Photoluminescence by Interstellar Dust

Home , NGC 2023, Protoplanetary nebula, Red Rectangle nebula

A Dissertation entitled Photoluminescence by Interstellar Dust

by Uma Parvathy Vijh

As partial fulﬁllment of the requirements for the Doctor of Philosophy Degree in Physics

Advisor: Prof. Adolf N. Witt

Graduate School

The University of Toledo August 2005

An Abstract of

Photoluminescence by Interstellar Dust

Uma Parvathy Vijh

Submitted in partial fulﬁllment of the requirements for the Doctor of Philosophy Degree in Physics

The University of Toledo August 2005

In this dissertation, we report on our study of interstellar dust through the process of photoluminescence (PL). We present the discovery of a new band of dust PL, blue luminescence (BL) with λpeak ∼ 370 nm in the proto-planetary nebula known as the

Red Rectangle (RR). We attribute this to fluorescence by small, 3-4-ringed polycyclic aromatic hydrocarbon (PAH) molecules. Further analysis reveals additional independent evidence for the presence of small PAHs in this nebula. Detection of BL using long-slit spectroscopic observations in other ordinary reflection nebulae suggests that the BL carrier is an ubiquitous component of the ISM and is not restricted to the particular environment of the RR. We present the spatial distribution of the BL in these nebulae and find that the BL is spatially correlated with IR emission structures attributed to aromatic emission features (AEFs), attributed to PAHs.

The carrier of the dust-associated photoluminescence process causing the extended red emission (ERE), known now for over twenty ﬁve years, remains unidentiﬁed. We constrain the character of the ERE carrier by determining the wavelengths of the radiation that initiates the ERE – λ < 118 nm. We note that under interstellar

ii conditions most PAH molecules are ionized to the di-cation stage by photons with

E > 10.5 eV and that the electronic energy level structure of PAH di-cations is consistent with ﬂuorescence in the wavelength band of the ERE.

In the last few chapters of the dissertation we present ﬁrst results from ongoing work: i) Using narrow-band imaging, we present the optical detection of the circum- binary disk of the RR in the light of the BL, and show that the morphology of the BL and ERE emissions in the RR nebula are almost mutually exclusive. It is very suggestive to attribute them to diﬀerent ionization stages of the same family of carriers such as PAH molecules. ii) We also present a pure spectrum of the BL free of scattered light, resolved into seven molecular emission bands, superimposed upon a broad continuum. The relative intensity of the component bands varies with position within the nebula, suggesting an origin in a set of several related molecular species, most likely small PAHs.

iii Acknowledgments First and foremost I express my heartfelt gratitude for my advisor Prof. Adolf Witt for his guidance and support. His knowledge, erudition and enthusiasm for astrophysical problems has encouraged me and will continue to motivate me. I must also thank my parents for their support, their belief and pride in me. I hope I live up to their pride. I also thank my husband and friend Aarohi who in spite of having to go through graduate school and dissertation-writing/defense at almost the same time as myself has never failed to support and encourage me. I hope I have been as good a friend to him as he has to me. We would like to acknowledge Karl Gordon for his efforts towards the observations at Steward Observatory and Karl was also the PI of the HST proposal, observations from which lead to the determination of the ERE excitation wavelength. We also acknowledge Paul Sell who made the measurements for the above study as a NSF-REU student at the University of Toledo. We acknowledge Donald York and the observing team at Apache Point Observatory for the spectroscopic mapping and narrow-band imaging effort, preliminary results from which are reported is this dissertation. We would also like to thank Dr. Louis Allamandola for his valuable comments on our first discovery manuscript. We also thank David Malin and Hans van Winckel for providing us with blue images of the Red Rectangle, Laurent Verstraete and W. W. Jochims for supplying us with data on ionization cross sections of PAH molecules, and Lewis Hobbs, Theodore Snow and Donald York for constructive discussions about the Red Rectangle nebula and its central source. I would like to thank my committee members for their support and helpful comments. This research was made possible through a generous allocation of observing time at CTIO, KPNO, Steward Observatory and through grants from the US Na- tional Science Foundation. Financial support for this study was provided through NSF Grant AST0307307 to The University of Toledo. I would also like to acknowledge a CTIO thesis student travel grant, for travel to Chile. This research has made use of NASA’s Astrophysics Data System (ADS) Biblio- graphic Services and the SIMBAD database, operated at CDS, Strasbourg, France, and also of the Aladin image server. Some of the data presented in this paper was obtained from the Multimission Archive at the Space Telescope Science Institute (MAST). STScI is operated by the Association of Universities for Research in As- tronomy, Inc., under NASA contract NAS5-26555. Support for MAST for non-HST data is provided by the NASA Office of Space Science via grant NAG5-7584 and by other grants and contracts.

iv Contents

Abstract ii

Acknowledgments iv

Contents v

List of Figures x

List of Tables xiv

1 Introduction 1

1.1 Dust Grains and Nanoparticles ...... 2

1.2 Luminescence by Dust Grains: Extended Red Emission ...... 4

1.2.1 Observational Techniques for ERE Detection ...... 9

1.2.2 Observational Constraints ...... 13

1.2.3 Models for the ERE Carrier ...... 14

1.3 PAHs in the ISM ...... 22

1.3.1 Unidentiﬁed Infrared Bands or the Aromatic Emission Features 23

1.3.2 Optical Fluorescence ...... 25

v 1.3.3 Formation of PAHs in the ISM ...... 29

1.4 Outline ...... 30

2 Discovery of Blue Luminescence 31

2.1 Introduction ...... 32

2.2 Target and Observations ...... 34

2.3 Results ...... 38

2.3.1 Line-depth Technique ...... 38

2.3.2 Identiﬁcation ...... 41

2.4 Discussion ...... 44

3 Small PAHs in the Red Rectangle 45

3.1 Introduction ...... 46

3.2 Observations ...... 48

3.3 Analysis and Results ...... 50

3.3.1 Blue Luminescence in the Red Rectangle ...... 50

3.3.2 Attenuation of HD 44179 ...... 61

3.4 Discussion ...... 75

3.4.1 Identiﬁcation of the BL Carrier ...... 75

3.4.2 Spectral Variability of the BL ...... 77

3.4.3 Spatial Variation of the BL ...... 77

3.4.4 Attenuation of HD 44179 ...... 79

3.5 Conclusions ...... 81

vi 4 Detection of BL in Other Nebulae 84

4.1 Introduction ...... 85

4.2 Observations ...... 88

4.3 Results ...... 89

4.3.1 Ced 201 ...... 89

4.3.2 Ced 112 ...... 93

4.3.3 NGC 5367 ...... 100

4.3.4 NGC 2023 ...... 105

4.4 Discussion ...... 109

4.4.1 The BL Spectrum ...... 109

4.4.2 Spatial Distribution of BL ...... 113

4.4.3 BL in Hydrogen Ionization Regions ...... 118

4.4.4 Survival vs. In-situ formation of BL Carriers ...... 119

4.5 Conclusions ...... 122

5 The Excitation of Extended Red Emission 123

5.1 Introduction ...... 125

5.2 Observations and Data Reduction ...... 131

5.3 Results ...... 136

5.3.1 The ERE, H2, and z-Band Morphology of NW Filaments ... 136

5.3.2 Determination of RV in NW Filaments ...... 142

5.3.3 Width of the ERE Filaments ...... 145

5.3.4 Wavelength of ERE Initiation ...... 146

vii 5.4 Discussion ...... 149

5.4.1 Impact on Existing ERE Models ...... 150

5.4.2 Possible New ERE Carriers ...... 153

5.4.3 Consistency Check: The Red Rectangle ...... 158

5.4.4 Consistency Check: The High-|b| Galactic Cirrus ...... 162

5.5 Conclusions ...... 164

6 Optical Emission Band Morphologies of the Red Rectangle 167

6.1 Introduction ...... 167

6.2 Observations and Reductions ...... 168

6.3 Results and Discussion ...... 170

6.3.1 Blue Luminescence in the Red Rectangle ...... 170

6.3.2 Sharp Emission Feature at 5800 A˚ ...... 174

6.3.3 ERE ...... 177

6.4 Conclusions ...... 181

7 Spectral Characteristics of BL in the Red Rectangle 183

7.1 Introduction ...... 184

7.2 Observations and Reductions ...... 186

7.3 Results ...... 189

7.3.1 The Band Structure of the BL ...... 189

7.3.2 Variation of BL Spectra with Position in the RR ...... 191

7.3.3 Variation of BL Intensities with Position in the RR ...... 193

7.3.4 Morphology of BL in the RR ...... 194

viii 7.4 Discussion ...... 197

7.4.1 Identiﬁcation of the BL Carrier ...... 197

7.4.2 Temperature of the BL Carrier ...... 199

7.5 Conclusions ...... 200

8 Summary 201

8.1 What did we learn? ...... 201

8.2 Future Work ...... 206

A Determination of BL Intensity at the Balmer Jump 207

References 225

ix List of Figures

1-1 Interstellar dust emission ...... 3

1-2 Schematic energy level diagram of a photoluminescing system .... 5

1-3 Line-depth technique ...... 11

1-4 Structures of some small PAHs ...... 23

1-5 UIR/AEF spectrum ...... 25

1-6 Jablonski diagram ...... 27

1-7 Solution dependence of PAH ﬂuorescence ...... 28

1-8 Temperature dependence of PAH ﬂuorescence ...... 28

1-9 Formation routes for PAHs in the ISM ...... 29

2-1 Variation of ﬂuorescence peak with molecular size ...... 34

2-2 The RR nebula at diﬀerent optical wavelengths ...... 37

2-3 Filling in of a Balmer line due to ﬂuorescence ...... 40

2-4 Comparison of BL spectrum with solid-state and PAH candidates .. 43

3-1 Blue Image of the RR with overlaid apertures ...... 51

3-2 Red HST image of the RR with overlaid apertures ...... 52

x 3-3 Spectral variation of the BL ...... 54

3-4 Ratio of Band-IBL to Band-Isc and normalized distributions ...... 56

3-5 Spatial distribution of the BL and correlations with other emissions . 60

3-6 Observed UV/optical SED of HD 44179 compared to model SED .. 65

3-7 UV/optical attenuation curve for HD 44179 ...... 67

3-8 Attenuation curve for HD 44179, a diﬀerent representation ...... 69

3-9 Ionization potential of PAHs as a function of molecular mass ..... 72

4-1 Image of Ced 201 with overlaid slits ...... 90

4-2 BL spectrum in Ced 201 ...... 92

4-3 Image of Ced 112 with overlaid slits ...... 94

4-4 BL spectrum in Ced 112 ...... 96

4-5 ISOCAM image of Ced 112 with overlaid apertures ...... 98

4-6 Band-IBL and Band-Isc distribution in Ced 112 ...... 99

4-7 Ratio of Band-IBL to Band-Isc in Ced 112 ...... 100

4-8 Image of NGC 5367 with overlaid slit ...... 102

4-9 Bl spectra in NGC 5367 ...... 104

4-10 Image of NGC 2023 with overlaid slit ...... 106

4-11 BL spectra in NGC 2023 ...... 108

4-12 Average BL spectra in Ced 112 and Ced 201 compared to BL in the RR111

4-13 Banded BL spectrum in Ced 201 ...... 112

4-14 Band-IBL and Band-Isc in Ced 201 and Ced 112 ...... 115

4-15 Ratio of Band-IBL to Band-Isc in Ced 201 and Ced 112 ...... 116

xi 4-16 Average BL spectrum in NGC 2023 and NGC 5367 ...... 119

5-1 F606W WFPC2 image of NGC 7023 ...... 131

5-2 The ACS ﬁlter and the continuum-subtracted ERE and Hα images . 133

5-3 NICMOS images and continuum-subtracted ERE and H2 ...... 135

5-4 Four cuts through the ERE and H2 images ...... 138

5-5 The density-dependence of dust optical depth of a PDR ...... 141

5-6 CCM extinction relations ...... 143

5-7 Histogram of ratios of ﬁlament widths ...... 144

5-8 Allowed ratio of extinctions ...... 146

5-9 The limits of ERE excitation ...... 148

5-10 First and second ionization potentials of PAHs ...... 155

5-11 Model spectra of representative PAH di-cations...... 157

6-1 Image at 3934 A˚ with contours tracing the BL ...... 173

6-2 Image at 4050 A˚ with scattered light contours ...... 174

6-3 5800 A˚ image divided by 5700 A˚ image...... 176

6-4 5855 A˚ image divided by 5700 A˚ image...... 177

6-5 6400 A˚ image divided by 5700 A˚ image...... 180

6-6 6400 A˚ image divided by 4050 A˚ image...... 181

7-1 HST red image superimposed with slits and extraction windows ... 188

7-2 Nebular spectra at diﬀerent oﬀset from the central star ...... 190

7-3 Normalized BL spectra from diﬀerent nebular positions ...... 192

xii 7-4 Band-IBL and Isc distribution ...... 194

7-5 Deep blue image of the RR overlaid with extraction windows ..... 196

7-6 Fluorescence spectrum of anthracene ...... 199

xiii List of Tables

2.1 Fluorescence Intensity at two positions in the RR nebula...... 41

4.1 Details of objects observed, their illuminating stars and slits used. .. 89

6.1 Narrow-band observations of the RR nebula...... 170

7.1 RR Observations at APO ...... 187

7.2 Molecular Bands in the BL spectrum of the Red Rectangle ...... 191

xiv Chapter 1

Introduction

Our universe is a dusty universe, and dust grains constitute the dominant form of solid matter in the universe. Dust is seen in a wide variety of astrophysical environments, ranging from circumstellar envelopes around cool red giants to supernova ejecta, from diffuse and dense interstellar clouds and star-forming regions to debris disks around main-sequence stars, from comets to interplanetary space to distant galaxies and quasars. Longward of the Lyman limit of hydrogen, grains represent the main source of continuum opacity within galaxies, resulting in absorption and scattering of starlight at UV/visible wavelengths and thermal emission in the IR and sub-mm range. In the high-z universe, dust grains obscure a large fraction of the ongoing star formation, while actually promoting star formation by inducing necessary conditions of low temperature and high density in molecular clouds. They play a central role as a photoelectric heating agent for the interstellar gas. Within the interstellar medium (ISM), dust grains constitute a reservoir for a significant fraction of all atoms more massive than helium and play the role of an efficient catalyst for

1 2

the formation of H2.

Interstellar dust grains span a wide range in sizes: from a few angstroms to a

few micrometers. They play a vital role in the evolution of galaxies as an absorber,

scatterer, and emitter of electromagnetic radiation. Dust grains act as drivers for

the mass loss of evolved stars and are essential in the star and planet formation

process. The most prominent visual eﬀect of interstellar dust is as an agent shaping

the spectacular appearance of dusty systems such as protostars, young stellar objects,

evolved stars and galaxies.

1.1 Dust Grains and Nanoparticles

Despite the importance of interstellar grains in many astrophysical processes, their

nature, i.e. size distribution, their composition, their structure , as well as their total

mass, is not well known. Interstellar nanoparticles form the smallest and the most

numerous of these interstellar grains. The existence of this component has only been

recognised within the past two decades. Nanoparticles, after absorbing star light,

emit this absorbed energy not only as thermal radiation but also through speciﬁc

vibrational bands and electronic photoluminescence transitions. This fact suggests

the grains in the nanometer size range still retain their pure chemical and structural

characteristics.

The collective UV/optical absorption cross-section of nanoparticles is comparable to that of classical interstellar grains in the ISM. See Figure 1-1 .They are most likely

the structural building blocks of larger, composite grains and as such hold the key to 3 understanding the nature of interstellar grains in general. The study of interstellar nanoparticles is particularly promising, because they exhibit a wealth of observable diagnostics and in addition, their large surface/volume ratio makes them the dominant reservoir of solid surfaces in the ISM, enabling the formation of molecular hydrogen and other simple molecules as well as complex organic molecules which may lead to the origins of life.

Figure 1-1 Diﬀerent components of emission from interstellar dust (from Desert et al., 1990, ﬁg.4) 4

1.2 Luminescence by Dust Grains: Extended Red

Emission1

Our knowledge of the existence of interstellar dust, its spatial distribution, its

chemical composition and size distribution, and its mass relative to the rest of the

interstellar medium is almost exclusively based upon observations of the interactions

of interstellar dust grains with radiation. The processes studied in the past included

mainly interstellar extinction and polarization of starlight at optical, near-IR and ul-

traviolet wavelengths, followed by studies of dust emission over a wide range of wave-

lengths from the near-IR to the sub-mm and microwave region. In this thesis we are

reporting on the study of interstellar dust through the process of photoluminescence

(PL), a process in which absorptions of photons at ultraviolet/optical wavelengths are

followed by electronic transitions associated with the emission of longer-wavelength

optical and near-IR photons. Interstellar dust in nebulae and in the diﬀuse interstellar

medium of galaxies contains a component which responds to illumination by ultra-

violet photons with eﬃcient luminescence in the 500 nm to 1000 nm spectral range,

known as Extended Red Emission (ERE) (see Witt & Vijh, 2004, for a detailed re-

view). In the next section of this chapter we review the techniques of detection of

the ERE and the constraints that any proposal for the ERE carrier must confront.

Many models have been advanced over the past two decades to explain the ERE phe-

1Based in part on “Extended Red Emission: Photoluminescence by Interstellar Nanoparticles”

a review that appeared in Astrophysics of Dust 2004, ASP Conf. Proc. 309, Authors: A. N. Witt

& U. P. Vijh 5

nomenon, but despite promising progress on several fronts, no completely satisfactory

model for the ERE carrier/process exists at this time. 2

Series of rotational/vibrational transitions

ν h exc

Figure 1-2 Schematic energy level diagram of a photoluminescing system

Figure 1-2 schematically illustrates the physics of a PL process. The ground state

(1) represents the electronic ground state of a large molecule or molecular ion, or a state near the top of the valence band in a semiconductor particle. The excitation of the PL process under astrophysical conditions results from the absorption of a single

UV/optical photon, leading to an electronic transition from state (1) to state (2).

State (2) typically is a bound, high-lying vibrational/rotational level of the first or 6 second electronically excited state of a molecule or molecular ion, or a high state in the conduction band of a semiconductor particle. The excited system relaxes through a series of vibrational/rotational transitions until the electron finds itself in state (3), from where an optical electronic transition back to the ground state (1) is possible. In a polycyclic aromatic hydrocarbon (PAH) molecule, for example, state (3) can either be the lowest state in the singlet or triplet vibrational/rotational manifold of the first excited electronic level. The resulting emissions are referred to as fluorescence and phosphorescence, respectively. In a semiconductor particle, state (3) is a state near the lower edge of the conduction band, and the downward transition across the band- gap is simply referred to as photoluminescence. For a more complete review of PL in aromatic molecules and in semiconductor nanoparticles, the reader should consult the works of Birks (1970) and Yoffe (2001).

Different photoluminescing systems can be characterized by the energy difference between the exciting and PL photons, known as the Stokes shift, and by the quantum yield or photon conversion efficiency. Here we define quantum yield as the ratio of the number of PL photons to the number of exciting photons needed to produce the photoluminescence. Most luminescing systems exhibit quantum yields of substantially less than 100%, but if the energy of the exciting photon is more than twice the energy of the PL photon, highly isolated molecules or molecular ions are expected to yield two (or more) PL photons per excitation, in a process known as Poincaré

fluorescence (Leger et al., 1988). Thus, quantum yields well in excess of 100% may result in systems with large Stokes shifts under astrophysical conditions, although no specific identifications invoking this process have been made. In most natural 7

systems, ionization or the creation of defects either quenches or shifts the PL into

another spectral region, e.g. the near-IR, which permits relaxation to occur without

the emission of an optical photon, thus reducing the quantum yield to well below

100%.

Early suggestions that the high surface brightness of reﬂection nebulae in the

visible might be due in part to ﬂuorescence by the nebular dust grains were advanced

by Struve & Swings (1948) and by Aller (1956). Indeed, many minerals now thought

to be part of the composition of interstellar grains do ﬂuoresce under illumination by

ultraviolet light or following particle bombardment (Koike et al., 2002). However, a

first sensitive search for the existence of such fluorescence in several bright reflection

nebulae at wavelengths shortward of 490 nm by Rush & Witt (1975) failed to reveal

any detectable sign of dust ﬂuorescence, leading to the conclusion that the nebular

surface brightness at these wavelengths was a result of scattering with a relatively

high grain albedo.

The advent of new, sensitive detectors in the red region of the visible spectrum in the 1970’s brought the ﬁrst detections of PL, although they were not immediately recognized as such. The detection of the broad luminescence band in the spectrum of the Red Rectangle nebula (Cohen, 1975) was followed by its analysis in terms of unusual dust scattering properties by Greenstein & Oke (1977), although the possi-

bility of ﬂuorescence was mentioned. Subsequently, with much better observational

data, Schmidt et al. (1980) suggested a molecular-emission origin for the broad band

of excess radiation in the Red Rectangle. The belief that this object was unique in

exhibiting this excess radiation and the lack of a context of other sources showing 8

the same phenomenon clearly contributed to our inability to recognize the true na-

ture of the luminescence process in the Red Rectangle. For example, the spectrum

of red excess radiation detected in the Galactic dark nebula Lynds 1780 by Mattila

(1979) appeared suﬃciently diﬀerent that its connection to the Red Rectangle was

not recognized until several years later (Chlewicki & Laureijs, 1987).

A breakthrough occurred in this ﬁeld with the realization that dust luminescence

in the red part of the spectrum, extended red emission or ERE, as it soon became

known, was a common feature in many dusty environments that are illuminated

by ultraviolet photons. Detections of the ERE in reﬂection nebulae (Witt et al.,

1984a; Witt & Schild, 1986, 1988) soon led to the observation of ERE in high-latitude cirrus clouds (Guhathakurta & Tyson, 1989) and to the discovery of the presence of

ERE in the continuum spectrum of planetary nebulae (Furton & Witt, 1990, 1992).

Perrin & Sivan (1992) started a series of discoveries of the ERE band in HII regions with the Orion nebula, followed by the 30 Doradus nebula in the LMC (Darbon et al.,

1998) and the HII region Sh 152 (Darbon et al., 2000). The ERE detections in other

external galaxies include the halo of M82 (Perrin et al., 1995) and the prominent

dust lane in NGC 4826, the Evil Eye galaxy (Pierini et al., 2002). The detection

and measurement of ERE in the diﬀuse interstellar medium of the Milky Way Galaxy

over a wide range of Galactic latitudes by Gordon et al. (1998) represents a particular

milestone for three reasons. It demonstrated that the ERE carrier is a component

of interstellar dust on a Galaxy-wide scale; it represented detections of the ERE

intensities at a level three orders of magnitude fainter than observed in the Red

Rectangle; and ﬁnally, it permitted the ﬁrst reliable estimate of the quantum yield or 9

photon conversion eﬃciency of the ERE process.

1.2.1 Observational Techniques for ERE Detection

Several factors contribute to the challenges one faces when attempting to observe

ERE. First, ERE appears only in spatially extended objects, usually of relatively low surface brightness. Second, ERE is a very broad emission, spanning the wavelength range from 540 nm to at least 950 nm. And third, it always occurs in conjunction with other sources of diﬀuse emission or scattered light, most of which frequently are much brighter than the ERE. For example, in reﬂection nebulae the competing radiation is dust-scattered light; in HII regions dust-scattered light, atomic recombination continua as well as line emissions occur in the same spectral region, while in external galaxies integrated starlight is a contender as well. Attempts to measure

ERE as part of the high-Galactic-latitude background require careful subtractions of atmospheric foregrounds such as airglow and atmospheric scattered light, zodiacal light, integrated starlight and diﬀuse Galactic light, before the ERE can be isolated.

As a result, diﬀerent observational techniques must be applied to diﬀerent types of sources.

Color-Difference Technique The color-difference technique is an efficient way to

probe the presence of ERE in extended objects, provided that broadband colors

of the source are readily predictable in the absence of ERE, and the presence

of suﬃcient ERE changes these colors. This is the case when the competing

radiation is dust-scattered starlight, e.g. in reﬂection nebulae (RN) or in the the 10

diﬀuse Galactic background radiation. In an optically thin RN, the well-known

wavelength dependence of the scattering optical depth leaves the nebular radi-

ation bluer than that of the illuminating star throughout the optical spectrum.

A simple approach for estimating the color diﬀerence between scattered and

illuminating radiation is given by Witt (1985). If, therefore, the V-R color of a

nebular region is redder than that of the star while the B-V color is bluer, this

is a strong indication for the presence of ERE.

Long-Slit Spectroscopy The spectroscopy of ERE in low-surface-brightness, ex-

tended sources is strongly aﬀected by terrestrial airglow. Long-slit spectro-

graphs provide the best approach to simultaneously record and subsequently

subtract the foreground sky spectrum. Also, in sources whose spectrum is

dominated by emission lines, e.g. HII regions and planetary nebulae, long-slit

spectroscopy oﬀers a possibility for studying the faint continuum, including the

broad ERE band, between strong emission lines. The large extent of the ERE

in wavelengths requires a combination of relatively low wavelength resolution

and broad spectral coverage. Most of the work presented in this dissertation is

based on this technique.

Line-Depth Technique The presence of quasi-continuous luminescence may be dif-

ﬁcult to discern, if it lacks a distinctive spectral feature, as is the case with the

ERE band. In reﬂection nebulae, however, one can easily distinguish dust lumi-

nescence from scattered radiation by virtue of the fact that scattering faithfully

reproduces relative line depths or equivalent widths of spectral lines present in 11

the spectrum of the illuminating source. The presence of a certain fraction of lu-

minescence at the positions of such lines is revealed by a proportional reduction

in the relative line depths in the nebular spectrum compared to corresponding

relative line depths in the illuminating star, as illustrated in Figure 1-3.

This method is particularly useful in the search for luminescence in the blue

part of the spectrum, where strong hydrogen Balmer lines dominate the spectra

of most reﬂection nebulae. It is important to note that the spectral resolution

of both the stellar and the nebular spectrum must be identical to avoid a false-

positive result. This method was ﬁrst employed by Rush & Witt (1975), who

used photographic spectroscopy and narrow-band photoelectric photometry to

place useful upper limits on the contribution made by possible dust photolumi-

nescence to the observed surface brightness of several reﬂection nebulae.

Figure 1-3 Illustration of the line-depth technique. In the nebular spectrum, consisting of scattered light and PL, the relative line depth (CN − LN )/CN is smaller than the corresponding value (CS − LS)/CS in the stellar spectrum.

Unsharp Masking There exists a high degree of spatial variability of the ratio of 12

ERE to scattered light within a single nebula. In particular the ERE distri-

bution seems to be in bright ﬁlamentary structures. These bright structures

are a result of a favorable viewing geometry and a high local opacity for the

stellar radiation responsible for ERE excitation, thus leading to exceptionally

large column densities of ERE emitters, conﬁned to a space with narrow lateral

dimensions. The resulting ERE structures can be imaged with the unsharp

masking technique (Malin & Zealey, 1979), which permits the simultaneous

display of sharp features ranging in intensity over four orders of magnitude.

Witt & Malin (1989) applied this technique with good success to NGC 2023.

The location, width and orientation of numerous ERE ﬁlaments was revealed,

enabling subsequent studies of spatial correlations with ﬁlamentary emissions

by ﬂuorescent molecular hydrogen (Field et al., 1998) and by CN (Fuente et al.,

1995).

Imaging Spectropolarimetry Scattered light in reﬂection nebulae is substantially

linearly polarized at optical and near-IR wavelengths, typically at the 20% to

30% level. Watkin et al. (1991), through the use of imaging polarimetry, discov-

ered that localized ERE ﬁlaments in NGC 7023 are associated with a signiﬁcant

reduction of linear polarization in the R- and I-bands, consistent with a dilution

of polarization through the presence of unpolarized emission with an intensity

equal to that of the ERE, as determined by independent spectroscopic and

photometric techniques. This demonstration that ERE is unpolarized makes

attempts to explain the ERE in terms of scattering with unusual dust charac- 13

teristics (e.g. Greenstein & Oke, 1977) unsustainable. Rather, the ERE must

be the result of an unpolarized emission process.

1.2.2 Observational Constraints

The identiﬁcation of the ERE carrier remains an outstanding challenge. At this stage, an extensive set of observational data exist which any ERE carrier candidate must meet. Finding some material which produces approximately the correct spectrum of some ERE source is simply not suﬃcient. We now summarize the observational information about the ERE and try to deduce the resulting constraints for models.

• ERE is a PL process. • The ERE has been observed in many astrophysical environments: reflection nebulae, HII regions, dark nebulae (illuminated by the ISRF), high-latitude Galactic diffuse ISM and cirrus, dusty ISM of external galaxies, carbon-rich planetary nebulae • Carriers must consist of cosmically abundant refractory elements that are de- pleted from the gas phase in interstellar space and are capable of forming photoluminescent materials. Such elements include C, O, Si, Fe, and Mg. • Carrier particles must survive under a wide range of interstellar and circumstellar conditions and have a Galaxy-wide distribution. • The PL occurs in a broad, unstructured band with a peak wavelength varying between 600 nm and > 900 nm, within given sources and from source to source, in response to increasingly dense and hard illumination by UV photons. • The width and the peak wavelength of the ERE band are positively correlated. • Under conditions of low radiation density the ERE quantum yield is 10%. • The ERE is unpolarized with an isotropic radiation pattern. • The ERE peak wavelength and quantum yield are strongly affected by the density and hardness of the local UV photon field. 14

• The ERE is observed in the absence of scattering and appears therefore not associated with sub-micron-sized scattering grains.

• The ERE appears to be uncorrelated with the strength of the interstellar 2175 A˚ extinction feature and the ubiquitous UIR emission bands.

• The ERE carrier particles appear to originate in such dust-forming environments as the proto-planetary Red Rectangle nebula and in C-rich planetary nebulae.

1.2.3 Models for the ERE Carrier

None of the currently discussed “uniﬁed” models for interstellar dust (Draine,

2004; Dwek et al., 2004) speciﬁcally predict or account for the existence of ERE and

its observed characteristics. Only one attempt (Zubko et al., 1999) has been made

to simultaneously meet requirements posed by the wavelength dependence of extinc-

tion, abundance constraints, and ERE observations by including silicon nanoparticles

as a small-grain component into standard dust models. Lack of adequate knowl-

edge of the size-dependent dielectric functions of likely materials in the nanopar-

ticle regime is currently preventing further work in this direction. Most of the

past eﬀort has been focused purely on identifying classes of carrier particles capa-

ble of producing the ERE. These have included carbon-based solids, such as hydro-

genated amorphous carbon (HAC) (Duley, 1985), quenched carbonaceous composite (QCC) (Sakata et al., 1992), coal (Papoular et al., 1996), and bacterial pigments

(Hoyle & Wickramasinghe, 1999), as well as large carbon-based molecules, such as

PAHs (d’Hendecourt et al., 1986; Leger et al., 1988), C60 (Webster, 1993) and PAH-

clusters (Seahra & Duley, 1999). Non-carbon bearing ERE candidates such as silicon

nanoparticles (Witt et al., 1998; Ledoux et al., 1998) and silicates (Koike et al., 2002) 15 have also received close attention.

Hydrogenated Amorphous Carbon

Hydrogenated amorphous carbon (HAC) is an organic refractory material consisting of a mixture of sp2 and sp3 coordinated hydrocarbons with a variable band-gap.

The HAC ERE model attracted wide support early on, because not only did HAC exhibit PL in apparent agreement with observed ERE spectra (Witt & Schild, 1988), it also appears to be an interstellar dust component required to explain the interstellar 3.4 µm C-H stretch feature (Pendleton & Allamandola, 2002; Furton et al., 1999) and is widely employed as an integral component of models for interstellar dust. The early HAC PL spectra were obtained with exciting radiation near 500 nm wavelength and they displayed luminescence only in the red region with a spectrum similar to that of the ERE. Under UV illumination, however, the HAC PL spectrum extends well into the blue and near-UV, consistent with a band-gap of over 3 eV (Robertson,

1996; Rusli et al., 1996). Under astrophysical illumination conditions with ample UV photons, HAC would therefore exhibit PL in the 400 nm to 500 nm spectral range, contrary to observation. The HAC band-gap can be narrowed by irradiation or annealing, bringing it closer to the observational constraints, but this also results in a dramatic decrease of the PL quantum yield by three or four orders of magnitude. As a result, HAC is not able to reproduce the spectral characteristics and the observed quantum yield of the ERE simultaneously, and the HAC ERE model is no longer considered viable. 16

PAH Molecules and Ions

Large PAH molecules and their ions are considered likely candidates for explaining the ubiquitous UIR emission bands (discussed in more detail in §§ 1.3.1) in the ISM

and in many of the nebular sources in which ERE is observed as well. Given that

PAHs are also known to ﬂuoresce eﬃciently (Berlman, 1965), their role as a potential

source of ERE was proposed early on (d’Hendecourt et al., 1986). Under collision-free

conditions and with far-UV photons available for excitation, PAHs and their ions are

even expected to yield several luminescence photons per single excitation, resulting

in eﬃciencies of several 100% (Leger et al., 1988). However, neutral PAHs luminesce

predominantly in the UV/optical wavelength range (300 nm - 600 nm) (Birks, 1973;

Peaden et al., 1980), with similar eﬃciencies in the gas phase (Stockburger, 1973;

Sgro et al., 2001; Reyl´e& Br´echignac, 2000) as in solution (Berlman, 1965). Studies

of the charge state distribution of PAHs in diﬀerent interstellar environments by

Bakes & Tielens (1994) and Weingartner & Draine (2001) agree that a signiﬁcant

fraction of PAHs should still be neutral in radiation environments typical of reﬂection

nebulae. A size distribution of PAH molecules would therefore reveal its presence by

luminescing not only in the red ERE range but in the blue wavelength region as well.

Laboratory studies of the PL characteristics of larger PAH molecules, especially when

excited by far-UV photons in a collision-free environment, are needed to evaluate the

likelihood that PAH molecules are the origin of the ERE.

Relatively little is known about the ﬂuorescence characteristics of PAH cations.

The comparatively small energy diﬀerence (∼ 1 eV) between the electronic ground 17

state, D0, and the ﬁrst excited electronic state, D1, in most PAH cations (Leach,

1987) causes their potential ﬂuorescence to occur in the near-IR and not in the optical

range (Crawford et al., 1985). PAH cations are reported to absorb predominantly at

energies > 7.75 eV (< 160 nm) (Robinson et al., 1997), which results in an increased

probability of internal conversion to highly excited vibrational levels of the electronic

ground state D0, from where relaxation occurs essentially by emission of IR photons

(Leach, 1995). This suggests that quantum yields for ﬂuorescence by PAH cations

may be quite low. The dications are similar to neutral molecules having closed-

shell structures and singlet and triplet energy levels. Even less is known about the

ﬂuorescence properties of PAH dications.

The absence of consistent spatial correlations between ERE and UIR band emissions attributed to PAHs is not by itself conclusive evidence against the PAH origin of ERE, because PAHs in diﬀerent size ranges, degrees of hydrogenation, and charge states, all of which are expected to be environment-dependent, could be responsible for UIR bands and ERE to a diﬀerent extent (Duley, 2001).

PAH Clusters

Spectroscopic data on ﬂuorescence by speciﬁc PAH molecules are limited to sys-

tems with atomic weight < 500 amu or about 40 carbon atoms. Duley & Seahra

(1998) suggested that interstellar PAHs exist as much larger molecules or in the form

of stacks and aggregates of PAHs of up to 700 carbon atoms, which could be respon-

sible for the interstellar 2175 A˚ absorption feature as well as the UIR band emissions.

They also proposed (Seahra & Duley, 1999) that these same PAH clusters will pro- 18 duce PL emission with a principal band centered at 700 nm, which would account for the ERE. This model suﬀers from several weaknesses. The model predicts two emission side bands, with peaks at 500 nm and 1000 nm wavelength, to be associated with the main ERE band. A search for these additional bands in two of the strongest

ERE ﬁlaments in NGC 7023 by Gordon et al. (2000) failed to produce evidence for either one. The model predicts a remarkably constant peak wavelength for the main

ERE band for a wide range of size distributions of PAH clusters, contrary to the observational evidence. Finally, the model predicts correlations between the strengths of the 2175 A˚ band and the ERE band, which are not supported by observations.

The most attractive feature of this model is the eﬃcient utilization of the same dust component to explain several interstellar phenomena, namely the 2175 A˚ band, ERE,

UIR bands. It would be highly desirable to obtain laboratory data on PAH clusters which could support the multiple claims made for this model (Duley, 2001).

Other Carbonaceous Carrier Models

The C60 molecule exhibits a PL spectrum resembling that of the ERE in a few reﬂection nebulae, leading Webster (1993) to suggest C60 as the source of ERE. How-

+ ever, sensitive searches for the presence of C60 and C60 in interstellar space and in

NGC 7023 (Snow & Seab, 1989; Moutou et al., 1999; Herbig, 2000) failed to ﬁnd observable traces of this species. Furthermore, the measured quantum yield of C60 PL is only 8.5×10−4 (Kim et al., 1992), missing the ERE requirement by several orders of magnitude.

Discharges through gas mixtures containing hydrocarbons result in non-volatile 19

carbonaceous residues with interesting PL characteristics. Wdowiak et al. (1989), us-

ing a gas mixture containing CO, CH4, N2, H2O, and Ar, with relative abundances

approximating those believed to apply to the Red Rectangle, produced luminescing

residues of high quantum yield. However, under UV illumination the PL spectrum

consistently peaked at wavelengths < 600 nm, with substantial parts of the lumines-

cence spectrum extending through the 400 - 500 nm range. The spectral mismatch

with the observations plus the presence of blue PL appear to rule out this particular

candidate.

Quenched carbonaceous composite (QCC), a ﬁlmy residue produced through a

microwave discharge in a low-density methane plasma by Sakata et al. (1983), was

shown to be a more promising ERE carrier candidate. This material exhibits PL spec-

tra with variable peak wavelengths spanning the 680 - 725 nm range, depending on the

substrate temperature during deposition. Mass spectroscopy of QCC (Sakata et al.,

1983) indicates that QCC does not contain large PAH molecules but appears to con-

sist of a random agglomeration of 1- to 4-ring PAHs, small aliphatic molecules and

radicals. When dissolved in liquid freon, QCC exhibits blue luminescence, reminis-

cent of the ﬂuorescence of small PAHs. A more complete characterization of QCC

through additional laboratory work would clearly be desirable.

Silicon Nanoparticles

In recent years, silicon nanoparticles (SNPs) (Witt et al., 1998; Ledoux et al.,

1998, 2000, 2001, 2002) have emerged as an alternative to the carbonaceous ERE carrier candidates discussed above. The SNP model meets the observational con- 20

straints posed by ERE observations in respect to spectral variability and quantum

yield better than most of the carbonaceous candidates. The photophysics of SNPs

(Smith & Witt, 2002) also appears to be consistent with the observed variations of

ERE peak wavelength and ERE quantum yield with environmental conditions. The

balance between photo-ionization and recombination with free electrons, combined

with the fact that charged SNPs are much less likely to luminesce, explains the varia-

tion of the ERE quantum yield with UV-radiation density, while the expected photo

fragmentation of multiply charged SNPs leads to the erosion of the SNP size distribu-

tion, starting with the smallest sizes, and a corresponding shift to much longer ERE

peak wavelengths in astrophysical environments with the highest radiation densities.

Studies of the optical characteristics of SNPs (Amans et al., 2003) show the absorp-

tion coeﬃcient still rising at the wavelength of 200 nm, where current measurements

stop. This is consistent with ERE excitation requirements, but measurements at

shorter wavelengths are urgently needed.

The highly eﬃcient PL exhibited by SNPs is the result of quantum conﬁnement in nanocrystals and passivation of surface dangling bonds with atoms such as H,

O, N, C, and Fe. Most existing laboratory studies have been done with O- and

H-passivation, and it has been found that the surface composition aﬀects the PL spectrum (Wolkin et al., 1999; Zhou et al., 2003), in the sense that SNPs with H-

passivation and diameters of < 2.5 nm luminesce at blue and near-UV wavelengths,

while O-passivated SNPs luminesce in the red only. Witt et al. (1998) suggested that

interstellar SNPs would most likely be O-passivated. Once produced in a circumstellar

outﬂow, SNPs with oxygen passivation would persist even in the ISM with abundant 21

H-atoms, because the Si-O bond is energetically two to three times stronger than the

Si-H bond. Such particles would contribute to the widely observed 9.7 µm vibrational Si-O absorption band, generally referred to as the “silicate” feature. However,

Li & Draine (2002) have pointed out that stochastically heated O-passivated SNPs

would produce a 20 µm emission band well in excess of currently established observational limits. This objection might be overcome, if SNPs with O-passivation were either attached to or embedded in larger grains. Experiments, in which SNPs were produced by ion implantation in solids with subsequent annealing, have demonstrated that SNPs preserve their luminescent characteristics even when embedded

(e.g. Iwayama et al., 2002). Another way to overcome the objection by Li & Draine

(2002) is to consider other forms of passivation. For example, recent laboratory work

(Mavi et al., 2003) has shown that passivation with Fe-atoms yields even more per-

sistent and eﬃcient photoluminescence in SNPs than O-passivation. Carbon should

be able to provide passivation equally well, although experiments in this direction

have yet to be done. The existence of Fe- or C-passivated SNPs could also overcome

the observational constraint, which so far has restricted ERE detections to C-rich

environments among planetary nebulae (Furton & Witt, 1992). The SNP model,

while promising, also suﬀers from a number of loose ends. The study of the opti-

cal properties of SNPs must be pushed into the vacuum ultraviolet, the full range

of possible passivation schemes must be explored, the process of SNP formation in

stellar outﬂows must be examined, and the question of the long-term persistence of

the PL of SNPs under astrophysical conditions over astronomical times scales must

be investigated. 22

Thus the ERE is a well-characterized photo-luminescence property of dust existing

in a wide range of astrophysical environments. The carrier of the ERE is a major

component of interstellar dust: it intercepts ∼20% of the UV/optical photons in the diﬀuse ISM, if the intrinsic quantum yield is ∼50%, and a still higher fraction of the

UV/optical photons, if the yield is lower than ∼50%. Yet, no current comprehensive model for interstellar dust explicitly accounts for the existence of the ERE.

1.3 PAHs in the ISM

Polycyclic aromatic hydrocarbons (PAHs) are a class of very stable organic molecules made up of only carbon and hydrogen. These molecules are ﬂat, with each carbon having three neighboring atoms much like graphite. The structures of a variety of representative PAHs can be seen in Figure 1-4. On earth, they are a standard product

of combustion from automobiles and airplanes and some (such as benzo[a]pyrene) are

present in charcoal broiled foods. These molecules are considered highly carcinogenic

but they are also very common and stable in terrestrial environments. 23

Figure 1-4 Structures of some small PAHs

1.3.1 Unidentiﬁed Infrared Bands or the Aromatic Emission

Features

PAHs are ubiquitous in space and form the class of largest known molecules in

the ISM. The infrared spectra of a wide variety of objects associated with dust and

gas including the diﬀuse interstellar medium (ISM), the edges of molecular clouds,

reﬂection nebulae, young stellar objects, H ii regions, star forming regions, some

C-rich Wolf-Rayet stars, post-AGB stars, planetary nebulae, novae, normal galaxies, starburst galaxies, most ultra-luminous infra-red galaxies and AGNs are dominated by emission features at 3.3, 6.2, 7.7, 8.6, 11.2 and 12.7 µm (see Peeters et al., 2004, and 24

references therein). Approximately 20-30% of the Galactic IR radiation is emitted in

these UIR bands and 10-15% of the interstellar carbon is contained in the UIR carriers

(Snow & Witt, 1995), indicating that the carriers represent an abundant component of the ISM. Since the carriers of these features remained a mystery for almost a decade, the initial name for these features, the Unidentiﬁed Infrared (UIR) emission features, is still in use. We shall henceforth refer to these bands as the aromatic emission features (AEFs). The AEFs coincide with the vibrational modes characteristic of aromatic materials (Duley & Williams, 1981) and while the vibrational spectrum of

(almost) any exclusively aromatic material can provide a global ﬁt to the observed

AEFs, they are generally attributed to mostly large aromatic hydrocarbon species.

Figure 1-5 shows examples of rich AEF spectra. It should be noted that not all

sources show all these emission features at the same time and their peak position and

relative strength vary. 25

Figure 1-5 The ISO-SWS spectra of the planetary nebula NGC 7027 and the Photo- Dissociation region at the Orion Bar illustrate the richness and variety of the UIR/AEF spectrum. Also indicated are the aromatic mode identiﬁcations of the major UIR features (from Peeters et al., 2004).

1.3.2 Optical Fluorescence

Although the AEFs are attributed to PAHs, their speciﬁc sizes and ionization

states remain elusive. The vibrational transitions responsible for the AEFs are largely

independent of size, structure and ionization state of the molecule. However, elec-

tronic fluorescence, a transition from the upper excited level to the ground state, is 26 more specific (Reylé& Bréchignac, 2000). The processes which occur between the absorption and emission of light are best illustrated by a Jablonski diagram as shown in

Figure 1-6. The S’s represent singlet states and the T’s triplet levels. Each electronic level has many vibrational and rotational levels, only vibrational levels are shown in the figure. The absorption of a photon leads to an electron being raised from the ground state to one of several vibrational levels of the upper excited states. With a non-zero probability, this electron loses its excess vibrational energy by relaxing to the first excited state with the emission of infra red photons. If there are no other competing channels of relaxation, the molecule fluoresces as the electron relaxes from the zero vibrational level of the first excited state to one of the vibrational levels of the ground state. Provided the lifetime of the excited singlet state S1 is sufficiently long for the excited molecules to attain thermal equilibrium, the fluorescence emission occurs primarily from the lowest vibrational levels of S1. The fluorescence spectrum thus displays the vibrational spacing of the ground state. These are fast transitions and the intensities of the bands are governed by the Franck-Condon principle. Because electronic excitations do not greatly alter nuclear geometry, the vibrational energy levels of the excited state are similar to those in the ground state, and the absorption and fluorescence spectra display symmetric patterns. Another general property of

ﬂuorescence is that the same ﬂuorescence spectrum is generally observed irrespective of the excitation wavelength. This is known as Kasha’s rule. 27

IR IR : infrared emission IC ISC : inter−system crossing S 2 IC : internal conversion

IC ISC S1 IR ISC T absorption 1

IR IR fluorescence phosphorescence

Figure 1-6 Jablonski diagram

Laboratory Measurements of PAH Fluorescence

Laboratory spectra of PAHs are mostly obtained in solution. Gas-phase spectra are more relevant for comparison with astronomical comparison, and there are significant differences between spectra obtained in different ways as demonstrated by

Chi et al. (2001a,b). Figure 1-7 illustrates the differences in case of anthracene and pyrene. Spectra obtained at different temperatures also exhibit systematic differences as shown in Figure 1-8. Therefore only gas-phase spectra obtained at appropriate temperatures should be used for comparison with astronomical fluorescence spectra. 28

Figure 1-7 Fluorescence spectrum of anthracene and pyrene in gas-phase and in solution (from Chi et al. (2001a))

Figure 1-8 Fluorescence spectrum of anthracene and pyrene at diﬀerent temperatures (from Chi et al. (2001a)) 29

1.3.3 Formation of PAHs in the ISM

Most of the carbon in the outﬂows from C-rich AGB stars is tied up in CO and

C2H2. Since CO is stable, acetylene and its radical derivatives are likely to be the dominant precursor molecules from which PAHs are formed. Figure 1-9 depicts two routes of PAH formation under diﬀerent pressure/temperature conditions.

Figure 1-9 Formation routes for PAHs in the ISM (from Keller, 1987) 30

1.4 Outline

The chapters in this dissertation follow the format of papers in the Astrophysical

Journal, with sections on introduction, observation, results and discussion. The refer-

ences are consolidated at the end and footnotes to chapter titles provide information

on the publication details. Chapter 2 presents the discovery of the blue luminescence

(BL) emission in the Red Rectangle (RR) and its tentative identiﬁcation as ﬂuores-

cence by small, neutral PAHs. Chapter 3 presents details of further investigation of

the BL in the RR, the spatial correlation of the BL distribution with the 3.3 µm

PAH emission and a study of the attenuation characteristics of the central star in the

RR. These provided further evidence for the identiﬁcation of the BL as ﬂuorescence

from 3-, 4- ringed PAHs. In Chapter 4 we present the detection of BL in other or-

dinary reﬂection nebulae, establishing the BL as a more general photoluminescence

phenomenon exhibited by interstellar dust. In Chapter 5 we present our investigation

of the excitation wavelength of the ERE, which points to PAH dications as possible

carriers of the ERE. In Chapter 6 we report on the results of a narrow-band imaging of the RR, which reveals the changing morphology of the nebula as probed by the diﬀerent emissions and in Chapter 7 we present the detection of the resolved BL spectrum uncontaminated by scattered light at large oﬀsets from the central star in

the RR. Finally we summarize the main results of this work on Chapter 8. Chapter 2

Discovery of Blue Luminescence in

the Red Rectangle: Possible

Fluorescence from Neutral

Polycyclic Aromatic Hydrocarbon

Molecules?1

Here we report our discovery of a band of blue luminescence (BL) in the Red

Rectangle (RR) nebula. This enigmatic proto-planetary nebula is also one of the

brightest known sources of extended red emission as well as of unidentiﬁed infra-

1Appeared in The Astrophysical Journal Letters 2004, Vol. 606, pg. 65, Authors: U. P. Vijh,

A. N. Witt, & K. D. Gordon

31 32

red (UIR) band emissions. The spectrum of this newly discovered BL is most likely

ﬂuorescence from small neutral polycyclic aromatic hydrocarbon (PAH) molecules.

PAH molecules are thought to be widely present in many interstellar and circumstellar

environments in our galaxy as well as in other galaxies, and are considered likely

carriers of the UIR-band emission. However, no speciﬁc PAH molecule has yet been

identiﬁed in a source outside the solar system, as the set of mid-infra-red emission

features attributed to these molecules between the wavelengths of 3.3 µm and 16.4 µm is largely insensitive to molecular sizes. In contrast, near-UV/blue ﬂuorescence of

PAHs is more speciﬁc as to size, structure, and charge state of a PAH molecule. If the carriers of this near-UV/blue ﬂuorescence are PAHs, they are most likely neutral

PAH molecules consisting of 3-4 aromatic rings such as anthracene (C14H10) and pyrene (C16H10). These small PAHs would then be the largest molecules speciﬁcally identiﬁed in the interstellar medium.

2.1 Introduction

The family of emission bands at 3.3, 6.2, 7.7, 8.6, 11.2, & 12.7 µm (called the unidentified infra-red (UIR) bands) is found in almost all astrophysical environments including the diffuse interstellar medium (ISM), the edges of molecular clouds, reflection nebulae, young stellar objects, HII regions, star forming regions, some C-rich Wolf-Rayet stars, post-AGB stars, planetary nebulae, novae, normal galaxies, starburst galaxies, most ultra-luminous infra-red galaxies and AGNs (see

Peeters et al., 2004, and references therein). Approximately 20-30% of the Galac- 33

tic IR radiation is emitted in these UIR bands and 10-15% of the interstellar car-

bon is contained in the UIR carriers (Snow & Witt, 1995), indicating that the car-

riers represent an abundant component of the ISM. These UIR bands are the signa-

tures of aromatic C-C and C-H fundamental vibrational and bending modes, and are

generally attributed to a family of PAH molecules containing 50-100 carbon atoms

(Hony et al., 2001; Cook & Saykally, 1998; Leger & Puget, 1984; Verstraete et al.,

2001; Allamandola et al., 1985; Sellgren, 1984). Although the presence of PAHs in

space is widely accepted by most astronomers, their speciﬁc sizes and ionization states

remain elusive. On absorption of a far-UV photon, a PAH molecule usually undergoes

a transition to an upper electronic state. If the molecule undergoes iso-energetic tran-

sitions to highly vibrationally excited levels of the ground state, then the molecule

relaxes through a series of IR emissions in the C-C and C-H vibrational and bending

modes. These transitions are largely independent of size, structure and ionization

state of the molecule. However, electronic fluorescence, a transition from the upper excited level to the ground state, is more specific (Reylé& Bréchignac, 2000). In

particular, the wavelength of the ﬁrst electronic transition seen in neutral PAHs is

closely dependent on the size of the molecular species. In general, this ﬂuorescence

wavelength increases with the molecular weight of the molecule (Fig. 2-1). An ob-

servation of ﬂuorescence in an astronomical source in the UV/visible range oﬀers the

possibility of estimating the size of the PAH molecules, which the observation of the

UIR band emission in the same source does not. 34

5000

4500

4000 (Å) peak em λ 3500

3000

condensed PAHs non−condensed PAHs 2500 50 100 150 200 250 300 350 400 450 500 Mol. wt. (a.m.u.)

Figure 2-1 The ﬁrst electronic transition in the ﬂuorescence spectra (Berlman, 1965; Peaden et al., 1980; Ewald & Garrigues, 1985, 1988) of neutral PAH molecules as a function of their molecular weight. We distinguish between condensed and non- condensed PAHs, with condensed PAHs considered to be more stable under interstellar conditions.

2.2 Target and Observations

The Red Rectangle (RR) (Cohen, 1975; van Winckel, 2003) is a unique proto- planetary nebula. Its bipolar structure is a result of mass loss from an evolved central binary star, HD 44179, directed by a circumbinary disk. Diﬀerent aspects of the nebula become apparent in images at diﬀerent wavelength regions shown in

Figure 2-2. At blue wavelengths (3800 A˚ <λ< 4900 A˚ bottom panel) the nebular structure is dominated by a bright spherical blob of about 8” diameter, embedded in a faint rectangular envelope with approximate dimensions of 1200 × 2400. The blue 35 spectrum of the nebula is dominated by dust-scattered light of the central A-type star. At red wavelengths (6220 A)˚ the nebula appears to have sharp radial structures, representing the projected walls of a bi-conical outﬂow cavity, as shown in the high-resolution Hubble Space Telescope WFPC2 image in the middle panel. The nebular spectrum in the red region (λ > 5400 A)˚ is dominated by Extended Red

Emission (ERE) (Witt & Vijh, 2004; Witt & Boroson, 1990), following a spatial distribution (Schmidt & Witt, 1991) totally diﬀerent from that of the dust-scattered light at blue wavelengths (λ < 5000 A).˚ The carrier of the ERE is still under investigation (Witt & Vijh, 2004). The top panel shows a high-resolution short-exposure

WFPC2 near-IR image of the central source, demonstrating in fact that the star itself is not directly visible from Earth. Instead, we see two blobs of scattered light above and below a disc (Bujarrabal et al., 2003), which obscures the star at visible wavelengths.

The RR nebula has the distinction of being the brightest known source of UIR band emission (Russell et al., 1978; Geballe et al., 1985) which is attributed to PAHs.

In addition, the central star HD 44179 is thought to be a post-AGB star in its first stage of evolution toward a planetary nebula. It is in this stage that substantial mass- loss takes place and a bi-polar structure develops. This stage in stellar evolution is considered a stage of active dust-production (Whittet, 2003). Therefore, the RR nebula is a likely location in which to find observational evidence for small PAH molecules, which are the building blocks for larger PAH structures (Keller, 1987) and carbon grains. In most other sources of UIR band emission, the emitting particles are a component of interstellar dust, which is likely to have experienced substantial 36 processing by UV radiation, interstellar shocks and chemical modification over a very long time span of existence since leaving the environment of its original formation. It is well established that PAH molecules with less than 40 C-atoms are unlikely to survive the harsh radiation conditions of interstellar space (Jochims et al., 1999), while they may well be abundant in the more benign, UV-poor radiation field of HD 44179, especially when partially shielded by the opaque disk apparent in Figure 2-2. 37 ted gths ith the d on the bottom. In 2.5" WFPC2 F814W ˚ A), obtained using the 4-m AAT. The ˚ A - 4900 ˚ A. The geometry of the central illuminating source, of 6220 d CTIO 2, respectively. e image (3800 able structures of the RR nebula at different optical wavelen 30" ation and width of the two CTIO spectrograph slits are indica umstellar disc, is shown in the right-most panel, obtained w ˚ A. The angular scales of each of the three frames are indicate 80" AAT Blue WFPC2 F622W on both the bottom and middle panels, and are labeled CTIO 1 an obscured from direct visibility by an optically opaque circ HST WFPC2 at a wavebandall centered panels at north 8140 is to the right and east to the top. The orient middle panel shows a HST WFPC2 image centered at a wavelength and different angular scales. The left-most panel shows a blu Figure 2-2 The RR nebula. These images illustrate the observ 38

We obtained low-resolution, long slit spectra of the RR with the R-C (Cassegrain) spectrograph at the Cerro Tololo Inter-American Observatory (CTIO) 1.5-m telescope. These observations were made on March 26, 2003, using a 200.5 wide slit, 7.07

long. The grating #09 with 300 l mm−1, blazed at 4000 A,˚ provided a 8.6 A˚ resolution

and a spectral coverage of 2600 A.˚ Using a CuSO4 ﬁlter to select the grating’s ﬁrst

order, the setup covered a wavelength range of 3400-6000 A˚ including the range from

Hβ to the Balmer discontinuity. The Loral IK CCD detector yielded a spatial scale

of 1.300 pixel−1 along the slit. We used a coronographic decker assembly to minimize

scattered light from the star while probing as much of the inner nebula as possible.

All observations were made using the full extent of the 7.07 long slit to get simulta-

neous sky observations. We obtained spectra at two nebular locations, 200.5 and 500

south of the central star HD 44179 with the slit in east-west direction. The nebular

exposures were bracketed by exposures of the central star. Individual exposures were

limited to 5-10 minutes on the nebula and 1 minute for the star, and 4-5 exposures

were obtained for each orientation. Data Reductions were carried out with IRAF 2.12

EXPORT, and all spectra were ﬂux calibrated via observation of standard stars.

2.3 Results

2.3.1 Line-depth Technique

We used the line-depth technique to initially detect and measure the ﬂuorescence

spectrum at each of the two nebular positions. This technique allows the identiﬁ- 39

cation and measurement of any continuous emission in a reﬂected spectrum, such

as represented by the blue spectrum of the RR, relying upon the comparison of the

depths of nebular and stellar spectral absorption lines. It works particularly well with

the strong hydrogen Balmer lines at 6562, 4861, 4340, 4101, 3970, 3889, 3835, 3798,

3770 A˚ and so on. In the absence of any emission processes such as continuous ﬂuores-

cence by PAH molecules, the relative line-depths in the reﬂected spectrum would be

identical to the corresponding ones in the spectrum of the illuminating star, assum-

ing that the wavelength-dependent scattering properties of the dust, e.g albedo and

phase function asymmetry, remain unchanged over the few-angstrom width of each

individual line. The presence of any continuous emission manifests itself by the lines

in the reﬂected spectrum having smaller line depths than those in the illuminating

source spectrum. The illuminating source in this case is the star HD 44179, having

an ideal spectral type A with strong Balmer lines. Also, at Teﬀ ≤ 9000 K the star is

not likely to produce Balmer line emission in the surrounding nebula. The decrease

in the depth of a line can be directly related to the amount of underlying continuous

emission (in this case, the ﬂuorescence intensity). As the star is enshrouded by the

disk and as the stellar spectrum is actually seen in reﬂected light, if there were any

ﬂuorescence in this zero-oﬀset spectrum, the line-depth technique would reveal any

additional ﬂuorescence at the oﬀset positions in the nebula.

As an example we show in Figure 2-3 the normalized spectra at the zero oﬀset

00 00 position and at an oﬀset 2.5 south, 3.9 east of Hζ at 3889 A.˚ One can see that the

line-depth at the nebular position is substantially diminished. 40

1.1

0.9

0.8

0.7 Normalised Intensity

0.6

0˝ 2.5˝S, 3.9˝E 0.5 3870 3880 3890 3900 3910 λ (Å)

Figure 2-3 Normalized spectra from two positions, zero offset and 200.5 south, 300.9 east at 3889 A,˚ illustrating the filling in of the Balmer line due to fluorescence. Error bars are photon-statistics errors.

In the presence of any continuous emission the relative line depth in the nebula

RN =(CN − LN )/CN is smaller than the relative line depth RS =(CS − LS)/CS of the same line in the illuminating star, where C and L refer to continuum and line center intensities, respectively, and where the subscripts N and S refer to the nebular and stellar spectrum, respectively. The ﬂuorescence intensity relative to the scattered nebular continuum can therefore be expressed as RS/RN − 1 (see Witt & Vijh, 2004, for a review).

Given the nature of our measuring technique, we can determine the ﬂuorescence intensity only at the wavelengths of the resolvable Balmer lines as well as the Balmer discontinuity, i.e. at ten wavelength positions with our spectral resolution in our 41

spectral range. By measuring the line-depths of each of the Balmer lines in the

star and the nebula, we extracted the ﬂuorescence intensities at those wavelengths.

Table 2.1 shows the ﬂuorescence intensities at the two oﬀsets: 200.5 south, 300.9 east

and 500 south, 600.5 east. Independent observations by the authors using the 90-in.

Bok telescope (Steward Observatory) and the B&C spectrograph show corroborating

evidence for this UV/blue ﬂuorescence in the RR (U. Vijh, A. Witt, & K. Gordon,

in preparation).

Table 2.1 Fluorescence Intensity at two positions in the RR nebula.

Wavelength Intensity (erg cm−2 s−1A˚−1sr−1)

200.5 south, 300.9 east 500 south, 600.5 east

4861.3 4.059±0.026 × 10−6 9.528±0.145 × 10−8 4340.5 1.209±0.008 × 10−5 1.571±0.025 × 10−6 4101.7 2.914±0.020 × 10−5 7.032±0.123 × 10−7 3970.1 3.565±0.023 × 10−5 3.441±0.048 × 10−6 3889.1 3.817±0.024 × 10−5 3.164±0.045 × 10−6 3835.4 3.517±0.022 × 10−5 1.881±0.029 × 10−6 3797.9 4.438±0.024 × 10−5 3.892±0.047 × 10−6 3770.1 5.428±0.023 × 10−5 5.483±0.048 × 10−6 3749.8 6.417±0.019 × 10−5 5.666±0.034 × 10−6 3570.0 1.000±0.019 × 10−5 1.180±0.034 × 10−6

Errors are combined photon-statistics errors

2.3.2 Identiﬁcation

Hydrogenated amorphous carbon (HAC), silicon-nanoparticles (SNP), SiC grains, and small PAHs are all known to luminesce at blue wavelengths and have also been considered as likely candidates to explain the ERE (Witt & Vijh, 2004). Figure 2-4 42 shows ﬂuorescence spectra of HAC, SNP and SiC (Li et al., 2000; Huang et al., 2002;

Liao et al., 2002; Belomoin et al., 2000; Patrone et al., 2000). It reveals that the observed BL in the RR is unlike the luminescence spectrum of HACs, SNPs, or SiC.

Given the exceptionally strong UIR-band emission in the RR, the most likely source of the newly discovered BL may then be ﬂuorescence by PAH molecules. The nebular BL spectrum suggests a peak in the ﬂuorescence intensity corresponding the

ﬁrst electronic transition S1 - S0 near 3750 A.˚ Comparison with Figure 2-1 suggests

PAH molecules in the mass range of 170 to 270 amu as the likely sources. Shown in

Figure 2-4 are comparisons with laboratory ﬂuorescence spectra (Chi et al., 2001a,b;

Reylé& Bréchignac, 2000) of gas-phase PAHs in that mass range with the BL spectrum of the RR at an offset of 200.5 south, 300.9 east. 43

1 RR−2.5"S, 3.9"E + HAC film from CH3 0.9 HAC film H−capped SNP 0.8 SNP − Patrone SiC 0.7

0.6

0.5

0.4

Normalized Intensity 0.3

0.2

0.1

0 3000 3500 4000 4500 5000 5500 6000 6500 7000 Wavelength λ (Å) 1 RR−2.5"S, 3.9"E Pyrene 0.9 Anthracene Phenanthrene 0.8 Chrysene Naphthalene 0.7

0.6

0.5

0.4

Normalized Intensity 0.3

0.2

0.1

0 3000 3200 3400 3600 3800 4000 4200 4400 4600 4800 5000 Wavelength λ (Å)

Figure 2-4 The UV/blue fluorescence spectra of the RR nebula compared to PAH and solid-state luminescence. (a), comparison with some solid-state candidates and (b), comparison with small neutral PAHs. The intensities are normalized for easy comparison of the spectra with existing laboratory fluorescence spectra. The dashed lines representing the RR fluorescence spectrum are drawn to guide the eye (measurements exist only at specific wavelengths indicated by solid circles). 44

Laboratory spectra obtained with PAHs in the gas-phase are more likely to be com-

parable to astrophysical spectra. These gas-phase spectra show considerable variation

with temperature, becoming smoother, losing band-structure and the peak shifting

to longer wavelengths with increasing temperatures (Chi et al., 2001a,b). Though

anthracene suggests a good match to the observed spectrum, we do not have suf-

ﬁcient spectral resolution to completely rule out other possible candidates, as the

measurement technique works only at the discrete wavelengths of the Balmer lines.

2.4 Discussion

Another important factor which should be considered when identifying possible carriers is fluorescence efficiency. Although many different fluorescing species may exist with similar abundances, the spectra of the most efficient ones will be dominant.

Laboratory studies have shown that on comparison of emission rates per molecule, from PAH molecules placed at 1 AU from the Sun, anthracene and pyrene show

10-100 times more eﬃciency than naphthalene (C10H8) and phenanthrene (C14H10)

(Br´echignac & Hermine, 1994). Additional support for small PAHs as sources of the

observed BL is found in the spatial correlation between the BL and the PAH 3.3 µm emission in the RR and the presence of the distinct PAH ionization discontinuity in the spectrum of the central source, HD 44179 (U. Vijh. A. Witt, & K. Gordon, in preparation). Chapter 3

Small PAHs in the Red Rectangle 1

Following our initial discovery of blue luminescence in the spectrum of the Red

Rectangle (RR) and its identiﬁcation as ﬂuorescence by small three- to four-ringed

polycyclic aromatic hydrocarbon (PAH) molecules, we report on the spatial correla-

tion between the blue luminescence and the 3.3 µm emission, commonly attributed to small, neutral PAH molecules, and on the newly-derived UV/optical attenuation curve for the central source of the RR, HD 44179. Both results provide strong additional evidence for the presence of small PAH molecules with masses of less than

250 amu in the RR, which supports the attribution of the blue luminescence to ﬂuo- rescence by the same molecules. We contrast the excellent spatial correlation of the two former emissions with the distinctly diﬀerent spatial distribution of the extended red emission (ERE) and of the dust-scattered light within the RR. The UV/optical attenuation curve of the central star is unlike any interstellar extinction curve and is

1Appeared in The Astrophysical Journal, Vol. 619, pg. 368, Authors: U. P. Vijh, A. N. Witt,

& K. D. Gordon

45 46

interpreted as resulting from circumstellar opacity alone. Major contributions to this

opacity are absorptions in broad bands in the mid-UV, contributing to the electronic

excitation of the luminescing PAH molecules, and a sharp ionization discontinuity

near 7.5 eV in the far-UV, which places a sharp upper limit on the masses of the

PAH molecules that are responsible for this absorption. The strength of the far-UV

absorption leads to an abundance of the PAH molecules of 10−5 relative to hydrogen

in the RR. Such small PAHs are perhaps unique to the environment in the RR, where

they are shielded from harsh radiation by the dense circumstellar disk.

3.1 Introduction

Recently, we reported the discovery of blue luminescence (BL)(Vijh et al., 2004)

in the spectrum of the bipolar, proto-planetary Red Rectangle (RR) nebula (Cohen,

1975; Cohen et al., 2004). The emission, peaking near 378 nm, is contained in a

band of FWHM ∼ 45 nm. The band-integrated intensity of the BL is comparable to that of the scattered light intensity and of the intensity of the extended red emission (ERE), which is also exceptionally strong in this object (Schmidt et al., 1980;

Witt & Boroson, 1990). This suggests that the BL is produced by a carrier of sub-

stantial abundance in the gas/dust medium in the RR.

Through a comparison of the observed spectrum with laboratory photolumines-

cence spectra of a collection of likely dust components and with gas-phase ﬂuores-

cence spectra of polycyclic aromatic hydrocarbon (PAH) molecules, Vijh et al. (2004)

identiﬁed neutral three- and four-ring PAH molecules such as pyrene, anthracene 47

and phenanthrene as the most likely sources of the BL. This identiﬁcation also ap-

pears consistent with several other important facts. Firstly, the RR is the brightest

known source of emission in the unidentiﬁed infrared (UIR) features (Russell et al.,

1978; Geballe et al., 1985), which are commonly attributed to PAH molecules and

ions (Bakes et al., 2004); this suggests high relative abundances and optimal excitation conditions for such molecules. Also, as a post-AGB star (van Winckel, 2003),

the carbon-rich central star of the RR, HD 44179, is in an active mass-losing and

dust-producing stage of its life, suggesting that PAH molecules are currently con-

densing in the carbon-rich outﬂow (Keller, 1987; Cherchneﬀ et al., 1992). This in-

cludes, in particular, smaller PAH molecules that are able to exist in the relatively

benign radiation environment of the RR, but which would not be viable after being

ejected into the much harsher radiation ﬁeld of interstellar space, where three- and

four-ringed PAHs are not expected to survive (Jochims et al., 1994, 1999). And ﬁ-

nally, anthracene and pyrene have exceptionally high ﬂuorescence eﬃciencies, which

would cause their spectra to dominate any ﬂuorescence spectrum of a mixture of PAH

molecules (Br´echignac & Hermine, 1994). If conﬁrmed by further observations, the

detection of three- and four-ringed PAH molecules would represent the detection of

the largest speciﬁcally identiﬁed molecules so far observed outside the solar system.

In this paper we report new observations and analysis dealing with the spatial

distribution of the BL in the RR nebula. We show that the spatial distribution

of the BL is distinctly diﬀerent from that of the scattered light distribution in the

same region of the spectrum and also distinctly diﬀerent from that of the ERE.

Signiﬁcantly, there is an exceptionally close correlation between the distribution of 48

the BL and the distribution of the UIR-band emission at 3.3 µm, attributed to the

C–H stretch transition in neutral PAH molecules, which provides strong support for the identiﬁcation of the BL as ﬂuorescence from relatively small PAH molecules. In the second part of this paper we discuss the newly-derived attenuation curve for the central star, HD 44179. The attenuation is shown to be circumstellar rather than interstellar in nature and to be dominated by absorptions attributable to electronic bound-bound and bound-free transitions in small PAH molecules.

3.2 Observations

Low-resolution, long slit spectra of the RR were obtained with the R-C (Cassegrain) spectrograph at the Cerro Tololo Inter-American Observatory (CTIO) 1.5 m telescope. These observations were made on 2003 March 26 and 29, using a 200.5 wide slit,

7.07 long. For observations in the blue (340 - 600 nm), grating #09 with 300 l mm−1

was used which is blazed at 4000 A,˚ provides a 8.6 A˚ resolution and a spectral cov-

erage of 260 nm. Using a CuSO4 ﬁlter to select the grating’s ﬁrst order, the setup

covered a wavelength range of 340-600 nm, including the range from Hβ to the Balmer

discontinuity. Grating #11 with 158 l mm−1, blazed at 8000 A,˚ with a 16.4 A˚ reso-

lution was used in the ﬁrst order with a GG495 cut-on ﬁlter for the red observations

with a wavelength coverage from 500-1000 nm. The Loral IK CCD detector yielded a

spatial scale of 100.3 pixel−1 along the slit. A coronographic decker assembly was used

to minimize scattered light from the star while probing as much of the inner nebula

as possible. All observations were made using the full extent of the 7.07 long slit to 49

get simultaneous sky observations. Spectra were taken at two nebular locations, 200.5

and 500 south of the central star HD 44179 with the slit in E-W direction (PA = 90◦).

The nebular exposures were bracketed by exposures of the central star. Individual

exposures were limited to 5-10 minutes on the nebula and 1 minute for the star, and

4-5 exposures were obtained for each orientation. Data reductions were carried out

with IRAF 2.12 EXPORT, and all spectra were ﬂux calibrated via observation of

standard stars.

The 3.3 µm data are reproduced from Kerr et al. (1999). The spectroscopic ob-

servations were made at the United Kingdom Infrared Telescope (UKIRT) on 1995

December 11. The spectra were taken with a slit 9000 long, 100.2 wide, placed 500 south of the central star HD 44179 (PA = -85◦).

For the FUV spectral energy distribution (SED) of HD 44179 archival observations from the International Ultraviolet Explorer (IUE) LWP 22416, SWP 38188, and LWR

04273 were used.

David Malin (David Malin Images) provided us with one of the few existing deep blue images of the RR for the comparison of our spatially resolved spectra with morphological details of the nebular structure. The image, reproduced in Figure 3-1,

was taken by David Malin with the AAT at the Anglo-Australian Observatory. The

exposure time was 15 min, the detector was a Kodak IIa-O plate (blue sensitive) with

a GG385 cut-on ﬁlter. The eﬀective bandpass extends from 390 nm to 480 nm. A red

image of the RR taken with the HST has been reproduced to aid in the correlation

of the observed features with the nebular structure. 50

3.3 Analysis and Results

3.3.1 Blue Luminescence in the Red Rectangle

Nebular spectra were extracted in two-pixel bins along the two slits with 7 aper-

tures along the 200.5 south slit and 8 apertures along the 500 south slit with a spatial

scale of 200.6 per aperture. Thus, each spectrum represents a 200.5 × 200.6 region of the

nebula. Figure 3-1 shows the blue image of the nebula overlaid with the BL aper-

tures and the 3.3 µm slit and Figure 3-2 is a red image overlaid with the central ERE apertures. The line-depth technique (Vijh et al., 2004; Witt & Vijh, 2004) was used

to detect and measure the BL at the positions of the hydrogen Balmer lines for each

spectrum. Measurement of the Balmer discontinuity in the nebular spectra compared

to that in the stellar spectrum is used to infer the BL at the Balmer discontinuity

(see Appendix A for details). Measurement of the equivalent widths of the absorp-

tion lines could also be used for a similar analysis, however a small uncertainty in

the determination of the continuum levels over the steep spectral shape results in

a relatively large uncertainty in the equivalent width. In contrast the errors in the

line-depth measurements are essentially limited to the photon noise in the spectrum.

The co-added nebular spectra had a signal-to-noise ratio per pixel of ∼ 10 in the outermost regions to ∼ 250 closer to the central star in the blue and a range of ∼

20 - 225 in the red region of the spectrum. The stellar spectra had a signal-to-noise

ratio per pixel of ∼ 500 in both spectral ranges. 51

Figure 3-1 Blue image of the Red Rectangle with overlaid extraction apertures. The size of the image is 3000 × 3000. The white dot is an overlaid pixel to indicate the position of the central source. The dashed lines indicate the slit used for the 3.3 µm observation. North is to the top, east to the left. 52

Figure 3-2 Red image of the RR with overlaid extraction apertures. The size of the image is 1900.4 × 2000.2. The HST image (Credits: NASA, ESA, Hans Van Winckel (Catholic University of Leuven, Belgium), and Martin Cohen (University of Califor- nia, Berkeley)) has been reproduced to aid the identiﬁcation of the nebular structure with the extracted apertures. North is to the top, east to the left.

Spectral Variability in the BL

The BL spectrum is not identical in diﬀerent regions of the nebula and certain trends can be noted. The primary peak in the spectrum is near 378 nm (Vijh et al.,

2004, Fig. 4). A secondary peak around 397 nm starts to develop as we probe regions farther from the central star, more prominently along the 500 south slit. Also, farther 53 from the star, the BL intensities at longer wavelengths start to become stronger compared to the peak intensity. Interestingly, spectra that probe regions inside the conical outﬂow indicate that the BL in these regions has another peak at much shorter wavelengths, shortward of 360 nm. Figure 3-3 depicts the BL spectrum at three such representative locations: 200.5 south, 700.8 east, a position in the outer regions of the nebula; 200.5 south, 200.6 east, a region close to the central source, inside the bipolar outﬂow; 500 south, 500.2 west, a region along the farther slit just outside the cone wall.

Based on the existing correlation between the wavelength of peak intensity of the

ﬂuorescence spectrum and the molecular size of PAH molecules (Vijh et al., 2004), these trends can be interpreted to suggest a size variation in the BL carrier at different locations in the nebula: closer to the central source we see an emergence of a smaller population of PAHs and farther out we see increasing contributions from larger emitters. Given the limited spectral resolution inherent in our detection technique and only fair spatial resolution these trends are at best an indication of the size distribution of the BL emitters. We defer a more complete analysis of the spectral variation of the BL spectrum until after completion of an ongoing program in col- laboration with D.G. York which maps the complete RR nebula with dense spatial coverage using the 3.5-m telescope at Apache Point Observatory. 54

1 2".5s−7".8e 2".5s−2".6e 5"s−5".2w 0.8

0.6

0.4 Normalised Intensities 0.2

3400 3600 3800 4000 4200 4400 4600 4800 5000 λ (Å)

Figure 3-3 Spectral variation of the BL. The solid line traces the intensity at 200.5 south, 700.8 east, the dashed line at 200.5 south, 200.6 east and the dash-dotted line at 500 south, 500.2 west positions. Solid circles indicate wavelength positions associated with the 200.5 south slit and solid squares indicate wavelength positions associated with the 500 south slit. Intensities are normalized to unity at the peak positions.

Gradients of BL/Scattered Light

Figure 3-4 (upper panel) shows the distribution of the ratio of band-integrated BL to the total scattered light in the same band at various positions along the two slits.

Figure 3-4 (lower panel) illustrates the spatial distribution of the BL and the scattered

light intensities on a normalized scale. Due to the anisotropy of the phase function,

scattered light is strongly forward directed and falls oﬀ steeply as we probe regions

where our line of sight penetrates regions of the nebula where larger scattering angles

predominate. The BL, on the other hand is due to ﬂuorescence and is an isotropic

emission and thus falls off less steeply. Thus, the spatial variation of the BL and 55 scattered light along the two slits reveals the isotropic nature of the BL. Therefore, the ratio IBL/IScat increases away from the star, whereas the actual intensities fall away from the star as the density of emitters and the exciting radiation both decrease. It is equally illuminating to note that the BL intensities have similar profiles along the two slits, whereas the distribution of the scattered light it quite different; along the slit closer to the star, the scattered light falls off more steeply than along the second slit as the relative change in the distance from the central star varies more slowly in the latter case. 56 10 2.5s-BL/Scat 5s-BL/Scat

BL/Scattered Light 0.1

0.01 -15 -10 -5 0 5 10 15 offset (arcsec) 10 2.5s-BL 5s-BL 2.5s-Scat 5s-Scat

0.1 Normalised Intensities 0.01

0.001 -15 -10 -5 0 5 10 15 offset (arcsec)

Figure 3-4 Ratio of Band-IBL to Band-Isc and normalized distributions. upper panel: Ratio of band-integrated BL to scattered light along the two slits. Solid line traces the ratio along the 200.5 south slit and the dashed line that along the 500 south slit. lower panel: Normalized BL and scattered light intensities. Intensities are normalized to unity at 200.6 east position. In both ﬁgures, solid circles denote spatial positions along the 200.5 south slit and the solid squares those along the 500 south slit. Positive and negative oﬀsets indicate positions to the west and east of the star HD 44179 respectively. 57

Correlations with Other Emissions

Figure 3-5 displays the spatial distribution of the band-integrated intensities of

ERE and the BL along the two slits and the 3.3 µm UIR-band emission along a 500 south (PA=-85◦) slit (Fig. 3-1 and 3-2 show the apertures). The 3.3 µm intensities

(obtained from Kerr et al. (1999)) have been normalized to the intensity of the BL at

the 0 east-west, 500 south position. This ﬁgure reveals a number of interesting facts.

The 3.3 µm UIR-band emission and the BL are almost perfectly correlated, which we will discuss in more detail in the next paragraph. The spatial distribution of the

BL is distinctly different from that of the ERE, the ERE is strongly peaked close to the central source and on the X-shaped walls of the outflow cone while the BL has a broader, more diffuse distribution.

The 3.3 µm emission correlates exceptionally well with the BL along the 500 south

slit, and the small diﬀerences are probably attributable to the slightly diﬀerent widths

and PA of the two slits (see Fig. 3-1). The 3.3 µm emission feature belongs to the distinctive set of mid-infrared emission features which have been attributed to emission from PAH molecules, which become vibrationally excited upon absorption of a

UV photon and subsequently relax with the emission of IR photons (Bakes et al.,

2004; Schutte et al., 1993), in addition to electronic relaxation through ﬂuorescence

or phosphorescence transitions. Most of the vibrational transitions are insensitive

to molecular sizes, but the 3.3 µm emission attributed to C–H stretch from PAH molecules is sensitive to the size of the molecular species. PAHs having ∼ 20 carbon atoms show 3.3 µm emission intensities ∼ 50 times stronger than that from PAHs 58

having ∼ 100 carbon atoms (Schutte et al., 1993). Earlier observations of the RR

showed a pronounced lack of similarity between the 3.3 µm and the 11.3 µm surface brightness distributions (Bregman et al., 1993), the latter being attributed to

C–H out-of-plane bending modes in relatively large PAH molecules. Furthermore,

the 3.3 µm emission traces neutral PAH species (Bakes et al., 2004), and becomes ex-

tremely weak in ionized PAHs. The fact that the BL follows the 3.3 µm distribution so closely indicates that the BL emitters and the 3.3 µm emitters are similar small, neutral PAHs and thus strengthens our identiﬁcation of the BL as ﬂuorescence by small neutral PAHs.

The ERE, BL and the 3.3 µm emission are all isotropic emissions and yet the spatial distributions of the ERE is dissimilar to those of the BL and the 3.3 µm emission.

The ERE falls oﬀ much more steeply than the other two that have a broader distribution. The ERE is largely conﬁned to the bipolar, X-shaped structure of the nebula

(see also Schmidt & Witt, 1991, Fig. 2,), whereas the BL is the dominant emission

in the outer regions. The variation in the spatial distributions of these diﬀerent emis-

sions must be attributed the distribution of the carriers and the respective exciting

radiations. The BL carriers, most likely small PAHs, are probably ionized inside the

bipolar outﬂow cone, where they are exposed to direct illumination from the central

RR source. Thus ﬂuorescence in these regions is eﬀectively quenched. Farther out,

outside the cones, the molecules/carriers are shielded from the far-UV ionizing radia-

tion but still receive through scattering the mid-UV exciting radiation and thus emit

with fairly high intensities. The ERE carriers on the other hand seem more robust

as the ERE intensities are high closer to the star inside the bipolar outﬂow cone and 59 along the walls (X-shaped structure). This could, in fact be interpreted as pointing toward some unknown ionized species as the carrier of the ERE. Outside the outﬂow region the ERE intensities drop by over three orders of magnitude, suggesting that either the carriers of the ERE do not exist in such regions or if they do exist they do not receive the exciting radiation needed to produce the ERE. 60

1 Band-ERE-2".5s Band-BL-2".5s

0.1 E W

0.01

Intensities 0.001

0.0001

1e-05 -15 -10 -5 0 5 10 15 offset (arcsec) 1 Band-ERE-5"s Band-BL-5"s 3.3-5"s 0.1 E W

0.01

Intensities 0.001

0.0001

1e-05 -15 -10 -5 0 5 10 15 offset (arcsec)

Figure 3-5 Spatial distribution of the BL and correlations with other emissions. Band- integrated intensities of BL and ERE at positions along the two slits (upper panel: along the 200.5 south slit and lower panel: along the 500 south slit) are plotted. The 3.3 µm emission is normalized to the value of the BL at 500 south, zero oﬀset position. Filled circles indicate positions along the 200.5 south slit and ﬁlled squares those along the 500 south slit. The intensities are band-integrated and are given in units of ergs cm−2 s−1sr−1 for the BL and the ERE. 61

3.3.2 Attenuation of HD 44179

Visual Reddening and Attenuation

The central source of the RR, HD 44179, is a single-line spectroscopic binary

star (Waelkens et al., 1992, 1996; Van Winckel et al., 1995; Men’shchikov et al., 2002)

that is totally hidden from direct view at optical wavelengths by a circum-binary disk

seen nearly edge on. Diﬀraction-limited imaging (Cohen et al., 2004) reveals the

central source as a pair of scattered-light lobes above and below the disk, directing

light from the hidden central stars via scattering by dust toward the Earth. The

photometrically dominant component of the binary was classiﬁed as spectral type

B9/A0 (Cohen, 1975) originally, but is now recognized as a post-AGB star with an

eﬀective temperature near 8000 K and a highly metal-deﬁcient atmosphere typical

of post-AGB stars (van Winckel, 2003; Men’shchikov et al., 2002). The low metal

abundance, e.g. [Fe/H] = -3, gives this star the spectroscopic appearance of a much

hotter B9/A0 star, but with an intrinsic (B−V ) color characterized by its actual eﬀective temperature. By convolving the appropriate Kurucz atmosphere model (Teﬀ

= 8250 K, log g = 1.5, [Fe/H] = -3.0) with the B and V pass bands, we determine an intrinsic (B−V ) = -0.04 mag for the photometrically dominant component of

HD 44179. We note that the nature of the secondary component is not known directly from observation, but its inferred mass of 0.35 M and the fact that the innermost region of the RR contains a small H ii region suggest that it is most likely a hot white dwarf with a luminosity of less than 2% of that of the primary (Men’shchikov et al.,

2002). 62

The observed (B−V ) color of HD 44179 has been inferred from broadband photometry of the RR by Cohen (1975) as 0.39 mag., leading to an estimated color excess

of E(B−V) = 0.35 mag. The peculiar nature of HD 44179, its location within an optically thick disk, and its uncertain distance make it diﬃcult to estimate the total visual attenuation of its light as seen from Earth. Here we adopt the latest model

(Men’shchikov et al., 2002) which assigns a distance of 710 pc and a total luminosity

of 6050 L to the central object, with uncertainties of 10% and 20%, respectively.

Given a visual brightness of V = 8.83 (Cohen, 1975), we then estimate a visual at-

tenuation of AV = 4.2 mag. Here, we deliberately distinguish between attenuation

and extinction. By extinction we understand the partial obscuration of direct star light due to absorption and scattering by dust along the direct line-of-sight to the star, while attenuation measures the reduction of observable ﬂux due to absorption and scattering in a complex geometry, where the residual observable ﬂux may indeed only be scattered light re-directed into the line of sight while the direct light is totally absorbed, as is the case in the RR.

The attenuation of the light of HD 44179 is characterized by a ratio RV =

AV/E(B−V ) = 11.9, much larger than the average value of RV ∼ 3.1 encountered with interstellar extinction. The attenuation in the UV/visible portion of the spectrum of HD 44179 is therefore much grayer than typical interstellar extinction, an eﬀect attributed in the past to extinction by larger grains in the RR

(Men’shchikov et al., 2002). We consider it more likely that the large value of RV arises from the fact that only scattered light from HD 44179 can be observed at optical wavelengths. Upon ﬁrst scattering, scattered light is initially bluer than that 63 of the illumination source, but upon transfer through an optically thick medium and multiple scattering, scattered light also reddens but at a slower rate than expected from the optical path length covered, as shown in the example of the colors of reﬂec- tion nebulae by Witt (1985). The high value of the estimated visual attenuation of

◦ AV = 4.2 mag and the fact that scattered light must be re-directed by 90 require that multiple scattering is the dominant process for the transfer of optical radiation from the center of the RR.

UV/Optical Attenuation Curve

In Figure 3-6 we display the observed (attenuated) UV/optical SED of HD 44179 over the wavelength interval from 130 nm to 600 nm. We produced this SED by combining the observed ﬂux distributions in the near- and far-UV obtained from the IUE archive (LWP 22416, SWP 38188, and LWR 04273) with our optical SED obtained from CTIO spectroscopy. No adjustments or corrections were applied; the continuity between the ground-based and space-based observations is almost perfect, indicating that the two sets of absolute calibrations are in excellent accord. Also shown in Figure 3-6 is a model SED for a metal-deﬁcient, giant-star using a Kurucz

(1993) stellar atmosphere for [Fe/H] = -3 , log g = 1.5 and Teﬀ = 8250 K, which we propose as a suitable representation of the intrinsic SED for the photometrically dominant component of HD 44179. The model SED is normalized to the observed

SED at V. The eﬀective temperature of this model atmosphere is higher than that proposed by Men’shchikov et al. (2002) by 500 K, but our choice was dictated by the requirement that the model and the observed SED should have the same relative 64 strength of the Balmer discontinuity such that the ratio of these two SEDs would be continuous across the Balmer discontinuity. Model SEDs with temperatures diﬀering from 8250 K by as little as 250 K exhibit noticeable discontinuities, when divided into the observed SED. 65

2.5e−12 Observed SED Model SED at 8250 K

2e−12

1.5e−12

1e−12 Flux

5e−13

−5e−13 1500 2000 2500 3000 3500 4000 4500 5000 5500 6000 λ (Å)

Figure 3-6 Observed UV/optical SED of HD 44179 compared to model SED. The observed SED has been obtained by combining the archived flux distributions in the near- and far-UV from IUE and optical flux distribution obtained by our CTIO spectroscopic measurements. The model SED is from a Kurucz (1993) stellar atmosphere for [Fe/H]=-3, log g =1.5, and Teff =8250 K, normalized to the observed SED at the V-band. The stellar fluxes are displayed in units of ergs cm−1s−1A˚−1. 66

Figure 3-7 displays the wavelength dependence of the UV/optical attenuation suﬀered by the light of HD 44179. The attenuation curve is normalized to E(B−V )

= 1, and it is shown in comparison with the average Galactic interstellar extinction curve with identical normalization. Several distinct diﬀerences between the two curves are immediately apparent. With the slopes forced to coincide in the B−V range, the observed attenuation curve rises more steeply in the near-UV compared to the interstellar extinction curve. There is no sign of the 217.5 nm absorption peak; instead we see a very much broader attenuation hump reaching its peak near λ−1 ∼ 5 µm−1.

Finally, near λ−1 ≈ 6.0 µm−1 the attenuation curve exhibits a sharp discontinuous rise, which is much more abrupt than the gentle rise shown by the extinction curve. 67

12 HD 44179 Attenuation Avg. IS extinction pyrene

6 - V)/E(B V) λ

E( 4 λβ

λβ′ λ p 2 λα

-2 1 2 3 4 5 6 7 λ-1 (µm)-1

Figure 3-7 UV/optical attenuation curve for HD 44179. The average interstellar extinction curve is shown for comparison with the solid line. Also shown is the expected contribution of PAH ionization to the far-UV part of the attenuation curve by pyrene as a representative example with a dashed line. Indicated schematically are the four characteristic PAH absorption bands, the α, para, β and the β0. The curves are normalized to E(B−V )=1. 68

As is well known, the perception of anomalous attenuation/extinction curves is strongly dependent upon their normalization (Fitzpatrick, 2004). In Figure 3-8, therefore, we display the attenuation curve for HD 44179 normalized at V, scaled to

AV = 1.0, and compare it to Galactic extinction curves with RV = 3.1, representative of the diﬀuse ISM, and RV = 5.5, representative of a molecular cloud environment and a grain size distribution dominated by larger grains. While the overall level of UV attenuation in HD 44179 relative to the attenuation at longer wavelengths is roughly comparable to the RV = 5.5 extinction curve and consistent with earlier

ﬁndings of large dust grains in the RR (Men’shchikov et al., 1998; Jura et al., 1997), the two features that make the HD 44179 attenuation truly unique remain clearly visible: the absence of the 217.5 nm absorption peak and the sharp discontinuous rise in attenuation at λ−1 ≈ 6.0 µm−1. 69

3.5 HD 44179 Attenuation Avg. IS extinction, RV = 3.1 CCM extinction curve, RV = 5.5

2.5

2 V /A λ A 1.5

0.5

0 0 1 2 3 4 5 6 7 λ-1 (µm)-1

Figure 3-8 The UV/optical attenuation curve for HD 44179, normalized to AV = 1 mag. For comparison, two Galactic extinction curves for RV = 3.1 and RV = 5.5 are shown as well. 70

In two subsequent sections we will interpret the features just described in more detail. Given the overall appearance of the attenuation curve for HD 44179, we can conclude, however, that the suggestion by Men’shchikov et al. (2002) that most, if

not all, of the reddening of HD 44179 is caused by interstellar instead of nebular

material intrinsic to the RR is almost certainly incorrect. The hydrogen column

density (N(H I) ≤ 1×1020 cm−2) implied by the equivalent width of the 578 nm diﬀuse

interstellar band in the spectrum of HD 44179 (Hobbs et al., 2004) translates to an

extinction of AV < 0.05 mag. Also noteworthy is the absence of narrow, interstellar

absorption lines Na i and Ca ii in the spectrum of HD 44179 (Hobbs et al., 2004).

The probability that the line-of-sight to one of the most unusual objects in the Galaxy

should also be characterized by one of the most unusual interstellar extinction curves

ever observed is vanishingly small. Given the large internal optical depth of the

RR implied by the visual attenuation found in §§ 3.3.2, we suggest that the unique

characteristics of the attenuation curve of HD 44179 must be explained in terms of

the optical characteristics of the molecular and dust constituents of the RR.

Far-UV Rise

The abrupt and discontinuous rise in the observed attenuation curve beginning

at λ−1 ≈ 6 µm−1, corresponding to a wavelength of 167 nm or a photon energy of

7.4 eV, is not characteristic of dust extinction but rather suggests the onset of strong

absorption, e.g. the absorption associated with the ionization discontinuity in small

PAH molecules, which occurs at just this energy. Verstraete et al. (1990) published

detailed ionization cross-sections for two PAH molecules, pyrene and coronene, which 71 exhibit a sharp onset at λ−1 ≈ 6.0 µm−1, similar to the observed rise in the HD 44179 attenuation curve. When normalized to cross-section per C-atom, the absolute values of the ionization cross-sections and their wavelength dependencies are almost indistin- guishable. A typical value is 1.35×10−18 cm2 C-atom−1 at 143 nm or λ−1 ≈ 7.0 µm−1.

The work by Jochims et al. (1996) shows similar ionization cross-section spectra for naphthalene, azulene, anthracene, phenanthrene, and benz(a)anthracene, all with an sharp onset near 7.5 eV, corresponding to λ−1 ≈ 6.07 µm−1. However, the ionization potential (IP) of PAH molecules is generally size-dependent in a manner similar to the size-dependence of the peak wavelength of ﬂuorescence (Vijh et al., 2004). We demonstrate this dependence of the IP on molecular size by plotting known IP values

(Eilfeld & Schmidt, 1981) against their molecular mass in Figure 3-9. This graph illustrates that the rise in the far-UV attenuation curve is consistent with the ionization of PAHs, provided their molecular mass is not greater than about 250 amu. A similar upper limit of the masses and sizes of PAH molecules was also deduced by Vijh et al.

(2004) on the basis of the observed blue ﬂuorescence seen in the RR. The attenuation curve suggests that absorption by molecules of this size and smaller could contribute to the sudden far-UV rise but that molecules of larger size cannot be present with signiﬁcant abundance close to the central star; if they were, the discontinuity would occur at longer wavelengths. 72

Acenes & Benzologs Perylene, Coronene & Bianthene Series Terrylene & Peropyrene Series Pyrene Series 9 Calulated Values

Red Rectangle

7 Ionisation Potential (eV)

5 I.P. Graphite

0 100 200 300 400 500 600 Molecular Wt. (amu)

Figure 3-9 Ionization potential of diﬀerent families of PAHs as a function of molecular mass. Also shown for reference are the IP of graphite (Clar & Schmidt, 1975; Burns & Yelke, 1969) and the ionization discontinuity apparent in the attenuation curve of the central source of the RR. 73

An important question is whether the observed increase in attenuation in the

RR over the wavelength interval from λ−1 ≈ 6.0 µm−1 to λ−1 ≈ 7.0 µm−1 is in agreement with the likely carbon column density and the measured ionization cross- section at λ−1 ≈ 7.0 µm−1. The observed increase in attenuation over this interval is

2.0 mag , when reduced to the actual color excess of E(B−V ) = 0.35 of HD 44179, corresponding to an optical depth increase by ∆τ = 1.85. With an ionization cross- section σ = 1.35 × 10−18cm2 C − atom−1 at 143 nm wavelength, the corresponding

18 −2 carbon column density is NC = 1.37 × 10 cm for carbon atoms tied up in small

PAH molecules with masses less than 250 amu. If the visual attenuation of 4.2 mag is associated with a mixture of gas and dust with relative elemental abundances as found

21 −2 −1 in the Sun, then N(H i + 2H2)/AV =1.87×10 cm mag (Bohlin et al., 1978), and

21 −2 N(H i + 2H2)=7.85 × 10 cm . With a solar carbon abundance ratio of NC/NH =

−4 3.4 × 10 , this leads to a total expected column density of carbon in all forms of NC

= 2.7 × 1018cm−2. Hence, in order to produce the observed far-UV rise by ionization of small PAHs, we require that about 50% of all carbon atoms reside in small PAHs, if normal solar abundances prevail. Thus, the strength as well as the wavelength dependence of the observed far-UV rise in the attenuation curve are fully consistent with an origin in the ionization of small PAH molecules. In addition, HD 44179 is recognized as a carbon-rich post-AGB star whose relative carbon abundance must exceed that of the Sun by factors of two to six (Cherchneﬀ & Barker, 1992), making the case even stronger. In Figure 3-7 we have indicated the expected contribution of PAH ionization to the far-UV part of the attenuation curve for pyrene as an representative example. The fact that the observed rise is slightly steeper than the 74

individual curve of a speciﬁc PAH molecule can be understood easily, if still smaller

PAH molecules with higher IP values contribute to the absorption at the shorter

wavelengths.

Mid-UV Hump

Given the facts that blue ﬂuorescence from PAH molecules with an intensity com-

parable to that of the ERE contributes strongly to the the nebular emission of the

RR and that the ionization of these molecules in the inner part of the RR contributes

strongly to the far-UV attenuation at λ−1 & 6 µm−1, it appears most likely to us that

the mid-UV hump in the attenuation curve is produced by absorption related to elec-

tronic bound-bound transitions in these same molecules. With a sharp onset of the

ﬂuorescence near 360 nm, the absorption powering this emission must occur at shorter

wavelengths. According to Clar (1972), the absorption spectra of PAH molecules con-

sist of four characteristic bands, the α, para, β, and β0 bands, located at progressively

shorter wavelengths in the mid-UV. They are schematically indicated in the mid-UV

portion of Figure 3-7, with arrows starting at the band edges and the height indicat-

ing the relative strengths of the bands (See also Schmidt, 1977). The β-band which typically occurs at a wavelength 1/1.35 of that of the α band, is typically two orders of magnitude stronger than the α-band, while the intermediate para-band is typically one order of magnitude stronger than the α-band. The wavelength positions of these bands are closely related to the IPs of the corresponding molecules (Clar & Schmidt,

1975, 1977, 1979), so that a continuum of sizes of small PAH molecules with IP >

7.4 eV can be expected to produce a broad, quasi-continuous absorption band as ob- 75 served in the attenuation curve of HD 44179. Verstraete & Leger (1992) produced a generic mid- and far-UV absorption spectrum for interstellar PAH molecules, which exhibits the mid-UV hump and the far-UV rise so clearly apparent in the attenuation curve of HD 44179.

3.4 Discussion

3.4.1 Identiﬁcation of the BL Carrier

When we first attributed the BL in the RR to fluorescence by three- and four- ringed neutral PAH molecules (Vijh et al., 2004), the identification of the likely carrier was tentative for several reasons. Our observational technique for separating the BL spectrum from the underlying scattered light spectrum produces measurements of the

BL intensities only at the positions of the hydrogen Balmer lines, yielding data with very limited and uneven spectral resolution that is incapable of revealing fine spectral features needed for specific identifications of individual molecular species. In addition, the fluorescence spectra of several likely molecular PAH species occupy nearly the same wavelength interval, contributing to the observed spectrum in proportion to their relative abundances and their fluorescence quantum yields. Finally, the gas- phase fluorescence spectra of PAHs vary as a function of temperature, with sub- bands losing their sharpness with increasing temperature (Chi et al., 2001a). For these reasons, the identification of the carrier of the BL can only be made in part on the basis of spectral comparison and must rely largely upon other supporting evidence 76 such as presented in this paper.

There are prospects for improving the spectral resolution of the BL spectrum by using other techniques for separating the BL spectrum from the scattered light spectrum. One such technique is spectropolarimetry. It relies on the fact that the scattered nebular light is highly linearly polarized (Perkins et al., 1981) while the BL is expected to be unpolarized. The presence of the BL in the composite spectrum of the RR would then be revealed by a proportional reduction in the linear polarization of the combined light. The only two published sets of spectropolarimetric observations at optical wavelengths of the RR nebula (Schmidt et al., 1980; Reese & Sitko, 1996) do indeed reveal a sudden reduction of the linear polarization shortward of 410 nm consistent with the onset of the BL band. Unfortunately, the existing data lack the spectral resolution and adequate spectral coverage at shorter wavelengths that would be necessary. A program of spectropolarimetric observations of the RR is being planned.

At this stage, the strongest supporting evidence for the correctness of our iden- tiﬁcation of the carrier of the BL are the close spatial correlation with the 3.3 µm

C–H stretch emission that is highly specific to small, neutral PAH molecules, and the evidence obtained from the analysis of the attenuation curve of the central source that exhibits strong features identifiable with electronic-band absorptions and an ionization discontinuity specifically pointing toward PAH molecules with masses of 250 amu or less. 77

3.4.2 Spectral Variability of the BL

Our initial results showing evidence of systematic spectral variability of the BL with position within the bi-polar structure of the RR are intriguing but clearly need more observations with the highest possible spatial and spectral resolution. Such observations would have the potential of probing the variation in the dominant sizes of the PAH molecules as a function of distance from the central source and as a function of position within and outside of the bi-polar outflow cones and within the shadow of the circum-source disk. Our initial results are consistent with an increase in the size of the most strongly fluorescing PAH molecules with increasing distance from the central source, with a superimposed azimuthal dependence with respect to the outflow axis of the RR. This could reflect the growth of the molecules with increasing distance, but it could also be a result of the different excitation requirements for the fluorescence of PAH molecules of different sizes. The wavelengths of the absorption bands of

PAH molecules are related to their respective ionization potentials (Clar & Schmidt,

1975, 1977, 1979) and occur at shorter wavelengths for smaller PAHs, and at longer wavelengths for larger PAHs. As a consequence, exciting radiation for the smallest

PAHs may be lacking at increased distances as a result of the high opacity of the intervening nebular material.

3.4.3 Spatial Variation of the BL

The close spatial correlation between the BL and the 3.3 µm emission along two nearly identically placed slits, positioned 500 south of HD 44179 is remarkable in it- 78

self by supporting the interpretation that both emissions are produced by the same

molecules. It is further remarkable in that both emissions show a distinctly diﬀerent

distribution than the distribution of the scattered light in the RR. The latter is pro-

duced by dust with a strongly forward-directed scattering phase function, which when

coupled with an embedded source, produces a strongly peaked brightness distribution

of the resulting scattered light. This is clearly apparent when viewing the blue image

of the RR (Fig. 3-1), which is dominated by a bright circular source of about 800 diameter containing most of the scattered light. The PAH ﬂuorescence, both electronic in the form of the BL as well as vibrational in the form of the 3.3 µm emission, by contrast is much less centrally peaked, consistent with an isotropic emission process for both. It is further remarkable that the BL brightness distribution does not reﬂect the crossing of the X-shaped arms or whiskers of the RR, which are so prominent at red wavelengths (Cohen et al., 2004). This indicates that the BL is neither associated

with the interior of the outﬂow cones nor with the walls of the outﬂow cavity but

rather is associated with a population of emitters residing in a shell surrounding the

inner bi-polar structure. This would place them into the shadow of the circum-source

disk and thus outside the reach of ionizing radiation, consistent with our attribution

of the BL to neutral PAHs.

The diﬀerence between the spatial distributions of the BL and the ERE is strik-

ing. It has been known for some time that the ERE is strongly concentrated in the

X-shaped arms of the RR (Schmidt & Witt, 1991), which are thought to be the pro- jections of the walls of the bi-polar outﬂow cavity upon the plane of the sky. Our data

leave little doubt that the two emitters involved require totally diﬀerent excitation 79

conditions. Given the ubiquity of the ERE in a wide range of interstellar environments

(Witt & Vijh, 2004), it is not likely that the concentration of the ERE in the walls of

the RR outﬂow cavity is the result of an abundance enhancement but rather reﬂects

the speciﬁc excitation requirement of the ERE process. As reviewed by Witt & Vijh

(2004), there are several independent indicators suggesting that photon with energies

E > 7.25 eV are required to initiate the ERE. These energies are much higher than

those required for the excitation of ﬂuorescence in neutral PAH molecules and the

number of such high-energy photons are much more limited in the SED of HD 44179.

Thus, they can barely penetrate the walls of the outﬂow cavity before their number

is exhausted, thus limiting the spatial extent of the ERE-producing emitters. The

topic of the ERE excitation will be the subject of another publication by the present

authors.

3.4.4 Attenuation of HD 44179

The attenuation curve of HD 44179 appears highly unusual when compared to

the wavelength dependence of interstellar extinction, but it appears to be fully con-

sistent both qualitatively and quantitatively with expectations, if the scattered light

of HD 44179 has to penetrate a large column density (AV = 4.2 mag) of circumstellar

material ﬁlled with neutral PAH molecules with atomic masses less than 250 a.m.u.

With a column density of carbon in the form of small PAH molecules estimated from the strength of the far-UV absorption related to PAH ionization, we can estimate a molecular abundance of such PAH molecules with typically 16 carbon atoms 80 per molecule relative to hydrogen in the RR of about 10−5. This agrees well with the canonical abundances of PAH molecules derived from the analysis of the strength of the mid-IR aromatic emission features, commonly referred to as the unidentiﬁed infrared emission bands (Leger & d’Hendecourt, 1987). This makes PAH molecules the most abundant species of interstellar molecules after H2 and CO. The fact that despite the strong indications for a presence of small PAH molecules there is no clear counterpart for the interstellar 217.5 nm absorption feature in the attenuation curve of HD 44179, suggests that proposals to associate the 217.5 nm feature with PAH absorption in current dust models should be viewed with caution. The current observations appear to prohibit at least a role for small PAHs with masses < 250 amu.

This still leaves the possibility that larger PAH structures could be involved in contributing to the interstellar 217.5 nm feature, because they would have the advantage of greater thermodynamic stability under interstellar conditions. The small PAH molecules apparent in the RR are able to form and survive in the relatively benign radiation environment of the RR, especially in the shadow of the optically thick dust ring blocking direct radiation from the central source. It is in these parts of the

RR nebula where the BL intensity is strongly enhanced in comparison to the scattered radiation and the ERE. These small PAH molecules, however, are not likely to survive the harsher radiation conditions prevailing in the diﬀuse interstellar medium

(Jochims et al., 1994, 1999), where they would be subject to photo-dissociation. As a consequence, it does not appear likely that BL with spectral characteristics and intensities similar to those observed in the RR will be seen in dusty star-forming regions such as reﬂection nebulae and H ii regions (Rush & Witt, 1975). 81

The conditions deduced for the RR environment also places interesting constraints on the carriers of the diﬀuse interstellar bands (DIBs). Current ideas for the carriers of these much-studied but still unidentiﬁed interstellar absorption features concentrate on carbonaceous molecules, their ions in particular (see Snow, 1995, for a review).

PAH ions have been considered as the most likely carriers by a number of investiga-

tors. Given the strong indicators for a large column density of PAH molecules and

their ions in the RR, it is remarkable that only one DIB λ5780 is detected in the spectrum of HD 44179 (Hobbs et al., 2004). T. P. Snow (2004; private communica-

tion) has proposed that this fact might be explained, if DIBs are produced by PAH

cations with masses well in excess of 250 a.m.u. Another fact that must be taken into

consideration is the absence of absorptions resulting from the D2 ←D0 transition of

anthracene+ (708.76±0.13 nm), naphthalene+ (670.73±0.06 nm) (Sukhorukov, 2004),

and acenaphthene+ (646.3 nm)(Biennier et al., 2003) which might be expected in the

spectrum of HD 44179, if these cations are present in the light path. Some of these

discrepancies could be resolved by considering the ionization of these PAHs to the

dication stage.

3.5 Conclusions

We summarize our conclusions:

1. Our spatially resolved observations of the BL in the RR have yielded the ﬁrst

tentative evidence that the detailed spectrum of the BL is spatially variable,

consistent with a change in the dominant size of the ﬂuorescing molecules with 82

position relative to the central source.

2. The spatial distribution of the band-integrated BL in the RR diﬀers fundamen-

tally from those of the dust-scattered radiation and of the ERE, consistent with

an origin from a nebular volume located mainly outside the outﬂow cavity and

concentrated in the shadow of the disk obscuring the central source. This sup-

ports our earlier conclusion that the BL is produced by ﬂuorescence from small,

neutral PAH molecules.

3. There exists an excellent spatial correlation between the BL and the distri-

bution of 3.3 µm C–H stretch emission from PAHs. The latter emission is

predominantly produced by small, neutral PAH molecules, and the close spatial

correlation supports the suggestion that the BL has its origin in the similar

molecules.

4. By combining our observations with existing IUE data for HD 44179, we have

determined the UV/optical attenuation curve for the central source in the RR.

The attenuation curve is characterized by an exceptionally high value of RV =

AV/E(B−V ) = 11.9, a remarkable absence of the 217.5 nm absorption bump,

the latter being a characteristic of Galactic interstellar extinction, and a strong,

discontinuous rise in attenuation near λ−1 ≈ 6.0 µm−1. These unusual charac-

teristics suggest strongly that the reddening observed in the light from HD 44179

is a result of radiative transfer in the RR and not due to interstellar extinction.

5. The far-UV rise in the attenuation curve of HD 44179 is qualitatively and quanti- 83

tatively diﬀerent from the much more gradual far-UV rise in Galactic extinction

curves. The rise is, however, consistent with the onset of photo-ionization of

small PAH molecules with masses of less than 250 amu, if the abundance of

such PAH molecules relative to hydrogen is about 10−5. This represents a fully

independent conﬁrmation of the presence of small PAH molecules in the RR

environment, consistent with our conclusion regarding the origin of the BL in

the RR.

6. The broad mid-UV hump in the attenuation curve of HD 44179 whose appear-

ance is quite unlike that of the familiar 217.5 nm interstellar absorption band,

can be explained satisfactorily as resulting from a superposition of the (typi-

cally) four mid-UV absorption bands characteristic of neutral PAH molecules.

An observable consequence of this absorption and the resulting electronic exci-

tation is de-excitation via ﬂuorescence, which we indeed observe.

7. We do not expect that BL with comparable relative intensities as seen in the RR

will be seen in star-forming regions such as reﬂection nebulae and H ii regions,

where the mid-IR UIR band emissions are otherwise strong. The small PAH

molecules required by the BL are not expected to survive in the radiation ﬁelds

typical of star-forming regions.

Note added in proof in the Journal: Recent observations of other reﬂection nebulae have produced evidence for the presence of BL in some instances, although at lower relative intensities compared to the RR (see next Chapter for details). Chapter 4

Blue Luminescence and the

Presence of Small Neutral PAHs in the ISM 1

Blue Luminescence (BL) was ﬁrst discovered in a proto-planetary nebula, the Red

Rectangle (RR) surrounding the post-AGB star HD 44179. BL has been attributed to ﬂuorescence by small, 3-4 ringed neutral polycyclic aromatic hydrocarbon (PAH) molecules, and was thought to be unique to the RR environment, where such small molecules are actively being produced and shielded from the harsh interstellar radiation by a dense circumstellar disk. In this paper we present the BL spectrum detected in several ordinary reﬂection nebulae illuminated by stars having temperatures between 10,000 – 23,000 K. All these nebulae are known to also exhibit the infrared

1To appear in The Astrophysical Journal November 2005, Vol 633, Authors: U. P. Vijh, A. N.

Witt, & K. D. Gordon

84 85

emission features called aromatic emission features (AEFs) attributed to large PAHs.

We present the spatial distribution of the BL in these nebulae. In the case of Ced 112,

the BL is spatially correlated with mid-IR emission structures attributed to AEFs.

These observations provide evidence for grain processing and possibly for in-situ for-

mation of small grains and large molecules from larger aggregates. Most importantly,

the detection of BL in these ordinary reﬂection nebulae suggests that the BL car-

rier is an ubiquitous component of the ISM and is not restricted to the particular

environment of the RR.

4.1 Introduction

We discovered the presence of small, 3-4 ringed, neutral polycyclic aromatic hydrocarbons (PAHs) in the Red Rectangle (RR) through their electronic ﬂuorescence, called Blue Luminescence (BL) (λpeak ∼ 3750 A)˚ (Vijh et al., 2004, 2005b). The

observed spectrum was attributed to small PAHs on the basis of a comparison with

laboratory spectra of PAH molecules ﬂuorescing in the gas phase (Vijh et al., 2004),

the spatial correlation between the BL and the 3.3 µm emission, commonly attributed to small, neutral PAH molecules, and the newly-derived UV/optical attenuation curve for the central source of the RR, HD 44179 (Vijh et al., 2005b). The latter two results

provided strong independent evidence for the presence of small PAH molecules with

masses of less than 250 amu in the RR, which supports the attribution of the blue lu-

minescence to ﬂuorescence by the same molecules. Such ﬂuorescence is excited by the

absorption of near-ultraviolet photons, followed by the relaxation to the lowest level 86

of the ﬁrst excited electronic state via a series of mid-IR vibrational transitions. The

subsequent electronic transitions to various vibrational levels of the electronic ground

state are then observed as optical ﬂuorescence, i.e. the BL. The BL is much more

speciﬁc in terms of size and ionization state of the emitter than vibrational transitions

observed in the mid-IR (Reyl´e& Br´echignac, 2000). In particular, the wavelength of

the most energetic electronic transition seen in neutral PAHs is closely dependent on

the size of the molecular species. In general, this ﬂuorescence wavelength increases

with the molecular weight and size of the molecule (Vijh et al., 2004).

Other emission bands that are most commonly attributed to PAHs are the family

of features at 3.3, 6.2, 7.7, 8.6, 11.3, & 12.7 µm (referred to as aromatic emission features (AEF) or unidentiﬁed infra-red (UIR) bands) (Hony et al., 2001; Cook & Saykally,

1998; Leger & Puget, 1984; Verstraete et al., 2001; Allamandola et al., 1985; Sellgren,

1984). These AEFs are the signatures of aromatic C-C and C-H fundamental vibra-

tional and bending modes and the speciﬁc sizes and ionization states of the PAH

molecules or ions cannot be ascertained from these spectra alone. On absorption of

a far-UV photon, a PAH molecule or ion usually undergoes a transition to an upper

electronic state. If the molecule or ion undergoes iso-energetic transitions to highly vi-

brationally excited levels of the ground state, then the molecule or ion relaxes entirely

through a series of IR emissions in the C-C and C-H vibrational and bending modes.

These transitions are largely independent of size, structure and ionization state of

the molecule and are identiﬁed with the AEFs. The AEFs are found in almost all

astrophysical environments including the diﬀuse interstellar medium (ISM), the edges

of molecular clouds, reﬂection nebulae, young stellar objects, HII regions, star form- 87

ing regions, some C-rich Wolf-Rayet stars, post-AGB stars, planetary nebulae, novae,

normal galaxies, starburst galaxies, most ultra-luminous infra-red galaxies and AGNs

(see Peeters et al., 2004, and references therein). The presence of the AEF carriers

in such a wide range of environments suggests that they must include a wide range

of sizes of aromatic structures in several ionization states. Therefore, an observation

of BL in astronomical sources in the UV/visible range is of particular value because

it oﬀers the possibility of tracing the speciﬁc presence of small, most likely neutral

PAHs in the ISM.

The central star in the RR (where the BL was ﬁrst discovered), HD44179, is a post-AGB star that is in an active dust-producing stage. Small dust grains and

PAH molecules in a wide range of sizes are currently being produced in its circumstellar outﬂow environment. The discovery of BL in such an environment was therefore not a surprise. However, several detailed investigations (Jochims et al., 1999;

Le Page et al., 2003, and references therein)(see also Allain et al., 1996a,b) have led

to the conclusion that PAH molecules with fewer than 30 carbon atoms, i.e. the ones

detected in the RR, do not survive in the interstellar environment, once they are

ejected from their place of formation. Photo-dissociation by the interstellar radiation

ﬁeld is expected to limit their life time severely. However, as mentioned earlier, the

AEFs are seen in almost all astrophysical environments including those with harsh

interstellar radiation ﬁelds. Do the small PAHs that produce the BL survive under

interstellar conditions or are the AEFs primarily due to larger PAHs? The presence

or absence of BL in environments with diﬀering levels of UV radiation can help an-

swer these questions. Presence of BL preferentially traces the small, neutral PAH 88

molecules, as the larger molecules, even if they were neutral, have generally lower

fluorescence efficiencies, and ionized PAH molecules do not fluoresce at optical wave-

lengths. Small, partially dehydrogenated PAH radicals or radical ions having a singlet

ground state, however, could contribute to the BL as well (W.W. Duley, Private com-

munication). To test the viability of such small PAH structures in the general ISM far

from C-rich stellar outﬂows, we obtained long-slit spectra of several reﬂection nebulae

and star-forming regions where AEFs and other emission features like the extended

red emission (ERE) have been detected. The temperatures of the illuminating stars

in these regions cover a range of 8250 – 23000 K, and thus oﬀer a chance to study

the distribution of small PAH molecules as a function of environment.

4.2 Observations

The objects observed for this investigation are listed in Table 7.1. The observations

of the nebula Ced 201 were made at the Steward Observatory 2.3-m Bok telescope

on 2002 November 5 and those for NGC 2023, NGC 5367 and Ced 112 were made at

the Cerro Tololo Inter-American Observatory (CTIO) 1.5-m telescope on 2003 March

26 and 28. At the Bok telescope we used the Boller and Chivens spectrograph with

a grating having 300 lines mm−1, blazed at 376 nm in the ﬁrst order with a L3800

cut-on ﬁlter. The detector used was a 1200 × 800 CCD (ccd20) yielding a spatial

scale of 000.8 pixel −1. The slit used was 4 0 long and 400.5 wide. At the CTIO 1.5-m we used the R-C spectrograph with the grating #09 having 300 lines mm−1, blazed

at 400 nm in the ﬁrst order with a CuSO4 ﬁlter. The detector used was the Loral 89

IK CCD and the spectral coverage was from 340 nm to 587.5 nm. The slit used was

7 0 long and 200.5 wide and yielded a spatial scale of 100.3 pixel −1. A coronographic

decker assembly was used to minimize scattered light from the star while probing

as much of the inner nebula as possible at CTIO. For all observations the nebular

exposures were bracketed by exposures of the illuminating stars and the long-slit

enabled simultaneous exposures of the nebula and the sky. Table 7.1 gives the details

of exposure times and slit positions. Data reductions were carried out with IRAF

2.12 EXPORT, and all spectra were ﬂux-calibrated via observation of standard stars.

Table 4.1 Details of objects observed, their illuminating stars and slits used.

Object Observatory IlluminatingStar SlitPositions(Oﬀset) Exposures (s)

Name Spectral Type Temperature (K)

00 00 Ced 201 RN Bok 2.3 m BD +69 1231 B9.5 V 10,000 15 north, 20 south 4×1800 00 00 00 Ced 112 RN CTIO 1.5 m HD 97300 B9 V 10,500 0 ,6 north, 6 south 2× 1800 00 Ced 112 RN CTIO 1.5 m HD 97300 B9 V 10,500 10 north 3× 1800 00 NGC 5367 RN CTIO 1.5 m Her 4636 Herbig Ae/Be 18,700 5 north 2×1800 00 NGC 2023 RN, PDR CTIO 1.5 m HD 37903 B1.5 V 23,000 0 2×1800

4.3 Results

4.3.1 Ced 201

Ced 201 is a rather compact object at a distance of 420 pc (Casey, 1991), on the

edge of a molecular cloud. It is excited by the B9.5V star BD +69 1231. Witt et al.

(1987) note that the radial velocity of this star diﬀers from that of the molecular cloud

by 11.7±3.0 km s−1 so that Ced 201 is probably the result of an accidental encounter of the star with the molecular cloud, while for most other reﬂection nebulae the exciting star was born in situ. An arc-like structure located between the star and 90

the denser parts of the cloud to the north and the north-east at about 1800 from the

star may represent a shock due to the supersonic motion of the star. The mid-IR

ISO spectra of Ced 201 give evidence for transformation of very small carbonaceous

grains into the carriers of the AEFs, due to the radiation ﬁeld of the illuminating

star and/or to shock waves created by its motion (Cesarsky et al., 2000). Ced 201

exhibits ERE (Witt & Boroson, 1990) and AEFs (Cesarsky et al., 2000).

Figure 4-1 DSS.J.POSSII image of Ced 201, overlaid with our slits showing the extracted apertures. The white central pixel indicates the position of the star BD +69 1231. The dash at the bottom left is 1000. The right-most extraction window on the northern slit captures the spectrum of the arc-like structure seen on this image.

We obtained nebular spectra along two slits (PA 90◦), 1500 north and 2000 south of the central star. Figure 4-1 shows an optical image of the nebula overlaid with the 91

extracted apertures. We extracted 9 spectra along the 1500 north slit and 6 spectra

along the 2000 south slit each covering 68.4 sq. arcsec of the nebula. BL was measured

using the line-depth technique (Vijh et al., 2004) at each of the six hydrogen Balmer

lines from Hβ to Hη in each of these apertures. The line-depth technique detects the BL intensity at the wavelength positions of strong absorption lines by virtue of the fact that the presence of BL reduces the line depths in the dust-scattered nebular spectrum compared to the depths of the identical lines in the spectrum of the illuminating star. Figure 4-2 shows representative BL spectra in four such apertures

and the ratio of BL to the continuum in the same apertures. The use of a cut-

on ﬁlter at 3800 A˚ restricted the measurement of the last Balmer line measured to

3835 A.˚ Longward of 3800 A˚ these spectra are similar to the BL spectra seen in the

RR (Vijh et al., 2004, 2005b) and other nebulae discussed in further sections of this

paper. The spectra from different offsets do differ, as was also noted by Vijh et al.

(2005b) and can probably be attributed to diﬀering radiation environments. Given

the poor spectral resolution of BL spectra measured using the line-depth technique,

in this paper, we will not characterize the diﬀerences in the spectra but will simply

refer to them as BL spectra. 92

2.1e−05 20"S, 15.2"E 20"S, 0" 1.8e−05 15"N, 7.6"E 15"N, 7.6"W )

−1 1.5e−05 sr −1

Å 1.2e−05 −1 s −2 9.0e−06

(erg cm 6.0e−06 BL I

3.0e−06

0.0e+00 20"S, 15.2"E 0.32 20"S, 0" 15"N, 7.6"E 15"N, 7.6"W 0.28

0.24

0.20 sc. /I BL

I 0.16

0.12

0.08

0.04

0.00 3400 3600 3800 4000 4200 4400 4600 4800 5000 wavelength (Å)

Figure 4-2 BL spectra (upper panel) and ratio of BL to the continuum (lower panel) at 4 diﬀerent positions in Ced 201. 93

4.3.2 Ced 112

Ced 112 is a prominent reﬂection nebula in the cloud Chamelion I at a distance

of ∼ 160 pc (Whittet et al., 1997). It is illuminated by the pre-main sequence Herbig

AeBe star HD 97300, spectral type B9 (Rydgren, 1980) with a luminosity of ∼ 35 L

(van den Ancker et al., 1997). Its pre-main sequence evolutionary status as a Herbig

AeBe star is based on indirect clues, namely the association with a reﬂection nebu-

losity, the presence of an infrared excess at λ> 5 µm, and its location on the ZAMS in the HR diagram (Whittet et al., 1997). There is no evidence of signiﬁcant infrared

excess at wavelengths shorter than 5 µm nor of Hα emission. HD 97300 is very likely a relatively old object among HAeBe stars, representative of the latest stages of the pre-main sequence evolution. There is general agreement that the star is an embedded member of the young stellar population and is thus situated at a distance coincident with that of the cloud itself. Infrared emission from the nebula Ced 112 is dominated by AEFs centered at 6.2, 7.7, 8.6, 11.3, and 12.7 µm with very small contributions from continuum emission at longer wavelengths (Siebenmorgen et al.,

1998). Siebenmorgen et al. (1998) also report the detection of an extended, ring-like

structure around HD 97300, whose emission is dominated by the AEFs. 94

Figure 4-3 AAO.R.DSS2.038 image of Ced 112, overlaid with our slits showing the extracted apertures. The central, white pixel indicates the position of the star HD 97300. The dash at the bottom left is 1000.

We obtained nebular spectra along four slits (PA 90◦), 000 oﬀset, 600 north, 600 south and 1000 north of the central star. Figure 4-3 is an optical image of the nebula overlaid with the extracted apertures. We extracted 10 spectra along the 600 north and

600 south slits and 6 spectra along the slit through the center (000 oﬀset) each covering

16.25 sq. arcsec of the nebula. Along the 1000 north slit we extracted 4 spectra each

covering 32.5 sq. arcsec. BL was measured using the line-depth technique in each of

these spectra. Along the 600 south slit we extracted 14 spectra with ﬁner apertures 95 each covering 8.125 sq. arcsec to enable comparison with ISOCAM data discussed in the next section. Figure 4-4 show representative BL spectra in three such apertures and the ratio of BL to scattered light in the same apertures. 96

2.4e−05 6"S, 21.125"E 6"S, 8.125"W 2.1e−05 6"S, 14.625"W )

−1 1.8e−05 sr −1 1.5e−05 Å −1 s

−2 1.2e−05

9.0e−06 (erg cm BL I 6.0e−06

3.0e−06

0.0e+00 6"S, 21.125"E 6"S, 8.125"W 6"S, 14.625"W 1.00

0.80 sc.

/I 0.60 BL I

0.40

0.20

0.00 3400 3600 3800 4000 4200 4400 4600 4800 5000 wavelength (Å)

Figure 4-4 BL spectra (upper panel) and ratio of BL to the continuum (lower panel) at 4 diﬀerent positions in Ced 112. These spectra correspond to extractions along the slit shown in Figure 4-5. 97

Spatial Correlation of BL with mid-IR emission

As mentioned earlier, Siebenmorgen et al. (1998) reported the detection of an

extended, ring-like structure around HD 97300 in the mid-IR images of the nebula

taken with ISOCAM. They obtained images in four narrow-band ﬁlters centered at

6.0 µm (lw4), 6.8 µm (lw5), 11.3 µm (lw8) and 14.9 µm (lw9). All four images show an extended emission centered on the star and an elliptical ring of size 5000 × 3600 around it. They also report that the spectra taken with the circular variable ﬁlter between 5.8 and 13.8 µm roughly aligned along a line intersecting the star, the emission minimum, and the ring seen in the SE direction (PA = 142.4◦) show strong AEF features. Our

600 south oﬀset slit extends over the nebula and the ring. Figure 4-5 shows the lw9

image overlaid with our slit, with the extraction apertures marked and Figure 4-6

shows a plot of the distribution of the band-integrated BL and dust-scattered light

in the apertures along the slit. The integrated BL clearly shows enhancements in the

regions where the slit crosses the ring on both sides, while the scattered light does

not. It is also interesting to note that in the spectra shown by Siebenmorgen et al.

(1998) the ﬂux at 7.8 µm between two positions on the ring and outside away from the star varies by a factor of ∼ 2.4 while the band integrated BL intensity varies by a factor of 2.6 for similar locations. Figure 4-7 shows the distribution of the ratio of

the band-integrated BL to the band-integrated scattered light. The enhancement of

the BL at the positions of the ring is clearly seen in this representation as well. 98

Figure 4-5 ISOCAM image (lw9) of Ced 112 from Siebenmorgen et al. (1998) overlaid with our slit showing the extraction apertures. The white pixel is overlaid to indicate the position of the central star HD 97300. The arrows indicate where the slit crosses the ring-like structure seen in the image. 99

0.1 Band-BL-6" S Band-Scat-6" S

0.01 Intensities

0.001

0.0001 -30 -20 -10 0 10 20 30 offset (arcsec)

Figure 4-6 Band-integrated BL and scattered light intensities in the apertures indicated in Figure 4-5. A central-blocker used in the slit prevented measurements close to zero offset. Offsets are along the slit, with zero offset being exactly 600 south of the star. 100

100 sc

/Band-I 10-1 BL Band-I

10-2 -30 -20 -10 0 10 20 30 offset (arcsec)

Figure 4-7 Ratio of Band-IBL to Band-Isc in the apertures indicated in Figure 4-5. Oﬀsets are along the east-west, 600 south slit with zero being exactly south of the star.

4.3.3 NGC 5367

NGC 5367 surrounds the double star Her 4636 in the head of the prominent

Cometary Globule CG 12 and lies at high galactic latitude (b = +21◦). It is an exam-

ple of a relatively isolated low-to-intermediate-mass star formation region (van Hill et al.,

1975), which is thought to have formed as a result of a nearby supernova event 10 –

20 million years ago (Williams et al., 1977), and has converted about 20% of its gas

mass into stars (White, 1993). The most probable distance to NGC 5367 is 630 pc,

and the mass of the surrounding molecular cloud is ∼ 120 M (Williams et al., 1977).

Williams et al. (1977), using optical and infrared photometry, showed the likeliest 101

masses of the two stars in the binary system to be 4.5 and 8 M , with spectral type

B7 and B4, respectively, with the latter being surrounded by a dust-shell, emitting as a 1600 K blackbody. Reipurth & Zinnecker (1993) also obtained spectra of the two

stars (called N and S for north and south) as part of their visual pre-main sequence

binary program. They show that the northern component has strong Hα emission and the southern component, the B7 star, has no such emission, and they conclude that both are probably Herbig Ae/Be stars. Vieira et al. (2003) also include Her 4636

in their study of Herbig Ae/Be stars and estimate a rotational velocity of 158 km s−1 and an eﬀective temperature of ∼ 18,700 K. 102

Figure 4-8 AAO.R.DSS2.325 image of NGC 5367, overlaid with our slits showing the extracted apertures. The central, white pixel indicates the position of the star Her 4636. The dash at the bottom left is 2000.

We obtained nebular spectra 500 north of the illuminating stars. Figure 4-8 shows an optical image of the nebula overlaid with the extracted apertures. We extracted 8 spectra each covering 16.25 sq. arcsec of the nebula along the slit. BL was measured using the line-depth technique in each of these apertures. A combined spectrum of both the stars was used for the stellar measurements. Figure 4-9 show representative

BL spectra in three such apertures and the ratio of BL to scattered light in the same apertures. The red rise in the spectra is a result of emission from hydrogen recom- 103

bination ﬁlling in the hydrogen Balmer lines, in addition to the BL. Characteristic

hydrogen recombination spectra alone would result in a spectral intensity steadily

decreasing with decreasing wavelength. Not only do these spectra show a peak at

λ ∼ 370 nm they also show a secondary peak at λ ∼ 400 nm, which is a feature seen in many BL spectra. 104

5.0e−06 5"N, 24.3"E 5"N, 28.6"W 5"N, 41.6"W 4.0e−06 ) −1 sr −1 3.0e−06 Å −1 s −2 2.0e−06 (erg cm BL I 1.0e−06

0.0e+00 5"N, 24.3"E 5"N, 28.6"W 1.00 5"N, 41.6"W

0.80

sc. 0.60 /I BL I

0.40

0.20

0.00 3400 3600 3800 4000 4200 4400 4600 4800 5000 wavelength (Å)

Figure 4-9 BL spectra (upper panel) and ratio of BL to the continuum (lower panel) at three diﬀerent positions in NGC 5367. These spectra are contaminated with emission from hydrogen recombination, but the peak ∼ 400 nm, ∼ 377 nm and the rise towards the Balmer discontinuity at λ= 357 nm are indicative of BL. 105

4.3.4 NGC 2023

NGC 2023 is a bright reﬂection nebula which is a part of the Orion L1630 molecular cloud at a distance of 450 to 500 pc (de Boer, 1983). This nebula, also known as

vdB 52, is illuminated by the bright B1.5 V star HD 37903 (Racine, 1968). NGC 2023

is an example of photo-dissociation regions (PDRs) that exist on the surfaces of molec-

ular clouds in the vicinity of hot, young stars. The interaction of the newborn stars

with their parental molecular environment leads to intense emission of ﬁne-structure

lines of carbon and oxygen and of H2 ro-vibrational transitions. NGC 2023 has been well studied at many wavelengths, including the mid- and near IR region showing

AEFs attributed to PAHs (Sellgren et al., 1985; Verstraete et al., 2001; Moutou et al.,

1999). NGC 2023 also exhibits ERE (Gorodetskii & Rozhkovskii, 1978; Witt et al.,

1984b; Witt & Schild, 1988). 106

Figure 4-10 SERC.ER.DSS2.768 image of NGC 2023, overlaid with our slits showing the extracted apertures. The central, white pixel indicates the position of the central star HD 37903. The dash at the bottom left is 2000.

We obtained nebular spectra with a slit (PA 90◦)at000 oﬀset with a central decker blocking out the star. Figure 4-10 shows an optical image of the nebula overlaid

with the extracted apertures. We also obtained spectra of the star HD 37903. We

extracted 5 spectra each covering 32.5 sq. arcsec of the nebula along the slit. BL was

measured using the line-depth technique in each of these apertures. Figure 4-11 show representative BL spectra in three such apertures and the ratio of BL to scattered light in the same apertures. As in NGC 5367, the red rise in the spectra is a result 107 of hydrogen recombination ﬁlling in the hydrogen Balmer lines in addition to the

BL. In this case the origin of this recombination emission is the rather extended region of photo-ionized hydrogen gas overlying the Orion constellation. There is no measurement of BL at Hβ because the nebular Hβ (λ ∼ 4861 A)˚ emission more than

ﬁlls the scattered stellar Hβ absorption line. Characteristic recombination spectra alone would result in a spectrum steadily decreasing with decreasing wavelength.

These spectra show a upturn at λ ∼ 400 nm and a pronounced peak at λ ∼ 380 nm, which we interpret as evidence for the presence of BL. 108

2.0e−05 0", 26"E 0", 13"E 0", 13"W 1.6e−05 ) −1 sr −1 1.2e−05 Å −1 s −2 8.0e−06 (erg cm BL I 4.0e−06

0.0e+00 0", 26"E 0", 13"E 1.00 0", 13"W

0.80

sc. 0.60 /I BL I

0.40

0.20

0.00 3400 3600 3800 4000 4200 4400 4600 4800 5000 wavelength (Å)

Figure 4-11 BL spectra (upper panel) and ratio of BL to the continuum (lower panel) at three diﬀerent positions in NGC 2023. These spectra are heavily contaminated with emission from hydrogen recombination, but the peak ∼ 380 nm and the rise towards the Balmer discontinuity at λ = 357 nm are indicative of BL. 109

4.4 Discussion

We have presented optical long-slit spectra of reflection nebulae illuminated by stars having a range of temperatures between 10,000 – 23,000 K, and thus covering a wide range of radiation environments. The long-slit technique enables us to extract many individual spectra in each nebula that differ in local radiation density and probably size-distribution of emission carriers but do not differ considerably in chemical composition or grain-processing histories. Comparison of BL spectra from different nebulae will help us study the evolution of the carriers with changing environments and also see what effects different grain processing have on the BL emitters.

4.4.1 The BL Spectrum

The BL spectra in different apertures along the slit in the same nebula varied in intensity with offset from the respective stars but were fairly similar in spectral shape, and therefore we produced one average spectrum for each of the slits. This allows us to identify differences between the BL spectrum in different nebulae and

ﬁnd possible correlations to varying environments and local ISM histories. Figure 4-12

shows the average BL spectrum from the two slits in Ced 201, four slits in Ced 112

and two BL spectra from the RR (reported earlier in Vijh et al. (2004)). In this

section we only discuss Ced 201 and Ced 112. The other two nebulae NGC 2023 and

NGC 5367 will be discussed in a later section (§ 4.4.3). The RR BL spectrum from

the slit 200.5 south oﬀset has been scaled down by a factor of 5 to aid comparison.

Four differences in these spectra may be noted. One, the BL intensity in the RR-200.5 110 south is about 10 times higher than in Ced 112 but the RR-500 south is similar in intensity. Two, the RR spectra have peaks at λ ∼ 377 nm similar to Ced 112. Three, the RR spectra have a long-wavelength component that is absent in Ced 112, but present in Ced 201. And finally, the Ced 201 BL spectrum is unlike that of Ced 112, but could be similar to that of the RR, given the short wavelength rise, but lack of data short-ward of 380 nm makes it difficult to classify the spectrum. All the spectra also show a secondary peak at λ ∼ 400 nm, which could be an indication of a population of efficient fluorescing molecules of slightly larger sizes. The nebular spectra in Ced 201 show some interesting band-like emission features, which are not present in the stellar spectrum and therefore must be nebular in origin. Figure 4-13 shows the nebular spectrum at 1500 north, 3800 west divided by the stellar spectrum.

This resulting spectrum is free of any stellar features and the peaks at the positions of the Balmer lines indicate ﬁlling in of the nebular lines by the BL. The two broad bands at λ ∼ 4050 A˚ and λ ∼ 4500 A˚ are new features and are likely to be signatures of slightly larger PAHs (> 14 C atoms) than the 3-4 ringed ones that produce the blue peak at λ ∼ 377 nm. 111

1.4e−05 RR − 2.5"S offset (× 0.2) RR − 5"S offset Ced 112 − 0" offset 1.2e−05 Ced 112 − 6"N offset Ced 112 − 6"S offset Ced 112 − 10"N offset ) 1e−05 Ced 201 − 15"N offset −1 Ced 201 − 20"S offset sr −1

Å 8e−06 −1 s −2 6e−06 (erg cm BL

I 4e−06

2e−06

0 3400 3600 3800 4000 4200 4400 4600 4800 5000 wavelength (Å)

Figure 4-12 Average BL spectra from the four slits in Ced 112, and two slits in Ced 201 compared to two spectra from the RR. 112

0.01 Ced 201 − 15"N, 38"W 0.009

0.008

0.007

0.006 *

/F 0.005 neb F 0.004

0.003

0.002

0.001

0 3400 3600 3800 4000 4200 4400 4600 4800 5000 wavelength (Å)

Figure 4-13 Ratio of the nebular spectrum to the stellar spectrum in Ced 201 from a position 1500 north and 3800 west of the central star. The two broad bands at λ ∼ 4050 Aand˚ λ ∼ 4500 A˚ are new features and are likely to be signatures of slightly larger PAHs (> 14 C atoms).

Kemper et al. (1999) model the far-IR and sub-millimeter observations of Ced 201 and conclude that heating by the ionization of PAHs (carbonaceous particles with linear size > 15 A)˚ is required to explain the observed properties of Ced 201. Creation of

PAHs that produce the observed mid-IR features by processing of small carbonaceous grains is also invoked in this nebula (Cesarsky et al., 2000). Cesarsky et al. (2000) suggest that some of carbonaceous material that emits the mid-IR continuum and broad bands far from the star is processed through the eﬀect of the star that moves through the molecular cloud, producing AEF carriers: the continuum is only 2 times 113 fainter 1200 from the star than close to the star, demonstrating the partial disappear- ance of its carriers near the star while the AEFs become very strong close to the star.

They also infer that these carbonaceous grains must be very small (radius of the order of 1 nm) since they are heated transiently by visible photons to the temperatures of

' 250 K necessary to emit the observed mid-IR features. It is also interesting to note that Witt et al. (1987) ﬁnd from UV/visible scattering studies of Ced 201 that the grains responsible for the visible scattering have a narrow size distribution skewed towards big, wavelength - sized grains. This is mostly seen within about 2000 from the star. On the other hand the exceptionally strong 2175 A˚ extinction band observed in the spectrum of BD +69 1231, the illuminating star, suggests the presence of many smaller carbonaceous grains. All this is evidence for grain processing by the radiation of the star, or/and by shattering of grains by a shock wave associated with its motion.

The banded spectrum (Figure 4-13) is of the arc-like structure seen toward the NW in the DSS image of the nebula shown in Figure 4-1. This arc is probably a result of a shock and we suggest that the BL carriers are produced in situ from the larger

AEF carriers being photo-dissociated in the stellar UV-ﬁeld.

4.4.2 Spatial Distribution of BL

We observed little variation in spectral shape in the BL spectra extracted from diﬀerent apertures along a given slit, in all the nebulae, but considerable variation in the intensity of both the BL and the scattered light. We plot band-integrated BL and scattered light intensities as a function of oﬀset from the central illuminating star 114

for each slit position. We integrate from Hβ to the Balmer discontinuity (486 nm –

357 nm) for both the BL and the scattered light in case of Ced 112 and till 383.5 nm

at the lower wavelength limit for Ced 201 because of the wavelength ﬁlter that was

used for those observations. Figure 4-14 shows band-integrated BL and scattered

light intensities as a function of oﬀset from the central star along the six slits in the

two nebulae. The distribution of BL in Ced 112 seems to be ﬂatter than that in

Ced 201, but the overall BL-intensities in Ced 201 are about ﬁve times higher than

those in Ced 112. Compared to the maximum intensity of the BL in the RR (200.5

south slit) the intensities in Ced 112 and Ced 201 are 20 times and 4 times lower

respectively. 115

100 Ced201 15"N − BL Ced201 20"S − BL 15"N − sc 20"S − sc 10−1

10−2

10−3

10−4

10−5 −80 −40 0 40 80 −80 −40 0 40 80 Ced112 0" − BL Ced112 6"S − BL

) 0" − sc 6"S − sc −1 −1 10 sr −1 s −2 −2 10

10−3

10−4 Intensities (erg cm

10−5 Ced112 6"N − BL Ced112 10"N − BL 6"N − sc 10"N − sc 10−1

10−2

10−3

10−4

10−5 −30 −10 10 30 −30 −10 10 30 offset (arcsec)

Figure 4-14 Band-integrated BL and scattered light intensities along the various slits in Ced 201 (upper panel) and Ced 112 (middle and lower panel). 116

0.40 Ced201 15"N Ced201 20"S 0.35

0.30

0.25

0.20

0.15

0.10 Sc

0.05 /Band−I

BL 0.00 −90 −60 −30 0 30 60 90 0.35 Band−I Ced112 0" 6"S 0.30 6"N 10"N

0.25

0.20

0.15

0.10

0.05

0.00 −40 −30 −20 −10 0 10 20 30 40 offset (arcsec)

Figure 4-15 Ratio of the band-integrated BL to band-integrated scattered light in Ced 201 (top panel) and Ced 112 (bottom panel). On large scales the BL distribution is uniform compared to the scattered light in Ced 112 but shows asymmetries in Ced 201. 117

Our BL observations of Ced 112 also trace out the ring of of AEF emitters seen

in mid-IR observations by Siebenmorgen et al. (1998). However, taking the complete

distribution over all four slits into account, the large-scale distribution of both BL

and scattered light is fairly uniform. This is also borne out by the distribution of the

ratio of the band-integrated BL to the scattered light intensity shown in Figure 4-15

(bottom panel). The ratio of BL to scattered light remains fairly constant over large

areas of the nebula.

The spatial distribution of the BL in Ced 201 is interesting. The BL seems asym- metrically distributed with respect to the star; this is primarily due to the asymmetry in the nebulosity itself. Figure 4-1 shows that the nebula is much more extended to the east than the west, toward the molecular cloud that the star is ploughing into.

BL can be detected up to 8000 into the cloud, though the intensity falls oﬀ rapidly.

On the west the BL and the scattered light both fall oﬀ more quickly, but where

the 1500 north slit crosses the arc-like shock-structure, the nebular spectrum shows

the distinct banded features discussed earlier in § 4.4.1.This complex distribution is

true of other tracers as well. The CO(2–1) map of the molecular cloud is distinctly

asymmetric with double peaked proﬁles to the north and single peaked proﬁles to

the south whereas the AEF spectra on the other hand are remarkably symmetric

(Cesarsky et al., 2000). The relative intensity of BL to the scattered light is shown

in the upper two panels of Figure 4-15. The forward-directed scattered light falls

oﬀ much more rapidly than the more isotropic BL, and thus the ratio rises at larger

oﬀsets. 118

4.4.3 BL in Hydrogen Ionization Regions

NGC 5367 and NGC 2023 are both nebulae where the detected spectrum is heav-

ily contaminated by emission from hydrogen recombination. Again as the individual

spectra in each slit do not vary in spectral character we present an average BL spec-

trum for each of the slits in the two nebulae in Figure 4-16. The two show similar

red spectra with intensities decreasing towards the higher-order Balmer lines char-

acteristic of recombination emission from ionized hydrogen but with peaks slightly

short-ward of λ ∼ 380 nm. The NGC 2023 spectrum is more severely contaminated than NGC 5367 because of a stronger ionizing radiation ﬁeld, and the NGC 5367

spectrum also shows a secondary peak at λ ∼ 400 nm similar to other BL spectra, in addition to the peak short-ward of 380 nm. The measurement at the Balmer-jump

(λ = 357 nm) in the other BL spectra is a measure of diﬀerence in absorption, but in these spectra it is a measure of emission from ionized hydrogen. Nevertheless, the peak in the spectrum at λ < 380 nm and the secondary feature at 400 nm are signatures of BL, albeit marginal. 119

1.2e−05

1e−05 )

−1 8e−06 sr −1

Å 6e−06 −1 s

−2 4e−06 NGC 2023 − 0" offset 2e−06

1.5e−06

1e−06 Intensities (erg cm

5e−07 NGC 5367 − 5"N offset 0 3400 3600 3800 4000 4200 4400 4600 4800 5000 wavelength (Å)

Figure 4-16 Average BL spectrum in NGC 2023 (top panel) and NGC 5367 (bottom panel). Due to contamination by emission from hydrogen recombination, these are only marginal detections.

Observation of extended emission, especially in the blue region of the spectrum, is challenging because of competing contribution from scattered light, and emission from hydrogen recombination succeeds in swamping it even further. We would advise other observers to avoid regions with ionized hydrogen in future searches of BL.

4.4.4 Survival vs. In-situ formation of BL Carriers

BL was ﬁrst discovered in the RR, a proto-planetary nebula surrounding the evolved post-AGB star HD 44179. This is a region of active dust production, and therefore it is not surprising to see BL from small molecules with high ﬂuorescence 120

eﬃciencies like anthracene or pyrene, provided there are environments where such

species can remain neutral. The regions in the shadow of the circum-binary disk

of the central binary provides just such a sheltered environment, and this is where

the BL in the RR is predominantly observed. Therefore, a “bottom-up” production

mechanism where the PAH molecules are built from smaller carbon structures is

feasible in the RR. Such small PAH molecules are not expected to survive in the

ISM, being photo-dissociated under interstellar radiation conditions (Jochims et al.,

1999; Le Page et al., 2003, and references therein). Therefore we did not expect to see

BL in other nebulae and predicted its absence in Vijh et al. (2005b). However, this

prediction is valid only if the only possible means for the existence of these molecules

in the other nebulae was travel through the ISM: time-scales of such travel are much

longer than the expected molecular lifetimes in such environments.

Therefore, the detection of BL in other, non-carbon-rich nebulae came as a sur-

prise. However, the formation of PAH molecules through a “top-down” process

emerges as a likely alternative formation mechanism of small PAHs in these loca-

tions. The presence of hydrogenated amorphous carbon (HAC) solids in interstellar

space is well established through the observation of the interstellar absorption feature

at 3.4 µm (see Pendleton, 2004, for a recent review). In an important laboratory in-

vestigation, Scott & Duley (1996) demonstrated that heating of HAC solids will lead to the evolution of aromatic structures, distinguished by their characteristic 3.29 µm

C-H stretch feature. Subsequently, Scott et al. (1997) showed that HAC solids will

decompose under intense UV radiation into an extended sequence of aromatic carbon

clusters with a wide mass range, with small molecules (< 20 atoms) appearing ﬁrst 121

under the lowest-ﬂuence conditions. These laboratory investigations followed ear-

lier theoretical work by Duley et al. (1992) and Taylor & Williams (1993), suggesting

that even smaller hydrocarbons, in particular CH+, could be produced through the

erosion of carbonaceous dust in shocks, cloud-intercloud interfaces and PDRs. We

envision that a similar process could be responsible for the appearance of small PAH

molecules in reﬂection nebulae. Furthermore, larger PAH molecules that are able to

survive conditions in the diﬀuse ISM are likely to undergo photo-fragmentation, as

recently suggested by Rapacioli et al. (2005) to explain the variation in the sizes of

of AEF emitters across a PDR in the reﬂection nebula NGC 7023.

We conclude, therefore, that the most likely source of the small PAHs in reﬂection

nebulae, detected through the BL, are larger carbonaceous structures such as HAC

grains or large, stable PAH clusters. The likely cause is a process of photo-dissociation

and photo-fragmentation, driven by the strongly enhanced radiation ﬁelds of the

nearby early-type illuminating stars. There is other ancillary evidence for grain pro-

cessing in these environments, in particular in Ced 201 (Cesarsky et al., 2000). New

molecules are being formed as the star BD +69 1231 rams into the molecular cloud,

exposing previously sheltered environments to UV radiation. In this case, we are wit-

nessing a dynamic process where these molecules are both being created from larger

species and destroyed by impinging UV radiation. In answer to the question posed

in the introduction: Do small PAHs survive in the ISM? or are the AEFs primarily produced by larger PAHs?: Yes, small PAHs survive as parts of larger entities and

are liberated again in environments of enhanced radiation density. They do exist in

regions where AEFs are observed and should be taken into account in models that 122

try to explain the AEFs. However, since BL is strongly biased toward small PAH

molecules, the detection of a BL signal does not preclude the presence of a much

wider distribution of sizes of aromatic structures.

4.5 Conclusions

We have presented here BL spectra in reﬂection nebulae illuminated by stars

having temperatures between 10,000 – 23,000 K. Of the nebulae studied, Ced 201

and NGC 2023 have previously been shown to exhibit ERE as well (Witt & Boroson,

1990; Witt et al., 1984b; Witt & Schild, 1988). ERE has also been seen in the other two nebulae NGC 5367 and Ced 112 (U. Vijh et al., in preparation). Many of these

nebulae also show the AEFs, attributed to PAHs. Our BL spectra which we have

previously identiﬁed as ﬂuorescence from small, neutral PAHs provide evidence that

these small molecules are present in the ISM. Given that these small PAHs (with

less than 20 carbon atoms) are not expected to survive harsh interstellar conditions,

we believe that the BL spectra provide evidence for the production of these small

molecules from larger structures by the impinging UV radiation from the central

stars. Chapter 5

The Excitation of Extended Red

Emission: New Constraints on its

Carrier from HST Observations of

NGC 7023 1,2

The carrier of the dust-associated photoluminescence process causing the extended red emission (ERE) in many dusty interstellar environments remains unidentiﬁed.

1Based on observations made with the NASA/ESA Hubble Space Telescope, obtained at the

Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555. These observations are associated with program #9471. 2Submitted to The Astrophysical Journal, May 2005, Authors: A. N. Witt, K. D. Gordon, U. P.

Vijh, P. H. Sell, T. L. Smith, & Rui-Hua Xie

123 124

Several competing models are more or less able to match the observed broad, un-

structured ERE band. We now constrain the character of the ERE carrier further

by determining the wavelengths of the radiation that initiates the ERE. Using the

imaging capabilities of the Hubble Space Telescope, we have resolved the width of

narrow ERE ﬁlaments appearing on the surfaces of externally illuminated molecular

clouds in the bright reﬂection nebula NGC 7023 and compared them with the depth

of penetration of radiation of known wavelengths into the same cloud surfaces. We

identify photons with wavelengths shortward of 118 nm as the source of ERE initi-

ation, not to be confused with ERE excitation. There are strong indications from

the well-studied ERE in the Red Rectangle nebula and in the high-|b| Galactic cirrus that the photon flux with wavelengths shortward of 118 nm is too small to actually excite the observed ERE, even with 100% quantum efficiency. We conclude, therefore, that ERE excitation results from a two-step process. The first step, involving far-UV photons with E > 10.5 eV, leads to the creation of the ERE carrier, most likely through photo-ionization or photo-dissociation of an existing precursor. The second step, involving more abundant near-UV/optical photons, consists of the optical pumping of the previously created carrier, followed by subsequent de-excitation via photoluminescence. The latter process can occur many times for a single particle, depending upon the lifetime of the ERE carrier in its active state. While none of the previously proposed ERE models can match these new constraints, we note that under interstellar conditions most polycyclic aromatic hydrocarbon (PAH) molecules are ionized to the di-cation stage by photons with E > 10.5 eV and that the electronic energy level structure of PAH di-cations is consistent with fluorescence in the wave- 125

length band of the ERE. Therefore, PAH di-cations deserve further study as potential

carriers of the ERE.

5.1 Introduction

Extended Red Emission (ERE) results from a dust-related optical photoluminescence process in the interstellar medium. It was ﬁrst observed in the peculiar bi-polar Red Rectangle nebula (Cohen, 1975; Schmidt et al., 1980). Soon af-

ter its initial discovery, ERE was shown to be present in many other dusty in-

terstellar environments, albeit at mostly lower intensities, e.g. in reﬂection nebu-

lae (Witt et al., 1984b; Witt & Boroson, 1990), HII regions (Perrin & Sivan, 1992;

Darbon et al., 2000), carbon-rich planetary nebulae (Furton & Witt, 1990, 1992), as well as the diﬀuse interstellar medium of the Milky Way (Gordon et al., 1998) and

other galaxies (Perrin et al., 1995; Pierini et al., 2002). The ERE intensity in many

dusty sources is proportional to the local density of the illuminating radiation ﬁeld.

This provides a strong argument in favor of the suggestion that the ERE is photon-

driven, i.e. the ERE is a photoluminescence process. The spectroscopic signature of

the ERE is a broad (60 to 100 nm FWHM), unstructured emission band, typically

extending from 540 nm to beyond 900 nm in wavelength. The peak wavelength of the

ERE band varies from somewhat longward of 600 nm to beyond 800 nm in response

to varying environmental conditions, in particular the density of the illuminating ul-

traviolet (UV) radiation ﬁeld (Smith & Witt, 2002). This is a deﬁning characteristic

of the ERE, which distinguishes it from other emission features with essentially in- 126

variable emission wavelengths. At high Galactic latitudes, the intensity of the ERE is

comparable to that of the dust-scattered diﬀuse galactic light (DGL), which has led

to an estimated lower limit of the ERE quantum yield of 10 ± 3% and the conclusion

that the ERE carrier must be a major contributor to the absorption part of inter-

stellar extinction at UV/visible wavelengths in the Milky Way (Gordon et al., 1998).

Its ubiquitous presence in radiation environments ranging in strength over more than

ﬁve orders of magnitude testiﬁes to the relative robustness of the ERE carrier. More

details about ERE studies done over the past three decades may be found in a recent

review by Witt & Vijh (2004).

None of the current models for interstellar dust (e.g. Draine, 2004; Zubko et al.,

2004) predict ERE, nor do they provide a satisfactory post-facto explanation for the observed characteristics of the ERE. As reviewed by Witt & Vijh (2004), nu-

merous ad-hoc models have been proposed to account for the ERE. Most of them

suggest ERE carriers in the form of large molecular structures or nanometer-sized

grains, in which electronic excitations by shorter-wavelength photons are followed by

eﬃcient electronic radiative transitions across a bandgap of < 2 eV. Well-studied ex-

amples of such processes are ﬂuorescence and phosphorescence in organic molecules

and the photoluminescence in semiconductor nanoparticles. Such particles may be

able to meet the constraints posed by the observed spectral characteristics, includ-

ing their environment-dependent variations, and by the inferred quantum eﬃciency

(Smith & Witt, 2002). However, all carrier models proposed so far suﬀer from the

following dilemma.

The excitation of fluorescence or photoluminescence generally occurs with opti- 127 mum efficiency at energies just slightly above the bandgap energy, i.e. photons at wavelengths just shortward of the ERE band should be most suitable for the initiation of the ERE. The difference between the optimum excitation energy and the energy at which the luminescence emerges, the so-called Stokes shift, is typically of the order of 1 eV or less in typical photoluminescent agents. At first glance, this does not appear to be the case for the ERE in view of two previous studies of the excitation requirements of the ERE. The first study by Witt & Schild (1985), based upon multi-band surface brightness photometry of several reflection nebulae, concluded that the ratio of the intra-nebular absorption of the ERE-exciting radiation,

A(exc), to the color excess E(B-V) suﬀered by the intra-nebular radiation reaching the same nebular location had values A(exc)/E(B-V) ∼ 7-9. When related to the observed UV/optical extinction curves of the illuminating stars for the three nebulae included in that study, it was concluded that the required ratio of A(exc)/E(B-V) could only be found in the immediate vicinity of the 217.5 nm absorption peak of the respective extinction curves or at wavelengths shortward of 150 nm, where the far-

UV rise in the extinction again provides the required opacity. While the multi-valued nature of the extinction curve prevented a unique estimate of the excitation energy by this method, the results indicated that ultraviolet rather than optical photons are responsible for ERE excitation.

The second study related to ERE excitation by Darbon et al. (1999) compiled the records of positive ERE detections as well as the results of negative searches for

ERE in a large number of nebulae and related them to the eﬀective temperatures of the exciting stars. While ERE is present in abundance (with few exceptions) among 128

nebulae illuminated by stars with Teﬀ > 10, 000 K, not a single nebula with stars of

Teff < 7000 K exhibits ERE, although dust is present with significant optical depths in the nebulae associated with the latter group of stars. Stars with Teff = 7000 K

have spectral energy distributions (SEDs) with a far-UV cutoﬀ near 170 nm, cor-

responding to an energy of 7.3 eV, while stars with Teﬀ = 10, 000 K have their

corresponding far-UV ﬂux cutoﬀ near 110 nm, corresponding to photon energies of

11.2 eV (Kurucz et al., 1974). The results of Darbon et al. (1999) can thus be in-

terpreted to mean that photons with energies somewhere within the range between

7.3 eV and 11.2 eV represent the minimum photon energy required to initiate the

ERE process.

The results of Witt & Schild (1985) are fully consistent with this conclusion.

However, the Darbon et al. (1999) results eliminate the 217.5 nm absorption peak

of the interstellar extinction curve as a possible source of ERE excitation, because

Teff = 7000 K stellar atmospheres still have significant flux at this wavelength, yet

they are not found to be able to excite ERE. Further supporting observations come

from the more recent work of Darbon et al. (2000), who compared the spatial distri-

butions of the ERE and the emissions in various unidentiﬁed infrared emission (UIR)

bands, attributed to aromatic hydrocarbon structures, in the HII region Sh 152.

While the UIR bands are strongest at the interfaces between the ionized volume of

the nebula and beyond, the ERE is strongly correlated with the distribution of the

ionized hydrogen located in the central region of the HII region. This provides evi-

dence for the fact that far-UV photons are needed for the initiation of the ERE as

well as proof of the robust nature of the ERE carrier, while the aromatic hydrocarbon 129 structures can apparently be successfully excited by radiation at longer wavelengths, which is able to penetrate into the nebular environments outside the ionization front.

This latter picture is also supported by the recent ﬁndings by Smith et al. (2004) that the 3.3 µm emitter does not require UV photons for excitation.

The results discussed above suggest a Stokes shift in excess of 6 eV for the ERE process, if in fact the far-UV photons required to initiate the ERE are also the source of the actual ERE excitation. Another, more likely possibility is that the far-UV photons simply create and maintain the ERE emitter, for example by photo-ionization or -dissociation of a precursor, which could then be capable of photoluminescence upon excitation by abundant lower-energy photons. The ERE excitation could thus be a two-step process, with a far-UV photon required to produce the carrier, a particle capable of strong absorption in the optical/near-UV region in which lower-energy- photons closer in energy to the 2 eV band gap would excite the actual ERE.

To decide among these possibilities, it is essential to determine the wavelength of the ERE initiation as accurately as possible. This must be followed by a comparison of the number density of photons with wavelengths shortward of this limit in a given system with the number density of ERE photons generated within the system. If the latter is greatly in excess of the former, a two-step excitation process is the most likely explanation. The observations reported in this chapter provide a basis for a determination of the critical wavelength of ERE initiation. We employ the fact that the ERE appears in narrow filamentary structures (e.g. Witt & Malin, 1989) on the surfaces of molecular cloud clumps in reflection nebulae, in particular in the bright reflection nebula NGC 7023. We assume that the physical width of a cloud edge as seen in 130

the light of the ERE is determined by the depth of penetration of the photons giving

rise to this particular nebular emission. We will compare these penetration depths

with those of photons of known wavelengths giving rise to the appearance of the

same cloud edges in the light of photo-excited 1-0 S(1) H2 vibrational ﬂuorescence at

2.12 µm(Lemaire et al., 1996; Field et al., 1998; Takami et al., 2000; An & Sellgren,

2003) and in the light of scattered radiation in the z-band. In NGC 7023, the former

is the result of photo-excitation in the Lyman and Werner bands in H2 (Takami et al.,

2000), which occurs near 110 nm and shortward in wavelength, while the latter is the

result of simple scattering by dust at the eﬀective wavelength of observation near

900 nm. The penetration depths of the respective exciting radiations are inversely

proportional to the respective extinction coeﬃcients for the exciting radiations, with

local gas/dust densities canceling when ratios are considered. The unknown wave-

length region of ERE initiation can then be estimated from the derived ratios of

extinction coeﬃcients, taking into account the opacity sources present in the nebular

environment.

§ 5.2 of this chapter presents the observations and reductions; in § 5.3 we present

the results. The discussion of the implications of our results are contained in § 5.4. In particular, we apply the wavelength constraint for ERE initiation to two well-studied

ERE sources to assess the likelihood of a two-step excitation process for the ERE.

This is followed by a set of conclusions in § 5.5. 131

5.2 Observations and Data Reduction

Figure 5-1 Location of the regions shown in Figs. 5-2 and 5-3 on a F606W WFPC2 image of NGC 7023 (Gordon et al., 2000). 132

The prominent northwest ﬁlament in NGC 7023 (see Figure 5-1) was imaged with

the Advanced Camera for Surveys (ACS) and the Near Infrared Camera and Multi

Object Spectrometer (NICMOS) on the Hubble Space Telescope (HST) as part of

HST program #9741. The ACS images were taken in F475W, F625W, F850LP, and

F656N which correspond to SDSS g, SDSS r, SDSS z, and Hα ﬁlters. The centers

of the ACS images were oﬀset to avoid the bright central star of NGC 7023. The

three broad-band ACS images were taken to allow for the creation of a continuum-

subtracted ERE image and the narrow band F656N ﬁlter to measure the contribution

from scattered Hα emission to this ERE image. The origin of the scattered Hα

radiation is the strong Hα emission line in the spectrum of the central illuminating

star, HD 200775. The NICMOS images were taken with the NIC2 camera and the

F212N and F215N ﬁlters which measure the 1-0 S(1) H2 emission line wavelength

range and associated red continuum at 2.121 µm and 2.15 µm respectively.

The observations in each ACS filter were taken split between two images with a large enough dither to fill in the gap between the two chips. The total exposure times were 1000, 1000, 1200, and 1560 s for the F475W, F625W, F850LP, and F658N filters, respectively. The two images per filter were combined with the online multidrizzle and additionally processed using LACOSMIC (van Dokkum, 2001) to identify cosmic

rays which were not removed with multidrizzle. Residual cosmic rays are present in

the ﬁnal mosaics, especially in the gap between the two ACS chips where we only have

a single measurement. These residual cosmic rays do not aﬀect our measurements as

they rely on structures much larger than a cosmic ray. 133

N W

SDSS g SDSS r ACS F475W ACS F625W

SDSS z ACS F850LP ACS F658N

ERE H−alpha

Figure 5-2 A 1.05 × 1.00 region, rotated 37◦.25 from North, and centered on 21h01m3200.505, +68◦1002500.86 (2000) is shown for the four observed ACS ﬁlters and continuum-subtracted ERE and Hα images. This region encompasses the northwest ﬁlament of NGC 7023.

A region extracted from the ACS images focusing on the northwest ﬁlament is

shown in Fig. 5-2 for all four observed ﬁlters and the continuum-subtracted ERE and

Hα images. The ERE image was created by first creating a continuum image due to dust scattered light using a linear combination of the F475W and F850LP images. These two filters have central wavelengths (474 nm and 905 nm, respectively) to the blue and red of the ERE peak seen in this filament which is at ∼ 650 nm 134

(Gordon et al., 2000). The resulting continuum image was subtracted from the observed F625W image to produce the ERE image. The success of the continuum

subtraction can be seen by the much reduced strength of the clump southwest of the

ﬁlament (brightest extended source in all four observed ACS bands). This clump is

composed of only scattered light as it has a linearly decreasing spectrum across the

three broad ACS bands as measured in a 30×30 pixel box centered on the clump. The

Hα image was created by subtracting the F625W image from the F658N image after scaling the F625W image by the ratio of ﬂuxes of stars measured in both images.

The prominence of ERE in this ﬁlament is clearly seen as the F625W image is much sharper than the F475W and F850LP images. If the F625W image were dominated by scattered light and not ERE, it would have an intermediate appearance between the F475W and F850LP. One possible reason for this sharper appearance could be

Hα emission. This is not the case as the sharpness of the continuum-subtracted

Hα image more closely matches the F425W and F850LP images than the F625W or

ERE images. Note that the Hα image has the same morphology as the scattered- light dominated F425W and F850LP images. This is because the central star of

NGC 7023, HD 200775, has a strong Hα emission line (Gordon et al., 2000). Thus,

the continuum-subtracted Hα image traces the scattered Hα light of the central star,

not Hα emission from extended gas in the nebula. 135

NICMOS NICMOS F212N F215N

ERE H2 ACS F850LP

Figure 5-3 A 0.045 × 0.045 region, rotated 40◦.0 from North, and centered on 21h01m3200.644, +68◦1002800.17 (2000) is shown for the two observed NICMOS ﬁlters, continuum-subtracted ERE, continuum-subtracted H2, and ACS F850LP images. This region encompasses the bright portion of the northwest ﬁlament of NGC 7023.

The observations in the two NICMOS filters were taken split between 4 images with a spiral dither pattern. The total exposure times were 352 s in both the F212N and F215N filters. The 4 separate images were combined using custom scripts to allow for careful masking of bad pixels and the coronographic spot. A single star was present in each image allowing for accurate registration between images and to the ACS observations. The mosaics of the F212N and F215N observations are shown in the top row of Fig. 5-3 displayed with the same scale. It is clear from these two images that this filament has very strong H2 emission which was known from existing near-infrared spectroscopy (Gordon et al., 2000). The H2 image was created by subtracting the calibrated F215N image from the calibrated F212N image. 136

The accuracy of the subtraction of the continuum can be easily seen as the stellar

image disappears completely as well as the diﬀraction spikes from the central star of

NGC 7023. In addition to the NICMOS images, the continuum-subtracted ERE and

ACS F850LP images are shown for the same region in Fig. 5-3.

5.3 Results

5.3.1 The ERE, H2, and z-Band Morphology of NW Fila-

ments

We examined the ﬁlamentary morphology of the NW PDR in NGC 7023 on two

spatial scales. The first scale is defined by the more limited field of view of NICMOS

and focuses on the complex, multi-ﬁlament structure in the center part of the NW

PDR in NGC 7023 shown in Figure 5-3. Our observations of the 1-0 S(1) emission from molecular hydrogen are limited to this ﬁeld of view, and we are presenting this image together with the ERE and z-band images limited to this ﬁeld. The wider

ﬁeld of view of the ACS deﬁnes the second scale shown in Figure 5-2; it includes

the southern and northern extensions of the ﬁlaments in Figure 5-3, for which we

can compare the ERE structure only with the corresponding structure in the z-band.

The data shown in Figures 5-2 and 5-3 represent the highest-resolution images of the

NGC 7023 NW PDR to date. The location of these images within the larger extent

of NGC 7023 is indicated in Figure 5-1.

The structure of the NW PDR shown in the more limited ﬁeld of NICMOS (Fig- 137 ure 5-3) is highly complex and consists of a series of narrow ﬁlaments with typical

00 widths of 1 when viewed in the light of ERE and the light of H2. These ﬁlaments are fully resolved in both images. The angular resolution is marginally higher in the ACS images compared to the NICMOS images. This is based on comparisons of stellar images which yielded ratios FWHM(ACS)/FWHM(NICMOS) = 0.93, with stellar images having diameters with FWHM of 0.1025 and 0.110200 respectively. When viewed in the z-band, which is dominated by dust-scattered light, the same structures appear much more diﬀuse, although they were imaged with the same angular resolution.

This comparison suggests that the sharpness of the structures as seen in ERE and H2 is not a result of a limited physical extent of the structures, e.g. thin sheets viewed edge-on, but that they have a physical width that is wider than indicated by the ERE and H2 ﬁlament images. The narrow ERE and H2 ﬁlaments, therefore, may represent just the edges of more extended molecular cloud clumps illuminated by the central

B3Ve star in NGC 7023, HD 200775. 138

Case 4 Case 1

ERE H2

4 3 1 2 4 3 1 2

Case 3 Case 2

Figure 5-4 Four cuts through the ERE (solid line) and H2 (dotted line) images are given illustrating the 4 cases seen. These cases are 1) near-perfect match between ERE and H2 filaments, 2) evidence for H2 self-shielding, 3) H2 filaments broader than ERE, and 4) H2 filaments narrower and in front of ERE filaments. The x-axes of the graphs increase in the direction of propagation of the exciting radiation.

There is independent evidence from the interferometric HCO+ observations by

Fuente et al. (1996) for the existence of at least two but possibly four high-density molecular clumps in our field of view that can be distinguished on the basis of their differing radial velocities. In particular, the extended filament that begins on the 139 right-hand side of Figure 5-2 and turns upward, passing near the embedded star (star

K in the map of An & Sellgren, 2003), appears to be a coherent structure moving at 2.4 km s−1, while the group of narrow ﬁlaments shown in Figure 5-3 closest to the illuminating source belongs mostly to a structure moving at 4.0 km s−1. The apparent complexity of the ﬁlamentary structure of the NW PDR is most likely the result of the line-of-sight superposition of several independently moving molecular clouds, whose surfaces facing HD 200775 are currently being photo-dissociated.

A comparison of the morphology of the H2 and ERE filaments in Figure 5-3 reveals two significant facts. One, the ERE filaments appear to be about as sharp as the corresponding H2 filaments. We will investigate this aspect further in §§ 5.3.3; for now we can conclude that the radiation initiating the H2 fluorescence and the ERE, respectively, apparently faces similar optical depths per unit mass of molecular cloud environment. Two, the spatial correlation between the distribution of the ERE and the H2 emission in Figure 5-3 is not exceptionally strong. ERE is present wherever

H2 emission is seen, but there are bright ERE ﬁlaments, especially in the upper part of Figure 5-3, for which the corresponding H2 features are either much weaker or absent altogether. A similar qualitative relationship between ERE and H2 emission has previously been noted by Field et al. (1994) in a study of the reﬂection nebula

NGC 2023 and was conﬁrmed in a more detailed study of NGC 7023 by Lemaire et al.

(1996).

When comparing the detailed proﬁle shapes and positions of ERE and H2 ﬁla- ments, we can identify four classes, illustrated in Figure 5-4. These can be described as follows. Case 1 represents locations where ERE and H2 exhibit nearly identical 140

structures. Out of 36 cuts examined, seven fall into this category. Class 2 includes

cases (17/36) where the ﬁrst H2 ﬁlament is very narrow and recessed with respect to

the corresponding ERE ﬁlament. A second ERE ﬁlament is not matched by H2. We

believe these to be instances of H2-self-shielding. Case 3 represents instances (9/36)

where the H2 ﬁlament is broader than the corresponding ERE ﬁlament, while Case

4 illustrates the rather rare (3/36) circumstance where the H2 ﬁlament is narrower

than the ERE ﬁlament and peaks slightly in front of the ERE ﬁlament.

The near-absence of H2 emission (Case 2) in regions exhibiting strong ERE can-

not be explained by diﬀerences in line-of-sight extinction. We just concluded that

both types of emission appear to be initiated by photons of similar energy; the re-

sulting outgoing 1-0 S(1) H2 emission faces an opacity only 1/8 as large as the re-

sulting ERE. Hence, if one type of emission should be missing due to extinction, it

would be the ERE, not the H2 emission, contrary to our observations. As noted by

Field et al. (1994) and Lemaire et al. (1996), this result provides a strong argument against models in which the ERE carrier is initiated chemically by hydrogenation of

carbonaceous dust in the hot atomic hydrogen gas found in H2 photo-dissociation fronts (Witt & Schild, 1988; Duley & Williams, 1990). Lack of carriers also cannot

be the explanation for the absence of H2 luminescence; H2 molecules are the principal constituents of molecular clouds in the ISM. 141

Figure 5-5 The density-dependence of the dust optical depth of a PDR (solid line) and the eﬃciency of converting far-UV photons into near-IR 1-0 S(1) H2 photons (dotted line) is shown. The NGC 7023 NW PDR involves molecular gas with densities ranging 4 −3 6 −3 from n(H2) ≥ 10 cm to n(H2) ≥ 10 cm (An & Sellgren, 2003).

We notice that ERE filaments in the NW PDR that lack H2 counterparts of corresponding surface brightness tend to be broader than those that have bright H2 counterparts. If this broader, more diffuse appearance is a reflection of the penetration depth of the radiation initiating the ERE, this then indicates that the density in the

ERE filaments lacking H2 counterparts is substantially lower. Under these conditions, the UV pumping efficiency as measured by the number of 1-0 S(1) photons emitted per incident photon in the wavelength range 110.4 – 91.2 nm is significantly reduced

(Draine & Bertoldi, 1996), because dust is absorbing an increasing fraction of the photons needed for the excitation of H2. This is because molecular hydrogen is unable to be re-formed as quickly under lower-density conditions, the rate being proportional 142

to the square of the density. As we had seen earlier, the ERE carrier is part of the

dust and a likely competitor for photons in the far-UV range. Thus, with competition

lacking from molecular hydrogen, the ERE carrier is able to absorb a greater fraction

of the available photons under lower-density conditions. We illustrate the dependence

of the dust optical depth of a PDR and the eﬃciency of converting absorbed far-UV

photons into near-IR photons in the 1-0 S(1) H2 transition in Figure 5-5, calculated

from relations provided by Draine & Bertoldi (1996). The dust optical depth of the

PDR is measured from the front of the molecular cloud to the point at which 50% of

the H2 is photo-dissociated.

It appears likely, therefore, that the lack of a detailed spatial correlation between

ERE ﬁlaments and the 1-0 S(1) emission of H2 is a result of density variations by about one order of magnitude between high-density ﬁlaments, where both ERE and

H2 are bright, and lower density cloud faces, where the ERE is dominant.

5.3.2 Determination of RV in NW Filaments

When viewing the filaments in the light of different emission mechanisms, e.g. the fluorescent 1-0 S(1) vibrational emission of H2 at 2.12 µm and the much more diffuse appearance of the same structures in the light of dust-scattered light in the z-band (Figure 5-3), we are comparing the relative penetration into the cloud surfaces

of the respective exciting radiations. In the case of the ﬂuorescent H2 emission, the

exciting radiation is Lyman and Werner band photons absorbed by H2 molecules

at wavelengths λ ≤ 110.4 nm. The opacity restricting the transfer of such photons 143 into the cloud faces is most certainly extinction by dust and, depending on local density conditions, self-shielding by H2 (Draine & Bertoldi, 1996). In the z-band, dust extinction is the only significant opacity source. The relative penetration depth along a given line-of-sight from the central star into a molecular cloud is determined by the inverse ratio of these two opacities. Taking into account the possible contribution by H2 self-shielding to the opacity, the ratio of the widths of the H2 and the z filaments at identical locations, then, is an upper limit to the ratio of the dust opacities at the two effective wavelengths involved, 110.4 nm and 900 nm, respectively. This ratio can now be used to identify the appropriate dust extinction curve for the molecular cloud fronts in the NGC 7023 NW PDR.

Figure 5-6 Extinction curves generated from the CCM RV -dependent relationship are shown for a range of RV values. 144

Figure 5-7 Histogram of the ratios FWHM(z)/FWHM(H2) and FWHM(z)/FWHM(ERE) of ﬁlaments widths.

Cardelli et al. (1989, hereafter CCM) showed that the wide range of extinction

curves encountered in diﬀerent galactic environments can be reproduced as a family of

functions dependent upon a single parameter, the ratio of total-to-selective extinction,

RV , as illustrated in Figure 5-6. The low-density diﬀuse ISM in the Milky Way is most

commonly represented by CCM curves with RV = 3.1 while dense molecular cloud material exhibits CCM extinction curves with RV in the range from 5 to 6. As long as we adopt the CCM formalism, these curves can be distinguished from each other by measuring a single ratio of extinctions at two widely separated wavelengths. As we noted before, such extinction ratios can be estimated from the relative penetration depth at the respective wavelengths. We measured the FWHM of the H2 and z-band

ﬁlaments (Figure 5-3) at 22 positions where such measurements could be carried out 145

without interference from blends by closely spaced H2 ﬁlaments, whose z-counterparts

would merge into a single proﬁle. A histogram of the measurements of the ratio of

FWHM(z)/FWHM(H2) is shown in Figure 5-7 by a dotted line. The average of the measured ratios of FWHM(z)/FWHW(H2) is 2.74 ± 0.61. In Figure 5-8 we show

the ratio of the dust extinction cross sections at 110.4 nm and 900 nm, the latter

being the eﬀective wavelength of the z-band, as a function of RV for CCM extinction

curves, and we have entered our measured ratio as a horizontal bar. We conclude

+0.61 from this ﬁgure that a CCM extinction curve with RV =5.62−0.48 is a good estimate

for the dust extinction curve in the molecular clouds making up the NW PDR in

NGC 7023. If H2 self-shielding is a signiﬁcant contributor to the total extinction for

the H2-exciting radiation, this value of RV is a lower limit. The fact that RV =5.62

is very typical of molecular cloud environments (RV = 5.5) suggest, however, that

H2 self-shielding is a minor contributor to the opacity for the H2-exciting radiation

in the ﬁlaments measured in this experiment.

5.3.3 Width of the ERE Filaments

Employing the same rationale as in the previous section, we now compare the

relative widths of the ERE and z ﬁlaments in Figures 5-2 and 5-3 in identical locations.

A total of 41 measurements were made in the larger ﬁeld (Fig. 5-2) and 26 such

measurements were made in the smaller ﬁeld (Fig. 5-3). A histogram of the resulting

ratios FWHM(z)/FWHM(ERE) is shown in Figure 5-7 (solid line). The two sets of

measurements resulted in average ratios of FWHM(z)/FWHM(ERE) of 2.52 ± 0.65 146

and 2.32 ± 0.64, respectively, with a weighted average of 2.40 ± 0.65. Comparing this value to the corresponding ratio of FWHM(z)/FWHM(H2)=2.74±0.61, we conclude that on average the ERE ﬁlaments are about 10% broader than the corresponding H2

ﬁlaments, if they are present, and that the opacity facing the radiation that initiates

the ERE is, therefore, about 10% less than the opacity encountered by the radiation

that excites the H2 1-0 S(1) radiation.

Figure 5-8 The ratio of extinctions at 110.4 nm and 900 nm is plotted as a solid line. The allowed region from the measurement of the FWHM(z)/FWHM(H2) ratio is shown as the shaded region.

5.3.4 Wavelength of ERE Initiation

In order to identify the wavelengths of the radiation that appears to be required

for ERE initiation, we plotted a CCM extinction curve for RV =5.62, plus two dotted

extinction curves representing the uncertainty limits in the RV value, found to be ap- 147

propriate for the dust in the molecular material in NGC 7023, in Figure 5-9. Assuming

a ratio A(z)/A(V) = 0.59 for the RV = 5.62 case, we place our measured ratio of

FWHM(z)/FWHM(ERE) onto this graph as a horizontal line at A(λ)/A(V) = 1.43.

This constraint implies that only those wavelength ranges in which the extinction exceeds this horizontal limit exhibit adequate dust opacity to be consistent with our data in Figure 5-7. Resulting from the multi-valued nature of the extinction curve, this identiﬁes two spectral regions: the range from 250 nm to 189 nm, where the horizontal line crosses the 217.5 nm bump in the extinction curve, and the region shortward of 118 nm in the far-UV, where the far-UV rise of the extinction curve reaches values in excess of our opacity constraint. Thus, similar to the investigation by Witt & Schild (1985), our method alone does not result in the identiﬁcation of a

unique spectral region responsible for ERE initiation, but it clearly excludes near-UV

or optical radiation slightly shortward of the ERE band from consideration. 148

Figure 5-9 The limits on the ERE excitation wavelength are shown on top of a CCM RV = 5.62 extinction curve. The constraint imposed by the ERE filament widths is discussed in §§ 5.3.3; the constraints due to the limits on the effective temperatures of the exciting stars are discussed in §§ 5.3.4; and the two vertical lines labeled < 100% and < 50% related to the insufficiency of far-UV photons for the actual excitation of the ERE are discussed in §§ 5.4.4

We can now narrow down the choice between the two UV spectral ranges identiﬁed

above by applying the results of Darbon et al. (1999) as a further constraint. In short,

Darbon et al. (1999) found that stars with Teff ≥ 10, 000 K can excite ERE and very commonly do, while stars with Teff ≤ 7000 K do not. We entered the FUV cut-offs of the SEDs of atmospheres with these two temperatures in Figure 5-9 as two vertical

lines. The long-wavelength limit of the spectral region in which ERE is initiated

radiatively is deﬁned by the point at which our applicable extinction curve enters from

below the cross-hatched area limited by the three constraint lines. For RV = 5.62,

this wavelength is 118 nm, corresponding to a photon energy near 10.5 eV. Therefore, 149

the application of the constraints by Darbon et al. (1999) eliminates the 217.5 nm

extinction band as the source of excitation of the ERE. Our new constraint based upon

the penetration of the exciting radiation into the cloud face considerably narrows the

acceptable wavelength range for ERE initiation allowed by the Darbon et al. (1999)

results, however. Both sets of constraints are consistent with the possibility that

photons with wavelengths between 118 nm and 91.2 nm, the latter being set by

absorption by HI, can contribute to ERE initiation.

5.4 Discussion

Having obtained a fairly accurate estimate of the wavelengths of the photons

needed to initiate the ERE will be critical for determining further details of the ERE

process as well as the ERE carrier. As we stated earlier, our present observations do

not allow us to determine whether photons with E > 10.5 eV are required to excite the ERE carrier, with (at best) one ERE photon resulting from each absorption of each far-UV photon, or whether the presence of far-UV photons is merely required for the creation of the ERE carrier, e.g. by either photo-ionization or photo-dissociation of an ERE-precursor, followed by pumping of the newly created ERE carrier by lower- energy optical photons. However, this issue can be resolved directly by counting the number of photons available in the 118 nm to 91.2 nm range that are absorbed in a given system and comparing this to the number of ERE photons emitted. We will return to this question in §§ 5.4.3 and §§ 5.4.4. 150

5.4.1 Impact on Existing ERE Models

While the nature of the ERE as a dust-related interstellar photoluminescence pro-

cess has been undisputed, many diverse proposals have been advanced for possible

identiﬁcations of the ERE carrier (see Witt & Vijh, 2004, for a review). With the de-

termination of the minimum wavelength of the radiation required for ERE initiation,

many of these proposals can now be ruled out. In discussing these proposals below,

we will focus primarily on the excitation requirements and not on respective failures

to meet other observational constraints.

Neutral PAH Molecules

d’Hendecourt et al. (1986) proposed that ﬂuorescence, i.e. electronic transitions

from the S1 state to vibrational levels of the S0 ground state, or phosphorescence,

i.e. electronic transitions from the T1 state to vibrational levels of the S0 ground state, in neutral polycyclic aromatic hydrocarbons (PAHs) could be responsible for the ERE. The same species are considered the likely source of the ubiquitous mid-

IR aromatic emission features (AEF) (Bakes et al., 2004; Onaka, 2004; Peeters et al.,

2004). In the laboratory, many PAH molecules are known to ﬂuoresce eﬃciently

at optical wavelengths. However, as noted by d’Hendecourt et al. (1986), PAH ﬂu-

orescence generally occurs at much shorter wavelengths than the ERE band while

phosphorescence frequently occurs in the same spectral region as the ERE. Although

the T1 - S0 transitions are forbidden, isolated PAH molecules in the ISM could con-

ceivably be strong sources of phosphorescence, provided that the T1 state can attain 151

high populations. The T1 state, which is always at a lower energy compared to the S1 state in neutral PAH molecules, is populated by intersystem crossings from excited states of the singlet manifold to highly-excited vibrational levels of the T1 state. This

implies, therefore, that the excitation of phosphorescence occurs via the same absorp-

tion bands that excite the ﬂuorescence in these molecules. The absorption spectra of

neutral PAH molecules are well known; their absorption bands are found immediately

shortward in wavelength of where their S1 - S0 ﬂuorescence occurs, with the strongest

bands found generally near 200 nm. Thus, the wavelengths of radiation required for

the excitation of PAH ﬂuorescence and PAH phosphorescence are two to three times

longer than what was found in the present investigation. Recently, neutral PAH ﬂuo-

rescence (Vijh et al., 2004, 2005b) has been detected in the Red Rectangle, an object

in which ERE is present as well, and the spatial distributions of the two emissions

were found to be distinctly diﬀerent. This is only further support for our conclusion

that neutral PAH phosphorescence is not the source of ERE.

PAH Clusters

Seahra & Duley (1999) suggested that PAH clusters in the form of stacks or ag-

gregates of PAH molecules with up to 700 carbon atoms could produce ERE-like pho-

toluminescence in a band centered near 700 nm. The same authors (Duley & Seahra,

1998) had previously proposed that the same PAH clusters could be responsible for

the interstellar absorption band at 217.5 nm as well as the mid-IR AEF bands. The

energy for both the ERE and the AEF in this model was to come from the absorp-

tion in the 217.5 nm band. While the present penetration study does not exclude this 152

possibility, the application of the Darbon et al. (1999) constraint (Figure 5-9) does.

The PAH cluster model is therefore not supported by the ERE excitation results.

Hydrogenated Amorphous Carbon

Photoluminescence by grains consisting of hydrogenated amorphous carbon (HAC)

or coated by HAC mantles was an early candidate for the ERE process (Duley, 1985;

Witt & Schild, 1988; Duley & Williams, 1990; Jones et al., 1990). With a band gap

of over 3 eV (Robertson, 1996), HAC not only exhibits eﬃcient photoluminescence at wavelengths much shorter than observed in ERE, it also is eﬃciently excited at

optical/near-UV wavelengths (Watanabe et al., 1982), which is in conﬂict with our

present results. The latter objection can also be raised against a material closely re-

lated to HAC, referred to as quenched carbonaceous composite (QCC) (Sakata et al.,

1992). While the band gap of various HAC materials can be reduced by a variety

of treatments, the most eﬃcient excitation occurs always just shortward of the emis-

sion band resulting from HAC photoluminescence. This puts these models in serious

conﬂict with our new excitation constraints.

Silicon Nanoparticles

An ERE model that received intensive study in recent years involves photoluminescence by oxygen-passivated silicon nanoparticles (SNPs), containing between about 200 and 6000 silicon atoms (Witt et al., 1998; Ledoux et al., 1998, 2000, 2001,

2002; Smith & Witt, 2002; Li & Draine, 2002). The band gap of SNPs is size depen-

dent due to quantum conﬁnement and with suitable limits on the nanoparticle sizes it 153

readily matches the observed wavelengths of the ERE band. However, again as in all

instances of typical semiconductor photoluminescence, the most eﬀective excitation

occurs at energies just slightly above the band gap energies, i.e. slightly above 2.3 eV.

In addition, SNPs are photo-ionized when exposed to photons with energies in excess

of 5.1 eV (Fuke et al., 1993), with the result that photoluminescence from subsequent

excitations is quenched (Nirmal & Brus, 1999; Smith & Witt, 2002). Consequently,

our new constraints regarding ERE initialization also do not favor the SNP model.

5.4.2 Possible New ERE Carriers

All previously considered ERE carrier models appear to be inconsistent with the

requirement that photons with energies in excess of 10.5 eV are required to initiate

the ERE in astrophysical environments. There are several reasons for assuming that

the energy limit of 10.5 eV is not directly related to the actual excitation of the ERE.

Below, we will argue that the number density of photons with energies in excess

of 10.5 eV in typical ERE-producing environments is insuﬃcient to account for the

number of ERE photons emitted from those environments, if the E > 10.5 eV photons were to be the direct source of excitation, even if a photon conversion eﬃciency of

100% is assumed. A possible solution to this dilemma is provided by models in which the E > 10.5 eV photons simply create the ERE carrier, with a small number of FUV photons required for its maintenance. The ERE carrier thus created needs to have the ability for eﬃcient photoluminescence in the 540 – 900 nm wavelength range and strong absorption throughout the optical/UV regions of the spectrum, where ample 154 photons for the pumping of the ERE are available. Furthermore, the carrier species involved in such a process must be abundant in the ISM in order to account for the contribution of the ERE carrier to the absorption in the ISM, amounting to about

10% of the photons of the Galactic interstellar radiation ﬁeld in the 540 – 91.2 nm wavelength range (Gordon et al., 1998). 155

16 I1 I2 15

14 H Ionization Opacity

12 Dust Extinction Opacity + H2 Self-shielding Opacity 11 Dust Extinction Opacity

10 Ionization Potential (eV)

6 0 100 200 300 400 500 600 Molecular Wt. (amu)

Figure 5-10 First ionization potentials of PAH molecules (filled circles) (Eilfeld & Schmidt, 1981) and their corresponding cations (open circles) (Leach, 1996) versus molecular weights. The three horizontal lines correspond to the cutoff of the far-UV radiation field due to H-ionization opacity (top), the onset of the energy regime in which dust extinction and H2 self-shielding opacity compete for photons (middle), and the lower limit on the energy for ERE initiation (lowest), respectively. 156

In terms of abundance and total cross section, interstellar PAHs meet these requirements. However, as shown in §§ 5.4.1, neutral PAHs cannot account for the observed ERE. By contrast, doubly ionized PAH molecules oﬀer interesting prospects.

In Figure 5-10 we have plotted the ﬁrst ionization potentials of neutral PAH molecules and of PAH mono-cations as a function of the size of the molecules as measured by their molecular weights. We note that most PAHs with up to 50 carbon atoms can be readily ionized into the mono-cation stage with photons with energies in the 6 eV to

8 eV range, and can subsequently be sent into the di-cation state with photons in the energy range from 10.5 to 13.6 eV. We should, therefore, expect to ﬁnd PAHs in the di-cation state in environments characterized by the presence of neutral or molecular hydrogen, as previously suggested by Leach (1987, 1996). PAH di-cations are closed- shell systems similar to neutral PAH molecules with S1 - S0 transitions occurring at optical wavelengths longer than the corresponding transitions in neutral molecules.

Fluorescence could thus be expected in the ERE range. The viability of such a model requires laboratory spectroscopy on PAH di-cations that could reveal the presence of optical ﬂuorescence, following the excitation through near-UV/optical absorptions.

No such experimental data exist at this time. 157

anthracene++ coronene++ pyrene++ terrylene++ ovalene++ optical absorption (a.u.) 0 2 4 6 8 10 12 14 photon energy (eV)

Figure 5-11 Model spectra of representative PAH di-cations.

We can, however, obtain some general idea of the expected spectroscopic prop-

erties of PAH di-cations, using density-functional theory modeling calculations (TD-

DFT) (Onida et al., 2002). This theory incorporates electronic screening and relevant correlation eﬀects (Xie et al., 2004) for vertical electronic excitations in the ground-

state geometry and thus represents a fully ab initio formalism for computing excited

states. Over the past decade, TD-DFT has given promising results for ﬁnite systems

(atoms, molecules, clusters) (Xie et al., 2004; Onida et al., 2002; Xie et al., 2005).

In this work, we use the TD-DFT implemented in the Gaussian 03 package (Frisch,

2003). Based on the ground-state structures for all clusters, which were optimized

by using DFT (Kohn & Sham, 1965) with the nonlocal hybrid B3LYP functional

(Becke, 1993), we compute the excited energies of both singlet and triplet states

(Lewis & Kasha, 1944), the oscillator strengths, and the optical absorption gap (gen- 158

erally defined by the energy of the first dipole-allowed transition for finite systems

(Xie et al., 2004)). To ensure the accuracy of our calculations, we choose 6-31G(d,p)

Pople-type basis set which adds polarization functions to the atoms (Xie et al., 2005).

In Figure 5-11 we show the composite absorption spectra of ﬁve PAH di-cations, in-

cluding results for pyrene++ and anthracene ++, whose likely presence in the Red

Rectangle (HD 44179) has been demonstrated by both their blue ﬂuorescence as

well as their spectroscopic signature of the ﬁrst ionization (Vijh et al., 2004, 2005b).

When compared with available experimental data for neutral and once-ionized PAH

species, these types of calculations agree remarkably well with experimental measure-

ment. The spectra in Figure 5-11 show that representative PAH di-cations absorb

strongly in the energy range from 2 eV to 7.5 eV, suggesting a characteristic band

gap of 2 eV. Photoluminescence across this band gap could be responsible for the

ERE. A still controversial issue is the question of the stability of twice-ionized PAH

molecules. A recent study of the fragmentation pathways of the di-cation of benzene

by Rosi et al. (2004) revealed an unexpectedly high degree of stability. This result suggests that the di-cations of multi-ringed PAHs could exhibit a still higher degree of stability than benzene and thus play a substantial role in interstellar chemistry and photophysics.

5.4.3 Consistency Check: The Red Rectangle

The two principal results emerging from the study of the excitation of the ERE in the NGC 7023 NW PDR are: (1) the ERE initiation requires photons with energies 159

> 10.5 eV, e.g. for the creation of the ERE carrier; (2) the ERE process is most likely a two-step process, in which the ﬁrst step creates the carrier via ionization or dissociation, requiring in excess of 10.5 eV, and the second step produces the actual excitation of the ERE carrier via longer-wavelength photons. In this section we will examine the Red Rectangle, one of the two environments where ERE is being produced with very high eﬃciency, the other being the high-|b| Galactic cirrus.

Two aspects are of greatest signiﬁcance, the morphology of ERE sources and the

ERE energetics. The Red Rectangle is a well-studied bi-polar protoplanetary nebula with a circum-stellar disk obscuring the central source (Men’shchikov et al., 2002;

Cohen et al., 2004). The central source consists of a L ∼ 6000 L AGB star with

a Teﬀ ∼ 8250 K atmosphere (Vijh et al., 2005b) and a low-mass (M ≈ 0.35 M )

hot white dwarf of unknown luminosity, which could amount to at most ∼ 100 L

(Men’shchikov et al., 2002; Driebe et al., 1998). This upper limit to the the white

dwarf’s luminosity would apply only, however, during the ∼ 105 yr duration of the

stellar core’s contraction from the AGB region of the Hertzsprung-Russell diagram to

the top of the white dwarf cooling sequence. Once the stellar core has reached a size

corresponding to the white dwarf mass-radius relation, a more typical luminosity of

a Teﬀ = 60,000 K hot white dwarf is about 6.5 L . The relevance of these luminosity

limits will become apparent presently.

Spatially, the ERE in the Red Rectangle is conﬁned to geometrically thin regions

within the walls of the bi-conical outﬂow cavities (Schmidt & Witt, 1991; Cohen et al.,

2004, Vijh et al., in prep.), leading to the appearance of the X-shaped bipolar struc-

ture, while scattering in the 500 to 350 nm wavelength range produces a spherical 160

reﬂection nebula devoid of the bipolar structure (Cohen et al., 2004). The conﬁne-

ment of the ERE to the walls of the outﬂow cavity is in agreement with our ﬁnding

of an initiation of the ERE by high-energy UV photons. These photons, stemming

predominantly from the hot white dwarf, are spatially conﬁned to the outﬂow cones

and its walls by the optically thick circum-binary disk and the steeply rising far-UV

opacity of the nebular dust in the Red Rectangle (Vijh et al., 2005b). At wavelengths

longward of 170 nm (E< 7.3 eV), the nebular opacity is greatly reduced. Thus, the mid-UV, near-UV and optical photons, which are produced almost entirely by the cooler but extremely luminous AGB star, are able to penetrate the walls of the out-

ﬂow cavity and produce the nearly spherical blue reﬂection nebula component of the

Red Rectangle (Cohen et al., 2004). They are also able to eﬃciently pump the ERE

carriers produced in the thin wall regions of the outﬂow cones.

The ERE energetics of the Red Rectangle is particularly useful in supporting

the concept of a two-step process of ERE excitation. Schmidt et al. (1980) re-

ported the band-integrated ERE luminosity of the entire Red Rectangle nebula as

2 LERE ≈ 0.6 (D/280 pc) L . Using the currently accepted distance of 710 pc

(Men’shchikov et al., 2002; Hobbs et al., 2004), we arrive at an estimate of LERE ≈ 3.9 L

for the Red Rectangle. This value is uncorrected for attenuation by the optically thick

circum-binary disk surrounding the stellar sources as well as for any interstellar ex-

tinction along the line of sight. Given the unusual color characteristics of the RR and

the non-isotropic morphology of the RR, these attenuations are diﬃcult to estimate.

The light of the central AGB star alone suﬀers an attenuation in excess of AV = 4

magnitudes (Vijh et al., 2005b). We, therefore, consider it a conservative estimate 161

that the total ERE in the RR suﬀers an attenuation by about 1 magnitude, resulting

in a total estimated ERE luminosity of LERE ≈ 10 L . If photons with E > 10.5 eV

were responsible for the actual excitation of the ERE in a single-step process, and if

the excitation process were 100% eﬃcient, with one ERE photon of 1.9 eV produced

for every absorption of a E> 10.5 eV photon, the total luminosity in E > 10.5 eV photons required for this excitation would be LFUV > 55 L . Thus, with somewhat

more realistic assumptions about the ERE excitation eﬃciency of possibly 30% to

50%, we would require well in excess of 100 L in far-UV photons to produce the ob-

served ERE luminosity of the Red Rectangle, well in excess of what is likely available

from the hot white dwarf companion.

If, by contrast, the far-UV photons emitted by the hot white dwarf companion of

the central AGB star are used solely for creating and maintaining the ERE emitters,

the AGB star (Teﬀ ≈ 8250 K; L ≈ 6000 L ) has more than enough energy to excite

LERE ≈ 10 L with mid-UV and optical photons. Assuming a Stokes shift of 1 eV

typical for photoluminescence, absorption of an excitation energy of less than 15 L

would suﬃce to generate LERE ≈ 10 L , assuming 100% eﬃciency. (Gordon et al.,

1998) have estimated the ERE photon conversion eﬃciency to be > 10%, which

would lead to an upper limit for the ERE generating luminosity of 150 L . This is

still only about 2.5% of the luminosity of the AGB star. Thus, a two-step process of

ERE excitation is entirely feasible, given the energetics of the Red Rectangle nebula

system. 162

5.4.4 Consistency Check: The High-|b| Galactic Cirrus

Next to the Red Rectangle, the high-|b| Galactic cirrus is another environment in which the intensity of the ERE is roughly equal to that of the dust-scattered light, for very similar reasons. First, in both instances the scattering geometry involves mainly large-angle scattering, which does not lead to high scattered-light intensities.

Second, the illuminating radiation ﬁeld in both cases contains both a strong far-UV component and a strong optical component (see Witt & Johnson, 1973, for a spectrum

of the interstellar radiation ﬁeld), a condition which appears to be associated with

highly eﬃcient production of ERE. The interstellar medium giving rise to the high-|b|

Galactic cirrus is optically thin in the far-UV, where the ERE is initiated, as well as at the ERE wavelength range itself. This, coupled with a reasonably accurate knowledge of the spectrum of the exciting interstellar radiation ﬁeld (Mathis et al.,

1983), allows a reliable estimate of the ERE intensity per hydrogen atom along an

average line of sight to be made, as well as an estimate of a lower limit of the photon

conversion eﬃciency of the ERE process (Gordon et al., 1998).

Gordon et al. (1998) found an ERE intensity of (1.43 ± 0.31) × 10−29 ergs s−1

A˚−1 sr−1 H-atom−1 for the Galactic cirrus at |b| > 20◦, independently conﬁrmed

by Szomoru & Guhathakurta (1998). With the well-established observed relation between hydrogen column density and extinction in the Galaxy (Diplas & Savage,

1994), Gordon et al. (1998) estimated that (10±3)% of all photons that are absorbed

from the interstellar radiation ﬁeld by interstellar dust must be absorbed by the ERE

carrier to produce the observed ERE intensity with a conversion eﬃciency of 100% 163

on a photon-per-photon basis. If, as our current results suggest, far-UV photons

are required to initiate the ERE, we must ﬁrst consider the far-UV portion of the

interstellar radiation ﬁeld. If we consume photons through absorption by the ERE

carrier, starting at the 91.2 nm limit set by the interstellar hydrogen opacity and

proceed toward longer wavelengths, we ﬁnd that we must use all photons shortward

of 130 nm in wavelength or with energies down to 9.5 eV. This limit is indicated by a

vertical dashed line in Figure 5-9, labeled 100%. We note that this limit violates the

ERE initiation limit of 10.5 eV set by the opacity constraint found in NGC 7023. If

we assume a more realistic conversion eﬃciency of 50% for the photons absorbed by

the ERE carriers, all photons absorbed between 91.2 nm and 185 nm in wavelength

would be required to generate the observed ERE intensity. This second limit is shown

in Figure 5-9 as a dashed vertical line labeled 50%.

Therefore, in close analogy to the case of the Red Rectangle discussed earlier, we

ﬁnd that that the interstellar radiation ﬁeld does not contain enough photons with

energies E > 10.5 eV to pump any ERE carriers sufficiently to generate the ERE intensities observed from the high-|b| Galactic cirrus. Again, the problem would meet with a simple solution, if the high-energy photons are used solely to produce and maintain a population of ERE carriers which can then be pumped by the far more abundant optical/near-UV photons of the interstellar radiation field. Thus, a two- step process of ERE excitation is also being suggested by the data available for the high-|b| diffuse interstellar medium. 164

5.5 Conclusions

We summarize our conclusions:

1. We have imaged sharp molecular cloud edges in the NW PDR of the reﬂection

nebula NGC 7023 in the light of the ERE band and in the light of emission from

the H2 1-0 S(1) transition, using the HST ACS and HST NICMOS instruments,

respectively. The cloud edges appear as narrow ﬁlaments with widths as small

as 0.300 in both bands, while the appearance of the same cloud edges is broader

and more diﬀuse in images showing dust-scattered radiation at Hα and in the

z-band.

2. In general, every H2 ﬁlament is matched by a corresponding ERE ﬁlament, while

the reverse correlation does not hold to the same degree. We interpret this one-

sided correlation as resulting from a density-dependent conversion eﬃciency

of far-UV photons into near-IR 1-0 S(1) photons. In lower-density regions,

absorption of far-UV photons by dust dominates over the absorption by H2

molecules, leading to ERE ﬁlaments without corresponding H2 ﬁlaments, while

in higher-density regions absorptions of far-UV photons by dust and by H2

molecules are comparable.

3. By comparing the widths of H2 ﬁlaments with the corresponding widths of

the cloud edges in the z-band, we estimated the wavelength-dependence of the

dust extinction in the molecular clouds comprising the NW PDR of NGC 7023.

Assuming that the extinction follows the RV -dependent relationship developed 165

by CCM, we found a value of RV ≈ 5.6.

4. By comparing the widths of the ERE ﬁlaments with the corresponding widths of

the cloud edges in the z-band, we found the ERE ﬁlaments on average about 10%

wider than the corresponding H2 ﬁlaments. This indicates that the dust opacity

encountered by photons that initiate the ERE is only slightly smaller than the

dust opacity encountered by photons exciting the H2 molecules. Given that the

H2 emission in NGC 7023 is excited by far-UV photons with λ> 110.4 nm, the

CCM-extinction curve of RV ≈ 5.6 identiﬁes two regions of the spectrum where

the dust opacity matches the required value, the range around the mid-UV

extinction bump (250 nm to 189 nm) and the region shortward of 118 nm.

5. We used the observation by Darbon et al. (1999) that stars with Teff ≤ 7000 K

do not excite ERE to eliminate the possibility that ERE is initiated by mid-

UV photons in the 250 nm to 186 nm range. This leaves far-UV photons at

wavelengths shortward of 118 nm (∼10.5 eV) as the ultimate source of ERE

excitation.

6. We showed that all existing ERE models fail in the face of this new requirement.

7. With the examples of two environments where the ERE excitation is highly ef-

ﬁcient and where the number of ERE photons has been determined, we demon-

strated that the number of far-UV photons with energies in excess of 10.5 eV

is insuﬃcient to generate the number of observed ERE photons, assuming rea-

sonable values for the photon conversion eﬃciencies. 166

8. We concluded that the ERE excitation must therefore be a two-step process.

The ﬁrst step requires E > 10.5 eV photons to produce the actual ERE carrier,

probably by a process of photo-ionization or photo-dissociation of a suitable

precursor. A modest far-UV ﬂux would then suﬃce to maintain a population of

ERE carriers thus created. The second step of the ERE excitation then requires

eﬃcient pumping of the ERE carrier by abundant optical/near-UV photons with

energies above the ERE band-gap of ≈ 2 eV.

9. A possible ERE carrier must therefore be a system that (1) has an ioniza-

tion potential in excess of 10.5 eV, that (2) exhibits strong absorption in

the optical/near-UV spectral region, and that (3) is capable of eﬃcient photo-

luminescence in the wavelength range of the ERE band.

10. We suggest that PAH dications with masses ≤∼ 500 amu appear to meet the

ﬁrst two of these three requirements. Laboratory experiments designed to study

the possibility of ﬂuorescence by PAH dications are needed to test their potential

as ERE sources. Chapter 6

Optical Emission Band

Morphologies of the Red

Rectangle1

6.1 Introduction

The optical spectrum of the Red Rectangle (RR) nebula consists of several com-

ponents of comparable intensities, and it is therefore possible to use narrow-band

imaging to isolate the morphological substructures where particular emissions are

dominant. i) Dust-scattered radiation: to ﬁrst order, this spectrum is a somewhat

bluer version of that of the more luminous star in the central binary system, HD 44179,

a Teﬀ ∼8250 K A-giant AGB star. ii) Blue luminescence (BL) band: this emission

1To be submitted to The Astrophysical Journal Letters, July 2005, Authors: U. P. Vijh, A. N.

Witt, D. G. York, V. V. Dwarkadas, B. E. Woodgate, & P. Palunas

167 168

band with a peak near 380 nm (FWHM ∼ 45 nm) has been attributed to ﬂuores-

cence by neutral PAH molecules with 14 - 18 C-atoms. (Vijh et al., 2004, 2005b)

iii) Extended red emission (ERE): this band peaks near 670 nm (FWHM ∼ 180 nm)

and it is the primary cause for the unusual red color of the RR. The carrier of the

ERE is yet unidentiﬁed (Witt & Vijh, 2004). iv) Sharp red emission features: this is a set of relatively narrow emission bands due to as yet unidentiﬁed molecules.

They appear in a wavelength range largely coincident with that of the ERE band

(Van Winckel et al., 2002). It has also been suggested that some of these bands are

emission counterparts of corresponding Diﬀuse Interstellar Bands (DIBs). Spectro-

scopic studies have demonstrated that the intensities of the emission bands ii), iii),

and iv) are comparable to those of the underlying scattered light continuum.. This

opens the possibility of constraining the physical state(s) of the respective emitters,

because the almost-edge-on geometry of the bi-polar outﬂow allows one to associate

speciﬁc physical conditions with regard to density, radiation ﬁeld, and ionization

equilibrium with speciﬁc nebular regions.

6.2 Observations and Reductions

The data were obtained at the Apache Point Observatory on January 7, 8, 9, 10

and 11 2005 using the Goddard Fabry Perot on the 3.5-m telescope. Images of the

RR were obtained using narrow-band ﬁlters centered at 3934 A˚ ( FWHM ∼ 27 A),˚

4050 A˚ (FWHM ∼ 57 A),˚ 4596 A˚ (FWHM ∼ 57 A),˚ 5700 A˚ (FWHM ∼ 135 A),˚

6400 A˚ (FWHM ∼ 135 A).˚ The etalon was used to obtain images centered at 5800 A˚ 169

and 5855 A,˚ each with a 15 A˚ eﬀective width. The nights were partially cloudy and

only relative ﬂuxes can be ascertained from these images.

Data reduction was carried out using standard IRAF tasks. All images were

trimmed to a size of 561 × 561 pixels or 3.04 × 3.04. The darks, twilight ﬂats and

object frames were corrected using an average zero frame. The ﬂats and object frames

were then dark-corrected using a combined dark frame. The ﬂat frames for each ﬁlter

were obtained by moving 3000 between successive exposures to avoid coherent addition

of bright stars. These ﬂat frames were then median-combined for each ﬁlter. Dark

corrected object frames were then flat fielded using the combined twilight flat for

each ﬁlter. Object exposures were taken in sets of ﬁve, in a pentagram pattern of

pointings. Object frames with the best seeing were selected and median-combined

using tasks imalign and imcombine. Table 6.1 shows the exposure times for each

ﬁlter. Stars in the frames were used to register and shift the images: this worked

for all ﬁlters except the 3934 A˚ ﬁlter. As no stars other than the central extended source were visible on these frames, the center of the extended source was used to

shift and co-add the images. This approach was shown to be correct, as one of the

brightest stars in the ﬁeld now appeared in the co-added image albeit with extremely

low counts. All images were further cropped to 200 × 200 pixel or 7300 × 7300 to best

show the RR. Images were scaled using counts of bright stars before computing the

ratio images discussed in the next section. 170

Table 6.1 Narrow-band observations of the RR nebula.

Filter Filter width (A)˚ Exposure (s) Comment

3934 27 5×900 BL band 4050 57 5×120, 1×180 Scattered 4596 57 5× 120 Scattered 5700 135 5× 120 Scattered + ERE 5800 15 25 ×40 Sharp emission feature 5855 15 17×40 Sharp emission feature 6400 135 4×60, 1×180 ERE + Scattered

6.3 Results and Discussion

6.3.1 Blue Luminescence in the Red Rectangle

The two narrow-band images at 3934 A˚ and 4050 A˚ show the appearance of the

RR in the wavelength region of the BL band. The former is much closer to the peak of the BL emission, while the latter covers the long-wavelength tail of the BL band and is therefore comparatively more dominated by the scattered light component of the spectrum. Figure 6-1 is an image of the RR with the 3934 filter superimposed with contours. The most remarkable feature of these images is the fact that the inner isophotes of the 3934 A˚ image are consistently elongated in the E-W direction, apparently coincident with the orientation of the disk surrounding HD 44179. A check on stellar images in the 3934 A˚ frames confirmed that this elongation is real and not the result of guiding errors in the individual 900 s exposures. This image confirms results from long-slit spectroscopy at CTIO (Vijh et al., 2005b) which showed that 171

the BL emission exhibits a strong preference for the parts of the RR that are located

in the shadow of the disk. This is consistent with the identiﬁcation of the BL as

originating from ﬂuorescence by small (3- to 4-ring) neutral PAH molecules. Such

molecules are photo-excited by stellar photons of ∼ 4 eV energy, which can diﬀuse

into the outer disk regions via scattering. The same molecules would be photo-ionized

by 7.5 eV photons, which are present in regions of the RR with direct lines-of-sight

to the central sources, e.g. the outﬂow cones. We believe, therefore, that we are

looking at emission from the outer portions of the optically thick disk surrounding

the central stars in the RR, i.e regions where the local radiation ﬁeld causes the

ionization equilibrium to be preferentially neutral. We note that the isophotes of the

4050 A˚ image, shown in Figure 6-2 are distinctly circular in the inner region of the

image, in contrast to the elongated 3934 A˚ isophotes. This reﬂects the fact that dust- scattered light is dominant here. The outer isophotes of the 4050 A˚ image become

almost square, with the sides being parallel to the polar axis of the RR. While the

BL image at 3934 A˚ shows slight intensity enhancements at the base of the outﬂow

walls on the W-side of the nebula, no such enhancements are apparent in the 4050 A˚

image, despite the fact that this is a deeper exposure. The dimensions of the RR

disk implied by the narrow-band imaging observations are in excellent agreement

with the results of Bujarrabal et al. (2003) who mapped this disk in CO emission with the IRAM interferometer array. Their map shows the disk having an angular extent of 900. When de-convolved with their beam pattern, the inferred disk diameter is 5.200. The 3934 A˚ data show elongated isophotes over an extent of 6.500, which is in perfect agreement, taking into account the ∼100 seeing disk of the optical data. 172

A perplexing aspect of the 3934 A˚ image is the orientation of the disk (PA ∼ 90◦) compared to the orientation of the polar axis of the RR (PA ∼ 11◦). If the disk were exactly perpendicular to the polar axis, a PA of 101◦ would have been expected.

However, HST images of the innermost portions of the circum-binary disk of the RR

(Cohen et al., 2004, Fig. 3) also seem to indicate that the disk is not perpendicular to the nebular outﬂow axis, and in fact that it is oriented in E-W direction. If conﬁrmed, we have independently determined the orientation of the RR disk. 173

Figure 6-1 Image at 3934 A˚ with overlaid contours showing elongated proﬁles along the disk tracing the BL distribution. 174

Figure 6-2 Image at 4050 A˚ with overlaid contours showing circular scattered light proﬁles.

6.3.2 Sharp Emission Feature at 5800 A˚

The 5700 A˚ image contains mainly dust-scattered nebular light and a contribution from ERE, mainly in the outlying regions. The narrow-band 5800 A˚ and 5855 A˚

images contain a major contribution from the 5797 A˚ and 5855 A˚ sharp emission

features, in addition to scattered light and a slightly larger contribution of ERE com- 175 pared to the 5704 A˚ image. The 5700 A˚ image is slightly overexposed, leading to bleeding from the stellar image and to the appearance of diﬀraction spikes that are not visible on the other images. The 5800/5700 and the 5855/5700 ratio images,

Figures 6-3 and 6-4, show that there is excess emission at 5800 A˚ and 5855 A˚ in two locations to either side of the central source, on a line approximately perpendicular to the nebular axis. The X-shaped dark markings on in the ratio images corresponding to the walls of the outflow cones are the result of a systematic displacement of the intensity maxima of the ERE and the sharp emission features. The narrow emission features reach peak intensities slightly inside the cones with respect to the ERE maxima (Schmidt & Witt, 1991). The rung-like structures inside the cone across the polar axis of the nebula also appear dark on these ratio images. This experiment should definitely be repeated with images of of comparable depth and neither overexposed at the center. With those conditions, it would be possible to derive significant quantitative information about the spatial variation of the 5800 A˚ band intensity. 176

Figure 6-3 5800 A˚ image divided by 5700 A˚ image. 177

Figure 6-4 5855 A˚ image divided by 5700 A˚ image.

6.3.3 ERE

The 6400 A˚ image is dominated by the ERE. When compared to images at wave-

lengths < 5000 A,˚ this image demonstrates that the X-shaped structures (variably

referred to as “arms” or “whiskers” in the RR literature) are a result of the ERE 178

band. We interpret these structures as the walls of the internally illuminated bi-polar

outﬂow cavities, seen in projection. When divided by the 5700 A˚ “continuum” image,

the resulting 6400/5700 ratio image, Figure 6-5 convincingly demonstrates that the

ERE is preferentially emitted by material in and within the walls of the outﬂow cavities, and is extremely weak outside these regions. This is demonstrated even more clearly in the ratio image 6403/4050, Figure 6-6. This is basically a display of the

ratio of ERE to dust-scattered radiation in the RR. We can interpret the relative

conﬁnement of the ERE to the outﬂow cavities and their walls in one of two ways:

either the carrier of the ERE exists only in the outﬂow cavities and its walls, or the

excitation requirements for the ERE restrict this emission process to regions which

can receive the required photons from the central sources. Both of these ways can

actually be combined in the light of the results of a recently completed investiga-

tion (Witt et al., 2005) of the ERE in the reﬂection nebula NGC 7023. This study

concluded that the ERE results from a two-step process:

Step 1 Creation of the ERE carrier. Ionization of the ERE carrier progenitor by

photons with energies 10.5 eV < E < 13.6 eV creating the actual ERE carrier.

This step can proceed only inside the outﬂow cavity and its walls with a direct

line of sight to the central stars, in particular to the hot white dwarf companion

of HD 44179. The latter object is the suspected source of the far-UV photons

in the RR.

Step 2 Pumping of the ERE carrier, followed by ﬂuorescence. Excitation of the car-

rier ion by optical/near-UV photons from the bright A-giant HD 44179 followed 179

by ﬂuorescence in the ERE band.

While the latter photons permeate other parts of the nebula as well, they can excite the ERE only in those places, where the earlier ionization has created the carriers. Neutral PAHs, such as those seen in the RR disk shadow will also absorb these optical/near-UV photons and respond with blue ﬂuorescence. If PAHs are also responsible for the ERE, the ionization potential constraint found in NGC 7023 suggests that PAHs are ionized to the di-cation stage before they luminesce in the

ERE band. Quantum-chemical calculations of the electronic energy level structure of

PAH di-cations show the presence of strong absorption bands shortward of ∼ 6000 A.˚

Whether or not PAH di-cations ﬂuoresce in the ERE band (> 6000 A)˚ is currently subject of laboratory investigations. 180

Figure 6-5 6400 A˚ image divided by 5700 A˚ image. 181

Figure 6-6 6400 A˚ image divided by 4050 A˚ image.

6.4 Conclusions

1. The morphology of the BL and ERE emissions in the RR nebula are almost

mutually exclusive.

2. As the total intensities from these two emissions when summed over the nebula 182

are of comparable magnitude, it is very suggestive to attribute them to diﬀerent

ionization stages of the same family of carriers such as PAH molecules.

3. We have presented the optical detection of the circum-binary disk of the RR in

the light of the BL.

4. The sharp emission features at 5800 A˚ and 5855 A˚ have spatial distributions

similar to the ERE, but they reach peak intensities inside the cones – in front

of the ERE and show excess emission in the disk.

Lastly, since the BL most likely arises in small neutral PAH molecules, the detection of

the disk in the light of the BL places in question earlier claims by Waters et al. (1998)

that the disk in the Red Rectangle is made up of oxygen-rich material produced in an

earlier mass-ejection epoch, while the bi-polar outﬂow lobes are carbon-rich matter

produced during the more recent outﬂow. Chapter 7

Spectral Characteristics of Blue

Luminescence in the Red

Rectangle1

We report the successful resolution of the spectrum of the recently discovered blue

luminescence in the Red Rectangle into a spectrum consisting of seven molecular

emission bands, superimposed upon a broad continuum. The relative intensity of

the component bands varies with position within the nebula, suggesting an origin in

a set of several related molecular species. The band widths and positions resemble

those of ﬂuorescence spectra of neutral aromatic molecules with between three and

seven carbon rings. The successful identiﬁcation of the carrier molecules will require

gas-phase ﬂuorescence spectra of a suitable range of aromatic molecules, which avoid

1Submitted to The Astrophysical Journal Letters June 2005, Authors: U. P. Vijh, A. N. Witt,

D. G. York, J. C. Barentine, R. McMillan, T. P. Snow, & D. .F. Malin

183 184

the solvent-dependent wavelength shifts aﬀecting the great bulk of currently available

laboratory spectra.

7.1 Introduction

The Red Rectangle (RR) is a well-known proto-planetary nebula in which the mass-losing stellar source, HD 44179, is surrounded by an optically thick disk, seen almost edge-on from Earth. Constrained by the massive disk, the mass outﬂow from the central post-AGB star has created a bi-polar nebula of exceptional symmetry

(Cohen et al., 2004). The presence of the disk prevents direct star light from reaching

the Earth while the nebular matter above and below the disk is being fully illuminated

by the stellar radiation. Given that dust-scattered light is subject to a strongly

forward-directed phase function, relatively little dust-scattered light from the RR is

reaching the Earth. On the other hand, nebular emissions excited by the stellar illumination are radiated isotropically and are therefore comparatively intense and easily observable. For this reason and because the carbon-rich RR environment is actively forming dust and molecules, the RR has been a most productive laboratory for the discovery of optical emissions from these nebular constituents. Prominent examples are the broad extended red emission (ERE) (Cohen, 1975; Schmidt et al.,

1980), the sharp molecular emission features near 580 nm (Schmidt & Witt, 1991;

Van Winckel et al., 2002), and, most recently, the blue luminescence (BL) reported

by Vijh et al. (2004, 2005b). Following the discovery of BL in the RR, Vijh et al.

(2005a) also detected BL in several other ordinary reﬂection nebulae, albeit with 185

much lower intensities. This suggests that the BL carrier is a ubiquitous component

of the interstellar medium and is not restricted to the particular environment of the

RR.

The identiﬁcation of the carrier of the BL is currently still a matter of controversy.

While Vijh et al. (2005b) have advanced several observational arguments which favor an explanation in terms of ﬂuorescence by small, neutral polycyclic aromatic hydrocarbon molecules, Nayfeh et al. (2005) have proposed small silicon nanoparticles as the carrier of the BL. The main reason why it is diﬃcult to decide among these two

(and possibly other) alternatives rests with the method with which the initial de-

tection of the BL spectrum was accomplished. The BL detection was achieved with

the Balmer line-depth technique (Vijh et al., 2004), which permits the determination

of the BL intensities in the presence of dominant dust-scattered light only at the

wavelengths of the Balmer absorption lines. This results in a poorly resolved spectral

profile with insufficient detail for a definitive identification.

In this Letter we are reporting on a new observational eﬀort to gain a more complete view of the BL spectrum in the RR and its variation with position in the nebula. We are taking advantage of the fact shown by Vijh et al. (2005b) that

the gradient of the dust-scattered light in the RR with oﬀset from the center is

much steeper than the gradient in the BL, a direct consequence of the forward-

directed phase function for the former and the isotropic emission pattern for the

latter. By obtaining deeply exposed nebular spectra at larger angular oﬀsets we

have succeeded for the ﬁrst time in obtaining well-resolved BL spectra without any signiﬁcant contribution from dust-scattered light. 186

In § 7.2 we will summarize the observations presented in this paper. The results

will be presented in § 7.3, where we will compare BL spectra obtained at diﬀerent

locations within the RR and where we will relate these results to a deep blue image of

the RR obtained with the 3.9-m Anglo-Australian Telescope to gain insight into the

spatial distribution of the BL-emitting molecules. In § 7.4 we will address the complex

issue of identiﬁcation of the BL carrier(s) and we will summarize our conclusions in

§ 7.5.

7.2 Observations and Reductions

The observations of the RR and HD 44179 were made at the Apache Point Obser-

vatory 3.5-m telescope, using the dual imaging spectrograph (DIS). Table 7.1 gives

details of the observations. The DIS is a medium dispersion double spectrograph

with two independent collimators and cameras for the blue and red sides. In this

letter we report on the results from the blue observations only. The high-blue grating

has 1200 lines mm−1 with a central wavelength of 440 nm and a linear dispersion of 0.62 A˚ pixel−1. The blue detector was a Marconi 2048 × 1024 back illuminated

CCD yielding a pixel scale of 0.4200 pixel−1. A slit 2.05 long and 100.5 wide was used.

All observations of the central star were carried out at the parallactic angle. Nebular observations with the slit passing through the star on Nov 26 were also carried out at the parallactic angle which translated to a mean PA of -24◦. On April 02, the nebula was observed with a slit at PA 90◦,ata200.5 southern oﬀset to obtain tie-in data with previous observations made at CTIO (reported in Vijh et al., 2004, 2005b)). 187

Table 7.1 RR Observations at APO

Date Object Slit Position Exposure (s)

Nov 26, ’04 RRa Par.angle 300×4 Apr02,’05 HD44179 Par.angle 2×120,2×60 Apr02,’05 RR 200.5 south, PA 90◦ 6×300

Data reductions were carried out with IRAF 2.12 EXPORT. Spectra of standard stars were obtained on both nights, but as the observing conditions were non- photometric, spectra are not calibrated to absolute flux, only to relative flux. In contrast, the CTIO observations, with which we will compare these data, were obtained under photometric conditions and at a much lower airmass. The Nov. observations were corrected for differential atmospheric refraction by using the stellar observation as the trace. At CTIO, the airmass was low and as the slit used was wider (200.5)

diﬀerential refraction was not an issue. These observations were used to correct the

data from the E-W slit for diﬀerential refraction to a limited degree. We extracted 20

spectra along each slit each covering 100.26 × 100.5 of the nebula. The long-slit enabled

simultaneous sky observation which was subtracted out. Figure 7-1 shows a high

resolution HST image of the nebula superimposed with the two slits and shows where

the extracted spectra shown in Figure 7-2 were obtained. 188

Figure 7-1 HST Red image with superimposed slit images at PA -24◦ and PA 90◦. The image is 3000 × 3000, north is to the top and east to the left. The boxes indicate apertures, spectra from which are shown in Figure 7-2. The sharp nebular structures in this image are primarily due to ERE. 189

7.3 Results

7.3.1 The Band Structure of the BL

In Figure 7-2, we illustrate the changing composition of the nebular light in theRR

with increasing oﬀset from the center by displaying three nebular spectra, extracted

from the PA -24 slit data at oﬀsets 1.8900 south, 4.4100 south, and 10.7100 south. These data cover a range in intensity of the nebular light of a factor of 17.5. The presence of strong Balmer absorption lines in the top spectrum illustrates the dominance of dust-scattered light originating in the AIII star located in the central cavity of the

RR disk. The second spectrum, down in intensity by a factor of about six, begins to show the increasing dominance of several broad emission bands with only a weak contribution from dust-scattered light, evidenced by the greatly reduced relative depth of the Balmer absorption lines. The lowest spectrum at an oﬀset of (9.7800 south, 4.3600

east) shows only a spectrum of emission bands superimposed upon a broad continuum

without any measurable contribution from scattered light. North of the star along

the slit, a similar progression of spectra are seen as well. 190

Figure 7-2 Nebular spectra at different offsets from the star along the PA -24 slit. These illustrate the changing composition of the nebular light in the RR with increasing offset from the center. Vertical bars at the bottom indicate the positions of the bands in the BL spectrum tabulated in Table 7.2.

In the last spectrum free of scattered light, we detect the presence of seven distinct bands. Their wavelengths and energies are listed in Table 7.2. We identify these seven

bands as the components of the BL. 191

Table 7.2 Molecular Bands in the BL spectrum of the Red Rectangle

Band# Peak Wavelength (nm) Energy (cm−1)

1 377.5 26,483.2 2 392.8 25,459.5 3 407.7 24,530.9 4 423.7 23,599.4 5 438.8 22,791.5 6 451.6 22,144.0 7 465.2 21,495.2

With the peak wavelengths of the BL bands known, we again examine the RR spectrum at 1.8900 south, and we notice that the presence of bands 3, 4, 6, and 7 is clearly recognizable by eye. The banded structure of the BL spectrum does, therefore, persist over almost the entire range of oﬀsets along the PA-24 slit.

7.3.2 Variation of BL Spectra with Position in the RR

The pure BL spectrum is observable only in the outer fringes of the RR; at positions closer to the nebular center we need to separate the scattered light component from the BL, which is possible with the present data only through the line-depth technique (Vijh et al., 2004). In order to obtain comparable data for all parts of the

RR, we evaluate the BL intensities at the wavelengths of the Balmer lines in those

spectra which are free of scattered light. In Figure 7-3 we display several BL spectra

obtained during the present observing program at APO as well as during the previous

observing program at CTIO (Vijh et al., 2004, 2005b). The spectra are normalized

at Balmer Hη (383.54 nm) to aid inter-comparison. Note that the spectral coverage 192

of the DIS at APO does not extend to as short a wavelength as the CTIO spectro-

graph. Principal uncertainties in the data stem from the uncertainty of placing the

continuum level for the line-depth measurements and from diﬀerential atmospheric

refraction, which aﬀected the APO spectrum at (2.500 south; 4.4100 west), and is corrected to only a limited degree. We will, therefore, only highlight certain qualitative trends apparent in these data.

Figure 7-3 Several BL spectra obtained during the present observing program at APO as well as during the previous observing program at CTIO (Vijh et al., 2004, 2005b). The spectra are normalized at Balmer Hη (383.54 nm) to aid inter-comparison. Ver- tical bars at the bottom indicate the positions of the bands in the BL spectrum tabulated in Table 7.2.

We notice high BL intensities at wavelengths shortward of Hη at small southern but large western oﬀsets, where the BL was initially detected. This suggest that the 193

emission in band #1 (see Tab. 7.2) is dominant in these locations. These positions are well outside the walls of the RR outﬂow cones, i.e. in the disk-shadow region.

All spectra appear to exhibit a shoulder or secondary maximum in the 390 - 420 nm wavelength range. There is also an increase in emission at longer wavelengths with increasing oﬀset. Given the band-nature of the underlying BL spectrum, we interpret these trends in terms of varying relative strengths of the seven BL bands. This may be interpreted as observational support for a model that attributes the BL bands not to a single molecule but rather to a set of two or more molecular carriers. A shift to increasing emission in the longer-wavelengths bands would be consistent with an increasing molecular size with distance from the center.

7.3.3 Variation of BL Intensities with Position in the RR

In Figure 7-4 we are comparing the band-integrated (360 nm - 480 nm) BL intensities along the PA-24 APO slit with the band-integrated intensities of dust-scattered radiation. For large oﬀsets, where there is very little scattered light we show intensities calculated in two ways: un-corrected for scattered light, by integrating the nebular spectra, and two: over-correcting for scattered light by assuming all the light at Hβ is scattered. As observed previously (Vijh et al., 2004, 2005b; Cohen et al.,

2004), the scattered-light halo of the RR dominates in the innermost parts of the nebula as a result of the strongly forward-directed phase function coupled with an embedded source. The BL varies with distance more gradually and extends to larger angular oﬀsets than the scattered light. 194

101 Band-BL: line-depths Band-BL: sc over-corrected Band-BL: sc un-corrected Band-sc ) -1 0 sr 10 -1 s -2

10-1 Intensities (erg cm

10-2 -15 -10 -5 0 5 10 15 offset (arcsec)

Figure 7-4 Band-integrated BL and the scattered light distribution along the PA -24 slit. The central star’s position is at 000 oﬀset, negative oﬀsets indicate positions in the southern part of the slit.

7.3.4 Morphology of BL in the RR

In Figure 7-5 we display a deep blue image of the RR taken by D.F. Malin with

the 3.9-m Anglo-Australian telescope. The eﬀective spectral range covered by the

plate/ﬁlter combination in this image is about 385 nm - 485 nm, and therefore includes

six of the brightest BL bands. The slit positions at which nebular spectra are shown

in Figure 7-3 are indicated as white boxes. The positions showing essentially pure

BL band spectra without contributions from scattered light are located outside the circular central halo. This makes it very likely that the faint, fuzzy nebular extension 195

surrounding the central scattered light halo derives its luminosity primarily from BL.

This extended distribution clearly lacks the sharply deﬁned structure visible in the

light of the ERE (Fig. 7-1). The BL and the ERE, while comparable in intensity, are

not correlated in their spatial distributions. The ERE is enhanced strongly on and

near the walls of the outﬂow cones of the RR and deﬁnes the distinctive X-shaped

appearance of the RR in red-wavelength images of the RR shown in Figure 7-1. We believe that the fluorescing molecules of the BL are physically located outside the mantles of the RR outflow cones where they are protected from direct illumination by the central post-AGB giant and its hot He-white dwarf companion. This would be consistent with the attribution of the BL to fluorescence by neutral molecules. 196

Figure 7-5 Deep blue image of the RR taken by D.F. Malin with the 3.9-m Anglo- Australian telescope. The image is 3000 × 3000, north is to the top and east to the left. White boxes indicate slit positions at which nebular spectra (shown in Fig. 7-3) have been extracted. 197

7.4 Discussion

7.4.1 Identiﬁcation of the BL Carrier

The presence of seven distinct BL bands with well-determined wavelengths should provide adequate constraints for the identiﬁcation of the BL carrier(s). The complex band structure of the BL spectrum appears to clearly eliminate silicon nanoparticle luminescence (Nayfeh et al., 2005) as a potential explanation of the BL. Instead, lack-

ing laboratory ﬂuorescence spectra of other potential candidates (e.g. carbon chains),

the observed bands resemble typical ﬂuorescence bands observed in PAH molecules,

and the wavelength range occupied by them would indicate molecules with three to

seven carbon rings (Vijh et al., 2004). However, the variation of the BL spectrum with location in the Red Rectangle strongly suggests that several molecules must contribute to the observed spectrum. Furthermore, the molecules in the RR nebula are most likely in the gas phase. Large existing data bases of ﬂuorescence spectra of organic molecules, including PAHs (Berlman, 1965; Karcher, 1985), include only

spectra obtained with the ﬂuorescing molecules in solution. It is well known that

the interaction of a solvent with the luminescing molecules introduces substantial

solvent-dependent and molecule-dependent red shifts (e.g. 10 nm to 19 nm for an-

thracene in cyclohexane) of the bands as well as changes in their relative intensities,

as again recently demonstrated by Chi et al. (2001a,b) and Lakowicz (1999). It is,

therefore, not surprising that a detailed comparison of our BL band spectrum with

existing solvent-based ﬂuorescence spectra of single molecular species (Berlman, 1965;

Karcher, 1985) did not produce a satisfactory match despite the qualitative similarity 198 of many of the spectra. On the other hand, it is interesting to note that fluorescence spectroscopy of complex aromatic mixtures (Apicella et al., 2004) including between ten and twenty molecular species in cyclohexane solution still produces spectra with between five and seven distinct bands of comparable intensity in the spectral range covered by our observations. This provides some confidence that spectral synthesis using as input the gas-phase fluorescence spectra of between ten and twenty aromatic species with between three and seven rings will produce a laboratory spectrum that will match our observed BL spectrum in the Red Rectangle.

There are usually one or two dominant bands in a given molecule, with other bands being of greatly reduced intensity. As an example, in Figure 7-6 we show a ﬂuorescence spectrum of anthracene dissolved in cyclohexane, excited at 350 nm

(Du et al., 1998) 2. While there are five bands observed, only the first two are of high intensity. By contrast, the seven bands of the BL spectrum are of comparable intensity. This, coupled with the observations showing the relative intensity of the different bands varying with position throughout the RR nebula, suggests to us that the origin of the BL spectrum is associated with more than one molecule.

2Data available at (http://omlc.ogi.edu/spectra/PhotochemCAD/html/anthracene.html) 199

2.5e+06 Anthracence in cyclohexane

2.0e+06

1.5e+06

1.0e+06 Fluorecence Intensity (a.u.)

5.0e+05

0.0e+00 3400 3600 3800 4000 4200 4400 4600 4800 5000 wavelength (Å)

Figure 7-6 Fluorescence spectrum of anthracene dissolved in cyclohexane, excited at 3500 A˚ (Du et al., 1998).

7.4.2 Temperature of the BL Carrier

Chi et al. (2001a,b) have reported gas-phase ﬂuorescence spectra of of PAH molecules

with 14 to 20 carbon atoms for temperatures ranging from 400 K to 920 K. While

discrete bands are easily distinguished at the lower temperatures, spectra obtained

at T > 620 K no longer show distinguishable vibrational structure. The BL bands

observed in the outer parts of the RR show very prominent bands. Even though it

is not really possible to designate temperatures to individual molecules and given

that molecules in the ISM are stochastically heated, the band spectrum implies that

the BL carrier molecules at large oﬀsets are at temperatures considerably lower than 200

620 K.

7.5 Conclusions

The spectrum of the BL in the RR has now been resolved into seven component emission bands. This spectrum resembles that of fluorescence spectra of small polycyclic aromatic hydrocarbon molecules with between three and seven rings. No single known molecule matches the observed band spectrum, in part because almost the entire existing data base of fluorescence spectra have been obtained with molecules in solution, which introduces shifts in the wavelengths of bands as well as changes in their relative intensities. We suggest that several different molecular species contribute to the spectrum. The spectrum clearly is not that of the PL of silicon nanoparticles.

The relative strength of the component bands of the BL varies with position. In the disk, the band at 377.5 nm appears to be the dominant band. At larger oﬀsets the band at 407.7 nm as well as the bands at 423.7 nm and 438.8 nm gain in relative strength, while the 377.5 nm band weakens. A data base of gas-phase ﬂuorescence spectra of aromatic molecules with between three and seven aromatic rings will be needed before a satisfactory match to the observed spectrum can be achieved. Chapter 8

Summary

8.1 What did we learn?

Despite the importance of interstellar grains in many astrophysical processes, their nature, i.e. their size distribution, composition, structure , as well as their total mass, is not well known. Interstellar nanoparticles form the smallest and most numerous of these interstellar grains. In this work we have progressed further on this quest to gain a better understanding of the processes and composition of interstellar grains.

A new emission band of blue luminescence (BL) was discovered in the Red Rect- angle (RR) nebula (see Chapter 2). This enigmatic proto-planetary nebula is also one

of the brightest known sources of extended red emission (ERE) as well as of aromatic

emission features (AEFs). We attribute the BL to be ﬂuorescence from small neutral

polycyclic aromatic hydrocarbon (PAH) molecules. PAH molecules are thought to be

widely present in many interstellar and circumstellar environments in our galaxy as

well as in other galaxies, and are considered likely carriers of the AEFs in the near and

201 202

mid- infrared. However, no speciﬁc PAH molecule has yet been identiﬁed in a source

outside the solar system, as the set of mid-infra-red emission features attributed to

these molecules between the wavelengths of 3.3 µm and 16.4 µm is largely insensitive to molecular sizes. In contrast, near-UV/blue fluorescence of PAHs is more specific as to size, structure, and charge state of a PAH molecule. The spectrum of the BL suggests that the carriers are most likely neutral PAH molecules consisting of 3-4 aromatic rings such as anthracene (C14H10) and pyrene (C16H10). These small PAHs would then be the largest molecules specifically identified in the interstellar medium.

Following our initial discovery of blue luminescence in the spectrum of the RR and its identiﬁcation as ﬂuorescence by small three- to four-ringed PAH molecules, we illustrated the spatial correlation between the blue luminescence and the 3.3 µm emission, commonly attributed to small, neutral PAH molecules (see Chapter 3).

We contrast the excellent spatial correlation of the two former emissions with the

distinctly diﬀerent spatial distribution of the ERE and of the dust-scattered light

within the RR. By combining our observations with existing IUE data for HD 44179,

we have determined the UV/optical attenuation curve for the central source in the

RR. The attenuation curve is characterized by an exceptionally high value of RV

= AV/E(B−V ) = 11.9, a remarkable absence of the 217.5 nm absorption bump, the latter being a characteristic of Galactic interstellar extinction, and a strong, discontinuous rise in attenuation near λ−1 ≈ 6.0 µm−1. These unusual characteristics

suggest strongly that the reddening observed in the light from HD 44179 is a result

of radiative transfer in the RR and not due to interstellar extinction. The far-UV

rise in the attenuation curve of HD 44179 is qualitatively and quantitatively diﬀerent 203

from the much more gradual far-UV rise in Galactic extinction curves. The rise is,

however, consistent with the onset of photo-ionization of small PAH molecules with

masses of less than 250 amu, if the abundance of such PAH molecules relative to

hydrogen is about 10−5. This represents a fully independent conﬁrmation of the pres-

ence of small PAH molecules in the RR environment, consistent with our conclusion

regarding the origin of the BL in the RR. The broad mid-UV hump in the attenuation

curve of HD 44179, whose appearance is quite unlike that of the familiar 217.5 nm

interstellar absorption band, can be explained satisfactorily as resulting from a su-

perposition of the (typically) four mid-UV absorption bands characteristic of neutral

PAH molecules. An observable consequence of this absorption and the resulting elec-

tronic excitation is de-excitation via ﬂuorescence, which we indeed observe. Thus we

now have strong additional evidence for the presence of small PAH molecules with

masses of less than 250 amu in the RR, which supports the attribution of the blue

luminescence to ﬂuorescence by the same molecules

BL, which was ﬁrst discovered in the proto-planetary nebula RR around the post-

AGB star HD 44179, was thought to be unique to the RR environment where such

small molecules are actively being produced and shielded from the harsh interstellar

radiation by a dense circumstellar disk. However, the BL spectrum was detected in

several ordinary reﬂection nebulae illuminated by stars having temperatures between

10,000 – 23,000 K (see Chapter 4). All these nebulae are also known to exhibit the infrared emission features called aromatic emission features (AEFs) attributed to large

PAHs. We also present the spatial distribution of the BL in these nebulae, and in

the case of Ced 112, the BL is spatially correlated with mid-IR emission structures 204

attributed to AEFs. These observations provide evidence for grain processing and

possibly for in-situ formation of small grains from larger aggregates. Most impor-

tantly, the detection of BL in these ordinary reﬂection nebulae suggests that the BL

carrier is an ubiquitous component of the ISM and is not restricted to the particular

environment of the RR.

Though the carrier of the dust-associated photoluminescence process causing the

ERE remains unidentiﬁed, we have constrained the character of the ERE carrier further by determining the wavelengths of the radiation that initiates the ERE (see

Chapter 5). Using the imaging capabilities of the Hubble Space Telescope, we have

resolved the width of narrow ERE ﬁlaments appearing on the surfaces of externally

illuminated molecular clouds in the bright reﬂection nebula NGC 7023 and compared

them with the depth of penetration of radiation of known wavelengths into the same

cloud surfaces. We identify photons with wavelengths shortward of 118 nm as the

source of ERE initiation. There are strong indications from the well-studied ERE

in the Red Rectangle nebula and in the high-|b| Galactic cirrus that the photon ﬂux with wavelengths shortward of 118 nm is too small to actually excite the observed

ERE, even with 100% quantum eﬃciency. We conclude therefore, that ERE excitation results from a two-step process. The ﬁrst step, involving far-UV photons with

E > 10.5 eV, leads to the creation of the ERE carrier, most likely through photo- ionization or photo-dissociation of an existing precursor. The second step, involving more abundant near-UV/optical photons, consists of the optical pumping of the previously created carrier, followed by subsequent de-excitation via photoluminescence.

The latter process can occur many times for a single particle, depending upon the 205

lifetime of the ERE carrier in its active state. While none of the previously pro-

posed ERE models can match these new constraints, we note that under interstellar

conditions most PAH molecules are ionized to the di-cation stage by photons with

E > 10.5 eV and that the electronic energy level structure of PAH di-cations is consis-

tent with ﬂuorescence in the wavelength band of the ERE. Therefore, PAH di-cations

deserve further study as potential carriers of the ERE.

The narrow-band imaging study of the RR (see Chapter 6) opens the possibility

of constraining the physical state(s) of the respective emitters, because the almost-

edge-on geometry of the bi-polar outﬂow allows one to associate speciﬁc physical

conditions with regard to density, radiation ﬁeld, and ionization equilibrium with

speciﬁc nebular regions. We have presented the optical detection of the circum-

binary disk of the RR in the light of the BL. The morphology of the BL and ERE

emissions in the RR nebula are almost mutually exclusive. As the total intensities

from these two emissions when summed over the nebula are of comparable magnitude,

it is very suggestive to attribute them to diﬀerent ionization stages of the same family

of carriers such as PAH molecules.

Finally, we present the spectrum of the BL in the RR which has now been resolved into seven component emission bands (see Chapter 7). This spectrum resembles

that of ﬂuorescence spectra of small polycyclic aromatic hydrocarbon molecules with

between three and seven rings. No single known molecule matches the observed

band spectrum, in part because almost the entire existing data base of ﬂuorescence

spectra have been obtained with molecules in solution, which introduces shifts in the

wavelengths of bands as well as changes in their relative intensities. We suggest that 206

several diﬀerent molecular species contribute to the spectrum. The spectrum clearly

is not that of the PL of silicon nanoparticles. The relative strength of the component

bands of the BL varies with position. In the disk, the band at 377.5 nm appears to

be the dominant band. At larger oﬀsets the band at 407.7 nm as well as the bands at

423.7 nm and 438.8 nm gain in relative strength, while the 377.5 nm band weakens. A

data base of gas-phase ﬂuorescence spectra of aromatic molecules with between three

and seven aromatic rings will be needed before a satisfactory match to the observed

spectrum can be achieved.

8.2 Future Work

The immediate need is for laboratory data on gas-phase ﬂuorescence spectroscopy on neutral and ionized PAHs and also other possible carriers of the optical photoluminescence bands to supplement the rich observational data. The data and analysis presented in Chapters 6 and 7 are the ﬁrst results from an ongoing project to obtain

deep images and a complete spectroscopic map of the RR with the aim to completely

characterize the nebular emission from this enigmatic object. What other revelations

and surprises this nebula will present is left to the imagination of the reader. Appendix A

Determination of BL Intensity at

the Balmer Jump

Assume that the ratio of the scattered light intensity at λ less than the Balmer

sc? sc? jump (BJ) to the scattered light intensity at λ greater than the BJ, IBJ is

determined by the stellar spectrum. We measure λBJ at 397 nm.

Assuming a constant albedo over the wavelength range and an optically thin scenario,

the relative increase in the scattered ﬂux at 357 nm compared to that at 397 nm is

1.08 (Savage & Mathis, 1979). This factor is an upper limit as it will be smaller in

an optically thick case. Thus, in the stellar spectrum

sc? I357 sc? =0.2334 I397

207 208

After correction for wavelength dependent scattering,

sc? I357 sc? =0.2520 I397

=⇒

scN scN I357 =0.2520 ∗ I397

where the superscript N denotes the values for the nebular spectra. Also,

totalN BL scN I357 = I357 + I357

BL scN = I357 + I397 ∗ 0.2520

Therefore, the intensity of the BL at 357 nm,

BL totalN scN I357 = I357 − I397 ∗ 0.2520

which can be determined by knowing the value of the scattered light intensity at

397 nm, which in turn is previously determined by subtracting the BL intensity

(measured using the line-depth technique) at that wavelength from the total intensity

at that wavelength. References

Allain, T., Leach, S., & Sedlmayr, E. 1996a, A&A, 305, 602

—. 1996b, A&A, 305, 616

Allamandola, L. J., Tielens, A. G. G. M., & Barker, J. R. 1985, ApJ, 290, L25

Aller, L. 1956, Gaseous Nebulae (Wiley, New York)

Amans, D., Callard, S., Gagnaire, A., Joseph, J., Ledoux, G., & Huisken, F. 2003,

Journal of Applied Physics, 93, 4173

An, J. H., & Sellgren, K. 2003, ApJ, 599, 312

Apicella, B., Ciajolo, A., & Tregrossi, A. 2004, Anal. Chem., 76, 2138

Bakes, E. L. O., Bauschlicher, C., & Tielens, A. G. G. M. 2004, in ASP Conf. Proc.

309: Astrophysics of Dust, ed. A. N. Witt, G. C. Clayton, & B. T. Draine (San

Francisco: ASP), 731

Bakes, E. L. O., & Tielens, A. G. G. M. 1994, ApJ, 427, 822

Becke, A. D. 1993, J. Chem. Phys., 98, 5648

209 210

Belomoin, G., Therrien, J., & Nayfeh, M. 2000, Applied Physics Letters, 77, 779

Berlman, I. B. 1965, Handbook of Fluorescence Spectra of Aromatic Molecules (New

York: Academic Press)

Biennier, L., Salama, F., Allamandola, L. J., & Scherer, J. J. 2003, J. Chem. Phys.,

118, 7863

Birks, J. B. 1970, Photophysics of Aromatic Molecules (Wiley, London)

—. 1973, Organic Molecular Photophysics, Vol. 1 (Wiley, New York), 55

Bohlin, R. C., Savage, B. D., & Drake, J. F. 1978, ApJ, 224, 132

Br´echignac, P., & Hermine, P. 1994, in AIP Conf. Proc. 312: Molecules and Grains

in Space, ed. I. Nenner (New York: American Institute of Physics Press.), 613

Bregman, J. D., Rank, D., Temi, P., Hudgins, D., & Kay, L. 1993, ApJ, 411, 794

Bujarrabal, V., Neri, R., Alcolea, J., & Kahane, C. 2003, A&A, 409, 573

Burns, J., & Yelke, E. 1969, Rev. Sci. Instrum., 40, 1236

Cardelli, J. A., Clayton, G. C., & Mathis, J. S. 1989, ApJ, 345, 245

Casey, S. C. 1991, ApJ, 371, 183

Cesarsky, D., Lequeux, J., Ryter, C., & G´erin, M. 2000, A&A, 354, L87

Cherchneﬀ, I., & Barker, J. R. 1992, ApJ, 394, 703

Cherchneﬀ, I., Barker, J. R., & Tielens, A. G. G. M. 1992, ApJ, 401, 269 211

Chi, Z., Cullum, B. M., Stokes, D. L., Mobley, J., Miller, G. H., Hajaligol, M. R., &

Vo-Dinh, T. 2001a, Spectrochemica Acta A, 57, 1377

—. 2001b, Fuel, 80, 1819

Chlewicki, G., & Laureijs, R. J. 1987, in NATO ASIC Proc. 191: Polycyclic Aromatic

Hydrocarbons and Astrophysics, ed. A. Leger, L. d’Hendecourt, & N. Boccara, 335

Clar, E. 1972, The Aromatic Sextet (London: John Wiley & Sons)

Clar, E., & Schmidt, W. 1975, Tetrahedron, 31, 2263

—. 1977, Tetrahedron, 33, 2093

—. 1979, Tetrahedron, 35, 1027

Cohen, M., Van Winckel, H., Bond, H. E., & Gull, T. R. 2004, AJ, 127, 2362

Cohen, M. e. a. 1975, ApJ, 196, 179

Cook, D. J., & Saykally, R. J. 1998, ApJ, 493, 793

Crawford, M. K., Tielens, A. G. G. M., & Allamandola, L. J. 1985, ApJ, 293, L45

Darbon, S., Perrin, J.-M., & Sivan, J.-P. 1998, A&A, 333, 264

—. 1999, A&A, 348, 990

Darbon, S., Zavagno, A., Perrin, J.-M., Savine, C., Ducci, V., & Sivan, J.-P. 2000,

A&A, 364, 723 de Boer, K. S. 1983, A&A, 125, 258 212

Desert, F.-X., Boulanger, F., & Puget, J. L. 1990, A&A, 237, 215

d’Hendecourt, L. B., Leger, A., Olofsson, G., & Schmidt, W. 1986, A&A, 170, 91

Diplas, A., & Savage, B. D. 1994, ApJ, 427, 274

Draine, B. T. 2004, in ASP Conf. Proc. 309: Astrophysics of Dust, ed. A. N. Witt,

G. C. Clayton, & B. T. Draine (San Francisco: ASP), 691

Draine, B. T., & Bertoldi, F. 1996, ApJ, 468, 269

Driebe, T., Schoenberner, D., Bloecker, T., & Herwig, F. 1998, A&A, 339, 123

Du, H., Fuh, R. A., Li, J., A. Corkan, A., & Lindsey, J. S. 1998, Photochemistry and

Photobiology, 68, 141

Duley, W. W. 1985, MNRAS, 215, 259

—. 2001, ApJ, 553, 575

Duley, W. W., Hartquist, T. W., Sternberg, A., Wagenblast, R., & Williams, D. A.

1992, MNRAS, 255, 463

Duley, W. W., & Seahra, S. 1998, ApJ, 507, 874

Duley, W. W., & Williams, D. A. 1981, MNRAS, 196, 269

—. 1990, MNRAS, 247, 647

Dwek, E., Zubko, V., Arendt, R. G., & Smith, R. K. 2004, in ASP Conf. Proc. 309:

Astrophysics of Dust, ed. A. N. Witt, G. C. Clayton, & B. T. Draine, 499 213

Eilfeld, P., & Schmidt, W. 1981, J. Electron Spectrosc. Rel. Phen., 351, 101

Ewald, M., & Garrigues, P. 1985, Spectral Atlas of Polycyclic Aromatic Compounds,

ed. W. Karcher (Dordrecht, Boston, Lancaster: D. Reidel), 34

—. 1988, Spectral Atlas of Polycyclic Aromatic Compounds, ed. W. Karcher (Dor-

drecht, Boston, Lancaster: D. Reidel), 58

Field, D., Gerin, M., Leach, S., Lemaire, J. L., Pineau Des Forets, G., Rostas, F.,

Rouan, D., & Simons, D. 1994, A&A, 286, 909

Field, D., Lemaire, J. L., Pineau des Forets, G., Gerin, M., Leach, S., Rostas, F., &

Rouan, D. 1998, A&A, 333, 280

Fitzpatrick, E. L. 2004, in ASP Conf. Proc. 309: Astrophysics of Dust, ed. A. N.

Witt, G. C. Clayton, & B. T. Draine (San Francisco: ASP), 33

Frisch, M.J., e. a. 2003, , Gaussian 03 (Gaussian, Inc., Pittsburgh PA)

Fuente, A., Martin-Pintado, J., & Gaume, R. 1995, ApJ, 442, L33

Fuente, A., Martin-Pintado, J., Neri, R., Rogers, C., & Moriarty-Schieven, G. 1996,

A&A, 310, 286

Fuke, K., Tsukamoto, K., Misaizu, F., & Sanekata, M. 1993, J. Chem. Phys., 99, 7807

Furton, D. G., Laiho, J. W., & Witt, A. N. 1999, ApJ, 526, 752

Furton, D. G., & Witt, A. N. 1990, ApJ, 364, L45

—. 1992, ApJ, 386, 587 214

Geballe, T. R., Lacy, J. H., Persson, S. E., McGregor, P. J., & Soifer, B. T. 1985,

ApJ, 292, 500

Gordon, K. D., Witt, A. N., & Friedmann, B. C. 1998, ApJ, 498, 522

Gordon, K. D., Witt, A. N., Rudy, R. J., Puetter, R. C., Lynch, D. K., Mazuk, S.,

Misselt, K. A., Clayton, G. C., & Smith, T. L. 2000, ApJ, 544, 859

Gorodetskii, D. I., & Rozhkovskii, D. A. 1978, Luminosity of diﬀerent types of dif-

fuse nebula in the near-IR region Integrated photometry of NGC 2023 (Physical

Processes in Gas-Dust Nebulae), 52

Greenstein, J. L., & Oke, J. B. 1977, PASP, 89, 131

Guhathakurta, P., & Tyson, J. A. 1989, ApJ, 346, 773

Herbig, G. H. 2000, ApJ, 542, 334

Hobbs, L. M., Thorburn, J. A., Oka, T., Barentine, J., Snow, T. P., & York, D. G.

2004, ApJ, 615, 947

Hony, S., Van Kerckhoven, C., Peeters, E., Tielens, A. G. G. M., Hudgins, D. M., &

Allamandola, L. J. 2001, A&A, 370, 1030

Hoyle, F., & Wickramasinghe, N. C. 1999, Ap&SS, 268, 321

Huang, X., Xu, J., Li, W., & Chen, K. 2002, Thin Solid Films, 422, 130

Iwayama, T. S., Hama, T., Hole, D. E., & Boyd, I. W. 2002, Surf. & Coatings Tech.,

712, 158 215

Jochims, H. W., Baumg¨artel, H., & Leach, S. 1996, A&A, 314, 1003

—. 1999, ApJ, 512, 500

Jochims, H. W., Ruhl, E., Baumg¨artel, H., Tobita, S., & Leach, S. 1994, ApJ, 420,

307

Jones, A. P., Duley, W. W., & Williams, D. A. 1990, QJRAS, 31, 567

Jura, M., Turner, J., & Balm, S. P. 1997, ApJ, 474, 741

Karcher, W., ed. 1985, Spectral Atlas of Polycyclic Aromatic Compounds: Includ-

ing Data on Occurrence and Biological Activity, Vol. 1 – 3 (Dordrecht: Kluver

Academic for the Commission of the European Communities)

Keller, R. 1987, in NATO ASIC Proc. 191: Polycyclic Aromatic Hydrocarbons and

Astrophysics, ed. A. Leger, L. d’Hendecourt, & N. Boccara, 387

Kemper, C., Spaans, M., Jansen, D. J., Hogerheijde, M. R., van Dishoeck, E. F., &

Tielens, A. G. G. M. 1999, ApJ, 515, 649

Kerr, T. H., Hurst, M. E., Miles, J. R., & Sarre, P. J. 1999, MNRAS, 303, 446

Kim, D., Lee, M., Suh, Y. D., & Kim, S. K. 1992, J. Am. Chem. Soc., 114, 4429

Kohn, W., & Sham, L. J. 1965, Physical Review, 140, 1133

Koike, C., Chihara, H., Koike, K., Nakagawa, M., Okada, M., Tsuchiyama, A., Aok,

M., Awata, T., & Atobe, K. 2002, Meteoritics and Planetary Science, 37, 1591 216

Kurucz, R. 1993, ATLAS9 Stellar Atmosphere Programs and 2 km/s grid. Kurucz

CD-ROM No. 13. Cambridge, Mass.: Smithsonian Astrophysical Observatory,

1993., 13

Kurucz, R. L., Peytremann, E., & Avrett, E. H. 1974, Blanketed model atmospheres

for early-type stars (Washington : Smithsonian Institution : for sale by the Supt. of

Docs., U.S. Govt. Print. Oﬀ., 1974.), 37

Lakowicz, J. R. 1999, Principles of Fluorescence Spectroscopy, 2nd edn. (New York:

Kluwer Academic/Plenum Publications)

Le Page, V., Snow, T. P., & Bierbaum, V. M. 2003, ApJ, 584, 316

Leach, S. 1987, in NATO ASIC Proc. 191: Polycyclic Aromatic Hydrocarbons and

Astrophysics, ed. A. Leger, L. d’Hendecourt, & N. Boccara, 99

Leach, S. 1995, Planet. Space Sci., 43, 1153

Leach, S. 1996, Zeitschr. f. Physik. Chem., 195, 15

Ledoux, G., Ehbrecht, M., Guillois, O., Huisken, F., Kohn, B., Laguna, M. A., Nen-

ner, I., Paillard, V., Papoular, R., Porterat, D., & Reynaud, C. 1998, A&A, 333,

L39

Ledoux, G., Gong, J., Huisken, F., Guillois, O., & Reynaud, C. 2002, Applied Physics

Letters, 80, 4834

Ledoux, G., Guillois, O., Huisken, F., Kohn, B., Porterat, D., & Reynaud, C. 2001,

A&A, 377, 707 217

Ledoux, G., Guillois, O., Porterat, D., Reynaud, C., Huisken, F., Kohn, B., & Pail-

lard, V. 2000, Phys. Rev. B, 62, 15942

Leger, A., & d’Hendecourt, L. 1987, in NATO ASIC Proc. 191: Polycyclic Aromatic

Hydrocarbons and Astrophysics, ed. A. Leger, L. d’Hendecourt, & N. Boccara

(Dordrecht: D. Reidel Publishing Co.), 223

Leger, A., d’Hendecourt, L., & Boissel, P. 1988, Physical Review Letters, 60, 921

Leger, A., & Puget, J. L. 1984, A&A, 137, L5

Lemaire, J. L., Field, D., Gerin, M., Leach, S., Pineau des Forets, G., Rostas, F., &

Rouan, D. 1996, A&A, 308, 895

Lewis, G. N., & Kasha, M. 1944, J. Am. Chem. Soc., 66, 2100

Li, A., & Draine, B. T. 2002, ApJ, 564, 803

Li, G., Burggraf, L. W., Shoemaker, J. R., Eastwood, D., & Stiegman, A. E. 2000,

Applied Physics Letters, 76, 3373

Liao, M., Feng, Z., Chai, C., Yang, S., Liu, Z., & Wang, Z. 2002, Journal of Applied

Physics, 91, 1891

Malin, D. F., & Zealey, W. J. 1979, S&T, 57, 354

Mathis, J. S., Mezger, P. G., & Panagia, N. 1983, A&A, 128, 212

Mattila, K. 1979, A&A, 78, 253 218

Mavi, H. S., Rasheed, B. G., Shukla, A. K., Soni, R. K., & Abhi, S. C. 2003, Mat.

Sci. and Engg. B, 97, 239

Men’shchikov, A. B., Balega, Y. Y., Osterbart, G., & Weigelt, G. 1998, New Astron-

omy, 3, 601

Men’shchikov, A. B., Schertl, D., Tuthill, P. G., Weigelt, G., & Yungelson, L. R. 2002,

A&A, 393, 867

Moutou, C., Verstraete, L., Sellgren, K., & Leger, A. 1999, in ESA SP-427: The

Universe as Seen by ISO, ed. P. C. M. F. Kessler, 727

Nayfeh, M. H., Habbal, S. R., & Rao, S. 2005, ApJ, 621, L121

Nirmal, M., & Brus, L. 1999, Acc. Chem. Res., 32, 407

Onaka, T. 2004, in ASP Conf. Proc. 309: Astrophysics of Dust, ed. A. N. Witt, G. C.

Clayton, & B. T. Draine (San Francisco: ASP), 163

Onida, G., Reining, G. L., & Rubio, A. 2002, Rev. Mod. Phys., 74, 601

Papoular, R., Conard, J., Guillois, O., Nenner, I., Reynaud, C., & Rouzaud, J.-N.

1996, A&A, 315, 222

Patrone, L., Nelson, D., Safarov, V. I., Sentis, M., Marine, W., & Giorgio, S. 2000,

Journal of Applied Physics, 87, 3829

Peaden, P. A., Lee, M. L., Hirata, Y., & Novotny, M. 1980, Anal. Chem., 52, 2268 219

Peeters, E., Allamandola, L. J., Hudgins, D. M., Hony, S., & Tielens, A. G. G. M.

2004, in ASP Conf. Proc. 309: Astrophysics of Dust, ed. A. N. Witt, G. C. Clayton,

& B. T. Draine (San Francisco: ASP), 141

Pendleton, Y. J. 2004, in Astronomical Society of the Paciﬁc Conference Series, 573

Pendleton, Y. J., & Allamandola, L. J. 2002, ApJS, 138, 75

Perkins, H. G., Scarrott, S. M., Murdin, P., & Bingham, R. G. 1981, MNRAS, 196,

635

Perrin, J.-M., Darbon, S., & Sivan, J.-P. 1995, A&A, 304, L21

Perrin, J.-M., & Sivan, J.-P. 1992, A&A, 255, 271

Pierini, D., Majeed, A., Boroson, T. A., & Witt, A. N. 2002, ApJ, 569, 184

Racine, R. 1968, AJ, 73, 233

Rapacioli, M., Joblin, C., & Boissel, P. 2005, A&A, 429, 193

Reese, M. D., & Sitko, M. L. 1996, ApJ, 467, L105

Reipurth, B., & Zinnecker, H. 1993, A&A, 278, 81

Reyl´e, C., & Br´echignac, P. 2000, European Physical Journal D, 8, 205

Robertson, J. 1996, Phys. Rev. B, 53, 16302

Robinson, M. S., Beegle, L. W., & Wdowiak, T. J. 1997, ApJ, 474, 474

Rosi, M., Bauschlicher, C. W., & Bakes, E. L. O. 2004, ApJ, 609, 1192 220

Rush, W. F., & Witt, A. N. 1975, AJ, 80, 31

Rusli, R., Robertson, J., & Amaratunga, G. A. J. 1996, Journal of Applied Physics,

80, 2998

Russell, R. W., Soifer, B. T., & Willner, S. P. 1978, ApJ, 220, 568

Rydgren, A. E. 1980, AJ, 85, 444

Sakata, A., Wada, S., Narisawa, T., Asano, Y., Iijima, Y., Onaka, T., & Tokunaga,

A. T. 1992, ApJ, 393, L83

Sakata, A., Wada, S., Okutsu, Y., Shintani, H., & Nakada, Y. 1983, Nature, 301, 493

Savage, B. D., & Mathis, J. S. 1979, ARA&A, 17, 73

Schmidt, G. D., Cohen, M., & Margon, B. 1980, ApJ, 239, L133

Schmidt, G. D., & Witt, A. N. 1991, ApJ, 383, 698

Schmidt, W. 1977, J. Chem. Phys., 66, 828

Schutte, W. A., Tielens, A. G. G. M., & Allamandola, L. J. 1993, ApJ, 415, 397

Scott, A., & Duley, W. W. 1996, ApJ, 472, L123

Scott, A., Duley, W. W., & Pinho, G. P. 1997, ApJ, 489, L193

Seahra, S. S., & Duley, W. W. 1999, ApJ, 520, 719

Sellgren, K. 1984, ApJ, 277, 623 221

Sellgren, K., Allamandola, L. J., Bregman, J. D., Werner, M. W., & Wooden, D. H.

1985, ApJ, 299, 416

Sgro, L. A., Minutolo, P., Basile, G., & DAlessio,´ A. 2001, Chemosphere, 42, 761

Siebenmorgen, R., Natta, A., Kruegel, E., & Prusti, T. 1998, A&A, 339, 134

Smith, T. L., Clayton, G. C., & Valencic, L. 2004, AJ, 128, 357

Smith, T. L., & Witt, A. N. 2002, ApJ, 565, 304

Snow, T. 1995, Observational Constraints on the Carriers of the Diﬀuse Bands, ed.

A. G. G. M. Tielens & T. P. Snow (Dordrecht: Kluwer Academic Publishers), 325

Snow, T. P., & Seab, C. G. 1989, A&A, 213, 291

Snow, T. P., & Witt, A. N. 1995, Science, 270, 1455

Stockburger, M. 1973, Organic Molecular Photophysics, Vol. 1 (Wiley, New York),

Struve, O., & Swings, P. 1948, PASP, 60, 61

Sukhorukov, O. e. a. 2004, Chem. Phys. Lett., 386, 259

Szomoru, A., & Guhathakurta, P. 1998, ApJ, 494, L93

Takami, M., Usuda, T., Sugai, H., Kawabata, H., Suto, H., & Tanaka, M. 2000, ApJ,

529, 268

Taylor, S. D., & Williams, D. A. 1993, MNRAS, 260, 280 222

van den Ancker, M. E., The, P. S., Tjin A Djie, H. R. E., Catala, C., de Winter, D.,

Blondel, P. F. C., & Waters, L. B. F. M. 1997, A&A, 324, L33 van Dokkum, P. G. 2001, PASP, 113, 1420 van Hill, H., Loren, R., & Davis, J. 1975, ApJ, 198, 235 van Winckel, H. 2003, ARA&A, 41, 391

Van Winckel, H., Cohen, M., & Gull, T. R. 2002, A&A, 390, 147

Van Winckel, H., Waelkens, C., & Waters, L. B. F. M. 1995, A&A, 293, L25

Verstraete, L., & Leger, A. 1992, A&A, 266, 513

Verstraete, L., Leger, A., d’Hendecourt, L., Defourneau, D., & Dutuit, O. 1990, A&A,

237, 436

Verstraete, L., Pech, C., Moutou, C., Sellgren, K., Wright, C. M., Giard, M., L´eger,

A., Timmermann, R., & Drapatz, S. 2001, A&A, 372, 981

Vieira, S. L. A., Corradi, W. J. B., Alencar, S. H. P., Mendes, L. T. S., Torres,

C. A. O., Quast, G. R., Guimar˜aes, M. M., & da Silva, L. 2003, AJ, 126, 2971

Vijh, U. P., Witt, A. N., & Gordon, K. D. 2004, ApJ, 606, L65

—. 2005a, ApJ, 633, in press

—. 2005b, ApJ, 619, 368

Waelkens, C., Van Winckel, H., Trams, N. R., & Waters, L. B. F. M. 1992, A&A,

256, L15 223

Waelkens, C., Van Winckel, H., Waters, L. B. F. M., & Bakker, E. J. 1996, A&A,

314, L17

Watanabe, I., Hasegawa, S., & Kurata, Y. 1982, Jap. J. Appl. Phys., 21, 856

Waters, L. B. F. M., Cami, J., de Jong, T., Molster, F. J., van Loon, J. T., Bouwman,

J., de Koter, A., Waelkens, C., Van Winckel, H., Morris, P. W., Cami, J., Bouwman,

J., de Koter, A., de Jong, T., & de Graauw, T. 1998, Nature, 391, 868

Watkin, S., Gledhill, T. M., & Scarrott, S. M. 1991, MNRAS, 252, 229

Wdowiak, T. J., Donn, B., Nuth, J. A., Chappelle, E., & Moore, M. 1989, ApJ, 336,

838

Webster, A. 1993, MNRAS, 264, L1+

Weingartner, J. C., & Draine, B. T. 2001, ApJS, 134, 263

White, G. J. 1993, A&A, 274, L33

Whittet, D. C. B., ed. 2003, Dust in the galactic environment

Whittet, D. C. B., Prusti, T., Franco, G. A. P., Gerakines, P. A., Kilkenny, D.,

Larson, K. A., & Wesselius, P. R. 1997, A&A, 327, 1194

Williams, P. M., Brand, P. W. J. L., Longmore, A. J., & Hawarden, T. G. 1977,

MNRAS, 180, 709

Witt, A. N. 1985, ApJ, 294, 216

Witt, A. N., Bohlin, R. C., & Stecher, T. P. 1984a, ApJ, 279, 698 224

Witt, A. N., & Boroson, T. A. 1990, ApJ, 355, 182

Witt, A. N., Gordon, K. D., & Furton, D. G. 1998, ApJ, 501, L111

Witt, A. N., Gordon, K. D., Vijh, U. P., Sell, P. H., Smith, T. L., & Xie, R.-H. 2005,

ApJ, submitted

Witt, A. N., Graﬀ, S. M., Bohlin, R. C., & Stecher, T. P. 1987, ApJ, 321, 912

Witt, A. N., & Johnson, M. W. 1973, ApJ, 181, 363

Witt, A. N., & Malin, D. F. 1989, ApJ, 347, L25

Witt, A. N., & Schild, R. E. 1985, ApJ, 294, 225

—. 1986, ApJS, 62, 839

—. 1988, ApJ, 325, 837

Witt, A. N., Schild, R. E., & Kraiman, J. B. 1984b, ApJ, 281, 708

Witt, A. N., & Vijh, U. P. 2004, in ASP Conf. Proc. 309: Astrophysics of Dust, ed.

A. N. Witt, G. C. Clayton, & B. T. Draine (San Francisco: ASP), 115

Wolkin, M. V., Jorne, J., Fauchet, P. M., Allan, G., & Delerue, C. 1999, Physical

Review Letters, 82, 197

Xie, R., Bryant, G. W., Sun, G., Kar, T., Chen, Z., Smith, V. H., Araki, Y., Tag-

matarchis, N., Shinohara, H., & Ito, O. 2004, Phys. Rev. B, 69, 201403

Xie, R., Bryant, G. W., Zhao, J., Kar, T., & Smith, V. H. 2005, Phys. Rev. B, 71,

125422 225

Yoﬀe, A. D. 2001, Advances in Physics, 50, 1

Zhou, Z., Brus, L., & Friesner, R. 2003, Nano Lett, 3, 163

Zubko, V., Dwek, E., & Arendt, R. G. 2004, ApJS, 152, 211

Zubko, V. G., Smith, T. L., & Witt, A. N. 1999, ApJ, 511, L57