astron. Astrophys. 419, 501. Förster-Schreiber, N., Sauvage, M., Charmandaris, V., Laurent, O., Gallais, P., Mirabel, I. F., et al.: 2003, Astron. Astrophys. 399, 833. Galliano, F., Madden, S. C., Jones, A. P., Wilson, C. D., Bernard, J.-P., et al.: 2003, Astron. Astrophys. 407, 159. Gerhard, O. and Silk, J.: 1996, Astrophys. J. 472, 34. Guélin,M., Zylka, R., Mezger, P. G., Haslam, C. G. T., Kreysa, E., et al.: 1993, Astron. Astrophys. 279, L37. Haas, M.: 1998, Astron. Astrophys. 337, L1. Haas, M., Klaas, U., and Bianchi, S.: 2002, Astron. Astrophys. 385, L23. Haas, M., Lemke, D., Stickel, M., Hippelein, H., Kunkel, M., et al.: 1998, Astron. Astrophys. 338, L33. Haas, M., Siebenmorgen, R., Leipski, C., Ott, S., Cunow, B., Meusinger, H., et al.: 2004, Astron. Astrophys. 419, L49. Helou, G.: 1986, Astrophys. J. 311, 33. Helou, G., Khan, I. R., Malek, L., and Boehmer, L.: 1988, Astrophys. J. Suppl. 68, 151. Helou, G., Malhotra, S., Beichman, C. A., Dinerstein, H., Hollenbach, D. J., Hunter, D. A., et al.: 1996, Astron. Astrophys. 315, L157. Helou, G., Malhotra, S., Hollenbach, D. J., Dale, D. A., and Contursi, A.: 2001, Astrophys. J. 548, L73. Helou, G., Roussel, H., Appleton, P., Frayer, D., Stolovy, S., et al.: 2004, Astrophys. J. Suppl., 154, 253. Helou, G., Soifer, B. T., and Rowan-Robinson, M. R.: 1985, Astrophys. J. 298, L7. Higdon, S. J. U., Higdon, J. L., van der Hulst, J. M., and Stacey, G. J.: 2003, Astrophys. J. 592, 161. Hippelein, H., Haas, M., Tuffs, R. J., Lemke, D., Stickel, M., et al.: 2003, Astron. Astrophys. 407, 137. Hoffman, G. L., Helou, G., Salpeter, E. E., and Lewis, B. M.: 1989, Astrophys. J. 339, 812. Hodge, P.: 1973, Astrophys. J. 182, 671. Hoffman, G. L., Salpeter, E. E., Farhat, B., Roos, T., Williams, H., and Helou, G.: 1996, Astrophys. J. Suppl. 105, 269. Hunter, D. A., Kaufman, M., Hollenbach, D. J., Rubin, R. H., Malhotra, S., et al.: 2001, Astrophys. J. 553, 121. Jura, M., Kim, D. W., Knapp, G. R., and Guhathakurta, P.: 1987, Astrophys. J. 312. Kessler, M. F., et al.: 1996, A&A 315, L27. Knapp, G. R., Guhathakurta, P., Kin, D. W., and Jura, M.: 1989, Astrophys. J. Suppl. 70, 329. Kraemer, K. E., Price, S. D., Mizuno, D. R., and Carey, S. J.: 2002, Astron. J. 124, 2990. Krause, O., Lisenfeld, U., Lemke, D., et al.: 2003, Astron. Astrophys. 402, L1. Krügel,E., Siebenmorgen, R., Zota, V., and Chini, R.: 1998, Astron. Astrophys. 331, L9. Leech, K. J., Völk,H. J., Heinrichsen, I., Hippelein, H., Metcalfe, L., et al.: 1999, Mon. Not. R. Astron. Soc. 310, 317. Lemke, D., et al.: 1996, A&A 315, L64. Li, A. and Draine, B. T.: 2002, Astrophys. J. 572, 232. Lonsdale-Persson, C. J. and Helou, G.: 1987, Astrophys. J. 314, 513. Lord, S. D., Malhotra, S., Lim, T., et al.: 1996, Astron. Astrophys. 315, L117. Lu, N., Helou, G., Werner, M. W., Dinerstein, H. L., Dale, D. A., Silbermann, N. A., et al.: 2003, Astrophys. J. 588, 199. Madden, S. C.: 2000, New Astron. Rev. 44, 249. Madden, S. C., Geis, N., Genzel, R., Herrmann, F., Jackson, J., et al.: 1993, Astrophys. J. 407, 579. 352 M. SAUVAGE ET AL.

Malhotra, S., Helou, G., Hollenbach, S., Kaufman, M., Lord, S., et al.: 1999, in: Cox, P. and Kessler, M. F. (eds.), Proceedings of the Conference “The Universe as Seem by ISO”, ESA-SP 427, p. 937. Malhotra, S., Helou, G., Stacey, G., Hollenbach, S., Lord, S., et al.: 1997, Astrophys. J. 491, L27. Malhotra, S., Helou, G., van Buren, D., Kong, M., Beichman, C. H., Dinerstein, H., et al.: 1996, Astron. Astrophys. 315, L161. Malhotra, S., Kaufman, M. J., Hollenbach, D., Helou, G., Rubin, R. H., et al.: 2001, Astrophys. J. 561, 766. Martin, C., et al.: 2005, Astrophys. J. Lett. 619, L1. Melisse, J. P. M. and Israel, F. P.: 1994, Astron. Astrophys. 285, 51. Misiriotis, A., Popescu, C. C., Tuffs, R. J., and Kylafis, N. D.: 2001, Astron. Astrophys. 372, 775. Negishi, T., Onaka, T., Chan, K.-W., and Roellig, T. L.: 2001, Astron. Astrophys. 375, 566. O’Halloran, B., Metcalfe, L., Delaney, M., McBreen, B., Laureijs, R., Leech, K., et al.: 2000, Astron. Astrophys. 360, 871. Ott, S., Siebenmorgen, R., Chartel, N., and Vo, T.: 2003, in: Gry, C., Peschke, S., Matagne, J., Garcia- Lario, P., Lorente, R., Salama, A. (eds.), Proceedings of the Conference “Exploiting the ISO Data Archive. Infrared Astronomy in the Internet Age”, ESA SP-511, p. 159. Pagani, L., Lequeux, J., Cesarsky, D., Donas, J., Milliard, B., Loinard, L., and Sauvage, M.: 1999, Astron. Astrophys. 351, 447. Pierini, D., Leech, K. J., Tuffs, R. J., and Völk,H. J.: 1999, Mon. Not. R. Astron. Soc. 303, L29. Pierini, D., Leech, K. J., and Völk,H. J.: 2003a, Astron. Astrophys. 397, 871. Pierini, D., Lequeux, J., Boselli, A., Leech, K. J., and Völk,H. J.: 2001, Astron. Astrophys. 373, 827. Pierini, D. and Möller, C.: 2003, Mon. Not. R. Astron. Soc. 346, 818. Pierini, D., Popescu, C. C., Tuffs, R. J., and Völk,H. J.: 2003b, Astron. Astrophys. 409, 907. Pfenniger, D. and Combes, F.: 1994, Astron. Astrophys. 285, 94. Pfenniger, D., Combes, F., and Martinet, L.: 1994, Astron. Astrophys. 285, 79. Popescu, C. C., Misiriotis, A., Kylafis, N. D., Tuffs, R. J., and Fischera, J.: 2000b, Astron. Astrophys. 362, 138. Popescu, C. C. and Tuffs, R. J.: 2002a, Mon. Not. R. Astron. Soc. 335, L41. Popescu, C. C. and Tuffs, R. J.: 2002b, Rev. Mod. Astron. 15. Popescu, C. C. and Tuffs, R. J.: 2003, Astron. Astrophys. 410, L21. Popescu, C. C., Tuffs, R. J., Fischera, J., and Völk,H. J.: 2000a, Astron. Astrophys. 354, 480. Popescu, C. C., Tuffs, R. J., Kylafis, N. D., and Madore, B. F.: 2004, Astron. Astrophys. 414, 45. Popescu, C. C., Tuffs, R. J., Madore, B. F., Gil de Paz, A., Völk,H., et al.: 2005, Astrophys. J. Lett. 619, L75. Popescu, C. C., Tuffs, R. J., Völk,H. J., Pierini, D., and Madore, B. F.: 2002, Astrophys. J. 567, 221. Quillen, A. C., Rieke, G. H., Rieke, M. J., Caldwell, N., and Engelbracht, C.: 1999, Astrophys. J. 518, 632. Reach, W. T., Boulanger, F., Contursi, A., and Lequeux, J.: 2000, Astron. Astrophys. 361, 895. Rice, W., Lonsdale, C. J., and Soifer, B. T., et al.: 1988, Astrophys. J. Suppl. 68, 91. Roussel, H., Sauvage, M., Vigroux, L., and Bosma, A.: 2001a, Astron. Astrophys. 372, 427. Roussel, H., Sauvage, M., Vigroux, L., Bosma, A., Bonoli, C., Gallais, P., et al.: 2001b, Astron. Astrophys. 372, 406. Roussel, H., Vigroux, L., Bosma, A., Sauvage, M., Bonoli, C., Gallais, P., et al.: 2001c, Astron. Astrophys. 369, 473. Sauvage, M., Blommaert, J., Boulanger, F., Cesarsky, C. J., Cesarsky, D. A., Desert, F. X., et al.: 1996, Astron. Astrophys. 315, L89. Sauvage, M., et al.: 2006, in preparation. Schmidtobreick, L., Haas, M., and Lemke, D.: 2000, Astron. Astrophys. 363, 917. NORMAL NEARBY GALAXIES 353

Siebenmorgen, R., Krügel,E., and Chini, R.: 1999, Astron. Astrophys. 351, 495. Silva, L., Granato, G. L., Bressan, A., and Danese, L.: 1998, Astrophys. J. 509, 103. Smith, B. J.: 1998, Astrophys. J. 500, 181. Smith, B. J. and Madden, S. C.: 1997, Astron. J. 114, 138. Soifer, B. T. and Neugebauer, G.: 1991, Astron. J. 101, 354. Stacey, G. J., Geis, N., Genzel, R., Lugten, J. B., Poglitsch, A., et al.: 1991, Astrophys. J. 373, 423. Stacey, G. J., Swain, M. R., Bradford, C. M., et al.: 1999, in: Proceedings of the Conference “The Universe as seen by ISO”, 1998, ESA SP-427, p. 973. Stickel, M., Bregman, J. N., Fabian, A. C., White, D. A., and Elmegreen, D. M.: 2003, Astron. Astrophys. 397, 503. Stickel, M., Lemke, D., Klaas, U., Beichman, C. A., Rowan-Robinson, M., et al.: 2000, Astron. Astrophys. 359, 865. Stickel, M., Lemke, D., Klaas, U., Krause, O., and Egner, S.: 2004, Astron. Astrophys. 422, 39. Swaters, R. A., Sancisi, R., and van der Hulst, J. M.: 1997, Astrophys. J. 491, 140. Temi, P., Brighenti, F., Mathews, W. G., and Bregman, J.: 2004, Astrophys. J. Suppl. 151, 237. Trewhella, M., Davies, J. I., Alton, P. B., Bianchi, S., and Madore, B. F.: 2000, Astrophys. J. 543, 153. Tuffs, R. J., Lemke, D., Xu, C., et al.: 1996, Astron. Astrophys. 315, L149. Tuffs, R. J. and Gabriel, C.: 2003, Astron. Astrophys. 410, 1075. Tuffs, R. J. and Popescu, C. C.: 2003, in: Gry, C., et al. (eds.), “Exploiting the ISO Data Archive. Infrared Astronomy in the Internet Age”, Siguenza, Spain 24–27 June, 2002, ESA SP-511, p. 239. Tuffs, R. J., Popescu, C. C., Pierini, D., Völk,H. J., Hippelein, H., et al.: 2002a, Astrophys. J. Suppl. 139, 37. Tuffs, R. J., Popescu, C. C., Pierini, D., Völk,H. J., Hippelein, H., et al.: 2002b, Astrophys. J. Suppl. 140, 609. Tuffs, R. J., Popescu, C. C., Völk,H. J., Kylafis, N. D., and Dopita, M. A.: 2004, Astron. Astrophys. 419, 821. Uchida, K. I., Sellgren, K., Werner, M. W., and Houdashelt, M. L.: 2000, Astrophys. J. 530, 817. Valentijn, E. A. and van der Werf, P. P.: 1999a, in: Cox, P., and Kessler, M. F. (eds.), Proceedings of the Conference “The Universe as seen by ISO”, ESA-SP 427, p. 821. Valentijn, E. A. and van der Werf, P. P.: 1999b, Astrophys. J. 522, L29. Valentijn, E. A., van der Werf, P. P., de Graauw, Th., and de Jong, T.: 1996, Astron. Astrophys. 315, L145. van Dishoeck, E. F.: 2004, Annu. Rev. Astron. Astrophys. 42, 119. Völk, H. J. and Xu, C.: 1994, Infrared Phys. Technol. 35, 527. Vogler, A., Madden, S. C., Beck, R., Sauvage, M., Vigroux, L., and Ehle, M.: 2005, 441, 491. Walsh, W., Beck, R., Thuma, G., et al.: 2002, Astron. Astrophys. 388,7. Wilke, K., Klaas, U., Lemke, D., Mattila, K., Stickel, M., et al.: 2004, Astron. Astrophys. 414, 69. Wilke, K., Stickel, M., Haas, M., Herbstmeier, U., Klaas, U., et al.: 2003, Astron. Astrophys. 401, 873. Wunderlich, E. and Klein, U.: 1988, Astron. Astrophys. 206, 47. Wunderlich, E., Wielebinski, R., and Klein, U.: 1987, Astron. Astrophys. Suppl. 69, 487. Xilouris, E. M., Madden, S. C., Galliano, F., Vigroux, L., and Sauvage, M.: 2004, Astron. Astrophys. 416, 41. Xu, C. and Buat, V.: 1995, Astron. Astrophys. 293, L65. Young, L. M. and Lo, K. Y.: 1997, Astrophys. J. 476, 131. OBSCURED ACTIVITY: AGN, QUASARS, STARBURSTS AND ULIGs OBSERVED BY THE INFRARED SPACE OBSERVATORY

APRAJITA VERMA1,∗, VASSILIS CHARMANDARIS2,3,6, ULRICH KLAAS4, DIETER LUTZ1 and MARTIN HAAS5 1Max-Planck Institut fur¨ extraterrestrische Physik, Postfach 1312, D-85741 Garching, Germany 2Department of Astronomy, Cornell University, 106 Space Sciences Building, Ithaca, NY 14853, U.S.A. 3Chercheur Associe,´ Observatoire de Paris, F-75014, Paris, France 4Max-Planck-Institut fur¨ Astronomie, Konigstuhl¨ 17, D-69117 Heidelberg, Germany 5Astronomisches Institut, Ruhr-Universitat,¨ Universitatsstr.¨ 150, D-44780 Bochum, Germany 6Present address: Department of Physics, University of Crete, GR-71003 Heraklion, Greece (∗Author for correspondence: E-mail: [email protected])

(Received 4 November 2004; accepted 15 November 2004)

Abstract. Some of the most ‘active’ galaxies in the Universe are obscured by large quantities of dust and emit a substantial fraction of their bolometric luminosity in the infrared. Observations of these infrared luminous galaxies with the Infrared Space Observatory (ISO) have provided a relatively unabsorbed view to the sources fuelling this active emission. The improved sensitivity, spatial resolution and spectroscopic capability of ISO over its predecessor Infrared Astronomical Satellite (IRAS) of enabled significant advances in the understanding of the infrared properties of active galaxies. ISO surveyed a wide range of active galaxies which, in the context of this review, includes those powered by intense bursts of star formation as well as those containing a dominant active galactic nucleus (AGN). Mid-infrared imaging resolved for the first time the dust enshrouded nuclei in many nearby galaxies, while a new era in infrared spectroscopy was opened by probing a wealth of atomic, ionic and molecular lines as well as broad band features in the mid- and far- infrared. This was particularly useful, since it resulted in the understanding of the power production, excitation and fuelling mechanisms in the nuclei of active galaxies including the intriguing but so far elusive ultraluminous infrared galaxies. Detailed studies of various classes of AGN and quasars greatly improved our understanding of the unification scenario. Far-infrared imaging and photometry revealed the presence of a new very cold dust component in galaxies and furthered our knowledge of the far-infrared properties of faint starbursts, ULIGs and quasars. We summarise almost nine years of key results based on ISO data spanning the full range of luminosity and type of active galaxies.

Keywords: galaxies: active, galaxies: starburst, galaxies: Seyfert, galaxies: quasars, infrared: galaxies

1. Introduction

“Activity” is a common feature of galaxies that emit a large fraction of their energy in the infrared. Existing primarily in a state of quiescence, during certain periods of their evolution galaxies experience short-lived phases of extreme star formation and/or black hole activity. From early ground-based measurements (e.g. Low Space Science Reviews (2005) 119: 355–407 DOI: 10.1007/s11214-005-8072-z C Springer 2005 356 A. VERMA ET AL. and Kleinmann, 1968; Kleinmann and Low, 1970a,b) to surveys performed by the Infrared Astronomical Satellite (IRAS, Beichman et al., 1988), many previously known active galaxies (e.g. optical- or radio-selected quasars, active galactic nuclei (AGN), and starbursts) were found to be as or more powerful in the infrared than in the visible wave length range. In addition, a new population of optically faint, in- 11 frared luminous galaxies (luminous infrared galaxies (LIGs, with LIR > 10 L)) were discovered that emit most of their bolometric luminosity in the infrared (Houck et al., 1984). AGN and massive bursts of star formation are the only known mechanisms that are capable of generating such luminosities. However, without spectroscopic information of the obscured components, classifying the underlying power source has proved to be difficult. Furthermore, quantitatively assessing the contribution to the total power from the two possible mechanisms is non-trivial. It is known that dust is formed during the late stages of stellar evolution and it is destroyed by intense radiation fields and shocks. Consequently, a detailed understanding of the physical mechanisms responsible for the observed infrared emission is an essential first step into revealing the nature of the underlying exciting source. As a result, since the most ‘active’ regions in our Universe appear to be enshrouded in large quantities of gas and dust, addressing the aforementioned points would have serious consequences in estimating star formation and black hole activity in the Universe. In a cosmological context, both the reported increase in the star formation density of the Universe for 0.1 z 2 (e.g. Madau et al., 1996; Barger et al., 2000) and the high frequency of starbursts hosted in galaxies with disturbed morphologies or in interacting/merging systems, imply that starbursts play an important role in galaxy formation and evolution. In addition, black hole activity could be buried in any 12 luminous infrared system. Ultraluminous infrared galaxies (ULIGs, LIR > 10 L) have been proposed to be an evolutionary stage in the life of a quasar and the local analogues of the z > 2 (sub-)mm population. Quantifying the amount of obscured activity at earlier times and its contribution to the infrared and X-ray backgrounds remain open issues requiring a sound understanding of the properties of local active galaxies. Following on from the coarse, photometric legacy of IRAS, ISO provided the first means to investigate the physical conditions of local active galaxies in significant detail at wavelengths where the obscured, and often most active, regions could be probed. Equipped with imaging, photometric and spectroscopic capabilities, ISO surveyed a wide range of emission properties from broad-band spectral energy distributions (SEDs), through imaging to detailed spectral analysis. We discuss the infrared emission properties of obscured active galaxies that ISO surveyed during its lifetime. In the context of this chapter, ‘active’ refers to both star formation and black hole activity and encompasses the range of non-normal galaxies: interacting/merging galaxies; starbursts; radio galaxies; AGN; quasars; low ionisation nebular emission regions (LINERs); and the full suite of LIGs–ULIGs and 13 hyperluminous infrared galaxies (HyLIGs > 10 L). OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 357

2. Activity Manifest in the Infrared

High energy UV (and visible) photons emanating from active regions can heat or excite environmental dust, located around active sites or distributed throughout the interstellar medium (ISM), resulting in re-radiation in the infrared. Grains ranging from the very small (radius a 10 nm) to the large (a ∼ 30 μm) contribute to emission over the 2.5–200 μm range. The type, size and distribution of dust grains together with the incident interstellar radiation ﬁeld (ISRF) shape the observed spectral features and the overall form of the SED which arise from both thermal and non-thermal processes. At long infrared wavelengths (λ 25 μm), emission predominantly arises from grains that re-radiate in thermodynamic equilibrium. In young radio galaxies, thermal bremsstrahlung may also provide a contribution to the infrared continuum. Emission related to transient rather than steady-state heating of dust grain complexes can dominate the SEDs of some active galaxies giving rise to continua and features over 3–18 μm. Non-thermal components include the featureless far-infrared tail (increases with increasing wavelength) of the synchrotron radiation spectrum in strong radio galaxies.

2.1. INFRARED CONTINUA:THERMAL AND NON-THERMAL

The mid-infrared spectra of galaxies display two continua: A flat continuum for λ 5 μm and, in some galaxies, an unrelated continuum for λ 11 μm (Helou et al., 2000). These continua are not produced by grains heated to thermal equilibrium by the radiation field but arise from the stochastic heating of small dust grains (amorphous silicates and graphites, a ∼ 0.01–0.1 μm) which are transiently heated by the absorption of a single photon (Sellgren, 1984; Beichman, 1987; Boulanger et al., 1998b; Draine and Li, 2001; Draine, 2003). A single photon with energy of only a few electron volts (λ 0.4 μm), is sufficient for the onset of infrared emission (Uchida et al., 1998). The steeply rising (near-thermal) continuum longward of 11 μm, which is strong in some active galaxies (but weak in quiescent galaxies), is thought to originate from the excess transient heating of very small fluctuating dust grains (VSG, a ≤ 10 nm). This component appears to be a characteristic of intensely star-forming regions (Desert et al., 1990; Verstraete et al., 1996; Cesarsky et al., 1996b; Laureijs et al., 1996). Under this transient heating regime, the mid-infrared continuum is directly proportional to the underlying radiation field over several orders of magnitude and varies with the conditions of the H II regions from which it originates (Boulanger et al., 1998a; Laurent et al., 2000; FörsterSchreiber et al., 2004). For example, in the close vicinity of the dense starburst knots in the Antennae, this continuum is shifted to shorter mid-infrared wavelengths implying higher temperatures (Vigroux et al., 1996). The 3 μm <λ<5 μm continuum is likely to arise from a featureless fluctuating component (Helou et al., 2000). Although often weak in active galaxies, the direct 358 A. VERMA ET AL. stellar continuum can also contribute (λ<5 μm) and should be considered in careful spectral decomposition analysis. The strength of this component increases with decreasing wavelength and thus the transition between direct stellar to interstellar dust emission occurs within the near- to low-mid-infrared (Boselli et al., 1997, 1998; Alonso-Herrero et al., 2001). The mid-infrared also includes the transition from the transient to the steady-state heating regime. The wavelength at which this transition occurs is determined by the incident radiation field above which grains of a certain size no longer suffer large temperature changes. Dust grains are heated by the general ISRF and attain a ‘steady state’ where emission and absorption reach an equilibrium. The grain temperature at which this equilibrium occurs charac- terises the emission giving rise to a thermal black-body continuum. Qualitatively, the mid- to far-infrared continua of active galaxies are often approximated by a series of superposed black-bodies, each originating from a body of dust heated to a characteristic temperature (Haas et al., 1998a; Ivison et al., 1998; Klaas et al., 2001; Bendo et al., 2003; Stevens et al., 2005). This approximation provides only an indication of typical temperatures and is not an accurate description of the complex heating of a variety of dust grains. More accurate descriptions are achieved by considering a multi-grain dust model and complex radiation fields using, for example, semi-empirical models (e.g. Dale et al., 2001a) or radiative transfer theory (e.g. Acosta-Pulido et al., 1996; Silva et al., 1998; Alexander et al., 1999a; Siebenmorgen et al., 1999b; Efstathiou et al., 2000; Verma et al., 2002; Farrah et al., 2003; Freudling et al., 2003; González-Alfonso et al., 2004; Siebenmorgen et al., 2004a). For sources that are spatially unresolved by the ISO instruments, SED decomposition (using one of the methods described earlier) is the only means to investigate the underlying physical processes. Some AGN-hosting galaxies display a relatively broad and flat (in ν Fν) hot continuum that dominates their mid-infrared SEDs. This emerges from grains in the putative torus that are heated up to the grain sublimation temperature (Tsub ∼ 1500 K) by the strong radiation field of the accretion disk of the central AGN. The ‘warm’ IRAS galaxy criterion (S60/S25 > 0.2) selects galaxies displaying this component which are mostly identified with AGN (de Grijp et al., 1985). The extended wavelength range of ISO over IRAS enabled (a) the precise determination of the turnover in the far-infrared SEDs of active galaxies and (b) the identification of two dust components that comprise the far-infrared emission that are associated with large grains in thermal equilibrium. The first (or ‘cold’) component represents dust surrounding the star-forming regions and thus the strength of the component is directly related to the star formation rate (however, see Section 4.5). While the SEDs of normal quiescent galaxies peak at ∼150 μm, the enhanced heating in actively star-forming regions produces a strong far-infrared peak at ∼60– 100 μm that dominates the infrared SED and corresponds to dust heated to 30–60 K (e.g. Calzetti et al., 2000; Dale et al., 2001a; Spinoglio et al., 2002). This component also enhances the mid-infrared continuum (e.g. see Figure 1 in Sanders and Mirabel, 1996). In AGN, the fractional contribution of this ‘cold’ far-infrared OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 359 component can be weaker in comparison to the warmer mid-infrared component (λ ∼ 30 μm) than in starbursts, reflective of the intense heating power of the AGN. The second (‘very cold’ or ‘cirrus’) component is dust that is spatially more extended and is associated with the H I disk or the molecular halo (Trewhella et al., 2000; Radovich et al., 2001; Stickel et al., 2004). Located far and, in some galaxies, shielded (Calzetti et al., 2000) from the intense emission from the central engine, the extended component is heated by the ambient interstellar radiation field to temperatures of 10–20 K (Odenwald et al., 1998; Alton et al., 1998; Davies et al., 1999; Haas et al., 1998b; Haas, 1998; Trewhella et al., 2000). These temperatures are consistent with the theoretically predicted heating/cooling of dust grains in thermal equilibrium by a diffuse radiation field. Both profile measurements for resolved sources and far-infrared-sub-mm SED analysis for more distant galaxies reveal that a very cold component is required to explain all the emission beyond 100 μm (Rodr´ıguezEspinosa et al., 1996; Radovich et al., 1999; Siebenmorgen et al., 1999a; Domingue et al., 1999; Calzetti et al., 2000; Haas et al., 2000a, 2003; Klaas et al., 2001; PérezGarc´ıaand Rodr´ıguezEspinosa, 2001). This component, that IRAS was insensitive to, may be a significant contributor to the far-infrared luminosity in some active galaxies (e.g. 60% in local starbursts, Calzetti et al., 2000).

2.2. FEATURES

The advent of ISO spectroscopy enabled the first sensitive and high-resolution mid- to far-infrared spectroscopic measurements of active galaxies providing a better understanding of the physical conditions within visually obscured regions of systems such as Circinus (Moorwood et al., 1996; Sturm et al., 2000), the starburst M82 (Sturm et al., 2000) and the Seyfert 2 NGC 1068 (Lutz et al., 2000b). The full 2.5–200 μm spectra of these three galaxies are shown in Figure 1, and display the richness of the infrared range and the differences between these active galaxy templates (Sturm et al., 2000). The composite source Circinus displays broad mid- infrared features, fine structure lines of low excitation, H-recombination lines and a thermal continuum that all predominantly trace the starburst component these are intermingled with the strong VSG and mid-infrared AGN continua and high excitation fine structure lines that are predominantly caused by the hard radiation field of an AGN. In addition, warm and cold molecular gas and warm atomic gas are probed by rotational lines of H2, molecular absorption features (e.g. OH, CH, XCN and H2O) and very low excitation (<13.6 eV) fine structure lines (e.g. [O I], [C II]), respectively.

2.2.1. Unidentiﬁed Infrared Bands The mid-infrared spectra of many active galaxies are dominated by characteristic broad emission features, the most prominent of which are located at 3.3, 6.2, 7.7, 8.6 and 11.3 μm. They are invariably observed together and are commonly referred 360 A. VERMA ET AL.

Figure 1. Composite SWS+LWS spectra of Circinus, M82 and NGC 1068 displaying the spectral differences between the different source types (E. Sturm, private communication). OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 361 to as unidentified infrared bands (UIBs). These features have been detected in a wide range of celestial sources (e.g. Allamandola et al., 1995; Cesarsky et al., 1996a,b; Verstraete et al., 1996; Mattila et al., 1996, 1999; Lemke et al., 1998) and are commonly detected in the spectra of normal galaxies. UIBs may originate from C C and C H bending and stretching vibrational modes in aromatic hydrocarbons (Leger and Puget, 1984; Allamandola et al., 1985; Puget and Leger, 1989) and the favoured carriers are polycyclic aromatic hydrocarbons (PAHs). The features are attributed to PAH clusters (Boulanger et al., 1998b), while individual PAHs can contribute to the 5–10 μm continuum and may be responsible for the non-linear continuum rise seen in the far UV (Siebenmorgen and Kruegel, 1992). Alternative carriers are small amorphous carbon grains that emit when exposed to moderately intense UV/visible radiation (Puget and Leger, 1989; Boulade et al., 1996; Uchida et al., 1998; Pagani et al., 1999). Since the exact molecular structure of the carriers remains under debate, the features are referred to as UIBs throughout this text (for a more thorough review on the structure of UIBs see the chapter of Peeters et al. in this volume or Peeters et al. (2003)). Where strongly detected, the carriers are likely to provide the majority of the flux in the 3–12 μm band and show near- invariant profiles in a variety of galaxies (Rigopoulou et al., 1999; Helou et al., 2000; Förster Schreiber et al., 2003). The widths of the lines are thought to arise from temperature broadening of the vibrational bands and have been fit using Lorentzian rather than Gaussian profiles (Boulanger et al., 1998b; Verstraete et al., 2001) but see the discussion in Peeters et al. (this volume) for further details on this topic. The spatial correlation between mid-infrared emission with atomic and molecular hydrogen suggests that UIB emission originates at the interface of H II and molecular clouds, the so-called photo-dissociation regions (PDRs, e.g. Verstraete et al., 1996; Cesarsky et al., 1996b; Pagani et al., 1999). Mid-infrared starburst spectra are generally dominated by these UIB features and show a strong continuum from very small grains. The relative strengths of the features are sensitive to local radiation fields, distribution of grains throughout the galaxy system and extinction (Verstraete et al., 1996; Sturm et al., 2000; Lutz et al., 1998b; Laurent et al., 2000). UIB emission in starburst regions originates from UV-heated ionised UIB carriers. A significant contribution can also arise in heating of neutral UIB carriers in the extended regions by the diffuse ISRF, i.e. by less-energetic (visible) photons, which is mostly seen in normal galaxies and is proposed in some active galaxies (Mattila et al., 1999; Smith, 1998; Pagani et al., 1999; Haas et al., 2002). In intense radiation fields (∼105 times that of the local ISRF) such as in metal poor environments or in the vicinity of super star clusters or AGN, UIBs are weak or absent (Contursi et al., 2000; Thuan et al., 1999; Rigopoulou et al., 1999; Sturm et al., 2000; Santos-Lleó et al., 2001). Small grains and UIB carriers are suscepti- ble to high temperature fluctuations and photons with sufficient energy may cause grain sublimation or destruction. This is reflected in the UIB-to-mid-infrared continuum strength ratio (7/15 μm flux density) that is seen to be lower (>0.5) in active 362 A. VERMA ET AL. regions and galaxies than in more quiescent regions such as in normal galaxies and the disks of some active galaxies (ratio ∼1 with a large dispersion, see Figure 1 of Vigroux et al. (1999) and Figure 12 in Dale et al. (2000)). Weak detections of UIBs in galaxies with AGN probably originate from the cooler material such as circumnuclear starburst regions or diffusely heated extended regions rather than from the immediate vicinity of the AGN torus. In nearby Seyferts, spatially resolved mid-infrared spectroscopy suggested that the absence or suppression of UIB emission is due to the fact that the dust is predominantly heated by processes related to the central AGN (e.g. in NGC 1068 (Le Floc’h et al., 2001), Circinus (Moorwood, 1999), NGC 4151 (Sturm et al., 1999), Mrk279 (Santos-Lleó et al., 2001)). This was confirmed through a comparison of ISO and high-spatial resolution ground-based data of AGN and starbursts, that demonstrated the absence of UIBs in the nuclei of AGN-hosting galaxies (Siebenmorgen et al., 2004b). For unresolved sources, the VSG continuum may be so strong that the UIB features are diluted by the over- powering continuum resulting in their non-detection (Lutz et al., 1998b). Thus, at high redshifts, where low-metallicity- and/or AGN-dominated systems may be more prevalent, the intrinsic weakness of UIB features seen in local systems of these types would imply that the use of UIB features as a tracer of similar systems at high redshift may be problematic.

2.2.2. Fine Structure Lines Fine structure lines are tracers of nebular conditions such as excitation and the intrinsic ionising spectrum. By virtue of their differing excitation potentials and critical densities they provide an insight into the energetics and chemical composition of the regions from which they originate. Starburst galaxies commonly show low excitation fine structure lines (e.g. [Ar II], [Ar III], [Fe II], [Ne II] and [S III]) that have ionisation potentials between 13.6 and ∼50 eV and mostly arise from H II regions that have been photoionized by their central massive stars. A PDR or shock origin may also contribute to the flux of some low-excitation lines. High-excitation lines with potentials exceeding 50 eV (e.g. [Ne V, VI], [S IV], [Mg V, VII, VIII][OIV], [Si V, VI, VII]) are present in the spectra some low-metallicity galaxies but are more commonly found in AGN where high-energy photons, produced by non-thermal processes related to the central black hole, photoionise the gas. Excitation to these high-energy levels (>50 eV) is difficult to attain by photons emitted by stars in typical H II regions. However, the mid-infrared spectra of some starbursts do show the high-excitation [O IV] and [S IV] lines. The presence of these lines is associated to collisionally excited, shocked or coronal gas (e.g. Lutz et al., 1998a). Active galaxy spectra also show fine structure lines with excitation potentials below the ionisation potential of hydrogen (e.g. [O III]5288 μm, [O I]63,145 μm and [C II]158 μm), that trace mostly cool atomic and molecular clouds. [C II]158 μm and [O I]63 μm are important cooling lines in PDRs but also originate in widespread atomic gas (see Section 4.6). OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 363

As infrared ﬁne structure lines are fairly insensitive to uncertainties of the effective temperature of the gas, and because they are less affected by extinction than optical lines, they provide an unequalled insight into the properties of the ISM (e.g. electron temperature and density, ISRF, metallicity) of active galaxies and their embedded power source. Ratios of two lines of the same species provide information on the nebular conditions such as electron temperatures and densities. Ratios of high- to low-excitation states of the same element probe the hardness of the ISRF and ratios of high- to low-excitation lines of different elements have been used as tracers of AGN presence and thus contribute to nebular diagnostics (see Section 6.1). ISO-spectroscopy enabled a major advance over IRAS by allowing the dominant fuelling mechanism in a given source to be identiﬁed (Genzel et al., 1998; Laurent et al., 2000; Sturm et al., 2000).

2.2.3. Molecular Features Molecular features trace the cold dense molecular material in active galaxies. Sev- eral active gas-rich starbursts, Seyferts and ULIGs display a range of molecular lines: the rotational lines of molecular hydrogen (Section 2.2.3.2), OH, H2O, NH3 and CH are seen in absorption or emission (see Figure 1 of Fischer et al., 1999b; also refer to Hur et al., 1996; Fischer et al., 1997, 1999a,b; Bradford et al., 1999; Colbert et al., 1999; Sturm et al., 2000; Rigopoulou et al., 2002; Spoon et al., 2002; Spinoglio et al., 2005). Molecular absorption features appear to be more common (and prominent) in sources of high luminosity (e.g. ULIGs) than those of lower luminosity (Fischer et al., 1999a,b; Spoon et al., 2002). Spectra displaying molecular absorption features show a wide variety of excitation: from starburst galaxies like M82 which display features indicating that most of the O and C molecules occupy the ground state, to objects such as ULIGs where a signiﬁ- cant population exists at higher energy levels (Fischer et al., 1999a,b). Additional absorption features at 6.85 and 7.25 μm are seen in the spectra of some active galaxies and ULIGs (Spoon et al., 2002). They are attributed to CH-deformation modes of carbonaceous material because of their similarity to features seen along lines of sight towards Sgr A* (Chiar et al., 2000). ULIG and Seyfert spectra also exhibit an absorption feature at 3.4 μm that is likely to be aliphatic CH stretch absorptions of refractory carbonaceous material (Imanishi, 2002; Dartois et al., 2004). 2.2.3.1. Silicate Absorption. The 9.7 and 18 μm absorption features arise from the stretching and bending modes of SiO. The extinction suffered by a galaxy can be estimated by the depth of these lines. For example, the extreme 9.7 μm silicate absorption feature seen in the starburst-ULIG Arp220 provides a lower limit to the optical depth of AV 45 (Charmandaris et al., 1997), but see Spoon et al. (2004) for a more detailed analysis of this system). Though it is possible that the depth of the 9.7 μm may be overestimated in the feature-dense mid-infrared spectra of starbursts exhibiting UIB emission (Sturm et al., 2000; F¨orsterSchreiber et al., 2003). The 364 A. VERMA ET AL. extreme extinction of the starburst M82 (AV = 15–60) derived by Gillett et al. (1975) was ascribed to this problem as ISO data could not support such a high extinction (Sturm et al., 2000).

2.2.3.2. Molecular Hydrogen and Warm Molecular Gas. Rotational transitions of H2 that originate in dense molecular media were detected in the spectra of several active galaxies: NGC 3256 (Rigopoulou et al., 1996); NGC 891 (Valen- tijn and van der Werf, 1999); NGC 1068 (Lutz et al., 2000b); NGC 4945 (Spoon et al., 2000); NGC6 240 (Lutz et al., 2003). Low rotational transitions originate from warm gas (T ∼ 100–200 K) and the higher and rovibrational lines from much hotter gas (T 1000 K). Rigopoulou et al. (2002) investigated a sample of starbursts and Seyferts observed by ISO presenting pure rotational lines from S(7) to S(0). Irrespective of starburst or Seyfert type, temperatures of ∼150 K are derived from the S(1)-to-S(0) ratio. A similar temperature was found for NGC 4945 (Spoon et al., 2000). While ∼10% of the galactic gas mass in a starburst system can be accounted for by the warm component, on average higher fraction (18%, range 2–35%) are observed for Seyferts (Rigopoulou et al., 2002). The temperature of the molecular hydrogen implies that it arises from a combination of emission from fairly normal PDRs (at least for starbursts), low velocity shocks and X-ray heated gas around the central AGN. The latter is more probable in Seyfert galaxies which accounts for the larger fraction of warm gas found in these systems. Shock heating appears to be the most likely origin of the extremely strong rotational H2 lines detected in the merging ‘double active nucleus’ system NGC 6240 (Lutz et al., 2003). While, the origin of shock heating can be largely associated to winds originating from massive stars and supernovae, recently Haas et al. (2005) have found the ﬁrst observational evidence of pre-starburst shocks in the overlap region of the early merger, the Antennae, (see Section 4.1). From a re-analysis of archival ISOCAM-CVF data, they ﬁnd that the strongest molecular hydrogen emission (traced by the H2S(3) line) is displaced from the regions of active star formation. This, together with the high line luminosity (normalised to the far-infrared luminosity) indicates that the bulk of excited H2 gas is shocked by the collision itself.

2.2.3.3. Ice Absorption and Cold Molecular Gas. An interesting discovery made by ISO was the first extragalactic detection of absorption features due to ices (H2O, CH4 and XCN) present in cold molecular components of starbursts, Seyferts and predominantly ULIGs. Sturm et al. (2000) first reported the detections of weak water ice absorption features at 3 μm in the spectra of local starbursts M82 and NGC 253. Very strong water ice, as well as the first detections of absorptions due to CO and CO2 ices, were detected in the actively star-forming galaxy NGC 4945 (Spoon et al., 2000). The presence of ices is related to the presence of cold material and the radiation environment. One would expect cold molecular clouds in starbursts to be likely hosts, while more intense environments, such as OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 365 those predominantly fuelled by AGN, to be less favoured. This is corroborated by the lack of absorption features in the spectrum of NGC 1068 (Sturm et al., 2000) and the limits on the strength of the CO absorption feature set from ISOSWS spectra of 31 AGN (Lutz et al., 2004b). However, water ice absorption has been detected in Seyfert galaxies, e.g. in the UIB-free spectrum of the Seyfert galaxy NGC 4418 (Spoon et al., 2001). The absorption-dominated spectrum bears strong similarities to the spectra of embedded protostars, and the depth of absorption features implies that its Seyfert nucleus is so deeply embedded that ices are shielded from the AGNs intense radiation field. In a heterogeneous sample of 103 active galaxies with high signal-to-noise mid-infrared spectra, approximately 20% display absorption features attributed to ices (Spoon et al., 2002). While ices are weak or absent in the spectra of starbursts and Seyferts, in ULIGs they are strong, considerably stronger than the weak features found previously in M82 and NGC 253. This implies that rather than the radiation environment, it is the amount of cold material that determines their presence in these most luminous and massive infrared galaxies that are known to contain large concentrations of molecular material (e.g. Solomon et al., 1997; Downes and Solomon, 1998). Although plausible, more sensitive (and spatially resolved) mid-infrared spectroscopy is required to understand these differences, and the distribution and energetic conditions of cold molecular material. More- over, higher spectral resolution over a wider mid-infrared wavelength range is necessary for accurate decomposition of the spectra into their various components, which will aid the identification of subtle and/or blended features (e.g. Spoon et al., 2004). Spoon et al. (2002) proposed an evolutionary sequence for their ice spectra ranging from a highly obscured beginning of star formation where ices dominate the spectra, to a less obscured stage of star formation where the UIBs become stronger. At any stage, the star formation may be accompanied by AGN activity.

2.2.3.4. OH Features and (Mega-)Masers. Transitions of the OH molecule were also detected in the ISOSWS and ISOLWS spectra of some active galaxies (Fischer et al., 1999b), such as in the starbursts NGC 253 (Bradford et al., 1999) and NGC 4945 (Brauher et al., 1997) permitting estimation of the OH column density and abundance. The depths of the OH features detected in the ISOLWS far-infrared spectra of galaxies hosting known OH (mega-)masers were found to be sufﬁcient to provide all of the infrared photons required to radiatively pump the maser activity (e.g. in the starburst-ULIG Arp220 (Skinner et al., 1997; Suter et al., 1998); AGN-ULIG Mrk231 (Suter et al., 1998); the merging double systems ULIG IRAS 20100-4156 and LIG 3 Zw 35 (Kegel et al., 1999); the starburst NGC 4945 (Brauher et al., 1997)). In Arp220, the fact that the far-infrared continuum source is extended and that the maser must be located in front of it favours the interpretation that a starburst provides the continuum photons for its maser (Skinner et al., 1997). The high pump rate determined for this source suggests that pumping mechanisms other 366 A. VERMA ET AL. than radiative pumping (such as collisional pumping) may contribute (He and Chen, 2004).

2.3. MAGNETICALLY ALIGNED DUST GRAINS

Polarised emission detected from active galaxies is ascribed to dichroic absorption by the ISM of a galaxy i.e. by elongated dust grains that are magnetically aligned. This has been observed in galactic sources where the position angle of the measured polarisation indicates the orientation of the projected magnetic field. Polarisation is an important probe of the physical conditions in active galaxies, as it can provide constraints on the formation of the putative torus (Alexander et al., 1999b). Such polarisation studies were performed for some ULIGs and revealed 3–8% of the 5–18 μm flux to be polarised (Siebenmorgen and Efstathiou, 2001). The archetypal starburst-like ULIG Arp220 has the lowest polarisation, while the warm AGN-like ULIG Mrk231 has the highest polarisation. The first polarisation profile of a starburst galaxy was made at 6 μm for NGC 1808 showing a 20% increase towards the outer regions (Siebenmorgen et al., 2001). This, together with a complementary measurement at 170 μm (unresolved), are best explained by the large-scale (∼500 pc) magnetic alignment of (non-spherical) large grains (≥100 A).˚

3. Very Cold Dust and Diffuse Radiation Fields: ISM

ISOPHOT-data clearly confirmed the presence of very cold dust (T ∼ 10–20 K) in a variety of infrared galaxies. Often in conjunction with supplementary (sub-)mm photometry, this component was detected in spatially resolved maps or intensity profiles of local galaxies (e.g. Haas et al., 1998b), and was isolated as a significant contributor to the far-infrared luminosity in SED analyses of unresolved active galaxies (Calzetti et al., 2000; Klaas et al., 2001; Spinoglio et al., 2002; Haas et al., 2003). Although cold ‘cirrus’ or ‘disk’ components in galaxies were indicated by IRAS colours, IRAS was insensitive to this very cold component due to its shorter wavelength coverage (<100 μm). As a result, the gas-to-dust ratios of distant galaxies based solely on the IRAS data were overestimated, whereas the addition of ISO measurements resulted in values which are closer to that of the Galactic value (∼160) (Haas et al., 1998b). The very cold component often accounts for a substantial fraction of far-infrared emission and is most commonly seen in early type and disk galaxies (see Sauvage et al., this volume). With a scale-length in excess of the stars, less than the atomic HI gas and similar to that of the molecular gas, it is likely that the dust is associated to this cold gas components (e.g. Davies et al., 1999; Popescu and Tuffs, 2003; Popescu et al., 2004). Very cold dust is also seen in both starburst and AGN-hosting active galaxies. NGC 253, which hosts a nuclear starburst, has been shown to contain extended cold OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 367 dust along the major and minor axes. While the dust along the minor axis is attributed to outflows, the cold emission along the major axis is shown to have a scale-length 40% greater than that of the stars (Radovich et al., 2001). The correlation of the 90 μm emission profile of four Seyfert galaxies with the R-band extent but not with the star formation tracer Hα suggests that a large contribution of the far-infrared emission from Seyferts does not arise from the (nuclear) star-forming regions but from a cooler far-infrared disk component within which normal stars heat dust grains (PérezGarc´ıa et al., 2000). The detection of substantial cold far-infrared emission (T < 20 K) from the star-forming obscured overlap region of the merging Antennae galaxy (see Section 4.1) indicates the presence of cold dense dust within which stars are forming (Haas et al., 2000a). SED analysis using black or grey-body decomposition suggested the need for more than one dust component, including a very cold dust component, to explain the emission above 100 μm of many galaxies. However, by allowing both temperature and dust grain emissivity to be free parameters, it is possible to approximate the entire far-infrared SED (λ>50 μm) using a single modified black-body. Invariably, for active galaxies, this was achieved with temperatures ∼30–50 K and emissivity values of 1.5–2 (e.g. Colbert et al., 1999; Siebenmorgen et al., 1999a). However, doing so raised conflicts with other observational constraints (see Klaas et al. (2001) for a full description). Moreover, by leaving the dust emissivity as a free parameter presupposes that intrinsic dust properties differ from source to source, with no physical basis supporting this assumption. By choosing a Galactic value of γ = 2 for the grain emissivity (i.e. assuming that dust properties do not vary between sources), single black-bodies fail to explain all the far-infrared emission – additional cold components (T ∼ 10–20 K) are required that are associated with the spatially resolved extended very cold components described earlier. The SEDs of active galaxies are warmer than those of normal galaxies and the far-infrared part is generally fit by two components representing active star formation and very cold (diffusely heated) material. The very cold dust can be a significant contributor in terms of galaxy mass and energy output. For example, two of the three components used to fit the SED of the LINER NGC 3079 have very cold temperatures of 12 and 20 K, respectively (Klaas and Walker, 2002). Calzetti et al. (2000) find that two dust components are required at T = 40–55 K and T ∼ 20–23 K to explain the far-infrared SEDs of eight local starbursts. The latter comprises up to 60% of the total flux and the associated mass is as much as 150 times that of the former. For the most actively star-forming galaxies or those with strongest radiation fields, the contribution to the bolometric luminosity from cold emission (as measured in the 122–1100 μm range) was shown to be smaller in actively star-forming galaxies (5%) than in cirrus-dominated galaxies (∼40%) (Dale et al., 2001a). SED analysis of a wide range of active galaxies indicates that very cold components are common whether they are centrally fuelled by AGN or starbursts, suggesting that the central activity does not directly affect the far-infrared emission. 368 A. VERMA ET AL.

4. Cold Dust and Low Excitation Radiation Fields: Starburst Regions

The T = 40–45 K temperature component found in local starbursts by Calzetti et al. (2000) mentioned earlier represents the emission from dust surrounding sites of active star formation. For starburst galaxies, this emission (peaking between 60 and 100 μm) dominates the infrared SED and is the major contributor to the infrared luminosity. This ‘cold’ infrared emission is generally less extended than the very cold component described in the preceding section. In some nearby galaxies, the far-infrared morphology traces star-forming regions well. For example, the strongly far-infrared-emitting dark molecular cloud coincides with a cluster of star- forming regions in the galaxy NGC 4051 (PérezGarc´ıa et al., 2000). Similarly, far- infrared emission could be resolved for the pair members of interacting spirals where enhanced star formation has been established (Xu et al., 2000; Gao et al., 2001). The correlation of the warmer far-infrared emission with the R-band morphology of local Seyferts hosting compact nuclei and circumnuclear star formation further supports that this far-infrared component is related to star formation activity (Rodr´ıguez Espinosa and PérezGarc´ıa,1997). Moreover, far-infrared emission has also been spatially associated with regions exhibiting signs of strong star formation (e.g. strong UIBs). Roussel et al. (2001) demonstrated that the integrated far-infrared flux and the star formation related mid-infrared flux are strongly correlated in the disks of normal star-forming galaxies, indicating that the grains responsible for the mid- and far-infrared share a common heating source.

4.1. THE IMPORTANCE OF INTERACTIONS/MERGERS

Interactions and mergers enhance starburst activity, both nuclear and extra-nuclear (see review by Struck, 1999). The final merger stage plays a critical role in how active galaxies form stars and may influence the appearance of AGN and the onset of the ULIG phase. The fraction of disturbed or interacting systems in infrared selected samples of galaxies appears to increase with luminosity with nearly 100% of all ULIGs displaying evidence of interaction (see Sanders and Mirabel, 1996, and references therein). Numerical simulations have also helped in establishing that interactions and mergers play an important role in the formation and evolution of such galaxies (Mihos and Hernquist, 1996). The induced gravitational instabilities form bars which strip in-falling gas of angular momentum and enable radial inflows to the nuclear regions that can feed starbursts or AGN (Combes, 2001). Moreover, mergers appear to be the triggering mechanism for the ultraluminous phase. Dynamical shocks and tidal forces are the only mechanisms that can cause transference and compression of sufficient quantities of matter to fuel luminosities 11 in excess of 10 L. ISOCAM’s spatial scale was sufficient to map the morphologies of local interacting active systems (see Mirabel and Laurent, 1999 for a review) and agree with OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 369 results from ground-based high-resolution images that (optically faint or invisible) compact sources located in deeply obscured nuclear regions or interaction interfaces dominate the mid-infrared emission (e.g. Soifer et al., 2000, 2001; Charmandaris et al., 2002; Gallais et al., 2004; Siebenmorgen et al., 2004b). For some sources, diffuse emission between these ‘hot-spots’ can also contribute significantly to the total mid-infrared emission (e.g. Le Floc’h et al., 2002; Gallais et al., 2004). This diffuse emission can be enhanced in merging systems that are common among active galaxies. By probing through the famous characteristic dust lane of the AGN-hosting elliptical galaxy Centaurus A (NGC 5128), imaging with ISOCAM revealed a barred ‘mini-spiral’ located in its central 5 kpc. The mini-spiral was proposed to be tidal debris from a gas-rich object that was accreted in the last gigayear (Mirabel et al., 1999; Block and Sauvage, 2000). Evidence of accreted companions in Centaurus A exists from previous studies of its stellar and gas kinematics (Charmandaris et al., 2000, and references therein). ISOPHOT contributed to the accretion scenario through the detection of the northern cold dust shell that is possibly the ISM remnant of a captured disk galaxy (Stickel et al., 2004). As a result, Centaurus A is a prime example of a giant elliptical that must have experienced several minor mergers. It is likely that, as the interactions took place, tidal forces funnelled gas into the dynamical centre and boosted the star formation rates in the gas-rich components forming the dust observed in the infrared and visible.

4.1.1. Nuclear and Extra-Nuclear Starbursts Disturbed morphological signatures such as rings, tidal tails, bridges and arcs are clearly detected in ISOCAM observations of numerous merging galaxies (e.g. the Antennae (Mirabel et al., 1998) Figure 2; UGC12915, 12914 (Jarrett et al., 1999); NGC 985 (Appleton et al., 2002); Arp299 (Gallais et al., 2004)). The mid-infrared morphologies depend on the mass ratio of the progenitor galaxies and the merger stage. As well as enhancing star formation in the nuclear regions, interactions cause a general increase in the widespread activity of the parent galaxies. At the early stages of interaction, wide scale star formation can be induced and dust grains are heated by a generally warmer radiation ﬁeld. For example, ISOCAM observations of the VV 114 galaxy show that 60% of the mid-infrared ﬂux arises from a diffuse component that is several kpc in size (Le Floc’h et al., 2002). Nearly 40% of the 7 and 15 μm emission detected in the early-stage merger (nuclei separation 5 kpc) galaxy pair Arp299 (NGC 3690 and IC 694) is diffuse and originates from the interacting disks (Gallais et al., 2004). In more advanced mergers, shocks and gravitational instabilities encourage the onset of star formation. The transfer of matter to nuclear regions and compression in interaction zones gives rise to enhanced mid-infrared emission in nuclear and often extra-nuclear sites. A clear example of extra-nuclear star formation was uncovered by ISOCAM in the intra-group medium of Stephan’s Quintet (Xu et al., 1999). 370 A. VERMA ET AL.

Figure 2. This composite ﬁgure of the Antennae is from Mirabel et al. (1998) showing the ISOCAM- CVF contours overlaid on a HST V and I band image. Half of the mid-infrared emission arises from starburst activity that is completely obscured in the optical, including the brightest mid-infrared knot (knot A). The CVF spectrum of this knot is shown in the lower part of the image displaying strong signatures of massive star formation – the [Ne III] line and a strong rising continuum above 10 μm.

The best example of ISOCAM’s capability in probing through highly extincted regions is imaging of the spectacular major merger in the Antennae (Arp244, NGC 4038/4039). The completely obscured overlap region in the HST-WFPC2 image (AV ∼ 35 mag; Verma et al., 2003) was shown by ISOCAM to contain a OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 371 massive star-forming region (50 pc in radius) comprising more than 15% of the total 12–17 μm luminosity which outshines the two nuclei by a factor of 5 (Mirabel et al., 1998). The warm mid-infrared emission traces the molecular gas that fuels star formation (Wilson et al., 2000, 2003), whereas the cold (T ∼ 30 K) and very cold dust (T ∼ 20 K) are traced by far-infrared ISOPHOT and sub-mm observations (Haas et al., 2000a). Haas et al. (2000a) find that the sub-mm knots, in particular K2, are likely to be in a pre-starburst phase and the simultaneous presence of powerful off-nuclear starbursts and pre-starbursts may be a general feature of colliding galaxies. The authors postulate that once star formation has commenced in these clouds, the Antennae may evolve from a luminous IR galaxy into an ultraluminous one. An interesting feature for the Antennae (also seen for Centaurus A) is that the most luminous 15 μm knots are not coincident with the darkest, most prominent dust absorption lanes seen in the HST image. Presumably the darkest optical dust lanes trace the coldest emission (i.e. that probed by the far-infrared) rather than the warm dust from which 15 μm emission arises. This is plausibly the cause of the displacement. Alternatively, projection effects may be responsible for this offset (Mirabel et al., 1998). Extranuclear overlap starbursts emitting in the mid-infrared are not unique to the Antennae. In a sample of eight spiral–spiral pairs with overlapping disks, Xu et al. (2000) found extranuclear, overlap region starbursts in five galaxies classified as advanced mergers or having severely disturbed morphologies (but not in less- disturbed systems). However, the starbursts in the bridges or tidal tails are generally fainter than those in the nuclear regions.

4.1.2. Collisional Rings Collisional ring galaxies are produced when a compact galaxy passes through the centre of a larger disk galaxy giving rise to a radial density wave that creates a massive star-forming ring (see review by Appleton and Struck-Marcell, 1996). If the collision is off centre then the ring morphology is asymmetric e.g. as in VII Zw 466 (Appleton et al., 1999). These systems are excellent local examples of induced star formation in colliding galaxies. The effects of collisions on the ISM can be investigated in detail, and may be relevant for higher redshift galaxies including ULIGs (Appleton et al., 2000). Azimuthal variations in over-density in the rings are often noted as well as extranuclear star-forming knots that show strong UIBs (Charmandaris et al., 1999, 2001b; Appleton et al., 1999, 2002). When detected, star formation in the ring contributes at a level of 10% of the bolometric output (Appleton et al., 1999, 2000). The mid-infrared emission is clearly enhanced in the nucleus and the expanding ring, even though diffuse emission from the intervening medium was also observed (Charmandaris et al., 1999, 2001b). The diffuse components coincide with regions where only weak Hα had been previously measured. The mid-infrared colours measured in collisional ring galaxies are similar to those of late-type galaxies, indicative of more modest star formation than is found 372 A. VERMA ET AL. for extranuclear starbursts (e.g. in the nucleus inner ring and spokes of Cartwheel, the entirety of Arp10 and the ring of Arp118 (Charmandaris et al., 2001b)). The morphological agreement between the mid-infrared and Hα or radio emission is good, although differences have been noted for VII Zw 466 (Appleton et al., 1999). A deviation from this picture is seen for the extranuclear starburst knot in the archetypal collisional ring galaxy, the Cartwheel, that coincides with the strongest Hα and radio knot. The regions have an unusually red mid-infrared colour (15/7 μm ∼ 5.2) which is the highest among the extragalactic regions observed with ISOCAM and it is even more extreme than the powerful starburst region (knot A) revealed in the Antennae (Charmandaris et al., 1999). Interestingly, this region also hosts 41 the remarkable hyperluminous X-ray source (L0.5–10 keV 10 ergs/s; Gao et al., 2003) which dominates the X-ray emission from the Cartwheel ring.

4.2. HOT STAR POPULATION,STARBURST EVOLUTION AND METALLICITY

Starburst galaxies, defined as having star formation rates that cannot be sustained over a Hubble time, are likely the birthplace of a large fraction of massive stars. Locally, four starburst galaxies can account for 25% of the local massive star population within 10 Mpc (Heckman et al., 1998) and are important as the providers of the Lyman continuum photons that ionise hydrogen. ISO spectroscopy enabled the stellar populations of starbursts to be probed using fine structure lines that are sensitive to the nebular conditions from which they originate (see FörsterSchreiber et al. (2001) for a detailed analysis of the stellar populations and radiation environment of the local starburst M82). The ratio of [Ne III]15.6 μm(Ep = 41 eV) and [Ne II]12.8 μm(Ep = 22 eV) is sensitive to the hardness of the stellar energy distributions of the OB stars which excite them. In a survey of 27 starburst galaxies, Thornley et al. (2000) confirmed earlier ground-based results that the excitation measured by the [Ne III]15.6 μm/[Ne II]12.8 μm ratio was lower in starburst galaxies than measured for Galactic H II regions, indicating a deficiency of massive stars in the starbursts. Thornley et al. suggest a lower upper mass cut-off to the initial mass function (IMF) or ageing of the massive star population to be responsible for the low excitation. Through extensive quantitative photo-ionisation modelling, Thornley et al. found the ratios to be consistent with the formation of massive stars (50–100M). This, together with the known presence of very massive stars in local starburst regions, led Thornley et al. to prefer an ageing scenario to explain the low excitation. This explanation was further supported by modelling presented in Giveon et al. (2002). The modelling of Thornley et al. indicates a short starburst timescale of only a few O star lifetimes (106–107 yrs) and suggests that strong negative feedback regulates the star-forming activity. Disruption due to stellar winds and SNe may cease the starbursts before the depletion of gas and suggests that periodic starburst events are likely (e.g. in M82 FörsterSchreiber et al. (2001)). OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 373

Verma et al. (2003) confirmed that for a given nebular abundance, local starburst nuclei are of lower excitation than Galactic and Magellanic cloud H II regions. This was shown for the primary products of nucleosynthesis argon, sulphur and neon confirming the Thornley et al. (2000) result was not restricted to Ne only. Additionally, the observed excitation was confirmed to be metallicity dependent, with the lowest metallicity compact dwarfs showing the highest excitations, at a similar level to local H II regions. Ageing, a modified IMF and metallicity are all possible contributors to the metallicity-excitation relation. In the same study, the abundance of sulphur was observed to be lower than expected for a primary product of nucleosynthesis. Wolf–Rayet and BCD galaxies did not show such an under- abundance (Verma et al., 2003; Nollenberg et al., 2002). This under-abundance and weakness in the sulphur line strength implies that (a) the fine structure lines of neon are favoured over the those of sulphur as a star formation tracer in future spectroscopic surveys of galaxies and (b) that sulphur is an unsuitable tracer of metallicity.

4.3. WOLF–RAYET STARBURSTS

The Wolf–Rayet phase is a signature of massive star formation as the progenitors are believed to be massive young O stars (3 tsb 8 Myr, M 20M). Wolf–Rayet features are found in a wide range of galaxies such as BCDs, H II galaxies, AGN, LINERs and starbursts (see Schaerer et al., 1999 for a compilation). Recently, L´ıpari et al. (2003) identiﬁed possible Wolf–Rayet features in the spectra of some PG quasars. The identifying Wolf–Rayet signatures are in the visible region and thus are only representative of the unobscured components. Nevertheless, a correlation between the star formation indicator UIB 7.7 μm/15 μm continuum with Hα in Wolf–Rayet galaxies suggests that the infrared and visible trace, at least partly, the same components (e.g. in Haro 3 Metcalfe et al., 1996). The short-lived Wolf– Rayet phase is commonly detected in low metallicity systems with a strong ISRF such as blue compact dwarfs (BCDs, see Section 4.4). ISO observations of known Wolf–Rayet galaxies reveal differences in their infrared properties compared to non-Wolf–Rayet galaxies. In comparison to starbursts, Wolf–Rayet galaxies are of higher excitation and lower abundance contrary to expectation of a more important role of Wolf–Rayetstars at higher metallicities from stellar evolution models (Verma et al., 2003). The Wolf–Rayet galaxy NGC 5253 has a weaker excitation determined from mid-infrared ISO data than expected by stellar models (Crowther et al., 1999) implying that the UV spectrum may not be as hard as commonly thought. However, Crowther et al. (1999) show that the effect is due to a young compact starburst surrounded by an older starburst both of which lie within the SWS aperture. As a result, measured line strengths are diluted and the apparent effective temperature of the emitting region is lowered, thus reducing the measured excitation. This effect, which is fully consistent with an ageing/decaying starburst scenario suggested by 374 A. VERMA ET AL.

Thornley et al. (2000), can explain the low excitation of starbursts described in Section 4.2 (Thornley et al., 2000; Verma et al., 2003).

4.4. BLUE COMPACT DWARFS

A number of low metallicity (Z Z) blue compact (<5 kpc) dwarf galaxies were also observed by ISO. These include I Zw 36 (Mochizuki and Onaka, 2001; Nollenberg et al., 2002), Haro 11 (Bergvall et al., 2000), Haro 3 (Metcalfe et al., 1996), NGC 5253 (Crowther et al., 1999), II Zw 40 (Schaerer and Stasińska,1999; Madden, 2000), NGC 1569 (Madden, 2000; Haas et al., 2002; Lu et al., 2003; Galliano et al., 2003), NGC 1140 (Madden, 2000). Super-star clusters (SSCs) have been identified within some of them which are likely to be sites of massive star production that generate the bulk of the ionising flux (Gorjian et al., 2001; Vacca et al., 2002; Vanzi and Sauvage, 2004). These clusters can also account for a substantial fraction of the emerging mid-infrared emission (e.g. Vacca et al., 2002). The high obscuration, hot infrared emission and thermal radio spectra (measured for some) imply these are embedded young clusters. The properties of BCDs are similar to low metallicity regions in the SMC (Contursi et al., 2000). Some BCDs also exhibit Wolf–Rayet features in their nuclear regions indicative of young massive stars producing a significant amount of ionising UV photons (see Section 4.3). ISO observations (see Madden, 2000 for a review) reveal varying characteristics between BCDs although fine structure lines common to starbursts are present e.g. [O IV], [S III, IV] and [Ne II, III] with ratios that indicate high excitation (e.g. Crowther et al., 1999). The low metallicity and dust mass of BCDs in comparison to normal galaxies permit photo-ionisation on galaxy wide scales (Madden, 2000). UIBs are weak or absent due to destruction of their carriers by the intense radiation field and the mid-infrared is dominated by a strong VSG continuum. [C II]/(CO) emission (and [C II]/FIR to a lesser degree) can be greatly enhanced (×10) in comparison to normal metallicity starburst galaxies (Bergvall et al., 2000; Madden, 2000). Along with the hot dust detected in BCDs (e.g. Metcalfe et al., 1996), very cold dust components were surprisingly observed in some BCD members of the Virgo cluster (Popescu et al., 2002) and in NGC1569 (Galliano et al., 2003). In the latter, a very cold component (T = 5–7 K) contributes between 40 and 70% of the total dust mass of the galaxy. Molecular gas (CO) and infrared emission indicate that, although relatively metal poor, these dwarfs have high dust and gas densities for their optical sizes (Mochizuki and Onaka, 2001). An exceptional BCD is SBS 0335-052, one of the most metal-poor galaxies currently known (1/41Z). This dwarf harbours obscured, central SSCs (Dale et al., 2001b; Hunt et al., 2001) and shows a young thermal radio spectrum (Takeuchi et al., 2003; Hirashita and Hunt, 2004). SBS 0335-252 is unusually bright in the mid-infrared. No UIBs have been detected and the [S IV] and [Ne II] lines are weak in comparison to the strong continuum (Thuan et al., 1999). Estimates of the optical depth vary between authors, from optically thin (Dale et al., 2001b) to AV ∼ 10–30 OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 375

(Thuan et al., 1999; Plante and Sauvage, 2002). The current optical, radio and infrared measurements make it difﬁcult to present a consistent scenario for this intriguing galaxy. Further observations are needed to elucidate how such as small galaxy can produce such strong infrared emission without polluting its ISM with metals (Thuan et al., 1999).

4.5. STAR FORMATION RATES

The non-linear correlation (with large dispersion) between the far-infrared and Hα emission in normal galaxies (Sauvage and Thuan, 1992) suggests the use of the far-infrared as a direct SFR estimator has limitations (Kennicutt, 1998). At least for normal galaxies, both a contribution from the nuclear regions (Roussel et al., 2001) as well as from cold dust that is unrelated to the starburst activity may be responsible for the non-linear behaviour. At the highest infrared luminosities, such as in a sample of ULIGs, this very cold component comprises only 10% of the total far-infrared flux (Klaas et al., 2001). However, for less-luminous sources that often have more dominant very cold components (e.g. the starbursts in Calzetti et al., 2000), the contribution to the far-infrared from heating from non-ionising, older stars should be considered when SFRs are determined (Lonsdale Persson and Helou, 1987). A more basic problem is that SFR estimates based on the far- infrared flux are prone to errors arising simply from poorly determined far-infrared luminosities due to inadequate sampling or constraints on the far-infrared SEDs. This affects often the faintest most luminous sources detected in IRAS and ISO flux limited surveys where for some sources only one or two far-infrared data points are available. The mid-infrared also provides an alternative measure of the SFR. Based on ISO data for a large sample of non-AGN-hosting galaxies, Dale et al. (2001a) highlighted that, rather than the commonly used far-infrared flux (42–122 μm), the mid-infrared SED between 20 and 42 μm is most sensitive to the level of star formation activity. Additionally, the correlation of the mid-infrared flux 7/15 μm flux ratio with optical tracers of star formation such as Hα in a range of galaxies (e.g. in normal galaxies (Sauvage et al., 1996; Roussel et al., 2001), Wolf–Rayet galaxies (Metcalfe et al., 1996), mixed pair local interacting galaxies (Domingue et al., 2003, 2005), collisional ring galaxies (Charmandaris et al., 2001b), other interacting galaxies (Le Floc’h et al., 2002)) and the mid-infrared H-recombination lines, that are correlated to the Lyman continuum, shows that the mid-infrared flux can be used as a SFR indicator (Vigroux et al., 2001; FörsterSchreiber et al., 2004). FörsterSchreiber et al. (2004) provide an empirical calibration of the UIB flux between 5 and 8.5 μm range as a SFR estimator which is valid for starbursts up to the rate of M82, close to which the ISRF causes sublimation of the UIB carriers. The VSG continuum represented by the monochromatic 15 μm flux, that has been shown to correlate well with the [Ar II]6.99 μm line (FörsterSchreiber et al., 2003), can withstand harder radiation fields providing a SFR estimator for 376 A. VERMA ET AL. higher redshift, more luminous sources. A potential scatter in this correlation can be introduced by an AGN contribution to the VSG continuum; an event that may become more common in galaxies of high luminosity (Section 6.5). In this case, the mid-infrared may not be a direct indicator of the SFR unless the AGN contribution can be accurately subtracted (Laurent et al., 2000).

4.6. FIR SPECTROSCOPY AND THE [C II]DEFICIT

The far-infrared spectrum of the starburst galaxy M82 (Colbert et al., 1999) shows several fine structure lines including those that are commonly detected in active galaxies: [O I] 63,145 μm; [O III] 52,88 μm; [N II] 122 μm; [N III]57μm and [C II] 158 μm. These features are superimposed on a strong thermal continuum. [C II] 158 μm, along with [O I] 63,145 μm and [N II] 122 μm are important cooling lines of star-forming regions and have been shown to correlate with the dust continuum over 4 orders of magnitude in the Galaxy (Baluteau et al., 2003). For many galaxies, especially those with low metallicities, [C II] 158 μm is the strongest far- infrared line, corresponding to 0.1–1% of the far-infrared flux (Malhotra et al., 1997; Madden, 2000; Bergvall et al., 2000). The spatial correlation of the [C II] 158 μm emission with obscured regions and the fact that it scales linearly with UIB and CO relative strength suggests it originates in PDRs (Unger et al., 2000; Malho- tra et al., 1997, 2001; Masegosa et al., 2001). Fitting of the LWS spectrum of M82 confirms this view with 75% of the line produced in PDRs and the remainder from H II regions (Colbert et al., 1999) while, for a sample of nearby galaxies, Negishi et al. (2001) find half to originate in PDRs and half in low density ionised gas. The spectra of active galaxies range from those exhibiting strong ionic fine structure emission lines (such as M82) to molecular absorption line dominated systems (such as Arp220) (Fischer et al., 1999b; Fischer, 2000). The fine structure lines that are predominant in starburst galaxies are absent or weak in the spectra of higher luminosity (Malhotra et al., 1997; Fischer et al., 1997, 1999a,b). In particular, the [C II]158 μm/FIR continuum ratio decreases with increasing luminosity (Malhotra et al., 1997) and in ULIGs is only 10% of that measured for starbursts (Luhman et al., 1998, 2003; Fischer, 2000) and is even lower in LINERs (Sanei and Satyapal, 2002). As an important cooling line of warm atomic gas, [C II] was expected to be high for strongly star-forming galaxies with a high UV radiation field. Contrary to intuition, the absence of the fine structure lines and the strong molecular absorption detected (e.g. OH, H2O, CH, NH3 and [O I]) imply a weaker radiation field in ULIGs in comparison to nearby starbursts and AGN that is probably caused by shielding (Fischer, 2000). For sources deficient in [C II]158 μm, the [OI]63,145 μm lines appear to become the preferred means of cooling star-forming gas (Malhotra et al., 1999; Brauher and Lord, 2000). In a systematic survey of LWS spectra of 36 local galaxies, Negishi et al. (2001) found a strong decrease in the [C II]158 μm/FIR and [N II]122 μm/FIR cooling line ratios with bluer far-infrared colours. However, no such trend was seen OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 377 for [OI]63 μm/FIR implying that the oxygen lines do not suffer from the [C II] deficiency. A number of explanations have been proposed for the [C II] deficit. Tracing [C II] towards the Sgr B2 complex, which has an intense radiation field and gas density similar to active galaxies, raises the possibility that self-absorption may contribute to the observed deficiency (Cox et al., 1999). The [C II]158 μm line is seen both in emission (originating from Sgr B2) and absorption (Vastel et al., 2002) consistent with the presence of intervening (cold) clouds along the line of sight. This indicates the [C II]158 μm can be optically thick. However, the weakness of the [O I]145 μm line observed in Arp220 and Mrk231 and the large extinctions implied (AV 400 that exceed most other extinction measurements but see Haas et al. (2001) for Arp220) pose problems for solely a self-absorption explanation of the [C II]/FIR deficiency (Luhman et al., 2003). Lower efficiency of photoelectric gas heating for high-radiation fields, such as those in the PDRs of active galaxies, has been suggested (Malhotra et al., 1997), but an accurate analysis of this scenario remains to be performed. Additionally, the radiation field and gas density may play a role. Ageing bursts were discounted because of the detection of the high- excitation lines [Ne III] and [S III] and because older starbursts are unlikely to fuel such high luminosities (Fischer et al., 1999a). Increased collisional de-excitation, rather than the decrease in the photoelectric heating efficiency, was supported by Negishi et al. (2001). Luhman et al. (1998) proposed high UV flux incident on moderate density PDRs to explain the deficiency. However, in a revision of this work, Luhman et al. (2003) preferred a non-PDR contribution to the far-infrared to explain the deficiency in the [C II]/FIR ratio. The authors suggest dust-bounded photo-ionisation regions as a possible origin, since such regions contribute to the far-infrared (i.e. increased UV absorption by dust) but do not generate significant amounts of [C II]. This interpretation is corroborated by the increasing fraction of dust-absorbed photons with ionisation parameter (Voit, 1992). This scenario is also preferred in the radiative-transfer models of Arp220, although the effects of extinction cannot be ruled out (González-Alfonso et al., 2004). The [C II] deficiency affects possible future searches of [C II]158 μm line emission from high-redshift sources that have similar rest-frame far-infrared colours to local ULIGs. Moreover, if the stronger cooling line in warmer infrared galaxies is [O I]63 μm, then the recent non-detections in four z = 0.6–1.4 ULIG-quasars (Dale et al., 2004) also indicate a pessimistic view on either of these lines as high-redshift tracers of active galaxies. The [O I]63 μm line may also be self-absorbed.

4.7. GAS-TO-DUST RATIO AND STARBURST EVOLUTION

The gas-to-dust ratio is commonly used as an indicator of the fraction of gas that is yet to be turned into stars. Dust mass estimates determined from ISO SED analyses are used for this determination. The gas-to-dust ratios determined for starbursts, 378 A. VERMA ET AL.

ULIGs, and quasars, are similar or only slightly larger (∼150–400) than those determined for the Milky Way (∼150) (Calzetti et al., 2000; Haas et al., 2003). The highest gas-to-dust ratios are seen in BCD galaxies indicating young stellar evolution (Dale et al., 2001b). Low gas-to-dust ratios and low dust temperatures indicate a post-starburst system (e.g. the radio galaxy MG1019+0535 (Manning and Spinrad, 2001) and HyLIG F15307+3253 (Verma et al., 2002)).

5. Warm Dust and Hard Radiation Fields: Nuclear AGN

5.1. AGN IN MERGERS

ISO’s low spatial scale could not discern emission from the torus and circumnuclear starburst regions even in the most nearby active galaxies. As AGN emission is unresolved in ISOCAM images, in this section we discuss the environment of the AGN, i.e. the heating of the ISM of the host galaxy. The collisional ring galaxies NGC 985 and Arp118 both host Seyfert nuclei. In NGC 985, near-infrared images resolve the bulge of a small, second galaxy (whose accretion likely led to the formation of the ring) near the active bright nucleus (Appleton et al., 2002). The nuclei in both of these galaxies show a hot mid-infrared continuum that can wholly be attributed to dust heating by active nuclei. UIB emission is suppressed around the nuclei, suggesting they are absent in the vicinity of the AGN. In the case of Arp118, the Seyfert nucleus accounts for 75% of the mid-infrared emission from the galaxy. Two other LIGs, VV 114 and Arp299, show strong and hot mid-infrared continuum emission, i.e. that hot dust is present (a 230–300 K black-body describes the VSG hot continuum), indicating that embedded AGN may reside within these galaxies (Le Floc’h et al., 2002; Gallais et al., 2004). Seyfert nuclei seem to have high mid- infrared luminosities, as well as contributing signiﬁcantly to the total mid-infrared emission of the system. In three merging pairs of ULIGs within which one member contains an AGN, the AGN hosting galaxy is invariably more infrared luminous than the other (Charmandaris et al., 2002). Mirabel and Laurent (1999) note that the most luminous sources are those that exhibit both AGN and starburst activity, reﬂecting the well-known trend of increasing AGN importance with infrared luminosity (see Section 6).

5.2. WARM SEDS

ISO observations of the 12 μm IRAS galaxy sample (Rush et al., 1993), which consists predominantly of Seyfert 1s and 2s but includes starbursts and normal galaxies, displays a range of SED shapes that can be expected for these source types (see Figure 2 in Spinoglio et al. (2002)). These average SEDs show clear differences between normal galaxies, starbursts and the Seyferts in the optical to mid-infrared range, but they appear indistinguishable in the far-infrared. OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 379

By approximating the Seyfert SEDs using black-body functions, authors such as Rodr´ıguezEspinosa et al. (1996) or Antón et al. (2002) found that two components best fit their SEDs where the warmer component was associated to the active nucleus and the colder component with star-forming or extended regions. Similar associations were made for the two (or three) components implied by a novel inverse Bayesian method to analyse the SEDs of 10 Seyfert galaxies (Pérez Garc´ıa et al., 1998). Prieto et al. (2001) showed that the fluxes of high excitation coronal lines detected in the ISOSWS spectra of Seyfert galaxies, which are ascribed to excitation by the intense radiation field of the active nucleus, were found to be correlated to this mid-infrared component but not to the far-infrared component. This implies that the mid-infrared component arises from dust heated by processes related to the active nucleus.

5.3. IONISATION IN AGN AND SEYFERTS:THE BIG BLUE BUMP

A thin accretion disk that fuels a black hole displays a generic signature (quasi- thermal photons at T ∼ 105) in the extreme UV (EUV) commonly called the big blue bump (BBB). The bump is likely to extend from the UV towards the soft X-ray bump but, due to galactic and intrinsic absorption between the Lyman edge and a few hundred electron volts, the peak of the bump is not directly detectable. A novel method to reconstruct the BBB is to use the high ionisation potential infrared fine structure lines (coronal lines, 50–300 eV) that originate in the narrow-line region of Seyferts and quasars and directly probe the EUV SED. Reconstruction of the BBB was attempted using photo-ionisation modelling of the Seyferts NGC 4151 (Seyfert 1.5), NGC 1068 (Seyfert 2) and Circinus (Seyfert 2). The reconstruction is not only dependent on the ionising continuum but also on the geometry of the NLR and the ionisation parameter. It was possible to constrain the presence of this bump for these galaxies by complementing the ISO data with UV, optical and near- infrared lines. All three investigated sources (Moorwood et al., 1996; Alexander et al., 1999c, 2000) require the presence of a thermal blue bump, a single power-law continuum fails to reproduce the observed narrow-line ratios. Circinus provides the clearest example of an EUV bump that peaks at 70 eV and contains half of the AGN luminosity. The predicted SEDs of NGC 4151 and NGC 1068 display clear troughs falling sharply beyond the Lyman limit and rising sharply above 100 eV. The preferred explanation is that it is due to absorption by neutral hydrogen of 19 −2 column density NH ∼ 5 × 10 cm . In the case of NGC 4151 this absorbing gas was identified through an analysis of UV absorption lines (Kriss et al., 1995). This scenario allows for the presence of (obscured) BBB in all of the sources investigated (see also Prieto et al. (2001)). However, Binette et al. (1997) find that photoionization by a simple power-law suffices to explain Circinus, if an optically thin (density-bounded) gas component is included. In a detailed metallicity analysis, Oliva et al. (1999) find that either approach can equally explain the observations, and 380 A. VERMA ET AL. note that the ionising continuum is mostly dependent on the gas density distribution. More complex geometry and distribution of gas (e.g. 3-D multi-cloud models) and emission mechanisms (shock heating and photo-ionisation) (Contini et al., 1998, 2002; Martins et al., 2003), than an assumed single cloud with filling factor <1 (Alexander et al., 1999c), are more appropriate assumptions to model the narrow- line emission properties surrounding AGN (Martins et al., 2003). While such detailed modelling was not possible for less well-observed sources, a comparison of spectroscopic tracers, such as electron density and excitation, to photo-ionisation models enabled at least a consistency check to be performed between power-law or black-body continua. However, neither Sturm et al. (2002) nor Prieto et al. (2001) were able to distinguish the two for their sources with the range of lines they used. The metallicity, geometry and density of obscuring clouds and extinction all play a role in these analyses.

5.4. THE DUSTY TORUS AND VARIABILITY

If near- and mid-infrared emission from the inner region of AGN is thermal re-radiation by dust surrounding the central engine then, over sufficiently long timescales, variability seen in the optical/UV should be detectable in the infrared. This would provide convincing proof and constraints on the presence, temperature and distribution of dust grains from the variable source. However, ISO monitoring of variable AGN, that were contemporaneous to other multi-wavelength campaigns, failed to show any convincing thermal variability (e.g. Mrk279 (Santos-Lleó et al., 2001); BL Lac PKS 2155−304 (Bertone et al., 2000); BL Lac 2007+777 (Peng et al., 2000); OVV quasar 3C 446 (Leech et al., 2002)). Rather, the variability where detected was consistent with non-thermal components such as jets (Peng et al., 2000; Leech et al., 2002). These findings are consistent with variability traced by polarisation measurements of the OVV quasar 3C 279 at 170 μm (Klaas et al., 1999) showing that the polarisation detected was aligned to the radio axis and thus consistent with synchrotron emission. The rapid variability noted in two measurements separated by ∼1 yr confirms that the far-infrared emission is compact and originates from the core.

5.5. UNIFICATION

Under the standard uniﬁcation scheme (Antonucci, 1993), the orientation relative to the line of sight of a dusty putative torus surrounding the accretion disk of the central black hole, may connect various types of active galaxies: radio galaxies to quasars and Seyfert 1s to Seyfert 2s. A check of these schemes is offered through their infrared emission: the dust torus intercepts optical-UV photons released by the AGN and re-emits them in the infrared. At λ>30 μm the infrared emission is largely optically thin and hence isotropic. As a result, the infrared SEDs should OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 381 be similar for all members of a uniﬁed class, irrespective of the orientation of the obscuring torus to the line of sight. In this section, we discuss how ISO results have contributed to understanding this issue.

5.5.1. Radio Galaxies and Quasars Orr and Browne (1982) and Barthel (1989) proposed a scheme that unifies quasars and radio galaxies as intrinsically similar sources, with observed differences ascribed to the viewing angle with respect to the black hole. Figure 3 illustrates the manifestation of unified schemes in the infrared SEDs of quasars. A quasar with a flat radio spectrum is observed with a viewing or aspect angle of 0◦ to the line of sight (pole-on) such that a beamed jet with superluminal motion causes the bright radio and far-infrared appearance (e.g. FR 0234+28). At intermediate angles between the pole- and edge-on orientations, a quasar with a steep radio spectrum is seen. The shallower the aspect angle, the more obvious the thermal bump of the torus becomes. At large orientations (90◦ or perpendicular) to the line-of-sight, we observe the radio galaxy with the torus edge-on (e.g. Cyg A, Fornax A, Centau- rus A), whereby the nucleus that contains an “optical quasar” is obscured by the dusty torus. Powerful Fanaroff–Riley (FR) II radio galaxies are “edge-on” quasars viewed at high inclination (Haas et al., 1998a). Overwhelmingly, the SEDs of radio galaxies and radio quasars appear indistinguishable in the far-infrared (Haas et al., 1998a, 2004; van Bemmel et al., 2000; Meisenheimer et al., 2001; Andreani et al., 2002) confirming the expectation from unified schemes i.e. that the dust is

Figure 3. Illustration of the uniﬁed scheme of quasars complemented by its manifestation in the infrared SEDs of quasars (Haas et al., 1998). 382 A. VERMA ET AL. isotropic and optically thin. This is contrary to indications from IRAS data that implied radio galaxies have lower mid-far infrared emission than radio quasars. The lower redshift range probed by IRAS may account for this (Meisenheimer et al., 2001).

5.5.2. Narrow Line Radio Galaxies; Broad Line Radio Galaxies and Quasars The 3C sample of quasars selected at 178 MHz contains two types of galaxies: (1) those exhibiting broad lines: steep-spectrum quasars and broad line radio galaxies; and (2) those exhibiting narrow lines: FR II narrow line radio galaxies. Under the unification scenario, sources of identical isotropic lobe power should have the same far-infrared isotropic power. Haas et al. (2004) confirm earlier results of Meisenheimer et al. (2001) for all 3CR sources present within the ISO Data Archive,1 providing clear evidence in favour of geometric unification. The considerable dispersion in lobe power can be attributed to environment and evolution. Siebenmorgen et al. (2004a) confirm earlier work of Freudling et al. (2003) and find that the SEDs of narrow-line radio galaxies to be colder (λpeak ∼ 100 μm) than the broad-line counterparts (λpeak ∼ 40 μm). This result is fully consistent with unification schemes where, in broad-line radio galaxies, the BLR is exposed and so is hot dust that causes the SED shape to appear warmer.

5.5.3. FR I and FR II Galaxies These extended radio sources are separated by the fact that FR II galaxies show powerful edge-brightened extended double lobes, whereas FR Is have diffuse lobes that are brightest close to the nucleus. Moreover, FR Is are of lower luminosity than FR IIs. In a recent study of 3CR sources, Müller et al. (2004) found that the low radio power FR I galaxies have low mid- and far-infrared luminosities. In contrast, the FR II sources are infrared luminous. The high far-infrared-to-mid-infrared ratio and the temperature of the far-infrared peak which is comparable to that of quiescent radio-quiet elliptical galaxies, implies that heating by the ISRF of the host galaxy is probable. This, together with the finding that the dust masses of FR Is and FR IIs are similar but with low gas fractions in the nuclear regions, implies that the black holes in FR I galaxies are fed at a lower rate than FR II galaxies (Müller et al., 2004).

5.5.4. Flat Spectrum and Steep-Spectrum Quasars Flat spectrum quasars show a strong synchrotron component (e.g. FR 0234+28), smoothly concatenating the mm ﬂux to that observed in the mid-infrared (Haas et al., 1998a). A thermal far-infrared component may equal the synchrotron in strength (e.g. 3C 279), but in most cases only upper limits to a thermal contribution can be derived due to the dominating synchrotron emission. This prevalence of far-infrared synchrotron emission is consistent with Doppler boosted synchrotron emission of

1http://www.iso.vilspa.esa.es. OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 383 a relativistic jet seen almost end-on and does not contradict this uniﬁcation (Haas et al., 1998a).

5.5.5. Radio-Loud and Radio-Quiet Quasars A complete range of radio-loud, radio-intermediate and radio-quiet quasars was analysed by Polletta et al. (2000). Except for ﬂat radio sources, the mid- and far- infrared emission are of thermal origin. The far-infrared emission is predominantly explained by heating from the AGN itself and, while starburst components are present, they contribute to the infrared emission at different levels but always less than the AGN ( 27%).

5.5.6. Seyfert 1s and Seyfert 2s Within the unified scheme, Seyfert 1s and Seyfert 2s differ only in viewing angle of the torus obscuring the central black hole. In Seyfert 1s the broad line region is exposed, whereas in Seyfert 2s it is hidden from view due to obscuration by the torus. The central regions where dust is heated to high temperatures gives rise to a mid-infrared peak in the continuum emission. Aspect-oriented unification is consistent with the result that the mid-infrared dust temperatures of Seyfert 1s are higher than in Seyfert 2s, where the torus is partially optically thick up to the mid- infrared (PérezGarc´ıaand Rodr´ıguezEspinosa, 1998). Spectroscopic tracers such as high-excitation lines and broad lines are indicative of the presence and orientation of the putative torus. In a large sample of Seyferts, Clavel et al. (2000) showed that UIB equivalent widths are larger in Seyfert 2s than in Seyfert 1s. An interpretation of this result is that UIB features originate from extended regions and hence emit isotropically, while the AGN continuum is absorbed by the torus and must emit anisotropically (PérezGarc´ıaand Rodr´ıguezEspinosa, 2001), since it is dependent on the orientation of the torus to the line of sight. This interpretation thus provides direct proof of the unification of the Seyfert classes. A test of unification schemes is offered via observations in hard X-rays that can penetrate an obscuring (Compton-thin) torus. Thus, in the simplest unification models, if X-ray emission is not beamed, the mid-infrared continuum arising directly from AGN heating of the torus or dust in the narrow-line region is determined by the geometric distribution of the obscuring material and the aspect angle. After isolating the mid-infrared continuum due to the active nucleus alone, Lutz et al. (2004a) found no distinction between Seyfert 1s and Seyfert 2s with the X-ray-to- mid-infrared flux ratio, contrary to the predictions of unified schemes. The authors suggest the apparent uniformity is due to the dilution of anisotropic emission by a predominant mid-infrared component on extended scales that emits isotropically. Another result that is difficult to reconcile with the torus model arises from ISO SED analysis of a sample of hard X-ray selected sources comprising both type 1 and 2 Seyfert galaxies. Kuraszkiewicz et al. (2003) showed that while Seyfert 2s are redder in their optical colours than Seyfert 1s, their mid- to far-infrared emission, as traced by the 25–60 μm colour, remains similar. More complex geometries for the 384 A. VERMA ET AL. obscuring material are required to explain both the difference in optical colours and the similarity in the infrared. Clumpy and extended models are suggested where a disk of only a few hundred parsecs can reproduce the IR SED without invoking a starburst component (see Elitzur et al., 2003 and Section 6.6).

5.6. EVOLUTION OF COMPACT STEEP-SPECTRUM RADIO GALAXIES

ISO data contributed towards understanding whether the compactness of steep- spectrum radio galaxies was due to age or evolution of environmental conditions that inhibits the growth of the radio lobes. It was expected that dense material causing this inhibition should be accompanied by far-infrared emitting dust. However, Fanti et al. (2000) found no far-infrared excess in compact steep-spectrum radio galaxies, in comparison to a control sample of 16 extended radio galaxies, which favours a youthful scenario to explain their compactness.

5.7. CONSTRAINTS FROM LINE PROFILES

5.7.1. Probing the NLR The emission line profiles of Seyfert galaxies are known to exhibit blue asymmetries and blue-shifts with respect to the systemic velocity in their optical spectra. The profiles provide constraints on the NLR. Narrow lines detected in the mid-infrared have the advantage of being much less sensitive to extinction than optical lines and therefore provide a relatively unobscured view to the NLR. The similarity in line shapes between the optical and the mid-infrared provides constraints on the nature of the NLR and models how such asymmetries were produced. Sturm et al. (1999) are able to rule out simple radial flows plus dust scenarios and they propose that a central geometrically thin, but optically thick, obscuring screen located close to the nucleus is present in NGC 4151 to explain the observed asymmetry. The mid-infrared line profiles of NGC 1068 are due to a combination of an highly ionised outflow and extended emission of lower excitation, and remaining blue shifts attributed to an intrinsic asymmetry in the NLR or very high column density in the line of sight (Lutz et al., 2000b). For a larger sample, Sturm et al. (2002) show that the differences in line profiles depend on the ionisation potential. Circinus shows symmetric low-ionisation and asymmetric high-ionisation lines, while the spectrum of NGC 7582 shows the reverse. These differences suggest that the NLR properties affecting individual galaxies are particular to that galaxy. Differences including extinction, asymmetric distributions of NLR clouds, orientation must all influence the line shapes (Sturm et al., 2002).

5.7.2. Probing the BLR Lutz et al. (2000a) searched for broad components in the H-recombination lines of NGC 1068, since the detection would imply that at least part of the BLR is accessible OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 385 in the mid-infrared (looking through the obscuring putative torus). The lack of detection places a lower limit in both extinction (AV > 50 mag) and hydrogen 23 −2 column density (NH > 10 cm ). Although little can be stated regarding favoured torus models or constraining gas properties based on this information, Brα was identiﬁed as being the preferred indication for future studies and Lutz et al. (2002) found broad line components in approximately one-fourth of their sample of Seyfert galaxies from ground-based high-resolution spectroscopy.

6. Ambiguous Sources: Ultra- and Hyperluminous Infrared Galaxies

As we have previously discussed, ULIG samples contain galaxies of both starburst and AGN type. Establishing whether these types are linked in evolution or nature is a key issue in the study of active galaxies raising several questions: what is the role of star formation and black holes; which is/are present; do super-massive black holes and starburst activity concurrently exist in all ULIGs and how much do they contribute to the bolometric luminosity? Prior to ISO, few observational tools were available to probe through the dust obscuration to reach the source of the infrared power. ISO’s spectroscopic and photometric capabilities enabled signiﬁcant advances in this ﬁeld, largely elucidating the previously termed ‘starburst-AGN controversy’ which is discussed in this section.

6.1. QUANTITATIVE SPECTROSCOPY

6.1.1. Fine Structure Line Diagnostics Fine structure lines directly trace the excitation states of the ISM from which they originate. The radiation fields generated by AGN/QSO greatly exceed that of starbursts and thus can excite interstellar gas to higher ionisation species (as described in Section 2.2.2). The difference in excitation between AGN and starbursts is the basis of fine structure line diagnostic diagrams. Such excitation diagrams are well established in the optical (Veilleux and Osterbrock, 1987; Kewley et al., 2001), for which Sturm et al. (2002) constructed the first infrared analogues (Figure 4) that successfully distinguish starbursts from AGN. Diagnostic diagrams from Sturm et al. (2002) are shown in Figure 4 both involving the fine structure line [O IV]25.91 μm. While this line was expected in the high ionisation potential of AGN, it was also often weakly detected in star-forming galaxies and ULIGs (Lutz et al., 1998a). In our galaxy, the [O IV]25.91 μm line has not been observed in H II regions surrounding young, hot, massive stars. However, its detection in ISO spectra of starbursts and the fact that it was spatially resolved in the low excitation starburst M82 implies it is plausibly produced by ionising shocks or extremely hot ionising stars (Lutz et al., 1998a; Schaerer and Stasińska,1999). Its diagnostic capability thus arises from the fact that the measured strength of the line in starburst galaxy spectra is at least 10 times fainter than those measured in 386 A. VERMA ET AL.

Figure 4. Diagnostic diagrams based on fine structure lines from Sturm et al. (2002). (a) [Ne VI]/[O IV] vs. [Ne VI]/[Ne II]: (diamonds) Seyfert 1s, (asterisks) Seyfert 2s: the composite Seyfert galaxies that exhibit UIBs are encircled and are mostly separated from the remaining galaxies. (b) [O IV]26 μm/Brβ vs. [Si II]34 μm/Brβ: (diamonds) AGN, (asterisks) starbursts: the starbursts occupy a different regions of the diagnostic diagram than the AGN. spectra of AGN. This absolute difference in the strength of [O IV] relative to a low excitation starburst line such as [Ne II] provides a straightforward indicator of AGN versus starburst activity (Sturm et al., 2002). Combining the high:high excitation fine structure line ratio [Ne VI]/[O IV] with the high:low excitation fine structure line ratio [Ne VI]/[Ne II], provides a basic diagnostic of composite sources (Figure 4a). By using the limiting values of the high:low excitation ratio [O IV]/[Ne II] (always <0.01 in starbursts and 0.1 < AGN < 1.0) Sturm et al. (2002) were able to estimate the fractional contribution of an AGN to the bolometric luminosity through a simple mixing model. Figure 4b shows OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 387 how the high- and low-ionisation indicators normalised to the H-recombination line Brβ (an indicator of the Lyman continuum rate) provide an analogue to the Veilleux and Osterbrock (1987) diagnostics that are insensitive to obscuration and will be extremely useful for classifying ULIGs observed with future infrared spectrographs. An equivalent diagnostic at longer wavelengths can be constructed using [C II]158 μm/[O I]63 μm versus [O III]88 μm/[O I]63 μm that can separate starburst, AGN and PDR contributions (Spinoglio et al., 2003). As the spectroscopic sensitivity of ISO permitted the detection of H- recombination lines or useful limits on the [O IV] emission in only the brightest sources, alternative tools using more commonly detected spectral features (e.g. UIBs and bright fine structure lines) were used during earlier ISO science analysis. Genzel et al. (1998) combined excitation and UIB strength to asses the nature of a large sample of AGN, starbursts and ULIGs. Laurent et al. (2000) constructed a diagnostic based on the mid-infrared continuum shape to UIB strength (Section 6.1.2). However, the use of UIBs is limited and presents problems in interpretation (Section 6.1.3).

6.1.2. UIB-Based Diagnostics The predominance of UIBs in starbursts (and their deficiency in AGN, or low metallicity environments) naturally lends the ratio of UIB feature flux (normalised to the mid-infrared continuum) to be a suitable tracer of star formation activity (Verstraete et al., 1996; FörsterSchreiber et al., 2004; Peeters et al., 2004). For normal and starburst galaxies, the UIB-to-continuum ratio is seen to be high in the nuclear star-forming regions and to decrease with radial distance (Mattila et al., 1999; Dale et al., 2000), indicating a weakening ISRF. Conversely, a strong enhancement in the continuum to line ratio is seen in the vicinity of AGN, where the continuum is strong and UIBs weak as the features suffer dilution due to the absolute high level of the continuum and/or destruction of their carriers (Lutz et al., 1998b). The effectiveness of this indicator is clearly seen in Centaurus A, where spatially resolved ISOCAM-CVF observations of the AGN-dominated nucleus and the extended starburst region (Mirabel et al., 1999) show striking differences in their mid-infrared emission. Together with the absence of UIB emission features, the spectral index of the continuum is steeper in AGN than in the surrounding circumnuclear star-forming regions. These spectral differences enable, at least in principle, discrimination between the emission characteristics of the two distinct physical mechanisms. Genzel et al. (1998) demonstrated the effectiveness of the 7.7 μm UIB-to- continuum ratio versus a tracer of hardness (such as the ratio of the high excitation potential fine structure line [O IV] to the low-excitation line [Ne II] (or [S III])) to classify ambiguous sources. This diagram (Figure 5a) enabled the classification of ULIGs for which the underlying fuelling mechanisms were unclear in the pre-ISO era. The separation between AGN, starbursts and the intervening ULIGs is unequivocally depicted. The well-known low-excitation starburst M82 is 388 A. VERMA ET AL.

Figure 5. UIB diagnostics from Genzel et al. (1998) (a) and Laurent et al. (2000) (b). Both of these diagrams show the mixing levels of starbursts and AGN in active sources. Finally, the MIR/FIR diagnostic of Peeters et al. (2004) is shown for comparison (c). OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 389 far separated from the AGN NGC4151. These extremities more or less define an appropriate mixing line that begins with 0% AGN in the lower right corner to 100% in the upper right. Thus, placing a ULIG on this plane provides an indication of its AGN fraction. The line of equivalent mixing is marked on the diagram, 4 of the 13 ULIGs shown fall to the left of that line, implying a bolometrically dominant AGN is needed to explain their emission. Similarly, using the ratios of 6.2 μm UIB-to-continuum versus warm-to-hot continuum, Laurent et al. (2000) constructed a diagnostic diagram that can quantitatively estimate the fractional contribution from AGN, H II regions and PDRs. Laurent et al. (2000) showed how a mid-infrared spectrum may be quantitatively decomposed into these three components through spectral template fitting (see Fig- ure 5b), using the pure H II spectrum of M17 to fit the strong VSG continuum in starbursts (Cesarsky et al., 1996b), an isolated PDR spectrum of the reflection nebula NGC 7023 (Cesarsky et al., 1996a) and the nuclear spectrum of Centaurus A representing the hot AGN spectrum (Laurent et al., 2000). Such diagnostics are extremely valuable for assessing the fuelling mechanisms of ambiguous sources such as ULIGs (Laurent et al., 2000; Tran et al., 2001) and can reveal the presence of obscured AGN in the mid-infrared. This diagnostic tool is most successful when used with spectra of high signal-to-noise ratio and with sufficient coverage of the mid-infrared range (i.e. to best determine the continua shortward (<5 μm) and longward (>12 μm) of the family of strong UIB features) that enables accurate decomposition of the H II PDR and AGN templates and reduces problems associated with blended mid-infrared absorption and emission lines (Spoon et al., 2002). A follow up of this method, on a larger sample of galaxies was presented by Peeters et al. (2004) with similar results. The authors also construct a new mid- and far-infrared combined diagnostic that is useful for differentiating obscured sources (Figure 5c). In integrated galaxy spectra, where no spatial information is available, e.g. for compact and/or distant sources, decomposition of the spectral characteristics is the only means to identify the power sources constituting the mid-infrared emission (e.g. Vigroux, 1997; Rigopoulou et al., 1999; Lutz et al., 1998b; Mouri et al., 1998; Genzel et al., 1998; Laurent et al., 2000).

6.1.3. Limitations The use of the UIB strength as a diagnostic has several limitations. One originates from the fact that measuring the strength of the 7.7 μm UIB feature in the wavelength range of ISOPHOT-S (2–11 μm) was challenging as ULIGs also display strong silicate absorption at 9.7 μm and tracing the continuum was often extinction dependent. Another limitation of relying on UIB strength originates from its sensitivity to the radiation environment. Its absence in hard radiation fields in the vicinity of massive stars or AGN weakens its effectiveness as a tracer of star formation. Fi- nally, as discussed by Laurent et al. (1999, 2000), ISO’s limited spatial resolution in the mid-infrared, may introduce a scatter into the diagnostic diagram, especially for distant sources with angular size smaller than the instrument beams. This results 390 A. VERMA ET AL. in a dilution of the AGN signatures as a larger fraction of the circumnuclear star- forming regions of the host galaxy disk enter into the beam. For these reasons, while UIB diagnostics were relatively successful in particular for faint sources. The diagnostics formed purely from fine structure lines, which directly trace differing excitation regions within galaxies, offer far cleaner and more effective means of source-type segregation and should be preferred for classification of faint and high-redshift galaxies that will be observed by future telescopes.

6.1.4. Broad-Band Colour Diagnostics The well-known diagnostics based on the broad band infrared colours determined from IRAS data (e.g. the IRAS 60 μm/25 μm colour represented the warm-to-cold emission ratio) have been extended to include ISO photometry (Spinoglio et al., 2002) and show the separation of ‘warm’ AGN (Section 5.2), ‘cold’ starburst (Sec- tion 4) and ‘very cold’ (Section 3) dust-dominated systems (see their Figure 14). While in the far-infrared, the SEDs of Seyfert 1s and Seyfert 2s are indistinguishable, the diagram of Spinoglio et al. (2002) is based on the fact that the SEDs of Seyfert 1s are ﬂatter in the mid- to far-infrared than Seyfert 2s’. Using averaged spectra from this sample, the authors were able to produce redshift tracks on colour–colour diagrams for the Spitzer MIPS and IRAC ﬁlters as a means of determining infrared photometric redshifts. Even though these tracks were calculated without evolutionary considerations, the diagrams are useful tools to differentiate AGN (Seyfert 1 and Seyfert 2s, QSOs) dominated sources from starbursts.

6.2. THE NATURE OF ULIGS

Using these techniques, the majority of ULIGs are found to be predominantly powered by starbursts. For large samples of ULIGs Lutz et al. (1998b), Rigopoulou et al. (1999), and Tran et al. (2001) corroborate the early result of Genzel et al. (1998) and find that ∼75–85% of their sample do not contain energetically dominant AGN. Starburst activity can explain the bulk of the bolometric luminosity, although the presence of a weak or heavily obscured AGN cannot be excluded. While the fraction of AGN-dominated sources is 15–25%, the fraction of AGN hosting ULIGs is 20–25%. The AGN-dominated sources have ‘warm’ IRAS colours (S25/S60 0.2), confirming previous IRAS results on ULIGs with optical or near-infrared spectroscopic classification. Moreover, a comparison of the mid-infrared classification with good quality optical spectroscopy is remarkably consistent providing both optically classified H II regions and LINERs are both ascribed to the mid-infrared starburst class (Lutz et al., 1999).

6.3. THE NATURE OF LINERS

The nuclei of a large fraction (as much as one-third) of local galaxies display spectra with weak emission lines of lower ionisation than is seen in typical AGN, e.g. OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 391

Seyfert nuclei. Furthermore, the low-ionisation species present can not be explained through photo-ionisation by normal stars. The population of low ionisation nuclear emission-line regions – LINERs (Heckman, 1980) is well known to be mixed. Because of their prevalence, establishing the fraction of AGN or star-forming region like LINERs has strong implications on our understanding of galaxies. ISO spectroscopic results clearly indicate that, at least in the infrared, their properties are more similar to starbursts (Lutz et al., 1999). In this case, the low ionisation emission lines are ascribed to ionising shocks related to starburst-driven winds (Lutz et al., 1999; Sugai and Malkan, 2000). Interestingly, though, in an infrared study of LINERs, Satyapal et al. (2004) found as much as two-thirds of the sample with hard X-ray counterparts have compact X-ray morphologies consistent with those found for AGN. The remainder with scattered X-ray emission are brighter in the infrared than the compact ones. The authors found that most LINERs displaying high-excitation mid-infrared lines also show compact X-ray morphology, even though a few systems deviate from this (i.e. NGC 404 and NGC 6240). These studies indicate that infrared bright LINERs are most likely members of a shock- dominated population, whereas optical selection may trace the low luminosity AGN component. The fact that LINERs occupy a locus in ionisation diagrams situated between classical starbursts and AGN suggests an evolutionary connection whereby LINERs may represent a phase when the AGN is deeply buried and the unrelated starburst dominates the detected emission (Lutz et al., 1999).

6.4. ULIG SEDS

An alternative classification method arises from an analysis of ULIG SEDs that include ISO data. This includes fainter and/or more luminous ULIGs and HyLIGs that were inaccessible to ISO’s spectrometers. Often in conjunction with multiwavelength data from other telescopes, ISO data has enabled classification of ULIGs through decomposition of their SEDs. A statistical study of 20 ULIGs (Klaas et al., 2001) clearly showed a progression of SED shapes. These were broadly divided into two types. The first (type A) show a power-law increase from 1 to 50 μm and are classified as Seyferts from optical spectroscopy. They also have warm infrared colours ( f25 μm/f60 μm > 0.2). Type B sources display a relatively flat SED from1to10μm, that is followed by a step rise towards the peak of emission in the far-infrared. These SED characteristics are displayed by ULIGs with optical spectroscopic classifications of Seyfert 2, LINER, H II/starburst. These are ‘cool’ ULIGs with infrared colour f25 μm/f60 μm < 0.2. The grouping of LINERs in the ‘cool’ starburst-like ULIG class is consistent with the analysis of mid-infrared versus optical spectroscopic classification performed by Lutz et al. (1999) discussed in the previous section. This, together with the near-infrared colours (see Figure 6 in Klaas et al. (2001)) of ULIGs, suggest that the far-infrared emission from LINER ULIGs is of starburst origin. Klaas et al. (2001) interpreted the emission of ‘warm’ ULIGs (type A) to arise from dust heated to high temperatures 392 A. VERMA ET AL. directly from the central AGN, which in turn dominates the mid-infrared emission. The second SED feature is ascribed due to thermal re-radiation by dust heated by starburst activity. The indistinguishable far-infrared SEDs of AGN- and starburst-dominated ULIGs suggests that the far-infrared emission largely comes from less active or shielded regions that are generally not heated by the AGN (Klaas et al., 2001). Thus, Klaas et al. (2001) proposed a three-stage dust model consisting of hot AGN-heated dust, warm starburst-heated dust (50 K > T > 30 K) and cold, pre-starburst or cirrus-like dust (30 K > T > 10 K). The general consen- sus from ISO spectroscopic results of warm ULIGs being AGN dominated and cold ULIGs starburst dominated is corroborated by the findings of Klaas et al. (2001).

6.5. HyLIGs AND THE INCREASING ROLE OF AGN AT HIGH LUMINOSITY

The increasing importance of AGN with increasing luminosity (Shier et al., 1996; Veilleux et al., 1997) has been confirmed by ISO observations of ULIGs (Lutz et al., 1998b; Tran et al., 2001; Charmandaris et al., 2002). Tran et al. (2001) quantified that the change from starburst- to AGN-dominated systems occurs at ∼ . 12 log10(LIR) 12 4–12.5L. While ISO results show that 10 L can be powered by star formation alone (Genzel et al., 1998; Lutz et al., 1998b; Rigopoulou et al., 1999; 13 Tran et al., 2001), beyond 10 L an AGN contribution seems inevitable. Moreover, the production of such a luminosity would require star formation rates greater than 3 −1 or approximately few ×10 M yr . Such rates require huge concentrations of molecular gas to be present (most likely the result of mergers) and highly ineffective negative feedback (Takagi et al., 2003). The prevalence of AGN in HyLIG samples partially arises from a selection effect. These rare luminous sources have been discovered with heterogeneous selection methods, most commonly through a correlation of known quasar or radio catalogues with IRAS or ISO surveys resulting in the known population being biased towards AGN. Constraining the properties of this sample and full consideration of the biases involved is unfortunately not possible due to the paucity of known HyLIGs. Nevertheless, ISO observed several hyperluminous galaxies in a range of programs. Only a few HyLIGs were sufficiently bright to permit low resolution spectroscopic observations. Analyses of IRAS 09104+4109 (Tran et al., 2001) and IRAS F15307+3252 (Aussel et al., 1998) showed that the mid-infrared emission predominantly originates from very compact regions <100 pc (Taniguchi et al., 1997) with lower limits√ on the size of the emission region of 6 pc for F15307 (Aussel et al., 1998) and 5/ m (m: magnification) for the Cloverleaf. These quasars show hot gas at temperatures in excess of 400 K. They exhibit low UIB-to-continuum ratios expected for AGN. For F15307 SED modelling showed that a starburst component is required to fit this source fuelling 45% of its luminosity, although this poses a problem for the non-detection of CO (Verma et al., 2002). In a sample of OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 393 four HyLIGs including HyLIGs without clear AGN features, a SED analysis indicated an AGN was required to explain the IRAS-ISO SEDs (Verma et al., 2002). One of the highest redshift sources observed by ISO was the powerful radio galaxy 8C 1435+635 (z = 4.25). It was shown to be hyperluminous and massive but the infrared-mm emission was solely ascribed to star formation and extremely high star formation rates were implied (Ivison et al., 1998). Several HyLIGs are also well known PG quasars that do not conform to being dust-free naked quasars (Haas et al., 2000b). Haas et al. (2003) show a schematic view of the evolution of SEDs encompassing various types. A significant AGN contribution is seen, but PG quasars lack the cooler extended dust component seen in lower luminosity sources. Patchy or clumpy disk torus models may explain all of the far-infrared emission, without invoking starburst contributions.

6.6. STARBURST COMPONENTS IN AGN-HOSTING GALAXIES

Is concurrent star formation a requisite to explain the SEDs of QSOs displaying a far-infrared excess? Even the cold far-infrared-sub-mm emission of AGN- dominated sources may be fuelled by the central obscured active nuclei itself. While ISO data cannot spatially resolve such emission, invoking strong star formation is not always necessary to describe the SEDs of AGN and ultraluminous quasars (see later). A clumpy distribution of obscuring material around the black hole allows high-energy photons to escape and heat extended dust (Nenkova et al., 2002; Elitzur et al., 2003). Alternatively warped, thick and extended disks (on par- sec scales) have been proposed to explain the emission (e.g. Granato and Danese, 1994). There is clear evidence that some AGN do not host bolometrically important starbursts. For example, the archetypal Seyfert 2 galaxy NGC 5252 shows no signs for starburst activity and the bulk of its far-infrared SED can be explained by the heating of the cold dust by the general ISRF, while a small fraction is ascribed to a very large disk (kpc scale) surrounding the active nucleus (Prieto and Acosta- Pulido, 2003). The infrared emission of narrow-line Seyfert 1 galaxies measured by ISOPHOT can be accounted for by re-radiation of the soft X-ray component (Polletta and Courvoisier, 1999a,b). Similarly, at higher AGN power, starburst activity is not a necessary requisite to explain the full SEDs of a large sample of 3CR sources (Siebenmorgen et al., 2004a) or PG quasars (Haas et al., 2000b, 2003) for which AGN heating plus extended cold dust sufﬁces (also see Rocca-Volmerange and Remazeilles, 2005). In a sample of 22 quasars, starburst emission does contribute to the far-infrared emission but to a level no more than 27% (Polletta et al., 2000). Thus, while concurrent starbursts may be present, ISO SEDs of optically classi- ﬁed AGN and quasars indicate that they are typically not bolometrically dominant. The detection of spatially resolved starburst tracers in the host galaxies of quasars or AGN-like ULIGs and HyLIGs would clarify this issue. 394 A. VERMA ET AL.

6.7. AGN COMPONENTS IN STARBURST-DOMINATED GALAXIES

Infrared diagnostics were largely successful in assessing the nature of most infrared sources. However, how strong a limit can one place on the absence of an AGN in the mid-infrared? At short mid-infrared wavelengths <10 μm the strong VSG continuum related to the star formation in starburst-dominated sources will overshadow the weaker AGN continuum, particularly in Seyfert 2s where the continuum due to the AGN will suffer high obscuration (Laurent et al., 2000; Schulz et al., 2000). This results in an AGN being identiﬁed through the Laurent et al. (2000) diagnostic only if it dominates over star formation. Galaxies such as IRAS 00183-7111 and Arp220 are highly obscured even in the mid-infrared (Sturm et al., 1996; Tran et al., 2001; Spoon et al., 2004), Arp220 may be optically thick even in the far-infrared (Fischer et al., 1997; Haas et al., 2001; Gonz´alez-Alfonso et al., 2004)). This substantial obscuration by dust may cause any AGN to be completely hidden (Schulz et al., 2000), a fact which is compounded by the possibility that AGN tracers could in some cases be completely overshadowed by circumnuclear activity (Laurent et al., 1999; Dennefeld et al., 2003). Without spatial resolution, the possibility of detecting extremely obscured or weak AGN in, at least the mid-infrared, appears to be low (Laurent et al., 2000; Schulz et al., 2000).

6.8. NEWLY DETECTED ULIGS

As well as the ULIGs and HyLIGs described here which were known prior to the launch of ISO, several were discovered primarily in ISO surveys. Among the many ULIGs detected in the deep ISO surveys, a hyperluminous QSO (Morel et al., 2001) was discovered in ELAIS, and a highly extincted (AV ∼ 1000) gas and dust- rich elliptical was detected in the ISOPHOT Serendipitous Survey (Krause et al., 2003). One might expect the population of heavily obscured extremely red objects (EROs) to be luminous in the infrared. Indeed, ultraluminous emission (and high star formation rates) was determined from the ISO SEDs of the ERO HR10 with a SED similar to that of Arp220 (Elbaz et al., 2002) and a sub-mJy radio ERO from the Phoenix survey (where coeval star formation and AGN are probable (Afonso et al., 2001)) through targeted observations. Tentative infrared detections of a gamma ray burst were also reported to be at ULIG luminosities (Hanlon et al., 1999, 2000) originating from either the host galaxy or the GRB infrared afterglow.

6.9. EVOLUTION OF ULIGS

The basic similarity in luminosity and space density of ULIGs and quasars, and the fact that unbiased samples of ULIGs contained both star-forming and AGN members, led to the suggestions that the two populations are linked through evolution (Sanders et al., 1988, also see Sanders and Mirabel, 1996 and references OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 395 therein). The example of the composite galaxy NGC 985 is used to discuss possible ULIG formation and destruction scenarios (Appleton et al., 2002). Some popular evolutionary models include the following that were recently well summarised by L´ıpari et al. (2003):

(1) merger − giant shocks − super-starbursts + galactic winds − elliptical galaxies (Joseph, 1999), (2) merger − H2 inflow (starbursts) − cold ULIRGs − warm ULIRGs + QSOs − optical QSOs (Sanders, 1999), (3) merger/s − extreme starburst + galactic wind (inflow + outflow) − Fe II/BAL composite IR-QSOs − standard QSOs and ellipticals (L´ıpari et al., 2003).

6.9.1. Merger Stage and Luminosity Merging activity that initiates associated starbursts leading to the ultraluminous behaviour in the early stages is common to all of the scenarios listed above. Thus, the merger stage should be related to the luminosity evolution of ULIGs. Furthermore, if quasars are the end product of ULIG evolution and the merger stage indicates an evolutionary status, then more advanced mergers should appear more AGN-like. The near- to mid-infrared energy distributions of ULIGs are dominated by compact nuclear regions (Charmandaris et al., 2002). This confirms results from high- resolution ground-based observations (Soifer et al., 2000) that show 30–100% of the mid-infrared flux can be accounted for by the central 100–300 pc. Based on ISO imaging of a merging sequence of active galaxies, including starbursts, LIGs and ULIGs, Charmandaris et al. (2001a) suggest an evolutionary sequence that is traced by the 15/7 μm colour that monotonically increases with star formation activity, from 1 (quiescent) to 5 (extreme), as galaxies evolve from pre-starburst to the merging starburst phase. The ratio decreases again to 1 for the post-starburst phase in evolved merger remnants (cf. local less luminous post-starburst colours of NGC 7714 (O’Halloran et al., 2000) and NGC 1741 (O’Halloran et al., 2002)). The fact that the extra-nuclear activity typically contributes only a minor fraction to the bolometric luminosity of ULIGs (in comparison to LIGs (Hwang et al., 1999)) implies that simple scaling in mass and luminosity does not link the two populations, but evolutionary differences (such as merger stage, or presence of an AGN) could play a role. It is unclear whether this hardening of the mid-infrared continuum is due to an AGN contribution. The amount of molecular gas may also affect the spectral feature differences seen between the populations, such as the presence of ice features (Spoon et al., 2002). Increasing compactness of sources with increasing merger stage shown by Hwang et al. (1999) implies that compact ULIGs are in a more advanced merger state than LIGs. However, based on a comparison of the NIR morphologies of merging ULIGs and infrared spectroscopic tracers, Lutz et al. (1998b) and Rigopoulou et al. (1999) do not find any relation of increasing AGN-like activity with merger stage. In addition, ULIGs appear to not show a trend of increasing luminosity with merger 396 A. VERMA ET AL. stage again contrary to expectations of these evolutionary scenarios (Rigopoulou et al., 1999). Indeed, Rigopoulou et al. (1999) suggest that early merger stage ULIGs may be exhibiting maximal luminosities (also see Murphy et al., 2001). Moreover, the mass of the warm, cold and very cold dust components found in ULIG SEDs show no relation to merger stage (Klaas et al., 2001). The evolution or merger stage may not be traced well simply by nuclear separation alone as a merger stage indicator. More complex differences in the mass distribution and dynamics of ULIGs more likely influence the luminosity and appearance of AGN.

6.9.2. Dust Evolution in Quasars If ULIGs are related to quasars as the dust-obscured members of that class, understanding the dust properties of quasars also serve to test the evolutionary connec- tions between quasars and ULIGs. In the Sanders et al. (1988) scenario a quasar is preceded by a dusty ULIG-phase and as a result it is unlikely that the obscuring dust will suddenly disappear, some traces in the infrared should remain. Haas et al. (2003) investigated a sample of 64 PG quasars and identified a range of SED shapes that could mark a transition between states. The SED analysis showed that the central regions of PG quasars suffer low extinction (AV < 0.3) with an optical slope αopt that is independent of infrared properties like the near- to mid-infrared slope αIR. Therefore, with regard to the unification schemes, a nearly face-on view onto the PG quasars was assumed. The observed variety of SEDs were grouped into physically meaningful classes that reflect the amount and distribution of the reprocessing dust around the AGN according to evolution of the quasar (Haas et al., 2003). Their new evolutionary scenario begins with dissipative cloud–cloud collisions and angular momentum constraints that create the organisation of dust clouds into a torus/disk-like configuration. Starburst activity is initiated giving rise to a far-infrared bump, but it becomes increasingly overpowered by the central AGN which results in increasing mid-infrared emission and steepening of the mid- to far-infrared spectral slope. As the AGN continues to grow, the mid-infrared remains high and the mid- to far-infrared slope decreases until the black hole begins to starve and a a decline in the mid-infrared and far-infrared emission is observed. Figure 12 in Haas et al. (2003) depicts this scheme, the SEDs are arranged along the expectations for such an evolutionary scheme.

6.9.3. Relation of Broad Absorption Line Quasars to ULIGs Broad absorption line quasars are a rare class of the optically selected quasar population (∼12%) but have a higher frequency in far-infrared luminous, strong Fe II and weak [O III] emission line selected samples. Both orientation and age arguments are consistent with the origin of the broad absorption features with no clear resolution which is the most likely cause. Elvis (2000) suggest the broad absorption arises from viewing a single funnel-shaped outflow of absorbing material directly down the flow. It has also been proposed that these sources are quasars viewed at an orientation directly along the radial surface of the torus, where winds along the surface OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 397 give rise to the observed absorption. Alternatively, they may be a young phase in quasar evolution. The latter together with the high occurrence of BALQSO in IRAS selected samples suggests that they may be a transition stage between ULIGs and QSOs (Voit et al., 1993; L´ıpari, 1994; Egami et al., 1996; Canalizo and Stockton, 2001; L´ıpari et al., 2003). In this scenario, very young QSOs are ejecting their gaseous envelopes at very high velocity close to the initiation of the active phase of the black hole (Hazard et al., 1984). Two BALQSOs were serendipitously discovered in ISO surveys: the FeLoBAL2 ISO J005645.1-273816 was discovered in an ISOCAM distant cluster survey (Duc et al., 2002). 00 37 14.3 −42 34 55.6 was discovered in ELAIS (Alexander et al., 2001). For the former, high absorption and hot dust emission are implied by the high mid-infrared-to-UV ratio, which may indicate the mid-infrared is an efficient means to detect LoBALQSOs (Duc et al., 2002). The sources 00 37 14.3 −42 34 55.6 was the only BALQSO then known to have both infrared and hard X- ray detections. While its hard X-ray-to-MIR flux ratio is similar to high-redshift quasars, the ratio differs from low-redshift BALQSOs. Alexander et al. attribute this to absorption. As a consequence, they suggest that absorbed BALQSOs at high redshift should be easily detected in the hard X-ray bands. In addition, the only radio-loud BALQSO known, 1556+3517, was observed to investigate the relation between radio-loud quasars and the BALQSO phenomenon. Mid-infrared detections implied 4M of dust lie along the line of sight which most likely arises from the molecular torus, the NLR, or a recent starburst rather than from the BALQSO wind itself (Clavel, 1998). ISO data confirms that even in this diverse, small sample, the BALQSO phenomenon and obscuration (i.e. infrared emission) are connected. L´ıpari et al. (2003) speculate on an evolutionary connection between BALQSOs and ULIGs.

6.10. SUMMARY

Evaluating which of the three scenarios mentioned earlier in this section is most realistic will continue to be addressed by future missions. However, the ISO results clearly suggest that a simple Sanders et al. (1988) evolutionary scenario of a merger triggered starburst event with an evolving obscured-to-naked quasar scenario to be unlikely. Rather, it seems that the nature of the diverse infrared luminous ULIGs and HyLIGs has a far more complex history. Whether a starburst or an AGN dominates the bolometric luminosity of a ULIG is most probably determined by local conditions, evolution and obscuration by dust. Nearby ULIGs are mostly starburst-dominated systems with ‘cold’ colours. This cold fraction decreases at higher luminosities where AGN contributions become prevalent and galaxy SEDs are ‘warmer’. ULIGs are often linked to high-redshift (z ∼ 2)

2FeLoBALs show absorption lines from many excited states of Fe II and having low ionisation lines and are heavily absorbed in X-rays. 398 A. VERMA ET AL.

(sub-)mm emitting galaxies which comprise a signiﬁcant fraction of the infrared background and contribute signiﬁcantly to the star formation density at that epoch. If ULIGs and sub-mm galaxies are members of the same population, then strong evolution is implied, since the locally unimportant ULIGs, that constitute less than 1% of the local star formation density must dominate the infrared number counts at higher redshifts (Elbaz and Cesarsky, 2003). The strongly evolving z 1.5 LIGs discovered in ISOCAM surveys (see Elbaz et al., this volume; Oliver et al., this volume) that resolve 10–60% of the far-IR-sub-mm background may provide a step between.

7. Active Galaxies as Seen by ISO: A Summary

The Infrared Space Observatory has produced a wealth of information regarding the infrared emission properties of the full range of active galaxies. Greatly enhancing the IRAS legacy, it has provided a physical basis on which the next generation of infrared telescopes such as Spitzer, SOFIA, ASTRO-F and Herschel will build. Most of ISO’s power was due to the improved sensitivity and extended wavelength coverage. The addition of high-resolution spectroscopy and development of new diagnostic diagrams have enabled us to probe the ongoing physical conditions of the ISM in active galaxies and to quantitatively assess the presence and bolometric contribution of star formation or an AGN in infrared luminous sources. Analysis of fine structure lines provided an excellent set of diagnostic tools that will be better populated and applied to a wider variety of sources in future missions with even higher sensitivity. The weakness of the sulphur fine structure lines in starburst galaxies and of the cooling line [C II]158 μm in the most luminous active galaxies suggests that other strong lines (e.g. [Ne II]12.8 μm) should be the are preferred indicators of star-forming objects. The finding of IRAS that galaxies with “cold” infrared colours are starburst dominated and “warm” are AGN dominated was broadly confirmed with the new ISO data. Although, it was shown that spatially resolved spectroscopic information is necessary to address the issue of the origin of the far-infrared emission in AGN and the role of concurrent star formation. ISO results on AGN in general support unification scenarios, but suggest that a clumpy distribution of obscuring material around the active nucleus may be a better representation than the commonly assumed torus. While the majority of ULIGs are starburst dominated, it now appears that an increasing fraction of AGN at higher infrared luminosities is present. Mergers are indeed the trigger of ULIG activity, but the simple generic evolution of a luminous system from a ULIG to an optical- QSO phase appears unlikely with more local factors influencing the presence of an AGN. With increases in sensitivity and spatial resolution, it is certain that the suite of current and forthcoming infrared telescopes will enhance our knowledge of the obscured active universe by probing fainter and to higher redshifts as well as in more spatial detail. OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 399

Acknowledgements

We would like to thank Matthias Tecza for his comments on this text. We are also grateful to Eckhard Sturm for providing the SWS + LWS spectra shown in Figure 1 and for discussion. This work was supported by the DLR grant FKZ: 50QI 0202. This work is based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, The Netherlands and the United Kingdom) and with the participation of ISAS and NASA.

References

Acosta-Pulido, J. A., Klaas, U., Laureijs, R. J., Schulz, B., Kinkel, U., Abraham, P., et al.: 1996, A&A 315, L121. Afonso, J., Mobasher, B., Chan, B., and Cram, L.: 2001, ApJ 559, L101. Alexander, D. M., Efstathiou, A., Hough, J. H., Aitken, D. K., Lutz, D., Roche, P., et al.: 1999a, MNRAS 310, 78. Alexander, D. M., La Franca, F., Fiore, F., Barcons, X., Ciliegi, P., Danese, L., et al.: 2001, ApJ 554, 18. Alexander, D. M., Ruiz, M., and Hough, J. H.: 1999b, ESA SP-435: Workshop on ISO Polarisation Observations, 1 pp. Alexander, T., Lutz, D., Sturm, E., Genzel, R., Sternberg, A., and Netzer, H.: 2000, ApJ 536, 710. Alexander, T., Sturm, E., Lutz, D., Sternberg, A., Netzer, H., and Genzel, R.: 1999c, ApJ 512, 204. Allamandola, L. J., Sandford, S. A., Hudgins, D. M., and Witteborn, F. C.: 1995, ASP Conference Series 73: From Gas to Stars to Dust, pp. 23–32. Allamandola, L. J., Tielens, A. G. G. M., and Barker, J. R.: 1985, ApJ 290, L25. Alonso-Herrero, A., Quillen, A. C., Simpson, C., Efstathiou, A., and Ward, M. J.: 2001, AJ 121, 1369. Alton, P. B., Trewhella, M., Davies, J. I., Evans, R., Bianchi, S., Gear, W., et al.: 1998, A&A 335, 807. Andreani, P., Fosbury, R. A. E., van Bemmel, I., and Freudling, W.: 2002, A&A 381, 389. Ant´on,S., Thean, A. H. C., Pedlar, A., and Browne, I. W. A.: 2002, MNRAS 336, 319. Antonucci, R.: 1993, ARA&A 31, 473. Appleton, P. N., Charmandaris, V., Gao, Y., Combes, F., Ghigo, F., Horellou, C., et al.: 2002, ApJ 566, 682. Appleton, P. N., Charmandaris, V., Horellou, C., Mirabel, I. F., Ghigo, F., Higdon, J. L., et al.: 1999, ApJ 527, 143. Appleton, P. N., Charmandaris, V., Horellou, C., Mirabel, I. F., and Laurent, O.: 2000, LNP Vol. 548: ISO Survey of a Dusty Universe, 232 pp. Appleton, P. N. and Struck-Marcell, C.: 1996, Fundam. Cosmic Phys. 16, 111. Aussel, H., Gerin, M., Boulanger, F., Desert, F. X., Casoli, F., Cutri, R. M., et al.: 1998, A&A 334, L73. Baluteau, J., Damour, F., Caux, E., Gry, C., and Caccarelli, C.: 2003, ESA SP-511: Exploiting the ISO Data Archive. Infrared Astronomy in the Internet Age, 209 pp. Barger, A. J., Cowie, L. L., and Richards, E. A.: 2000, AJ 119, 2092. Barthel, P. D.: 1989, ApJ 336, 606. Beichman, C. A.: 1987, ARA&A 25, 521. Beichman, C. A., Neugebauer, G., Habing, H. J., Clegg, P.E., and Chester, T. J.: 1988, NASA RP-1190, Vol. 1, 0 pp. 400 A. VERMA ET AL.

Bendo, G. J., Joseph, R. D., Wells, M., Gallais, P., Haas, M., Heras, A. M., et al.: 2003, AJ 125, 2361. Bergvall, N., Masegosa, J., Ostlin,¨ G., and Cernicharo, J.: 2000, A&A 359, 41. Bertone, E., Tagliaferri, G., Ghisellini, G., Treves, A., Barr, P., Celotti, A., et al.: 2000, A&A 356,1. Binette, L., Wilson, A. S., Raga, A., and Storchi-Bergmann, T.: 1997, A&A 327, 909. Block, D. L. and Sauvage, M.: 2000, A&A 353, 72. Boselli, A., Lequeux, J., Contursi, A., Gavazzi, G., Boulade, O., Boulanger, F., et al.: 1997, A&A 324, L13. Boselli, A., Lequeux, J., Sauvage, M., Boulade, O., Boulanger, F., Cesarsky, D., et al.: 1998, A&A 335, 53. Boulade, O., Sauvage, M., Altieri, B., Blommaert, J., Gallais, P., Guest, S., et al.: 1996, A&A 315, L85. Boulanger, F., Abergel, A., Bernard, J. P., Cesarsky, D., Puget, J. L., Reach, W. T., et al.: 1998a, ASP Conference Series 132: Star Formation with the Infrared Space Observatory,15pp. Boulanger, F., Boisssel, P., Cesarsky, D., and Ryter, C.: 1998b, A&A 339, 194. Bradford, C. M., Stacey, G. J., Fischer, J., Smith, H. A., Cohen, R. J., Greenhouse, M. A., et al.: 1999, ESA SP-427: The Universe as Seen by ISO, 861 pp. Brauher, J. and Lord, S.: 2000, Bull. Am. Astron. Soc. 32, 1286. Brauher, J., Lord, S., Unger, S., Wolﬁre, M., and Fischer, J.: 1997, Bull. Am. Astron. Soc. 29, 1246. Calzetti, D., Armus, L., Bohlin, R. C., Kinney, A. L., Koornneef, J., and Storchi-Bergmann, T.: 2000, ApJ 533, 682. Canalizo, G. and Stockton, A.: 2001, ApJ 555, 719. Cesarsky, D., Lequeux, J., Abergel, A., Perault, M., Palazzi, E., Madden, S., et al.: 1996a, A&A 315, L309. Cesarsky, D., Lequeux, J., Abergel, A., Perault, M., Palazzi, E., Madden, S., et al.: 1996b, A&A 315, L305. Charmandaris, V., Combes, F., and van der Hulst, J. M.: 2000, A&A 356, L1. Charmandaris, V., Laurent, O., Le Floc’h, E., Mirabel, I. F., Sauvage, M., Madden, S. C., et al.: 2002, A&A 391, 429. Charmandaris, V., Laurent, O., Mirabel, I. F., and Gallais, P.: 2001a, Astrophys. Space Sci. Suppl. 277, 55. Charmandaris, V., Laurent, O., Mirabel, I. F., Gallais, P., Sauvage, M., Vigroux, L., et al.: 2001b, Ap&SS 276, 553. Charmandaris, V., Laurent, O., Mirabel, I. F., Gallais, P., Sauvage, M., Vigroux, L., et al.: 1999, A&A 341, 69. Charmandaris, V., Mirabel, I. F., Tran, D., Laurent, O., Cesarsky, C. J., Gallais, P., et al.: 1997, Extragalactic Astronomy in the Infrared, 283 pp. Chiar, J. E., Tielens, A. G. G. M., Whittet, D. C. B., Schutte, W. A., Boogert, A. C. A., Lutz, D., et al.: 2000, ApJ 537, 749. Clavel, J.: 1998, A&A 331, 853. Clavel, J., Schulz, B., Altieri, B., Barr, P., Claes, P., Heras, A., et al.: 2000, A&A 357, 839. Colbert, J. W., Malkan, M. A., Clegg, P. E., Cox, P., Fischer, J., Lord, S. D., et al.: 1999, ApJ 511, 721. Combes, F.: 2001, Astrophys. Space Sci. Suppl. 277, 29. Contini, M., Prieto, M. A., and Viegas, S. M.: 1998, ApJ 505, 621. Contini, M., Viegas, S. M., and Prieto, M. A.: 2002, A&A 386, 399. Contursi, A., Lequeux, J., Cesarsky, D., Boulanger, F., Rubio, M., Hanus, M., et al.: 2000, A&A 362, 310. Cox, P. P., Roelfsema, P. R., Baluteau, J. P., Peeters, E., Martin-Hernandez, L., Tielens, A. G. G. M., et al.: 1999, ESA SP-427: The Universe as Seen by ISO, 631 pp. Crowther, P. A., Beck, S. C., Willis, A. J., Conti, P. S., Morris, P. W., et al.: 1999, MNRAS 304, 654. OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 401

Dale, D. A., Helou, G., Brauher, J. R., Cutri, R. M., Malhotra, S., and Beichman, C. A.: 2004, ApJ 604, 565. Dale, D. A., Helou, G., Contursi, A., Silbermann, N. A., and Kolhatkar, S.: 2001a, ApJ 549, 215. Dale, D. A., Helou, G., Neugebauer, G., Soifer, B. T., Frayer, D. T., and Condon, J. J.: 2001b, AJ 122, 1736. Dale, D. A., Silbermann, N. A., Helou, G., Valjavec, E., Malhotra, S., Beichman, C. A., et al.: 2000, AJ 120, 583. Dartois, E., Marco, O., Mu noz-Caro, G. M., Brooks, K., Deboffle, D., and d’Hendecourt, L.: 2004, A&A 423, 549. Davies, J. I., Alton, P., Trewhella, M., Evans, R., and Bianchi, S.: 1999, MNRAS 304, 495. de Grijp, M. H. K., Miley, G. K., Lub, J., and de Jong, T.: 1985, Nature 314, 240. Dennefeld, M., Boller, T., Rigopoulou, D., and Spoon, H. W. W.: 2003, A&A 406, 527. Desert, F.-X., Boulanger, F., and Puget, J. L.: 1990, A&A 237, 215. Domingue, D. L., Keel, W. C., Ryder, S. D., and White, R. E.: 1999, AJ 118, 1542. Domingue, D. L., Sulentic, J. W., and Durbala, A.: 2005, ArXiv Astrophysics e-prints. Domingue, D. L., Sulentic, J. W., Xu, C., Mazzarella, J., Gao, Y., and Rampazzo, R.: 2003, AJ 125, 555. Downes, D. and Solomon, P. M.: 1998, ApJ 507, 615. Draine, B. T.: 2003, ARA&A 41, 241. Draine, B. T. and Li, A.: 2001, ApJ 551, 807. Duc, P.-A., Hall, P. B., Fadda, D., Chanial, P., Elbaz, D., Monaco, P., et al.: 2002, A&A 389, L47. Efstathiou, A., Rowan-Robinson, M., and Siebenmorgen, R.: 2000, MNRAS 313, 734. Egami, E., Iwamuro, F., Maihara, T., Oya, S., and Cowie, L. L.: 1996, AJ 112, 73. Elbaz, D. and Cesarsky, C. J.: 2003, Science 300, 270. Elbaz, D., Flores, H., Chanial, P., Mirabel, I. F., Sanders, D., Duc, P.-A., et al.: 2002, A&A 381, L1. Elitzur, M., Nenkova, M., and Ivezic, Z.: 2003, ArXiv Astrophysics e-prints. Elvis, M.: 2000, New Astron. Rev. 44, 559. Förster Schreiber, N. M., Genzel, R., Lutz, D., Kunze, D., and Sternberg, A.: 2001, ApJ 552, 544. Förster Schreiber, N. M., Roussel, H., Sauvage, M., and Charmandaris, V.: 2004, A&A 419, 501. Förster Schreiber, N. M., Sauvage, M., Charmandaris, V., Laurent, O., Gallais, P., Mirabel, I. F., et al.: 2003, A&A 399, 833. Fanti, C., Pozzi, F., Fanti, R., Baum, S. A., O’Dea, C. P., Bremer, M., et al.: 2000, A&A 358, 499. Farrah, D., Afonso, J., Efstathiou, A., Rowan-Robinson, M., Fox, M., and Clements, D.: 2003, MNRAS 343, 585. Fischer, J.: 2000, ESA SP-456: ISO Beyond the Peaks: The 2nd ISO Workshop on Analytical Spec- troscopy, 239 pp. Fischer, J., Lord, S. D., Unger, S. J., Bradford, C. M., Brauher, J. R., Clegg, P. E., et al.: 1999a, ESA SP-427: The Universe as Seen by ISO, 817 pp. Fischer, J., Luhman, M. L., Satyapal, S., Greenhouse, M. A., Stacey, G. J., Bradford, C. M., et al.: 1999b, Ap&SS 266, 91. Fischer, J., Satyapal, S., Melnick, G., Cox, P., Cernicharo, J., Luhman, M. L., et al.: 1997, Bull. Am. Astron. Soc. 29, 1356. Freudling, W., Siebenmorgen, R., and Haas, M.: 2003, ApJ 599, L13. Gallais, P., Charmandaris, V., Le Floc’h, E., Mirabel, I. F., Sauvage, M., Vigroux, L., et al.: 2004, A&A 414, 845. Galliano, F., Madden, S. C., Jones, A. P., Wilson, C. D., Bernard, J.-P., and Le Peintre, F.: 2003, A&A 407, 159. Gao, Y., Wang, Q. D., Appleton, P. N., and Lucas, R. A.: 2003, ApJ 596, L171. Gao, Y., Xu, C., Lu, N. Y., Tuffs, R., and Sulentic, J.: 2001, Bull. Am. Astron. Soc. 33, 835. 402 A. VERMA ET AL.

Genzel, R., Lutz, D., Sturm, E., Egami, E., Kunze, D., Moorwood, A. F. M., et al.: 1998, ApJ 498, 579. Gillett, F. C., Kleinmann, D. E., Wright, E. L., and Capps, R. W.: 1975, ApJ 198, L65. Giveon, U., Sternberg, A., Lutz, D., Feuchtgruber, H., and Pauldrach, A. W. A.: 2002, ApJ 566, 880. González-Alfonso,E., Smith, H. A., Fischer, J., and Cernicharo, J.: 2004, ApJ 613, 247. Gorjian, V., Turner, J. L., and Beck, S. C.: 2001, ApJ 554, L29. Granato, G. L. and Danese, L.: 1994, MNRAS 268, 235. Haas, M.: 1998, A&A 337, L1. Haas, M., Chini, R., and Klaas, U.: 2005, A&A 433, L17. Haas, M., Chini, R., Meisenheimer, K., Stickel, M., Lemke, D., Klaas, U., et al.: 1998a, ApJ 503, L109. Haas, M., Klaas, U., and Bianchi, S.: 2002, A&A 385, L23. Haas, M., Klaas, U., Coulson, I., Thommes, E., and Xu, C.: 2000a, A&A 356, L83. Haas, M., Klaas, U., Müller, S. A. H., Bertoldi, F., Camenzind, M., Chini, R., et al.: 2003, A&A 402, 87. Haas, M., Klaas, U., Müller, S. A. H., Chini, R., and Coulson, I.: 2001, A&A 367, L9. Haas, M., Lemke, D., Stickel, M., Hippelein, H., Kunkel, M., Herbstmeier, U., et al.: 1998b, A&A 338, L33. Haas, M., Müller, S. A. H., Bertoldi, F., Chini, R., Egner, S., Freudling, W., et al.: 2004, A&A 424, 531. Haas, M., Müller, S. A. H., Chini, R., Meisenheimer, K., Klaas, U., Lemke, D., et al.: 2000b, A&A 354, 453. Hanlon, L., Laureijs, R. J., Metcalfe, L., McBreen, B., Altieri, B., Castro-Tirado, A., et al.: 2000, A&A 359, 941. Hanlon, L., Metcalfe, L., Delaney, M., Laureijs, R., McBreen, B., Smith, N., et al.: 1999, A&AS 138, 459. Hazard, C., Morton, D. C., Terlevich, R., and McMahon, R.: 1984, ApJ 282, 33. He, J. H. and Chen, P. S.: 2004, New Astron. 9, 545. Heckman, T. M.: 1980, A&A 87, 152. Heckman, T. M., Robert, C., Leitherer, C., Garnett, D. R., and van der Rydt, F.: 1998, ApJ 503, 646. Helou, G., Lu, N. Y., Werner, M. W., Malhotra, S., and Silbermann, N.: 2000, ApJ 532, L21. Hirashita, H. and Hunt, L. K.: 2004, A&A 421, 555. Houck, J. R., Soifer, B. T., Neugebauer, G., Beichman, C. A., Aumann, H. H., Clegg, P. E., et al.: 1984, ApJ 278, L63. Hunt, L. K., Vanzi, L., and Thuan, T. X.: 2001, A&A 377, 66. Hur, M., Stacey, G. J., Fischer, J. A., Smith, H. A., Unger, S., Lord, S. D., et al.: 1996, Bull. Am. Astron. Soc. 28, 1392. Hwang, C., Lo, K. Y., Gao, Y., Gruendl, R. A., and Lu, N. Y.: 1999, ApJ 511, L17. Imanishi, M.: 2002, ApJ 569, 44. Ivison, R. J., Dunlop, J. S., Hughes, D. H., Archibald, E. N., Stevens, J. A., Holland, W. S., et al.: 1998, ApJ 494, 211. Jarrett, T. H., Helou, G., Van Buren, D., Valjavec, E., and Condon, J. J.: 1999, AJ 118, 2132. Joseph, R. D.: 1999, Ap&SS 266, 321. Kegel, W. H., Hertenstein, T., and Quirrenbach, A.: 1999, A&A 351, 472. Kennicutt, R. C.: 1998, ARA&A 36, 189. Kewley, L. J., Dopita, M. A., Sutherland, R. S., Heisler, C. A., and Trevena, J.: 2001, ApJ 556, 121. Klaas, U., Haas, M., Müller, S. A. H., Chini, R., Schulz, B., Coulson, I., et al.: 2001, A&A 379, 823. Klaas, U., Laureijs, R. J., and Clavel, J.: 1999, ApJ 512, 157. Klaas, U. and Walker, H. J.: 2002, A&A 391, 911. OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 403

Kleinmann, D. E. and Low, F. J.: 1970a, ApJ 161, L203. Kleinmann, D. E. and Low, F. J.: 1970b, ApJ 159, L165. Krause, O., Lisenfeld, U., Lemke, D., Haas, M., Klaas, U., and Stickel, M.: 2003, A&A 402, L1. Kriss, G. A., Davidsen, A. F., Zheng, W., Kruk, J. W., and Espey, B. R.: 1995, ApJ 454, L7. Kuraszkiewicz, J. K., Wilkes, B. J., Hooper, E. J., McLeod, K. K., Wood, K., Bjorkman, J., et al.: 2003, ApJ 590, 128. Laureijs, R. J., Acosta-Pulido, J., Abraham, P., Kinkel, U., Klaas, U., Castaneda, H. O., et al.: 1996, A&A 315, L313. Laurent, O., Mirabel, I. F., Charmandaris, V., Gallais, P., Madden, S. C., Sauvage, M., et al.: 2000, A&A 359, 887. Laurent, O., Mirabel, I. F., Charmandaris, V., Gallais, P., Vigroux, L., and Cesarsky, C. J.: 1999, ESA SP-427: The Universe as Seen by ISO, p. 913. Le Floc’h, E., Charmandaris, V., Laurent, O., Mirabel, I. F., Gallais, P., Sauvage, M., et al.: 2002, A&A 391, 417. Le Floc’h, E., Mirabel, I. F., Laurent, O., Charmandaris, V., Gallais, P., Sauvage, M., et al.: 2001, A&A 367, 487. Leech, K. J., Saxton, R., Barr, P., Santos-Lleo, M., Jiminez-Bailon, E., Diaz, R., et al.: 2002, Publ. Astron. Soc. Aust. 19, 152. Leger, A. and Puget, J. L.: 1984, A&A 137, L5. Lemke, D., Mattila, K., Lehtinen, K., Laureijs, R. J., Liljestrom, T., Leger, A., et al.: 1998, A&A 331, 742. L´ıpari, S.: 1994, ApJ 436, 102. L´ıpari, S., Terlevich, R., D´ıaz,R. J., Taniguchi, Y., Zheng, W., Tsvetanov, Z., et al.: 2003, MNRAS 340, 289. Lonsdale Persson, C. J. and Helou, G.: 1987, ApJ 314, 513. Low, J. and Kleinmann, D. E.: 1968, AJ 73, 868. Lu, N., Helou, G., Werner, M. W., Dinerstein, H. L., Dale, D. A., Silbermann, N. A., et al.: 2003, ApJ 588, 199. Luhman, M. L., Satyapal, S., Fischer, J., Wolfire, M. G., Cox, P., Lord, S. D., et al.: 1998, ApJ 504, L11. Luhman, M. L., Satyapal, S., Fischer, J., Wolfire, M. G., Sturm, E., Dudley, C. C., et al.: 2003, ApJ 594, 758. Lutz, D., Genzel, R., Sturm, E., Tacconi, L., Wieprecht, E., Alexander, T., et al.: 2000a, ApJ 530, 733. Lutz, D., Kunze, D., Spoon, H. W. W., and Thornley, M. D.: 1998a, A&A 333, L75. Lutz, D., Maiolino, R., Moorwood, A. F. M., Netzer, H., Wagner, S. J., Sturm, E., et al.: 2002, A&A 396, 439. Lutz, D., Maiolino, R., Spoon, H. W. W., and Moorwood, A. F. M.: 2004a, A&A 418, 465. Lutz, D., Spoon, H. W. W., Rigopoulou, D., Moorwood, A. F. M., and Genzel, R.: 1998b, ApJ 505, L103. Lutz, D., Sturm, E., Genzel, R., Moorwood, A. F. M., Alexander, T., Netzer, H., et al.: 2000b, ApJ 536, 697. Lutz, D., Sturm, E., Genzel, R., Spoon, H. W. W., Moorwood, A. F. M., Netzer, H., et al.: 2003, A&A 409, 867. Lutz, D., Sturm, E., Genzel, R., Spoon, H. W. W., and Stacey, G. J.: 2004b, A&A 426, L5. Lutz, D., Veilleux, S., and Genzel, R.: 1999, ApJ 517, L13. Müller, S. A. H., Haas, M., Siebenmorgen, R., Klaas, U., Meisenheimer, K., Chini, R., et al.: 2004, A&A 426, L29. Madau, P., Ferguson, H. C., Dickinson, M. E., Giavalisco, M., Steidel, C. C., and Fruchter, A.: 1996, MNRAS 283, 1388. Madden, S. C.: 2000, New Astron. Rev. 44, 249. 404 A. VERMA ET AL.

Malhotra, S., Helou, G., Stacey, G., Hollenbach, D., Lord, S., Beichman, C. A., et al.: 1997, ApJ 491, L27. Malhotra, S., Kaufman, M., Hollenbach, D., Helou, G., Rubin, R., Brauher, J., et al.: 1999, Bull. Am. Astron. Soc. 31, 1455. Malhotra, S., Kaufman, M. J., Hollenbach, D., Helou, G., Rubin, R. H., Brauher, J., et al.: 2001, ApJ 561, 766. Manning, C. and Spinrad, H.: 2001, AJ 122, 113. Martins, L. P., Viegas, S. M., and Gruenwald, R.: 2003, ApJ 587, 562. Masegosa, J., Márquez,I., Sempere, M. J., and Cernicharo, J.: 2001, ESA SP-460: The Promise of the Herschel Space Observatory, p. 459. Mattila, K., Lehtinen, K., and Lemke, D.: 1999, A&A 342, 643. Mattila, K., Lemke, D., Haikala, L. K., Laureijs, R. J., Leger, A., Lehtinen, K., et al.: 1996, A&A 315, L353. Meisenheimer, K., Haas, M., Müller, S. A. H., Chini, R., Klaas, U., and Lemke, D.: 2001, A&A 372, 719. Metcalfe, L., Steel, S. J., Barr, P., Clavel, J., Delaney, M., Gallais, P., et al.: 1996, A&A 315, L105. Mihos, J. C. and Hernquist, L.: 1996, ApJ 464, 641. Mirabel, I. F. and Laurent, O.: 1999, Ap&SS 269, 349. Mirabel, I. F., Laurent, O., Sanders, D. B., Sauvage, M., Tagger, M., Charmandaris, V., et al.: 1999, A&A 341, 667. Mirabel, I. F., Vigroux, L., Charmandaris, V., Sauvage, M., Gallais, P., Tran, D., et al.: 1998, A&A 333, L1. Mochizuki, K. and Onaka, T.: 2001, A&A 370, 868. Moorwood, A. F. M.: 1999, ESA SP-427: The Universe as Seen by ISO, p. 825. Moorwood, A. F. M., Lutz, D., Oliva, E., Marconi, A., Netzer, H., Genzel, R., et al.: 1996, A&A 315, L109. Morel, T., Efstathiou, A., Serjeant, S., Márquez,I., Masegosa, J., Héraudeau,P., et al.: 2001, MNRAS 327, 1187. Mouri, H., Taniguchi, Y., Sato, Y., and Kawara, K.: 1998, A&A 334, 482. Murphy, T. W., Soifer, B. T., Matthews, K., and Armus, L.: 2001, ApJ 559, 201. Negishi, T., Onaka, T., Chan, K.-W., and Roellig, T. L.: 2001, A&A 375, 566. Nenkova, M., Ivezić, Z.,ˇ and Elitzur, M.: 2002, ApJ 570, L9. Nollenberg, J. G., Skillman, E. D., Garnett, D. R., and Dinerstein, H. L.: 2002, ApJ 581, 1002. Odenwald, S., Newmark, J., and Smoot, G.: 1998, ApJ 500, 554. O’Halloran, B., Metcalfe, L., Delaney, M., McBreen, B., Laureijs, R., Leech, K., et al.: 2000, A&A 360, 871. O’Halloran, B., Metcalfe, L., McBreen, B., Laureijs, R., Leech, K., Delaney, M., et al.: 2002, ApJ 575, 747. Oliva, E., Marconi, A., and Moorwood, A. F. M.: 1999, A&A 342, 87. Orr, M. J. L. and Browne, I. W. A.: 1982, MNRAS 200, 1067. Pérez Garc´ıa,A. M. and Rodr´ıguezEspinosa, J. M.: 1998, Ap&SS 263, 103. Pérez Garc´ıa,A. M. and Rodr´ıguezEspinosa, J. M.: 2001, ApJ 557, 39. Pérez Garc´ıa,A. M., Rodr´ıguezEspinosa, J. M., and Fuensalida, J. J.: 2000, ApJ 529, 875. Pagani, L., Lequeux, J., Cesarsky, D., Donas, J., Milliard, B., Loinard, L., et al.: 1999, A&A 351, 447. Peeters, E., Allamandola, L. J., Hudgins, D. M., Hony, S., and Tielens, A. G. G. M.: 2003, ArXiv Astrophysics e-prints. Peeters, E., Spoon, H. W. W., and Tielens, A. G. G. M.: 2004, ApJ 613, 986. Peng, B., Kraus, A., Krichbaum, T. P., Müller,S. A. H., Qian, S. J., Quirrenbach, A., et al.: 2000, A&A 353, 937. Pérez Garc´ıa,A. M., Rodr´ıguezEspinosa, J. M., and Santolaya Rey, A. E.: 1998, ApJ 500, 685. OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 405

Plante, S. and Sauvage, M.: 2002, AJ 124, 1995. Polletta, M. and Courvoisier, T. J. L.: 1999a, ESA SP-427: The Universe as Seen by ISO, p. 953. Polletta, M. and Courvoisier, T. J.-L.: 1999b, A&A 350, 765. Polletta, M., Courvoisier, T. J.-L., Hooper, E. J., and Wilkes, B. J.: 2000, A&A 362, 75. Popescu, C. C. and Tuffs, R. J.: 2003, A&A 410, L21. Popescu, C. C., Tuffs, R. J., Kylafis, N. D., and Madore, B. F.: 2004, A&A 414, 45. Popescu, C. C., Tuffs, R. J., Völk,H. J., Pierini, D., and Madore, B. F.: 2002, ApJ 567, 221. Prieto, M. A. and Acosta-Pulido, J. A.: 2003, ApJ 583, 689. Prieto, M. A., PérezGarc´ıa,A. M., and Rodr´ıguezEspinosa, J. M.: 2001, A&A 377, 60. Puget, J. L. and Leger, A.: 1989, ARA&A 27, 161. Radovich, M., Kahanpää, J., and Lemke, D.: 2001, A&A 377, 73. Radovich, M., Klaas, U., Acosta-Pulido, J., and Lemke, D.: 1999, A&A 348, 705. Rigopoulou, D., Kunze, D., Lutz, D., Genzel, R., and Moorwood, A. F. M.: 2002, A&A 389, 374. Rigopoulou, D., Lutz, D., Genzel, R., Egami, E., Kunze, D., Sturm, E., et al.: 1996, A&A 315, L125. Rigopoulou, D., Spoon, H. W. W., Genzel, R., Lutz, D., Moorwood, A. F. M., and Tran, Q. D.: 1999, AJ 118, 2625. Rocca-Volmerange, B. and Remazeilles, M.: 2005, A&A 433, 73. Rodr´ıguezEspinosa, J. M. and PérezGarc´ıa,A. M.: 1997, ApJ 487, L33. Rodr´ıguezEspinosa, J. M., PérezGarc´ıa,A. M., Lemke, D., and Meisenheimer, K.: 1996, A&A 315, L129. Roussel, H., Sauvage, M., Vigroux, L., and Bosma, A.: 2001, A&A 372, 427. Rush, B., Malkan, M. A., and Spinoglio, L.: 1993, ApJS 89,1. Sanders, D. B.: 1999, Ap&SS 266, 331. Sanders, D. B. and Mirabel, I. F.: 1996, ARA&A 34, 749. Sanders, D. B., Soifer, B. T., Elias, J. H., Neugebauer, G., and Matthews, K.: 1988, ApJ 328, L35. Sanei, M. and Satyapal, S.: 2002, Bull. Am. Astron. Soc. 34, 1294. Santos-Lleó,M., Clavel, J., Schulz, B., Altieri, B., Barr, P., Alloin, D., et al.: 2001, A&A 369, 57. Satyapal, S., Sambruna, R. M., and Dudik, R. P.: 2004, A&A 414, 825. Sauvage, M., Blommaert, J., Boulanger, F., Cesarsky, C. J., Cesarsky, D. A., Desert, F. X., et al.: 1996, A&A 315, L89. Sauvage, M. and Thuan, T. X.: 1992, ApJ 396, L69. Schaerer, D., Contini, T., and Pindao, M.: 1999, A&AS 136, 35. Schaerer, D. and Stasińska,G.: 1999, A&A 345, L17. Schulz, B., Clavel, J., Altieri, B., Barr, P., Claes, P., Heras, A. M., et al.: 2000, ESA SP-456: ISO Beyond the Peaks: The 2nd ISO Workshop on Analytical Spectroscopy, p. 245. Sellgren, K.: 1984, ApJ 277, 623. Shier, L. M., Rieke, M. J., and Rieke, G. H.: 1996, ApJ 470, 222. Siebenmorgen, R. and Efstathiou, A.: 2001, A&A 376, L35. Siebenmorgen, R., Freudling, W., E. Krügel,and Haas, M.: 2004a, A&A 421, 129. Siebenmorgen, R., Krügel,E., and Chini, R.: 1999a, A&A 351, 495. Siebenmorgen, R., Krügel,E., and Laureijs, R. J.: 2001, A&A 377, 735. Siebenmorgen, R., Krügel,E., and Spoon, H. W. W.: 2004b, A&A 414, 123. Siebenmorgen, R., Krügel,E., and Zota, V.: 1999b, A&A 351, 140. Siebenmorgen, R. and Kruegel, E.: 1992, A&A 259, 614. Silva, L., Granato, G. L., Bressan, A., and Danese, L.: 1998, ApJ 509, 103. Skinner, C. J., Smith, H. A., Sturm, E., Barlow, M. J., Cohen, R. J., and Stacey, G. J.: 1997, Nature 386, 472. Smith, B. J.: 1998, ApJ 500, 181. Soifer, B. T., Neugebauer, G., Matthews, K., Egami, E., Becklin, E. E., Weinberger, A. J., et al.: 2000, AJ 119, 509. 406 A. VERMA ET AL.

Soifer, B. T., Neugebauer, G., Matthews, K., Egami, E., Weinberger, A. J., Ressler, M., et al.: 2001, AJ 122, 1213. Solomon, P. M., Downes, D., Radford, S. J. E., and Barrett, J. W.: 1997, ApJ 478, 144. Spinoglio, L., Andreani, P., and Malkan, M. A.: 2002, ApJ 572, 105. Spinoglio, L., Malkan, M. A., Smith, H. A., and Fischer, J.: 2003, ASP Conf. Ser. 290: Active Galactic Nuclei: From Central Engine to Host Galaxy, p. 557. Spinoglio, L., Malkan, M. A., Smith, H. A., Gonz´alez-Alfonso,E., and Fischer, J.: 2005, ApJ 623, 123. Spoon, H. W. W., Keane, J. V., Tielens, A. G. G. M., Lutz, D., and Moorwood, A. F. M.: 2001, A&A 365, L353. Spoon, H. W. W., Keane, J. V., Tielens, A. G. G. M., Lutz, D., Moorwood, A. F. M., and Laurent, O.: 2002, A&A 385, 1022. Spoon, H. W. W., Koornneef, J., Moorwood, A. F. M., Lutz, D., and Tielens, A. G. G. M.: 2000, A&A 357, 898. Spoon, H. W. W., Moorwood, A. F. M., Lutz, D., Tielens, A. G. G. M., Siebenmorgen, R., and Keane, J. V.: 2004, A&A 414, 873. Stevens, J. A., Amure, M., and Gear, W. K.: 2005, MNRAS 357, 361. Stickel, M., van der Hulst, J. M., van Gorkom, J. H., Schiminovich, D., and Carilli, C. L.: 2004, A&A 415, 95. Struck, C.: 1999, Phys. Rep. 321,1. Sturm, E., Alexander, T., Lutz, D., Sternberg, A., Netzer, H., and Genzel, R.: 1999, ApJ 512, 197. Sturm, E., Lutz, D., Genzel, R., Sternberg, A., Egami, E., Kunze, D., et al.: 1996, A&A 315, L133. Sturm, E., Lutz, D., Tran, D., Feuchtgruber, H., Genzel, R., Kunze, D., et al.: 2000, A&A 358, 481. Sturm, E., Lutz, D., Verma, A., Netzer, H., Sternberg, A., Moorwood, A. F. M., et al.: 2002, A&A 393, 821. Sugai, H. and Malkan, M. A.: 2000, ApJ 529, 219. Suter, J. P., Harvey, V., Cox, P., Fischer, J., Lord, S. D., Melnick, G. J., et al.: 1998, Bull. Am. Astron. Soc. 30, 1383. Takagi, T., Arimoto, N., and Hanami, H.: 2003, MNRAS 340, 813. Takeuchi, T. T., Hirashita, H., Ishii, T. T., Hunt, L. K., and Ferrara, A.: 2003, MNRAS 343, 839. Taniguchi, Y., Cowie, L. L., Sato, Y., Sanders, D. B., Kawara, K., Joseph, R., et al.: 1997, A&A 328, L9. Thornley, M. D., Schreiber, N. M. F., Lutz, D., Genzel, R., Spoon, H. W. W., Kunze, D., et al.: 2000, ApJ 539, 641. Thuan, T. X., Sauvage, M., and Madden, S.: 1999, ApJ 516, 783. Tran, Q. D., Lutz, D., Genzel, R., Rigopoulou, D., Spoon, H. W. W., Sturm, E., et al.: 2001, ApJ 552, 527. Trewhella, M., Davies, J. I., Alton, P. B., Bianchi, S., and Madore, B. F.: 2000, ApJ 543, 153. Uchida, K. I., Sellgren, K., and Werner, M.: 1998, ApJ 493, L109. Unger, S. J., Clegg, P. E., Stacey, G. J., Cox, P., Fischer, J., Greenhouse, M., et al.: 2000, A&A 355, 885. Vacca, W. D., Johnson, K. E., and Conti, P. S.: 2002, AJ 123, 772. Valentijn, E. A. and van der Werf, P. P.: 1999, ApJ 522, L29. van Bemmel, I. M., Barthel, P. D., and de Graauw, T.: 2000, A&A 359, 523. Vanzi, L. and Sauvage, M.: 2004, A&A 415, 509. Vastel, C., Polehampton, E. T., Baluteau, J.-P., Swinyard, B. M., Caux, E., and Cox, P.: 2002, ApJ 581, 315. Veilleux, S. and Osterbrock, D. E.: 1987, ApJS 63, 295. Veilleux, S., Sanders, D. B., and Kim, D.-C.: 1997, ApJ 484, 92. Verma, A., Lutz, D., Sturm, E., Sternberg, A., Genzel, R., and Vacca, W.: 2003, A&A 403, 829. OBSCURED “ACTIVITY”: AGN, QUASARS, STARBURSTS AND ULIGs 407

Verma, A., Rowan-Robinson, M., McMahon, R., and Efstathiou, A. E.: 2002, MNRAS 335, 574. Verstraete, L., Pech, C., Moutou, C., Sellgren, K., Wright, C. M., Giard, M., L´eger, A., et al.: 2001, A&A 372, 981. Verstraete, L., Puget, J. L., Falgarone, E., Drapatz, S., Wright, C. M., and Timmermann, R.: 1996, A&A 315, L337. Vigroux, L.: 1997, Extragalactic Astronomy in the Infrared,p.63. Vigroux, L., Aussel, H., Charmandaris, V., Cesarsky, C., Elbaz, D., Fadda, D., et al.: 2001, Astrophys. Space Sci. Suppl. 277, 565. Vigroux, L., Charmandaris, V., Gallais, P., Laurent, O., Madden, S. C., Mirabel, I. F., et al.: 1999, ESA SP-427: The Universe as Seen by ISO,p.805 Vigroux, L., Mirabel, F., Altieri, B., Boulanger, F., Cesarsky, C., Cesarsky, D., et al.: 1996, A&A 315, L93. Voit, G. M.: 1992, ApJ 399, 495. Voit, G. M., Weymann, R. J., and Korista, K. T.: 1993, ApJ 413, 95. Wilson, C. D., Scoville, N., Madden, S. C., and Charmandaris, V.: 2000, ApJ 542, 120. Wilson, C. D., Scoville, N., Madden, S. C., and Charmandaris, V.: 2003, ApJ 599, 1049. Xu, C., Gao, Y., Mazzarella, J., Lu, N., Sulentic, J. W., and Domingue, D. L.: 2000, ApJ 541, 644. Xu, C., Sulentic, J. W., and Tuffs, R.: 1999, ApJ 512, 178. . . . AND BEYOND THE EUROPEAN LARGE AREA ISO SURVEY

SEB OLIVER1,∗ and FRANCESCA POZZI2 1Astronomy Centre, Department of Physics and Astronomy, University of Sussex, Brighton, BN1 9QH, U.K. 2Dipartimento di Astronomia, Universita di Bologna, via Ranzani 1, ID40127 Bologna, Italy (∗Author for correspondence: E-mail: [email protected])

(Received 14 July 2004; Accepted in ﬁnal form 21 October 2004)

Abstract. The European Large Area ISO Survey (ELAIS) was the largest Open Time survey on the Infrared Space Observatory (ISO). It was designed to explore obscured galaxies and hence quantify the recent star-formation history of the Universe. The ﬁnal reanalysis of the data has been completed and a band-merged catalogue with associations across many wavelengths compiled and released the data to the global astronomical community (http://astro.imperial.ac.uk/Elais/). This paper summarises some of the key results.

Keywords: infrared, galaxies, evolution, cosmology, observations

1. The ELAIS Project

The European Large Area ISO Survey (ELAIS) was the largest non-serendipitous project on ESA’s Infrared Space Observatory (ISO, Kessler et al., 1996) using ISO- CAM (Cesarskey et al., 1996) and ISOPHOT (Lemke et al., 1998). The survey covered around 12 square degrees, with ISO data at 6.7, 15, 90, and 175 μm. To reduce the effect of cosmic variance, ELAIS was split into three fields of comparable sizes, two in the north (N1 and N2) and one in the south (S1), plus six smaller area. The aims of the project and a detailed description of the observations can be found in Oliver et al. (2000). A Preliminary Analysis was undertaken to provide early scientific results and to provide follow-up source lists at 6.7 and 15 μm (Serjeant et al., 2000), 90 μm (Efstathiou et al., 2000) – the 175 μm data was processed by the FIRBACK consortium (Dole et al., 2001). To extract the maximum information from the ISO data a very careful Final Analysis of the 15 μm data (Lari et al., 2001; Vaccari et al., 2005), and at 90 μm(Héraudeau et al., 2004) was undertaken. The infrared catalogues have also been cross-correlated with each other and associated with complementary optical, NIR and radio catalogues (including La Franca et al., 2004; Gonzales-Solares et al., 2005; Väisänen et al., 2002; Väisänenand Johansson, 2004; Ciliegi et al., 1999; Gruppioni et al., 2001). The final ISO and Radio selected compilation comprises nearly 4,000 sources with Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, The Netherlands, and the United Kingdom), and with the participation of ISAS and NASA. Space Science Reviews (2005) 119: 411–423 DOI: 10.1007/s11214-005-8069-7 C Springer 2005 412 S. OLIVER AND F. POZZI limiting flux densities at 6.7, 15, 90, and 175 μmof∼1.0, 0.7, 70 and 223 mJy (Rowan-Robinson et al., 2004), these “band-merged” catalogues have been made public (http://astro.imperial.ac.uk/Elais/).

2. Motivation

The original motivation for studying galaxy evolution in the infrared comes from our understanding of star formation within our own galaxy. In these regions the intrinsically bright optical and UV emission from young massive stars is strongly affected by dust. The effects of obscuration are less signiﬁcant in the thermal infrared (1– 3 μm) and in the mid and far infrared we actually see the reprocessed emission from the dust. Thus, to properly understand star formation and hence galaxy formation we must study both optical and infrared wavelengths. Indeed, actively star forming galaxies such as M82 emit signiﬁcantly more power at far infrared wavelengths than in the optical and UV. The discovery that as much or more power in the extragalactic background emerges at far-infrared/sub-mm wavelengths as in optical/UV bands (Puget et al., 1996; Hauser and Dwek, 2001) emphasizes the cosmological importance of these far-infrared emitting populations. The unexpectedly numerous galaxies discovered by SCUBA (e.g., Hughes et al., 1998) are expected to be star forming galaxies at high redshifts where the rest-frame far-infrared galaxies has been shifted to sub-mm wavelengths. The ELAIS survey was designed to study the comparatively recent evolution of obscured star formation providing a bridge between local IRAS samples and SCUBA. It also aimed to understand the AGN phenomenon by exploring emission from the dusty tori.

3. Number Counts

The first evidence for galaxy evolution in any new survey inevitably comes from the number counts. We analysed the preliminary counts at 6.7, 15 μm (Serjeant et al., 2000), 90 μm (Efstathiou et al., 2000) and the FIRBACK team presented the 175 μm counts (Dole et al., 2001). The Final Analysis counts at 15 μminS1 and 90 μm have been discussed by Gruppioni et al. (2002) and Héraudeau et al. (2004) respectively. Whichever wavelength you choose the result is the same: the counts are inconsistent with any no evolution predictions, i.e., we detect significant evolution at all wavelengths. This is illustrated at 15 and 90 μm in Figure 1. The departure from no evolution is particularly evident at 15 μm. In this band, the large dynamic range in flux density spanned by ELAIS (0.5–100 mJy), has allowed us to sample the source counts between the flat-Euclidean IRAS counts and the deep, ISOCAM counts (Elbaz et al., 1999, 2004; Oliver et al., 1997, 2002; Aussel et al., 1999). Moreover, the large area of the ELAIS survey means it provides a very precise measurement of the shape of the counts near 1–2 mJy, the THE EUROPEAN LARGE AREA ISO SURVEY 413

Figure 1. Number counts from ELAIS and other ISO surveys. Both show significant departures from no evolution models. Left: Differential-Euclidean normalized 15 μm counts. Data points: ELAIS data from Gruppioni et al. (2002); new A2390 data from Metcalfe et al. (2003) and new Lockman data from Rodighiero et al. (2004). Other data points from Elbaz et al. (1999). Dashed, dotted, long-dashed and dot-dot-dot-dashed lines represent the model contribution for the spiral, the starburst, the Type-1 and the Type-2 AGN populations (Pozzi et al., 2004; Matute et al., 2004). The solid line is the total (adapted from Pozzi et al., 2004.) Right: Differential-Euclidean normalized 90 μm counts: ELAIS (solid symbol, Héraudeau et al., 2004); Lockman Hole (open symbols, Rodighiero et al., 2003) and IRAS PSC-z (crosses, Héraudeau et al., 2004) compared with a no-evolution model (solid line). The dashed and dotted line represent model A and E Guiderdoni et al., 1998. The dash-dotted line is the counts predicted by Rowan-Robinson (2001) (taken from Héraudeau et al. (2004)). critical flux density where the ISOCAM counts start diverging from no-evolution models. The ability to detect strong evolution in sources down to S ∼ 0.5 mJy within ELAIS has been possible, thanks to the sophisticated reduction tools based on a physical model of the detectors (Lari Method; Lari et al., 2001). These allowed us to observe at least a factor two deeper than expected.

4. Multi-Wavelength Identiﬁcation

To investigate what sort of galaxies are responsible for the excessive numbers of galaxies seen at faint fluxes requires data at other wavelengths. There have been extensive observing imaging survey campaigns or comparisons with archival data at radio (Ciliegi et al., 1999; Gruppioni et al., 2001, 2003), optical (La Franca et al., 2004; Gonzales-Solares et al., 2005; McMahon et al., 2001; Babbedge et al., 2004), NIR (Väisänen et al., 2002; Väisänenand Johansson, 2004) and X-ray wavelengths (Alexander et al., 2001; Basilakos et al., 2002). Eighty percent have optical identifications to R ∼ 23 (La Franca et al., 2004) and 92% to R ∼ 24 (Gonzales-Solares et al., 2005). Redshifts have been obtained for 500 galaxies of our infrared samples (47% of the 15 μm galaxies with identifications), primarily on AAT, ESO 3.6 m, WHT, and other smaller telescopes. The median redshift distribution is 0.17 at 15 μm and 0.1 at 90 and 175 μm. In Figure 2 (right) we show the 15 μm redshift distribution, as summarised in Rowan-Robinson et al. (2004). Besides the strong peak around 414 S. OLIVER AND F. POZZI

Figure 2. Left: Integral 15 μm number counts from the ELAIS S1 field broken down by optical classes. Notice that the starburst/AGN fraction increases at fainter fluxes and, naturally, the unidentified fraction increases sharply at the faintest fluxes. Taken from La Franca et al. (2004.) Right: Redshift distribution for ELAIS 15 μm sources. Solid curve: both photometric and spectroscopic redshifts (840 sources), dotted curve: spectroscopic redshifts only (468 sources), broken curve: effect of assigning blank fields uniformly to range 0.5 < z < 1.5. Long-broken curve: predicted redshift distribution from Rowan-Robinson (2001), dash-dotted curve: predicted redshift distribution from Pozzi et al. (2004). Predictions take into account both the identified objects and the blank fields (taken from Rowan-Robinson et al. (2004)).

0.2, a tail of sources at higher redshift, presumably the most luminous examples of the same sources which are seen in the deep surveys which have the redshift peak z ∼ 0.8 (Aussel et al., 1999; Flores et al., 1999, 2004). The optical spectra can be used to classify sources as AGN (Type-1 and Type-2), starburst galaxies and normal galaxies (La Franca et al., 2004; Pozzi et al., 2003) About 75% of the sources are starburst galaxies while ∼25% contain an AGN. About 20% of the galaxies show e(a) spectra (La Franca et al., 2004) with both weak emission (OII) and absorption (Hδ). The e(a) galaxies are also found in the deep ISO-CAM surveys (Flores et al., 1999) and are explained to be a result of selective dust extinction which affects the youngest, most massive, young stars in HII regions and they are believed to be associated with dust enshrouded starburst galaxies (Poggianti and Wu, 2000). We have dissected the number counts in S1 using these classifications (Figure 2 (left), La Franca et al., 2004). It is clear that the unidentified fraction increases as we go to fainter fluxes (where the excess over no evolution predictions becomes more extreme). Assuming that these unidentified sources are galaxies (as seen in deeper ISO surveys) rather than AGN then the galaxy fraction increases at fainter ISO fluxes. Overall we have an identified AGN fraction of around 25% (La Franca et al., 2004). The broad-band photometry from all the ISO instruments combined also provides a powerful constraint on the properties of the infrared emission processes. It appears that there are many luminous infrared galaxies with Arp220 like Spectral Energy Distributions (SEDs), but also, surprisingly, a large population of THE EUROPEAN LARGE AREA ISO SURVEY 415

11.5 luminous Lir > 10 L galaxies with cool (cirrus-type) SEDs (Rowan-Robinson et al., 2004).

5. Active Star Formation

As we have mentioned a significant fraction of the ELAIS sources have an Active Galactic Nuclei (AGN) (La Franca et al., 2004). It should be noted that the presence of an AGN does not necessarily mean the infrared luminosity is powered by the AGN; Rowan-Robinson et al. (2004) model the infrared emission for all 306 galaxies with two or more IR bands and for only 34 (11%) does the AGN luminosity dominate. So a large fraction of the ELAIS sources are forming stars. The optical photometry suggests a significant fraction of the sources have a high infrared to optical ratio and that this fraction increases as we move to fainter infrared fluxes and brighter infrared luminosities (La Franca et al., 2004; Gonzales-Solares et al., 2005; Rowan-Robinson et al., 2004). The infrared/optical ratio is explored in Figure 3. Since the infrared luminosity is roughly proportional to star-formation rate and the optical is roughly proportional to stellar mass the ratio is some measure of the star-formation activity or efficiency in a galaxy. At low LFIR this ratio is 10 roughly constant, i.e., quiescent star formation, but at about 10 L, the ratio starts increasing and is fitted as log(L15/L R) = 0.47 log(L15/L) − 5.02 with a dispersion of 0.32 dex (La Franca et al., 2004), where the total amount of star formation is related to the strength of the “burst” which bears little correlation to the total stellar mass in the host galaxy (see e.g., Figure 3, right).

Figure 3. Left: Infrared excess vs. IR correlation. Far-infrared luminosity estimated primarily from the 15 μm band and optical, from R. Over-plotted are lines of constant optical luminosity and the fit reported by La Franca et al. (2004). Right: The same information but plotted against LOpt to better illustrate the poor correlation between activity and stellar mass at high star-formation rates (figures constructed from the final band-merged ELAIS catalogue, following La Franca et al. (2004); Rowan-Robinson et al. (2004)) 416 S. OLIVER AND F. POZZI

Figure 4. Mid-IR radio correlation. Infrared radio correlation (taken from gruppioni et al. (2003).

A number of galaxies show very high rates of activity LFIR/LOpt ∼ 100 and very modest optical luminosities. Similarly (Johansson et al., 2004) identify a population of extreme mid-to-near-IR objects (EMNOs), characterized by very high MIR/NIR flux ratios. However, they argue that their selection specifically picks out obscured AGN and nearly all also appear to be ULIRGs, and some are EROs. In Figure 3, the 34 sources whose model infrared SED is dominated by an AGN are circled. As expected many of these fall in the optically luminous part of the diagram. Here there are other optically luminous sources which also presumably host AGN but which may not dominate the infrared luminosity. The infrared/optical ratio is less meaningful for any of these objects. The importance of star formation in these populations is also illustrated by an association between the 15 μm sources and the 1.4 GHz radio catalogues (Gruppioni et al., 2003) which has revealed a radio – mid IR correlation nearly as strong as the famous far infrared radio correlation (e.g., Condon, 1992). This correlation, shown in Figure 4, can be fitted as log(L1.4GHz/L) = 1.08±0.04 log(L15 μm/L)−6.7± 0.41 with a dispersion of 0.27 dex (relation found taking into account also the radio upper limits). Gruppioni et al. (2003) also estimate the contribution of the infrared starbursts to the radio source counts and show that they are responsible for 60% of the counts at S 0.05 mJy. ELAIS has been able to check the validity of the mid-IR-radio relation from the local Universe up to z ∼ 0.5; deeper surveys have found that this relation is still valid up to z ∼ 0.8–1.0 (Elbaz et al., 2002). A statistical understanding of the star-forming populations and their evolution comes from studies of the luminosity functions. We have calculated the 15 μm luminosity functions for galaxies (Pozzi et al., 2004). The infrared to optical ratio was adopted as new criterion to separate the quiescent, non-evolving population from the rapidly evolving starbursts. For the starburst population, a combination of luminosity (L ∝ (1 + z)Q, Q ∼ 3.5) and density evolution ρ ∝ (1 + z)3.8 provides the best fit, Figure 5. This model was derived from the ELAIS data but also fit THE EUROPEAN LARGE AREA ISO SURVEY 417

Figure 5. Left: z = 0 luminosity function at 90 μm from the ELAIS. The shaded area shows the ±1σ estimate from the Final Analysis of the northern ELAIS fields. Pure luminosity evolution of (1 + z)3 is assumed. Luminosities are converted to bolometric luminosities assuming νLν = constant. Also plotted is the predicted 90 μm local luminosity function from the PSC-z survey (Serjeant and Harrison, 2004). Taken from Serjeant et al. (2004). Right: Galaxy luminosity function at 15 μm. Points represent the observed space density at different redshift, while lines represent the best-fit model predictions (dashed, dotted and solid lines: starburst, normal and starburst + normal populations) data split into 0 < z < 0.2 and z > 0.2 intervals (taken from Pozzi et al. (2004)). Both plots show significant evolution. the full 15 μm number counts (Figure 1, right) and the redshift distributions from the deeper ISO-CAM surveys (Aussel et al., 1999; Flores et al., 1999, 2004). We have also calculated the 90 μm luminosity function from the Preliminary Analysis (Serjeant et al., 2001) and again from the Final Analysis with more spectroscopic redshifts in Serjeant et al. (2004), Figure 5 (right). This function is also significantly different from the redshift zero IRAS counterpart. The ELAIS ISO data has also helped in the understanding of the star-forming galaxies seen by SCUBA. Using constraints on the SEDs of 19S850 μm > 8 mJy SCUBA sources. Fox et al. (2002) concluded that all had z > 1 and half had z > 2.

6. Active Galactic Nuclei

As we have already said, SED modeling suggests around 11% of the infrared sources have SEDs dominated by emission from dusty tori around AGN. Also around 24% of the sources contain an AGN, even if it does not necessarily dominate the infrared luminosity. So it is clear that AGN are related, even if not causally, with the infrared galaxy phenomenon. It is interesting to note that the fraction of AGN seen in infrared galaxies increases with the infrared luminosity Figure 6 (left) an effect that had previously been noted in IRAS infrared galaxies at lower redshifts (Sanders and Mirabel, 1996). We have also calculated the 15 μm luminosity functions for AGN (Type-1 and 2, Matute et al., 2002, 2005). These show very signiﬁcant evolution (see Figure 7). 418 S. OLIVER AND F. POZZI

Figure 6. Left: Fraction of infrared galaxies containing (bottom, or not containing – top) an AGN as a function of infrared luminosity. Right: Infrared luminosity as a function of black hole mass. Taken from Afonso-Luis et al. (2004).

Figure 7. AGN luminosity functions at 15 μm. Points represent the observed space density at different redshift, while lines represent the best-fit model predictions. Left: objects spectroscopically classified as Type-1 AGN; Right: objects spectroscopically classified as Type-2 AGN (taken from Matute et al. (2005)).

As yet we have too few AGN to determine the luminosity function in different redshift bins, but the IRAS sample of Rush et al., (1993) gives a low-z counterpart. The Type-1 AGN evolution is consistent with pure luminosity evolution at a rate of (1 + z)Q with Q ∼ 2.6 similar to that already seen for this type of sources at other wavelengths and the infrared galaxies. A similar evolution scenario is found for Type-2 AGN with Q ∼ 2–2.6 (depending on the assumed SED) and this is the first time the shape and evolution for Type-2 AGN has been derived. Alexander et al. (2001) have investigated the 15 ELAIS sources that were also detected by ROSAT. Six of these were broad-line QSOs, four were narrow-line star-forming galaxies or Type-2 AGN, three were stars, and two were unidentified optically. They did not find any particularly hard sources, arguing against heavily obscured AGN. Basilakos et al. (2002) conducted a 2–10 keV survey over 2.5 square THE EUROPEAN LARGE AREA ISO SURVEY 419 degree in S1 using Beppo-SAX. They found 17 hard X-ray sources, 10 of which had 15 μm counterparts and eight of these have spectral classifications: six QSOs, one Type-2 AGN and one apparently normal galaxy. These observations were too faint to investigate the bulk of the ISO populations. Manners et al. (2004) have investigated the associations between 15 μm sources and deep Chandra observations in ELAIS N1 and estimate an AGN fraction of 19% with the emission from the other ISO sources consistent with star formation. A similar AGN fraction has been observed in deeper 15 μm fields using the same X/MIR correlation method (HDF-N and Lockman Hole: Fadda et al., 2002). To-date, the X-ray surveys are either too small or too shallow to address this AGN fraction definitively. Afonso-Luis et al. (2004) have investigated optically selected QSOs in the ELAIS N1 field and found that their optical properties are largely indistinguishable from the QSOs, 15 μm, selected sample. They do find an apparent correlation between the black hole mass and the infrared luminosity (Figure 6, right).

7. Hyper-Luminous Galaxies

Objects as luminous as those seen in the SCUBA surveys are very rare. The wide areal coverage and large volume (5×106 Mpc−3) of ELAIS makes it ideal for discov- ering rare objects such as these. So far we have discovered nine galaxies whose in- 13.22 frared luminosities probably exceed L > 10 L (Rowan-Robinson et al., 2004) (60% of these are based on photometric redshifts and await spectroscopic conﬁrma- tion). It is very likely that this number increases as the follow-up proceeds but this is already almost as many as have ever been discovered from the IRAS database. The bolometric emission from ﬁrst of these, ELAISP90 J1640101410502 (Morel et al., 2001), has been modeled with an almost equal contribution from star formation and accretion disk processes. Such objects may provide us with accessible “local” counterparts to the elusive SCUBA galaxies or, if they are unrelated, will provide interesting additional challenges to galaxy evolution models.

8. Large-Scale Structure

Very recently we have been able to analyze the clustering within ELAIS. In the southern ﬁeld, S1, we have estimated the angular correlation function (Figure 8, Gonzales-Solares et al., 2005). Using the median redshift z ∼ 0.2 (La Franca et al., 2004) and inverting Limber’s Equation we have estimated the 3D infrared galaxy- galaxy autocorrelation function with amplitude, encapsulated in the correlation = . +1.0 − length r0 4 3−0.6 h 1 Mpc. This is consistent with measurements from local infrared samples (r0 = 3.79 ± 0.14 at z ∼ 0.02, IRAS (Saunders et al., 1992); r0 = 3.7atz ∼ 0.02, PSC-z (Jing et al., 2002) and weaker than optical (r0 = 5.7 at z ∼ 0.05 – APM, Maddox et al., 1990; r0 = 4.92 ± 0.27 at z ∼ 0.08 – 2dFGRS, Norberg et al., 2001; r0 = 5.7 ± 0.2atz ∼ 0.18 – SDSS, Zehavi et al., 2002). 420 S. OLIVER AND F. POZZI

1.00

) 0.10 (

0.01

0.10 1.00 (deg)

Figure 8. The angular correlation function measured in our southern ﬁeld S1 (taken from Gonzales- Solares et al. (2005)).

Our results thus conform to the existing picture i.e., the clustering strength of infrared samples is less than optically selected samples. This result appears to contradict that presented by Elbaz et al. (2004) who found that in deep infrared surveys the infrared galaxies were more strongly clustered that the optical galaxies. However, this may suggest that within the ISO surveys we have already detected the evolution of clustering. Within the ELAIS survey itself Vaisanen and Johansson (2004) find a significant over-density of EROs around faint ELAIS mid-IR sources. A radial distribution of the EROs suggests a physical connection of these EROs and ISOCAM sources, which also suggests that the faintest ELAIS sources are more strongly clustered. This would be qualitatively consistent with the predictions of hierarchical models where star formation occurs first in the strongly clustered peaks of the density field and takes longer to initiate in the lower density, more weakly clustered “field.”

9. The ELAIS Legacy

The ELAIS fields were carefully chosen to avoid contamination from dust emission from our own Galaxy (Oliver et al., 2000). Coupled with the ever expanding multiwavelength data available the ELAIS fields are a prime target for new surveys. A number of surveys and projects have been carried out whose goals are relatively independent of the ISO observations themselves. Notable amongst these are the SCUBA 8 mJy (Scott et al., 2002) and MAMBO surveys (Greve et al., 2004) and X-ray surveys (Willott et al., 2003, 2004; Manners et al., 2003). These fields are particularly ideal for Spitzer and the ISO 15 μm band data (inaccessible to Spitzer) will be extremely valuable. For example the GOODS validation field (Chary et al., 2004) was executed in ELAIS fields. More significantly, three of the ELAIS fields are core components of the SWIRE fields (Lonsdale et al., 2003, 2004). THE EUROPEAN LARGE AREA ISO SURVEY 421

10. Discussion

We have demonstrated that infrared emitting galaxies evolve strongly. The indications are that the star forming galaxies make up the bulk of the excess over no evolution predictions. Low star formation rate (infrared luminosity) systems have similar infrared/optical ratios which is a measure of activity of the star formation (star-formation rate per unit stellar mass) so their total star-formation rate is determined by the total stellar mass. However, at high star-formation rates there is a correlation between the infrared/optical excess and the infrared luminosity (total star-formation rate) and a corresponding lack of correlation with the optical luminosity (total stellar mass), indeed very high infrared/optical excess galaxies are found with very modest optical luminosities. This suggests that active star formation is associated with factors external to the galaxy. AGN are found in 20–25% of the sources, but star formation may still be the dominant source of the infrared luminosity in all but 11% of the sources. The AGN and infrared phenomena are however, closely linked: the AGN fraction increases at higher luminosity; the infrared luminosity functions evolve at a similar rate; and the infrared luminosity is correlated to the black hole mass. This may imply a common factor in the AGN and star-formation phenomena. At the extreme end we have discovered nine hyper luminous galaxies, comparable with the total number discovered by IRAS. We detect clustering with an amplitude consistent with lower redshift infrared samples and lower than optical samples, this contrasts with the higher clustering of infrared galaxies in a deeper sample suggesting that an evolution of clustering may have been detected in the ISO surveys. It seems quite plausible that the environment provides the external factor determining active star formation and the common factor linking star formation and AGN activity, thus we might expect the evolution in the luminosity functions to be linked to an evolution in the clustering. The ELAIS samples will provide invaluable templates for understanding the deeper samples detected by the Spitzer GTO and legacy surveys (e.g., Lonsdale et al., 2003; Dickinson et al., 2004) and the ELAIS ﬁelds themselves provide a valuable legacy for surveys with Spitzer and other facilities. The complete ELAIS survey comprising a band-merged catalogue with multi-wavelength associations is now publicly available at http://astro.imperial. ac.uk/Elais/

Acknowledgements

The authors would like to thank Michael Rowan-Robinson and the whole ELAIS team for their enormous efforts in making the ELAIS survey a reality and “we” in the text refers to the ELAIS team. We would like to thank Phillippe H´eraudeau,Israel 422 S. OLIVER AND F. POZZI

Matute, Alejandro Afonso Luis and Evanthia Hatziminaoglou for providing figures in advance of publication and Petri Väisänenfor useful comments. This chapter was based in part on an original conference proceeding (Oliver et al., 2004).

References

Afonso-Luis, A., et al.: 2004, MNRAS 354, 961. Alexander, D., et al.: 2001, AJ 554, 18. Aussel, H., et al.: 1999, A&A 342, 313. Babbedge, T., et al.: 2004, MNRAS 353, 654. Basilakos, S., et al.: 2002, MNRAS 331, 417. Cesarskey, C., et al.: 1996, A&A 315, L32. Chary, R., et al.: 2004, ApJS 154, 80. Ciliegi, P., et al.: 1999, MNRAS 302, 222. Condon, J.: 1992, AnRA&A 30, 575. Dickinson, M., et al.: 2004, in A. Renzini and R. Bender (eds.), Proceedings of Multiwavelength Mapping of Galaxy Formation and Evolution Conference, Venice (Italy), October 2003. Dole, H., et al.: 2001, A&A 372, 364. Efstathiou, A., et al.: 2000, MNRAS 316, 1169. Elbaz, D., et al.: 2002, A&A, 384, 848. Elbaz, D., et al.: 1999, A&A 351, L37. Elbaz, D., et al.: 2004, in A. Renzini and R. Bender (eds.) Proceedings of Multiwavelength Mapping of Galaxy Formation and Evolution Conference, Venice (Italy), October 2003, astro-ph/0403209. Fadda, D., et al.: 2002, A&A 383, 838. Flores, H., et al.: 1999, ApJ 517, 148. Flores, H., et al.: 2004, A&A 415, 885. Fox, M., et al.: 2002, MNRAS 331, 839. Gonzales-Solares, E., et al.: 2004, MNRAS 352, 44. Gonzales-Solares, E., et al.: 2005, MNRAS 358, 333. Greve, T., et al.: 2004, MNRAS 354, 779. Gruppioni, C., et al.: 2001, MNRAS 335, 297. Gruppioni, C., et al.: 2002, MNRAS 335, 831. Gruppioni, C., et al.: 2003, MNRAS 341 1L. Guiderdoni, B., et al.: 1998, MNRAS 295, 877. Héraudeau, P., et al.: 2004, MNRAS 354, 924. Hauser, M. G., and Dwek, E.: 2001, ARA&A 39, 249. Hughes, D., et al.: 1998, Nature 394, 241. Jing, Y. P., Börner, G., and Suto, Y.: 2002, ApJ 564, 15. Johansson, P. H., Väisänen,P., and Vaccari, M.: 2004, A&A 427, 795. Kessler, M., et al.: 1996, A&A 315, L27. La Franca, F., et al.: 2004, AJ 127, 3075. Lari, C., et al.: 2001, MNRAS 325, 1173. Lemke, D., et al.: 1998, A&A 315, L64. Lonsdale, C., et al.: 2003, PASP 115, 897. Lonsdale, C., et al.: 2004, ApJS 154, 54. Maddox, S. J., et al.: 1990, MNRAS, 242, 43P. Manners, J., et al.: 2003, MNRAS 343, 293. Manners, J., et al.: 2004, MNRAS 355, 97. Matute, I., et al.: 2002, MNRAS 332, L11. THE EUROPEAN LARGE AREA ISO SURVEY 423

Matute, I., et al.: 2005, in preparation. McMahon, R. G., et al.: 2001, New Astron Rev 45(Issue 1–2), 97. Metcalfe, L., et al.: 2003, A&A 407, 791. Morel, T., et al.: 2001, MNRAS 327, 1187. Norberg, P., Baugh, C. M., Hawkins, E., et al.: 2001, MNRAS 328, 64. Oliver, S., et al.: 1997, MNRAS 289, 471. Oliver, S., et al.: 2000, MNRAS 316, 749. Oliver, S., et al.: 2002, MNRAS 332, 356. Oliver, S., et al.: 2004, in A. Renzini and R. Bender (eds.), Proceedings of Multiwavelength Mapping of Galaxy Formation and Evolution Conference, Venice (Italy), October 2003. Poggianti, B. and Wu, H.: 2000, ApJ, 529, 157. Pozzi, F., et al.: 2003, MNRAS 343, 1348. Pozzi, F., et al.: 2004, ApJ 609 122. Puget, J.-L., et al.: 1996, A&A 308L,5. Rodighiero, G., et al.: 2003, MNRAS 343 1155. Rodighiero, G., et al.: 2004, A&A 427, 773. Rowan-Robinson, M., et al.: 2004, MNRAS 351, 1290. Rowan-Robinson, M.: 2001, ApJ 549 745. Rush, B., Malkan, M. A., and Spinoglio, L.: 1993, ApJS, 89,1. Sanders, D. B. and Mirabel, I. F.: 1996, AnRA&A 34, 749. Saunders, W., Rowan-Robinson, M., and Lawrence, A.: 1992, MNRAS 258, 134. Scott, S., et al.: 2002, MNRAS 331, 817. Serjeant, S., et al.: 2000, MNRAS 316, 768. Serjeant, S., et al.: 2001, MNRAS 322, 262. Serjeant, S., et al.: 2004, MNRAS 355, 813. Serjeant, S. and Harrison, D.: 2004, in M. Plionis (ed.) Proceeding of the “Multi-Wavelength Cosmology” Conference, Mykonos, Greece, June 2003 (Kluwer). Väisänen,P., et al.: 2002, MNRAS 337, 1043. Väisänen,P. and Johansson, P. H.: 2004, A&A 421, 821. Väisänen,P. and Johansson, P. H.: 2004, A&A 422, 453. Vaccari, M., et al.: 2005, MNRAS 358, 397. Willott, C., et al.: 2003, MNRAS 339, 297. Willott, C., et al.: 2004, ApJ 610, 140. Zehavi, I., Blanton, M. R., Frieman, J. A., et al.: 2002, ApJ 571, 172. ISO’s CONTRIBUTION TO THE STUDY OF CLUSTERS OF GALAXIES∗

LEO METCALFE1,∗, DARIO FADDA2 and ANDREA BIVIANO3 1XMM-Newton Science Operations Centre, European Space Agency, Villafranca del Castillo, P.O. Box 50727, 28080 Madrid, Spain 2Spitzer Science Center, California Institute of Technology, Mail code 220-6, 1200 East California Boulevard, Pasadena, CA 91125, U.S.A. 3INAF – Osservatorio Astronomico di Trieste, via G.B. Tiepolo 11, 34131, Trieste, Italy (∗Author for correspondence: E-mail: [email protected])

(Received 10 November 2004; Accepted in ﬁnal form 20 December 2004)

Abstract. Starting with nearby galaxy clusters like Virgo and Coma, and continuing out to the furthest galaxy clusters for which ISO results have yet been published (z = 0.56), we discuss the development of knowledge of the infrared and associated physical properties of galaxy clusters from early IRAS observations, through the “ISO-era” to the present, in order to explore the status of ISO’s contribution to this ﬁeld. Relevant IRAS and ISO programmes are reviewed, addressing both the cluster galaxies and the still-very-limited evidence for an infrared-emitting intra-cluster medium. ISO made important advances in knowledge of both nearby and distant galaxy clusters, such as the discovery of a major cold dust component in Virgo and Coma cluster galaxies, the elaboration of the correlation between dust emission and Hubble-type, and the detection of numerous Luminous Infrared Galaxies (LIRGs) in several distant clusters. These and consequent achievements are underlined and described. We recall that, due to observing time constraints, ISO’s coverage of higher-redshift galaxy clusters to the depths required to detect and study statistically signiﬁcant samples of cluster galaxies over a range of morphological types could not be comprehensive and systematic, and such systematic coverage of distant clusters will be an important achievement of the Spitzer Observatory.

1. Introduction

1.1. ISO LOOKS DEEP

The most strongly star forming galaxies are heavily dust obscured, and estimates of their star formation rates (SFR) made at visual wavelengths often fall one or two orders of magnitude below their true values. Frequently, very active star forming galaxies occur in associations or groups. At the same time, it is believed that in the dense environments of galaxy clusters the interactions of galaxies with each other, with the cluster tidal ﬁeld, and with the intra-cluster medium (via

∗Based on observations with ISO, an ESA project with instruments funded by ESA Member States (especially the PI countries: France, Germany, the Netherlands, and the United Kingdom) and with the participation of ISAS and NASA. Space Science Reviews (2005) 119: 425–446 DOI: 10.1007/s11214-005-8065-y C Springer 2005 426 L. METCALFE ET AL. ram pressure) strip galaxies of their reserves of gas, and eventually suppress star formation. Much has already been learned about the evolution of the cosmic star formation rate in field galaxies from observations with the European Space Agency’s Infrared Space Observatory (ISO) satellite (Kessler et al., 1996). Thanks to deep surveys in the mid-infrared (MIR) and far-infrared (FIR) (e.g. Elbaz et al., 1999, 2002; Serjeant et al., 2000, 2004; Lari et al., 2001; Gruppioni et al., 2002; Metcalfe et al., 2003; Rodighiero et al., 2003; Sato et al., 2003; Kawara et al., 2004; Rowan- Robinson et al., 2004; and several others) conducted, respectively, with ISOCAM1 (Cesarsky et al., 1996) and ISOPHOT (Lemke et al., 1996), we now know that the comoving density of IR-bright galaxies has a very rapid evolution from z ∼ 0 to z ∼ 1. This evolution has been interpreted as the result of an increased rate of galaxy–galaxy interactions, coupled with an increase in the gas content of the galaxies (Franceschini et al., 2001). Dust obscuration plays an important role in concealing star formation in field galaxies. Its role in relation to the star formation activity of cluster galaxies is less known. Published ISO results to-date addressing the fields of galaxy clusters (Pierre et al., 1996; Lémonon et al., 1998; Altieri et al., 1999; Barvainis et al., 1999; Quillen et al., 1999; Soucail et al., 1999; Fadda et al., 2000; Duc et al., 2002, 2004; Metcalfe et al., 2003; Biviano et al., 2004; Coia et al., 2005a,b) concern a dozen clusters at z < 0.6, and point to an important role for dust, evolving with redshift, also in cluster galaxies. The fraction of IR-bright, star-forming cluster galaxies changes significantly from cluster to cluster, without a straightforward correlation with either the redshift, or the main cluster properties (such as the cluster mass and luminosity). The presence of dust in cluster galaxies may have hampered our efforts to fully understand the evolution of galaxies in clusters. Optical-band observations have provided a few, but very fundamental results, which cannot yet however be combined in a unique, well constrained scenario of cluster galaxy evolution.

1.2. SOME PROPERTIES OF CLUSTER GALAXIES

Perhaps the most fundamental phenomenology that all scenarios of cluster galaxy evolution (e.g. Dressler, 2004) must explain is the so called morphology density relation (hereafter MDR; Dressler, 1980), whereby early-type galaxies, i.e. ellipticals and S0s, dominate rich clusters, while late-type galaxies, i.e. spirals and irregulars, are more common in the ﬁeld. The fact that early-type galaxies seem to reside in the cluster centres since z ≥ 1–2, while the colour-magnitude relation (CMR, Visvanathan and Sandage, 1977) remains very tight even at z ∼ 1 suggests that the MDR is established at the formation of galaxy clusters and that the early-type

1Throughout this paper the ISOCAM filters having reference wavelengths 6.75 and 14.3 μm will respectively be referred to as the 7 and 15 μm filters. Both bandpasses are referred to as mid-infrared (MIR). ISO’s CONTRIBUTION TO THE STUDY OF CLUSTERS OF GALAXIES 427 galaxies defining the CMR are uniformly old and passively evolving since their formation redshift, zf > 2 (e.g. Ellis et al., 1997). Similar conclusions are obtained by analysing the fundamental plane (FP; Dressler et al., 1987), relating basic properties of early-type galaxies (their effective radius, internal velocity dispersion, and effective surface brightness). The FP, like the CMR, still holds for early-type galaxies in z ∼ 1 clusters, and its scatter is similar to that seen in nearby clusters (van Dokkum and Stanford, 2003). However, these conclusions could only be true on average. Independent analyses suggest that at least part of the cluster galaxies have undergone significant evolution over the last 3–8 Gyr. First and foremost is the observational evidence for an increasing fraction of blue cluster galaxies with redshift, the so called ‘Butcher- Oemler’ (BO) effect (Butcher and Oemler, 1978, 1984; Margoniner et al., 2001). Approximately 80% of galaxies in the cores of nearby clusters are ellipticals or S0s, i.e. red galaxies (Dressler, 1980), but the fraction of blue galaxies increases with redshift. These blue galaxies are typically disk systems with ongoing star formation (Lavery and Henry, 1988), with spectra characterized by strong Balmer lines in absorption (typically, EW(Hδ) > 3 A)˚ and no emission lines, and have been named ‘E+A’ (or also ‘k+a’) galaxies (Dressler and Gunn, 1983). Modelling of their spectra indicates that star formation stopped typically between 0.05 and 1.5 Gyr before the epoch of observation, in some cases after a starburst event (Poggianti et al., 1999, 2001). Similarly, the fraction of spirals increases with redshift (Dressler et al., 1997; Fasano et al., 2000; van Dokkum et al., 2001), at the expense of the fraction of S0s. Maybe, the colour and the morphological evolution of the cluster galaxy population are two aspects of the same phenomenon. Field spirals are being accreted by clusters (Tully and Shaya, 1984; Biviano and Katgert, 2004), and the accretion rate was higher in the past (Ellingson et al., 2001). It is therefore tempting to identify the blue galaxies responsible for the BO-effect in distant clusters with the recently accreted field spirals. Possibly, these evolutionary trends could be reconciled with the passive evolution inferred from the CMR and FP studies by taking into account the so-called ‘progenitor bias’ (van Dokkum et al., 2000), namely the fact that only the most evolved among cluster early-type galaxies are indeed selected in studies of the colour-magnitude and fundamental plane relations.

1.3. A HOSTILE ENVIRONMENT

There is no shortage of plausible physical mechanisms that could drive the evolution of a galaxy in a hostile cluster environment. Among these, the most popular today are ram pressure, collisions, and starvation from tidal stripping (Dressler, 2004), all in principle capable of depleting a spiral of its gas reservoir, thereby making it redder and more similar to a local S0. 428 L. METCALFE ET AL.

Ram pressure from the dense intra-cluster medium can sweep cold gas out of the galaxy stellar disk (Gunn and Gott, 1972) and induce star formation via compression of the gas that remains bound to the galaxy. Collisions or close encounters between galaxies generate tidal forces that tend to funnel gas towards the galaxy centre (Barnes and Hernquist, 1991) eventually fueling a starburst that ejects gas from the galaxy. The cumulative effect of many minor collisions (named ‘harassment’, Moore et al., 1996), can lead to the total disruption of low surface brightness galaxies (Martin, 1999). The collision of a group with a cluster can also trigger starbursts in cluster galaxies, as a consequence of the rapidly varying tidal field (Bekki, 2001). Finally, the so called ‘starvation’ mechanism (Larson et al., 1980) affects a galaxy’s properties by simply cutting off its gaseous halo reservoir. This can occur because of tidal stripping, a mechanism effective in galaxy–galaxy encounters, but also when galaxies pass through the deep gravitational potential well of their cluster. The common outcome of all these processes is galaxy gas depletion, ultimately leading to a decrease of the star formation activity for lack of fuel, and, hence, to a reddening of the galaxy stellar population. However, some of these processes induce a starburst phase before the gas depletion, and some do not. To date, it remains unclear which physical process dominates in the cluster environment. Useful constraints can be obtained by finding where the properties of cluster galaxies change with respect to the field, since the different processes become effective at different galaxy or gas densities. Recently it has been found (Kodama et al., 2001; Lewis et al., 2002; Gómez et al., 2003) that a major change in the star-formation properties of cluster galaxies occurs in the outskirts of clusters (at ∼1.5 cluster virial radii). These results would seem to exclude ram-pressure stripping as a major factor in cluster galaxy evolution, since the density of the intra-cluster medium is too low in the cluster outskirts. Recent observations of high-z clusters seem to have complicated, rather than simplified, the issue of cluster galaxy evolution. A surprisingly high fraction of red merger systems has been found in these distant clusters (van Dokkum et al., 2001). The red colours of these merging galaxies and the lack of emission lines in their spectra suggest that their stellar populations were formed well before the merger events, but the occurrence of relatively recent starburst events in these galaxies is instead suggested by detailed analyses of their spectra (Rosati, 2004).

1.4. AN UNCLUTTERED VIEW

A better understanding of the evolutionary processes affecting cluster galaxies can come from observations in the infrared. Dust, if present, is capable of obscuring most of a galaxy’s stellar radiation, making the observed galaxy red and dim at optical wavelengths, and affecting optical estimates of the galaxy star formation activity. The effects of dust can be particularly severe if the galaxy is undergoing a starburst (Silva et al., 1998), so that we might be missing a substantial part of the ISO’s CONTRIBUTION TO THE STUDY OF CLUSTERS OF GALAXIES 429 evolutionary history of cluster galaxies by observing them at optical wavelengths. Since the dust-reprocessed stellar radiation is re-emitted at IR wavelengths, the IR luminosity is a much more reliable indicator of a galaxy’s star formation activity (Elbaz et al., 2002). The plan of this review is as follows: in Section 2 the development of knowledge of the infrared properties of galaxy clusters from early IRAS observations, through the “ISO-era” to the present is described. Section 2.1 considers the accumulation of data on the Virgo cluster, while Section 2.2 addresses other nearby clusters, e.g. the Fornax, Hydra, Coma and Hercules clusters. Section 2.3 discusses the significant progress that has been possible with ISO in the study of cluster galaxy properties out to moderate redshifts (<0.6) and attempts to draw some comparison among the still rather heterogeneous sample of cluster observations. Section 2.4 reviews the status of attempts to directly observe diffuse intra-cluster dust in the infrared. Finally Section 3 summarises the current status of the field and remarks on the important opportunity represented by the Spitzer Observatory to decisively extend the field.

2. Cluster Galaxies in the Infrared

2.1. THE VIRGO GALAXY CLUSTER

Being the most nearby relatively massive galaxy cluster, the Virgo cluster has been studied extensively at all wavelengths. Already in 1983, Scoville et al. obtained 10 μm data with IRTF for 53 Virgo spiral galaxies and concluded that star formation is occurring in the nuclei of virtually all spiral galaxies independent of spectral type, with an average star formation rate (SFR) of 0.1M/yr. No correlation was found between 10 μm emission and gross properties (barred or normal, early or late spiral morphology, total optical luminosity), nor with location in the cluster. Using far- to mid-infrared flux correlations, they inferred a far-infrared (FIR) luminosity 10 ∼2 × 10 L for the brightest 10 μm source in the Virgo cluster, NGC4388. Leggett et al. (1987) made optical identifications of 145 IRAS galaxies in the 113 square degree field centred on the Virgo cluster and concluded that they were mostly seeing the spirals, and that the infrared properties of the Virgo cluster galaxies are indistinguishable from those of field galaxies at similar redshift. Virgo spirals were indistinguishable from field disc galaxies with normal SFRs. The typical infrared 9 luminosity for the sample was LIR ∼ 10 L, and the most luminous confirmed 10 cluster sources were found to have LIR = afew× 10 L. They concluded that the cluster environment has no effect on galaxy IR properties, even for very HI deficient galaxies. Their conclusion was however criticized by Doyon and Joseph (1989) who, using a sample of 102 Virgo spirals detected at 60 and 100 μm, were able to show that HI-deficient galaxies have lower IR fluxes, lower star formation activity, and cooler IR colour temperatures than those with normal HI content. 430 L. METCALFE ET AL.

If cluster environment affects the interstellar medium (ISM) of cluster galaxies, one might expect the FIR, radio and FIR–radio correlation to be affected. Niklas et al. (1995) looked at the FIR–radio correlation in Virgo galaxies and suggested that, in this respect, most of the Virgo galaxies behave like normal field galaxies. Only a few early type spirals with strong central sources as well as very disturbed galaxies show a high radio excess. These are in the inner part of the cluster, where galaxies have disturbed HI-distributions, and truncated HI-discs, as shown by 21 cm observations (Cayatte et al., 1990). Measurements of molecular gas emission in radio continuum (Kenney and Young, 1986, 1989) show a much less pronounced stripping effect. In the ISO era, Tuffs et al. (2002) and Popescu et al. (2002a,b) used ISOPHOT at 60, 100 and 170 μm, to study a luminosity- and volume-limited sample of 63 S0a or later-type galaxies in the core and periphery of the Virgo cluster. They reached sensitivities 10 times better than IRAS in the two shorter-wavelength bandpasses, and the confusion limit at 170 μm. These programmes sought to extend knowledge of FIR SEDs to lower limits covering a complete sample of normal2 late-type galaxies over a range of morphological type and star formation activity. A significant cold dust component (with a temperature of around 18 K) was found in all morphological classes of late-type galaxies, from early giant spirals to irregular galaxies and Blue Compact Dwarfs (BCDs), and which could not have been recognised by IRAS. These results required a revision of the masses and temperatures of dust in galaxies. On average, dust masses are raised by factors of 6–13 with respect to IRAS results. The FIR/radio correlation is confirmed for the warm FIR dust, and is found for the cold dust. The predominance of the very cold dust component (down to less than 10 K) in BCDs was remarkable, with a few 10s of percent of the UV/Optical component appearing in the cold dust emission. The cold emission might be due to collisional dust heating in dust swept up in the intergalactic medium (proto-galactic cloud) by a galactic wind from the BCD or, alternatively, to photon heating of the dust particles in an optically thick disk, indicative of a massive gas/dust accretion phase that makes BCDs sporadically bright optical/UV sources when viewed out of the disk equatorial plane. On average, 30% of the stellar light of spirals in Virgo is re-radiated by dust, with a strong dependence on morphological type, ranging from 15% for early spirals to 50% for some late spirals, and even more for some BCDs (Popescu and Tuffs, 2002). Figure 1, taken from Popescu and Tuffs (2002), shows the dependency of the dust contribution on the Hubble type in the Virgo cluster, and it illustrates a sequence of increasing FIR-to-total bolometric output running from normal to gas-rich-dwarf galaxies. Leech et al. (1999) studied a sample of 19 Virgo cluster spiral galaxies with ISO’s Long Wavelength Spectrometer (LWS) (Clegg et al., 1996) obtaining spectra

2A ‘Normal’ galaxy is understood to mean a galaxy not dominated by an active nucleus, with SFR sustainable for a substantial fraction of a Hubble time. ISO’s CONTRIBUTION TO THE STUDY OF CLUSTERS OF GALAXIES 431

Figure 1. From Popescu and Tuffs (2002): The ratio of dust luminosity to total bolometric luminosity as a function of Hubble Type, for Virgo cluster galaxies. Filled symbols are cluster core galaxies and open symbols refer to the cluster periphery. A correlation is established between Hubble type and the FIR contribution to bolometric luminosity. around the [CII] 157.7 μm fine-structure line for 14 of them. [CII] line radiation provides the most important gas cooling mechanism in normal, i.e. non-starburst, late-type galaxies, balancing photoelectric heating from grains. In field galaxies [CII] is typically 10−3 to 10−2 of total galaxy FIR. The sample, drawn from both the cluster core and the cluster periphery (and being a sub-sample of the Tuffs et al., 2002 sample), spanned the S0a–Sc morphological range to probe any difference in the [CII] emission between different types or between core and periphery galaxies. A good correlation was found between the strength of the [CII] line and the FIR flux, as measured by IRAS. Moreover, the [CII]-to-K -band flux ratio shows a two order of magnitude difference between early-type galaxies and late spiral types. Galaxies with large [CII]/FIR ratios tend to have later Hubble types. No apparent relationship was found between [CII] strength and galaxy position in the cluster, nor between [CII] and HI-mass surface density. Any influence of the Virgo cluster environment on the [CII] emission was found to be small compared with the strong dependence of the line emission on basic measurables such as morphology or bulk mass of the stellar component, as measured by the near-IR (K -band) luminosity. In a series of papers, Boselli et al. (1997a,b, 1998, 2003a,b, 2004) explored the properties of a large sample of ISO-detected spiral and irregular galaxies in the Virgo cluster, in measurements made with ISOCAM at 7 and 15 μm. 71 objects were in the cluster periphery (≥4 degrees from M87) and 28 in the core (≤2 degrees from M87) in order to allow study of the effects of the environment on dust emission and evolution. A further 24 Virgo cluster galaxies were serendipitously observed, bringing the full Virgo sample up to 123 galaxies. S0 and elliptical galaxies observed by chance in the ISO fields were used to establish the stellar contribution to the IR 7.4 10.1 emission. Objects had luminosities in the range 10 ≤ LFIR ≤ 10 L. 432 L. METCALFE ET AL.

Thirty four of the Virgo objects were fully resolved by ISOCAM, and MIR images could be presented for these, along with radial light and colour profiles and other morphological and structural information. Boselli et al. showed that the MIR emission of optically-selected, normal early- type galaxies is dominated by the Rayleigh–Jeans tail of the cold stellar component, while that of late-type galaxies is dominated by the thermal emission from dust, but partly contaminated by stellar emission, especially at 7 μm in early-type spirals (Sa). While the MIR emission (per unit mass) of the spirals and later systems are comparable, the average 7 and 15 μmtoK flux ratio of E, S0 and S0/a galaxies is significantly lower than that of spirals. At 15 μm, where the difference is clearer, spirals have on average a MIR emission per unit mass higher by more than one order of magnitude than E-S0/a. BCDs have, on average, a MIR emission per unit mass comparable to that of spirals. The IR emission carriers and their behaviour are consistent with quantitative expectations based on MIR studies of the ISM in our Galaxy. In spiral and irregular galaxies the 7 μm emission is almost entirely due to the UIB3 carriers. The MIR fluxes are proportional to the SFR when it is not too large, but fall off in the presence of a high SFR suggesting that the UIBs are destroyed by the UV field. However, in galaxies with a high SFR, there is an additional diffuse contribution (i.e. not well correlated with HII regions or Hα sources) to the 15 μm flux from very small, three-dimensional, grains. As a consequence, MIR dust emission is not an optimal tracer of star formation in normal, late-type galaxies, and MIR luminosities are better correlated with FIR luminosities than with more direct tracers of the young stellar population such as the Hα and the UV luminosity. This conclusion, valid for nearby normal, late-type galaxies, may not apply to luminous starburst galaxies (Luminous Infrared Galaxies, or LIRGs), such as those detected in the ISO deep surveys (Förster-Schreiber et al., 2004). The MIR emission traces well the FIR and bolometric emission (Boselli et al., 1998; Elbaz et al., 2002).

2.2. THE COMA,FORNAX,HYDRA,HERCULES &OTHER NEARBY (z < 0.1) CLUSTERS

Located at much larger distance than Virgo, but with a much larger mass, Coma has often been a favourite observational target, also in the IR. Wang et al. (1991) made optical identifications of a total of 231 IRAS point sources in the regions of the Fornax, Hydra and Coma clusters, and identified respectively 13, 29 and 26 cluster galaxies. They concluded that the cluster environment has no detectable influence on galaxy infrared properties. Bicay and Giovanelli (1987) came to the same conclusion, based on a sample of 200 FIR-emitting galaxies

3The Unidentified Infrared Bands (UIB) dominate the 5–12 micron MIR spectrum of a wide range of celestial sources, and are usually assumed to arise from polycyclic aromatic hydrocarbon molecules (PAHs). ISO’s CONTRIBUTION TO THE STUDY OF CLUSTERS OF GALAXIES 433 in seven nearby clusters (including Coma). However, they remarked that LIRGs were much less common in the clusters than in the field (see also Section 2.3). The most distant cluster studied with IRAS was A2151 (aka Hercules, at z = 0.036). Out of 41 sources detected at 60 μmina1.6 × 2.5 sq.deg. field, Young et al. (1984) correlated 24 with late-type spiral galaxies of the cluster remarking the total absence of IR emission from E and S0 galaxies. Odenwald (1986) found CO emission in 3 of the 9 most optically luminous galaxies in the spiral-rich Ursa Major I(S) galaxy group, seeking evidence for galaxy interactions in a cluster lacking an intra-cluster medium, but found little evidence that processes unrelated to the galaxies themselves had influenced their histories. Studying a sample of 200 galaxies in seven nearby clusters,4 Bicay and Giovanelli (1987) pointed out 11 the absence of luminous IR galaxies (LIRGs, i.e. galaxies with LFIR > 10 L). 10 The sample consisted almost entirely of IR normal galaxies (LFIR < 10 L)in contrast to the rather high percentage of LIRGs (20%) detected by IRAS in the field. Moreover, the lack of a strong correlation between galaxy HI content and IR emission led them to conclude that SFR is not enhanced by interaction with the ICM, and might even be quenched by it. On the contrary, the suppressed FIR- to-radio ratio of spiral galaxies found in rich clusters with respect to poor clusters (Andersen and Owen, 1995) seemed to suggest that ram pressure enhances the radio emission in rich clusters while galaxy–galaxy interactions play a more important role in poor clusters where velocity dispersion, and so encounter velocities, are smaller. Quillen et al. (1999) observed 7 E+A galaxies plus one emission-line galaxy at 12 μm with ISOCAM. They found that E+A galaxies have mid- to near-IR flux ratios typical of early-type quiescent galaxies, while the emission-line galaxy had enhanced 12 μm emission relative to the near-IR. Galaxies with ongoing star formation have a different velocity distribution in the cluster from galaxies with stopped SF, suggesting that the ongoing infall of field spirals into the cluster potential may first trigger and then quench star formation. Further observations of Coma cluster galaxies in the MIR came from Boselli et al. (1998) and Contursi et al. (2001). They also observed the cluster A1367, located in the Coma supercluster, detecting, in total, 18 spiral/irregular galaxies in the MIR and FIR with ISO. Confirming results found in Virgo galaxies, these authors concluded that most IR-detected Coma galaxies display diffuse MIR emission unrelated to their Hα emission. The aromatic carriers are not only excited by UV photons, but also by visible photons from the general ISM. When the UV radiation field is too intense, it can even destroy the aromatic carriers, and overall the MIR emission is dominated by photo-dissociation regions rather than HII-like regions. A cold dust component was detected in all galaxies, at a temperature of ∼22 K, more extended than the warm dust. Only a very weak trend was found between the total dust mass and the gas

4A262, Cancer, A1367, A1656 (Coma), A2147, A2151 (Hercules), and Pegasus. 434 L. METCALFE ET AL. content of the galaxies, even if some galaxies are very HI-deficient, and there was no detection of any relation between the MIR/FIR properties and the environment. All these results seemed to suggest very little (if any) dependence of the IR properties of galaxies on the environment. If anything, IR emission was thought to be quenched in the cluster environment. However, only nearby clusters had been studied, in which most of the galaxies are early type with little gas (and dust). With the launch of ISO it became possible for the first time to study galaxy clusters out to redshifts where a significant change in the composition of the cluster galaxy population had already been (Butcher and Oemler, 1984) or was soon to be (Dressler et al., 1999) observed in the optical.

2.3. GALAXY CLUSTERS AT INTERMEDIATE REDSHIFT

ISO’s mid-infrared camera, ISOCAM, with its vastly improved sensitivity and spatial resolution with respect to IRAS, has successfully observed several galaxy clusters out to redshifts at which signiﬁcant evolution might be expected to occur. At the same time, while studies of galaxies in nearby clusters are frequently done by targeting galaxies selected at other wavelengths, distant clusters can be completely surveyed due to their smaller angular size, producing an unbiased sample of infrared-emitting galaxies. Published ISO observations to date for clusters at redshifts above 0.1 (Pierre et al., 1996; L´emonon et al., 1998; Altieri et al., 1999; Barvainis et al., 1999; Fadda et al., 2000; Duc et al., 2002, 2004; Metcalfe et al., 2003; Biviano et al., 2004; Coia et al., 2005a,b) address seven clusters (see Table I)

TABLE I Clustersa in the redshift range 0.17 < z < 0.6 studied with ISOCAM.

No. 15 μm-only No. 7 μm-only No. 7-and-15 μm No. all Cluster z sources sources sources sources

A2218 0.175 5 18 4 27 A1689 0.181 3 20 9 32 A1732 0.193 – 4 0 4 A2390 0.23 1 11 3 15 A2219 0.228 3 – – 3 A370 0.37 1 5 0 6 Cl0024+1654 0.39 13 – – 13 J1888.16CL 0.56 6 – – 6

aThe content of the columns in the table is as follows: name and redshift of the cluster; number, respectively, of 15 μm-only, 7 μm-only and 7-and-15 μm confirmed cluster sources detected in each case, and total number of MIR sources detected. Results are gathered from Coia et al. (2005a,b), Biviano et al. (2004), Metcalfe et al. (2003), Duc et al. (2002, 2004) and Fadda et al. (2000). ISO’s CONTRIBUTION TO THE STUDY OF CLUSTERS OF GALAXIES 435 spanning the redshift range 0.17 < z < 0.6 and yield MIR data for around 110 cluster galaxies, slightly over 40 of these seen at 15 μm, and the rest only in the 7 μm bandpass. Almost half of the cluster galaxies detected at 15 μm prove to be LIRGs. (At the higher redshifts of the above sample only LIRGs fall above the sensitivity limit of the observations.) About 60% of these cluster galaxies were detected in observations originally intended to study distant field galaxies via the gravitational lensing amplification of the foreground clusters (Altieri et al., 1999; Barvainis et al., 1999; Metcalfe et al., 2003), and which, being generally very deep spatially- oversampled measurements, were able to provide insights about the lensing cluster galaxy populations (Biviano et al., 2004; Coia et al., 2005a,b). Figure 2, taken from Coia et al. (2005b) is a V-band image of the z = 0.39 galaxy cluster Cl0024+1564 overlaid with contours of an ISO 15 μm map. The capacity of ISOCAM to detect and assign MIR counterparts unambiguously to numerous galaxies in the field is evident. The first published ISOCAM observations of a distant cluster were those of the z = 0.193 cluster A1732 by Pierre et al. (1996), which they observed at 7 and 15 μmoveran8× 8 arcmin2 field. They found some evidence for a deficiency of spirals and star forming galaxies in the cluster, identifying only four cluster sources

Figure 2. From Coia et al. (2005b): Contours of a 15 μm map of the (gravitationally lensing) galaxy cluster Cl0024+1654 overlaid on a V-band FORS2 VLT image. 15 μm Sources numbered in order of increasing R.A. Dark circles identify spectroscopically conﬁrmed cluster galaxies. Greek letters denote four prominent gravitational lensing arcs. North is up and East to the left, with the centre of the map falling at R.A. 00 26 37.5 and DEC. 17 09 43.4 (J2000). 436 L. METCALFE ET AL.

(at 7 μm, no cluster sources were detected at 15 μm) and 10 MIR galaxies in total in the field, most of them judged to be foreground. Nevertheless, these were the faintest MIR extragalactic sources reported up to that point and underlined the need for ultra-deep observations to detect cluster members at 15 μm. Lémonon et al. (1998) reported evidence for an active star-forming region in a cooling flow (later ‘cool-core’) from 7 and 15 μm observations of the inner square arcmin of the well known lensing cluster A2390 (z = 0.23), with an attendant SFR −1 of as much as 80M yr in the central cD. But this was later found to be compatible with non-thermal emission from a jet associated with the cD (Edge et al., 1999). The estimated cluster mass deposition rates in cooling flows have since been lowered by one or two orders of magnitude (Böhringer et al., 2002). A1689 (z = 0.181) was the first distant cluster for which detailed ISO observations were reported. Fadda et al. (2000) detected numerous cluster members (30 at 7 μm and 16 at 15 μm) within 0.5 Mpc of the cluster centre, and they found a correlation between the B-15 μm colour and cluster-centric distance of the galaxies. The 15 μm galaxies are blue outliers with respect to the colour/magnitude relation for the cluster and become brighter going from the center to the outer parts of the cluster. Coupled with the systematic excess of the distribution of the B-15 μm colours with respect to nearby clusters (Virgo and Coma), this suggested the existence of an IR analogue of the Butcher–Oemler effect in A1689 (see Figure 3). A follow-up optical study of these infrared galaxies (Duc et al., 2002) showed that the morphology of the 15 μm sources in A1689 is generally spiral-like, with disturbances reminiscent of tidal interactions. No LIRGs were found in A1689. The highest 10 total IR luminosity found for a cluster galaxy was 6.2 × 10 L, corresponding

Figure 3. Fadda et al. (2000) found the [BT –7μm] colour distribution of A1689 cluster galaxies (filled circles) to be compatible with that of the nearby Virgo (crosses) and Coma (empty circles) clusters, but the [BT –15μm] colour distribution showed a systematic excess with respect to those nearer clusters. The cluster contains an excess of 15 μm sources relative to the field, suggesting that the environment of A1689 triggers starburst episodes in galaxies in the cluster outskirts that have similar IR luminosities and FIR/optical colours to those of field starburst galaxies. ISO’s CONTRIBUTION TO THE STUDY OF CLUSTERS OF GALAXIES 437

Figure 4. This plot, taken from Duc et al. (2002), compares the A1689 SFR derived from optical [OII] measurements, and from ISO MIR measurements. It illustrates the limitations of purely optical indicators of star formation rates. A major part, at least 90%, of the star formation activity taking place in Abell 1689 is hidden by dust at optical wavelengths.

to a SFR of ≈11M per year. The median SFR for A1689, derived from 15 μm measurements, was 2M per year, while the median found from [OII] (optical) measurements was only 0.2M per year. This paper revealed the importance of IR observations in the study of star- formation in clusters. About one-third of the 15 μm sources show no sign of star formation in their optical spectra. Moreover, comparing the star-formation estimates from IR and [OII] (see Figure 4), Duc et al. (2002) deduced that at least 90% of the star formation activity taking place in A1689 is obscured by dust. The ISO gravitational lensing survey programme (Metcalfe et al., 2003), and related work, led to deep observations of the core of several distant clusters. Large numbers of cluster galaxies were detected in the ﬁelds of A2218, A2390 and Cl0024+1654. So far, the cases of A2218 and Cl0024+1654 have been treated in dedicated papers. In the analysis of the galaxies in the ﬁeld of A2218, a rich cluster at z = 0.175, Biviano et al. (2004) found nine cluster members at 15 μm inside a radius of 0.4 Mpc. In contrast to the case of A1689, which is at almost identical redshift −1 (z = 0.181) and for which the median SFR is about 2M yr and median LIR is 10 around 10 L, only one of the A2218 MIR galaxies is a blue Butcher–Oemler galaxy. The MIR luminosity of A2218 galaxies is moderate and the inferred star 438 L. METCALFE ET AL.

TABLE II Summary of ISOCAM observations and results at 15 μm for ﬁve clustersa of galaxies.

Sensitivity (μJy) n scesc LIRGsf

2 b d e Cluster z Area ( )(5σ) min. CT Rvir Mvir Obs Exp Cl0024 0.39 37.8 140 141 13 35 0.94 6.42 10 – A370 0.37 40.5 350 208 1 20 0.91 5.53 1 8 A1689 0.18 36.0 450 320 11 18 1.1 5.7 0 1 A2390 0.23 7.0 100 54 4 28 1.62 20.35 0 1 A2218 0.18 20.5 125 90 6 46 1.63 18.27 0 1

aThe table is adapted from Coia et al. (2005b), and the data originates from Metcalfe et al. (2003) for Abell 370, Abell 2218 and Abell 2390, Fadda et al. (2000) and Duc et al. (2002) for Abell 1689, and Coia et al. (2005b) for Cl 0024+1654. The content of the columns in the table are as follows: name and redshift of the cluster, total area scanned, sensitivity reported at the 5σ level, ﬂux of the weakest reported source in μJy. Then number of cluster galaxies, total number of sources detected including sources without redshift and stars, virial radius of the cluster, virial 10 mass, number of sources with LIR > 9 × 10 L detected and expected. The expected number of sources was obtained by comparison with Cl0024+1654 as described in the text. Virial radii and masses are from Girardi and Mezzetti (2001) and King et al. (2002). bFaintest source considered in publication. cNumber of cluster sources, and total number of IR sources. dCluster virial radius (h−1 Mpc). − eCluster virial mass (h 1 1014 M). f 10 Number of LIRGs (or near LIRGs (LIR > 9×10 L)) detected vs. number expected if cluster were to be similar to Cl0024+1654.

−1 8 formation rate is typically less than 1M yr with a median LIR of only 6×10 L. The absence of a MIR BO effect in A2218 might be a consequence of the small area observed, about 20.5 square arcmin (r < 0.4 Mpc), and yet the area studied for A1689 was not much larger, at about 36 square arcmin. Coia et al. (2005b) suggest that the difference between these two clusters may be traced to their having different dynamical status. The same sort of comparison can be drawn between the clusters Cl0024+1654 (z = 0.39) and A370 (z = 0.37). These two clusters at similar redshift and mapped in almost identical ways, exhibit very different numbers of LIRGs (Table II), probably, according to Coia et al. (2005b), because an ongoing cluster merger gives rise to enhanced star-forming activity in Cl0024+1654. As can be appreciated from Figure 5 taken from Coia et al. (2005b), the median infrared luminosity of ISO-detected cluster galaxies in Cl0024+1654 is around 11 1 × 10 L. Star formation rates derived from the 15 μm data range from 8 to −1 −1 77M yr , with median(mean) value of 18(30)M yr . Because of the different sky areas mapped with ISO for several of the clusters discussed here, and the different sensitivities achieved, it is not straightforward ISO’s CONTRIBUTION TO THE STUDY OF CLUSTERS OF GALAXIES 439

Figure 5. The infrared luminosity distribution for MIR galaxies in Cl0024+1654. The median value 11 for LIR is ∼1.0 × 10 L. (From Coia et al., 2005b.) to compare the results for different clusters. A useful approach is to compare the number of LIRGs detected in each cluster, since these IR-bright galaxies would be expected to be seen for any of the observations considered. Such a comparison, taken from Coia et al. (2005b), is presented in Table II. For each cluster an “expected” number of LIRGs is derived to test the hypothesis that the cluster is similar to the LIRG-rich cluster Cl0024+1654.5 The LIRG count in Cl0024+1654 is multiplied by the ratios of (a) virial mass per unit area of the cluster to that of Cl0024+1654, the square of the respective distances to the cluster and to Cl0024+1654, and the observed solid-angle for the cluster to that of Cl0024+1654. The resulting column of the table can then be compared with the column listing the actual observed number of LIRGs for each cluster. The two most distant clusters observed with ISO are Cl0024+1654 (z = 0.39, Coia et al., 2005b) and J1888.16CL (z = 0.56, Duc et al., 2004). These are among the deepest ISOCAM observations and could detect several cluster members. (For Cl0024+1654, 13 out of 35 sources found at 15 μm are spectroscopically conﬁrmed to be cluster sources. For J1888.16CL, 6 out of 44 sources found at 15 μm are so conﬁrmed.) A common feature of these two clusters is the high star formation rate inferred from their MIR luminosities. These two observations were also the most extended cluster maps performed in terms of absolute cluster area covered at

5In fact, to avoid throwing away several sources close to the LIRG flux threshold of 1 × 1011 L and thereby degrading the statistics of the comparison, the flux threshold for the comparison of detected luminous sources was set to 9 × 1010 L. 440 L. METCALFE ET AL. the cluster, 2.3 × 2.3 and 1.3 × 6 square Mpc for Cl0024+1654 and J1888.16CL respectively. This fact, and evolutionary effects detectable around z ∼ 0.5 (see Dressler et al., 1999), may explain the high IR luminosity of the 15 μm sources found in these clusters. In the case of J1888.16CL, Duc et al. (2004) estimate star-formation rates ranging between 20 and 120M per year. At least six galaxies belong to the cluster and 11 have IR luminosities above 1.3 × 10 L. In Cl0024+1654, Coia et al. (2005b) 10 report 10 sources brighter than 9 × 10 L. The star formation rates inferred from the MIR flux are one to two orders of magnitudes greater than those based on the [OII] flux (though in this case the comparison was only possible for the three sources for which [OII] data was available.) This is compatible with the result in A1689 (Figure 4) and implies similar dust extinction characteristics. Interestingly, the galaxies emitting at 15 μm appear to have a spatial distribution and a velocity dispersion slightly different from the other cluster galaxies. Galax- ies in Cl0024+1654 are detected preferentially at larger radii, with the velocity dispersion of 15 μm sources being greater than that of the galaxies in the cluster. In J1888.16CL, Duc et al. (2004) estimate that to explain the number of sources detected on the basis of infall of galaxies from the field an infall rate of about 100 massive galaxies per 100 Myr is required, which seems unrealistic. Numerical simulations and X-ray observations show however that accretion onto clusters from the field is not a spherically symmetric process, but occurs along filaments or via mergers with other groups and clusters. One therefore cannot exclude the possibility that the LIRGs observed in these distant clusters belonged to such a recently accreted structure. An alternative possibility is that the collision with an accreted group of galaxies stimulated star formation in the galaxies of the group as a consequence of a rapidly varying tidal field (Bekki, 1999). This could be the case for Cl0024+1654 and A1689, clusters which show evidence of accreting groups of galaxies in their multi-modal velocity distributions. Cl0024+1654 is in the process of interacting with a smaller cluster.

2.4. A DIFFUSE INTRA-CLUSTER DUST COMPONENT?

The hot intra-cluster material contains metals, and so is not entirely primordial. Might not the stars which produced the metals also have deposited dust in the ICM? The first to note that emission from intra-cluster material might be observable were Yahil and Ostriker (1973), based on a galactic dust-to-gas ratio and the observed intra-cluster gas. Ostriker and Silk (1973) and Silk and Burke (1974) developed expressions for the lifetime of dust in a hot intra-cluster medium. Pustilnik (1975), drawing upon contemporaneous reports of optical absorption in clusters, attributed it to dust and estimated that cluster emission at 100 μm would be in the 103 to 104 Jy range for six nearby clusters. Voshchinnikov and Khersonskii (1984) also attributed claimed reddenning of galaxies in distant clusters to dust absorption, and estimated ISO’s CONTRIBUTION TO THE STUDY OF CLUSTERS OF GALAXIES 441 that the total FIR emission from the Coma or Perseus clusters should be 105– 106 Jy (tens of Jy/arcmin2) in the 50–100 μm range. They estimated the sputtering lifetime of intra-cluster dust grains to be up to 108 years. Hu et al. (1985), noting that intra-cluster dust must be short-lived, predicted FIR dust emission of a few Jy per square degree, close to the IRAS limit, for a sample of X-ray luminous clusters. IRAS measurements failed to bear out even the most modest of the above predictions. Kelly and Rieke (1990) co-added IRAS scans across 71 clusters with 0.3 ≤ z ≤ 0.92 to arrive at an average 60 μm value for cluster emission of 26 ± 5 mJy per cluster, and 46 ± 22 mJy at 100 μm. Dwek et al. (1990) refined models of intra-cluster dust and its interactions and calculated an upper limit of 0.2 MJy/sr for dust-emission from the Coma cluster, consistent with IRAS observations. Then total cluster emission would not be more than a few Jy at the peak wavelength (around 100 μm). They concluded that dust in the cluster centre could not explain the visual extinction, nor could cluster galaxies or their halos. Dust in the outskirts could, if it were un-depleted. But they saw no mechanism for the production of such dust. Wise et al. (1993) analysed 56 clusters at 60 and 100 μm from clusters with a range of X-ray emission, and some without cDs. For the only two clusters (A262 and A2670) showing a far-infrared excess lacking an immediate explanation (in terms of point sources or cirrus) they concluded that the result was likely to be due to discrete sources in the clusters. Averaged over the sample as a whole there was evidence of excess FIR at the 2-σ level. No large FIR excesses associated with cooling flows were found. Bregman et al. (1990) looked for evidence of star formation in 27 cD galaxies. In half of their sample of X-ray-bright clusters they found IR, X-ray and blue luminosities to be comparable, consistent with dust grains heated by the X-ray emitting gas, thereby suggesting that dust cooling can compare with thermal bremsstrahlung as a cooling mechanism for the intra-cluster gas. Cox et al. (1995) studied a much larger sample of 158 Abell clusters, again at 60 and 100 μm, and after making a more rigorous correction for spurious sources due to galactic cirrus, they concluded that only about 10% of cD galaxies in rich clusters have significant FIR emission, but with luminosities 10 times greater than the X-ray luminosities produced in the cores of clusters, a condition which they therefore regarded as transient for any individual cluster. If the FIR emission comes from dust heated by the intra-cluster thermal electrons, significant dust sputtering is expected on timescales of several 108 years. Dust must then be replenished to account for continuous IR emission, presumably through the mechanisms discussed for stripping material from cluster galaxies (see Section 1). By the launch of ISO one could imagine a “life-cycle” of dust in a cluster, tracing the flow of material – gas and dust – out of infalling galaxies, the destruction of dust in the high-temperature intra-cluster medium, and its possible eventual re- deposition, through the mechanism of cooling flows, into the cluster-dominant galaxies. But only upper-limits or occasionally, and with marginal significance, 442 L. METCALFE ET AL.

Figure 6. From Stickel et al. (2002): The overall zodiacal-light-subtracted surface brightness ratio I120 μm/I180 μm for A1656 (Coma) averaged over both scan position angles and all detector pixels. global or average cluster FIR emission, could constrain scenarios for the role of dust in the physics of the clusters as a whole. Stickel et al. (1998, 2002) used ISOPHOT to observe extended FIR emission of six Abell clusters. Strip scanning measurements were performed at 120 and 180 μm. The raw profiles of the I120 μm/I180 μm surface brightness ratio including zodiacal light show a bump towards Abell 1656 (Coma), dips towards Abell 262 and Abell 2670, and are without clear structure towards Abell 400, Abell 496, and Abell 4038. After subtraction of the zodiacal light and allowance for cirrus emission, only the bump towards Abell 1656 (Coma) is still present (Figure 6). This excess of ≈0.2 MJy/sr seen at 120 μm towards Abell 1656 (Coma) is interpreted as thermal emission from intra-cluster dust distributed in the hot X-ray emitting Coma intra-cluster medium. The integrated excess flux within the central region of 10 to 15 diameter is ∼2.8 Jy. Since the dust temperature is poorly constrained 7 only a rough estimate of the associated dust mass of MD ∼ 10 M can be derived. The associated visual extinction is negligible (AV 0.1 mag) and much smaller than claimed from optical observations. No evidence is found for intra-cluster dust in the other five clusters observed. Quillen et al. (1999) suggested integrated emission from the cluster galaxies as the most likely source for the detected signal at the centre of Coma. Stickel et al. (2002) replied that if this was indeed the case, the same signal should have been detected in all clusters observed. The absence of any signature for intra-cluster dust in five clusters and the rather low inferred dust mass in Abell 1656 indicate that intra-cluster dust is probably not responsible for the excess X-ray absorption reported in cooling flow clusters (White et al., 1991). These observations thus represent a further unsuccessful attempt to detect the presumed final stage of the cooling flow material. This agrees with numerous previous studies at other wavelengths, while spectroscopic observations with ESA’s XMM-Newton X-ray observatory have shown that intra-cluster gas ISO’s CONTRIBUTION TO THE STUDY OF CLUSTERS OF GALAXIES 443 does not cool beyond ∼1 keV (see, e.g., Molendi and Pizzolato, 2001), so dust deposition in cooling flows is not actually expected. Not surprisingly, then, further attempts by Hansen et al. (1999, 2000a,b) also failed to detect dust associated with the cooling flows presumed to exist at the centres of most galaxy clusters.

3. Conclusions

Building upon the legacy of IRAS, ISO could consolidate a number of important fundamental conclusions about the properties of galaxies in the Virgo, Coma and other nearby galaxy clusters. The correlation between galaxy Hubble type and mid- and far-infrared properties was firmly established. A major new cold dust component was identified by ISOPHOT observations, which could not have been found by IRAS. Moderate resolution FIR spectroscopy was possible for Virgo cluster galaxies using LWS, establishing a correlation between the strength of the [CII] line and the FIR flux and a two order of magnitude difference in [CII] to near-IR ratio between early type galaxies and late (spiral) types. Mid-infrared emission (5– 18 μm) was found to correlate with star formation, but to trace it less faithfully when star formation rates become high enough for UV photons to disrupt the infrared- emitting materials. However, the MIR emission traces well the FIR and bolometric emission. In general, the properties of galaxies in nearby clusters (z < 0.1) were found to exhibit little dependence on the cluster environment. ISO, for the first time, could extend mid-infrared observations to clusters beyond z = 0.1, and in so doing has detected over 100 galaxies in clusters in the redshift range 0.17 to almost 0.6. Although the collected observations on seven such clusters were rather heterogeneous, a number of important trends are found in the ISO results. There is a clear tendency for clusters at higher redshift to exhibit higher average rates of star formation in their galaxies, and numerous LIRGs have been observed in such clusters. There is tentative evidence for an association between the infrared luminosities found and galaxy infall from the field, but substantial evidence to link high levels of star-formation in cluster galaxies to the dynamical status of a cluster, and to interactions with other (sub-)clusters. It seems clear that star formation rates deduced for cluster galaxies from optical tracers often fall one to two orders of magnitude below rates derived from MIR emission levels, so that star formation in cluster galaxies may have been seriously underestimated in some cases, in the past. ISO found little evidence for widespread infrared emission from dust in the intra-cluster medium. Clearly, what ISO has not been able to supply has been systematic large area mapping of a substantial sample of galaxy clusters out to high redshift. Observing times to achieve such coverage with ISO would have been prohibitive, given the multi-position and heavily-overlapped rasters that would have been required. If galaxy evolution within clusters is to be further explored and related to galaxy 444 L. METCALFE ET AL. evolution in the field, clusters of different masses and dynamical status must be studied systematically in the MIR and FIR out to well over 1 virial radius in order to understand how and why the IR properties of galaxies change from cluster to cluster, and from cluster to field. Large area coverage is important since it is known that significant modifications of the galaxy properties already occur in the outskirts of galaxy clusters (Kodama et al., 2001; Lewis et al., 2002; Gómez et al., 2003). Several programmes underway or planned with the Spitzer Space Observatory should thoroughly address these challenges.

References

Altieri, B., Metcalfe, L., Kneib, J.-P., et al.: 1999, A&A 343, 65. Andersen, V., and Owen, F. N.: 1995, AJ 109, 1582. Barnes, J. E., and Hernquist, L. E.: 1991, ApJ 370, L65. Barvainis, R., Antonucci, R., and Helou, G.: 1999, AJ 118, 645. Bekki, K.: 2001, ApJ 546, 189. Bicay, M., and Giovanelli, R.: 1987, ApJ 321, 645. Biviano, A., and Katgert, P.: 2004, A&A 424, 779. Biviano, A., Metcalfe, L., McBreen, B., et al.: 2004, A&A 425, 33. B¨ohringer, H., Matsushita, K., Churazov, E., Ikebe, Y., and Chen, Y.: 2002, A&A 382, 804. Boselli, A., Lequeux, J., Contursi, A., et al.: 1997a, A&A 324, L13. Boselli, A., Tuffs, R., Gavazzi, G., Hippelein, H., and Pierini, D.: 1997b, A&A 121, 507. Boselli, A., Lequeux, J., Sauvage, M., et al.: 1998, A&A 335, 53. Boselli, A., Gavazzi, G., and Sanvito, G.: 2003a, A&A 402, 37. Boselli, A., Sauvage, M., Lequeux, J., Donati, A., and Gavazzi, G.: 2003b, A&A 406, 867. Boselli, A., Lequeux, J., and Gavazzi, G.: 2004, A&A 428, 409. Bregman, J. N., McNamara, B. R., and O’Connell, R. W.: 1990, ApJ 351, 406. Butcher, H., and Oemler, A.: 1978, ApJ 219, 18. Butcher, H., and Oemler, A.: 1984, ApJ 285, 426. Cayatte, V., van Gorkom, J. H., Balkowski, C., and Kotanyi, C.: 1990, AJ 100, 604. Cesarsky, C., Abergel, A., Agnese, P., et al.: 1996, A&A 315, L309. Clegg., P. E., Ade, P. A. R., Armand, C., et al.: 1996, A&A 315, L38. Coia, D., Metcalfe, L., McBreen, B., et al.: 2005a, A&A 430, 59. Coia, D., McBreen, B., Metcalfe, L., et al.: 2005b, A&A 431, 433. Contursi, A., Boselli, A., Gavazzi, G., Bertagna, E., Tuffs, R., and Lequeux, J.: 2001, AA 365, 11. Cox, C. V., Bregman, J. N., and Schombert, J. M.: 1995, ApJS 99, 405. Doyon, R., and Joseph, R. D.: 1989, MNRAS 239, 347. Dressler, A.: 1980, ApJ 236, 351. Dressler, A.: 2004, in Mulchaey, J. S., Dressler, A., and Oemler, A. (eds.), Clusters of Galaxies: Probes of Cosmological Structure and Galaxy Evolution, Carnegie Obs, Pasadena, CA. Dressler, A., and Gunn, J. E.: 1983, ApJ 270,7. Dressler, A., Lynden-Bell, D., Burstein, D., et al.: 1987, ApJ 313, 42. Dressler, A., Oemler, A., Couch, W., et al.: 1997, ApJ 490, 577. Dressler, A., Smail, I., Poggianti, B. M., et al.: 1999, ApJS 122, 51. Duc, P.-A., Poggianti, B. M., Fadda, D., et al.: 2002, A&A 382, 60. Duc, P.-A., Fadda, D., Poggianti, B., et al.: 2004, astro-ph/0404183. ISO’s CONTRIBUTION TO THE STUDY OF CLUSTERS OF GALAXIES 445

Dwek, E., Rephaeli, Y., and Mather, J. C.: 1990, ApJ 350, 104. Edge, A., Ivison, R., Smail, I., Blain, A., and Kneib, J.-P.: 1999, MNRAS 306, 599. Elbaz, D., Cesarsky, C. J., Chanial, P., et al., 1999, A&A 351, 37. Elbaz, D., Cesarsky, C. J., Chanial, P., et al., 2002, A&A 384, 848. Ellingson, E., Lin, H., Yee, H. K. C., and Carlberg, R. G.: 2001, ApJ 547, 609. Ellis, R., Smail, I., Dressler, A., et al.: 1997, ApJ 483, 582. Fadda, D., Elbaz, D., Duc, P.-A., et al.: 2000, A&A 361, 827. Fasano, G., Poggianti, B. M., Couch, W. J., et al.: 2000, ApJ 542, 673. Franceschini, A., Aussel, H., Cesarsky, C. J., Elbaz, D., and Fadda, D.: 2001, A&A 378,1. Girardi, M., and Mezzetti, M.: 2001, ApJ 548, 79. Gómez, P. L., Nichol, R., Miller, et al.: 2003, ApJ 584, 210. Gruppioni, C., Lari, C., Pozzi, F., et al.: 2002, MNRAS 335, 831. Gunn, J. E., and Gott, J. R.: 1972, ApJ 176,1. Hansen, L., Jørgensen, H., Nørgaard-Nielsen, H., Pedersen, K., Goudfrooij, P., and Linden-Vornle, M.: 1999, AA, 349, 406. Hansen, L., Jørgensen, H., Nørgaard-Nielsen, H., et al.: 2000a, A&A, 356, 83. Hansen, L., Jørgensen, H., Nørgaard-Nielsen, H., et al.: 2000b, A&A, 362, 133. Héraudeau, Ph., Oliver, S., del Burgo, C., et al.: 2004, MNRAS 354, 924. Hu, E. M., Cowie, L. L., and Wang, Z.: 1985, ApJS 59, 447. Kawara, K., Matsuhara, H., Okuda, H., et al.: 2004, A&A 413, 843. Kelly, D. M., and Rieke, G. H.: 1990, ApJ 361, 354. Kenney, J. D., and Young, J. S.: 1986, ApJ 301, 13. Kenney, J. D., and Young, J. S.: 1989, ApJ 344, 171. Kessler, M. F., Steinz, J. A., Anderegg, M. E., et al.: 1996, A&A, 315, 27. King, L. J., Clowe, D. I., and Schneider, P.: 2002, A&A 383, 118K. Kodama, T., Smail, I., Nakata, F., Okamura, S., and Bower, R.: 2001, ApJ 562,9. Lari, C., Pozzi, F., Gruppioni, C., et al.: 2001, MNRAS 325, 1173L. Larson, R. B., Tinsley, B. M., and Caldwell, C. N.: 1980, ApJ 237, 692. Lavery, R. J., and Henry, J. P.: 1988, ApJ 330, 596. Leech, K. J., Volk, H. J., Heinrichsen, I., et al.: 1999, MNRAS 310, 317. Leggett, S. K., Clowes, R. G., Kalafi, M., et al.: 1987, MNRAS 227, 563. Lemke, D., Klaas, U., Abolins, J., et al.: 1996, A&A 315, L64. Lémonon, L., Pierre, M., Cesarsky, C. J., et al.: 1998, A&A 334, L21. Lewis, I. J., Balogh, M., De Propris, R., et al.: 2002, MNRAS 334, 673. Margoniner, V., de Carvalho, R., Gal, R., and Djorgovski, S.: 2001, ApJ 548, L143. Martin, C. L.: 1999, ApJ 513, 156. Metcalfe, L., Kneib, J.-P., McBreen, B., et al.: 2003, A&A 407, 791. Molendi, S., and Pizzolato, F.: 2001, ApJ 560, 194. Moore, B., Katz, N., Lake, G., Dressler, A., and Oemler, A.: 1996, Nature 379, 613. Niklas, S., Klein, U., and Wielebinski, R.: 1995, A&A 293, 56. Odenwald, S. F.: 1986, ApJ 310, 86. Ostriker, J., and Silk, J.: 1973, ApJ 184, 113. Pierre, M., Aussel, H., Altieri, B., et al.: 1996, A&A 315, L297. Poggianti, B. M., Smail, I., Dressler, A., et al.: 1999, ApJ 518, 576. Poggianti, B. M., Bridges, T. J., Mobasher, B., et al.: 2001, ApJ 562, 689. Popescu, C. C., and Tuffs, R. J.: 2002, MNRAS 335, 41. Popescu, C. C., Tuffs, R. J., Volk, H. J., Pierini, D., and Madore, B. F.: 2002, ApJ 567, 221. Pustilnik, S. A.: 1975, SvAL 1, 49. Quillen, A. C., Rieke, G. H., Rieke, M. J., Caldwell, N., and Engelbracht, C. W.: 1999, ApJ 518, 632. Rodighiero, G., Lari, C., Franceschini, A., Gregnanin, A., and Fadda, D.: 2003, MNRAS 343, 1155. 446 L. METCALFE ET AL.

Rosati, P.: 2004, in Mulchaey, J. S., Dressler, A., and Oemler, A. (eds.), Clusters of Galax- ies: Probes of Cosmological Structure and Galaxy Evolution, Carnegie Obs., Pasadena, CA, http://www.ociw.edu/ociw/symposia/series/. Rowan-Robinson, M., Lari, C., Perez-Fournon, I., et al.: 2004, MNRAS 351, 1290. Sato, Y., Kawara, K., Cowie, L.L., et al.: 2003, A&A 405, 833. Scoville, N. Z., Becklin, E. E., Young, J. S., and Capps, R. W.: 1983, ApJ 271, 512. Serjeant, S., Oliver, S., Rowan-Robinson, M., et al.: 2000, MNRAS 316, 768. Serjeant, S., Carramiana, A., Gonzles-Solares, E., et al.: 2004, MNRAS.tmp, 475. Silk, J., and Burke, J. R.: 1974, ApJ 190, 11. Silva, L., Granato, G. L., Bressan, A., and Danese, L.: 1998, ApJ 509, 103. Soucail, G., Kneib, J.-P., Bzecourt, J., et al.: 1999, A&A 343, L70. Stickel, M., Lemke, D., Mattila, K., Haikala, L. K., and Haas, M.: 1998, A&A 329, 55. Stickel, M., Klaas, U., Lemke, D., and Mattila, K.: 2002, A&A 383, 367. Tuffs, R. J., Popescu, C. C., Pierini, D., et al.: 2002, ApJS 139, 37. Tully, R. B., and Shaya, E. J.: 1984, ApJ 281, 31. van Dokkum, P. G., and Stanford, S. A.: 2003, ApJ 585, 78. van Dokkum, P. G., Franx, M., Fabricant, D., Illingworth, G. D., and Kelson, D. D.: 2000, ApJ 541, 95. van Dokkum, P. G., Stanford, S. A., Holden, B. P., et al.: 2001, ApJ 552, L101. Visvanathan, N., and Sandage, A.: 1977, ApJ 216, 214. Voshchinnikov, N. V., and Khersonskii, V. K.: 1984, AdSpR 3, 443. Voshchinnikov, N. V., and Khersonskii, V. K.: 1984, ApSS 103, 301. Wang, G., Leggett, S. K., Clowes, R. G., MacGillivray, H. T., and Savage, A.: 1991, MNRAS 248, 112. White, D. A., Fabian, A. C., Johnstone, R. M., Mushotzky, R. F., and Arnaud, K. A.: 1991, MNRAS 252, 72. Wise, M., O’Connell, R., Bregman, J., and Roberts, M.: 1993, ApJ 405, 94. Yahil, A., and Ostriker, J.: 1973, ApJ 185, 787. Young, E., Low, F. J., Soifer, B. T., et al.: 1984, ApJ 278, L75.