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UNIVERSITA’ DEGLI STUDI DI MILANO-BICOCCA Scuola di Dottorato di Scienze Corso di Dottorato di Ricerca in Fisica e Astronomia XVIII ciclo UNIVERSITE´ DE PROVENCE AIX-MARSEILLE I Ecole Doctorale ”Physique et Sciences de la Mati`ere” Doctorat en Rayonnement et Plasmas A.A.2004-2005

ENVIRONMENTAL EFFECTS ON GALAXY EVOLUTION IN NEARBY CLUSTERS

Coordinatore del Dottorato: Prof. Claudio Destri Directeur de l’Ecole´ Doctorale: Prof. Jean-Jacques Aubert Tutore: Prof. Giuseppe Gavazzi Directeur de th`ese: Dott. Alessandro Boselli

Commissione-Jury: Dott. A. Boselli (Laboratoire d’Astrophysique de Marseille) Prof. V. Buat (Universit´e de Provence) Prof. G. Gavazzi (Universit`a di Milano - Bicocca) Prof. F. Haardt (Universit`a dell’Insubria) Rapporteurs: Prof. C. Balkowski (Observatoire Astronomique Paris-Meudon) Dott. B. Poggianti (Osservatorio Astronomico di Padova)

Tesi di Dottorato di: Luca Cortese Matricola R00280

”Objectivity cannot be equated with mental blankness; rather, objectivity resides in recognizing your preferences and then subjecting them to especially harsh scrutiny ...and also in a willingness to revise or abandon your theories when the tests fail (as they usually do).” Stephen Jay Gould

Acknowledgments

This work represents the end point of my student career. After approximately twenty one years from my first entrance in a class room (it was September 1984 in Phoenix, AZ), I’m finally going to attend my last ”school” examination. Therefore I want to seize this opportunity in order to briefly remember and to thank some of the friends met during this journey. First of all I must thank my advisor Peppo Gavazzi, my scientific father, for his precious guidance and his teachings especially at the beginning of my research carrier. Special thanks to Alessandro Boselli, my co-advisor, first of all for the the last year spent in Marseille: a splendid experience. Thanks also for all his helpful advices, comments and supports on this and other works during the last three years. Many people contributed, directly or indirectly, to this work, and I am grateful to all of them. Merci beaucoup to Samuel Boissier for all the interesting discussions and, above all, for his precious lessons of French. Thanks to Veronique Buat for his help during the year spent in Marseille and for having initiated me in the obscure secrets of dust. Thanks to Barry Madore for his hospitality at the Carnegie Observatories, for his kindness, support and, especially, for his help in improving my written English. Muchas gracias to Armando Gil de Paz for his precious help on making the GALEX data available to me: without his contribution a great part of this work would not have been possible. Many thanks to Bianca Poggianti for a careful reading of my thesis and for her useful comments and suggestions. I would like also to thank Monica Colpi for her scientific and, especially, financial support during these three years. Arigato to Tsutomu Takeuchi and Akio Inoue for useful discussions about dust and galaxy evolution, for their kindness and help during my stay in Marseille and for having introduced me to Japanese cuisine. Many friends made the last three years unique. At Milano University life wouldn’t have been as much fun without all Peppo’s students. In particular thanks to Ilaria, Lea and Paolo for their unique support and thanks also to Chri for having installed Linux on my laptop, making me able to write this work. In Marseille thanks a lot to all the ”Caf´e du Coin”: Helene, Claude, Kassem, Peter, Fabrice and the others. Thanks for all the coffees and cakes, and for having received me with open arms even if I wasn’t able to speak French. Thanks to Alexie, Jean-

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Baptiste, Hector who shared the office with me, and a special thanks to Celine for having borne my never ending phone calls with my advisors, for her kindness and for her precious help in understanding french bureaucracy. Life in Marseille would have been completely different without the volley matches with Raph, Patrick, Seb, Mika, Fabrice and all the others. Finally, nothing of this would have been possible without the constant support of my parents and my brother Claudio, who have always encouraged me to continue this beautiful adventure.

This research was partly supported the Universit`a Italo-Francese through the Vinci Programme and by the CNES through GALEX-Marseille. Abstract

The environmental effects on galaxy evolution in nearby clusters are investigated using a multiwavelength dataset. The present analysis is focused on the properties of three (Abell 1367, Virgo and Coma) among the best studied clusters in the local Universe. Due to the variety of their environmental conditions (e.g. spiral fraction, X-ray luminosity, evolutionary stage) they represent the most suitable ”laboratory” for comparative studies. By combining for the first time GALEX UV observations with optical, near and far infrared data, the evolutionary history of cluster galaxies is studied. The main goals of this thesis are: (a) The study of the dependence of the UV emission of galaxies from their morphological type, mass and the environment they inhabit, through the study of UV luminosity functions and color magnitude relations. (b) The study of UV dust extinction properties of local cluster galaxies and investigation of possible empirical relations useful to estimate the amount of UV attenuation in local and high redshift galaxies. (c) Investigation of the effect of large scale structures assembling on galaxy evolution through the dynamical analysis of Abell 1367, one of the best examples of a dynamically young local cluster of galaxies. (d) The characterization of the effects of different environmental mechanisms (i.e. gravitation interactions, ram pressure, preprocessing) on the evolutionary history of cluster galaxies in order to gain more insight on the origin of the morphology-density and star-formation-density relations. The observational evidences presented in this work suggest that: (I) Giant ellipticals are an old, homogeneous population showing no or little evolution at least in the past 8 Gyr; unlike dwarf ellipticals which still contain young stellar populations. (II) The importance of different environmental mechanisms has changed during the age of the Universe. Tidal interactions and preprocessing probably dominated the past Universe and shaped part of the morphology-density relation during the phase of cluster accretion of small groups. Ram pressure dominates in today clusters and is surely affecting the star formation history of galaxies but with less influence on their morphology. (III) The heterogeneous class of S0s galaxies, from bulge dominated to the disky S0s, is not the result of a single transformation mechanism: if ram pressure is able to produce disk dominated S0s, tidal interactions (and thus preprocessing) are required to account for bulge dominated S0s. (VI) Different observational evidences

vii viii confirm the presence of a correlation between the mean age of stellar populations and galaxy mass (downsizing effect). In the framework of the hierarchical model of galaxy formation, the origin of the downsizing effect remains unsolved. This clear observational evidences represents one of today’s main challenge for models of galaxy evolution. Riassunto

In questo lavoro vengono analizzati gli effetti dell’ambiente sull’evoluzione delle galassie, utilizzando una base di dati multi-lunghezza d’onda. In particolare tutta quest’analisi ´e focalizzata sullo studio di tre differenti ammassi di galassie dell’Universo Locale: Abell1367, Virgo, Coma. Questi tre ammassi sono tra i piu´ studiati nell’Universo locale e, date le loro differenti propriet´a (e.g. frazione di galassie a spirale, luminosit´a X, stadio evolutivo), rappresentano dei laboratori ideali per quantificare l’influenza dell’ambiente sull’evoluzione delle galassie. Combinando per la prima volta osser- vazioni ultraviolette del GALEX a dati ottici, in vicino e lontano infrarosso viene ricostruita l’evoluzione delle galassie d’ammasso. I principali obiettivi di questa tesi sono: (a) Studiare il legame tra le propriet´a dell’emissione UV delle galassie, il loro tipo morfologico, la loro massa e l’ambiente in cui esse si trovano, attraverso l’analisi delle funzioni di luminosit´a UV e delle relazioni colore-magnitudine. (b) Comprendere le propriet´a delle polveri interstellari respon- sabili dell’assorbimento della radiazione ultravioletta e ricavare relazioni empiriche utili per poter quantificare l’assorbimento della radiazione ultravioletta in assenza di osservazioni in lontano infrarosso. (c) Analizzare se e come lo stato dinamico di un ammasso ´e in grado di influenzare la storia evolutiva delle galassie, attraverso lo studio dell’ammasso di Abell1367: uno dei migliori esempi di ammasso locale, dinami- camente ancora giovane. (d) Quantificare l’influenza di diversi effetti d’ambiente (i.e. interazioni gravitazionali, ram-pressure, galaxy preprocessing) sull’evoluzione delle galassie d’ammasso, in modo da comprendere le origini del fenomeno di segregazione morfologica. Tutte le evidenze osservative presentate e analizzate in questo lavoro suggeriscono che: (I) Le ellittiche giganti rappresentano una popolazione vecchia, omogenea che non ha subito una significativa evoluzione negli ultimi 8 Gyr; al contrario dell’ellittiche nane che sono ancora oggi dominate da popolazioni stellari giovani. (II) L’influenza dell’ambiente sull’evoluzione delle galassie cambia sensibilmente con l’et´a dell’Universo. Le interazioni gravitazionali ed il galaxy preprocessing sono stati gli effetti dominanti nell’Universo passato e sembrano essere i responsabili, almeno in parte, del fenomeno di segregazione morfologica. La ram pressure sembra essere dominante negli ammassi di oggi. Questo meccanismo ´e sicuramente in grado di influenzare la storia di for-

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mazione stellare delle galassie, ma ha pochi effetti sulla loro morfologia. (III) Le galassie lenticolari (S0) risultano essere cos´ı il prodotto di processi completamente differenti: se oggi la ram pressure ´e in grado di trasformare una galassia a spirale in una lenticolare con piccolo bulge, sono necessarie interazioni gravitazionali per pro- durre i grandi bulge osservati in molte lenticolari nell’Universo locale. (IV) Diverse, e indipendenti, evidenze osservative confermano l’esistenza di una forte correlazione tra la massa degli oggetti e l’et´a media della loro popolazione stellare (downsizing effect). Nel quadro dei modelli gerarchici di formazione delle strutture, l’origine del downsiz- ing effect ´e tutt’ora sconosciuta. La comprensione di questo fenomeno rappresenta dunque una delle maggiori sfide per l’astronomia extragalattica. R´esum´e

Ce travaille est d´edi´e a` l’´etude des effets d’environnement sur l’´evolution des galaxies dans l’Univers voisin, en utilisant un ´echantillon multi-longueur d’onde. En partic- ulier toute cette analyse est focalis´ee sur les propri´et´es des trois diff´erents amas des galaxies: Abell1367, Virgo et Coma. Ces trois amas des galaxies sont parmi les mieux ´etudi´es dans l’Univers local et, en raison de la vari´et´e des leurs propri´et´es (par ex- emple fraction des galaxies a` spirale, luminosit´e X, ´etat dynamique), ils repr´esentent des laboratoires, les plus appropri´es, pour des ´etudes comparatives. En combinant pour la premi`ere fois des observations UV de GALEX a` des donn´ees en optique, en voisin et en lointain infrarouge j’ai d´etermin´e l’histoire evolutive des galaxies dans les amas. Les buts principales de cette th`ese sont: (a) Etudier´ la variation des propri´et´es UV des galaxies en fonction des propri´et´es de l’environnement ou elles se trouvent, de leur masse et type morphologique, en analysant les fonctions de luminosit´e en UV et les relations couleur-magnitude. (b) L’´etude du taux d’absorption des photons UV par les poussi´eres interstellaires, pour obtenir des relations empiriques tr´es utils pour quantifier l’attenuation par poussi´eres quand les donn´ees en infrarouge lointain sont absent´ees. (c) Analyser l’effet de la formation des amas sur l’´evolution des galaxies en ´etudiant l’amas d’Abell1367, un des meilleurs exemples d’amas voisin et dynamique- ment jeune. (d) Comprendre l’influence des diff´erents effets d’environnement sur l’histoire evolutive des galaxies d’amas, pour comprendre l’origine de la s´egr´egation morphologique dans les amas. Touts les r´esultats obtenus dans ce travaille montrent que: (I) La population des galaxies elliptiques g´eants est vieille et homog`ene. Elle ne montre pas d’´evolution au moins dans les derni´eres 8 Gyr; au contraire des elliptiques naines qui contien- nent toujours populations stellaires jeunes. (II) L’importance relative des diff´erents m´ecanismes d’environment varie avec l’ˆage de l’Univers. Les interactions de mar´ee et le prepocessing ont probablement domin´ees dans l’Univers pass´e et ont contribu´ees (au moins en partie) a` la s´egr´egation morphologique, pendant la formation des amas par des petits groupes des galaxies. La pression dynamique domine dans les amas d’aujourd’hui et elle affecte suremenˆ t l’histoire de formation des etoiles des galaxies avec moins d’influence sur leur morphologie. (III) La classe h´et´erog`ene des galaxies S0s (lenticuliers), n’est pas le r´esultat d’un seul m´ecanisme de transformation: si la

xi xii pression dynamique peut produire S0s domin´ees par le disque, les interactions de mar´ee (et le preprocessing) sont exig´ees pour expliquer les S0s domin´ees par le bulbe. (IV) Diff´erentes ´evidences sugg`erent la pr´esence d’une corr´elation entre l’ˆage moyen des populations stellaires et la masse des galaxies (downsizing effect). Dans le cadre du mod`ele hi´erarchique de formation des galaxies, l’origine de cet effet n’est pas en- core r´esolue. Il repr´esente aujourd’hui une des d´efis pour les mod`eles d’´evolution des galaxies. Contents

1 Introduction 5

2 GALEX & GOLDMINE: A multiwavelength window on the Local Universe 15 2.1 The Galaxy Evolution Explorer ...... 15 2.1.1 The Prime Mission ...... 16 2.1.2 Data collection mode ...... 17 2.1.3 Counts vs. magnitudes and fluxes conversions ...... 18 2.2 The Galaxy On Line Database Milano Network ...... 20

3 The FAUST-FOCA UV luminosity function of nearby clusters 23 3.1 Introduction ...... 23 3.2 The Data ...... 24 3.3 The UV luminosity functions ...... 25 3.3.1 The composite cluster luminosity function ...... 26 3.4 Discussion ...... 28

4 GALEX UV luminosity function of Abell1367 33 4.1 Introduction ...... 33 4.2 UV data ...... 33 4.3 The luminosity function ...... 37 4.4 Discussion ...... 40

5 Multiple merging in Abell1367 43 5.1 Introduction ...... 43 5.2 Observations and data reduction ...... 44 5.3 The global velocity distribution ...... 46 5.4 Localized velocity structures ...... 48 5.5 The cluster dynamics ...... 51 5.5.1 The North-West subcluster ...... 53 5.5.2 The South-East subcluster ...... 56

1 2 CONTENTS

5.6 Star formation activity in the infalling groups ...... 59 5.7 Cluster mass ...... 61 5.8 Two-Body Analysis ...... 62 5.9 Conclusions ...... 65

6 Unveiling the evolution of early type galaxies with GALEX. 69 6.1 Introduction ...... 69 6.2 Data ...... 70 6.3 The UV properties of early-type galaxies ...... 71 6.4 Discussion and conclusion ...... 75

7 UV dust attenuation in normal star forming galaxies 81 7.1 Introduction ...... 81 7.2 The Data ...... 84 7.2.1 The optically-selected sample ...... 84 7.2.2 The starburst sample ...... 86 7.3 The L /L β relation for normal star-forming galaxies . . . . 86 T IR F UV − 7.3.1 The dependence on the birthrate parameter ...... 89 7.4 A(Hα) ...... 90 7.4.1 Estimate of A(Hα) ...... 90 7.4.2 The β-A(Hα) relation ...... 92 7.5 Relations between dust attenuation and global properties...... 94 7.5.1 Metallicity ...... 94 7.5.2 Luminosity ...... 97 7.5.3 Surface brightness ...... 99

7.5.4 LHα/LF UV ratio ...... 101 7.6 A cookbook for determining LT IR/LF UV ratio ...... 103

8 High velocity interaction: NGC4438 in the Virgo cluster 107 8.1 Introduction ...... 107 8.2 Data ...... 108 8.3 The UV emission and the star formation history of NGC 4438 . . . . 110 8.4 Discussion and conclusion ...... 113

9 Ram Pressure stripping: NGC4569 in the Virgo cluster 117 9.1 Introduction ...... 117 9.2 Data and models ...... 118 9.3 The star formation history of NGC 4569: model predictions . . . . . 120 9.4 Discussion and conclusion ...... 121 CONTENTS 3

10 Galaxy Pre-processing: the blue group infalling in Abell1367 127 10.1 Introduction ...... 127 10.2 Observations ...... 128 10.2.1 HI observations ...... 128 10.2.2 UV to near-IR imaging ...... 131 10.2.3 Hα imaging ...... 131 10.2.4 MOS spectroscopy ...... 131 10.2.5 High Resolution spectroscopy ...... 134 10.3 Results ...... 135 10.3.1 Kinematics ...... 135 10.3.2 Hα properties ...... 138 10.3.3 HI properties ...... 143 10.3.4 The fate of the stripped gas ...... 148 10.3.5 The metal content ...... 149 10.3.6 Dating the starburst...... 151 10.3.7 CGCG97-120: simply a foreground galaxy, or an high velocity intruder? ...... 156 10.4 Discussion ...... 157 10.4.1 The evolutionary history of the Blue Infalling Group . . . . . 157 10.4.2 The contribution of preprocessing to cluster galaxies evolution. 158

11 Discussion & Conclusions 165 11.1 Discussion ...... 165 11.2 Conclusions ...... 175

A The extinction correction 179

B Estimate of the < 912A˚ flux from Hα + [NII] 181

Bibliography 183

Chapter 1

Introduction

Eighty-five years are as short as a jiffy compared to the whole history of humanity and science, but this is the brief time men needed to upset their view of the Universe they inhabit. Let us return for a moment at the beginning of this story: July 26, 1920, Harlow Shapley and Herber Curtis confront their positions on the size of the Universe and the nature of the spiral nebulae in talks later called the Great Debate (see Trimble 1995, for a review). Curtis argued that the Universe is composed of many galaxies like our own, which had been identified by astronomers of his time as spiral nebulae. Shapley argued that these spiral nebulae were just nearby gas clouds, and that the Universe was composed of only one big Galaxy: our Milky Way. The resolution of the debate came in the mid 1920’s. Using the 100 inch telescope at Mount Wilson, Edwin Hubble identified Cepheid variable stars in the Andromeda Galaxy (M31). These stars resulted far beyond the most distant stars known in our galaxy and allowed Hubble (1925) to show that M31 was a galaxy much like our own. With this discovery, the known universe expanded immensely and, in the same time, a new research area was born: extragalactic astronomy. Thanks to overwhelming technological progress, during its first 85 years of life, ∼ extragalactic astronomy has provided us with a detailed description of the Universe from our neighbours (the Local Group) to its observable edges (the Cosmic Microwave Background). We know that most of the visible matter in the Universe, in the form of stars, gas, and dust grains, is organized in galaxies. Galaxies come in many different forms and sizes (as clearly shown in Fig.1.1), but they can be broadly divided into two main species. Spirals, with a flattened, disk-like shape, blue colors, much gas and dust, and a widespread star formation activity that results in the presence within them of many young stars. Ellipticals, with a spheroidal shape, red colors, little or no gas and dust, and no star formation activity, thus containing exclusively old stars. We also know that the density of galaxies in the local Universe is not at all constant, but it spans from 0.2 ρ in voids to 5 ρ in superclusters and filaments, 100 ρ in ∼ 0 ∼ 0 ∼ 0 the cores of rich clusters, up to 1000 ρ in compact groups, where ρ is the average ∼ 0 0 5 6 1. Introduction

Figure 1.1: An example of the heterogeneous population of galaxies that inhabit our Universe. Mosaic of RGB (g,r,i) images adapted from Frei et al. (1996) 7

“field” density (Geller & Huchra 1989). It is well established that morphological type and local density are not independent quantities. In their analysis of 55 nearby clusters, Dressler (1980) and Whitmore et al. (1993) demonstrated that the fraction of spiral galaxies decreases from 60% in the “field” to virtually zero in the cores of rich clusters, compensated by an opposite increase of elliptical and S0 galaxies. This phenomenon, known as morphology segregation, is considered as the clearest observational signature of significant environmental dependences of the processes that govern the formation and the evolution of galaxies. Understanding the origin of this phenomenon (”Nature or Nurture?”) probably represents one of the major challenges of extragalactic astronomy. One possible way to overcome this problem is to take advantage of the effect provided by the finite speed of light. Observing today galaxies at different distances means observing them at different epochs in the history of the Universe, and thus with different ages. This investigative method is providing us with a sort of evolutionary sequence for galaxies: starting from the pioneering work by Butcher & Oemler (1978, 1984) we know that distant (and thus young) cluster of galaxies contain a much higher fraction of blue galaxies than nearby clusters. Recently Dressler et al. (1997) used high-resolution imaging with the Hubble Space Telescope (HST) to measure the morphology-density relation in the core regions of a sample of rich clusters at z 0.5. They found that the fraction of lenticular galaxies (S0s) in ∼ clusters declined by a factor of 2-3 between z = 0 and z = 0.5, and this evolution was accompanied by a corresponding increase in the fraction of star-forming spirals (see also Couch et al. 1998; Treu et al. 2003). Many research groups have suggested that the predominance of early type galaxies in local clusters is the result of physical processes that suppress star formation and eventually alter galaxy morphology. and several mechanisms have been proposed (see Boselli & Gavazzi 2006, for a detailed review):

Galaxy interaction with the intra-cluster medium (ICM). • Ram pressure stripping (Gunn & Gott 1972). As a galaxy orbits through a cluster, it experiences ram pressure from the ICM. When the ram pressure is greater than the binding force, the cold gas will be stripped (Abadi et al. 1999; Quilis et al. 2000; Vollmer et al. 2001). Before leading to complete gas abla- tion, ram-pressure could produces significant compression ahead of the galaxy temporally increasing its star formation activity (Bekki & Couch 2003). Even if it is well established that this phenomenon would finally lead to a gradual decrease in galaxy star formation activity, its effects on galaxy morphology are not yet completely understood (Fujita & Nagashima 1999; Mihos 2004a). Ram- pressure stripping is likely to be effective in the central region of clusters where the density of intra-cluster medium (ICM) is high. Viscous stripping (Nulsen 1982). In a galaxy travelling into the ICM the outer layers of the interstellar medium (ISM) experience a viscosity momentum trans- 8 1. Introduction

fer that could be sufficient for dragging out part of its gas. Thermal evaporation (Cowie & Songaila 1977). If the ICM temperature is high compared to the galaxy velocity dispersion, at the interface between the hot ICM and the cold ISM the temperature of the ISM rises rapidly, thus the gas evaporate and is not retained by the galaxy gravitational field. Starvation (or strangulation) (Larson et al. 1980a). This mechanism consists in the removal of the diffuse hot gas reservoir that is confined in the galaxy halo. Since this tenuous halo is less bound than the cold gas in the disk, its stripping is considerably easier (Bekki et al. 2002). A galaxy whose hot gas reservoir is removed slowly, exhausts its cold gas in more than one gigayear, because there is no supply of fresh gas from the surrounding hot gas. Note that while stripping gas from disks induces a truncation of star formation activity on a short timescale ( 107 yr), strangulation is expected to affect a galaxy star formation history on ∼a long time scale (> 1 Gyr) provoking a slowly declining activity which consumes the disk gas after the supply of cooling gas has been removed. All of the above mechanisms but starvation need relatively high den- sity of hot intra-cluster gas, and thus likely to happen in the central region of clusters. However Fujita (2004) has pointed out that ram pressure and thermal evaporation could not be negligible in cluster sub-clump regions (small groups around a cluster).

Galaxy-galaxy gravitational interaction. Collisions or close encounters between • galaxies can have a strong effect on their morphology and star formation rates. Various simulations have shown that major mergers between disk galaxies can produce galaxies resembling ellipticals as merger remnants (e.g.,Toomre & Toomre 1972; Barnes & Hernquist 1996) and that accretion of small onto spi- rals can transform the host spiral to S0 type (Walker et al. 1996). The tidal forces generated during the interaction tend to funnel gas toward the galaxy center. It is likely that this will fuel a central starburst, ejecting a large frac- tion of material. Gas in the outer part of the disk, on the other hand, will be drawn out of the galaxy by the encounter (Mihos 2004a). Although individual collisions are expected to be most effective in groups because the velocity of the encounters is too high for such mergers to be frequent (Ghigna et al. 1998; Okamoto & Habe 1999), Moore et al. (1996) showed that the cumulative effect of many weak high velocity interactions (i.e. galaxy harassment) can also be important in cluster of galaxies. However its influence is largely limited to low luminosity galaxies, while in bright spirals its effects are considerably milder (Mihos 2004a; Moore et al. 1996).

Galaxy-Cluster gravitational interaction. Tidal compression of galactic gas via • interaction with the whole cluster potential can effectively perturb cluster galax- 9

ies, inducing gas inflow, bar formation, nuclear and perhaps disk star formation (Merritt 1984; Miller 1986; Byrd & Valtonen 1990). On the other hand, gas can be hardly removed directly by the interaction (Boselli & Gavazzi 2006).

Although we have collected a plethora of observational evidences that at least some of these processes are playing a significant role on galaxy evolution we have not shed light on the origin of the morphology density relation. This is in part due to the fact that we do not yet know their detailed physics and the relative importance of each mechanism during the different phases of galaxy evolution.

Moreover the arduous effort of reconstructing the evolutionary history of galaxies would turn out to be completely useless if we did not take into account that the whole Universe is evolving, changing the physical condition of the environments populated by galaxies. In fact different and mostly independent observational evidences, as the Cosmic Microwave Background radiation (Kogut et al. 2003), the large scale structure (Hawkins et al. 2003) and supernovae observations (Tonry et al. 2003), are telling us that the Universe in not only expanding (Hubble & Humason 1931), but it is also accelerating. If theorists are right, this implies that the Universe is dominated by its energetic and a matter dark components, whose nature is still completely unknown. The dark energy term (usually indicated with the cosmological constant Λ) allows for the current accelerating expansion of the universe. Currently, 70% of the energy density of the Universe is supposed to be in this form. The dark matter∼ component of the Universe is supposed to be cold (i.e. not thermalized), non-baryonic, collisionless ”material”. This component makes up 26% of the energy density of the present ∼ Universe and only the remaining 4% is the matter and energy we directly observe. The only way to shade light on the∼properties of our, mostly obscure, Universe is thus through numerical simulations (e.g. Kauffmann et al. 1993; Springel et al. 2005). In particular hierarchical galaxy formation (White & Rees 1978) models within a Λ cold dark matter (ΛCDM) cosmogony are currently considered the most successful paradigm for understanding the evolution of matter in the Universe. In this scenario, structures grow hierarchically via gravitational instability from small perturbations seeded in the early epoch. The density of dark matter its component is a proxy for the epoch of initial collapse of a given structure: the most massive structures at any epoch represent the earliest that collapsed (Springel & Hernquist 2003). After their collapse, structures grow up through infall of smaller groups (Kauffmann 1995). However the typical size of the infalling groups increases with the age of the Universe but their infall rate considerably decreases (Ghigna et al. 1998; Okamoto & Habe 1999; Gnedin 2003). This means that clusters have accreted great part of their galaxy population in the past, through infalling of small groups. Today the accretion of new members is supposed to be rare and to happen mainly through the merging of big subclusters. Adding the well known observational evidence that star formation rapidly decreases 10 1. Introduction

with the age of the Universe (Lilly et al. 1996; Madau et al. 1998), we are facing a scenario that, at a first look, seems to suggest that studying star formation in rich clusters today is a melancholy affair. The Universe we inhabit today is old, and most of its star formation activity has gone out. In addition (and this is the worse part of the story) the Universe dramatically evolved itself, altering continuously the physi- cal conditions of the environments populated by galaxies. This implies that galaxies could have experienced different environmental effects during their history and that the dominant process in the local Universe could have been completely negligible in the early stages of its evolution, while the process shaping galaxy evolution could be less important in today clusters. Let us imagine, as predicted by models, that a great fraction of today cluster galaxies have infalled, within a compact group, into a cluster 5 Gyr ago. While in the group environment tidal interactions were very strong ∼and influenced significantly star formation activity and galaxy morphology; today, in cluster environment, gravitational interactions are less probable due the large relative velocities of cluster members (Ostriker 1980). Thus great part of galaxy evolution took place before the infalling into the cluster core. The discourage felt by a young student facing this music increases reading the re- cent review by Dressler (2004) on ”Star forming galaxies in clusters”: What we see in clusters today is only a faint echo of what once was... looking for star formation in today’s clusters is a little bit like searching for the last cashew in a picked-over nut-cup.[...] Star formation in rich clusters today is a pretty sad affair. Spirals are ”running down compared to half-a-Hubble time ago. The spirals that will be drawn into rich clusters in the future will die the death of a thousand cuts: in the rich group environment into which they have for so long been entrained, they are likely already to have had their fates sealed long ago. Thus, what has he to do? Give up and concentrate all his efforts on the study of the high redshift, still young, Universe? Obviously the answer is no; and not because this work would be useless. High redshift and local observations are complementary to give more insights on galaxy evolution and, until we will be able to understand all the physical mechanism influencing the present evolution of nearby galaxies (and we are still far from reaching this goal), it would be an error to concentrate all our efforts only at high redshift. Observations of the high redshift Universe approach us to the mechanisms that maybe shaped part of the morphology density relation; however today there is still insuffi- cient high-quality data to put strong constraint on different models (Dressler 2004). On the contrary in the local Universe, maybe we are missing most of the action, but we have the unique possibility to observe in detail galaxy properties over the whole range of sizes and masses, and study in detail the effects of different environmental mechanisms. In particular, what makes the local Universe still exciting? What can we learn about galaxy evolution that would still be impossible if we moved to higher redshifts? Owing 11

to the high quality images we can obtain for local galaxies, an extremely accurate and homogeneous morphological classification is possible down to MB 13, allowing a detailed discrimination among different subclasses of early-type galaxies≤ − (ellipticals, lenticulars, dwarfs) and among early-type galaxies and quiescent spirals (see the Virgo Cluster catalogue a sort of ”milestone” of the morphological classification, Binggeli et al. 1985). Accurate morphological classification becomes a difficult task just at the Coma cluster distance (z 0.025) and more or less impossible at higher redshift (Abraham et al. 1996a). The∼objects in the images are very small, thus it is very hard to detect the fine structure elements needed to distinguish different classes. In order to solve this problem alternative classifications based on structural parameters (Abraham et al. 1996b) or on spectral type (Madgwick et al. 2002) have been proposed but they are only useful to discriminate between a star forming disk and a quiescent bulge dominated galaxy, completely failing to distinguish between an elliptical and a lenticular or between an early-type galaxy and a bulge-dominated Sa spiral disks (Scodeggio et al. 2002; Gavazzi et al. 2002a). Thus at high redshift we can observe the evolution of the star formation-density relation (the Butcher-Oemler effect) but we cannot investigate morphological transformations that eventually affected galaxy evolution (Smail et al. 1997; Fabricant et al. 2000; Smith et al. 2005). Moreover in the local Universe we can study galaxies spanning all ranges of mass and 2 luminosity, reaching very faint (MB 13) low surface brightness ( 30 mag arcsec ) dwarf galaxies. This is crucial to study∼ −the (strong) dependence of∼galaxy evolution (Gavazzi et al. 2002a) and environmental effects with mass since the anti-hierarchical relation between star formation history and galaxy mass is one of the great challenge for models of galaxy evolution. Moreover dwarf galaxies today represent probably the major failure of hierarchical galaxy formation models: cold dark matter theory predicts that the groups and clusters of galaxies should contain many more dwarf objects than the observed number of dwarf galaxies (Klypin et al. 1999; Moore et al. 1999). Several explanation has been proposed (Somerville 2002), and even if no solu- tion has been found so far, it is indisputable that the only way to solve this problem is to understand the formation and evolution of dwarf galaxies, a task possible only in the local Universe. Another serious limit of high redshift observations is the quantification of star forma- tion activity in galaxies. The easiest and common way to estimate star formation rate for distant galaxies is through rest-frame ultraviolet (UV) observations. However ul- traviolet emission is strongly affected by dust attenuation: absorption by dust grains reddens the spectra at short wavelengths and modifies altogether the spectral energy distribution of galaxies. Since the UV radiation is emitted by young stars (t < 108 yr) that are deeply embedded in dust clouds than older stellar populations, rest-frame UV observations can lead to incomplete and/or biased reconstructions of the star formation activity and star formation history of galaxies. Moreover we have not yet a good characterization of the dust attenuation properties in galaxies and of their 12 1. Introduction

dependences with galaxy type (i.e. normal star forming galaxies vs. starburst) and no proper corrections have been achieved, having no possibility to correctly quantify the star formation rate at high redshift. As extensively discussed in Chapter 7 of this work, understanding dust properties and looking for empirical relations suitable for deriving dust attenuation corrections is today possible only for low redshift galaxies: in this case the study of the local Universe is mandatory to correctly interpret what we observe in distant galaxies. Finally, as remarked by Poggianti (2004a), in order to understand what happens to galaxies in clusters, two crucial pieces of information are 1) the gas content of cluster galaxies (i.e. the fuel for future star formation) and 2) the spatial distribution of the gas and of the star formation activity within each galaxy (i.e. differences from field galaxies are good indicators of environmental effects); and both can be achieved only in the local Universe. Neutral hydrogen (HI) and Hα observations observations1 are still a prerogative of nearby galaxies. In the near future, thanks to the advent of the Arecibo L-band Feed Array, it should be possible to detect an hydrogen mass 9 of 10 M at z 0.15, but only with a very high integration time ( 70 hours per∼beam). The few∼ examples shown above represent only the tip of the∼iceberg of the unique capability of local Universe observations to disclose the secret of galaxy evolution. We would lose too much, without any significant improvement, if we aban- doned observations of nearby galaxies in order to move our attention at high redshift galaxies.

The aim of this work Firmly convinced of the great significance of nearby Uni- verse observations, I have concentrated all my PhD work on the study of environmen- tal effects on the evolution of nearby clusters. In particular this thesis will focus on three different clusters: Abell1367, Virgo and Coma. These three clusters are among the best studied in the local universe and, due to the variety of their environmental conditions (e.g. spiral fraction, X-ray luminosity, evolutionary stage) they represent the most suitable ”laboratory” for comparative studies. The novelty of this work is that in addition to the optical and near infrared observations carried out during the last fifteen years by G.Gavazzi and A.Boselli (available through the GOLDMine database, Gavazzi et al. 2003a: http://goldmine.mib.infn.it) I will take for the first time advantage of recent UV observations by the Galaxy Evolution Explorer (GALEX, Martin et al. 2005). The use of a multiwavelength dataset is crucial to understand galaxy evolution since different galactic components such as old, new or evolved stars; active galactic nuclei; the interstellar medium contribute in different amounts to the observed emission at different wavelengths, from the radio to X-rays. Therefore, the comparison of global emission properties at a wide range of wavelengths can give

1The Hα Balmer emission (λ=6562.8 A)˚ is the most direct indicator of the current (< 4 106 yrs), massive (> 8 M ) star formation activity in galaxies (Kennicutt 1998)

13 us precious insight on the relative importance of these components, as well as on the origin of some parts of the emission spectrum. Since different emission bands have different sensitivities to absorption, their comparison may also give us insight into the dust content of the emitting regions. Moreover, comparison of global multi- wavelength emission properties of galaxies of different morphology can give us insight on the relative presence of different galactic components throughout the Hubble se- quence. While most of the studies of galaxies make use of individual energy bands, mainly the optical but also the radio and, more recently, the X-ray and infrared, it is rarer to find work comparing data from two or more emission windows. In particular the rest frame UV emission provides a powerful tool for measuring and understanding star formation in galaxies at all epochs. Ironically, the interpretation of high-redshift galaxies in the rest UV is most limited by the lack of large, systematic surveys of low- redshift UV galaxies serving as a benchmark. However, before the launch of GALEX, only a few experiments had observed the nearby Universe at ultraviolet wavelengths (Smith & Cornett 1982; Lampton et al. 1990; Kodaira et al. 1990). Among them, the FOCA experiment (Milliard et al. 1991) allowed the first determinations of the UV LF of local field galaxies (Treyer et al. 1998; Sullivan et al. 2000) and and the local rest-UV anchor point for the star formation history plot. However its low sensitivity and resolution and the small sky area covered. With its large field of view (diameter 1.2 degrees), high sensitivity and two ultraviolet filters, GALEX has opened a new ∼era in the UV astronomy, providing us for the first time with a large, complete and homogeneous dataset to study star formation activity in galaxies. Using this unique mine of data I will investigate the properties of galaxies from differ- ent points of view: one statistical, analyzing the global properties of the whole cluster sample; and another much more focused on the study of particular objects considered as prototypes of the different ways in which the environment could influence galaxy evolution. The comparison of all the observational results with models will be used to build up an evolutionary scenario for galaxies, linking the information I obtain in this work to what we know (or think to know) about the evolution of galaxies at higher redshifts.

The organization of the thesis In Chapter 2 I briefly describe the different datasets used in this work: the GALEX satellite and its mission, and the GOLDMine database. In Chapter 3 and 4 I start the statistical analysis of cluster galaxies, computing the UV luminosity function for nearby clusters. The analysis presented in Chapter 3 was performed before the launch of GALEX, thus I used data from the FOCA (Milliard et al. 1991) experiment of the three nearby clusters studied in this work. When GALEX was launched I had the possibility to extend my analysis two magnitudes deeper with higher quality data. First of all, this double estimate allow me to directly 14 1. Introduction

compare two different and independent datasets. Then the comparison between the cluster luminosity function and the field one is used to determine whether the envi- ronment affects the shape of the cluster luminosity function. In Chapter 5 I study the influence of the dynamical state of a cluster to the evolution of galaxies, performing a detailed dynamical study of the Abell cluster 1367. This cluster is considered as the prototype of a dynamically young local cluster, thus rep- resenting a good place to study the effects of a cluster’s assembly on galaxy evolution. Although X-ray, radio and optical observations suggest that Abell 1367 is dynami- cally young and it is still undergoing the process of formation, detailed spatial and dynamical analysis of this cluster has not been attempted so far. Since the dynamical state of a cluster is directly linked with its evolution this work will allow us to have a clear picture of the past, current and future assembly history of this structure and its galaxies. In Chapter 6 I focus my attention on the population of early-type galaxies in clusters, in particular studying the UV properties of giant and dwarf ellipticals and lenticu- lars in the Virgo cluster, in order to determine whether these different morphological types had the same evolutionary history or not. On the contrary from Chapter 7 till the end of this work I move my attention to the star forming cluster population. As discussed above if we want to use ultraviolet radiation to correctly estimate star formation we need to correct for dust attenuation. Thus in Chapter 7 I present an analysis of dust attenuation properties in nearby clus- ter star forming galaxies, obtaining a cookbook in order to estimate dust attenuation without far infrared observations. This analysis represents the tip of the iceberg and only a future comparison with different dust models will allow us to understand dust attenuation and to know how to correct UV observations of local and high redshifts galaxies. Thus, a statistical analysis of star formation activity in cluster galaxies us- ing UV data is still impossible. For this reason in Chapter 8, 9 and 10 I will focalize my attention on the study of three particular cluster galaxies considered as the proto- types of the three main environmental effects observed in clusters: tidal interaction, ram pressure stripping and preprocessing, respectively. These unique astrophysical laboratories will help me to understand the effects of different physical mechanisms on galaxy evolution in more depth. Finally in Chapter 11 I will summarize the evolutionary scenario for cluster galaxies which emerged from this work.

Great part of this thesis is published or submitted for publication on major astro- nomical refereed journals: Gavazzi et al. (2003b, 2006); Cortese et al. (2003a, 2004, 2005, 2006); Boselli et al. (2005a,b). Chapter 2

GALEX & GOLDMINE: A multiwavelength window on the Local Universe

2.1 The Galaxy Evolution Explorer

The Galaxy Evolution Explorer (GALEX) is a NASA Small Explorer class mission. It consists of a 50 cm-diameter, modified Ritchey-Chr´etien telescope with four op- erating modes: Far-UV (FUV) and Near-UV (NUV) imaging, and FUV and NUV spectroscopy. The telescope has a 3-m focal and the field of view is 1.2◦ circular (see Fig 2.1 and Table 2.1). Spectroscopic observations are obtained at multiple grism-sky dispersion angles, so as to mitigate spectral overlap effects. The FUV (1528A:˚ 1344- 1786A)˚ and NUV (2271A:˚ 1771-2831A)˚ imagers (see Fig.2.2) can be operated one at a time or simultaneously using a dichroic beam splitter. The FUV detector is pre- ceded by a blue-edge filter that blocks the night-side airglow lines of [OI]1304, 1356, and Lyα. The NUV detector is preceded by a red blocking filter/fold mirror, which reduces both zodiacal light background and optical contamination. The peak quan- tum efficiency of the detector is 12% (FUV) and 8% (NUV). The detectors are linear up to a local (stellar) count-rate of 100 (FUV), 400 (NUV) cps, which corresponds to mAB 14 15. The resolution of the system is typically 4.5/6.0 (FUV/NUV) arcseconds∼ (FWHM),− and varies by 20% over the field of view. Further detail ∼ about the mission, in general, and the performance of the satellite, in specific, can be found in Martin et al. (2005) and Morrissey et al. (2005), respectively. The mission is nominally expected to last 38 months; GALEX was launched into a 700 km, 29◦ inclination, circular orbit on 28 April 2003.

15 2. GALEX & GOLDMINE: A multiwavelength window on the Local 16 Universe

Figure 2.1: Cross section of the instrument portion of GALEX. The optical path is outlined in blue. Overall dimensions of the view shown are 1.5 m 1 m (adapted from Morrissey et al. 2005).

2.1.1 The Prime Mission GALEX is currently undertaking the first space UV sky-survey, including both imag- ing and grism surveys. The prime mission includes an all-sky imaging survey (AIS: 75-95% of the observable sky, subject to bright-star and diffuse Galactic background light limits) (m 20.5), a medium imaging survey (MIS) of 1000 deg2 (m 23), AB ' AB ' a deep imaging survey (DIS) of 100 square degrees (mAB 25), and a nearby galaxy survey (NGS). Spectroscopic (slit-less) grism surveys (R=100-200)' are also being un- dertaken with various depths and sky coverage. Many of the GALEX fields overlap existing and/or planned ground–based and space-based surveys being undertaken in other bands. All-sky Imaging Survey (AIS): The goal of the AIS is to survey the entire sky subject to a sensitivity of m 20.5, comparable to the POSS II (m =21 mag) and SDSS AB ' AB spectroscopic (mAB=17.6 mag) limits. Several hundreds to 1,000 objects are in each 1 deg2 field. The AIS is performed in roughly ten 100-second pointed exposures per eclipse ( 10 deg2 per eclipse). ∼ Medium Imaging Survey (MIS): The MIS covers 1000 deg2, with extensive overlap of the Sloan Digital Sky Survey. MIS exposures are a single eclipse, typically 1500 sec- onds, with sensitivity m 23, net several thousand objects, and are well-matched AB ' to SDSS photometric limits. 2.1. The Galaxy Evolution

Item FUV Band NUV Band Bandwidth: 1344 – 1786 A˚ 1771 – 2831 A˚ Effective wavelength (λeff ): 1528 A˚ 2271 A˚ Field of view: 1.28◦ 1.24◦ Zero point (m0): 18.82 mag 20.08 mag Image resolution (FWHM): 4.5 arcsec 6.0 arcsec Spectral resolution (λ/∆λ): 200 90 Detector background (typical): Total: 78 cnt sec−1 193 cnt sec−1 Diffuse: 0.66 cnt sec−1-cm−2 1.82 cnt sec−1-cm−2 Hotspots: 47 cnt sec−1 107 cnt sec−1 Sky background (typical): 2000 c-sec−1 20000 cnt sec−1 Limiting magnitude (5σ): AIS (100 sec): 19.9 mag 20.8 mag MIS (1500 sec): 22.6 mag 22.7 mag DIS (30000 sec): 24.8 mag 24.4 mag

Table 2.1: Selected Performance Parameters (Morrissey et al. 2005)

Deep Imaging Survey (DIS): The DIS consists of 20 orbit (30 ksec, m 25) ex- AB ' posures, over 80 deg2, located in regions where major multiwavelength efforts are already underway. DIS regions have low extinction, low zodiacal and diffuse galactic backgrounds, contiguous pointings of 10 deg2 to obtain large cosmic volumes, and minimal bright stars. An Ultra DIS of 200 ksec, mAB 26 mag is also in progress in four fields. ∼ Nearby Galaxies Survey (NGS): The NGS targets well-resolved nearby galaxies for −2 1-2 eclipses. Surface brightness limits are mAB 27.5 mag arcsec . The 200 targets are a diverse selection of galaxy types and environmen∼ ts (see Fig.2.3). Spectroscopic Surveys. The suite of spectroscopic surveys includes: the Wide-field Spectroscopic Survey (WSS), which covers the full 80 deg2 DIS footprint with com- parable exposure time (30 ksec), and reaches m 20 mag for S/N 10 spectra; the AB ∼ ∼ Medium Spectroscopic Survey (MSS), which covers the high priority central field in 2 each DIS survey region (total 8 deg ) to mAB=21.5-23.0 mag, using 300 ksec expo- sures; and the Deep Spectroscopic Survey (DSS) covering 2 deg2 with 1,000 eclipses, to a depth o f mAB=23-24 mag.

2.1.2 Data collection mode GALEX performs its surveys with plans that employ a simple operational scheme requiring only two observational modes and two instrument configurations. Each orbit GALEX collects data during night segments (eclipses) and visits to a single 2. GALEX & GOLDMINE: A multiwavelength window on the Local 18 Universe

pre-programmed target. Each target consists either of a single pointing (single visit observation) or multiple adjacent pointings (sub-visit observations). Currently sub- visits are only used for all-sky imaging survey (AIS) and in-flight calibration ob- servations. After removing instrument overhead, each eclipse typically yields up to 1700 seconds of usable science data. During any visit or sub-visit observation the spacecraft attitude is controlled in a tight, spiraled dither. A spiral dither is used to prevent ”burn-in” of the detector active area by bright objects and to average over high spatial frequency response variations. For each sub-visit the spiral dither pattern is restarted. Since celestial sources will move on the detector, the pipeline software will reposition the time-tagged photons to common sky coordinates based on the satellite aspect solution. As many as 12 sub-visits are allowed per eclipse period (typical for AIS), with all-sky survey sub-visits obtaining 100-110 s exposure time per leg. For plans with sub-visit targets, a 20 second slew time is required to move between each leg of the observation. For some survey plans (e.g. deep imaging, spectroscopy), a single visit is insufficient to build up the requisite signal-to-noise, so a series of visits are needed in order to obtain the minimum required exposure time.

2.1.3 Counts vs. magnitudes and fluxes conversions All GALEX data are normalized to their relative exposure time, thus each count (cnt) measured on a GALEX is in reality a cnt per sec (CPS). Below are given some equations useful to convert galaxies counts into fluxes or magnitudes. To convert from GALEX counts per sec (cps) to flux (erg cm−2 s−1 A˚−1):

F UV : F lux [erg cm−2 s−1 A˚−1] = 1.40 10−15 CPS (2.1) × × NUV : F lux [erg cm−2 s−1 A˚−1] = 2.06 10−16 CPS (2.2) × × To convert from GALEX counts per sec to magnitudes in the AB system (Oke 1974):

F UV : m(AB) = 2.5 log(CP S) + 18.82 (2.3) − × NUV : m(AB) = 2.5 log(CP S) + 20.08 (2.4) − × Thus to convert from flux to AB magnitudes:

F lux [erg cm−2 s−1 A˚−1] F UV : m(AB) = 2.5 log + 18.82 (2.5) − × 1.40 10−15 ×  F lux [erg cm−2 s−1 A˚−1] NUV : m(AB) = 2.5 log + 20.08 (2.6) − × 2.06 10−16 ×  The current estimates are that the zero-points defined here are accurate to within +/- 10% (1 sigma). 2.1. The Galaxy Evolution

Figure 2.2: The transmittance profile for the NUV and FUV GALEX filters. Different galaxy spectral energy distributions are superposed.

Figure 2.3: Example of GALEX image. GALEX NGS observation of NGC4631. In the color table, red-green (gold) is used for NUV, and blue for FUV. 2. GALEX & GOLDMINE: A multiwavelength window on the Local 20 Universe

2.2 The Galaxy On Line Database Milano Net- work

The Galaxy On Line Database Milano Network (http://goldmine.mib.infn.it) is de- signed to provide access to all the data collected by G.Gavazzi, A.Boselli (Tutor and Co-Tutor of this thesis) and collaborators during several observational campaigns, started in 1985 and still in progress, aimed at providing the phenomenology of local galaxies in the widest possible frequency range. The creation of the World Wide Web site and of the MySQL database has been performed by P. Franzetti and A. Donati and a detailed description of the database architecture can be found in Donati (2004). GOLDmine is focused on 9 local clusters of galaxies: A262 (Perseus-Pisces), Cancer, A1367, A1656 (Coma), Virgo, A2147, A2151, A2197, A2199 (Hercules). In addition it contains a filament of nearly isolated galaxies, the so called “Great Wall”, thus providing the ideal laboratory for comparative analyses of galaxies in different envi- ronments, spanning a factor of 20-100 in local galaxy density. Objects are selected in the above regions with strictly optical completeness criteria. Galaxies brighter than mp = 15.7 are taken from the Catalogue of Galaxies and of Clusters of Galaxies (CGCG) by Zwicky et al. (1961) in all clusters except Virgo where objects brighter than mp = 20.0 are taken from the Virgo Cluster Catalogue (VCC) by Binggeli et al. (1985). Obviously, due to the factor of 5 difference in distance between Virgo and ∼ the other clusters, this selection limit results in dwarf galaxies being included in our database only for the Virgo cluster. However globally GOLDmine covers the whole range (4 orders of magnitude) of luminosities spanned by real galaxies. GOLDmine contains 3649 galaxies. Extensive campaigns were carried out to observe as many as possible of the 3649 target galaxies through all possible observational windows, a task that we did not complete yet. The parameters listed in the GOLDmine database are divided into 5 categories: Gen- eral, Continuum and Line photometry, Dynamical and Structural. They can be obtained from GOLDmine by querying the database for an individual galaxy name or “by parameters”, “by near name or position” or “by available im- ages”. In this case all galaxies in a given range of photographic magnitude, and morphological type can be selected. General parameters include Catalogue designations, (J2000) celestial coordinates, op- tical diameters, photographic magnitude, redshift, distance, morphological type. Continuum parameters include: UV, U, B, V, J, H, K magnitudes computed at the optical radius (25th mag arcsec−2) (see Gavazzi et al. 1996); IRAS 60 and 100 mi- cron fluxes; radio continuum fluxes densities at 0.6 and 1.5 GHz. Line photometry includes: the atomic (HI) and molecular (H2) hydrogen mass; the Hα+[NII] line equivalent width and flux. Dynamical parameters include: the width of the HI line, with a quality flag; the width of the Hα line and the central velocity dispersion. 2.2. The Galaxy On Line Database Milano Network 21

Structural parameters include: the light concentration index (C31); the effective ra- dius Re; the effective surface brightness µe; the total asymptotic magnitude. These quantities (see Scodeggio et al. 2002) are given separately for the H, V and B bands. The novelty of GOLDmine consists of its image section, where images can be down- loaded in JPG and FITS format. Images include: Finding Charts from the Digitized Palomar Sky Survey for all galaxies. Broad band images obtained in the B, V, H and K bands. Narrow band images in the light of Hα and a red image of the underlying stellar continuum near Hα. RGB images. For some galaxies we combined several images to obtain “true” color pictures. Radial profiles of the light distribution as obtained on the available (B, V, H) images (see Gavazzi et al. 2000). When at least two radial profiles are available the color radial profile is also shown. Optical spectra integrated over the whole surface of the galaxy, obtained in drift-scan mode, i.e. by drifting the spectrograph slit over the galaxy extension (see Gavazzi et al. 2002a, 2004). Spectral Energy Distributions (SEDs) from the UV to the centimetric radio continuum obtained from broad-band photometry. The plotted data are total fluxes (extrapolated to the optical radii), unlike the individual aper- ture data given by NED. However they are given as observed, i.e. uncorrected for extinction from our Galaxy and for internal extinction (see Boselli et al. 2003a). It is our goal to provide a homogeneous set of keywords in all FITS header to character- ize the data, including: effective integration time, filter, telescope, WCS parameters, photometric effective zero point. This homogenization is not yet complete. As also remarked in Chapter 7, the high quality of data available through GOLDMine, make this datasample one of the most appropriate for studying the evolution of nearby galaxies.

Chapter 3

The FAUST-FOCA UV luminosity function of nearby clusters

3.1 Introduction

The study of the galaxy luminosity function (hereafter LF) provides us with a fun- damental tool for testing theories of galaxy formation and for reconstructing their evolution to the present. Recent determinations of the galaxy LF at various frequen- cies, in various environments (i.e.De Propris et al. 2003; Madgwick et al. 2002) and in a number of redshift intervals (i.e.Ilbert et al. 2004) have improved our knowledge of galaxy evolution and the role played by the environment in regulating the star formation activity of galaxies. The optical cluster LF is significantly steeper than that in the field (Trentham et al. 2005). This steepening is due to quiescent galaxies, more frequent at low luminosities in clusters, while the LF of cluster star forming objects is similar to that in the field (De Propris et al. 2003). The causes of this difference might reside in the density-morphology relation (Dressler 1980; Whitmore et al. 1993) and in particular in the overabundance of dwarf ellipticals in rich clusters (Ferguson & Sandage 1991), whose origin is currently debated in the framework of the environmental effects on galaxy evolution. The ultraviolet emission UV( 2000 A),˚ being dominated by young stars of interme- ∼ diate masses (2 < M < 5 M , Boselli et al. 2001) represents an appropriate tool to identify and quantify star formation activity. Although before the launch of GALEX, the shape of local field UV LF (Sullivan et al. 2000) was supposed to be well deter- mined, there was still a fair amount of uncertainty on the UV luminosity function of clusters. Its slope was undetermined due to the insufficient knowledge of the back- ground counts (Cortese et al. 2003b). Andreon (1999) proposed a very steep faint end (α 2.0, 2.2), significantly different from the field LF (α 1.5). However Cortese∼et−al. (2003b)− pointed out that this steep slope is likely∼caused− by an un-

23 24 3. The FAUST-FOCA UV luminosity function of nearby clusters

derestimation of the density of background galaxies and proposed a flatter faint-end slope (α 1.35 0.20). Unfortunately the statistical uncertainty was too high for making reliable∼ − comparisons between the cluster and the field LFs. In this chapter I re-compute the cluster UV luminosity function with two major improvements over previous determinations. We increase the redshift completeness of the UV selected sample using new spectroscopic observations of Coma and Abell 1367 (see Chapter 5 and Cortese et al. 2004), and compute for the first time the UV LF of the Virgo cluster. These improvements are not sufficient to constrain the LF of each individual cluster, however the UV composite luminosity function, constructed for the first time in this paper can be compared with that of the field. Doing so I try anticipating one of the main goals of the Galaxy Evolution Explorer (GALEX) which, as shown in the next Chapter, will help us shade light on the UV properties of galaxies and their environmental dependences. We assume a distance modulus µ= 31.15 for the Virgo cluster (Gavazzi et al. 1999a), µ=34.80 for Abell 1367 and µ=34.91 for the Coma cluster (Gavazzi et al. 1999b).

3.2 The Data

The sample analyzed in this chapter comprises the UV sources detected in Virgo, Coma and Abell 1367 clusters by the FOCA (Milliard et al. 1991) and FAUST (Lampton et al. 1990) experiments. The FOCA balloon-borne wide field UV camera (λ = 2000A;˚ ∆λ = 150A)˚ observed 3 square degrees ( 8 Mpc2) in the Abell 1367 ∼ ∼ (unpublished data) and Coma clusters (Donas et al. 1995) and 12 square degrees ( 1 Mpc2) in the Virgo cluster (data are taken from the extragalactic∼ database GOLDMine,∼ Gavazzi et al. 2003a). The FOCA observations of Virgo are not suf- ficient to compile a complete catalog: no sources brighter than mUV 12.2 were detected due to the small area covered. We thus complement the UV database∼ with the wide field observations performed by the FAUST space experiment (λ = 1650A;˚ ∆λ = 250A)˚ in the Virgo direction (Deharveng et al. 1994), covering 100 square degrees ( 8.8 Mpc2). The FAUST completeness limit is m 12.2∼(Cohen et al. ∼ UV ∼ 1994), significantly lower than the FOCA magnitude limit: mUV 18.5. However combining the two UV catalogs we hope to constrain the shape of the∼ UV luminosity function across 7 magnitudes. We use the FAUST observations for mUV < 12.2 and the FOCA observations for mUV 12.2. To account for the different response func- tion of FAUST and FOCA filters ≥we transform the UV magnitudes taken by FAUST at 1650A˚ assuming a constant color index: UV(2000) = UV(1650) + 0.2 mag (Dehar- veng et al. 1994, 2002). We think however that this difference does not bias the galaxy populations selected by the two experiments. The estimated error on the UV magni- tudes is 0.3 mag in general, but it ranges from 0.2 mag for bright galaxies, to 0.5 mag for faint sources observed in frames with larger than average calibration uncertain- 3.3. The UV luminosity functions 25

ties. The UV emission associated with bright galaxies is generally clumpy, thus it has been obtained by integrating the flux over the galaxy optical extension, determined at the surface brightness of 25 mag arcsec−2 in the B-band. The spatial resolution of the UV observations is 20 arcsec and 4 arcmin for FOCA and FAUST respectively. The astrometric accuracy is therefore insufficient for unambiguously discriminating between stars and galaxies. To overcome this limitation, we cross-correlate the UV catalogs with the deepest optical catalogs of galaxies available: the Virgo Cluster Catalog (VCC, Binggeli et al. 1985), complete to mB 18, for the Virgo cluster 0 ∼ and the r band catalog by Iglesias-P´aramo et al. (2003), complete to mr0 20, for Coma and Abell 1367. We used as matching radius the spatial resolution∼of each observation. In case of multiple identifications we select the galaxy closest to the UV position. The resultant UV selected sample is composed of 156 galaxies in Virgo, 140 galaxies in Coma and 133 galaxies in Abell 1367.

3.3 The UV luminosity functions

Unlike the VCC catalog, the Coma and A1367 r0 catalogs used for star/galaxy dis- crimination do not cover all the area observed by FOCA but only the cluster cores. This reduces our analysis to an area of 1 square degrees ( 2.6 Mpc2) in Coma ∼ ∼ and 0.7 square degrees ( 1.8 Mpc2) in Abell 1367. Including∼ new spectroscopic ∼observations (Cortese et al. 2004), the redshift complete- ness of the UV selected sample reaches the 65% in Abell 1367, the 79% in Coma and the 83% in Virgo. The redshift completeness per bin of magnitude of each cluster is listed in Table 3.1. We remark that for M 16.5 (corresponding to the FOCA UV ≤ − magnitude limit in Coma and Abell1367), the redshift completeness of the Virgo clus- ter sample is 98%. As discussed by Cortese et al. (2003b), the general UV galaxy counts (Milliard et al. 1992) are uncertain and cannot be used to obtain a reliable subtraction of the back- ground contribution from the cluster counts. Therefore, in order to compute the cluster LF, we use the statistical approach recently proposed by De Propris et al. (2003) and Mobasher et al. (2003). We assume that the UV spectroscopic sample is ’representative’, in the sense that the fraction of galaxies that are cluster members is the same in the (incomplete) spectroscopic sample as in the (complete) photometric sample. For each magnitude bin i we count the number of cluster members NM , the number of galaxies with a measured recessional velocity NZ and the total number of galaxies NT . The completeness-corrected number of cluster members in each bin is:

NM NT Ni = (3.1) NZ 26 3. The FAUST-FOCA UV luminosity function of nearby clusters

Table 3.1: Integral redshift completeness in bin of 0.5 magnitudes.

Redshift completeness M Virgo Coma Abell1367 UV ≤ 21.75 100% − − − 21.25 100% 100% −20.75 100%− 100% 100% −20.25 100% 100% 100% − 19.75 92% 100% 100% −19.25 95% 100% 100% − 18.75 97% 100% 100% −18.25 97% 97% 100% −17.75 97% 95% 95% − 17.25 98% 84% 80% −16.75 98% 79% 65% −

NT is a Poisson variable, and NM is a binomial variable (the number of successes in NZ trials with probability NM /NZ ). Therefore the errors associated with Ni are given by: 2 δ Ni 1 1 1 2 = + (3.2) Ni NT NM − NZ The completeness-corrected number of cluster members obtained from (3.1) are given in Table 3.2 and the luminosity functions for the four studied samples are shown in Fig.3.1. The two different datasets used for the Virgo cluster have only one magnitude bin (MUV = 18.75) overlap. In this bin the two LFs are in agreement and there is no indication− that a change of slope occurs. We thus feel comfortable combining them into a composite Virgo UV luminosity function across 7 magnitudes. In order to determine whether the LFs of the three clusters are in agreement we 2 2 2 perform a two-sample χ test. We obtain P (χ χobs) 82% for the Virgo and 2 2 ≥ ∼ Abell1367 LFs, P (χ χobs) 87% for the Virgo and the Coma cluster LFs and P (χ2 χ2 ) 98% for≥ the Coma∼ and Abell1367 LFs, pointing out that the three ≥ obs ∼ LFs are in fair agreement within their completeness limits.

3.3.1 The composite cluster luminosity function The uncertainties of each individual cluster luminosity function are too large to fit a complete Schechter (Schechter 1976) function to the data and compare it with the field UV LF. However combining the three data-sets analyzed in this paper we 3.3. The UV luminosity functions 27

Figure 3.1: The UV luminosity functions for the four analyzed data sets.

compute the UV composite luminosity function of 3 nearby clusters. The composite LF is obtained following Colless (1989), by summing galaxies in absolute magnitude bins and scaling by the area covered in each cluster. The number of galaxies in the jth absolute magnitude bin of the composite LF (Ncj) is given by: 1 N N = ij (3.3) cj m A j Xi i

where Nij is the completeness-corrected number of galaxies in the jth bin of the ith cluster, Ai is the area surveyed in the ith cluster and mj is the number of clusters contributing to the jth bin. The errors in Nij are computed according to:

1 δN 2 1/2 δN = ij (3.4) cj m A j h Xi  i  i where δNij is the error in the jth bin of the ith cluster determined in (3.2). The weight associated to each cluster is computed according to the surveyed area, instead of the number of galaxies brighter than a given magnitude, as used by Colless (1989). The UV composite luminosity function is given in Fig.3.2 in the full magnitude range. However since for magnitudes fainter than MUV 16.5 the only available data are the Virgo FOCA observations, we fit the composite∼ −luminosity function with 28 3. The FAUST-FOCA UV luminosity function of nearby clusters

Figure 3.2: The composite UV luminosity function of 3 nearby clusters. The solid line represents the best Schechter fit to the data for M 16.5. UV ≤ − the Schechter functional form (Schechter 1976):

∗ ∗ 0.4(M −M ) ∗ 0.4(M −MUV)(α+1) −10 UV φ(MUV) = 0.4 ln 10 φ 10 e only for MUV 16.5, that is the completeness limit in Coma and Abell 1367. ≤ − ∗ The resulting Schechter parameters are MUV = 20.75 0.40 and α = 1.50 0.10. The faint-end slope is consistent within 1 σ −with thelower limit for Coma− and A1367 recently proposed by Cortese et al. (2003b), but significantly flatter than the slope α 2.0, 2.2 found for Coma by Andreon (1999), suggesting that this very ∼ − − steep luminosity function was due to an underestimate of the density of background galaxies.

3.4 Discussion

Although the UV(2000 A)˚ radiation is dominated by young stars of intermediate masses (2

Figure 3.3: The UV bi-variate composite luminosity functions of nearby clusters. Red (UV B > 2) and blue (UV B < 2) galaxies are indicated with empty and filled circles−respectively. − ever, based on the spectral energy distributions computed by Gavazzi et al. (2002a), we can use the total color UV B, available for the 94% of galaxies in our sample, to discriminate between red elliptical− (UV B > 2) and blue spiral (UV B < 2) galaxies. B magnitudes are taken from the−VCC (Binggeli et al. 1985), the−Godwin et al. (1983) catalog and the Godwin & Peach (1982) catalog for Virgo, Coma and Abell 1367 respectively. The bi-variate composite luminosity function derived for galaxies of known UV B − color is shown in Fig.3.3. It shows that the star forming galaxies dominate the UV LF for M 18, as Donas et al. (1991) concluded for the first time. Conversely, UV ≤ − for MUV 17.5, the number of red and blue galaxies is approximately the same, pointing out≥ −that, at low luminosities, the UV emission must be ascribed not only to star formation episodes but also to Post-Asymptotic Giant Branch (PAGB) low mass stars in early type galaxies (Deharveng et al. 2002). Similarly, if we restrict the anal- ysis to the fraction ( 50 %) of objects with known morphological type, we find that ∼ late-types (Sa or later) dominate at bright UV luminosities, while early-type objects contribute at the faint UV levels. Since Virgo and Abell1367 are spiral-rich clusters while Coma is spiral-poor, one might expect that the LFs of the three clusters ob- tained combining all types should have different shapes, contrary to the observations. The point is that the combined LF of the two types is dominated, at high UV lumi- 30 3. The FAUST-FOCA UV luminosity function of nearby clusters

Table 3.2: The completeness-corrected differential number of galaxies per bin of mag- nitude

MUV Ni mag Virgo Virgo Coma Abell 1367 (Faust) (Foca) 21.75 0 0 0 1 −21.25 0 0 1 0 − 20.75 2 0 0 1 −20.25 1 0 5 1 − 19.75 7 0 3 4 −19.25 9 0 3 4 −18.75 13 2 5 3 − 18.25 0 2 8.6 6 −17.75 0 3 7.7 6.7 − 17.25 0 3 15.8 10.1 −16.75 0 4 18.6 12.7 −

nosity by the spiral component, while at low luminosity early- and late-type galaxies contribute similarly. The UV LF of the spiral component are similar in the three clusters. At faint UV luminosities also the number density of early-type galaxies is approximately the same in the three clusters. Only at relatively high UV luminosity the number density of early-type galaxies in the Coma cluster exceeds significantly that of the other two clusters, but it is still much lower than the one of the late-type component. Therefore the LF obtained by combining early- with late-type galaxies results approximately the same in the three clusters. The cluster composite luminosity function has identical slope and similar M ∗ as the ∗ UV luminosity function computed by Sullivan et al. (2000) for the field: MUV = 21.21 0.13, α = 1.51 0.10, as shown in Fig. 3.4. This result is quite sur- prising− since we have just− shown that at low luminosity the contribution of ellipticals is not negligible, and early-type galaxies are expected to be more frequent in high density environments. This result seems in contradiction with recent studies of clus- ter galaxies carried out in Hα (Iglesias-P´aramo et al. 2002) and B-bands (De Propris et al. 2003). They find that the LFs of star forming galaxies in clusters and in the field have the same shape, contrary to early type galaxies in clusters that have a brighter and steeper LF than their field counterparts (De Propris et al. 2003). In order to understand this apparent difference between optical and UV luminosity functions we needed to wait the launch of GALEX and higher quality (and more homogeneous) 3.4. Discussion 31

Figure 3.4: The cluster and the field UV luminosity functions. The composite cluster LF is given with filled circles. The solid line indicates the best Schechter fit of the field LF of Sullivan et al. (2000). The normalization is such that the two LFs match at M 19.25. UV ∼ −

UV observations.

Chapter 4

GALEX UV luminosity function of Abell1367

4.1 Introduction

As I have shown in the previous Chapter, before the launch of the Galaxy Evolution Explorer (GALEX), the FOCA experiment allowed the first determinations of the UV LF of local field galaxies (Treyer et al. 1998; Sullivan et al. 2000) and of nearby clusters (Donas et al. 1991; Andreon 1999; Cortese et al. 2003b). Combining the FOCA and FAUST data Cortese et al. (2003a) determined the first composite LF of nearby clusters. They found no significant differences with the LF in the field. However this early determination was affected by large statistical errors due to the uncertainty in the UV background counts (Cortese et al. 2003b). GALEX has opened a new era of extragalactic UV astronomy. In particular it provides for the first time precise UV photometry of galaxies over large stretches of the sky (Xu et al. 2005), thus making the background subtraction method more reliable than in the past. Moreover its higher sensitiveness, higher resolution, large field of view make GALEX observations a unique homogeneous sample for statistical analysis of galaxies UV properties.

4.2 UV data

GALEX provides far-ultraviolet (FUV; λeff = 1528A˚, ∆λ = 442A)˚ and near-ultraviolet (NUV; λ = 2271A˚, ∆λ = 1060A)˚ images with a circular field of view of 0.6 de- eff ∼ grees radius. The spatial resolution is 5 arcsec. The data analyzed in this Chapter consist of two GALEX pointings of the∼Abell cluster 1367, with a mean exposure time of 1460s, , centered at R.A.(J2000)=11:43:41.34 Dec(J.2000)=+20:11:24.0 (e.g. offset to the north of the cluster to avoid a star bright enough to threaten the detector, see Fig.4.1). Sources were detected and measured using SExtractor (Bertin & Arnouts

33 34 4. GALEX UV luminosity function of Abell1367

Figure 4.1: The GALEX observation of Abell1367. ROSAT X-ray contour are su- perposed in black. The tick rectangular region indicates the region covered by the optical catalogues used for the star/galaxy discrimination.

1996). The 100% completeness limit is mAB 21.5 both in FUV and NUV (Xu et al. 2005). As the NUV images are significan∼ tly deeper than the FUV, sources were selected and their parameters determined in the NUV. FUV parameters were extracted in the same apertures. We used a larger SExtractor deblending parame- ter compared to the standard GALEX pipeline, providing reliable MAGAUTO also for very extended sources. The calibration uncertainty of the NUV and FUV magnitudes is 10% (Morrissey et al. 2005). Magnitudes are corrected for Galactic extinction using∼ the Schlegel et al. (1998) reddening map and the Galactic extinction curve of Cardelli et al. (1989). The applied extinction corrections are of 0.18 and 0.17 mag for the NUV and FUV bands respectively. To avoid artifacts present at the edge of the 4.2. UV data 35

Figure 4.2: Comparison between FOCA (upper image) and GALEX (lower image) observation of the center of Abell1367. It emerges clearly the strong improvement in resolution and sensitiveness of new GALEX data. 36 4. GALEX UV luminosity function of Abell1367

Figure 4.3: Left: The comparison between FOCA and GALEX NUV (left) and FUV (right) magnitudes of galaxies in Abell1367. The continuum line shows the best linear fit to the data.

field, we considered only the central 0.58 deg radius from the field center. A reliable star/galaxy discrimination was achieved by matching the GALEX catalog against the deepest optical catalogs available for A1367 (B < 22.5 mag and r0 < 21 mag; Iglesias-P´aramo et al. 2003), using a search radius of 6 arcsec, as adopted by Wyder et al. (2005) for the estimate of the GALEX local field LF. The optical catalogs do not include a negligible part ( 0.09 square degrees) of the GALEX field. A total number of 292 galaxies in the FUV∼and of 480 galaxies in NUV with m 21.5 are detected AB ≤ in the 0.96 square degrees field ( 2.5Mpc2) analyzed in this Chapter. Great part of the ∼field observed by GALEX co∼vers the area studied in the previous Chapter with FOCA observations. The two observations of the cluster center are presented in Fig.4.2: emerges clearly the strong improvement in resolution and sensitiveness of new GALEX data. In Fig.4.3 (left) we compared the UV magnitudes determined from FOCA and from GALEX NUV observations for the 96 galaxies detected by both instruments. The two sets of measurements are in satisfactory agreement. The linear regression between the two datasets is:

M (2310A)˚ = (1.02 0.03) M (2000A)˚ + (1.74 0.51) (4.1) GALEX  × FOCA  M (1530A)˚ = (1.04 0.04) M (2000A)˚ + (1.71 0.70) (4.2) GALEX  × FOCA  with a mean dispersion of 0.23 and 0.32 mag in NUV and FUV bands respectively, consistent with the error assumed in the previous Chapter for FOCA observations. 4.3. The luminosity function 37

Figure 4.4: The redshift completeness per bin of UV magnitude in Abell 1367.

Band Sample Schechter P arameters M ∗ α NUV A1367 19.77 0.42 1.64 0.21 NUV F ield −18.23  0.11 −1.16  0.07 −  −  FUV A1367 19.86 0.50 1.56 0.19 FUV F ield −18.04  0.11 −1.22  0.07 −  −  UV(2000A)˚ Composite cluster 18.79 0.40 1.50 0.10 −  − 

Table 4.1: Best Fitting Parameters.

4.3 The luminosity function

The determination of the cluster LF requires a reliable estimate of the contribu- tion from background/foreground objects to the UV counts. This can be accurately achieved for mAB 18.5, since at this limit our redshift completeness is 90 % (Cortese et al. 2003b,≤ 2004; see Fig. 4.4). The redshift completeness drops∼rapidly at magnitudes fainter than mAB 18.5, thus requiring the contamination to be esti- mated statistically. Two methods∼are usually applied for the computation of cluster LFs. The first one is based on the statistical subtraction of field galaxies, per bin of UV magnitude, that are expected to be randomly projected onto the cluster area, as derived by Xu et al. (2005). Alternatively, the completeness corrected method proposed by De Propris et al. (2003) is to be preferred when the field counts have large uncertainties. It is based on the assumption that the UV spectroscopic sample 38 4. GALEX UV luminosity function of Abell1367

Figure 4.5: The GALEX NUV (left) and FUV (right) LF for Abell 1367. Open dots are obtained using the subtraction of field counts obtained by Xu et al. (2005); filled dots are obtained using the completeness corrected method. The solid line represents the best Schechter fit. The dotted line shows the composite nearby clusters 2000 A˚ LF by (Cortese et al. 2003a). The dashed line represents the GALEX local field LF (Wyder et al. 2005), normalized in order to match the cluster LF at M 17.80. AB ∼ −

Figure 4.6: The NUV (left) and FUV (right) bi-variate LF of A1367. Star-forming and quiescent galaxies are indicated with empty triangles and filled squares respectively. The dashed line represents the GALEX local field LF (Wyder et al. 2005), normalized as in Fig.4.5 4.3. The luminosity function 39

(e.g. membership confirmed spectroscopically) is ’representative’ of the entire cluster, i.e. the fraction of galaxies that are cluster members is the same in the (incomplete) spectroscopic sample as in the (complete) photometric one. For each magnitude bin i we count the number of cluster members NM (i.e. galaxies with velocity in the range 4000 0 A)˚ and quies- cent (EW (Hα) = 0 A)˚ objects. Unfortunately neither UV field counts for different morphological types nor a measure of EW (Hα) for all the UV selected galaxies are available. Thus we can only apply the completeness corrected method to determine the bi-variate LFs. We assume that in each bin of magnitude the fraction of star- forming and quiescent cluster members is the same in the (incomplete) spectroscopic sample as in the (complete) photometric sample. The bi-variate LFs derived by this method are shown in Fig.4.6. 40 4. GALEX UV luminosity function of Abell1367

4.4 Discussion

As shown in Fig.4.5, the GALEX LFs have a shape consistent with the composite LF of nearby clusters as constructed in the previous Chapter (see also Cortese et al. (2003a)). Conversely, whatever fitting procedure one adopts, they show a steeper faint-end slope and a brighter M ∗ than the GALEX field LF recently determined by Wyder et al. (2005). In fact the GALEX local field luminosity function shows a faintest bright-end and a flatter faint end than the previous determination by Sullivan et al. (2000), but the reason for this difference is not yet clearly understood. Wyder et al. (2005) argued that magnitudes estimated by FOCA are on average brighter than the GALEX one, with the difference becoming larger for fainter sources; sug- gesting that these offsets and nonlinearities in the FOCA photometry could account for a major part of the observed difference between the two field luminosity func- tions. However we have shown that this seems not the case at least for Abell1367 observations. On the contrary I think that part of the problem could be due not to different photometric estimates but to the different areas used by GALEX and FOCA to estimate the field LF. In the case of FOCA, Treyer et al. (1998) and Sullivan et al. (2000) used the pointing of Abell 1367 to estimate the field LF: thus a partial contam- ination of cluster galaxies could explain why the FOCA field and cluster LF results very similar. The brighter M ∗ observed in Abell1367 is probably to be ascribed to the particular galaxy population of this cluster. In fact Abell 1367 is a young cluster of galaxies composed of at least four dynamical units at the early stage of a multiple merging event (see Chapter 5 and Cortese et al. 2004). Some galaxies have their star formation enhanced due to interaction with the cluster environment, and it is this population that is responsible for the bright M ∗ observed in this cluster. Conversely the high faint-end slope observed in this cluster is due to the significant contribution of non star-forming systems at faint UV magnitudes. In fact, as shown in Fig.4.6, star-forming galaxies dominate the UV LF for M 17 mag, as Donas AB ≤ − et al. (1991) concluded for the first time. For MAB 16 mag however, the number of red galaxies increases very rapidly1. This result is≥consisten− t with an UV LF con- structed starting from the r0 LF computed by Iglesias-P´aramo et al. (2003): assuming a mean color NUV r0 1 mag and NUV r0 5 for star-forming and quiescent galaxies respectively−, we ∼are able to reproduce−the∼contribution, at low UV luminosi- ties, of elliptical galaxies. Moreover the difference observed between NUV and FUV cluster LFs can be understood looking at the FUV-NUV color magnitude relation (computed only for confirmed cluster members) shown in Fig.4.7. The star-forming objects dominate at high UV luminosities while the quiescent systems contribute more

1The bi-variate LFs cannot be compared with the ones computed by Treyer et al. (2005) for the field, since their samples do not contain ellipticals but only spiral galaxies. 4.4. Discussion 41

Figure 4.7: The FUV-NUV color magnitude relation for confirmed members of A1367. Symbols are as in Fig.4.6

Figure 4.8: The optical (r0-band) distribution for star forming (blue histogram) and quiescent (red histogram) galaxies in our sample. 42 4. GALEX UV luminosity function of Abell1367

at faint magnitudes. Their mean FUV-NUV color is 1.5 mag thus they influence the LF at higher luminosities in the NUV than in the FUV∼ (see Fig.4.6). The optical luminosity distribution of star forming and quiescent systems, presented in Fig. 4.8, points out clearly that early type galaxies contributing to the UV faint end slope are the giant, optically bright, galaxies that dominate the bright end of the optical luminosity functions. This means that, in UV, the steeper faint end slope observed in clusters is only due to the contribution of giant ellipticals and not of dwarf elliptical galaxies, as observed at optical wavelengths. We can thus conclude that, in clusters, a significant fraction of the low luminosity UV emission comes massive early type galaxies. This result is expected since in the field the fraction of quiescent systems is significantly lower than that of star forming objects (Dressler 1980; Whitmore et al. 1993), thus their contribution to the LF is negligible. Moreover, the UV emission of ellipticals has a different nature from the one emitted by star forming systems. In fact it does not arise from newly born stars but from low mass old post asymptotic giant branch stars (O’Connell 1999), as I will discuss in depth in Chapter 6. Finally, the LFs of cluster star-forming systems have a faint-end slope (α 1.25 ∼ −  0.2) consistent within the statistical uncertainties with the GALEX field LF. The similar shape observed in the LF of star forming galaxies in different environments goes in the same direction with recent studies of cluster galaxies carried out in Hα (Iglesias-P´aramo et al. 2002) and B-bands (De Propris et al. 2003). They find that the LFs of star forming galaxies in clusters and in the field have the same shape, contrary to early type galaxies in clusters that have a brighter and steeper LF than their field counterparts (De Propris et al. 2003). This indicates that, whatever mech- anism (i.e. ram pressure, tidal interaction, galaxy harassment) quenches/enhances the star formation activity in late-type cluster galaxies, it influences similarly and with a short time scale the giant and the dwarf components , so that the shape of their LF is unchanged and only the normalization is modified. Chapter 5

Multiple merging in Abell1367

5.1 Introduction

Clusters of galaxies represent the most massive gravitationally bound systems in the Universe. They provide us with valuable insights into the formation of large-scale structures, as well as into the formation and evolution of galaxies. The hierarchi- cal model predicts that galaxy clusters are formed by accretion of units of smaller mass at the nodes of large-scale filaments (West et al. 1991; Katz & White 1993). Statistical analyses of clusters have shown that even at low redshift a high fraction of clusters presents substructures, implying that clusters are still dynamically young units, undergoing the process of formation (Dressler & Shectman 1988). The Abell cluster 1367 (z 0.0216) lies at the intersection of two filaments, the first extending roughly 100 Mp∼c from Abell 1367 toward Virgo (West & Blakeslee 2000), the second connecting Abell 1367 to Coma (as a part of the Great Wall, Zabludoff et al. 1993). With its irregular X-ray distribution (Jones et al. 1979; Bechtold et al. 1983; Grebenev et al. 1995), high fraction of spiral galaxies and low central galaxy density, Abell 1367 can be considered as the prototype of a nearby dynamically young cluster. ASCA X-ray observations pointed out the existence of a strong localized shock in the intra-cluster medium (ICM) suggesting that Abell 1367 is experiencing a merging between two substructures (Donnelly et al. 1998). Moreover recent Chandra obser- vations (Sun & Murray 2002), and a preliminary analysis of the XMM data (Forman et al. 2003), indicate the presence of cool gas streaming into the cluster core, sup- porting a multiple merger scenario. Optical and radio observations also suggest that this cluster is currently experiencing galaxy infall into its center. Gavazzi et al. (1995, 2001a) discovered two head-tail radio sources associated with disk galaxies with an excess of giant HII regions on their leading edges, in the direction of the cluster center. The observational scenario

43 44 5. Multiple merging in Abell1367

is consistent with the idea that ram-pressure (Gunn & Gott 1972) is, for a limited amount of time, enhancing the star formation of galaxies that are entering the cluster medium. In addition Gavazzi et al. (2003b) pointed out the existence of a group of star bursting galaxies infalling into the cluster core. Although X-ray, radio and optical observations suggest that Abell 1367 is dynami- cally young and it is still undergoing the process of formation, detailed spatial and dynamical analysis of this cluster has not been attempted so far. Girardi et al. (1998) detected a secondary peak in the cluster velocity distribution, suggesting that Abell 1367 is a binary cluster, but their analysis was based on 90 redshifts, insufficient for drawing a detailed model of the cluster kinematics. ∼ Cortese et al. (2003b) carried out a deep (r0 < 20.5) spectroscopic survey of the cen- tral 1.3 square degrees of Abell 1367 adding 60 new spectra (33 members). Here I presen∼t new measurements for 119 galaxies (adding another 33 cluster members). In total 273 redshifts were measured in the region, out of which 146 are cluster members, allowing the first detailed dynamical analysis of Abell 1367.

5.2 Observations and data reduction

The cluster region analyzed in this Chapter covers an area of 1.3 square degrees centered at α(J.2000) = 11h44m00s δ(J.2000) = 19d43m30s. r0 ∼imaging material was used to extract a catalogue of galaxy candidates in Abell 1367 complete to r0 20.5 mag, and to select the targets of the present spectroscopic survey. Spectroscop∼ y of Abell 1367 was obtained with the AF2-WYFFOS multi fiber spectrograph at the 4.2m William Herschel Telescope (WHT) on La Palma (Spain) during 2003, March 27-29. WYFFOS has 150 science fibers of 1.6 arcsec diameter coupled to a bench- mounted spectrograph which relies on a TEK 1024 1024 CCD. The 316R grating was used, giving a dispersion of 240 A˚/mm, a resolution× of 6A˚ FWHM, and a total spectral coverage of 5600∼A.˚ The spectra were centered∼ at 6500A,˚ thus ∼ ∼ covering from 3600 A˚ to 9400 A.˚ We allocated typically 70 objects to fibers in a given configuration and, on average, 15 sky fibers. A total∼ of 4 configurations were executed, with an exposure time of 4x1800 sec for each configuration. Argon lamps for wavelength calibration were obtained for each exposure. The reduction of the multi fiber spectra was performed in the IRAF1 environment, using the IMRED package. After bias subtraction, the apertures were defined on dome flat-field frames and used to the spectra on the CCD. The arc spectra were extracted and matched with arc lines to determine the dispersion solution. The

1IRAF is distributed by the National Optical Astronomy Observatories, which is operated by the Association of Universities for Research in Astronomy, Inc., under the cooperative agreement with the National Science Foundation 5.2. Observations and data reduction 45

Observatory Dates N. gal. Spectrograph Dispersion Coverage CCD pix A˚/mm A˚ µm WHT March 03 98 AF2-WYFFOS 240 3600-9400 1024 × 1024 T EK 24 Cananea March 03 12 LFOSC 228 4000-7100 576 × 384 T H 23 Loiano March 03 - Feb. 04 9 BFOSC 198 3600-8900 1340 × 1300 EEV 20

Table 5.1: The spectrograph characteristics rms uncertainty of the wavelength calibration ranged between 0.1 and 0.3 A.˚ The lamps’ wavelength calibration was checked against known sky lines. These were found within 0.5 A˚ of their nominal position, providing an estimate of the systematic un- certaint∼y on the derived velocity of 25 km s−1. The object spectra were extracted, ∼ wavelength calibrated and normalized to their intensity in the interval 5400-5600 A.˚ A master sky spectrum, that was constructed by combining various sky spectra was normalized to each individual science spectrum and then subtracted from it. Unfor- tunately strong sky residuals were left after this procedure, limiting the number of useful spectra to 98 (as listed in Tab. 5.9). Nine additional long-slit, low dispersion spectra were obtained in March 2003 and in February 2004 using the imaging spectrograph BFOSC attached to the Cassini 1.5m telescope at Loiano (Italy). Another twelve spectra were taken with LFOSC at the 2.1m telescope of the Guillermo Haro Observatory at Cananea (Mexico). These ob- servations were performed using a 2.0 arcsec slit and the wavelength calibration was secured with exposures of HeAr and XeNe lamps at Loiano and Cananea respectively. The on-target exposure time ranged between 15 and 30 min according to the bright- ness of the targets. After bias subtraction, when 3 or more frames of the same target were obtained, these were combined (after spatial alignment) using a median filter to help cosmic rays removal. Otherwise the cosmic rays were removed using the task COSMICRAYS and/or under visual inspection. The lamps wavelength calibration was checked against known sky lines. These were found within 1 A˚ from their nominal position, providing an estimate of the systematic uncertain∼ty on the derived velocity of 50 km s−1. After subtraction of sky background, one-dimensional spec- tra were extracted∼ from the frames. The redshift were obtained using the IRAF FXCOR Fourier cross-correlation (Tonry & Davis 1979) task, excluding the regions of the spectra affected by night-sky lines. Moreover all the spectra and their best correlation function were visually examined to check the redshift determination. Table 5.2 lists the characteristics of the instrumentation in the adopted set-up. The 119 new velocity measurements presented in this Chapter are listed in Table 5.9 as follows: Column 1: Galaxy designation. 46 5. Multiple merging in Abell1367

Figure 5.1: Cumulative redshift distribution for galaxies in the studied region.

Column 2, 3: (J2000) celestial coordinates, measured with few arcsec uncertainty. Column 4: r0 band magnitude. Column 5: observed recessional velocity. Column 6: telescope (WHT=William Herschel Telescope; LOI=Loiano; CAN=Cananea) Combining the new set of 119 redshifts (given in Tab. 5.9) with the ones available from the literature (NED; Cortese et al. 2003b; Rines et al. 2003), we have the redshift for 273 galaxies of which 146 are cluster members (4000 km s−1 V 10000 km s−1). The cumulative redshift distribution, in the observed area, ≤as a ≤function of the r0 magnitude is shown in Fig.5.1. The completeness is 70% at r0 < 17.5, and it drops to 45% at r0 < 18.5. ∼ ∼

5.3 The global velocity distribution

The line of sight (LOS) velocity distribution for the 146 cluster members is shown in Fig. 5.2. The mean and standard deviation are known to be efficient estimators of the central location and scale when the underlying population is gaussian. Unfortu- nately they are not minimum variance estimators when the nature of the observed population is significantly non-Gaussian. The best location and scale estimators must be resistant to the presence of outliers and robust to a broad range of non-Gaussian underlying populations. Thus, following Beers et al. (1990), we consider the biweight estimator as the best estimator of location (CBI ) and scale (SBI ) of the cluster ve- 5.3. The global velocity distribution 47

Figure 5.2: Velocity histogram and stripe density plot for the members of Abell 1367. Arrows mark the location of the most significant weighted gaps in the velocity distribution. locity distribution. We find a location C = 6484 81 km s−1 and a scale S = 891 58 km s−1, BI  BI  in agreement with previous studies (e.g. Girardi et al. 1998; Struble & Rood 1999). Visual inspection of Fig. 5.2 suggests that the velocity distribution differs from a Gaussian, a deviation that should be quantified using appropriate statistical tests. We analyze the higher moments of the distributions using the kurtosis and the skew- ness shape estimators. Kurtosis indicates a difference in the tails length compared to a Gaussian (positive kurtosis is indicative of long tails). Skewness indicates asym- metry (positive skewness implies that the distribution is depleted from values lower than the mean location, conversely negative skewness denotes a depletion of values higher than the mean). In addition we calculate the asymmetry index (AI) and tail index (TI) introduced by Bird & Beers (1993) as alternatives to the distribution higher moments. These indicators measure the shape of a distribution but, contrary to skewness and kurtosis, which depend on the estimate of the location and the scale of the underlying distri- bution, they are based on the order statistics of the dataset. The AI measures the symmetry in a population by comparing gaps in the data on the left and right sides of the sample median. The TI compares the spread of the dataset at 90% level to the spread at the 75% level. The kurtosis, skewness and the TI reject a Gaussian distribution with a confidence 48 5. Multiple merging in Abell1367

Test Value Rejection of a gaussian AI -0.077 80 % ≤ TI 1.240 >99 % Skewness 0.269 >99 % Kurtosis 2.680 >99 % W 0.963 98.7 %

Table 5.2: 1D substructure indicators for the whole cluster sample level of 99%, suggesting that the cluster velocity distribution has longer tails than ≥ a Gaussian of the same dispersion. Moreover, in order to assess the normality of the velocity distribution, we use the Wilk - Shapiro (W) test (Yahil & Vidal 1977). Con- trary to the χ2 and Kolmogorov Smirnov, this test does not require any hypothesis on the mean and variance of the normal distribution. The W test rejects normality with a confidence level of 98.7%, in agreement with kurtosis, skewness and TI (see Table 5.3). The departure from a normal distribution could result from a mixture of several ve- locity distributions with different location and smaller velocity dispersion than the whole sample; thus, using the program ROSTAT (Beers et al. 1990), we investigate the presence of significant gaps (Beers et al. 1991) in the velocity distribution, indi- cating subclustering. A weighted gap is defined by:

1/2 yi = i(N i) (xi+1 xi)  − ∗ −  where N is the number of values in the dataset. A weighted gap is significant if its value, relative to the midmean (the mean of the central 50% of the dataset) of the other weighted gaps, is greater than 2.25. This value corresponds to a probability of occurrence in a normal distribution of less than 3%. We detected six significant weighted gaps in the Abell 1367 velocity distribution. The stripe density plot of radial velocities and the position of each gap (indicated with an arrow) are shown in Fig. 5.2. The velocity of the object preceding each gap, the normalized size of the gap and the probability of finding a normalized gap of the same size and position in a normal distribution are listed in Table 5.3.

5.4 Localized velocity structures

Given the non-Gaussian nature of the velocity distribution, we looked for spatially localized variations in the LOS velocity and velocity dispersion distributions. First of all we applied the three 3D tests commonly used to quantify the amount of substruc- 5.4. Localized velocity structures 49

Velocity Gap Significance km s−1 5742 2.53 1.40% 5835 2.66 1.40% 6619 2.90 0.60% 6880 2.64 1.40% 7059 3.01 0.20% 7542 2.33 3.00%

Table 5.3: The most significant weighted gaps detected in the velocity distribution of the whole cluster sample.

tures in galaxy clusters: the ∆ test (Dressler & Shectman 1988), the α test (West & Bothun 1990) and the  test (Bird 1994). The ∆ test is based on the comparison of the local mean velocity, Vlocal, and the velocity dispersion, σlocal, associated to each cluster member (computed using its 10 nearest neighbors) with the mean velocity V , and dispersion σ, of the whole galaxy sample. For each galaxy, the deviation is defined by: 11 δ2 = [(V V )2 + (σ σ)2] σ2 local − local − The observed cumulative deviation ∆, defined as the sum of the δ’s for the cluster members, is used to quantify the presence of substructures. As shown by Pinkney et al. (1996) for samples with no substructures, the value of ∆ is approximately equal to the total number of galaxies, while it is larger in the presence of substructures. The α test measures how much the centroid of the galaxy distribution shifts as a result of correlations between the local kinematics and the projected galaxy distribution. The centroid of the whole galaxy distribution is defined as:

1 N 1 N x = x y = y c N i c N i Xi=1 Xi=1 For each galaxy i and its 10 nearest neighbors in the velocity space, the spatial centroid is defined as:

11 11 xj/σj yj/σj xi = j=1 yi = j=1 c P 11 c P 11 j=1 1/σj j=1 1/σj P P 50 5. Multiple merging in Abell1367

Indicator Value Prob. of substructures ∆ 206.5 99.8 % α 0.161 Mpc 55.7 % 13  5.44 10 M 68.4 %

Table 5.4: 3D substructure indicators for our sample

where σj is the velocity dispersion for galaxy j and its 10 nearest neighbors in projec- tion. Finally the presence of substructures in the cluster sample is quantified using the α statistic defined as:

1 N α = [(xi x )2 + (yi y )2]1/2 N c − c c − c Xi=1 which represents the mean centroid shift for the galaxy cluster. The higher the value of α, the higher the probability of substructures. The  test quantifies the correlations between the position and the projected mass estimator (Heisler et al. 1985), defined as:

32 N M = v2 r P ME πGN zj j   Xj=1

where vzj is the radial peculiar velocity with respect to the nearest neighbors group (composed by a galaxy and its 10 nearest neighbors) and rj is the projected distance from the center of the nearest neighbor group. The substructure statistic is then defined as: 1 N  = M N P ME gal Xi=1 which represents the average mass of the nearest neighbors groups in the cluster. Since galaxies in the nearest neighbors groups have small projected separations,  is generally smaller than the global mass estimate.  is lower for a cluster with sub- structures than for a relaxed system. The value and the significance of the above tests are listed in Table 5.4. These statis- tical tests are calibrated using 1000 Monte Carlo simulations that randomly shuffle the velocity of galaxies, keeping fixed their observed position, thereby destroying any existing correlation between velocity and position. The probability of subclustering is then given as the fraction of simulated clusters for which the test value is lower (larger for the  test) than the observed one. Assuming that these tests reject the null 5.5. The cluster dynamics 51

Figure 5.3: Local deviations from the global kinematics for galaxies in Abell 1367 as measured by the Dressler & Shectman (1988) test. Galaxies are marked with open circles whose radius scales with their local deviation δ from the global kinematics. The ROSAT X-ray contours are shown with dotted lines. hypothesis if the confidence level is greater than 90%, only the ∆ test finds evidence of substructures (see Table 5.4). The local deviations from the global kinematics as measured by the ∆ test are shown in Fig 5.3. The positions of galaxies are marked with open circles whose radius scales with their local deviation δ from the global kinematics. The presence of a substructure with a high deviation from the global cluster kinematic is evident projected near the cluster core. More insights on the cluster dynamical state can be achieved by comparing the results of the one and three dimensional statistical tests with the N-body simulations per- formed by Pinkney et al. (1996). These authors analyzed how the significance level of statistical tests of substructure varies in different cluster merging scenarios. The deviation of the velocity distribution from a Gaussian and the detection of substruc- ture provided by the ∆ test suggest that Abell 1367 is in the early merging stage, 0.2 Gyr before core crossing. ∼

5.5 The cluster dynamics

The analysis of the galaxy distribution, of the local mean LOS velocity and of the velocity dispersion give further insight onto the cluster structure. The iso-density map of the cluster members (computed using the 10 nearest neighbors to each point) 52 5. Multiple merging in Abell1367

Figure 5.4: Palomar DSS image of the central region ( 1.3 square degrees) of Abell ∼ 1367 studied in this Chapter. The iso-density contours for the 146 confirmed cluster members are superposed. The lowest iso-density contour correspond to 3σ above the mean density in the field (left). The ROSAT X-ray contours are superposed in red (right). The straight line indicates the position of the abrupt gas temperature gradient detected by ASCA (Donnelly et al. 1998), used to divide our sample into two subclusters: the North-West and the South-East.

is shown in Fig.5.4 (left). The galaxy distribution appears elongated from north-west to south-east with two major density peaks. The highest density region corresponds approximately to the center of the NW X-ray substructure detected by ROSAT (Don- nelly et al. 1998), while the secondary density peak is slightly offset from the X-ray cluster center (α(J.2000) = 11h44.8m δ(J.2000) = 19d42m, Donnelly et al. 1998). Moreover the south galaxy density peak roughly coincides with the substructure de- tected by the ∆ test (see Fig.5.3) and with the infalling group of star-forming galaxies studied by Gavazzi et al. (2003b). The iso-density contours superposed on the ROSAT X-ray contours are shown in Fig.5.4 (right). The region between the two major density peaks coincides with the strong gradient in the gas temperature (see the straight line in Fig.5.4, right) observed for the first time by ASCA (Donnelly et al. 1998) and recently confirmed by Chandra (Sun & Murray 2002). This abrupt temperature change is strongly suggestive of a shock which has generated during a collision between two substructures, probably as- sociated with the SE and the NW galaxy density peaks. In fact N-body simulations show that temperature structures and X-ray morphology similar to the one observed in Abell 1367 are typical of clusters at an early merging phase ( 0.25 Gyr before core crossing) (Schindler & Mueller 1993; Gomez et al. 2002). ∼ The merging scenario is further supported by the LOS velocity and velocity disper- sion fields (computed using the 10 nearest neighbors to each point) shown in Fig. 5.5. The cluster dynamics 53

Figure 5.5: The LOS velocity field (left) and the velocity dispersion field (right) for the whole region studied in this Chapter. The LOS velocity and the velocity dispersion are computed using the 10 nearest neighbors to each pixel, whose size is 36 arcsec2. The iso-density contours for the 146 confirmed cluster members are superposed in black.

5.5. The SE subcluster has higher LOS velocity and velocity dispersion than the NW substructure. The region with the highest LOS velocity and velocity dispersion lies 6 arcmin N from the X-ray cluster center and it coincides with the substructure ∼detected by the ∆ test. This result points out the presence of a group of galaxies infalling in the SE cluster core (see Sec.5.5.2). Thus the NW subcluster appears as a relaxed system with the lowest velocity disper- sion among the whole sample; on the other hand the SE subcluster appears far from relaxation, and it is probably experiencing a multiple merging event. We use the position of the gas temperature gradient, shown by the straight line in Fig.5.4 (right), to divide our sample into two regions and to study separately the dynamical properties of the two subclusters. A sketch of the cluster dynamical model discussed in the next section is given in Fig.5.6.

5.5.1 The North-West subcluster The NW subcluster is composed of 86 galaxies and includes two density peaks: the highest and a secondary one located at the western periphery of the subcluster (la- beled as W subcluster in Fig.5.6), with a weak X-ray counterpart. It has a similar −1 −1 mean location (CBI = 6480 87 km s ) and a lower scale (SBI = 770 60 km s ) than the whole cluster.   54 5. Multiple merging in Abell1367

Figure 5.6: A 3D sketch of Abell 1367 summarizing the various sub-components described in Section 5.5. The cluster is viewed from its near side, as suggested by the eyeball indicating the observer’s position.

Fig.5.8 shows the LOS velocity distribution of this subcluster. The W test rejects the Gaussian hypothesis at a confidence level of 39%. Thus the LOS velocity dis- tribution is consistent with a Gaussian distribution, suggesting that this subcluster is a virialized system. Moreover its increasing velocity dispersion profile (see Fig. 5.9) is consistent with a relaxed cluster undergoing two body relaxation in the dense central region, with circular velocities in the center and more isotropic velocities in the external regions (Girardi et al. 1998). However this subcluster also shows some evidences of merging (see Fig.5.7). The brightest galaxy of this cloud CGCG97-095 (NGC3842), located 2 arcmin SE from ∼ the NW density peak, is a radio galaxy classified as a narrow-angle tail (NAT) (Bliton et al. 1998). The tail orientation (indicated with an arrow in Fig. 5.7) suggests that this galaxy (and the associated substructure) is moving from north-west to south- east, toward the main cluster core. Moreover two CGCG (Zwicky et al. 1961) galaxies, 97-073 and 97-079, show fea- 5.5. The cluster dynamics 55

Figure 5.7: Blow-up of the NW substructure of Abell 1367. The arrows indicate the direction of radio head tails associated with 97-079 and 97-073 and the orientation of the NAT radio galaxy 97-095. The dashed region shows the distribution of the diffuse cluster radio relic (Gavazzi 1978). The iso-density contours for the 146 confirmed cluster members are superposed. tures consistent with the infall scenario. Gavazzi et al. (1995, 2001a) found that both galaxies have their present star formation enhanced along peripheral HII re- gions which developed at the side facing the direction of motion through the cluster IGM. Their neutral hydrogen is significantly displaced in the opposite side (Dickey & Gavazzi 1991), where 50 kpc long tails are detected both in the light of the syn- chrotron radiation (Gavazzi & Jaffe 1987) and in Hα (Gavazzi et al. 2001a). The observational scenario is consistent with the idea that ram-pressure (Gunn & Gott 1972) is enhancing for a limited amount of time the star formation of galaxies that are entering the cluster medium for the first time. However these two galaxies appear not directly associated with the center of the NW subcluster since they lie at a projected distance of 0.34 Mpc from the main density peak (see Fig.5.7). Moreover their large distance (∼0.48 Mpc) from the shock front ∼ observed in X-ray between the NW and the SE substructure indicates that these ob- jects do not belong to the main galaxy density peak infalling into the cluster center. Conversely they are at a projected distance of only 0.08 Mpc from the center of the W subcluster, suggesting that they are associated with this subcloud. For these reasons we consider an alternative scenario in which these two galaxies be- long to a secondary substructure infalling into the NW substructure from the western side (see Fig. 5.6). This picture is supported by the presence of the extended radio relic detected both in X-ray and radio continuum in this region (Gavazzi 1978; Gavazzi & Trinchieri 1983). Cluster radio halos contain fossil radio plasma, the former outflow of a radio galaxy, that has been revived by shock compression during cluster merging 56 5. Multiple merging in Abell1367

Figure 5.8: The LOS velocity distribution for galaxies in the NW (upper) and in the SE (lower) subclusters.

(Enßlin et al. 1998; Enßlin & Bruggen¨ 2002). The radio relic observed in Abell 1367 extends, south-west to north-east, from 97-073 to 127-040 with a projected extent of 0.8 Mpc (see Fig.5.7). The age of its electrons is estimated to be 0.2 Gyr (Enßlin et al. 1998). The only plausible source of high energy electrons av∼ailable in this re- gion is the NAT galaxy 97-095, presently at 0.25 Mpc from the relic and whose tails ∼ point exactly in the relic direction. Assuming that the fossil radio halo originated from 97-095, we find that the infall velocity of this galaxy into the SE subcluster is V 1250 km s−1, consistent with the typical infall velocity of cluster galaxies. Thus the∼presence of the radio relic results consistent with a merging scenario in which the W subcluster, containing 97-079 and 97-073, is infalling into the NW substruc- ture, compressing the plasma ejected from 97-095 and re-accelerating the electrons to relativistic energies.

5.5.2 The South-East subcluster The SE cloud is composed of 60 galaxies associated with the X-ray cluster center. It has the highest LOS velocity and dispersion of the whole sample (see Fig.5.5) with a location C = 6596 137 km s−1 and a scale S = 1001 70 km s−1. Its velocity BI  BI  distribution, shown in Fig. 5.8, appears significantly non-Gaussian. The W test re- jects the Gaussian hypothesis at a confidence level of 96.8%, supporting the idea that 5.5. The cluster dynamics 57

Figure 5.9: The velocity dispersion radial profile of the NW (upper) and the SE (lower) subclusters. the cluster center is far from relaxation. This is in agreement with the decreasing velocity dispersion profile of this region (see Fig.5.9), consistent with isotropic veloci- ties in the center and radial velocities in the external regions, as expected in the case of galaxy infall onto the cluster (Girardi et al. 1998). The velocity distribution of Fig. 5.8 has three peaks at 5500 km s−1, 6500 km s−1 − ∼ ∼ and 8200 km s 1 respectively, probably associated with three separate groups. Moreo∼ver we remark that the galaxy gaps between the three peaks are fairly consis- tent with two of the most significant weighted gaps detected in the global velocity distribution (V 5800 km s−1 and V 7500 km s−1). In order to check∼for any position-velocit∼y segregation, we divide the SE subcluster in three groups according to their LOS velocity: galaxies with V < 5800 km s−1 belong to the low velocity group, galaxies with V > 7500 km s−1 belong to the high velocity group and galaxies with intermediate velocity belong to the SE subcluster. The projected distribution of the three groups is shown in Fig.5.10. The high-velocity group (V 8200 km s−1, triangles) appears segregated in the northern part of the SE ∼ cloud, extending 20 arcmin in right ascension but only 7 arcmin in declination. It is associated with∼ the substructure detected by the ∆ test∼ (see Fig. 5.3) and with the infalling group of star-forming galaxies recently discovered by Sakai et al. (2002) and by Gavazzi et al. (2003b). Its spatial segregation and high star formation activity suggest that this group is a separate unit infalling into the cluster, probably from the 58 5. Multiple merging in Abell1367

Figure 5.10: The distribution of galaxies belonging to the South-East subcluster. Triangles indicate galaxies with LOS velocity > 7500 km s−1, circles galaxies with LOS velocity < 5800 km s−1 and squares galaxies with LOS velocity comprises in the range 5800 km s−1 < V < 7500 km s−1. The ROSAT X-ray contours are shown. near side (see Fig. 5.6). It is remarkable that Sun & Murray (2002), using Chandra observations of the cluster center, discovered a ridge-like structure around the cluster center, 6 arcmin south from the center of the high velocity group, probably asso- ∼ ciated with a compact merging subcluster (perhaps this group) penetrating the SE cluster core. The low-velocity group (V 5500 km s−1, circles in Fig.5.10) seems segregated in ∼ the eastern part of the cloud, perhaps infalling from the eastern side into the cluster core (Fig. 5.6). This scenario is also supported by the detection of cool gas streaming into the cluster core from the eastern side (Forman et al. 2003), probably associated with this low velocity group of galaxies. Galaxies with V 6500 km s−1 (squares in Fig.5.10) are homogeneously distributed ∼ over the SE subcluster, representing its virialized galaxy population. However the brightest galaxy in this group 97-127 (NGC3862) is a NAT radio galaxy with very extended radio tails pointing in the direction of the low velocity group (Gavazzi et al. 1981), suggesting motion relative to the IGM. The velocity-space segregation observed in the SE subcluster suggests that the cluster center is experiencing multiple merging of at least two separate groups, supporting the idea that it is far from relaxation. This picture is consistent with the high gas 5.6. Star formation activity in the infalling groups 59

Figure 5.11: The LOS velocity distribution for emission line (upper) and non emission line galaxies (lower) in the whole cluster sample.

entropy in this region, since in absence of a cool dense core the substructures infalling into the major cluster can penetrate deep inside, disturbing the cluster core dynamics (Churazov et al. 2003). A sketch of the various substructures identified in Abell 1367 by the present study, is given in Fig. 5.6. Five substructures are detected. Two clouds, the NW and SE sub- clusters, are in the early merging phase, meanwhile three smaller groups are infalling into Abell 1367. The W subcloud, associated with the head-tail systems 97-073/79, is probably infalling into the NW subcluster, exciting the radio relic observed in be- tween the two structures. The other two groups are infalling into the SE subcluster: the low velocity group from the eastern side, while the high velocity group from the near side.

5.6 Star formation activity in the infalling groups

The dynamical study presented in the previous sections indicates that Abell 1367 is a dynamically young cluster in the early stage of a multiple merging event involving at least five substructures. Since merging is expected to trigger star formation in cluster galaxies (Bekki 1999), we study separately the spatial and velocity distribution of the star forming galaxies. Only 49 out of the 146 cluster members show recent star 60 5. Multiple merging in Abell1367

Figure 5.12: Projected density map of non emission line (left) and emission line (right) galaxies in Abell 1367. The iso-density contours of the 146 confirmed cluster members are superposed.

formation activity (e.g. Hα line in emission, Iglesias-P´aramo et al. 2002; Gavazzi et al. 2003a; Cortese et al., in preparation). Fig.5.11 shows the LOS velocity distribution of galaxies divided into emission line (upper panel) and non emission line (lower panel) galaxies. The star forming sample has higher location and scale (CBI = 6704 −1 −1  168 km s , SBI = 1076 76 km s ) than the quiescent sample (CBI = 6446 79 km s−1, S = 738 58km s−1). According to a two-sample Kolmogorov-Smirnov BI  test the two velocity distributions have only 5% probability of being consistent, suggesting a different origin and/or evolution.∼We remark that, if the star forming galaxies are infalling onto the cluster along radial orbits, their velocity dispersion should be √2 times the velocity dispersion of the relaxed sample, as observed in this case. This∼ result suggests that star forming systems are an infalling population while the non-star forming galaxies represent the virialized cluster population. The projected density distribution of star forming and non star forming is shown in Fig.5.12. The highest density of non emission line systems is observed near the center of the NW substructure. This morphological segregation further supports the idea that the NW cloud is a relaxed system merging for the first time into the SE subcluster. The emission line galaxies have a different distribution. The highest density of star forming systems is in the infalling groups, i.e. in the high velocity group infalling into the SE subcluster and in the W cloud infalling into the NW substructure, suggesting that their interaction with the cluster environment is triggering some star formation activity. Indeed in these systems the fraction of star forming galaxies lies between 64% and 36%, decreasing to 31% in the NW substructure and to 20% in the SE subcluster. 5.7. Cluster mass 61

Sample RH MV MPM 14 14 Mpc 10 M 10 M A1367 all types 0.41 7.04 0.90 7.82 2.50 A1367 non-star forming 0.37 4.35  0.70 5.11  0.90   A1367 NW all types 0.30 3.87 0.62 6.12 1.52 A1367 NW non-star forming 0.24 2.47  0.46 3.29  0.59 A1367 SE all types 0.27 5.80  0.88 6.87  1.20   A1367 SE non-star forming 0.26 3.90 0.83 5.58 0.74  

Table 5.5: Mass estimate for Abell 1367

5.7 Cluster mass

The virial theorem is the standard tool used to estimate the dynamical mass of galaxy clusters. Under the assumptions of spherical symmetry and hydrostatic equilibrium and if the mass distribution follows the distribution of the observed galaxies indepen- dent of their luminosity, the total gravitational mass of a cluster is given by 3π M = σ2R V G H

where σ is the galaxy velocity dispersion and RH is the cluster mean harmonic radius: N(N 1) RH = −−1 i>j Rij P where N is the total number of galaxies. An alternative approach is to use the projected mass estimator (Heisler et al. 1985), defined as 32 M = V 2R P M πGN i i Xi

where Vi is the observed radial component of the velocity of the i galaxy with respect to the systemic cluster velocity, and Ri is its projected separation from the cluster center. The numerical factor 32 assumes that galaxy orbits are isotropic. In case of purely radial or purely circular orbits this factor becomes 64 or 16 respectively. Mass estimates obtained using the two above methods and their uncertainties are listed in Table 5.7. We remark that these mass estimates are probably biased by the dynamical state of Abell 1367, which appears far from virialization. In particular the presence of substructures leads to an overestimate of the cluster mean harmonic radius and velocity dispersion, and thus of the virial mass (Pinkney et al. 1996). For 62 5. Multiple merging in Abell1367

this reason the mass derived for the whole cluster and for the SE and NW subclusters separately is probably overestimated. Assuming that the early type sample represents the virialized cluster population (see previous section), we also derive mass estimates for the three dynamical units using the non star forming systems only. For all the studied samples the virial mass estimates are affected by smaller uncer- tainties and yield smaller values than the projected mass estimates. This can be due to the contamination by interlopers (Heisler et al. 1985) or, more probably, to the assumption of isotropic orbits. Indeed assuming purely radial or circular orbits the mass estimate varies by a factor of 2, becoming consistent with the virial mass. The mass inferred from the non-star forming population are, as expected, systemati- cally lower than the ones obtained from all types. The value obtained for the whole sample is consistent with the mass estimates available in the literature (MV = 7.26 14 14 14  1.40 10 M Girardi et al. 1998; M = 6.07 0.93 10 M , M = 6.28 0.80 10 M V  P M  Rines et al. 2003).

5.8 Two-Body Analysis

In this section we investigate whether the two clouds A1367NW, A1367SE and the three groups infalling into the SE and NW subclusters form gravitationally bound systems. For each system we apply the two-body analysis described by Beers et al. (1991). The two subclumps are treated as point masses moving on radial orbits. They are assumed to start their evolution at time t=0 with zero separation, and are moving apart or coming together for the first time in their history. For bound radial orbits, the parametric solutions to the equations of motion are: R R = m (1 cos χ) 2 − R3 1/2 t = m (χ sin χ) 8GM  − 2GM 1/2 sin χ V = R (1 cos χ)  m  − where R is the components separation at time t, and V is their relative velocity. Rm is the separation of the subclusters at maximum expansion and M is the total mass of the system. Similarly, the parametric solutions for the unbound case are: GM R = 2 (cosh χ 1) V∞ − GM t = 3 (sinh χ χ) V∞ − 5.8. Two-Body Analysis 63

Figure 5.13: The bound and unbound orbit regions in the (Vrel, α) plane. The bound-incoming solutions (BIa and BIb), the bound-outgoing solutions (BO) and the unbound-outgoing (UO) solutions are indicated with solid lines. The dotted lines show the dividing line between bound and unbound regions. The vertical solid lines represent the observed Vrel and the dashed regions their associated 1σ uncertainty. 64 5. Multiple merging in Abell1367

sinh χ V = V∞ (cosh χ 1) − where V∞ is the asymptotic expansion velocity. The system parameters V and R are related to the observables Vrel (the LOS relative velocity) and Rp (the projected separation) by:

Vrel = V sin α, Rp = R cos α where α is the angle between the plane of the sky and the line joining the centers of the two components. The two systems are thus closed by setting the present time to t0 = 13 Gyr (the age of the Universe in a Ωm=0.3 and Ωλ=0.7 cosmology) and solved iteratively to determine the projection angle as a function of Vrel. We determine two solutions for each two-body model, assuming two extreme values for the total mass of each system ranging from the virial mass of the non-star forming population to the virial mass of the whole cluster. Table 5.8 summarizes the adopted parameters of the two-body analysis, and Fig. 5.13 shows the computed solutions in the (α, Vrel) plane. The vertical lines represent the observed values of Vrel and the dashed regions their associated 1σ uncertainties. The solutions have three different regimes: an unbound-outgoing regime (UO), a bound-outgoing regime (BO) and a bound-ingoing regime (BI). It is easy to show that the unbound solutions will lie in the region of the (α, Vrel) plane where:

V 2 R 2GM sin2 α cos α. rel p ≤ tot The dotted lines in Fig. 5.13 show the dividing line between bound and unbound regions. In the BO regime, the two subclumps are still separating and have not yet reached the maximum expansion. The BI regime describes the system after maximum expansion. For each Vrel, there are two corresponding values of α, a large and a small one. The large value assumes that the substructures are far apart, with low relative velocity, while the small value implies that the subclusters are close together near the plane of the sky (see Fig. 7 in Beers et al. 1991). Thus we split the BI regime into two branches, called BIa and BIb. The probability of each solution, computed following the procedure described by Beers et al. (1991), is given in Table 5.8. Our result is that the A1367NW/SE and the A1367SE/High Velocity group systems are bound with 100% probability and presently infalling with 96% and 100% probability respectively. The A1367NW/W and the A1367SE/Low Velocity group systems are bound at 99% and 96% probability respectively. We conclude that all systems constituting Abell 1367 are gravitationally bound at 96% probability. ≥ 5.9. Conclusions 65

System M V ∆V R Solution Probability tot rel  rel p BIa BIb BO UO 14 −1 10 M km s Mpc % % % % A1367NW/SE 7.04 84 162 0.45 57 40 3 0 4.35 84  162 0.45 55 41 4 0  A1367NW/W 7.04 500 200 0.37 57 40 2 1 2.47 500  200 0.37 56 41 2 1 A1367SE/Low Vel. gr. 7.04 1000 200 0.38 58 40 0 2  3.90 1000 200 0.38 57 39 0 4 A1367SE/High Vel. gr. 7.04 1500  200 0.08 56 44 0 0 3.90 1500  200 0.08 58 42 0 0 

Table 5.6: Two-body model parameters

5.9 Conclusions

I have presented a dynamical analysis of the central 1.3 square degrees of the galaxy cluster Abell 1367, based on 273 redshift of which ∼119 are new measurements. The LOS velocity distribution of the 146 cluster members is significantly non Gaussian, suggesting that the cluster is dynamically young. The member galaxies show an elongated distribution along the NW-SE direction with two major density peaks, consistent with the X-ray morphology. The strong difference in the LOS velocity and velocity dispersion of the two density peaks, the abrupt gas temperature gradient detected in X-rays and the 3D statistical tests support a merging scenario involving at least two subclusters. Moreover the dynamical properties of the NW and SE clouds suggest an even more complex picture, summarized in Fig. 5.6. At least another group of star forming galaxies (the high velocity group) infalling into the cluster core is detected, suggesting a multiple merging event. Furthermore our analysis suggests the presence of two other groups infalling into the cluster center. In the North-West part of Abell 1367 a group of galaxies (W subcluster), associated with the infalling galaxies 97-073/79 and with the radio relic observed in this region, is probably merging with the relaxed core of the NW subcluster. In the South part another group (the low velocity group) is infalling from the eastern side into the disturbed core of the SE subcluster. These three subgroups have a higher fraction of star forming galaxies than the cluster core, as expected during the early phase of merging events. The multiple merging scenario is consistent with the location of Abell 1367 being at the intersection of two filaments, the first extending roughly 100 Mpc from Abell 1367 toward Virgo (West & Blakeslee 2000) and the second extending between Abell 1367 66 5. Multiple merging in Abell1367

and Coma (as a part of the Great Wall, Zabludoff et al. 1993). As predicted by Katz & White (1993) this is the natural place for Abell 1367 to evolve into a rich relaxed cluster. 5.9. Conclusions 67

Name R.A. Dec. r’ V Tel. (J.2000) (J.2000) mag km s−1

114000+195426 114000.62 195426.7 15.98 10883 CAN 114159+193227 114159.52 193227.3 15.63 21228 CAN 114200+195846 114200.83 195846.0 17.09 6420 WHT 114208+191905 114208.01 191905.0 19.04 23456 WHT 114212+195650 114212.47 195650.3 17.73 20278 WHT 114213+193001 114213.87 193001.6 16.92 23641 WHT 114215+200427 114215.59 200427.0 19.20 6100 WHT 114219+200548 114219.15 200548.0 16.45 6841 CAN 114224+195329 114224.39 195329.8 18.29 31440 WHT 114224+191157 114224.48 191157.0 16.39 28546 CAN 114226+194317 114226.24 194317.1 17.50 23416 WHT 114230+191447 114230.62 191447.5 18.60 27304 WHT 114230+192553 114230.95 192553.8 17.80 45683 WHT 114238+194718 114238.24 194718.6 17.04 25610 WHT 114239+195145 114239.78 195145.9 19.06 53710 WHT 114240+195627 114240.26 195627.5 18.63 19946 WHT 114243+191615 114243.81 191615.8 18.72 5312 WHT 114249+193935 114249.85 193935.1 19.14 72429 WHT 114250+193955 114250.47 193955.7 19.22 13759 WHT 114252+195656 114252.17 195656.4 16.69 5936 WHT 114254+193851 114254.40 193851.3 17.17 6406 WHT 114254+194033 114254.93 194033.6 18.87 71389 WHT 114258+194321 114258.13 194321.1 18.98 6523 WHT 114258+194053 114258.37 194053.9 19.25 71436 WHT 114258+194644 114258.53 194644.2 19.00 88274 WHT 114258+195612 114258.94 195612.7 18.41 7059 WHT 114259+194801 114259.71 194801.1 18.92 71600 WHT 114300+192515 114300.65 192515.2 18.42 53145 WHT 114301+194758 114301.24 194758.9 18.67 72572 WHT 114301+195313 114301.97 195313.5 18.61 46935 WHT 114307+192807 114307.13 192807.3 17.37 32298 WHT 114307+193029 114307.16 193029.8 17.93 23763 WHT 114310+192526 114310.09 192526.4 16.62 19188 WHT 114310+191519 114310.29 191519.2 17.41 23578 WHT 114313+200747 114313.18 200747.9 16.40 5383 CAN 114314+194821 114314.49 194821.7 19.29 71433 WHT 114314+192534 114314.99 192534.3 15.76 23867 CAN 114317+195525 114317.25 195525.1 18.88 30273 WHT 114317+194658 114317.61 194658.2 15.69 6295 CAN 114318+201523 114318.05 201523.3 17.81 46170 WHT 114319+192520 114319.68 192520.9 15.99 6757 WHT 114320+193637 114320.44 193637.1 19.71 44171 WHT 114320+195206 114320.66 195206.2 18.15 52416 WHT 114322+195704 114322.06 195704.7 16.94 7909 WHT 114324+194121 114324.66 194121.4 18.33 35778 WHT 114332+201326 114332.24 201326.1 16.46 33438 CAN 114332+195108 114332.72 195108.2 19.07 14313 WHT 114335+200005 114335.47 200005.6 16.38 20600 CAN 114336+193930 114336.07 193930.8 19.24 44616 WHT 114337+193835 114337.17 193835.8 17.31 12502 WHT 114337+201533 114337.82 201533.5 20.19 11464 WHT 114339+193446 114339.09 193446.2 16.01 7477 WHT 114342+193636 114342.18 193636.3 19.26 71296 WHT 114343+195607 114343.12 195607.8 18.56 19711 WHT 114345+201252 114345.50 201252.2 19.27 20476 WHT 114350+195702 114350.16 195702.0 17.98 6848 WHT 114350+194138 114350.83 194138.0 19.13 72744 WHT 114353+195004 114353.42 195004.6 19.23 27946 WHT 114353+194422 114353.45 194422.2 15.66 6141 WHT 114353+194315 114353.61 194315.8 17.17 23578 WHT

Table 5.7: The 119 new redshift measurements 68 5. Multiple merging in Abell1367

Name R.A. Dec. r’ V Tel. (J.2000) (J.2000) mag km s−1

114356+201404 114356.80 201404.9 18.42 72058 WHT 114357+201122 114357.69 201122.7 17.06 5348 WHT 114358+195330 114358.86 195330.2 19.22 6200 WHT 114359+195630 114359.51 195630.8 20.37 6992 WHT 114402+194742 114402.65 194742.7 17.52 43665 WHT 114403+200552 114403.70 200552.6 15.80 5698 WHT 114404+192922 114404.17 192922.8 18.59 53335 WHT 114404+195956 114404.65 195956.6 17.33 33830 WHT 114407+193850 114407.21 193850.9 17.10 20877 WHT 114407+193143 114407.71 193143.1 18.44 53424 WHT 114412+195503 114412.22 195503.9 17.65 20916 WHT 114412+195633 114412.27 195633.4 17.02 6244 WHT 114412+201119 114412.92 201119.7 19.25 74731 WHT 114415+193037 114415.25 193037.5 16.58 6502 WHT 114415+193012 114415.33 193012.3 18.27 35227 WHT 114417+194543 114417.28 194543.9 18.14 66264 WHT 114422+194628 114422.16 194628.2 15.70 6527 CAN 114426+195951 114426.10 195951.5 16.98 30102 WHT 114430+194258 114430.30 194258.3 18.78 40347 WHT 114432+195341 114432.19 195341.6 18.89 42649 WHT 114432+194734 114432.98 194734.6 18.82 71100 WHT 114447+201248 114447.20 201248.5 18.17 6699 WHT 114449+195628 114449.72 195628.9 16.70 5539 WHT 114501+195504 114501.97 195504.5 18.79 45708 WHT 114503+193831 114503.00 193831.2 16.76 6193 LOI 114503+194743 114503.14 194743.9 17.91 23374 WHT 114504+201412 114504.25 201412.2 18.31 5477 WHT 114505+194057 114504.83 194056.9 15.67 6506 LOI 114506+200849 114506.38 200849.9 19.23 3822 WHT 114509+194845 114509.38 194845.4 17.49 19831 WHT 114509+193316 114509.40 193316.2 15.80 7409 LOI 114509+194526 114509.65 194526.9 16.90 19834 WHT 114516+193245 114516.18 193245.1 16.80 19669 WHT 114517+200120 114517.10 200120.7 15.32 14745 LOI 114517+201108 114517.29 201108.8 18.21 79253 WHT 114517+200110 114517.64 200110.0 15.46 14713 LOI 114520+194220 114520.33 194220.3 20.44 54544 WHT 114520+193259 114520.49 193259.4 17.48 4653 WHT 114522+195146 114522.62 195146.5 21.14 18012 WHT 114524+201239 114524.33 201239.3 18.73 44376 WHT 114526+201056 114526.27 201056.8 16.40 20134 CAN 114529+195658 114529.39 195658.2 16.29 24000 WHT 114530+193639 114530.37 193639.4 17.20 40000 LOI 114531+200217 114531.31 200217.5 19.78 45691 WHT 114533+194505 114533.88 194505.9 18.11 31440 WHT 114533+200028 114533.97 200028.7 17.56 35830 WHT 114536+194253 114536.19 194253.7 18.60 48966 WHT 114540+194302 114540.32 194302.8 17.74 5545 WHT 114543+193854 114543.65 193854.9 16.93 7828 LOI 114543+193905 114543.77 193905.9 16.30 7301 LOI 114544+194013 114544.86 194013.3 17.18 19487 WHT 114545+193151 114545.66 193151.4 18.57 6880 LOI 114545+201200 114545.78 201200.3 19.11 27431 WHT 114548+192708 114548.13 192708.4 16.72 30193 WHT 114549+195915 114549.88 195915.3 15.87 20035 CAN 114550+194824 114550.61 194824.6 18.85 41484 WHT 114602+194754 114602.12 194754.3 19.62 73746 WHT 114605+195151 114605.35 195151.0 18.86 46635 WHT 114620+194518 114620.85 194518.0 17.58 45683 WHT

Table 5.7: Continue Chapter 6

Unveiling the evolution of early type galaxies with GALEX.

6.1 Introduction

In Chapters 3 and 4 I have shown that at low UV luminosities the contribution of early-type quiescent galaxies is not negligible. This represents the first evidence of a morphology/star formation - density relation at ultraviolet wavelengths and demon- strates that we cannot blindly assume all UV selected galaxies are star-forming sys- tems, especially at low UV luminosities and in high density environments. This also points out the strong potential of ultraviolet observations for studying all cluster galaxies: not only star-forming systems in which UV emission traces the presence of newly born stars, but also early type galaxies whose emission is usually ascribed to low mass old post asymptotic giant branch stars. The excess ultraviolet radia- tion from giant early-type galaxies is in fact supposed to arise from hot low mass stars in late stages of stellar evolution (O’Connell 1999). All theoretical, spectral and imaging evidences have recently converged towards the view that the UV emission originates from He-burning, extreme horizontal branch stars, their post-HB progeny and post-AGB stars in the dominant, metal rich stellar population of elliptical galax- ies. However it is still unknown whether the UV emission of all early type galaxies is dominated by the contribution of old stellar populations independently from the galaxy morphology (i.e. ellipticals vs. lenticulars) and luminosity (i.e. dEs vs. gi- ant Es). In particular it would be interesting to know if the UV properties of dwarf elliptical galaxies differ from those of giants, as much as other structural (Gavazzi et al. 2005) and kinematic (van Zee et al. 2004) properties depend on luminosity, due to their different star formation histories (single episodic vs. burst) (Ferguson & Binggeli 1994; Grebel 2000). In fact, a recent burst of star formation would strongly contribute to the UV emission of an elliptical galaxy, even if its stellar population is

69 70 6. Unveiling the evolution of early type galaxies with GALEX.

dominated by old low mass stars. Due to morphological segregation (Whitmore et al. 1993), nearby clusters are the ideal targets for assembling complete, volume limited samples of early-type objects. As part of a study aimed at analyzing the environmental dependence of galaxy evo- lution, we observed large portions of the Virgo cluster with GALEX (Boselli et al. 2005a). Owing to the superior quality of the photographic material obtained by Sandage and collaborators, an extremely accurate and homogeneous morphological classification exists for Virgo galaxies, down to m 18 mag (M -13 assuming a B ≤ B ≤ distance of 17 Mpc), allowing a detailed discrimination among different subclasses of early-type galaxies (ellipticals, lenticulars, dwarfs) and from quiescent spirals. Fur- thermore a wealth of ancillary data for many Virgo members, covering a large portion of the electromagnetic spectrum from the visible to the infrared is available from the GOLDMine database (Gavazzi et al. 2003a).

6.2 Data

The analysis presented in this Chapter is based on an optically selected sample of early-type galaxies including giant and dwarf systems (E, S0, S0a, dE and dS0) ex- tracted from the Virgo Cluster Catalogue of Binggeli et al. (1985), which is complete to m 18 mag (M -13). The Virgo cluster region was observed in spring 2004 B ≤ B ≤ as part of the All Imaging Survey (AIS) and of the Nearby Galaxy Survey (NGS) carried out by the Galaxy Evolution Explorer (GALEX) in two UV bands: FUV (λeff = 1528A˚, ∆λ = 442A)˚ and NUV (λeff = 2271A˚, ∆λ = 1060A),˚ covering 427 ob- jects. See Chapter 2 and Martin et al. (2005) and Morrissey et al. (2005) for details on the GALEX instrument and data characteristics. The present sample includes all Virgo cluster early-type systems detected in the NUV GALEX band (264 objects, 194 from the NGS); of these, 126 (of which 74 from the NGS) have been also detected in the FUV. The resulting sample is thus ideal for the proposed analysis as it provides us with the first large volume-limited sample of ellip- tical, lenticular and dwarf galaxies spanning 4 dex in luminosity with homogeneous data. Whenever available, we extracted fluxes from the deep NGS images, obtained with an average integration time of 1500 sec, complete to mAB 21.5 in the NUV and FUV. Elsewhere UV fluxes hav∼e been extracted from the less∼deep AIS images ( 70 sq. degrees), obtained with an average integration time of 100 sec, complete to∼m 20 in both the FUV and NUV bands. The resulting sample,∼ although not AB ∼ complete in both UV bands, includes giants and dwarf systems: at a limiting magni- tude of MB -15, 71 % of the observed galaxies have been detected in the NUV, 46% in the FUV.≤All UV images come from the GALEX IR1.0 release. UV fluxes were obtained by integrating GALEX images within elliptical annuli of increasing diameter up to the optical B band 25 mag arcsec−2 isophotal radii consistently with the optical 6.3. The UV properties of early-type galaxies 71

and near-IR images. Independent measurements of the same galaxies obtained in different exposures give consistent photometric results within 10% in the NUV and 15% in the FUV in the AIS, and about a factor of two better for bright (NUV 16) ≤ galaxies. The statistical uncertainty in the UV photometry is on average a factor of 2 better in the NGS than in the AIS especially for fainter objects. ∼ UV data have been combined with multifrequency data taken from the GOLDMine database (http://goldmine.mib.infn.it; Gavazzi et al. 2003a). These are B and V imaging data, mostly from Gavazzi et al. (2005) and Boselli et al. (2003a), and near- IR H imaging from Gavazzi et al. (2000, 2001c). Optical and near-IR data have on average a photometric precision of 10%. Spectroscopic metallicity index Mg and ∼ 2 velocity dispersion data come from GOLDMine or from Golev & Prugniel (1998) and Bernardi et al. (2002). Galaxies analyzed in this Chapter are all bona-fide Virgo cluster members: given the 3-D structure of the cluster, distances have been assigned following the subcluster membership criteria of Gavazzi et al. (1999a). Owing to the high galactic latitude of Virgo, no galactic extinction correction was applied (A 0.05). B ≤

6.3 The UV properties of early-type galaxies

Despite the complex 3-D structure of Virgo (Gavazzi et al. 1999a), the uncertainty on the distance (hence on the luminosity) of the target galaxies, does not constitute a major source of dispersion in the determination of the color-magnitude (CMR) re- lation. Figure 6.1 shows various UV to optical and near-IR CMRs. Similar results are obtained if, instead of the mass-tracer H band luminosity (Zibetti et al. 2002), we use the B band absolute magnitude. The NUV to optical (Fig. 6.1b) and near-IR (Fig. 6.1a) CMRs are well defined and are similar to optical or near-IR CMRs, with brighter galaxies having redder colors, independent of their morphological type: the 8 color index (NUV V ) increases by 2 magnitudes from dwarfs (LH 10 LH ) to 11.5− ∼ ∼ giants (LH 10 LH ), while (NUV H) changes by 3 mag. A weak flattening ∼ 10 − ∼ of the relation appears for L 10 L . This behavior confirms the one reported H ≥ H by Ferguson (1994) in the (B V ) vs. MB CMR. On the contrary, the FUV to optical− (Fig. 6.1d) and near-IR (Fig. 6.1c) CMRs differ 9.5 systematically for dwarfs and giant systems: galaxies brighter than LH 10 LH 9.5 ∼ have similar red colors, while for LH 10 LH colors become progressively bluer. Even if this trend can be due to a selection≤ effect, (reddest dwarfs being undetectable in the FUV), it is indisputable that there exists a significant population of dEs with bluer colors than Es and S0s. The dichotomy between giants and dwarfs is even more apparent in the UV color index (F UV NUV ) (see Fig. 6.2). The (F UV NUV ) − − becomes redder with increasing luminosity for dwarf ellipticals while, on the contrary, it becomes bluer for giant ellipticals (Fig. 6.2a). The blueing relation is tight among 72 6. Unveiling the evolution of early type galaxies with GALEX.

Figure 6.1: The near-UV (left column) and far-UV (right column) to optical and near- IR color magnitude relations. Colors are in the AB magnitude system. Open circles are for ellipticals, filled circles for dwarfs, crosses for lenticulars (S0-S0a). Galaxies redder than the dashed line are undetectable by the present survey (at the NGS limit). Largest 1σ errors for luminous and dwarf systems are given. 6.3. The UV properties of early-type galaxies 73

Table 6.1: Main relations for early type galaxies

x y a b R rms Ellipticals1 L F UV NUV 0.30 0.14 +4.52 1.52 0.47 0.31 H − −   − LH F UV H 0.22 0.19 +10.55 2.10 0.28 0.43 L NUV − H −0.17 0.18 +4.85 1.85 −0.22 0.47 H −   LH F UV V 0.15 0.18 +8.38 1.88 0.21 0.38 L NUV − V −0.26 0.12 +2.55  1.30 −0.45 0.31 H −   B H F UV NUV 0.84 0.45 +3.22 0.98 0.43 0.32 σ − F UV − NUV −1.35  0.37 +4.39  0.89 −0.69 0.26 − −   − Lenticulars

LH F UV NUV 0.28 0.15 +4.40 1.62 0.31 0.45 L F UV− H −0.31 0.21 +0.75  2.00 −0.27 0.58 H −   LH NUV H 0.61 0.11 +0.51 1.17 0.65 0.36 L F UV − V 0.03  0.23 +6.62  2.38 0.03 0.59 H −   LH NUV V 0.49 0.09 +0.26 1.00 0.68 0.25 B H F UV −NUV 1.00 0.32 +3.70  0.70 0.49 0.42 σ − F UV − NUV −1.29  0.39 +4.28  0.84 −0.58 0.39 − −   − Dwarfs

LH F UV NUV 1.73 0.41 13.90 2.16 0.52 0.59 − ∗ − ∗ ∗ ∗ L F UV H 2.55 0.55 15.97 4.96 0.68 0.91 H −  −  LH NUV H 0.91 0.19 2.72 1.68 0.56 0.57 − ∗ − ∗ ∗ ∗ L F UV V 1.91 0.55 +11.35 4.93 0.60 0.87 H −   LH NUV V 0.63 0.17 +1.28 1.05 0.49 0.47 B H F UV −NUV 0.95  0.45 +0.12  0.73 0.40 0.60 − −   σ F UV NUV − − − − Notes to Table: Col. 1 and 2: x and y variables Col. 3 and 4: slope a and intercept b of the bisector linear fit with weighted variables Col. 5: Pearson correlation coefficient Col. 6: mean dispersion around the best fit 1: excluding VCC 1499 *: uncertain values because of the UV detection limit 74 6. Unveiling the evolution of early type galaxies with GALEX.

Figure 6.2: The relationship between the UV color index (F UV NUV ) and a) − the total H band luminosity, b) the B-H color index, c) the logarithm of the central velocity dispersion and d) the Mg2 index. Symbols are as in fig. 6.1. Labeled points indicate objects having unusual radio or optical properties (see Sect. 3). 6.4. Discussion and conclusion 75

ellipticals (see Table 1) and barely observed in lenticulars because of their higher dispersion 1. A similar behavior between ellipticals and lenticulars is observed in the (F UV NUV ) − color relation (Fig. 6.2b): this mixed giant population becomes bluer in the UV with increasing reddening in the (B H) color index. − The behavior of dwarf ellipticals is different: although with a huge dispersion, the (F UV NUV ) color index reddens as the (B H) and the other optical color indexes. The dic−hotomy between dwarf and giant systems− cannot be observed in the run of (F UV NUV ) vs. central velocity dispersion (which is directly related to the system total dynamical− mass; Fig. 6.2c) nor as a function of the metallicity sensitive (Pog- gianti et al. 2001) Mg2 Lick index (Fig. 6.2d) because these two parameters are not available for dwarfs. In ellipticals and lenticulars the UV color index (F UV NUV ) − depends on both the metallicity index Mg2 and σ in a way opposite to the behavior at optical wavelengths, where galaxies are redder when having higher Mg2 and velocity dispersions.

6.4 Discussion and conclusion

For the first time the UV properties of early-type galaxies have been studied down to MB -15 mag. The comparison with previous studies is thus limited to the brightest objects.∼ Our CMR can be compared with the one obtained by Yi et al. (2005) based on a complete sample of bright early-type objects (Mr -20 mag) extracted from the Sloan Digital Sky Survey (SDSS) by Bernardi et al. (2003).≤ The CMR presented by Yi et al. (2005) (NUV r vs. Mr) shows a significantly larger dispersion (σ 1.5 mag) than the one found− in Virgo (see Table 1). As discussed in Yi et al. (2005),≥ the large dispersion in their CMR can be ascribed to galaxies with a mild or residual star formation activity included in the Bernardi et al. (2003) sample. If restricted to the ”UV weak” sample, the dispersion in the Yi et al. relation drops to 0.58 mag, i.e. still larger than the one seen in the Virgo cluster in the same luminosity range. Despite possible larger distance uncertainties in the SDSS, the difference in the scat- ter between our and the Yi et al. (2005) CMR might arise from the classification in the SDSS that uses concentration indices and luminosity profiles in discriminating hot from rotating systems. It is in fact conceivable that the larger dispersion in the CMR of ”UV weak” galaxies of Yi et al. (2005) comes from the contamination of qui-

1The scatter in the blueing relation among ellipticals decreases significantly (from 0.31 to 0.10) if we exclude the misclassified post-starburst dwarf VCC 1499 (Gavazzi et al. 2001c; Deharveng et al. 2002), the radio galaxy M87, VCC 1297 (the highest surface brightness galaxy in the sample of Gavazzi et al. (2005)) and VCC 1146. Beside its extremely high surface brightness, making VCC 1297 a non standard object, we do not have any evidence indicating a peculiar star formation history or present nuclear activity in VCC 1297 and VCC 1146 that could justify their exclusion. 76 6. Unveiling the evolution of early type galaxies with GALEX.

escent, bulge-dominated Sa spiral disks, that have structural (concentration indices and light profiles) or population properties (colors and spectra) similar to ellipticals and lenticulars (Scodeggio et al. 2002). The monotonic increase of the (NUV V ) and (NUV H) colors with luminosity, similar to the one observed in the visible− bands by Ferguson− (1994) strongly suggests that both in dwarfs and giant systems the NUV 2310 A˚ flux is dominated by the same stellar population (main sequence low mass stars) emitting at longer wavelengths. On the contrary the different behaviour of the (F UV V ) and (F UV H) colors with − − luminosity, and the clear dichotomy observed in the (F UV NUV ) vs LH CMR strongly support a different origin for the FUV emission in dw−arf and giant systems. The reddening of the UV color index with luminosity observed in dwarf ellipticals, similar to the one observed in late type galaxies, indicates that the UV spectral energy distribution of low mass early type galaxies is shaped by the contribution of young stellar populations. This is shown in Fig.6.3 where the available optical spectra for our sample of dEs are shown. It clearly emerges that UV bluer systems have emission lines or strong Balmer line in absorptions witnessing present or recent star forma- tion activity. Moreover at increasing luminosity their (F UV NUV ) color index reddens as the optical colors confirming that in these systems −the FUV emission is dominated by the contribution of young main sequence stars. This is not the case for giant early type systems: the plateau observed in the FUV-optical CMRs and the blueing of the (F UV NUV ) color with luminosity (i.e. the UV upturn) suggest − that far ultraviolet emission comes from low mass old post asymptotic giant branch stars. This is also confirmed in Figs.6.5 and 6.4 where the optical spectra available for ellipticals and S0s in our sample are presented: as expected, all the spectra are dominated by the contribution of the old stellar populations. Moreover the observed trend between (F UV NUV ) and the metallicity sensitive Mg index, reproduced by − 2 models (Bressan et al. 1994; Yi et al. 1998), confirms the early IUE result of Burstein et al. (1988). Conversely Rich et al. (2005) did not find any correlation between the color index (F UV r) and Mg nor with the velocity dispersion σ in a large sample of − 2 SDSS early-type galaxies observed by GALEX. Their lack of correlation might derive −1 from insufficient dynamic range in Log σ (2.1-2.4 km s ) and Mg2 (0.18-0.30). The blueing of the UV color index with luminosity, metallicity and velocity dispersion indicates that the UV upturn is more important in massive, metal rich systems. This is consistent with stellar population models which predict that the strength of the UV upturn is mainly driven by stellar metallicity. The accurate morphological classification in our sample allow us to discriminate be- tween E and S0s and to study separately the two populations. The higher dispersion in the (F UV NUV ) vs. LH relation observed for the lenticulars compared to the ex- tremely tight−one for ellipticals (see Table 1), bears witness to a different evolutionary history for the two Hubble types: while cluster ellipticals represent an homogeneous population, S0s are a heterogeneous class probably formed by different independent 6.4. Discussion and conclusion 77

Figure 6.3: The relationship between the UV color index (F UV NUV ) and the − total H band luminosity. Symbols are as in fig. 6.1. The optical spectra available for dwarf ellipticals are presented. 78 6. Unveiling the evolution of early type galaxies with GALEX.

Figure 6.4: The relationship between the UV color index (F UV NUV ) and the − total H band luminosity. Symbols are as in fig. 6.1. The optical spectra available for ellipticals are presented. 6.4. Discussion and conclusion 79

Figure 6.5: The relationship between the UV color index (F UV NUV ) and the − total H band luminosity. Symbols are as in fig. 6.1. The optical spectra available for lenticulars are presented. 80 6. Unveiling the evolution of early type galaxies with GALEX.

physical mechanisms (see also Chapters 9 and 11), and with various star formation histories as also determined from kinematic and spectroscopic observations (Dressler & Sandage 1983; Neistein et al. 1999; Hinz et al. 2003). Using the available optical spectra we investigate the presence of residual star for- mation still present in our sample of giant early type galaxies. Only one S0 galaxy, VCC1003, shows a mild residual star formation activity (Hα in emission), while three ellipticals (VCC881,M87,VCC1619) and three S0s (VCC1030,VCC1062,VCC1253) have [NII] in emission and Hα in absorption, a typical feature of low ionization active galactic nuclei. This suggest that the difference observed between ellipticals and S0s cannot be ascribed to recent episodes of star formation but probably resides on their different past star formation history. Combining this result with the one obtained in Chapter 4, we can conclude that, at low UV luminosities, the significant contribution of giant early type systems to the ultraviolet luminosity function must be ascribed not to young stellar populations but to old low mass post-AGB stars. The newest result of this Chapter, shown in Fig. 6.2, addresses the question raised by O’Connell (1999) concerning the dependence of the UV properties on galaxy mor- phology. We have shown that a dichotomy exists between giant and dwarf ellipticals and, to a lesser extent, between ellipticals and lenticulars. The opposite behavior (reddening of the UV color index with luminosity) of dwarfs with respect to giants, similar to that observed for spirals, indicates that the UV spectra of low luminosity objects are shaped by the contribution of young stars, thus are more sensitive to the galaxy’s star formation history than to the metallicity. This implies that the stellar population of dwarfs has been formed in discrete and relatively recent episodes, as observed in other nearby objects (Grebel 2000). More evidences are building up that mass drives the star formation history in hot systems (Trager et al. 2000; Gavazzi et al. 2002a; Caldwell et al. 2003; Poggianti 2004b) as in rotating ones (Gavazzi et al. 1996; Boselli et al. 2001) and that the stellar population of massive ellipticals is on average older than that of dwarfs. Chapter 7

UV dust attenuation in normal star forming galaxies

7.1 Introduction

The use of ultraviolet emission in order to study the properties of star forming galax- ies is not an easy a rapid task. The presence of dust in galaxies represents one of the major obstacles complicating a direct quantification of the star formation activity in local and high redshift galaxies. Absorption by dust grains reddens the spectra at short wavelengths completely modifying the spectral energy distribution of galaxies. Since the UV radiation is emitted by young stars (t < 108 yr) that are generally more affected by attenuation from surrounding dust clouds than older stellar populations, rest-frame UV observations can lead to incomplete and/or biased reconstructions of the star formation activity and star formation history of galaxies affected by dust absorption, unless proper corrections are applied. In recent years our understanding of dust attenuation received a tremendous impulse from studies of local starburst galaxies (i.e.Calzetti et al. 1994; Heckman et al. 1998; Meurer et al. 1999; Calzetti 2001; Charlot & Fall 2000), that were based on three indicators: the ratio of the total infrared to far-ultraviolet emission (LT IR/LF UV ), the ultraviolet spectral slope β (determined from a power-law fit of the form f λβ to the UV continuum spectrum in the range 1300 and 2600 A,˚ Calzetti et al. 1994)∼ and the Balmer decrement. The total-IR (TIR) to UV luminosity ratio method (i.e. Buat 1992; Xu & Buat 1995; Meurer et al. 1995, 1999) is based on the assumption that a fraction of photons emitted by stars and gas are absorbed by the dust. The dust heats up and subsequently re-emits the energy in the mid- and far-infrared. The amount of UV attenuation can thus be quantified by means of an energy bal- ance. This method is considered the most reliable estimator of the dust attenuation in star-forming galaxies because it is almost completely independent of the assumed

81 82 7. UV dust attenuation in normal star forming galaxies extinction mechanisms (i.e. dust/star geometry, extinction law, see Buat & Xu 1996; Meurer et al. 1999; Gordon et al. 2000; Witt & Gordon 2000). When the spectrum is dominated by a young stellar population the ultraviolet spectral slope β, is found to have a weak dependence on metallicity, IMF, and star formation history (Leitherer & Heckman 1995). Thus the difference between the observed β and the one predicted by models can be entirely ascribed to dust attenuation (Meurer et al. 1999). However in systems with no or mild star formation activity the UV spectral slope can be strongly contaminated by the old stellar populations, whose contribution increases β (flattens the UV continuum, Boissier et al. 2005). Thus the spectral slope of mildly star form- ing systems could be intrinsically different from the one of starburst galaxies, even in the absence of dust attenuation (Kong et al. 2004). Meurer et al. (1999) have shown that in starburst galaxies the total far-infrared to ul- traviolet luminosity ratio correlates with the ultraviolet spectral slope, β (commonly referred to as the IRX-UV relation). They pointed out that this relation allows reli- able estimates of the attenuation by dust at ultraviolet wavelengths based on β. The Balmer decrement gives an estimate of the attenuation of ionized gas and not of the stellar continuum as in the previous two methods. It is based on the comparison of the observed Hα/Hβ ratio with its predicted value (2.86 for case B recombina- 4 −3 4 tion, assuming an electronic density ne 10 cm and temperature 10 K; e.g., Osterbrock 1989). Calzetti et al. (1994)≤found a significant correlation∼between the ultraviolet spectral slope β and the Balmer decrement Hα/Hβ. Starting from this empirical relation they obtained an attenuation law (known as the Calzetti attenua- tion law) often adopted to correct UV observations for dust attenuation in absence of both far-infrared observations and estimates of the ultraviolet spectral slope (Steidel et al. 1999; Glazebrook et al. 1999). Unfortunately the above empirical relations have been established only for starburst galaxies and they seem not to hold for normal star forming galaxies. Recently, Bell (2002) suggested that quiescent galaxies deviate from the IRX-UV relation of star- burst galaxies, because they tend to have redder ultraviolet spectra at fixed total far-infrared to ultraviolet luminosity ratio. Kong et al. (2004) confirmed this result and interpreted the different behaviour of starbursts and normal galaxies as due to a difference in the star formation histories. They proposed that the offset from the starburst IRX-UV relation can be predicted using the birthrate parameter b (e.g. the ratio of the current to the mean past star formation activity). However an inde- pendent observational confirmation of the correlation between the distance from the starburst IRX-UV relation and the birthrate parameter has not been obtained so far (Seibert et al. 2005). Even the Calzetti law does not seem to be universal. Buat et al. (2002) showed that for normal star forming galaxies the attenuation derived from the Calzetti law is 0.6 mag larger than the one computed from the F /F ratio and ∼ FIR UV their result has been recently confirmed by Laird et al. (2005). Why do normal star-forming galaxies behave differently from starbursts? Do normal 7.1. Introduction 83 galaxies follow different empirical relations that can be exploited to correct for dust attenuation in absence of far infrared observations? If this is the case, is there a transition between starburst and normal galaxies? Which physical parameters drive it? Answering these questions will be important for a better understanding of the interaction of dust and radiation specifically in nearby dusty star forming galaxies, but it also has direct consequences for our understanding and interpretation of galaxy evolution in a general context. Firstly it seems mandatory to characterize the dust at- tenuation properties of normal galaxies, to compare them with the ones of starbursts and to derive new recipes for the UV dust attenuation correction. This topic came once again to the fore with the launch of the Galaxy Evolution Explorer (GALEX). This satellite is delivering to the community an unprecedented amount of UV data on local and high redshift galaxies that require corrections for dust attenuation but currently lack far-infrared rest-frame data. The time is ripe to explore new methods for correction of these data, that might provide new insights on galaxy evolution. Whenever they can be combined with other data, GALEX observations provide the best available ultraviolet data for studying the dust attenuation properties of galaxies. Multiwavelength photometric and spectroscopic observations are in fact mandatory in order to: determine metallicity, ionized gas attenuation (A(Hα)), luminosity and mass, test the validity of the relations followed by starbursts (Heckman et al. 1998), explore relations that might prove useful to correct ultraviolet magnitudes and to compare them with various models of dust attenuation. Recent extensive spectro- scopic and photometric surveys, like the Sloan Digital Sky Survey (SDSS, Abazajian et al. 2005) and the Two Degree Field Galaxy Redshift Survey (2dF, Colless et al. 2001) have opened the path to studies of fundamental physical parameters based on enormous datasets. However, spectroscopic observations of nearby galaxies suffer from strong aperture effects, making these datasets not ideal for the purpose of the present investigation. In fact, Kewley et al. (2005) have recently shown that aperture effects produce both systematic and random errors on the estimate of star-formation, metallicity and attenuation. To reduce at least the systematic effects they suggest selecting only samples with fibres that capture > 20% of the light. This requires z > 0.04 and z > 0.06 for SDSS and 2dF respectively: too distant to detect both giant and dwarf star forming systems with GALEX and IRAS. Although significantly smaller than the SDSS, the dataset we have been building up over the last 10 years with data taken over a large stretch of the electromagnetic spec- trum for few thousand galaxies in the local universe (worldwide available from the site GOLDMine; Gavazzi et al. 2003a) turns out to be appropriate for the purposes of the present investigation. It includes drift-scan mode integrated spectra, narrow band Hα and broad band optical and near-infrared imaging for a volume limited sam- ple of nearby galaxies in and outside rich clusters. The combination of GALEX and IRAS observations with these ancillary data allows us to study the dust attenuation properties in a sizable sample of normal star forming galaxies not suffering from the 84 7. UV dust attenuation in normal star forming galaxies

aperture bias and to compare observations with model predictions. In this chapter I investigate the relations between dust attenuation and global galaxy properties and compare them with the ones observed in starburst galaxies. The aim of this work is to provide some empirical relations based on observable quantities (thus model independent) suitable for deriving dust attenuation corrections when far infrared data are not available. For this reason all relations obtained throughout this chapter will be given as a function of LT IR/LF UV , the observable that we consider the best dust attenuation indicator. We choose not to transform LT IR/LF UV into a (model dependent) estimate of the far ultraviolet extinction A(F UV ), leaving the reader free to choose his/her preferred dust model (i.e. Meurer et al. 1999; Buat et al. 1999, 2002, 2005; Gordon et al. 2000; Panuzzo et al. 2003; Burgarella et al. 2005, Inoue et al. in preparation). We assume that quantities are related linearly and residual plots are presented in order to test the validity of this hypothesis. Moreover, since we are looking for new recipies to estimate the LT IR/LF UV ratio, this quantity has to be considered as the dependent variable, implying the use of an unweighted simple linear fit to estimate the best fitting parameters (Isobe et al. 1990).

7.2 The Data

7.2.1 The optically-selected sample The analysis presented in this work is based on an optically selected sample of late- type galaxies (later than S0a) including giant and dwarf systems extracted from the Virgo Cluster Catalogue (VCC, Binggeli et al. 1985) and from the CGCG catalogue (Zwicky et al. 1961). The data include 300 square degrees covering most of the ∼ Virgo, Abell1367 and Abell262 clusters, the southwest part of the Coma cluster and part of the Coma-A1367 supercluster (11h30m < R.A. < 13h30m; 18◦ < decl. < 32◦) observed in spring 2004 as part of the All-sky Imaging Survey (AIS) and of the Nearby Galaxy Survey (NGS) carried out by GALEX in two UV bands: FUV (λeff = 1530A˚, ∆λ = 400A)˚ and NUV (λeff = 2310A˚, ∆λ = 1000A).˚ Details of the GALEX instrument and characteristics can be found in Martin et al. (2005) and Morrissey et al. (2005). Our sample has the quality of being selected with the criterion of optical completeness. All galaxies brighter than a threshold magnitude are selected in all areas. In Coma-A1367 supercluster and A262 cluster all galaxies brighter than mp=15.7 were selected from the CGCG catalogue (Zwicky et al. 1961). The Virgo region contains all galaxies brighter than mp=18 from the VCC catalogue (Binggeli et al. 1985). We thus consider our sample an optically selected, volume limited sam- ple. We include in our analysis all late-type galaxies, detected in both NUV and FUV GALEX bands and in both 60 µm and 100 µm IRAS bands (157 objects). When- 7.2. The Data 85

ever available, we extracted UV fluxes from the deep NGS images, obtained with a mean integration time of 1500 sec, complete to mAB 21.5 in the NUV and FUV. Elsewhere UV fluxes have∼been extracted from the shallo∼wer AIS images ( 70 sq. de- ∼ grees), obtained with a mean integration time of 100 sec, complete to mAB 20 in both the FUV and NUV bands. All UV images come∼ from the Internal Data Release∼ v1 (IR1.0). UV fluxes were obtained by integrating GALEX images within elliptical annuli of increasing diameter up to the optical B band 25 mag arcsec−2 isophotal radii, consistently with the optical and near-IR images. Independent measurements of the same galaxies obtained in different exposures give consistent photometric re- sults within 10 % in the NUV and 15% in the FUV in the AIS, and a factor of ∼ two better for bright (NUV 16) galaxies. The uncertainty in the UV photometry is on average a factor of 2 b≤etter in the NGS than in the AIS, particularly for faint objects. The typical uncertain∼ ty in the IRAS data is 15% (Boselli et al. 2003a). UV and far-infrared data have been combined to multifrequency data. These are opti- cal and near-IR H imaging (mostly from Gavazzi et al. 2000, 2005; Boselli et al. 2003a), optical drift-scan spectra (Gavazzi et al. 2004; Gavazzi et al. in prep.) and Hα imag- ing (Boselli & Gavazzi 2002; Boselli et al. 2002a; Gavazzi et al. 1998, 2002b; Iglesias- P´aramo et al. 2002; Gavazzi et al. in prep.), great part of which are available from the GOLDMine galaxy database (Gavazzi et al. 2003a) (http://goldmine.mib.infn.it). From the 157 galaxies selected we exclude Active Galactic Nuclei (AGN). AGNs have been selected using either the classification provided by NED, if available, or by in- spection to the integrated spectra of Gavazzi et al. (2004): we exclude galaxies with log([OIII]/Hβ) > 0.61/(log([NII]/Hα) 0.05)+1.3 (Kauffmann et al. 2003a). This criterion reduces the sample to 128 galaxies,− spanning a range of six magnitudes in B 1 band (-22< MB <-16) and of three orders of magnitude in mass (9 < M < 12 M ). Unfortunately ancillary data are not available for all galaxies observed by GALEX, we thus further divided the data in two subsamples. Sixty six galaxies in the primary sample have all the necessary complementary data (e.g. Hα photometry, Hα/Hβ ratio, metallicity, H-band photometry; see Gavazzi et al. 2000, 2002a,b, 2004 for the selection criteria adopted in each survey). The remaining 62 galaxies form the sec- ondary sample. We cannot exclude a possible contamination of AGN in the secondary sample, since no spectra are available for these objects. In all figures objects belonging to the primary sample will be indicated with filled circles while the secondary sample as empty circles. Since only galaxies belonging to the primary sample are present in all the plots analyzed in the presented work, all correlations will be quantified using only the primary sample. Data from UV to near-IR have been corrected for Galactic extinction according to Burstein & Heiles (1982). We assume a distance of 17 Mpc for the members of Virgo Cluster A, 22 Mpc for Virgo Cluster B, and 32 Mpc for objects in the M and W clouds (Gavazzi et al. 1999a).

1 Computed using the relation between LH and M by Gavazzi et al. (1996) 86 7. UV dust attenuation in normal star forming galaxies

Members of the Cancer, A1367, and Coma clusters are assumed to lie at distances of 65.2, 91.3, and 96 Mpc, respectively. Isolated galaxies in the Coma supercluster are −1 −1 assumed at their redshift distance, adopting H0 = 75 km s Mpc .

7.2.2 The starburst sample In order to compare the properties of our sample with starbursts, we compile a dataset of starburst galaxies observed by IUE from the sample of Calzetti et al. (1994). We consider 29 galaxies, excluding AGNs and galaxies that have not been observed by IRAS at 60 or 100 µm. Complementary data such as FIR, Hα fluxes and Balmer decrements are taken from Calzetti et al. (1995), metallicities come from Heckman et al. (1998) and H-band photometry (available only for 18 galaxies) from (Calzetti 1997). Excluding the far infrared fluxes, all these quantities are obtained within an apertures of 20 10arcsec2, consistent with IUE observations Calzetti et al. ∼ × (1994). Thus we stress that aperture effects could strongly affect any comparison with normal galaxies for which all data are homogeneously integrated values. First of all, if the UV emission is more extended than IUE field of view the LT IR/LF UV ratio is overestimated2. In addition, even when physical quantities are obtained in the same IUE apertures, the presence of age and metallicity gradients in galaxies makes not trivial any comparison with the integrated values obtained for normal star forming galaxies (Kewley et al. 2005). All the observables, but the ultraviolet spectra slope β, are calibrated in a consistent way with our sample of normal galaxy. The ultraviolet spectral slope of starbursts is obtained by fitting IUE spectra (Calzetti et al. 1994), while for GALEX observations it comes from the FUV-NUV color index (see next Section). However, as shown by Kong et al. (2004), these two calibrations are consistent each other and do not introduce any systematic difference between the two samples.

7.3 The L /L β relation for normal star- T IR F UV − forming galaxies

Meurer et al. (1999) have shown that the ratio of far infrared to far ultraviolet lu- minosity tightly correlates with the UV colors of starburst galaxies. This relation, known as the infrared excess-ultraviolet (IRX-UV) relation, is often presented as β vs. LT IR/LF UV relation. As discussed in the introduction, we will refer to the LT IR/LF UV ratio as the best indicator of UV dust attenuation and we will calibrate on it all the following relations. In order to determine the dust emission, we compute the total

2However Meurer et al. (1999) argued that the majority of UV flux for their starburst sample lies within the IUE aperture 7.3. The L /L β relation for normal star-forming galaxies 87 T IR F UV −

infrared flux emitted in the range 1-1000 µm, following Dale et al. (2001):

f60 log(fT IR) = log(fF IR) + 0.2738 0.0282 log( ) + − × f100 f f +0.7281 log( 60 )2 + 0.6208 log( 60 )3 + × f100 × f100 f +0.9118 log( 60 )4 (7.1) × f100 where fF IR is the far-infrared flux, defined as the flux between 42 and 122 µm (Helou et al. 1988): f = 1.26 (2.58 f + f ) 10−14 [Wm−2] (7.2) F IR × × 60 100 × and f60 and f100 are the IRAS fluxes measured at 60 and 100 µm (in Jansky). The total infrared luminosity is thus:

2 LT IR = 4πD fT IR (7.3)

The β parameter as determined from GALEX colors is very sensitive to the galaxy star formation history (see for example Calzetti et al. 2005). For this reason we assume β as defined by Kong et al. (2004):

log(f ) log(f ) β = F UV − NUV = 0.182 = 2.201 (F UV− NUV ) 1.804 (7.4) × − −

where fF UV and fNUV are the near and far ultraviolet observed fluxes respectively (in erg cm2 s−1 A˚−1), and FUV and NUV are the observed magnitudes. The relationship between the ratio of total infrared luminosity (LT IR) obtained from (7.1) to the far-ultraviolet fluxes and the UV spectral slope β (or the FUV-NUV color) for our sample of nearby star forming galaxies is given in Fig.7.1. Several functional forms of the L /L β relation can be found in the literature (i.e. Meurer et al. T IR F UV − 1999; Kong et al. 2004); we simply adopt a linear fit: log(LT IR/LF UV ) = a β + b. This functional form is consistent with other previously proposed for β > ×2, while it diverges for β < 2. Since the majority of normal and starbursts galaxies− have − β > 2 our choice is justified. This represents the simplest and less parameter dependen− t way to study the relation between two quantities.3 We find a strong correlation (Spearman correlation coefficient r 0.76 for the primary sample and s ∼ r 0.65 for the secondary sample, both corresponding to a probability P (r ) >99.9% s ∼ s 3We tested this hypothesis fitting our data with functional forms similar to the ones proposed by Meurer et al. (1999) and Kong et al. (2004): no significative improvement in the scatter of this relation is obtained. 88 7. UV dust attenuation in normal star forming galaxies

Figure 7.1: Ratio of the total infrared to far ultraviolet luminosity as a function of the ultraviolet spectral slope (lower x-axis) and the FUV-NUV color (upper x-axis). Open circles indicates our secondary sample while filled circles represent the primary sample. The dashed line represents the best linear fit to starburst IRX-UV relation. The solid line indicates the best bisector linear fit for our primary sample. The stars indicate the sample of IUE starbursts. Mean error bars for the plotted data are shown in the lower right corner, in this and subsequent figures. The residuals from the best linear fit for normal galaxies are shown in the bottom panel. 7.3. The L /L β relation for normal star-forming galaxies 89 T IR F UV −

that the two variables are correlated) between the total infrared to far ultraviolet ratio and the spectral slope, but significantly different from the one observed for starburst galaxies (dashed line in Fig.7.1; Meurer et al. 1999). A χ2 test rejects at a confidence level higher than 99.9%, the hypothesis that the two samples follow the same relation. The best linear fit for our primary sample (solid line in Fig.7.1) is:

L log( T IR ) = (0.70 0.06) β + (1.30 0.06) (7.5) LF UV  ×  The uncertainty in the estimate of the L /L using equation (7.5) is 0.26 0.02 T IR F UV ∼  dex for the primary sample but it increases to 0.35 0.03 dex, if we consider the whole sample (e.g. primary and secondary samples),∼ consistent with the mean uncertainty observed for starburst galaxies (Meurer et al. 1999). A large contribution ( 0.21 0.02 dex) to the observed scatter in Eq.(7.5) is due to the uncertainty on the estimate∼ of L /L and β. This result confirms once more that the L /L β T IR F UV T IR F UV − relation for normal galaxies deviates from the one observed for starbursts, as pointed out by previous studies of nearby galaxies (i.e. Bell 2002; Kong et al. 2004; Boissier et al. 2005; Buat et al. 2005; Seibert et al. 2005; Burgarella et al. 2005, Boissier et al. in prep.) and individual HII regions in nearby galaxies (Calzetti et al. 2005).

7.3.1 The dependence on the birthrate parameter

What physical mechanisms drive the difference observed in the L /L β be- T IR F UV − tween normal star forming galaxies and starbursts? Recently Kong et al. (2004) interpreted the offset as an effect of the different star formation history experienced by galaxies and proposed that the distance from the starburst IRX-UV can be pre- dicted using the birthrate parameter b (e.g. the ratio of the current to the mean past star formation activity, Kennicutt et al. 1994). In order to test if the perpendicular distance dS from the LT IR/LF UV β relation for starbursts correlates with the star formation history of normal galaxies,− we compute the birthrate parameter following Boselli et al. (2001): SF Rt (1 R) b = 0 − (7.6) L (M /L )(1 DM ) H tot H − cont where R is the fraction of gas that stellar winds re-injected into the interstellar medium during their lifetime ( 0.3, Kennicutt et al. 1994), t0 is the age of the galaxy (that we assume 12 Gyr),∼ DM is the dark matter contribution to the ∼ cont Mtot/LH ratio at the optical radius (assumed to be 0.5; Boselli et al. 2001). We compute the H-band luminosity following Gavazzi et al. (2002a):

log L = 11.36 0.4 H + 2 log(D) [L ] H − × × 90 7. UV dust attenuation in normal star forming galaxies

Figure 7.2: Relation between the birthrate parameter computed from the Hα emis- sion, and the distance from the L /L β relation for starbursts. The solid line T IR F UV − represents the best linear fit.

where D is the distance to the source (in Mpc), and the SFR from the Hα lumi- nosity (corrected for [NII] contamination and for dust extinction using the Balmer decrement) following Boselli et al. (2001):

LHα SF R = [M /yr] (7.7) 1.6 1041 × Fig.7.2 shows the relation between the birthrate parameter (eq.7.6) and the distance from the LT IR/LF UV β relation for starburst galaxies. The two quantities are correlated (r 0.40, corresp− onding to a correlation probability P (r ) 99.8%) but s ∼ s ∼ with a large scatter. Given the value of observational uncertainties, it is not worth trying to use the observed trend to reduce the dispersion in the LT IR/LF UV β relation for normal galaxies. This result confirms that part of the dispersion in−the LT IR/LF UV β relation for normal star forming galaxies appears an effect of the different star−formation history experienced by galaxies, as proposed by Kong et al. (2004).

7.4 A(Hα)

7.4.1 Estimate of A(Hα) The attenuation in the Balmer lines can be deduced from the comparison of the observed ratio LHα/LHβ with the theoretical value of 2.86 obtained for the recombi- nation case B, an electronic density n 104 cm−3 and temperature 104 K. The e ≤ ∼ 7.4. A(Hα) 91

variation of this value with density its negligible and with temperature is 5% (in the range between 5000 K and 20000 K, Caplan & Deharveng 1986). The ≤underly- ing absorption was deblended from the Hβ emission line using a multiple component fitting procedure. To do this the emission line is measured and subtracted from the spectra. The resulting absorption line is also measured with respect to a reference continuum. These two measurements are used as first guess in a fitting algorithm which fits jointly the emission and absorption lines to the reference continuum. For objects whose Hβ is detected in emission but the deblending procedure is not ap- plied (no absorption feature is evident) a mean additive correction for underlying absorption equal to -1.8 in flux and -1.4 Ain˚ EW is used. These values correspond to the fraction of the (broader) absorption feature that lies under the emission line. We adopt a dust screen geometry and the Milky Way extinction curve (e.g. Kenni- cutt 1983; Calzetti et al. 1994). Whereas varying the extinction curves has negligible effects in the visible, the dust screen assumption seems to under-estimate the extinc- tion by 0.2 mag compared with the amount deduced from the measurements of the thermal∼radio continuum (Caplan & Deharveng 1986; Bell & Kennicutt 2001). We do not apply any correction for Hα underlying absorption (Charlot & Longhetti 2001). However, since all the objects have EW (Hα + [NII]) > 3A,˚ the underestimate in the value of A(Hα) is negligible. In fact no change (at a 99% significance level) is observed comparing the best fits obtained in this work and the ones obtained adding to the Hα the same fixed underlying absorption used for Hβ when the underlying is not detected. We assume that the errors on A(Hα) are mainly due to the uncertainty on the Hβ flux. These errors represent in fact the lower limits because we do not ac- count for the uncertainty introduced by the fitting of the lines. They range from 0.01 to 0.43 mag and are found strongly anti-correlated with EW(Hβ) (see Gavazzi et al. 2004). Adopting the definition of the Balmer decrement as in Gavazzi et al. (2004):

1 LHα log( 2.86 L ) C1(Hβ) = × Hβ (7.8) 0.33 Since the A(Hα) attenuation is:

1 1 L A(Hα) = 1.086 ln( Hα ) (7.9) e 1 2.86 × L βα − Hβ From (7.8) and (7.9) we obtain: 1 A(Hα) = 1.086 0.33 C1(Hβ) ln(10) (7.10) e 1 × × βα − 92 7. UV dust attenuation in normal star forming galaxies

and assuming a galactic extinction law (eβα = 1.47) we derive:

A(Hα) = 1.756 C1(Hβ) (7.11) × A(Hα) = 0.85 mag is obtained on average, consistent with previous studies (e.g. Ken- nicutt 1983, 1992; Thuan & Sauvage 1992; Kewley et al. 2002). Eleven galaxies have Hβ undetected in emission but the underlying stellar absorption is clearly detected. For them we derive a 3 σlower limit to the Hβ flux (f ) using (Gavazzi et al. 2004): × Hβ

f < 3 rms − Hα(HW HM) (7.12) Hβ × (4500 4800) × assuming that Hα and Hβ emission lines have similar HWHM (Half Width Half Maximum). As shown in Eq.(7.8) a change in the theoretical value of the LHα/LHβ ratio would only produce a small ( 5%) constant over (or under) estimate of the ionized gas attenuation, thus leaving≤ unchanged the shape and dispersions of the observed relations, only affecting the values of the best fitting parameters.

7.4.2 The β-A(Hα) relation Calzetti et al. (1994) found a strong relationship between the ultraviolet spectral slope β and the Balmer decrement Hα/Hβ. For our starburst sample these two quantities are correlated (r 0.81) as follows (see also blue stars in Fig.7.3): s ∼ β = (0.75 0.10) A(Hα) (1.80 0.13) (7.13)  × −  This empirical relation was used by Calzetti et al. (1994) to deduce an attenuation law (the Calzetti law), often applied to high redshift galaxies (i.e. Steidel et al. 1999; Glazebrook et al. 1999). Contrary to the L /L β relation the Calzetti law T IR F UV − has not yet been tested for a sample of normal star forming galaxies. Buat et al. (2002) showed that for normal star forming galaxies the attenuation derived from the Calzetti law is 0.6 larger than the one computed from F IR/UV ratio. This ∼ result has been recently confirmed by Laird et al. (2005) on star forming galaxies at z 1. In order to check the Calzetti law on our sample we use the measure of the ∼ Hα/Hβ described in the previous subsection. Fig. 7.3 shows the relation between β and A(Hα) for our sample (empty and filled circles). For the primary sample we obtain r 0.58 (P (r ) >99.9%) and: s ∼ s β = (0.37 0.07) A(Hα) (1.15 0.08) (7.14)  × −  flatter than for starburst galaxies (see Fig.7.3). At low A(Hα) normal galaxies show on average a less steep ultraviolet spectral slope than starbursts. In addition normal galaxies with the same value of β span a range of 1 mag in A(Hα). At higher ∼ 7.4. A(Hα) 93

Figure 7.3: The relation between the ultraviolet spectral slope β and the Hα atten- uation obtained from the Balmer decrement. Symbols are as in Fig.7.1. Solid line represents the best linear fit to our primary sample (equation 7.14) while the dashed line indicate the best-fit for starburst galaxies obtained by Calzetti et al. (1994) (equa- tion 7.13). Arrows indicate galaxies for which the value of A(Hα) is a lower limit of the real value (i.e. Hβ undetected). The residuals from the best linear fit for normal galaxies are shown in the bottom panel. 94 7. UV dust attenuation in normal star forming galaxies

attenuation the two samples appear consistent. Our result suggest that the Calzetti law cannot be applied to normal galaxies. On the contrary, the relation between β and A(Hα) for normal galaxies, could be used to obtain a new attenuation law.

7.5 Relations between dust attenuation and global properties.

7.5.1 Metallicity Heckman et al. (1998) have shown that the ultraviolet spectral slope and metallic- ity of starbursts are well correlated. To determine the metal content of our galaxies we average five different empirical determinations based on the following line ratios: R ([OII]λ3727 + [OIII]λ4959, 5007)/Hβ (Zaritsky et al. 1994; McGaugh 1991), 23 ≡ [NII]λ6583/[OII]λ3727 (Kewley & Dopita 2002), [NII]λ6583/Hα (van Zee et al. 1998) and [OIII]λ5007/[NII]λ6583 (Dutil & Roy 1999). The mean uncertainty in the abun- dances is 0.10dex. In Fig. 7.4 we study the relationship between the gas metallicities and the LT IR/LF UV ratio (left-panel) and β (right-panel) for normal star forming and starburst galaxies. For normal galaxies the L /L ratio correlates (r 0.59, T IR F UV s ∼ P (rs) >99.9%) with the gas abundance: L log( T IR ) = (1.37 0.24) 12 + log(O/H) (11.36 2.11) (7.15) LF UV  × − 

with a dispersion of 0.35 0.03 in log(LT IR/LF UV ). As for the LT IR/LF UV β relation normal galaxies∼ differ from starbursts. At comparable metallicity normal− galaxies show a lower LT IR/LF UV (lower attenuation) than starbursts, in agreement with the recent result by Boissier et al. (2004) who studied radial extinction profiles of nearby late-type galaxies using FOCA and IRAS observations. Unexpectedly we find however that normal star forming galaxies follow exactly the same (signif- icant, r 0.58, P (r ) >99.9%) relationship between metallicity and ultraviolet s ∼ s spectral slope β determined for starbursts by Heckman et al. (1998) (see right panel of Fig.7.4). This might indicate that even though a normal and a starburst galaxy with similar gas metallicity have similar UV spectral slopes, they suffer from a sig- nificantly different dust attenuation, perhaps suggesting a different dust geometry (Witt & Gordon 2000). However we stress that this effect might occur due to aper- ture effects in the IUE data: while β is not significantly contaminated by aperture effects, the LT IR/LF UV ratio could be overestimated producing the observed trend (the total infrared luminosity is obtained by integrating the IRAS counts over the full galaxy extension, while the ultraviolet one is taken from IUE’s significantly smaller aperture 20 10 arcsec2). This idea could be supported by the correlation (r 0.49, × s ∼ 7.5. Relations between dust attenuation and global properties. 95

Figure 7.4: Relation between gas metallicity and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig.7.1. The solid lines show the best linear fit for our primary sample. The residuals from the best linear fits for normal galaxies are shown in the upper panels.

Figure 7.5: Relation between the galaxy size and the LT IR/LF UV ratio for starburst (left panel) and normal galaxies (right panel). Symbols are as in Fig. 7.1. Mean values and uncertainties in bins of 0.30 log(Diameter) are given. 96 7. UV dust attenuation in normal star forming galaxies

P (rs) >99.9% see Fig.7.5) observed between the starbursts’ optical diameters and the LT IR/LF UV ratio, completely absent in our sample of normal galaxies (rs 0.006, P (r ) 25%). GALEX observations of starburst galaxies will rapidly solve this∼ riddle. s ∼

Dust to Gas ratio The correlation between attenuation and metallicity can be interpreted assuming that the ultraviolet radiation produced by star forming regions suffers a dust attenuation increasing with the dust to gas ratio, which correlates with metallicity. (e.g. Issa et al. 1990; Inoue 2003). In order to check this hypothesis we compute the dust to gas ratio following Boselli et al. (2002b). In normal galaxies the dust mass is dominated by the cold dust emitting above 200 µm. The total dust mass can be estimated provided that the 100-1000 µm far-IR∼ flux and the cold dust temperature are known. Fitting β the SEDs of normal galaxies with a modified Planck law ν Bν (Td), with β = 2 (Alton et al. 2000), the total dust mass can be determined from the relation (Devereux & Young 1990): 2 a/T M = CS D (e dust 1) M (7.16) dust λ − where C depends on the grain opacity, Sλis the far-IR flux at a given wavelength (in Jy), D is the distance of the galaxy (in Mpc), Tdust is the dust temperature, and a depends on λ. Only IRAS data at 60 and 100 µm are available for our sample and, given the strong contamination of the emission at 60 µm by very small grains, the 60 to 100 µm ratio does not provide a reliable measure of Tdust (Contursi et al. 2001). Tdust determined by Alton et al. (1998) consistently with Contursi et al. (2001), seems to be independent of the UV radiation field, of the metallicity or of the total luminosity (Boselli et al. 2002b). Therefore we will adopt the average value Tdust = 20.8 3.2 K for all our galaxies introducing an uncertainty of 50% on the estimate of M ∼ dust (equation (7.16)). We then estimate the dust mass of the sample galaxies using (7.16) −1 −2 with C = 1.27 M Jy Mpc , consistent with Contursi et al. (2001), and a=144 K for Sλ = S100 µm (Devereux & Young 1990). The determination of the dust to gas ratio, in a way consistent with that obtained in the solar neighbourhood, requires the estimate of the gas and dust surface densities, thus of the spatial distribution of dust and gas over the discs. Unfortunately only integrated HI and H2 masses are available for our spatially unresolved galaxies. It is however reasonable to assume that the cold dust and the molecular hydrogen are as extended as the optical disc (Alton et al. 1998; Boselli et al. 2002b). To determine the mean HI surface density we adopt (Boselli et al. (2002b)):

log Σ = 20.92( 0.17) 0.65( 0.11) (def(HI)) cm−2 HI  −  × where def(HI) is the galaxy HI deficiency. Thus the dust to gas ratio is obtained from 7.5. Relations between dust attenuation and global properties. 97

Figure 7.6: Relation between the gas to dust ratio and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample.

the ratio of the dust surface density to the sum of molecular and neutral hydrogen surface densities. In Fig. 7.6 we compare the relation between the LT IR/LF UV ratio (left panel) and β (right panel) with the dust to gas ratio. The gas to dust ratio barely correlates with the LT IR/LF UV ratio (R 0.38). Contrary to metallicity, we do not find a significant correlation (R 0.11) with∼the ultraviolet spectral slope. This is ∼ probably due to the high uncertainty in our estimate of Mdust consequent to assuming a/Tdust the same temperature for all our galaxies (Mdust e , thus small errors ( 15%) on T propagate onto 50% errors on M ). ∝ ∼ dust ∼ dust

7.5.2 Luminosity Since it is well known that the metallicity of normal galaxies strongly correlates with galaxy luminosity (e.g. Skillman et al. 1989; Zaritsky et al. 1994) and mass (e.g. Tremonti et al. 2004), it is worth considering the correlation between attenuation and galaxy luminosity. Fig.7.7 shows the relationships between the dust attenuation indi- cators LT IR/LF UV and β and the H-band luminosity. The infrared to far ultraviolet ratio correlates (r 0.49, P (r ) >99.9%) with the total H-band luminosity: s ∼ s L L log( T IR ) = (0.34 0.10) log( H ) (2.66 0.88) (7.17) LF UV  × L −  The dispersion of this relation is 0.39 0.03 in log(L /L ). Since the H-band ∼  T IR F UV luminosity is proportional to the dynamical mass (Gavazzi et al. 1996), this implies a relationship between dust attenuation and dynamical mass. Also in starbursts the 98 7. UV dust attenuation in normal star forming galaxies

Figure 7.7: Relation between the H-band luminosity and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panel.

Figure 7.8: Relation between the TIR+FUV luminosity and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. 7.5. Relations between dust attenuation and global properties. 99

total H-band luminosity is correlated (rs 0.37, P (rs) 99.5%) with the LT IR/LF UV ratio and the great part of starbursts appear∼ offset (to 99%∼ confidence level) from the relation of normal galaxies. On the contrary, no difference is observed between the two samples in the β-LH plot, in agreement with what observed for metallicity. Fi- nally Fig.7.8 shows the relation between the bolometric luminosity (LTIR +LFUV) and the dust attenuation, computed assuming that the UV emission is absorbed by dust and emitted in the far infrared. The correlation coefficient (rs 0.31, P (rs) 98%) indicates that the two quantities correlate, as for starburst galaxies∼ (Heckman∼et al. 1998). This is not the case if we examine the relation between the ultraviolet spectral slope β and the bolometric luminosity (Fig.7.8 right panel): while there is no corre- lation (r 0.002, P (r ) 20%) for our sample of normal galaxies, a clear relation s ∼ s ∼ (rs 0.68, P (rs) >99.9%) holds for starbursts. Starbursts with higher bolometric luminosit∼ y (high TIR emission) show lower ultraviolet slope, consistent with the idea that high TIR emission corresponds to high attenuation (low β).

7.5.3 Surface brightness Wang & Heckman (1996) interpreted the increase of dust attenuation with rotational velocity (or mass) as due to the variations in both the metallicity and surface density of galactic disk with galactic size. Fig.7.9 shows the variation of the effective H-band surface brightness (defined as the mean surface brightness within the radius that contains half of the total galaxy light) and the dust attenuation. The two quantities are strongly anti-correlated (r -0.63, P (r ) >99.9%): s ∼ s

LT IR log( ) = ( 0.28 0.04) µe(H) + (5.92 0.81) (7.18) LF UV −  ×  with a scatter of 0.34 0.03 in log(L /L ): 1.2σ lower than the value ∼  T IR F UV ∼ obtained for H-band luminosity and consistent with the one obtained for the gas metallicity. Unfortunately in this case we cannot compare the behaviour of normal galaxies with the one of starbursts due to the lack of an estimate of µe for the star- bursts. Does this relation indicate that UV dust extinction depends on the thickness of stellar disk, or does it follows from the correlation between attenuation and star formation surface density? To attack this question we determine the SFR density (de- fined as the ratio between the SFR determined from Hα (eq.7.7) and optical galaxy area). Fig.7.10 shows the relation between the SFR density and log(LT IR/LF UV ). The two quantities are correlated (rs 0.44, P (rs) >99.9%) with a dispersion of 0.39 0.03 in log(L /L ), 1.2σ∼larger than the one observed for the mean ∼  T IR F UV ∼ H-band surface brightness4. Since the contribution of observational uncertainties to

4The difference between the two relations does not change if instead of the half-light radius, we use the total radius to estimate µe(H) 100 7. UV dust attenuation in normal star forming galaxies

Figure 7.9: Relation between the mean H-band surface brightness (µe) and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panel.

Figure 7.10: Relation between the star formation rate density and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample.The residuals from the best linear fit for normal galaxies are shown in the upper panel. 7.5. Relations between dust attenuation and global properties. 101

the scatter in the two relations is the same (0.18 0.02), our result might suggest that the UV attenuation is primarily∼ correlated with the thickness of stellar disk, supporting the hypothesis of Wang & Heckman (1996) that both gas metallicity and star surface density are directly connected with the physical properties of dust (i.e. quantity and spatial distribution).

7.5.4 LHα/LF UV ratio

Buat et al. (2002) suggested that the LHα/LF UV ratio could be another potential attenuation indicator but they found a scattered correlation between LHα/LF UV and A(F UV ), confirmed by Bell (2002). This correlation is expected since both Hα and UV emission are star formation indicators. The Hα luminosity comes from stars 7 more massive than 10 M and it traces the SFR in the last 10 yr while the UV ≤ luminosity comes from stars of lower mass (M 1.5 M ) and it can be used as an indicator of the SFR in the last 108 yr. This≥ means that under the condition that the star formation is approximately≈ constant in the last 108 yr the ratio ≈ LHα/LF UV (corrected for attenuation) should be fixed. Thus the ratio between the extinction corrected LHα and the observed LF UV should be a potential attenuation indicator. In Fig.7.11 we analyze the relationship between the dust attenuation and the LHα/LF UV ratio, where LHα is the Hα luminosity corrected for dust attenuation using the Balmer decrement and for the contamination of [NII]. The two quantities turn out to be strongly correlated (r 0.76, P (r ) >99.9%): s ∼ s L L log( T IR ) = (0.84 0.07) log( Hα ) (0.59 0.12) (7.19) LF UV  × LF UV − 

The dispersion around this relation is 0.24 0.02 in log(LT IR/LF UV ), consistent with the one observed for the log(L ∼/L ) β relation. The high correlation T IR F UV − and low scatter between the two quantities is expected since the two variables are mutually related: the FUV luminosity appears in both axes and LT IR and LHα are known to be correlated (Kewley et al. 2002), explaining why in the left panel of Fig.7.11 starbursts and normal galaxies show the same trend. The right-panel of Fig.7.11 shows the relation between the ultraviolet slope and the LHα/LF UV ratio. In this case starbursts and normal galaxies behave differently: at any given β starbursts have an higher LHα/LF UV than normal galaxies, consistent with what expected for galaxies experiencing a burst of star formation (Iglesias-P´aramo et al. 2004).

A secure determination of the Balmer decrement for large samples is still a hard task, especially at high redshift, thus we look for a relation similar to Eq.(7.19) using obs obs the observed Hα luminosity (LHα). The LHα/LF UV and log(LT IR/LF UV ) ratios are yet correlated (see Fig.7.12) but the correlation coefficient is lower than the previous 102 7. UV dust attenuation in normal star forming galaxies

Figure 7.11: Relation between the Hα and far ultraviolet luminosity and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. Hα luminosity is corrected for dust attenuation using the Balmer decrement, while the FUV flux is uncorrected. The solid lines show the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panels.

Figure 7.12: Relation between the observed Hα and far ultraviolet luminosity and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. Hα luminosity is the observed value not corrected for dust attenuation. The solid lines show the best linear fit for our primary sample.The residuals from the best linear fit for normal galaxies are shown in the upper panels. 7.6. A cookbook for determining LT IR/LF UV ratio 103

case (r 0.49, P (r ) >99.9%). The best linear fit gives: s ∼ s L Lobs log( T IR ) = (1.10 0.17) log( Hα ) (0.59 0.21) (7.20) LF UV  × LF UV −  with a mean absolute deviation of 0.34 0.03 ( 3.3σ higher than for Eq.7.20). ∼  ∼

7.6 A cookbook for determining LT IR/LF UV ratio

In this chapter we investigated the relations between dust attenuation, traced by the LT IR/LF UV ratio, and other global properties of normal star forming galaxies. Fur- thermore we compared the dust attenuation in normal and starbursts galaxies using multiwavelength datasets. The amount of dust attenuation is found to correlate with the UV colors, gas metallicity, mass and mean surface density but, generally speaking, differently for normal and starburst galaxies. Determine whether this difference is real or is due to aperture effects requires the analysis of GALEX observations for a sample of starburst galaxies. The dispersion in the LT IR/LF UV β relation correlates with the birthrate parameter b, suggesting that the observed scatter− is, at least partly, due to differences in the star formation history. These results stress that estimating the UV dust attenuation, and consequently the star formation rate of normal galaxies (at high redshift in particular) is highly uncertain ( 50%) when rest-frame far infrared observations are not available. Moreover the sample≥ selection criteria could strongly affect its properties, as recently pointed out by Buat et al. (2005) and Burgarella et al. (2005). They studied the dust attenuation properties and star formation activ- ity in a UV and in a FIR selected sample, showing that the former shows correlations with global galaxy properties, such as mass and bolometric luminosity, that the FIR selected sample does not. Their results stress that the dust attenuation properties are very heterogeneous and that LT IR/LF UV cannot be derived in a robust manner when far infrared observations are not available. However the present investigation has shown that among optically-selected samples of normal galaxies with no nuclear activity a number of empirical relations exists, allowing to derive the LT IR/LF UV ratio (and its uncertainty). Once the attenuation at UV is determined it can be transformed to any other λ, only knowing the shape of the attenuation law and dust geometry (i.e. Calzetti et al. 1994; Gavazzi et al. 2002a; Boselli et al. 2003a). In Table 7.1 we list all the relations, their associated r.m.s., mean absolute deviation from the best fit (m.a.d.)5 and the Spearman correlation coefficient.

5The mean absolute deviation is less sensitive to the contribution of outliers than the standard deviation. For a Gaussian distribution the mean absolute deviation (m.a.d.) is 2/π (r.m.s.), ∼ × while it is lower (higher) for a heavier (lighter) tailed distribution. As shown in Tpable 1 the values 104 7. UV dust attenuation in normal star forming galaxies

Before we proceed describing our recipes, we have to investigate whether the scatter in these relations is physical or is only driven by observational uncertainties. In the latter case, in fact, our cookbook would not be very useful, since it would be valid only for observations with the same uncertainties as our datasets. For H-band luminosity, obs H-band surface brightness, LHα/LF UV ratio and metallicity the contribution of obser- vational uncertainties to the observed scatter varies from 18% (r.m.s. 0.17 0.02) ∼ obs ∼  for LH to 40% (r.m.s. 0.21 0.02) for 12 + log(O/H) and LHα/LF UV : even ac- counting for∼ the contribution∼ ofmeasurements errors, the relative difference in the scatter of these relations does not change. On the contrary this confirms that the relation involving LH is the one with the highest ”physical” dispersion, while for the other three relations the scatter is similar. The situation is worse for the relations involving β and the LHα/LF UV ratio: the contribution of observational errors is 70-76% ( 0.21 0.02). Thus it is impossi- ∼ ∼  ble to determine which of these two relations has the lowest scatter and represents the best way to estimate dust attenuation without far infrared observations. We can conclude that observational errors could account for the difference scatter observed in the relations involving β and the LHα/LF UV ratio, but not for the difference observed in all the other relations. Our results can thus be used to suggest different ways to correct for UV dust attenuation. Ia) The LT IR/LF UV β relation still represents one of the best way to quantify dust attenuation. The uncertain− ty in the value of log(L /L ) is 0.26 0.03. Ib) T IR F UV ∼  If the UV spectral slope β is unknown but we know LHα (corrected for attenuation) we can obtain the ultraviolet attenuation using equation (7.19), with a r.m.s. of 0.24 0.02. This relation is valid under the assumption that the star formation rate  is approximately constant in the last 108 yr. obs ≈ IIa) If we know LHα, but no estimate of A(Hα) is available, we can use Eq.(7.20) (rms 0.34 0.03). IIb) ∼If neither β nor Hα luminosity are available we are left with the relations with H-band surface brightness6 (r.m.s. 0.34 0.03) and, in the worse case, ∼  III) H-band luminosity (rms 0.39 0.03 ). ∼  Summarizing, these relations allow us to estimate the value of the LT IR/LF UV ratio with an average uncertainties of 0.32 dex. This value corresponds approximately to ∼ σ(A(F UV )) 0.5 mag, assuming log(LT IR/LF UV ) = 1 (the mean value for our sam- ple) and using∼ the model of Buat et al. (2005). This is the lowest uncertainty on the estimate of the LT IR/LF UV ratio in absence of far infrared observations. However we caution the reader that this value holds only for an optically-selected sample and that samples selected according to different criteria, especially FIR-selected, could contain higher dispersions. The cookbook presented in this chapter is obviously insufficient obtained for r.m.s. and m.a.d. are consistent with the ones expected for a Gaussian distribution 6Since we need Hα flux to estimate metallicity, Eq.(7.15) cannot be used in this case. 7.6. A cookbook for determining LT IR/LF UV ratio 105

a b x a b m.a.d. rms rs β 0.70 0.06 1.30 0.06 0.20 0.02 0.26 0.02 0.76 12+log(O/H) 1.37  0.24 11.36 2.11 0.26  0.02 0.35  0.03 0.59  −    LH /L 0.34 0.10 2.66 0.88 0.29 0.03 0.39 0.03 0.49 µ (H) 0.28 0.04 −5.92 0.81 0.25  0.02 0.34  0.03 0.63 e −     − LHα/LF UV 0.84 0.07 0.59 0.12 0.19 0.02 0.24 0.02 0.76 Lobs /L 1.10  0.17 −0.59  0.21 0.27  0.02 0.34  0.03 0.49 Hα F UV  −    a: Mean absolute deviation from the best fit. b: Standard deviation from the best fit.

Table 7.1: Linear realtions useful to estimate the LT IR/LF UV ratio (log(L /L ) = a x + b). T IR F UV ×

to understand dust attenuation and know how to correct UV observations of local and high redshifts galaxies, but it represents only the tip of the iceberg. The next steps should be the folowings: a) compare all the relations obtained in this work with different models in order to try to determine the physical properties of dust b) use models and data in order to estimate a new attenuation law from the far-ultraviolet to the near-infrared valid for normal star forming galaxies, as the one obtained for starbursts by Calzetti et al. (1994). Only knowing the dust attenuation law we will be able to correct for dust extinction all our observations and thus to correctly estimate the star formation rate in galaxies.

Chapter 8

High velocity interaction: NGC4438 in the Virgo cluster

This analysis represents the tip of the iceberg and only a future comparison with different dust models will allow us to understand dust attenuation and to know how to correct UV observations of local and high redshifts galaxies. A statistical analysis of star formation activity in cluster galaxies using UV data is therefore still impossible. For this reason, in the last three chapter of this work, I will focalize my attention on the study of three particular cluster galaxies considered as the prototypes of the three main environmental effects observed in clusters: tidal interaction, ram pressure stripping and preprocessing, respectively. These unique astrophysical laboratories will be used to deeply understand the effects of different physical mechanisms on galaxy evolution..

8.1 Introduction

NGC 4438 (Arp 120) is the clearest example of an ongoing tidal interaction in a nearby cluster of galaxies. Apparently located close to the Virgo cluster center ( ∼ 300 kpc from M87), NGC 4438 is a bulge-dominated late-type spiral showing long tidal tails (30 kpc) thought to be induced by a recent dynamical interaction with the nearby SB0 galaxy NGC 4435. Multifrequency observations covering the elec- tromagnetic spectrum from X-rays (Kotanyi et al. 1983; Machacek et al. 2004) to radio continuum (Hummel & Saikia 1991), including both spectro-photometric and kinematical (Kenney et al. 1995; Chemin et al. 2005) data, have been carried out in the past to study the nature of this peculiar system. These observations have shown that the violent interaction between the two galaxies perturbed the atomic (Cayatte et al. 1990) and molecular (Combes et al. 1988) gas distribution, causing both gas infall toward the center which might have induced nuclear activity (Kenney et al.

107 108 8. High velocity interaction: NGC4438 in the Virgo cluster

1995; Kenney & Yale 2002; Machacek et al. 2004), and gas removal in the external parts displacing part of the gas in the ridge in between the two galaxies (Combes et al. 1988). Both multifrequency observational data (Kenney et al. 1995; Machacek et al. 2004) and model predictions (Combes et al. 1988; Vollmer et al. 2005) favor a recent ( 100 Myr) high-velocity, off-center collision between NGC 4435 and NGC ∼ 4438. Except for mild nuclear activity, it is still unclear whether the dynamical interac- tion between the two galaxies induced extra-nuclear star formation events: the low Hα/[NII] ratio and the similar X-ray and Hα morphology of NGC 4438 indicate that the Hα emission is in this case not due to the ionizing radiation but is probably due to gas cooling phenomena (Machacek et al. 2004). The UV emission is dominated by young stars of intermediate masses (2 < M < 5M ) and provides us with an alternative star formation tracer. As part of the Nearby Galaxy Survey (NGS), we have observed the central 12 deg2 of the Virgo cluster using the Galaxy Evolution Explorer (GALEX). A distance of 17 Mpc for Virgo is adopted.

8.2 Data

The GALEX data used in this work include far-ultraviolet (FUV; λeff = 1530A˚, ∆λ = 400A)˚ and near-ultraviolet (NUV; λeff = 2310A˚, ∆λ = 1000A)˚ images. The data consist of 2 independent GALEX pointings centered at R.A.(J2000)= 12h29m01.2s, Dec(J2000)= +13◦10’29.6” (819 sec) and R.A.(J2000)= 12h25m25.2s, Dec(J2000)= 13◦10’29.6” (1511 sec), for a total of 2330 sec of integration time. To study the star formation history of NGC 4438, the UV data have been combined with visible and near-IR images of the galaxy taken from the GOLDmine database (Gavazzi et al. 2003a), from the SDSS Data release 3 (Abazajian et al. 2005) from the 2MASS survey (Jarrett et al. 2003) and from the CFHT and SUBARU archives. These are Hα+[NII] (Boselli & Gavazzi 2002), B (Boselli et al. 2003a), K’ (Boselli et al. 1997b), u, g, r, i, z SDSS, R CFHT and SUBARU and H 2MASS images. For the main body of the galaxy (region 4 in Fig. 8.1, see next sect.) we added the integrated spectrum (3500-7000 A˚ ; Gavazzi et al. 2004). The current calibration errors of the NUV and FUV magnitudes are on the order of 10% (Morrissey et al. 2005), comparable to that at other frequencies. ∼ 8.2. Data 109

Figure 8.1: The combined NUV and FUV image of NGC 4438. The regions described in sect. 3 of the text are labeled 1 to 7. The horizontal line is 10 kpc long (assuming a distance of 17 Mpc). 110 8. High velocity interaction: NGC4438 in the Virgo cluster

8.3 The UV emission and the star formation his- tory of NGC 4438

Figure 8.1 shows the UV image of NGC 4438, obtained by combining together the NUV and FUV frames in order to increase the S/N. The UV emission of the galaxy is mostly due to compact, bright regions in the central part of the galaxy (marked as region 4 in Fig. 8.1), in the northern tidal tail (region 2) and in the section of the southern tail closest to the main body of the galaxy (region 5). The UV emission is mostly diffuse in the extended western part of the galaxy (region 3) and at the edge of the southern tidal tail (region 6). In addition Figure 8.1 shows the presence of two extended and patchy emission to the north-west of the galaxy ( 15-25 kpc from the nucleus, marked as region 1 and region 7). These features, previously∼ undetected in other visible and/or near-IR bands, are similar to a tidal tail: region 1 is 20 kpc long and 2 kpc wide, while region 7 is considerably smaller ( 2 kpc). The∼ RGB ∼ ∼ image of the galaxy obtained by combining the FUV, NUV and B frames (see Fig.8.4) shows the color of the different regions: while the edge of both the northern and the southern tidal tails (region 3 and 6) are red (thus dominated by relatively old stars), regions 2 and 5 as well as the newly discovered regions 1 and 7, have blue colors and seem therefore to be dominated by a younger population. The Hα+[NII] emission map, given in Fig. 8.2 as a contour plot superposed on the NUV image of NGC 4438, shows a lack of massive, ionizing young O-B stars (Kenni- cutt 1998). The Hα+[NII] emission observed in region 5 has a different morphology than the UV one; on the contrary its distribution is the same observed in X-ray as stated by Machacek et al. (2004) (see Fig.8.3). This evidence confirms the conclu- sions of Machacek et al. (2004) that the Hα+[NII] emission is not due to the ionizing radiation but is probably associated with the cooling gas. What is the nature of the newly discovered extragalactic UV emitting regions? The average NUV surface brightness of these features is 28.5 ABmag arcsec−2, while ∼ they are undetected both in the SUBARU R band (360 sec) image down to a sur- face brightness limit of 27.8 mag arcsec−2 and in Hα down to a surface brightness − ∼− − − limit of 5 10 17 erg s 1 cm 2 arcsec 2 (see Fig.8.2), implying a log(NUV/Hα) 0.3. Extra-planar∼ diffuse regions with an excess of UV over Hα flux ratio (i.e.≥ log(NUV/Hα) 1, as that observed at 11 kpc from the disk of M82, are often in- ≥ terpreted as due to the UV radiation produced by the central starburst and locally scattered by diffuse dust (Hoopes et al. 2005). It is unlikely that scattered light is responsible for the UV emission since the steep slope of the UV spectrum (β=-2.32 and -2.05, as defined by Kong et al. 2004) is typical of a recent unreddened star- burst (Calzetti 2001) and is unexpected in a scattering scenario since the dust albedo is greater in the NUV than in the FUV (Draine 2003) (i.e.β 1). Furthermore the lack of a powerful central starburst (as in M82) and the large≤ −distance of these 8.3. The UV emission and the star formation history of NGC 4438 111

Figure 8.2: The Hα+[NII] contours (red, in arbitrary scale, in between 8 10−17 and 6 10−16 erg cm−2 s−2 arcsec−2, with σ= 5 10−17 erg cm−2 s−2 arcsec−2, from Boselli & Gavazzi (2002)) are superposed to the NUV gray-level image of NGC 4438.

Figure 8.3: Chandra image of NGC 4438 in the 0.3-2 keV energy band with Hα contours superposed. Adapted from Machacek et al. (2004) 112 8. High velocity interaction: NGC4438 in the Virgo cluster

relatively patchy regions from the nucleus seem to exclude the scattering scenario. These data suggest that regions 1 and 7 are post starbursts, induced by the violent interaction with NGC 4435. In addition the absence of Hα emission associated with all the UV emitting regions suggests that the starburst lasted for a relatively short time, since it is not producing young, massive O-B stars any more. This is probably because the atomic and molecular gases, needed to feed star formation, have been removed during the interaction (Combes et al. 1988; Vollmer et al. 2005)1 . In order to date the starburst and reconstruct the star formation history of the galaxy, we have determined the spectral energy distribution (SED) of each region (see Fig.8.4) and then fitted it with a simple model of galaxy evolution. To this end we make the assumption that dust attenuating the SED is present only in region 4, where we cor- rect the UV to near-IR data using the far-IR to UV flux ratio as done in Boselli et al. (2003a) and described in Appendix A. This restricted application is reasonable since no dust emission has been observed in the tidal tails with ISOCAM (Boselli et al. 2003b); furthermore, in regions 1 and 7, dust is unexpected since it has not yet been produced by the young stellar population, as confirmed by the steep β parameter (see also Chapter 7). Assuming that NGC 4438 was a normal late-type object before interacting with NGC 4435, we use the models of Boissier & Prantzos (2000) in order to reconstruct its SED before the interaction. The two parameters of the model (spin λ and rotational velocity VC ) are constrained by the observed total H-band luminos- −1 ity and velocity rotation of NGC4438, leading to λ=0.01 and VC =290 km s . In Fig.8.4 we compare the model with the SED of the main body of the galaxy (region 4), composed by an old population with no significant contribution from the recent starburst. Both the total SED and the optical spectrum produced by the model are in good agreement, confirming that the adopted technique is able to reproduce the galaxy SED before the interaction. We then assume that the evolved stellar popula- tion of each region, if present, is the one given by the model and removed from the main body of the galaxy by the tidal interaction, while the younger population is produced by the induced starburst. For each region, we thus combine the SED of an evolved stellar population with the one produced by an instantaneous burst of star formation obtained using Starburst 99 (Leitherer et al. 1999) for a solar metallicity and a Salpeter IMF between 1 and 100 M . For each age and intensity of the burst, we determine the best combination of evolved population+ burst by fitting the FUV to K band SED and rejecting solutions in disagreement (i.e. too bright) with the upper limits. We then adopt the age corresponding to the lowest reduced χ2 2. The

1The upper limit of the HI surface density for these regions is 1 M pc−2 (Cayatte et al. 1990) 2All ages with χ2 < 1 are acceptable solutions. Given the small∼ num ber of photometric points available for regions 1 and 7 (2 GALEX bands), the fitted solution for a combination of a burst and an old population (two parameters) can be almost perfect (resulting in very low χ2, 10−2), as long as the obtained fit is in agreement with the limits at other wavelengths. Whenever≤the fit produces a SED not satisfying a detection limit, this solution is rejected. 8.4. Discussion and conclusion 113

results of our fitting procedure are presented in Fig.8.4. For each region (excluding region 4) two panels are given. The lower panels show the observed SED of each region (crosses, or arrows if are upper limits) and the best SED obtained from the fitting procedure (black line). The relative contribution of the evolved and young stellar populations to the observed SEDs are indicated in red and blue respectively. The burst luminosity contribution (for the age corresponding to the minimum χ2) in the band FUV, B and K is also given. In the upper panels the variation of the reduced χ2 parameter (black continuum line) and of the burst mass fraction (red dotted line) as a function of the age of the burst are given. This exercise gives an interesting result: the strong UV emission of regions 1 and 7 is due to a coeval starburst 6-20 ∼ Myr old. The age and the duration of the starburst are strongly constrained both by the lack of Hα emission and by the blue UV slope of the spectrum (lower limit to the age) and by the lack of an old stellar population (upper limit to the duration). The burst age for the other region cannot be determined with the same precision, but we can only put a lower limit to their age. Regions 2 and 5 are consistent with an older starburst ( 100 Myr, as suggested by their redder UV slope: β=-0.33 and -0.67 in ≥ regions 2 and 5 respectively) which probably ended 10 Myr ago as indicated by the lack of any Hα emission. Conversely the stellar ∼population in regions 3 and 6 appear not significantly affected by the recent burst. Moreover it is interesting to note that, while the fraction of stars produced by this burst is dominant in regions 1 and 7, the sum of the stars produced by the burst in all regions (including the inner part) contributes to the total galaxy stellar mass by less than 0.1 %, an extremely low value for such a violent interaction.

8.4 Discussion and conclusion

These observations have major consequences in constraining the evolution of cluster galaxies. A high-velocity off-center collision between two galaxies of relatively similar mass, whose violence is able to perturb the stellar distribution producing important tidal tails, is insufficient to induce a significant instantaneous starburst. This result might be representative only of the nearby Universe where encounters of gas-rich galaxies are probably rare since clusters are dominated by gas-poor early-type galax- ies such as the companion galaxy NGC 4435. It is conceivable, however, that at higher redshifts, where clusters are forming, stellar masses produced by a starburst induced by interactions predicted by the models of Moore et al. (1996) (galaxy harassment) might be more important given the higher fraction of gas-rich galaxies. The other interesting result is the long time differential between the age of the in- teraction ( 100 Myr as determined by dynamical simulations, (Combes et al. 1988; ∼ Vollmer et al. 2005) and the beginning of the starburst ( 10 Myr in regions 1 and 7, 100 Myr in regions 2 and 5). This result is totally consisten∼ t with the models of ∼ 114 8. High velocity interaction: NGC4438 in the Virgo cluster

Mihos et al. (1991) that predict for close-by encounters an enhancement of the star formation activity in the inner disk during some 100 Myr, stopping once the gas reser- voir is exhausted as in NGC 4438. In the tidal tails, on the contrary, star formation is expected to increase after 100 Myr, the time needed by the gas to re-collapse, but then ceasing after a few ∼Myr because the expansion of the tidal tail brings the gas surface density to subcritical values (no HI and CO has been detected in these regions). If these systems are dynamically stable and survive the interaction, they might be at the origin of some dwarf galaxies in the cluster similar to those observed in the Stephan’s Quintet by Mendes de Oliveira et al. (2004) or in other interacting systems (Neff et al. 2005; Hibbard et al. 2005; Saviane et al. 2004) Being produced by a single starburst, these gas poor systems might evolve into dwarf ellipticals, typical of rich clusters. Otherwise they will simply increase the fraction of unbound stars, contributing to the Virgo intracluster light (Willman et al. 2004). 8.4. Discussion and conclusion 115

Figure 8.4: The RGB (FUV=blue, NUV=green, B=red) color map of NGC 4438 and NGC 4435. The SED of each region defined in Fig. 1 are given in the lower plot of each frame. Crosses indicate the observed data, arrows upper limits (in mJy), the red dashed line the evolved population fit as determined by the model of Boissier & Prantzos (2000), the dotted blue line the starburst SED (from Starburst 99) and the dashed green line the combined fitting model. The burst luminosity contribution (for the age corresponding to the minimum χ2) in the band FUV, B and K is also given. The upper panel gives the variation of the reduced χ2 parameter (black continuum line, in logarithmic scale) and of the burst mass fraction (red dotted line) as a function of the age of the burst (in Myr). The lower panel of region 4 gives the integrated 3500 to 7000 A,˚ R=1000 spectrum of the main body of the galaxy (black continuum line) compared to the fitted model (red dashed line).

Chapter 9

Ram Pressure stripping: NGC4569 in the Virgo cluster

9.1 Introduction

Spiral disks can lose their atomic gas content during dynamical interactions with the hot and dense intergalactic medium (IGM) (Gunn & Gott 1972) and/or in direct interactions with nearby objects (Merritt 1983). These interactions can quench their star formation activity (Gavazzi et al. 2002c) leaving the objects to become anemic (van den Bergh 1976). To explain the well known morphological segregation effects Dressler (1980) it has been suggested that these quiescent spirals could evolve into lenticulars; however, observations and model predictions give still contradictory re- sults (see Boselli & Gavazzi 2005 for a review). Despite the on-going physical processes (tidal interactions were probably dominant at early epochs, while galaxies-IGM interactions are more important at present), it is clear that the fuel supply needed to feed star formation is more efficiently removed where the host-galaxy potential well is weakest, i.e., in the outer disk. Given the strong relation between the gas surface density and the star formation activity in spiral disks, commonly known as the Schmidt law (Kennicutt 1998; Boissier et al. 2003), it is expected that star formation will be quenched in the outer (lower density) portions of the disk. While interferometric observations of galaxies in Virgo have clearly shown that HI disks are less extended in those objects located close to the cluster center (Cayatte et al. 1990), the observational evidence for a truncation of the star forming disks has been proven by Hα imaging (Koopmann & Kenney 2004b,a). Although a truncation of the disk profile has been predicted (Larson et al. 1980b), we still do not know what the passive evolution of a stellar disk is once its gas is removed. In particular, it is unclear whether the progressive radial suppression of star formation is able to reproduce the structural properties of lenticulars, generally characterized as

117 118 9. Ram Pressure stripping: NGC4569 in the Virgo cluster

having higher surface brightness of their stellar disks and higher bulge-to-disk ratios than spirals (Dressler 1980). We have been collecting multi-frequency data for a large sample of late-type galax- ies in nearby clusters and in the field in order to undertake comparative statistical analyses of any systematic differences between cluster and field objects. Combined with multi-zone models for the chemical and spectrophotometric evolution of galax- ies (Boissier & Prantzos 2000), this unique database is helping us understanding the evolution of cluster spirals. As a first step during my thesis I studied the radial pro- files of the Virgo cluster galaxy NGC 4569 (M90). NGC 4569, the prototype anemic galaxy as defined by van den Bergh (1976), is extremely deficient in HI, having just one tenth of the atomic gas of a comparable field galaxy of similar type and di- ∼mensions. The galaxy has a truncated Hα and HI radial profile (at a radius of 5 kpc; see Fig.9.3) as firstly noticed by Cayatte et al. (1994) and Koopmann & Kenney∼ (2004a), witnessing a recent interaction with the cluster environment. NGC 4569 is located close ( 1 degree) to the cluster center. Being one of the largest galaxies ( 10 arcmin) in the∼ Virgo cluster, NGC 4569 is the ideal candidate for our study since∼ it can be spatially resolved at almost all wavelengths considered here. The combi- nation of the multi-frequency 2-D data with our spectrophotometric models allow us to study, for the first time, the radial evolution of the different stellar populations in this prototype, gas-stripped cluster galaxy with the aim of understanding whether its structural properties can evolve into those of a typical cluster lenticular (S0) galaxy.

9.2 Data and models

The large amount of spectrophotometric data available for NGC 4569, collected in the GOLDMine database (Gavazzi et al. 2003a), allow us to reconstruct its radial profile at different wavelengths: from the new GALEX UV bands (at FUV=1530 and NUV=2310 A),˚ to the visible B and V (Boselli et al. 2003a), Sloan u, g, r, i, z (Abaza- jian et al. 2005) and near-IR J, H and K bands (Boselli et al. 1997b; 2MASS Jarrett et al. 2003). Hα+[NII] narrow band imaging, used to trace the recent star formation activity, is available from Boselli & Gavazzi (2002). HI profiles are from Cayatte et al. (1994), while H2 profiles, determined from CO data using a luminosity dependent CO to H2 conversion factor (from Boselli et al. 2002b) are taken from the BIMA survey of Helfer et al. (2003) for the inner disk, and from Kenney & Young (1988) for the outer disk. The accuracy of the photometric imaging data is, on average 10 %. ∼ The galaxy rotation curve has been taken from Rubin et al. (1999). Unfortunately no metallicity gradient information is available for NGC 4569. The radial profiles have been constructed by integrating the available images within elliptical, concentric annuli. The ellipticity and position angles have been determined and then fixed using the deepest B band image following the procedure of Gavazzi 9.2. Data and models 119

et al. (2000) (see Fig.9.3). To avoid any possible contamination by the NW arm, whose kinematical properties indicate that it is not probably associated with the stel- lar disk but rather formed during the interaction with the ICM (Chung et al. 2005), the arm was masked in the construction of the radial profiles. If included, its con- tribution would be perceptible only in the FUV filter at radii > 8 kpc, increasing the surface brightness by < 0.5 mag. The UV to near-IR broadband images of the galaxy have been corrected for internal extinction using the recipe of Boissier et al. (2004), assuming a typical UV extinction gradient for a galaxy of the luminosity and scalelength of NGC 4569. In this case, in fact, the far-IR (IRAS) to UV flux ratio can- not be used to estimate the extinction because both fluxes are contaminated by the nuclear activity of the galaxy. The extinction in the other visible and near-IR bands has been determined using the prescription of Boselli et al. (2003a) and described in Appendix A. Hα+[NII] narrow-band imaging has been corrected for [NII] contami- nation and dust extinction (Balmer decrement) using the integrated spectroscopy of Gavazzi et al. (2004). To study the evolution of the disk of NGC 4569 at various radii, we have used the multi-zone chemo-spectrophotometric models of Boissier & Prantzos (2000), updated with an empirically-determined star formation law Boissier et al. (2003). These models have a resolution of 1 kpc, significantly lower than the ∼ one of our multiwavelength datasets (0.08-0.4 kpc). The errors in the surface bright- ness and color profiles have been computed following Gil de Paz & Madore (2005). For this reason we degraded all our images at the model resolution, and we extract the smoothed profiles used for the comparison between models and data. The nuclear emission due to the central AGN has been masked since the model is not able to reproduce the AGN activity (see Fig.9.4). The two model parameters (spin λ and rotational velocity VC ) are constrained by the H-band luminosity profile (determined assuming a distance of 17 Mpc) and the rotation curve of the galaxy, making the reasonable assumption that both of these observables are unperturbed during the in- −1 teraction. This gives λ=0.04 and VC = 270 km s (see Fig.9.1). To compute the Hα profile, the number of ionizing photons predicted by Version 5 of STARBURST 99 V´azquez & Leitherer (2005) for a single generation of stars distributed on the Kroupa et al. (1993) initial mass function (as used in our models) is convolved with our star formation history, and converted into a Hα flux as described in Appendix B. In addition to this model (valid for an unperturbed galaxy) we add an episode of ram pressure gas stripping. For simplicity, we adopt the plausible scenario of Vollmer et al. (2001): the galaxy has crossed the dense IGM only once, on an elliptical orbit. The ram pressure exerted by the IGM on the galaxy ISM varies with time (t) following a Lorentzian profile (see Fig.9.2):

(∆t)2  =  (9.1) 0 ((∆t)2 + (t t )2) − 0 120 9. Ram Pressure stripping: NGC4569 in the Virgo cluster

Figure 9.1: The radial profile of observed (open symbols) and extinction-corrected (filled symbols) H-band surface brightness (left) and of the rotational velocity (center) used to constrain the model without interaction (represented by the black solid line). The total gas radial profile (right) predicted by the unperturbed model (solid black line) is compared to the observed one (green filled circles), obtained by summing the HI component (red line) to the molecular one (blue and light blue) and correcting for Helium contribution ( 1.4), and to the model including the interaction (black dashed line). ×

where t0 is when the galaxy crosses the dense cluster core at high velocity and 0 is the value of ram pressure at t0. Following Vollmer et al. 2001 we assume a width profile ∆t = 9 107 years. In order to determine the amount of stripped gas we make the hypothesis× that the gas is removed at a rate that is directly proportional to the galaxy gas column density Σgas and inversely proportional to the potential of the galaxy, measured by the total (baryonic) local density Σpotential (provided by the Σgas model). The gas-loss rate adopted is then finally equal to  Σpotential . The two free parameters in our model are then t0 and 0. We make the further assumption that no extra star formation is induced during the interaction.

9.3 The star formation history of NGC 4569: model predictions

Once the width of the interaction event, ∆t, is fixed, it is possible to choose simul- taneously t0 and 0 because the amount of gas left and its radial distribution depend strongly on 0 while the resulting stellar light profiles depend mainly on t0 (see Fig.9.3 for some examples). If the cluster core crossing time is recent only the youngest stel- lar populations (emitting in Hα, whose age is 4 106 yrs, or far-UV, 108 yrs) have had time to feel the progressive radial suppression≤ of the star formation≤ activity. 9.4. Discussion and conclusion 121

Figure 9.2: Ram pressure stripping intensity (in arbitrary units) as a function of time (Eq.9.1). Adapted from Vollmer et al. (2001).

Comparing model predictions with the spectrophotometric radial profiles of cluster galaxies can thus be used to date the dynamical interaction with the IGM. We thus fitted the data with models for different values of t0 and 0. An important modi- fication applied to the usual χ2 test is that its value was artificially put to 100 for any model predicting surface brightnesses in disagreement with observational limits (non detections at relatively large radii) in order to reject these solutions. The model best matching the properties of NGC 4569 (Fig.9.4) is characterized by 0 = 1.2 M −2 −1 kpc yr and t0 = 100 Myr. This is largely consistent with the dynamical models of Vollmer et al. (2004a), who obtained t 300 Myr. Although not reproducing 0 ∼ perfectly the surface brightness profile, this model is able to qualitatively reproduce the truncation of the total gas disk profile (see Fig.9.1) and of the Hα and UV radial profiles (Fig.9.3) as well as the milder truncation observed at longer wavelengths. It is interesting to note that although older cluster core crossing epochs give more truncated disk profiles in the old stellar populations (B and i bands, blue dashed line), this is not the case in the gas profile which is modified by contributions from the recycled gas.

9.4 Discussion and conclusion

The present work gives the first quantitative estimate of the structural evolution of stellar disks in cluster galaxies due to gas removal caused by a dynamical interaction 122 9. Ram Pressure stripping: NGC4569 in the Virgo cluster

Figure 9.3: The radial profile of the observed (empty green circles) and extinction- corrected (filled green circles) total gas, Hα, FUV (1530 A),˚ NUV (2310 A),˚ B and i surface brightness. The yellow shaded area marks the range in between the observed (bottom side) and extinction-corrected (top side) surface brightness profiles. Surface brightnesses are compared to the model predictions without interaction (black solid line) or with interaction for several 0 and t0 parameters. Equal maximum efficiency −2 −1 (0=1.2 M kpc yr ) and different age: t0=100 Myr, red continuum line (the adopted model); t0=500 Myr, grey long dashed line, t0=1.5 Gyr, dashed magenta −2 line. Equal age (t0=100 Myr) and different maximum efficiency: 0=3 M kpc −1 −2 −1 yr , blue dotted line; 0=1/3 M kpc yr , orange dotted line. 9.4. Discussion and conclusion 123

of the galaxy with the IGM. Although the model only qualitatively reproduces the observed multi-wavelength radial profiles (the mismatch being attributed to resolu- tion effects) it delivers a strong message concerning the passive stellar evolution of stripped disks. First of all it is clear that the truncation of the total gas disk profile is soon reflected in the young population stellar disk, confirming the predictions of Larson et al. (1980b). As observed in NGC 4569, this gas-stripped galaxy has a color gradient opposite to that of normal, isolated spirals, which generally have bluer colors in their outer disks (see Fig.9.5 and Fig. 9.4). NGC 4569 is bluest towards the center. The trend is especially true for colors tracing the relatively young populations (< 108 yr); colors tracing populations older than the interaction event present the usual∼ gradient (i.e., redder towards the center). The inversion of the color gradient, here observed for the first time in a cluster galaxy, is well reproduced by our model. The consequence of these findings in the interpretation of the evolution of cluster spiral galaxies is significant. One of the most intriguing and still open question re- garding the effects of the environment on the evolution of galaxies is that of the origin of lenticulars, and their overabundance in the centers of rich clusters. Are lenticulars an independent population of galaxies formed in the primordial high-density environ- ments, or were they spiral disks whose star formation activity has been quenched once their gas reservoir was removed by the unfavorable cluster environment? Although the second interpretation seems logical, simple statistical considerations in the semi- nal work of Dressler (1980) show that this idea is not supported by the observations: spirals have lower surface brightnesses and bulge-to-disk ratios than lenticulars, and spirals are rotationally supported while lenticulars are dynamically hotter systems. Structural (and kinematical) modifications must thus be invoked if spirals are to be transformed into lenticulars. The present work has shown for the first time how a galaxy-cluster IGM interaction is able to remove most of the gas reservoir inducing important structural modifica- tions in the disk properties. We have in fact shown that, because of the differential radial stellar evolution of spiral disks, we expect that cluster spirals have (at least at short wavelengths) more truncated disk profiles, inverting the outer color gradient with respect to similar but unperturbed objects. The surface brightness of the disk, however, mildly decreases in Hα and in the UV bands while remains mostly constant at longer wavelengths even 5 Gyr after the interaction (Fig. 9.4d). The differential evolution of the stellar disk due to gas stripping alone is thus not able to reproduce the structural properties of present-day lenticulars. Gravitational perturbations, such as tidal interactions with other galaxies (Merritt 1983), interactions with the cluster potential well (Byrd & Valtonen 1990) or a mixture of both (called ‘galaxy harass- ment’ by Moore et al. 1996) must be invoked to reproduce the observed properties of nearby lenticulars. This new study and analysis is consistent with the idea that the present evolution of late-type galaxies in clusters differs from that at earlier epochs, where late-type galax- 124 9. Ram Pressure stripping: NGC4569 in the Virgo cluster ies were mostly perturbed by dynamical interactions (pre-processing and/or galaxy harassment; Dressler 2004, Moore et al. 1996) which were able to thicken the stel- lar disks thereby producing the present-day cluster lenticulars. We hope to confirm this original result in the near future once multi-frequency data come available for a statistical significant sample of late-type cluster galaxies. 9.4. Discussion and conclusion 125

(a) (b)

(c) (d)

Figure 9.4: The observed and model surface brightness (a), color (b) radial profiles of NGC 4569. In the model profiles the continuum lines are for models with gas removal, dashed lines for unperturbed models. c) the reduced χ2 as a function of the look-back −2 −1 time of the ram-pressure event for a few efficiencies 0 (M kpc yr ). Models were computed each 100 Myr for 0 < t0 < 500 Myr, 200 Myr for 1.5 < t0 < 0.5 Gyr and −2 −1 1 Gyr for 6.5 < t0 < 1.5 Gyr; and each 0.2 M kpc yr efficiencies between 0.4 and 1.6 (only the more relevant are shown here). d) the variation of the effective surface brightness (mean surface brightness within Re, the radius containing half of the total light) and radius due to differential variation of the star formation history of NGC 4569. Open triangles are for the unperturbed model, the other symbols for different ages of the interaction (100 Myr, 1.5 and 5.5 Gyr). 126 9. Ram Pressure stripping: NGC4569 in the Virgo cluster

Figure 9.5: The RGB (continuum subtracted Hα =blue, FUV=green, NUV=red) color map of NGC 4569 Chapter 10

Galaxy Pre-processing: the blue group infalling in Abell1367

10.1 Introduction

In the previous two chapters we have investigated the effects of the environment on the properties of galaxies inhabiting the core of the Virgo cluster. However galaxies interact with the harsh environment well before having reached the center of a cluster. In particular, if we believe that structures grow hierarchically, galaxy clusters form not by accreting individual galaxies randomly from the field, but rather through the infall of less massive groups falling in, along large scale filaments. Galaxy groups may therefore represent a natural site for a preprocessing stage in the evolution of cluster galaxies. These infalling groups have velocity dispersions that are significantly smaller than that of cluster, permitting the slow gravitational interaction typically observed in field galaxies. Moreover even in compact groups ram pressure seems to be able to displace the gas from the disk of galaxies (Fujita 2004; Roediger & Hensler 2005). This means that probably at least part of the morphological and star for- mation properties of cluster galaxies derives from earlier epochs and very different conditions than the ones observed in today clusters (Dressler 2004). Environmental interactions in the infalling groups may thus represent a preprocessing step in the evolution of cluster galaxies (Mihos 2004a). Unfortunately, witnessing preprocessing in local Universe is a real challenge since we live in a Λ-dominated Universe where the infall rate is significantly lower than in the past (Gottl¨ober et al. 2001). Today, we observe a plethora of clusters experiencing multiple merging (Gavazzi et al. 1999a; Donnelly et al. 2001; Cortese et al. 2004), but the structures involved are subclusters 14 with a mass 5 10 M , considerably higher than the typical mass of a compact 13∼ × group 10 M (Mulchaey 2000), as the North and South subclusters in Abell1367 studied∼in Chapter 5 (see Table 5.7). However Abell1367 represents a unique excep-

127 128 10. Galaxy Pre-processing: the blue group infalling in Abell1367 tion among local, dynamically young, clusters since in addition to massive evolved substructures it is also experiencing the merging of a compact group infalling directly into the cluster core. This group has a velocity dispersion of only 170km s−1, and ∼− it is infalling into the cluster core at a very high speed ( 1700km s 1). The rarity of this phenomenon could probably explain the unique prop∼erties observed in this group. In fact it was independently discovered by Iglesias-P´aramo et al. (2002) and Sakai et al. (2002) during two deep Hα surveys of nearby clusters, representing the region with the highest density of star forming systems ever observed in the local Universe. Sakai et al. (2002) argued that this group lies in the cluster background, having no interaction with the cluster environment. On the contrary the dynamical analysis presented in Chapter 5, is consistent with an infalling scenario, as also proposed by Gavazzi et al. (2003b). Moreover this picture is supported by X-ray observations: Sun & Murray (2002) (using Chandra observations) discovered extended gas features and a ridge near the SE cluster center. They proposed that these features are associated with a new merging component penetrating the SE subcluster. XMM clearly detects a cold front near the center of the SE subcluster, probably associated with a group infalling into the cluster core (A. Finoguenov, private comm.). All these observational evidences suggest that we are witnessing, for the first time in the local Universe, a compact group infalling into a core of a dynamically young cluster. It thus represents a unique laboratory to study with the great detail possible only in the local Universe, a physical process typically expected in clusters at high redshift. The study of this group could therefore help us shading light on the possible influence that preprocess- ing might have and have had on the past evolution of galaxies now populating high density environments. During the last few years we thus collected a great amount of multiwavelength spec- troscopic and imaging observations in order to try to reconstruct the history of this rare group of galaxies, which represents the only compact group infalling into the center of a galaxy cluster ever observed in the nearby Universe. Throughout this chapter I will refer to this group as the Blue Infalling Group (BIG), as defined by Gavazzi et al. (2003b)

10.2 Observations

10.2.1 HI observations Using the refurbished 305-m Arecibo Gregorian radio telescope we observed the BIG region in March 2005. We obtained observations for 4 different positions covering the group center and its NW outskirt (see Fig.10.1). Data were taken with the L-Band Wide receiver, using nine-level sampling with two of the 2048 lag subcorrelators set to each polarization channel. All observations were taken using the position-switching 10.2. Observations 129

Figure 10.1: The four Arecibo HI pointings obtained in the region of the BIG group, superposed to the r0 band image. The size of each circle correspond to the telescope beam. 130 10. Galaxy Pre-processing: the blue group infalling in Abell1367

Figure 10.2: GALEX NUV image of the Blue Infalling group (BIG).

technique, with each blank sky (or OFF) position observed for the same duration, and over the same portion of the telescope dish as the on-source (ON) observation. Each 5min+5min ON+OFF pair was followed by a 10s ON+OFF observation of a well- calibrated noise diode. The velocity resolution was 2.6 km s−1, the instrument’s beam at 21 cm is 3.05 3.01 and the pointing accuracy is about 1500. Flux density calibration corrections are×good to within 10% (and often much better), see the discussion of the errors given in O’Neil (2004). Using standard IDL data reduction software available at Arecibo, corrections were applied for the variations in the gain and system temperature with zenith angle and azimuth. A baseline of order one to three was fitted to the data, excluding those velocity ranges with HI line emission or radio frequency interference (RFI). The velocities were corrected to the heliocentric system, using the optical convention, and the polarizations were averaged. All data were boxcar smoothed to a velocity resolution of 12.9 km s−1 for further analysis. 10.2. Observations 131

10.2.2 UV to near-IR imaging The Blue Infalling Group has been observed by GALEX in April 2004, within the two pointings of the Abell cluster 1367. The observations are centered at R.A.(J2000)= 11:43:41.34 Dec(J.2000)=+20:11:24.0 (e.g. offset to the north of the cluster to avoid a star bright enough to threaten the detector), with a mean exposure time of 1460s, as described in Chapter 4. Fig.10.2 shows the GALEX NUV image of the Blue Infalling Group. UBVRH photometry for CGCG (Zwicky et al. 1961) galaxies is taken from Gavazzi et al. (2003a).

10.2.3 Hα imaging We observed BIG using the Device Optimized for the LOw RESolution (DOLORES) attached at the Nasmyth B focus of the 3.6m TNG in the photometric nights of 17th May and 18th June, 2004. The observations were taken through a [SII] narrow band filter centered at 6724A˚ and a width of 57A˚ covering the redshifted Hα ∼ ∼ and [NII] lines. The underlying continuum was taken through a broadband (Gunn) r0 filter. Images, split in 6 exposures of 1200 sec in the narrow band filter and 5 exposures of 300 sec in the r0 broadband filter, for a total of 2 hours and 30 minutes exposure respectively, were taken with a seeing of 1.2 arcsec. The photometric calibration was achieved by exposing the spectrophotometric∼ star Feige 34. After bias subtraction and flat-fielding, the images were combined. The intensity in the combined OFF-band frame was normalized to that of the combined ON-band one by the flux ratio of several field star. The NET image was obtained by subtracting the normalized OFF-band frame to the ON-band one. The resulting OFF and NET-band frames are shown in Figs. 10.6 and 10.7 respectively. Hα+[NII] fluxes and EWs are obtained as described in Boselli et al. (2002a).

10.2.4 MOS spectroscopy We observed the BIG region in MOS mode with the ESO/3.6m and with the TNG telescope. The ESO/3.6m observations were taken in the photometric nights of May 5th and 6th 2003 with the ESO Faint Object Spectrograph and Camera (EFOSC). We used the MOS mode of EFOSC to obtain the spectra of 9 of the emitting line knots. The EFOSC spectrograph was used with a 300 gr/mm grating and 2048 2048 thinned Loral CCD detector, which provided coverage of the spectral region 3860× 8070A.˚ − Slits width of 1.75” yielded a resolution of 19A.˚ We obtained eleven exposure of 1530 sec, for a total exposure time of 4.65∼hours. The TNG observations were taken in the∼ photometric nights of 26th March and 22nd April 2004 with DOLORES. We used the MOS mode of DOLORES to obtain the spectra of 8 of the emitting line knots, of the nuclear region of CGCG97-125 and of 14 132 10. Galaxy Pre-processing: the blue group infalling in Abell1367

Figure 10.3: High-contrast Hα+[NII] band frame of the BIG group. 10.2. Observations 133

Name R.A. Dec Velocity (J2000) (J2000) (km s−1) TNG ESO MOS Sakai02 Gavazzi03 − K1 114444.18 194816.0 8422 153 8265 117 8098   − DW3 d 114445.97 194744.4 8564 151 −  − − DW3 e 114445.97 194741.1 8072 124 −  − − DW3 a 114446.43 194741.2 8490 180 8266  − − 97-114b 114446.56 194640.3 8656 132 8504 8383 −  97-114a 114447.41 194649.8 8763 124 8425 −  − K2 a 114450.61 194605.1 8080 140 8070 8089 −  K2 b 114449.71 194604.7 8309 165  − − − DW2 c 114451.12 194718.7 8380 188  − − − DW2 b 114451.17 194717.5 8221 146 8077 −  − DW2 a 114451.67 194713.5 8253 292  − − − K5 114451.76 194752.7 8241 112 7995 −  − DW1 b 114453.78 194731.5 8070 − − − DW1 c 114454.29 194728.6 8343 223  − − − DW1 a 114454.64 194732.9 8265 136 8161 8067 −  97-125b 114454.89 194611.3 8261 191 8396 132 8170   − K3 114455.28 194803.3 8020 212  − − − 97-125a 114455.99 194628.0 8330 − − −

Table 10.1: Redshifts of the galaxies in the BIG group.

galaxies in the region. The DOLORES spectrograph was used with a grating which provided coverage of the spectral region 3200 8000A.˚ Slits width of 1.6” yielded a resolution of 17A.˚ We obtained six exposure−of 1800 sec, for a total exposure time ∼ of 3 hours. All the emitting line regions observed in MOS spectroscopy are shown in Fig.10.3 In addition we took spectra of the bright galaxy CGCG97-114 using the Loiano/1.52 m telescope. The BFOSC spectrograph attached at the Loiano telescope was used with a 300 gr/mm grating and 1300 1340 thinned EEV CCD detector, which pro- × vided a spectral coverage 3600 8900A.˚ A slit width of 2.00” yielded a resolution of 20A.˚ The observations were−taken in the ”drift-scan” mode, with the slit parallel ∼ to the galaxy major axis, drifting over the optical surface of the galaxy. The total exposure time was 2400 sec. The reduction of the spectra was carried out using standard tasks in the IRAF pack- age. Bias subtraction and flat-field normalization was applied using median of several bias frames and flat-field exposures. The various exposures were combined using a 134 10. Galaxy Pre-processing: the blue group infalling in Abell1367

median filter, thus removing the cosmic rays. The λ calibration was carried out using IDENT IF Y REIDENT IF Y F IT COOR on exposures of He/Ar lamps for each slit, and the calibration− was transferred− to the science frames using T RANSF ORM. Typical errors on the dispersion solution are of 0.5 1A,˚ as confirmed by the mea- surements of the sky lines. However, since the resolution∼ − of our spectra is 13A˚ we ≥ assume an rms of 3A˚ on our wavelengths calibration. The two-dimensional frames were sky subtracted using BACKGROUND. One-dimensional spectra were ob- tained integrating the signal along the slit using AP SUM. The apertures were lim- ited to regions where the signal intensity was above 1 σ of the sky noise. Spectra were flux-calibrated using the spectrophotometric standard star: ltt 3864 for the ESO, Feige 67 for the TNG and Feige 34 for the Loiano observations. The redshift of each knot was derived as the mean of the individual redshift obtained from each emission line. Our results are shown in Tab. 10.1 and compared with the previous measurements by Sakai et al. (2002) and Gavazzi et al. (2003b).

Line measurements All spectra were shifted to the rest frame wavelength and normalized to their intensity in the interval 5400-5600 A.˚ The flux-calibrated, normalized spectra are presented in Fig. 10.18. Under visual inspection of the spectra we carried out the measurement of the emission lines using SP LOT . This provided a list of fluxes and EWs with respect to a user defined continuum level. Hα (λ6563) is bracketed by the weaker [NII] doublet ([NII1] λ6548 and [NII2] λ6584). The three lines are not well resolved, thus using the task SPLOT we performed a two Gaussian fit to the blended emissions providing an estimate of the line ratio [NII]λ6584/(Hα + [NII]λ6548). The two bright galaxies CGCG97-125 and CGCG97-114 show evidence for underlying absorption in correspondence to emission lines. We de-blended the underlying absorption from the emission lines as discussed in Chapter 7. In order to compare our observations with the ones presented by Sakai et al. (2002) we re-measured, using the method described above, the spectra taken at the Stewart Observatory 2.3m Bok telescope and at the 6.5m MMT by these authors. The two sets of measurements presented in Tab. 10.2 are found in fair agreement.

10.2.5 High Resolution spectroscopy We obtained high dispersion long-slit spectra of CGCG97-125 and CGCG97-120 with the 1.93 mtelescope of the Observatoire de Haute Provence (OHP), equipped with the CARELEC spectrograph coupled with a 2048 512 TK CCD, giving a spatial scale of 0.54 arcsec per pixel. The observations were carried× out in the night of April 20, 2004 in approximately 2 arcsec seeing conditions through a slit of 5 arcmin 2 arcsec. The selected grism gives a spectral resolution of 33 A˚/mm or 0.45 A˚/pix×and a spectral 10.3. Results 135

Object Tel. C1 [OII] Hβ [OIII] [OIII2] Hα [NII2] K1 ESO (0.00) 3.68 1.00 0.94 2.53 − − K1 T NG 0.00 3.49 1.00 0.95 2.57 2.86 0.58 DW3 a T NG 0.02 8.26 1.00 0.31 1.06 2.86 0.59 DW3 d ESO 0.00 2.01 1.00 0.60 1.34 2.86 0.69 97-114b ESO 0.33 4.25 1.00 0.57 1.83 2.86 0.26 97-114a ESO 0.24 3.37 1.00 0.26 0.70 2.86 0.45 97-114 LOI 0.75 2.61 1.00 0.21 0.36 2.86 0.64 K2 a ESO 0.17 4.06 1.00 0.77 2.09 2.86 0.50 K2 b T NG 0.23 8.43 1.00 0.64 0.82 2.86 0.32 DW2 a T NG 0.16 5.45 1.00 0.73 0.99 2.86 0.60 DW2 b ESO > 0.1 > 5.18 1.00 < 1.00 < 0.99 2.86 < 0.61 K5 ESO 0.56 1.00 0.47 0.65 2.86 0.29 − DW1 b MMT 0.33 3.62 1.00 0.74 2.50 2.86 0.35 DW1 c T NG 0.00 2.76 1.00 0.47 1.53 2.86 0.66 DW1 a ESO 0.30 1.00 0.82 2.39 2.86 0.35 − DW1 a MMT 0.20 3.75 1.00 0.80 2.41 2.86 0.49 97-125b ESO 0.55 1.00 0.51 1.25 2.86 0.37 − 97-125b T NG 0.04 3.60 1.00 0.35 1.28 2.86 0.60 97-125 OHP 0.88 9.23 1.00 1.06 1.97 2.86 0.84 97-125 T NG 0.90 6.58 1.00 1.15 1.89 2.86 1.21 K3 T NG 0.00 2.87 1.00 0.25 0.97 2.86 0.89

Table 10.2: Line fluxes, corrected for internal extinction, of the galaxies in the BIG group. coverage in the region 6080-6990 A˚ containing the redshifted Hα ( λ 6562.8 A),˚ the [NII] doublet (λ 6548.1, 6583.4 A)˚ and the [SII] doublet (λ 6717.0, 6731.3 A).˚

10.3 Results

10.3.1 Kinematics Table 10.1 lists the positions and radial velocities of the objects that were measured spectroscopically. Our observations confirm the physical association of all the emit- ting line objects with the bright galaxies CGCG97-114 and CGCG97-125. On the contrary the brightest galaxy in this region CGCG97-120 seems not associated with this group, having a recessional velocity of 5635 km s−1 (see also Section 10.3.7). The −1 −1 velocity of galaxies in BIG (< V >= 8230 km s and σV = 170 km s ) exceeds −1 −1 significantly the mean cluster velocity of < V >= 6484 km s (σV = 891 km s ), 136 10. Galaxy Pre-processing: the blue group infalling in Abell1367

Figure 10.4: Upper panel: The position and the width (rectangular areas on the right) of the three slits obtained for CGCG97-125. The slits are superposed to the Hα + [NII] net image. Lower Panel: The three different rotations curves obtained for CGCG97-125. Letters indicate the different regions as labeled in the upper panel. 10.3. Results 137

Figure 10.5: The low resolution 2D spectrum obtained at ESO/3.6 for the knots DW3d (left) and DW3e (right), shows a significant difference ( 500km s−1) in the velocity of the two knots. ∼

suggesting that it is infalling at 1700 km s−1 into the cluster core. ∼ The high resolution spectra obtained at the OHP telescope give us more insights on the dynamical state of CGCG97-125. Velocity plots of CGCG97-125 were extracted from each spectrum by measuring the wavelength of the Hα line in each pixel along the slits. The three rotation curves so obtained are given in Fig.10.4. In each diagram the recessional velocity is plotted as a function of position along the slit (the spatial axis runs from E (left) to W (right)). All the three spectra show regions with mul- tiple velocity components, especially in correspondence to the galaxy center where two sudden velocity jumps of 100-150km s−1 are clearly present. It is interesting ∼ − to note that the velocity of these jumps decrease from 8400 km s 1 to 8150 km s−1 and their position moves to east, passing from the∼ north to the south∼ part of the galaxy. Even if several examples of kinematic disturbances has been observed in normal galaxies (Rubin et al. 1999; Haynes et al. 2000) and interacting systems (Jore et al. 1996; Duc & Mirabel 1998), the features observed in CGCG97-125 are extremely rare. To our knowledge, the only other galaxy with the same character- istics is UGC6697 (Gavazzi et al. 2001b), the merging systems in the NW part of Abell1367 (see Chapter 5). The velocity jumps observed in the rotation curve of CGCG97-125 are consistent with the idea that this galaxy has experienced a merging in the past; however its properties are unusual if compared with what expected from a similar phenomenon. During the accretion of a satellite, the gas falling into the galaxy center is expected to relax before the gas at the outskirts of the galaxy. The relaxation time is in fact R/V, where R is the radial distance from the center and ∝ V is the rotational velocity. On the contrary in this case, the major anomalies are observed near the galaxy center while in the outer part the rotation curve presents a 138 10. Galaxy Pre-processing: the blue group infalling in Abell1367

Name r0 Hα flux EW (Hα + [NII]) SF Ra −2 −1 −1 mag erg cm s A˚ M yr 97125 13.99 (1.33 0.29) 10−13 27 3 1.49  × −  97114 15.06 (6.59 0.71) 10 14 34 5 0.74  × −  DW1 17.93 (1.60 0.17) 10 14 128 15 0.09  × −  DW2 18.96 (3.67 0.86) 10 15 25 7 0.02  × −  DW3 19.11 (4.47 0.96) 10 15 56 15 0.02  ×  a: obtained using equation 7.7. L(Hα) corrected for [NII] contribution and extinction using values obtained from spectroscopy (see

Table 10.2).

Table 10.3: Properties of galaxies in BIG. typical S shape. Detailed dynamical simulations of a minor merger experienced by an S0 galaxy are thus mandatory to try to understand the particular features observed in this galaxy. The MOS spectroscopy collected at the ESO/3.6m and at the TNG telescopes gives us some information regarding the internal dynamic of DW3. The emitting line knots composing this system have considerably different recessional velocities, ranging be- tween 8000 km s−1 and 8600 km s−1. The western (DW3-b) and the eastern ∼ ∼ − (DW3-a) knots have a recessional velocity of 8250 8300 km s 1 significantly lower than the one observed in the northern knot DW3-d∼ ( −8564 km s−1) and 250 km s−1 higher than the redshift of the southern knot DW3-e.∼ This great difference∼ is clearly visible in Fig.10.5 where the two emitting line knots DW3-d and DW3-e are observed within the same slit (thus the relative offset is not affected by any uncertainty in the wavelength calibration). The observed high velocity gradient ( 500 km s−1) ∼ suggests that these five knots are probably not gravitationally bound, and thus that DW3 does not represents a dwarf virialized system.

10.3.2 Hα properties When observed in optical broad band images this group does not show any unex- pected feature if compared with other group of galaxies. The r0-band luminosity function of BIG in the interval -19.2 < Mr0 < -12.2 has a slope α -1, consistent 0 ∼ with the r -band luminosity function of Hickson compact groups (Hunsberger et al. 1998), suggesting that originally BIG was a normal compact group. On the opposite, BIG represents a true exception as far as its Hα properties. At least ten out of 12 star-forming regions are associated with dwarf systems (or extragalactic HII regions) with E.W.(Hα + [NII]) often exceeding 100 A(Ga˚ vazzi et al. 2003b). These galaxies, in spite of being 1000 times smaller than typical giant galaxies, are currently form- ing stars at a 10∼times higher rate (per unit mass) than normal galaxies of similar 10.3. Results 139

Figure 10.6: Stellar shells are seen around galaxy 97-125 in the r0 band image of BIG. No continuum emission is detected from the low brightness trails (except K2). 140 10. Galaxy Pre-processing: the blue group infalling in Abell1367

Figure 10.7: Extended low brightness trails appear in the Hα+[NII] NET frame of BIG. 10.3. Results 141 luminosity, as derived from their L(Hα) (see Table 10.3). As remarked by Sakai et al. (2002), it is the first time that such a high density of star-forming galaxies has been seen in a nearby cluster, in spite of having collected data over an area of A1367, Coma, and the Virgo Cluster approximately 500 times larger than the group size. Moreover new Hα images of BIG obtained last year reinforce the uniqueness of this group revealing a spectacular Hα filamentary structure on top of which the star form- ing knots observed by Sakai et al. (2002) and Gavazzi et al. (2003b) represent the tip of the iceberg. Multiple loops of ionized gas appear with a projected length exceeding 150 kpc, a typical transverse size of 5 kpc, among the most extended low-brightness Hα emission features ever detected (see Fig. 10.7). One stream (labeled NW in Fig. 10.7) extends from the northern edge of the frame to the dwarf galaxy DW3, with an extension of 100kpc. The second and brightest one (labeled W in Fig. 10.7) traces a loop around∼ galaxy 97-120 and seems connected to the bridge (labeled K2 in Fig.10.7) between 97-114 and 97-125, that was known from previous studies. If this is the case, the total projected extension of the NW and W trails would result 150 kpc. In addition to the filamentary features, at least two other diffuse Hα ∼ regions (labeled S and E in Fig.10.7) are detected. The total diffuse (Hα + [NII]) emission (e.g. excluding the contribution of the three bright galaxies and of the ten dwarfs/HII regions previously discovered) results 1.2 10−13 erg cm2 s−1 i.e. similar to the flux collected from one of the bright galaxies,∼ and×the typical surface brightness is 10−17.6 10−18.3 erg cm−2 s−1 arcsec−2. Along the filaments we detect typically an − E.W.(Hα + [NII]) 100 150 A.˚ The loop around 97-120≥ −alone contributes with 2.4 10−14 erg cm2 s−1, as ob- tained integrating the Hα + [NII] emission in a circular∼ × of 10 kpc radius and an annulus of 5 kpc, centered on 97-120. The derived line intensity is 2.05 Rayleigh (1 Rayleigh = 106/4π photons cm−2 s−1 sr−1), corresponding to an emission measure (EM) of 5.7 cm−6 pc. Assuming a torus geometry with a circular section of radius −2 −3 5 kpc and a filling factor of 1, the plasma density results ne 3.3 10 cm and ∼the ionized column density N 5 1020 cm−2 (the inferred densities∼ ×would be higher e ∼ × if the gas is in clumps or filaments, which is likely). The emission measure in the NW trails results lower ( 1.3 cm−6 pc) than in the loop around 97-120 and the plasma ∼−3 density is ne 1.1 cm . The trails geometry∼ is strongly suggestive of a rosetta orbit typical of tidal disruption of a satellite galaxy. However contrary to other known examples of tidal streams the features here observed show strong Hα emission and no continuum emission above −2 Σr0 = 26.8 mag arcsec (even though this limit is insufficient to rule out the pres- ence of stellar streams of brightness as low as observed for example in the M31 stream (Ibata et al. 2001)). The case offered by BIG seems therefore unique as it combines the, eventually present, faint stellar brightness of tidal streams with strong line emis- sions of tidal tails. What mechanism have produced a such unusual feature? In order to test the tidal disruption scenario we use the formalism of Johnston et al. (2001), 142 10. Galaxy Pre-processing: the blue group infalling in Abell1367

assuming: 1) that at least one of the gas trails is from a dwarf intruder merged into 97-125 and 2) that the geometry of the undetected stellar streams is the same of the observed gaseous trails. The intrinsic geometry of a streamer from a totally disrupted satellite can be used to estimate the mass m and age t of a young streamer:

3 2 11 w Rp vcirc m 10 M , (10.1) ∼ R 10 kpc 200 km/s and the time since its disruption

R R 200 km/s t 0.01 Ψ circ Gyr , (10.2) ∼ w  10 kpc  vcirc 

where w is the width of the streamer at radius R, Ψ is its angular length, Rp is the pericentric distance of the orbit and Rcirc is the radius of the circular orbit with the same energy as the true orbit. Of course, we cannot measure Rcirc directly, but we can approximate it as being halfway between the adopted apocenter and pericenter. Thus adopting a projected ratio of the loop width w, to the radius R, of 0.15, a ∼ pericentric distance of Rp 15 kpc, an orbit with the same energy of a circular orbit ∼ −1 of radius Rcirc 30 kpc and a rotation velocity Vcirc 298km s (Vogt et al. 2004), ∼ 9 ∼ we obtain a satellite mass 1 10 M and an age of the interaction 1.3 Gyr. The mean surface brightness∼of ×the tidal debris is then obtained using the∼ following equation:

10M /L ,ν 1 Gyr µ 0 (t) = 2.5 log f r −   Υ   t 

2.5 vcirc m 10 kpc 0 log 8 + 23.9 + M ,r ,(10.3) − 3 200 km/s10 M  R 

0 0 where M ,r0 is the r absolute magnitude of the Sun, Υ is the r mass to light ratio of the satellite and f is the mass fraction loses by the satellite. The mean surface brightness of the tidal debris in BIG is

3 f µ 0 = 26.8 2.5 log (10.4) r − Υ  1 

−2 consistent with the undetection of continuum emission above Σr0 = 26.8 mag arcsec . However I stress the reader that this simulation is based on the interaction between two field galaxies and not infalling into the cluster center as in this case. The most dramatic difference between field mergers and those in a cluster is in the evolution of tidal debris. In the field, most of the material stripped into tidal tails remains loosely bound to the host galaxy, forming a clear tracer of the gravitational interac- 10.3. Results 143 tion. In the cluster encounter, the cluster tidal field quickly strips the material from the galaxy, dispersing it throughout the cluster and making these tidal tracers very short-lived (Mihos 2004a). Thus we can assume the obtained value of µr as a lower limit for the surface brightness in the stellar trails. Although the undetection of stellar emission in the trails does not help us ruling out a tidal stream nature for these trails, their strong Hα emission makes BIG a unique ex- ample among interacting systems and compact groups. Conversely other known tidal tails have E.W.(Hα) ranging from zero (i.e. the Stephan’s quintet) to 20 A˚ (i.e. ∼ the Mice (NGC 4676) and some Hickson compact groups). Moreover tidal streams discovered in interacting systems (e.g. Shang et al. 1998; Forbes et al. 2003) are de- tected only in continuum with no Hα emission, even if associated to strong starburst merging systems (Wehner & Gallagher 2005). For these reasons, the unique features observed in BIG make us suppose that not only tidal interaction can produce the Hα trails but that probably the mutual influence of tidal and non-gravitational forces (e.g. ram pressure) can explain the physical properties of this group. In order to explain the properties of these trails we need a mechanism able to strip gas from galaxies with little or no influence on the stellar component: a condition respected only by galaxy interaction with the hot intracluster medium. This scenario is also supported by the discovery in the NW part of Abell1367 of two low surface brightness Hα cometary tails, with a total length of 75 kpc, associated with two star forming systems: CGCG 97-073 and 97-079 (Gavazzi et al. 2001a). In fact, the morphology and properties of the tails (which typical size and gas densities are similar to the trails observed in BIG) suggests that galaxies in the NW group are experiencing ram pres- sure due to their high velocity motion through the IGM. The only difference between these two cases is that CGCG97-079 and CGCG97-073 are infalling into the cluster as isolated systems, while galaxies in BIG are infalling within a compact group where gravitational interactions are not negligible.

10.3.3 HI properties HI observations give us additional hints on the properties of this unique group. CGCG97-125, the brightest member of BIG, has a normal hydrogen content: M(HI) 9 1 = 3.9 10 M (Sakai et al. 2002), implying an HI deficiency = -0.21. Its HI column densit×y distribution appears asymmetric, with the highest signal in the western side of the galaxy, as seen in the HI map obtained by Sakai et al. (2002) and reproduced in Fig.10.8, suggesting that this galaxy is strongly perturbed by an external agent. 8 On the contrary CGCG97-114 has an HI mass of only M(HI) = 3.0 10 M (Sakai × et al. 2002) with a resulting HI deficiency of 0.7. This low content is surprising if

1The HI deficiency is defined as the difference, in logarithmic units, between the observed HI mass and the value expected from an isolated galaxy with the same morphological type T and optical obs obs obs linear diameter D: HI DEF = < log MHI (T , D ) > logM (Haynes & Giovanelli 1984) opt − HI 144 10. Galaxy Pre-processing: the blue group infalling in Abell1367

Figure 10.8: HI column density distribution in BIG. Contours are 0.5, 1.0, 2.0, 3.0, 4.0, 5.0, and 6.020cm−2. Adapted from Sakai et al. (2002)

Figure 10.9: HI position-velocity diagram centered on CGCG 97-125. Adapted from Sakai et al. (2002) 10.3. Results 145

−1 compared with the high star formation activity (0.74 M yr , see Table 10.3). At this current SFR, the total HI mass of CGCG 97-114 would be depleted in 2.5 108 yr or in 9.4 108 if we add the total molecular gas mass (4 108) detected by Boselli× × × et al. (1997a) in this galaxy. This suggests that the galaxy is currently experiencing an intense, transient burst of star formation. In addition to the two detected CGCG galaxies, in Fig.10.8 there appears to be extended HI, mostly around CGCG 97-125. The HI extension appears to be a continuation of the Hα structure to the west of CGCG 97-125. This extended structure is typical of galactic merger remnants (Hib- bard & van Gorkom 1996) and suggests that a recent merger has affected this galaxy. The HI distribution around CGCG 97-125 is extended not only in the plane of the sky but in the velocity dimension. The position-velocity diagram presented by Sakai et al. (2002) centered on CGCG 97-125 is shown in Figure 10.9. The velocity distribution shows a regular gradient across the galaxy (the optical major axis of CGCG 97-125 is very close to east-west) ranging from 8090 up to 8490 kms−1, corresponding to a rotation speed of 298 km s−1 when corrected for inclination (Vogt et al. 2004). This value is exceptionally high for a galaxy of the same luminosity, which usually has a rotation speed of 200 km s−1. Thus, both the HI distribution and the HI kinematics yet available suggest∼ that CGCG 97-125 is a quite peculiar object. Addition information concerning the HI properties of this group can be obtained from Arecibo observations. In Fig.10.10 are shown the four spectra obtained for the dif- ferent pointings of BIG. Unfortunately three of the four pointings (97-125, 97-120, 97-114) are surely not independent due to the large overlap in the observed fields. The only, if any, independent observation is represented by the NW field, since it is far away from all bright galaxies and has a relatively small overlap with the field centered on 97-120. Since the side-lobes are located 5 arcmin from the field center, we can exclude a strong contamination of the NW poin∼ting from the bright galaxies in BIG. In addition to the strong HI emission in the velocity range 8000 8500km s−1, associated to the star forming galaxies, a new component in the velocity−range 7500 8000km s−1 is clearly present in all the four spectra. This emission is not associated−with any of the Hα emitting regions since no one of the star forming objects has a recessional velocity below 8000km s−1 (see Table 10.1). This fact strongly emerges in Fig.10.11 ∼ where we compared this spectrum with the mean spectrum obtained from the four pointings. A great fraction of the NW Hα trail described in the previous section lies exactly in the region observed by the BIG-NW pointing, suggesting that the HI emission is probably associated with this feature. This could mean not only that there is neutral hydrogen associated with these structures, but also that their re- cessional velocity is significantly lower than the mean group velocity (assuming that the low velocity component is associated the Hα trails also in the other three point- ings), strongly supporting a ram pressure stripping scenario. Infalling at 1700 km s−1 through the ICM, whose density is ρ 6 10−4atoms cm−3 at their present periph- eral location (A. Finoguenov, private∼comm.),× galaxies in this group will experience 146 10. Galaxy Pre-processing: the blue group infalling in Abell1367

Figure 10.10: The HI spectra obtained for each pointing. 10.3. Results 147

Figure 10.11: Comparison between the combined HI spectrum obtained from the four different Arecibo pointings, and the single pointing on the NW trail. It appears clearly the presence of a low velocity component not associated to the bright galaxies in BIG. 148 10. Galaxy Pre-processing: the blue group infalling in Abell1367

a ram pressure: P = ρv2 3 10−11[dyn cm−2] (10.5) ∼ −2 to an order of magnitude higher. Assuming a stellar surface densities σS 3 10 −2 −3 −2 ∼ × g cm and an interstellar gas surface densities σg 10 g cm , the restoring gravitational force of galaxies is: ∼

F = 2πGσ σ 1.3 10−11[dyn cm−2] S g ∼ Thus the restoring gravitational forces (pressure) at their interiors, are significantly smaller than the ram pressure. In the long run, the increasing ram pressure will fully strip their gaseous material leading to the complete ablation of their interstellar gas, thus suppressing the star formation because of fuel exhaustion. A stripped blob 8 of typical radius R of 2.5 kpc and mass M = 10 M ; might even experience a deceleration, (P F )πR2 a = − = 1.6 10−8[cm s−2] M with a consequent measurable velocity decrease of ∆V = 500 km s−1 in a time as short as 108 yrs, as observed in this case.

10.3.4 The fate of the stripped gas

Different predictions are made in the literature for what happens to the gas once it has been stripped. The large extent of the Hα trails and its associated HI gas indicates that it can survive for some 108 yr or even 1 Gyr. This may suggest that evaporation by the ICM is slow, e.g. because the heat flow is saturated and/or that a tangled magnetic field slows down the heat flow into the trail (Vollmer et al. 2001), as observed in the extended HI plume recently discovered in Virgo by Oosterloo & van Gorkom (2005). In spiral galaxies, if the HI column density is above a few times 1020 cm−2 , star formation almost invariably occurs (Boissier et al. 2003). The mean column densities in the trail is this value, suggesting that it could locally exceed ∼ this threshold. Hence, star formation could occur locally in the trails, provided the processes that regulate star formation for a cloud in the ICM are similar to those for gas clouds in spiral galaxies. The Hα emission in the trails could be signature of star formation in act, representing the most extended example of extragalactic star formation ever observed. However we have no evidence of stellar emission from the trails and the dynamical picture of BIG is consistent with the idea that at least part of the gas that has been stripped is just ionized by ram pressure. In this case the plasma density derived in section 10.3.2 implies an exceedingly short recombination −13 3 −1 time in the ionized trails τr = 1/Neαa 2-7 Myr, where αa = 4.2 10 cm s (Osterbrock 1989). Can their exceedingly∼ short recombination time×of few Myr be 10.3. Results 149

reconciled with an age between some 108 yr and 1.5 Gyr? We need a mechanism to sustain the ionization along the tail and the presence of the cluster IGM comes to help. The clouds stripped from a galaxy infalling onto the IGM might be kept ionized by X-ray bremsstrahlung emission of the IGM. Following Vollmer et al. (2001) and Maloney et al. (1996) the X-ray ionizing photon flux (φi) is:

ln(0.1/0.0136) F φ = X = 8.3 108 F photons cm−2 s−1 i 1.6 10−9 1.5 × X × where FX is the X-ray flux. Assuming a total cluster X-ray luminosity of 4 1043erg cm−2 (Donnelly et al. 1998) the X-ray flux at a projected distance of 125× − − −∼ kpc from the X-ray center (where BIG is observed) is 2.5 10 6erg cm 2s 1 and 3 −2 −1 ∼ × φi results 2.1 10 photons cm s . In equilibrium this gives rise to an ionized column densit∼ y N× = φ /α n . Using n = 10−2cm−3 we obtain N 5 1020cm−2, e i a e e e ∼ × consistent with value measured in the ionized tails. This simple calculation shows that the stripped gas can survive in the hostile IGM, being kept ionized by the X-ray photons.

10.3.5 The metal content

In order to determine the metal content of the observed emission line knots we followed the same procedure described in Chapter 7. The metallicities obtained from the different methods are shown in Tab.10.4. The uncertainty in the abundances is up to 0.2dex. All the star-forming regions in BIG are surprisingly metal-rich. Their metallicity lies in the range 8.5 < 12 + log(O/H) < 8.9. It is well known that irregular and spiral galaxies follow a ”metallicity - luminosity relation” (Skillman et al. 1989). Fig.10.12 shows the ”metallicity - luminosity relation” for galaxies in the Virgo cluster (empty circles, taken from Gavazzi et al. in prep., and obtained using the same methods and calibrations) and for the star-forming systems in BIG (triangles). The two bright galaxies CGCG97-114 and CGCG97-125 have a normal metal content for their luminosity. Conversely the star-forming knots show higher abundances for their intrinsic luminosities. If the faintest systems (K1, K5, K2, 114a and 114b) were isolated independently-evolved dwarf galaxies we would measure a metallicity 0.6-1.2 dex lower than the one observed in this case. Moreover their abundances are consistent with the values measured for tidal dwarf systems showing a metal content of 12 + log(O/H) 8.60 independent from their absolute magnitude (e.g. Duc & Mirabel 1999; Duc et∼ al. 2000). The HII regions DW1, DW2 and DW3 have a high-metal content but consistent, within the calibration uncertainties, with the abundances observed in dwarf galaxies of the same luminosity. However (as discussed in Section10.3.1) the star-forming knots in DW3 are probably not gravitationally 150 10. Galaxy Pre-processing: the blue group infalling in Abell1367

a b c d e Object Tel. R23 R23 NII/OII NII/Hα OIII/NII Mean Stdev K1 ESO 8.53 8.51 8.52 0.02 − − − K1 T NG 8.55 8.52 8.75 8.66 8.55 8.61 0.09 DW3 a T NG 8.62 8.66 8.70 8.66 0.04 − − DW3 d ESO 8.91 8.79 8.92 8.73 8.68 8.81 0.11 97-114b ESO 8.59 8.52 8.30 8.48 8.47 0.13 − 97-114a ESO 8.86 8.71 8.68 8.54 8.72 8.70 0.10 97-114 LOI 9.00 8.83 8.85 8.70 8.92 8.86 0.11 K2 a ESO 8.56 8.51 8.65 8.59 8.56 8.57 0.05 K2 b T NG 8.64 8.39 8.64 8.56 0.14 − − DW2 a T NG 8.53 8.44 8.61 8.67 8.71 8.59 0.11 K5 ESO 8.35 8.66 8.50 0.22 − − − DW1 b MMT 8.56 8.53 8.43 8.48 8.50 0.06 − DW1 c MMT 8.42 8.64 8.58 8.55 0.12 − − DW1 c T NG 8.81 8.70 8.84 8.71 8.66 8.75 0.08 DW1 a ESO 8.43 8.49 8.46 0.04 − − − DW1 a MMT 8.55 8.52 8.67 8.58 8.53 8.57 0.06 97-125 OHP 8.73 8.82 8.65 8.73 0.08 − − 97-125 T NG 8.49 8.77 8.98 8.72 8.74 0.20 − 97-125b ESO 8.45 8.59 8.52 0.10 − − − 97-125b T NG 8.76 8.64 8.75 8.67 8.67 8.69 0.05 K3 T NG 8.89 8.75 8.90 8.84 8.78 8.83 0.07 a: Zaritsky et al. 1994 b: McGaugh 1991 c: Kewley & Dopita 2002 d: Van Zee et al. 1998 e: Dutil & Roy 1999

Table 10.4: Metallicities of the galaxies in the BIG group. 10.3. Results 151

Figure 10.12: The relation between Metallicity and B-band Luminosity (with linear best-fit) for galaxies in nearby clusters (empty circles, adapted from Gavazzi et al. 2004). The triangles mark the mean metallicity obtained for the individual knots of BIG. bound, thus each knot should be considered as a single faint extragalactic HII region with a metallicity 0.8 dex higher than the one obtained from the metallicity- luminosity relation.∼ These results rule out an evolutionary scenario in which the faint HII region discovered in BIG are normal independently evolved dwarf galaxies, reinforcing the scenario of Sakai et al. (2002) who proposed that these systems formed from enriched material stripped by tidal interactions from the two brightest galaxies in BIG.

10.3.6 Dating the starburst. Contrary to the gaseous filaments, current star formation is clearly observed in all compact HII regions, dwarf and giant galaxies composing BIG, suggesting that bursts of star formation are presently taking place in this group. Do we have any hint on when the inset of the star bursting phase took place? The dwarf galaxy DW2, and in particular knots DW2b and DW2c show clear Post-Star-Burst signatures in their spectra, with low residual current star-formation. They have an extremely blue continuum (B-R 0.16), strong Balmer absorption (EW(Hδ) 8A)˚ and [OII] and Hα in emission. In∼ particular a clear age gradient is observable∼passing from DW2a, where star formation is still in place, to DW2c, that shows strong Balmer lines in absorption with some evidences of residual star formation (see Fig.10.18). These 152 10. Galaxy Pre-processing: the blue group infalling in Abell1367

features indicates that the starburst ended already 108 years ago (e.g. Poggianti & Barbaro 1997; Poggianti et al. 1999; Kauffmann et∼al. 2003b).

The star formation history of CGCG97125 The best piece of information for dating the interaction is provided by the brightest group galaxy: CGCG97-125. This galaxy is classified as S0a in the CGCG cata- logue (Zwicky et al. 1961), consistent with its red B-R color index ( 1.34, see also Fig.10.17) and with the shape of the continuum optical spectrum (Fig.10.13).∼ How- ever this system is far from being a normal early type galaxy. The presence of stellar shells around CGCG97-125 (see Fig.10.6) clearly indicates a past interaction/merger event, as also supported by the disturbed rotation curve analyzed in the previous sections. Numerical simulations predict that the stars from a satellite make a system of shells several 108 yr after the end of the merging event and then they last for more than 1 Gyr (Kojima & Noguchi 1997). The spectrum of 97-125 shows a continuum and absorption features typical of elliptical galaxies; however superimposed to it there are strong emission lines (see Fig.10.13) indicating that this galaxy is still experiencing a strong burst of star formation: a kind of rejuvenated early type galaxy. Using the blue line-strength indices to determine the age of the last star forming event (Longhetti et al. 1999) (Hδ/FeI 1.00 , H+K(CaII) 0.91 and ∆4000 1.78) we estimate that ∼ ∼ ∼ the age of the last starburst is 1 Gyr, in agreement with the prediction derived from the presence of the stellar shells.∼ However these models assume an instantaneous burst (SSP), that is clearly not the case of CGCG97125, the obtained age thus repre- sents only a lower limit of the real burst age. In Fig.10.13 we compare the drift-scan integrated spectrum and the nuclear spectrum of CGCG97125 obtained at the OHP and TNG telescope respectively: the integrated spectrum appears considerably bluer than the nuclear one. We can use this difference in order to try to reconstruct the re- cent star formation history of this galaxy. Therefore, assuming that 1)the continuum of the nuclear spectrum is dominated by the old stellar population with no significant contribution from the recent starburst while 2) the integrated one is strongly con- taminated by new stars produced during the burst, we can try to estimate the age of the interaction and the stellar mass produced during the burst. Tidal interactions and merging usually produce a sinking of the gas to the galaxy center triggering a burst of star formation, in contrast with our first assumption. Thus in order to test the validity of our hypothesis we used the SED fitting procedure proposed by Gavazzi et al. (2002a) and developed by Franzetti (2005). We assume a ”a la Sandage” star formation history (SFH):

t t2 SF H(t, τ) = exp( ) (10.6) τ 2 × −2τ 2 and the Bruzual & Charlot (2003) (BC03) population synthesis models. We fitted 10.3. Results 153

Figure 10.13: Comparison between the drift-scan integrated (blue) and nuclear (red) spectrum of CGCG97-125.

Spectrum Z Mass τ t Z log(M/M ) Gyr Gyr Nuclear 0.04 11.01 1.00 13 Starburst 0.04 9.27 0.80 1.4

Table 10.5: Best-fitting parameters for the nuclear and starburst component of CGCG97125. 154 10. Galaxy Pre-processing: the blue group infalling in Abell1367

the nuclear spectrum of CGCG97125, corrected for extinction2 assuming t=13 Gyr, a Salpeter IMF (α = 2.35 from 0.1 to 100 M ; Salpeter 1955) and exploring a parameter grid in metallicity (Z) and τ. Z is let free to vary from 1/50 to 2.5 Z in five steps: 0.0004, 0.004, 0.008, 0.02, and 0.05. τ varies from 0.1 to 25 Gyr in 45 approximately logarithmic steps. The best-fitting parameters obtained using the BC03 models are summarized in Table 10.5. The best value of τ is consistent with the one (τ 3.1Gyr) obtained by Gavazzi et al. (2002a) fitting a template of S0 galaxies. This ≤result val- idates our assumption that the continuum of the nuclear spectrum is dominated by an old stellar population of the same age expected for an unperturbed S0. By nor- malizing the obtained model to the observed H-band magnitude and subtracting it to the integrated UV to near-IR SED of CGCG97125, we have the possibility to esti- mate the starburst contribution to the galaxy emission, the burst age and the stellar mass produced during the star formation. The best-fitting parameters obtained for the starburst SED are summarized in Table 10.5. The burst age results 1.4 Gyr 9 ∼ and the stellar mass produced during the burst is 2 10 M , consistent with the values previously obtained from independent estimates∼ ×(i.e. dynamical models). The resulting best fitting SED for CGCG97-125 is shown in Fig.10.14 (black model). The model well reproduces the observations from the far-ultraviolet to the near infrared, with the exception of the near-ultraviolet. This disagreement does not depends on the model assumption but on the attenuation law used to correct for internal dust attenuation. In fact, as shown in Appendix II, we assume a Milky Way attenuation law (thus with a bump at 2175 A)˚ that seems not to be valid for normal star form- ing galaxies (see Chap. 7),∼producing an overestimate of the real galaxy emission in near-ultraviolet. However this does not influence our results as shown in Fig.10.14. We can thus conclude that the star burst in CGCG97-125 initiated 1-1.5 Gyr ago, 9 ∼ probably produced by a minor merging of a 2 10 M satellite, and is still taking ∼ × place. Our result points out that a minor merging able to disturb the morphology and the dynamics of a giant galaxy as CGCG97-125, seems not able to strongly mod- ify the mean age of the stellar population, producing only a small fraction ( 2%) of ∼ new stars (see also Fig.10.15). As in the case of NGC4438 (see Chapter 8), this result might probably be representative only for a minor merging into, today gas poor, early type galaxies.

2We corrected all the spectrophotometric data using the ultraviolet spectral slope β as suggested in Chapter 7 for FUV data and the method described in Appendix A for NUV and optical obser- vations. We assume that the nuclear and integrated spectrum are affected by the same amount of dust extinction, as supported by the similar value for the Hα/Hβ ratio obtained in the two spectra (see Table 10.2). 10.3. Results 155

Figure 10.14: The SED of CGCG97-125, corrected for internal extinction. Nuclear and drift-scan integrated spectra are shown in green. Black circles indicate photomet- ric observations and their relative uncertainties. Best fitting models for the nuclear spectrum (red) and for the starburst component (blue) are given. The resulting best fitting SED for CGCG97-125 is presented in black.

Figure 10.15: The star formation history of CGCG97-125 as obtained from the SED fitting procedure. 156 10. Galaxy Pre-processing: the blue group infalling in Abell1367

Figure 10.16: The 2D high resolution spectrum (left) and the optical rotation curve (right) of CGCG97-120

10.3.7 CGCG97-120: simply a foreground galaxy, or an high velocity intruder?

In the previous sections I never mentioned the brightest galaxy present in the BIG region: the spiral galaxy CGCG97-120. This massive system has a recessional veloc- ity of 5635 km s−1, thus blueshifted with respect to A1367 by approximately 800 −∼ km s 1. Observations of the neutral hydrogen line show that CGCG97-120 has lost approximately 90% of its original hydrogen content (HI deficiency = 0.9), suggesting that the galaxy has crossed the cluster core and that the ram pressure exerted by the dense intergalactic medium might have caused its hydrogen deficiency. The great velocity difference between this galaxy and BIG ( 2500km s−1) seems to rule out any ∼ association between the two systems, as argued by Sakai et al. (2002) and Gavazzi et al. (2003b). However the deep Hα images obtained at the TNG telescope repropose the question: one of the Hα trail traces in fact a perfect loop around CGCG97-120. Only a blind chance? As shown in Fig.10.6 and Fig.10.16 the galaxy morphology and kinematic are completely unperturbed, showing no signs of interaction. Moreover the scattering angle of an interaction between a satellite galaxy and CGCG97-120 would be of only 4-10 degrees, assuming a classical scattering model and an impact parameter of 10∼kpc: too small to produce the observed loop. Thus for the moment ∼ we have to suppose that the association between CGCG97-120 and the Hα trails is only a blind chance. 10.4. Discussion 157

10.4 Discussion

10.4.1 The evolutionary history of the Blue Infalling Group

The amount of information provided by the multiwavelength observations presented in this paper allow us to reconstruct the evolutionary history of BIG, during the last 1-2 Gyr. At the beginning of the story BIG was a normal compact group of galaxies with a typical dispersion velocity of 150-200 km s−1, composed of at least ∼ three galaxies: a massive evolved early type spiral (CGCG97-125), a massive late type spiral (CGCG97-114) and a gas rich dwarf galaxy (the satellite that has feed 9 CGCG97-125) with a stellar mass 10 M . Lying in the outskirts of Abell1367 it has been attracted by the cluster poten∼ tial starting its infall into the cluster core at a mean velocity of 1700 km s−1. During their journey, all galaxies are perturbed by ∼ gravitational interaction with members, as observed in all compact groups. Stars and gas are stripped, forming tidal tails, bridges (as K2), extragalactic compact HII regions (as K5 and K1) and tidal dwarfs (as DW1, DW2 and DW3). Tidal interactions lowered the restoring force by loosening the potential well of all galaxies in the group making easier stripping gas from the infalling galaxies by ram pressure and producing the unique Hα trails observed in BIG. In particular the gas rich satellite is partially dismantled by the combined action of tidal forces and ram pressure, and finally merged into CGCG97-125 producing stellar shells and a burst of star formation. The combination of gravitational forces and ram pressure is not only consistent with the evidence that BIG is a compact group that is infalling at 1700 km s−1 into ∼ the core of Abell1367, but is also necessary to try to explain all the aspects that make BIG so unique among other known interacting systems and merger remnants: i.e. the unexpectedly high star formation observed in this group, the presence of extended Hα trails and its associated neutral hydrogen, the lack of large-scale tidal tails and, as pointed out by Gavazzi et al. (2003b), the colors of dwarf objects DW1, DW2 and DW3 that are significantly bluer than tidal dwarfs observed in interacting systems (Weilbacher et al. 2000). The IGM compression is in fact able to trigger some star formation in the gas clouds contained within tidal structures (Bekki & Couch 2003), while ram pressure may push some of these clouds free of their parent galaxies, explaining the absence of tidal features and the extremely blue colors of the dwarf objects in BIG. Recently Mayer et al. (2005) have shown that gravitational tides can aid ram pressure stripping by diminishing the overall galaxy potential. The gas stripped along tails fragments into dense clouds and sheet-like structures pressure confined by the ambient medium with the approximately the same column density observed in our case. However their simulations are focused on the evolution of dwarfs (V 40 km s−1) systems orbiting around a Milk Way like galaxy, and it is not clear rot ∼ what would be the effects of the same mechanisms on a massive galaxy infalling into a cluster. 158 10. Galaxy Pre-processing: the blue group infalling in Abell1367

10.4.2 The contribution of preprocessing to cluster galaxies evolution.

Galaxy clusters formed not by accreting individual galaxies randomly from the field, but rather through the infall of small groups, falling in along large scale filaments; thus this group represents an unique laboratory reproducing the physical condition expected in a cluster still in formation. What can we learn about galaxy cluster evo- lution studying BIG? First of all, we are witnessing the first clear example of a well formed S0 galaxy infalling into the core of a cluster of galaxies. This observational evidence suggests that S0 galaxies can form outside clusters and subsequently fall into them: groups environment is in fact considered as the best place where gravi- tational interactions should operate efficiently and transform a normal spiral into an S0. Moreover gravitational interactions among group’s members are still in act, and CGCG97-125 has recently ( 1.5 Gyr) experienced a minor merging event. The burst of star formation, however, ∼is not able to strongly affect its global optical properties, since the mass of new stars produced is only 2% of the whole galaxy mass, con- sistent with the recent results obtained by Boselli∼ et al. (2005a) in the Virgo cluster. This suggests that the mechanism responsible of the transformation of CGCG97-125 into an S0 is older than 2 Gyr, corresponding to a redshift z 0.2. At its current −1 ≥ 9 SFR (1.49 M yr ) the total HI mass of CGCG 97-125 would be depleted in 2.6 10 yr, implying a total burst duration of 4 Gyr: consistent with the typical time-scale× ∼ of the Butcher-Oemler effect (Butcher & Oemler 1978, 1984). Tidal interactions within group members are not only able to produce morpholog- ical transformation in galaxies, but also to create new systems formed by gas and stars stripped from group’s members. This is the case of the extremely high number of metal rich star forming dwarfs/extragalactic HII region detected in the infalling group. What will be the future of these stripped systems? It is improbable that all the stripped clusters will infall into the main galaxies, rebuilding the gaseous disk as observed in field mergers: in fact cluster tides and ram pressure stripping act mutu- ally to strip off the material and to inhibit the disk resettling process. If they are dynamically bounded, they could be the progenitors of dwarf cluster galaxies as de- scribed in models by Kroupa (1998) and Duc et al. (2004). Simulations predict that young compact massive star clusters formed during the merger of gaseous disk galax- 7 9 ies coalesce within a few 100 Myr forming objects with masses of order 10 10 M , − as observed in this group, with negligible dark-matter content. However, till today major mergers were supposed not to have produced a significant fraction of the dwarf population, since each merger are expected to spawn only one or two tidal dwarf galaxies. Thus the discovery of a great number of extragalactic star forming knots in BIG ( 10 for one merging event) seems indicate that a tidal formation scenario ∼ for part of the dwarf population in cluster could be reasonable, especially at higher redshift, where groups like BIG are expected to infall at higher rate into the core of 10.4. Discussion 159

young clusters. Being produced by a single starburst, these systems might also evolve into dwarf ellipticals, typical of rich clusters. Otherwise, if they will disperse they stars and gas into the cluster their will simply in- crease the fraction of unbound stars, contributing to the Abell1367 intracluster light, supporting the idea that preprocessing could have had a strong contribution in the amount and distribution of intracluster light (Mihos 2004b). A strong contribution to the intracluster light in Abell1367 would also be provided by the Hα trails if some residual star formation is taking place. In this case these features would represent the most extended example of extragalactic star formation ever observed in the Universe. Surely the Hα trails are strongly contributing to the ICM enrichment, suggesting that a considerable amount of the cluster enrichment might derive from these late-type intruders, as opposed to winds from elliptical galaxies, commonly accepted as the major sources of pollution (e.g. Madau et al. 2001; Mori et al. 2002). This idea is strongly supported by the presence in the NW part of Abell1367 of other two galaxies with Hα trails (Gavazzi et al. 2001a), pointing out that this may not be a rare phe- nomenon in young clusters. Moreover in the last years an increasing number of X-ray (Hayakawa et al. 2004) and optical (Gavazzi et al. 2001a; Oosterloo & van Gorkom 2005) observations has shown that ram pressure stripping could have an important role on IGM enrichment and recent combined N-body and hydrodynamic simulations have pointed out that more of the 10% of the intracluster medium originated from gas stripped by ram pressure (Domainko et al. 2005). Thus, combined X-ray and optical studies of infalling groups should help us to shed more light on the effect of preprocessing not only on the evolution of cluster galaxies, but also on process of IGM enrichment, an issue which remains unsettled (Tornatore et al. 2004). The evolution- ary scenario here presented points out the great importance of groups like BIG not only for galaxy evolution but also for the evolution of clusters itself. Since the infall rate of these groups is considerably higher at high redshift, this analysis points out the strong contribution that small compact groups have probably had in shaping the properties of both galaxies and clusters of galaxies. Thus BIG represents a Rosetta Stone group, giving us the chance to shed light, with the great details possible only in the local Universe, on physical processes typically expected in young, far away, clusters. 160 10. Galaxy Pre-processing: the blue group infalling in Abell1367

Figure 10.17: B-R color map of BIG (Blue = B; Red = R). 10.4. Discussion 161

Figure 10.18: The observed smoothed (step 3) one dimensional spectra. The object identification and telescope are labeled on each panel. 162 10. Galaxy Pre-processing: the blue group infalling in Abell1367

Figure 10.18: Continue. 10.4. Discussion 163

Figure 10.18: Continue.

Chapter 11

Discussion & Conclusions

11.1 Discussion

In the Introduction to this thesis I have argued that try to recover galaxy evolution during the last 13 Gyr only from observations of today’s Universe represents a real, but fundamental, challenge. The Universe we inhabit is old, and most of the fun is over. In addition (and this is the worst part of the story) the Universe dramatically evolved itself, continuously altering the physical conditions of the environments pop- ulated by galaxies. However, as shown in this work, we can still achieve important pieces of information from the study of local galaxies and, combining this information with that obtained at higher redshift, we can try to paint a picture of our knowledge about the evolution of galaxies in clusters; exactly the effort that I’m going to attemp in this conclusion. What evolutionary scenario emerges from this work? The first estimates of the UV cluster luminosity functions from FOCA, FAUST and GALEX observations, here presented, point out that at these wavelengths the cluster LF is considerably steeper than the field one. The steepening of the UV LF from low to high density environment is due to the increasing contribution of early-type, non star forming galaxies, passing from the field to the cluster core. This represents the first evidence of a morphology/star formation - density relation at ultraviolet wave- lengths and demonstrates that we cannot blindly consider UV selected galaxies as star-forming systems, especially at low UV luminosities and in high density environ- ments. However this also point out the strong potential of ultraviolet observation in studying all cluster galaxies: not only star-forming systems which UV emission traces the presence of newly born stars, but also early type galaxies in which such emission must be ascribed to low mass old post asymptotic giant branch stars. So let me summarize what I have learned about the evolution of these different morphological types.

165 166 11. Discussion & Conclusions

The evolution of elliptical galaxies For the first time, in this work the UV prop- erties of early-type galaxies have been studied down to MB -15 mag. The newest result addresses the question raised by O’Connell (1999) concerning≤ the dependence of the UV properties on galaxy morphology. We have shown that a dichotomy ex- ists between giant and dwarf ellipticals and, to a lesser extent, between ellipticals and lenticulars. The blueing of the UV color index with luminosity, metallicity, and velocity dispersion indicates that the UV upturn is more important in massive, metal- rich systems. Since the UV upturn originates from a minority population of old hot helium-burning horizontal-branch (HB) stars, which emission becoms detectable after at least 10 Gyr (e.g. O’Connell 1999; Brown et al. 2000; Greggio & Renzini 1990; Tantalo et al. 1996), the relation found for giant ellipticals and its small dispersion suggest that clusters ellipticals represent an old, homogeneous population. This is also consistent with the dynamical analysis of Abell 1367 where we found evidence that elliptical galaxies have a Gaussian velocity distribution with a smaller velocity dispersions than the whole cluster sample, representing the virialized, old, cluster pop- ulation. This picture is supported by both higher redshift observations and N-body simulations. The population of elliptical galaxies in clusters show little evolution in their colors and no structural evolution since at least redshift of z 1 (Treu et al. ∼ 2005; Smith et al. 2005). The attempt to reproduce this observational evidence with N-body simulations (Springel & Hernquist 2003) results in the invalidation of the paradigm of elliptical formation by mergers of spiral galaxies. At the time the cluster ellipticals were formed in rich clusters, there were simply few if any spiral to merge. It appears clear from simulations, as well as observations of the high-z Universe, that large spiral galaxies as we know them today were very rare at z>2 (Driver et al. 1998; Dickinson et al. 2003; Trujillo et al. 2004; Conselice et al. 2005). It is specially true in dense environment where galaxies had too little time to form large disk from the accretion of high-angular momentum material. This is also supported by HST Deep fields: beyond z=1 the number density of large, well-formed spirals begins to rapidly diminish in favor of a smaller, chaotically arranged systems (Labb´e et al. 2003). In addition the recent Millennium simulation (Springel et al. 2005) has shown that the giant ellipticals observed in nearby clusters today, were already formed and massive at very high redshifts (z 16) and harboured in the center of regions where the first structures developed:∼the progenitors of rich clusters. We can thus conclude that both observations (at all redshifts) and simulations are consistent with the idea that clusters giant ellipticals are an old, homogeneous population showing no or little evolution at least in the past 8 Gyr. This is not the case for the cluster population of dwarf elliptical galaxies. The op- posite behavior in the UV color magnitude relations (reddening of the UV color index with luminosity) of dwarfs with respect to giant ellipticals, similar to that observed for spirals, indicates that the UV spectra of low luminosity objects are shaped by the contribution of young stars, thus presenting a very different star formation history. 11.1. Discussion 167

This implies that the stellar population of dwarfs has been formed in discrete and rel- atively recent episodes, as observed in other nearby objects (Grebel 2000). However this result is not sufficient to discriminate between different theoretical models for dE formation: primordial objects that lost their gas in a supernova-driven galactic wind (Yoshii & Arimoto 1987; Nagashima & Yoshii 2004), dwarfs irregular infalling into cluster and transformed by ram pressure (van Zee et al. 2004) and/or harassment (Moore et al. 1998), or tidal dwarfs (Kroupa 1998; Duc et al. 2004). The higher fre- quency of dwarf ellipticals in high density environments supports the idea that they are objects transformed by the harsh cluster environments. However the presence of observational evidence supporting at least the first two scenarios, and the very high dispersions observed in the UV color magnitudes and in structural and kinematic parameters (de Rijcke et al. 2005) seem to suggest that dwarf ellipticals are most likely a mixed population with primordial, and more recently transformed objects co-existing in the present day Universe. Moreover the reddening of the UV color index with luminosity is new evidence that mass drives the star formation history in hot systems (Trager et al. 2000; Gavazzi et al. 2002a; Caldwell et al. 2003; Poggianti 2004b) as in rotating ones (Gavazzi et al. 1996; Boselli et al. 2001, see also below). This phenomenon, today refereed as downsizing effect, is observed in both cluster and field and at least till z 0.8 indicating the presence of an ”anti-hierarchical” history for star formation in galaxies.∼ The presence of a downsizing effects in all galaxies, independent from their morphological type, represents today the major challenge for CDM models.

The evolution of lenticular galaxies Unlike the rather passive evolution ob- served in cluster ellipticals, much stronger evolution seems present in the population of cluster S0s. The dispersion observed in the UV color magnitude relation, consider- able higher than ellipticals’, bears witness to recent, minor episodes of star formation combined with an old stellar population, as determined also from kinematic and spectroscopic observations (Dressler & Sandage 1983; Neistein et al. 1999; Hinz et al. 2003). This result is consistent with recent studies of stellar population in early type galaxies which found significant differences between the ages of the stellar popula- tions of ellipticals and of the S0 galaxies, supporting the scenario of spirals evolving into S0s (Kuntschner & Davies 1998; van Dokkum et al. 1998; Terlevich et al. 1999; Poggianti et al. 2001; Smail et al. 2001). All these results are supported by the fact that the fraction of S0s in rich clusters has increased significantly since a redshift of z 1 (Smith et al. 2005), with a corresponding decrease of spiral fraction (Dressler et∼al. 1997). What is the mechanism responsible for the transformation of a gas rich spiral into a lenticular galaxy? The toy model presented in Chapter 9 has shown that ram pres- sure alone cannot account for all of the S0 population observed in nearby clusters. 168 11. Discussion & Conclusions

Galaxy-cluster IGM interaction is able to remove most of the gas reservoir inducing important structural modifications in the disk properties. We expect that cluster spi- rals have (at least at short wavelengths) more truncated disk profiles, inverting the outer color gradient with respect to similar but unperturbed objects, and then pro- ducing anemic spirals, similar to disk dominated S0s. The surface brightness of the disk, however, mildly decreases in Hα and in the UV bands while remaining mostly constant at longer wavelengths even 5 Gyr after the interaction. Excluding the inter- action with the ICM, the only mechanism able to produce a structural modification in spiral galaxies are gravitational interaction. Tidal interactions between galaxies affect both stars and gas. Stars respond by forming arms and bars, while the gas flows directly toward the central regions within about 108 yr after the initial collision. The sinking of the gas towards the galaxy center could trigger a burst of star formation and, on longer timescales, a truncation of the stellar disk (Iono et al. 2004), thus al- tering galaxy morphology. On the other hand we have to exclude harassment since its influence is largely limited to low luminosity galaxies, while in bright spirals its effects are much more limited (Mihos 2004a; Moore et al. 1996). Thus merger-driven S0 for- mation mechanisms appear not to work inside the cluster potential, since low velocity interactions are extremely rare. On the other hand, these processes should operate efficiently in the group environment, where the encounter velocities are smaller and cluster tides and the hot ICM are not important. The group environment can create S0s and feed them into the accreting cluster. Although the accretion into the cluster core is expected to happen at higher redshift, we have shown in Chapter 10 a clear example of this phenomenon observed in the local Universe. The starbursting group infalling into the core of Abell1367, represents probably the best example of galaxy preprocessing ever observed. The brightest member of this group is an S0 galaxy in strong gravitational interaction with the other group members. It is thus likely that many of these S0s were processed via mergers in the group environment before being incorporated into clusters; especially in the past where the groups’ infall rate was con- siderably higher than today. Moreover the discovery of other groups of S0 galaxies in strong interaction such as the one in the outskirts of the Ursa Major cluster strongly support this formation scenario (van Gorkom 2004). This formation scenario is sup- ported by the observational evidence that the bulk of S0 population in clusters was formed between z 0.2 and z 1, when the rate of infall of small group was the highest experienced by clusters∼ of galaxies∼ (Mihos 2004a). Finally, S0s are a heterogeneous class, from the bulge dominated to the disky S0s, and it should not be surprising that a single mechanism cannot fully account for the range of S0s types (Hinz et al. 2001, 2003): if ram pressure is able to produce disk dominated S0s (objects similar to the anemic galaxies of Van der Berg), tidal interaction (and thus preprocessing) are required to account for bulge dominated S0s. 11.1. Discussion 169

The evolution of the star formation activity in cluster spiral galaxies We can conclude that the bulk of the bulge dominated S0 cluster population was formed at higher redshifts, and in environments where tidal interactions were more proba- ble. However the morphology density relation, and in particular the star formation - density relation, as we observe it today is not fully established at high redshift, be- cause we observe how it evolves in clusters, with star forming late type spirals being transformed into anemic galaxies with quenched star formation of the same morpho- logical type. In order to have an idea of this phenomenon, look at Fig.11.1. It shows the distribution of median value of EW(Hα) as a function of the morphological type for Virgo galaxies with normal gas content (HI-def< 0.4; filled circles) and deficient galaxies (HI-def> 0.4; empty circles). The figure emphasize that, within each Hub- ble class, galaxies with normal HI content have EW(Hα) systematically higher by a factor two then their deficient counterpart. What is the major mechanism (if any) responsible of this reduction is still unknown: observational results are not always consistent each other, and their interpretation results not straightforward at all. Let me start from the results presented in Chapter 3 and 4. I compared the UV lu- minosity function of nearby clusters and local field showing that the shape of the LF for star-forming galaxies does not change significantly in different environments. The easiest interpretation of this result is that the dwarf to giant star forming galaxies ratio is independent from the environment; the only thing that changes is the absolute fraction of star forming galaxies (i.e. the normalization of the luminosity function). This is a very simple picture but consistent with the recent work of Balogh et al. (2004) who have shown that the distribution of Hα equivalent widths in star form- ing galaxies does not depend strongly on the local density, while the fraction of star forming galaxies is a steep function of the local density, in all environments. Under- standing the origin of these observed trends is one of the most interesting questions to be answered, since it probably include the key to shed light on the environmental influence on today’s galaxy evolution. First of all, these results seem to suggest that the mechanism that affects the star formation when a galaxy enters a dense environment, must work on a short time scale ( 107 108 yr), and must affect bright and faint galaxies in the same way, in order ≤ − to preserve the shape of the luminosity function and of the EW(Hα) distribution. Al- though this excludes strangulation from playing a major role in galaxy evolution, due to its high time scale ( 1 Gyr), it is not clear which mechanism dominates between ≥ ram pressure and tidal interactions. A lot of research groups (i.e. Dressler 2004; Balogh et al. 2004; Goto 2005) have proposed tidal interactions as the major mech- anism responsible of galaxy transformation. This idea is mainly supported by the fact that the decreasing of EW(Hα) with local density has approximately the same shape in all environments, from cluster to groups (see Fig. 11.2). Since in groups ram pressure stripping is supposed to be absent (even if Fujita 2004 has shown that ram pressure could be also important in groups), the only mechanisms available to quench 170 11. Discussion & Conclusions

Figure 11.1: The distribution of the individual HαE.W. measurements in the Virgo cluster along the Hubble sequence (small dots) and of the median EW(Hα) in bins of Hubble type. Error bars are drawn at the 25th and 75th percentile of the distribu- tion.Filled symbols represent HI-def< 0.4 (unperturbed) objects and open symbols HI-def> 0.4 (HI deficient) galaxies. 11.1. Discussion 171

10 1 10 1

30 30

20 20

10 10

0 0 0.1 1 10 0.1 1 10

Figure 11.2: The star formation rate as a function of density, comparing groups of galaxies with clusters. The upper and lower horizontal dashed lines show the 75% percentile and the median of the equivalent widths. The hashed region shows the relation for the complete sample, while the solid line shows the relation for systems with 500 km s−1 < σ < 1000 km s−1 (left) and σ < 500 km s−1 (right). The dependence on local density is identical irrespective of the velocity dispersion of the whole system. Figure taken from Bower & Balogh (2004). star formation are low velocity encounters: we return once again to the idea that the main mechanism responsible of the evolution of spiral galaxies is preprocessing in lit- tle groups. It is indisputable that this simple interpretation rules out the galaxy-ICM interaction; however as shown in this work, in local clusters the role of ram pressure seems significant. First of all in Chapter 5, we have shown that star forming galaxies in Abell1367 have an higher velocity dispersion than the quiescent population. This is observed also in other clusters (Sodre et al. 1989; Stein 1997; Biviano et al. 1997; Adami et al. 1998), and reflects the fact that spirals have an higher velocity disper- sion than ellipticals, and a velocity distribution hardly Gaussian (Boselli & Gavazzi 2006). By itself it provides evidence for infall of star forming galaxies into clusters. If a consistent fraction of star forming galaxies are still today passing from low to high density environments, their star formation activity will be soon quenched in order to reproduce the observed trends in luminosity function and EW(Hα) distributions. Since in today’s clusters tidal interaction are less probable, this result supports a ram pressure scenario. In addition van Gorkom (2004) has shown that the velocity dispersion of gas rich galaxies is far from Gaussian contrary to the one of HI deficient galaxies, suggesting that gas rich galaxies that enter the cluster center are likely to 172 11. Discussion & Conclusions

Figure 11.3: The ratio of the isophotal Hα and r0 radii as a function of the HI deficiency for galaxies in the Virgo cluster. be serious affected by interaction with the ICM. This is only the first, and if possible less strong evidence, of the role played by galaxy-intracluster medium interaction. In Chapter 9, we have argued that the population of anemic spirals in clusters, with truncated star forming disks, is produced by ram pressure stripping and that the time scale of the interaction is short ( 100 Myr). In addition a growing number of spiral ∼ galaxies are found with unusual morphology in HI, Hα and radio continuum, such as the head tail galaxies CGCG97-073 and CGCG97-079 in Abell1367 (see Chapter 5, Gavazzi et al. 1995, 2001a), NGC4522 (Vollmer et al. 2004b), NGC4388 (Yoshida et al. 2004), NGC4569 (see Chapter 9) in Virgo and CGCG160-055 and CGCG160- 095 (Bravo-Alfaro et al. 2000, 2001) in Coma. These are prime candidates for ongoing ram pressure stripping. In order to try to determine how important are ICM-ISM interactions for galaxy evo- lution, as a part of the undergraduate thesis of I.Arosio (Arosio 2005), we analyzed the morphological distribution of galaxies in Virgo and Coma showing that the ratio of the Hα to optical radius correlates with the HI deficiency (see Fig.11.3). This result is consistent with the increase of the fraction of galaxies with truncated star forming toward the center of the Virgo cluster, observed by Koopmann & Kenney (2004a) and 11.1. Discussion 173

Figure 11.4: The clustercentric radial distribution of the individual EW(Hα) mea- surements in the Virgo cluster. High and low (B-band) luminosity galaxies are given with open and filled dots respectively. Median in bins of 0.5 R/RV ir are given. Error bars mark the 25th and 75th percentile of the distribution. with the prediction of the ram pressure model presented in Chapter 9. In addition the strong correlation between EW(Hα) and HI deficiency observed in nearby clusters (Gavazzi et al. 2002c) completes the ram pressure supporting scenario. To summa- rize, in nearby clusters we observe galaxies that experience ram pressure stripping: the dominant effect on cluster disk galaxies is a reduction of the star formation rate, which goes hand in hand with the HI deficiency, and for most of the galaxies this seems due to ram pressure. How can we conciliate these results with the universal shape of the EW(Hα) vs den- sity relation presented by Balogh et al. (2004)? This apparent contradiction awaits an explanation. What I think emerges from this work is that the truth lies probably in the middle: it is indisputable that galaxy preprocessing (and in particular tidal interaction) has played a significant role, especially in shaping the properties of bright (giant) galaxies at higher redshift; but at the same time it is unquestionable that we observe today normal ”field-like” galaxies, not affected by any preprocessing, infalling alone (or in 174 11. Discussion & Conclusions

very loose groups), for the first time into clusters and on which ram pressure’s ef- fects are clearly evident, as shown in the dynamical study of Abell1367. In part, it might thus be correct to affirm that while tidal interactions have dominated in the past and have probably shaped the morphology-density relation for giant galaxies, ram pressure dominates in today clusters and is surely affecting the star formation history of galaxies but with less influence on their morphology. What is still far from being understood is the downsizing effect (i.e. the correlation of the mean age of stellar populations with the mass). In fact this effect is clearly present in clusters (Gavazzi et al. 2002a; Kodama et al. 2004; Poggianti et al. 2004) where, on the con- trary, environmental effects (whichever you prefer) are expected to be more efficient in quenching star formation in dwarfs than in giant galaxies, since gas and stars are less bounded to the galaxy. This effect is clearly evident in Fig.11.4 where is shown the dependence of the EW(Hα) on the cluster-centric distance in the Virgo cluster. While the decline in the star formation rate is clear for giant galaxies (passing from EW(Hα) 35A˚ at 2 virial radii to EW(Hα) 6A˚ in the cluster center), we do not identify an∼y significant trend for dwarf galaxies.∼ This result could be explained if we assume that a significant replenishment of dwarf galaxies is occurring into rich clusters at the present cosmological epoch. An high infall of dwarf systems is also supported by the fact that the velocity dispersion of dwarf star forming galaxies is considerably higher than the one of high luminosity spiral systems (Adami et al. 1998). Thus understanding how and at which rate galaxies infall and have infalled into cluster represents another important key to shade light on the evolution of star formation activity with clusters. If confirmed, a high infall rate for today’s dwarf galaxies will represent a new interesting challenge for hierarchical models of galaxy evolution, in their unfinished attempt of reproduce the Universe we inhabit. 11.2. Conclusions 175

11.2 Conclusions

In this thesis I have investigated the environmental effects on galaxy evolution in nearby clusters using a multiwavelength dataset. In particular this analysis has been focused on the properties of three different local clusters: Abell1367, Virgo and Coma. These three clusters are among the best studied in the local Universe and, due to the variety of their environmental conditions (e.g. spiral fraction, X-ray luminosity, evolutionary stage), they represent the most suitable ”laboratory” for comparative studies. By combining for the first time GALEX UV observations with optical, near and far infrared data we derived extensive observational evidence of cluster galaxy evolution.

I determined the first far-UV and near-UV luminosity functions for nearby • clusters finding that in clusters the faint end slope is steeper than in the field. This difference is entirely due to the contribution at low UV luminosities of non-star-forming, massive early-type galaxies that are significantly overdense in clusters; while the luminosity function of cluster star-forming galaxies is consistent with the field one. This indicates that, whatever mechanism affects the star formation activity in late-type cluster galaxies, it influences similarly and with a short time scale the giant and the dwarf components.

I investigated the dynamical state of Abell1367 showing that this cluster is still • a young cluster forming at the intersection of large scale filaments. At least two subgroups are currently infalling into the main cluster. They show a higher fraction of star forming galaxies than the cluster core, as expected during the early phase of merging events, confirming that the building up of large scale structures can strongly affects the evolutionary history of galaxies.

I studied for the first time the UV properties of a volume-limited sample of • early-type galaxies showing the presence of a clear dichotomy in the FUV-optical color magnitude relations between giant and dwarf ellipticals. For elliptical and lenticular galaxies, the (FUV-NUV) color becomes bluer with increasing lumi- nosity and with increasing reddening of the optical or near-IR color indices. For the dwarfs, the opposite trend is observed. These results are consistent with the idea that the UV emission is dominated by hot, evolved stars in giant systems, while in dwarf ellipticals residual star formation activity is more common.

While investigating the star formation history of galaxies in nearby clusters • using ultraviolet observations, it has been mandatory to study UV dust atten- uation properties of nearby galaxies in order to look for new recipe’s in oder to correct GALEX data. I confirmed that normal galaxies follow a LT IR/LF UV β relation offset from the one observed for starburst galaxies. The dispersion−of 176 11. Discussion & Conclusions

this relation is found to weakly correlate with the galaxy star formation history. I studied the correlation of dust attenuation with other global properties, such as the metallicity, dynamical mass, ionized gas attenuation, Hα emission and mass surface density providing some empirical relations from which the total infrared to far ultraviolet ratio (LT IR/LF UV ) can be estimated when far infrared data are absent. This result represents only the tip of the iceberg of a study of dust properties in normal galaxies. Only comparing data with models we will be able to properly correct data for dust extinction and thus to estimate the star formation rate in galaxies. Finally I studied in great details the star formation history of three different systems considered as the prototypes of the three main environmental mechanisms able to perturb galaxy evolution, namely: high velocity interactions, ram pressure stripping and galaxy preprocessing. We showed that in today’s cluster galaxies high-velocity tidal encounters be- • tween two galaxies of similar mass are able to perturb the stellar distribution and thus produce important tidal tails, but are not sufficient to significantly increase the star formation activity of cluster galaxies. Moreover we demonstrated that ram pressure stripping alone is not able to • transform a spiral galaxy into an S0, reproducing the structural properties of present-day lenticulars. Strong transformations in both morphology and star formation activity can be • produced by the mutual effects of low velocity encounters and ram pressure stripping in small groups infalling into the cluster core (preprocessing), as ob- served in Abell1367. Studying this unique example of preprocessing in the local Universe we showed that infalling groups could have a strong influence not only on galaxy evolution but also on the evolution of cluster galaxies, significantly contributing to the enrichment of the intracluster medium and to the intraclus- ter light. Considering all these observational results I conclude that Giant ellipticals are an old, homogeneous population showing no or little evolu- • tion at least in the past 8 Gyr unlike dwarf ellipticals which still contains young stellar populations. The importance of different environmental mechanism is directly linked with • the age of the Universe. Tidal interactions and prepocessing have probably dominated in the past Uni- • verse and has shaped part of the morphology-density relation during the cluster accretion of small groups. 11.2. Conclusions 177

Ram pressure dominates in today clusters and is surely affecting the star for- • mation history of galaxies with less influence on their morphology.

The heterogeneous class of S0s galaxies, from bulge dominated to the disky S0s, • is not the result of a single transformation mechanism: if ram pressure is able to produce disk dominated S0s, tidal interaction (and thus preprocessing) are required to account for bulge dominated S0s.

Different observational clues confirm the presence of a correlation between the • mean age of stellar populations and the mass of their parent galaxies (downsizing effect). In the framework of the hierarchical model of galaxy formation, the origin of the downsizing effect remains unsolved and represents one of the main challenges for models of galaxy evolution.

Appendix A

The extinction correction

Here we present the method used to correct for dust attenuation multiwavelength observations used in this work, and proposed by Boselli et al. (2003a) As discussed in Chapter 7 the observed stellar radiation of galaxies, from UV to near-IR wavelengths, is subject to internal extinction (absorption plus scattering) by the interstellar dust. Estimating the dust extinction at different λ in external galaxies is very difficult (it has been done only for the Magellanic clouds). This difficulty is mainly due to two reasons: a) the extinction strongly depends on the relative geometry of the emitting stars and of the absorbing dust within the disc of galaxies. The young stellar population are mostly located along the disc in a thin layer, while the old populations forms a thicker layer. This point is further complicated by the fact that different dust components (very small grains, big grains etc.), which have different opacities to the UV, visible or near-IR light, have themselves different geometrical distributions both on the large and small scales. b) it is still uncertain whether the Galactic extinction law is universal, or if it changes with metallicity and/or with the UV radiation field. Detailed observations of resolved stars in the Small Magellanic Cloud by Bouchet et al. (1985) indicate that the extinction law in the optical domain is not significantly different from the Galactic one in galaxies with a UV field 10 times higher and a metallicity 10 times lower than those of the Milky Way. A∼steeper UV rise and a ∼ weaker 2200 Abump˚ than in the Galactic extinction law have been however observed in the LMC and SMC (Mathis 1990). After the results of Chapter 7 the adoption of the Galactic extinction law for external galaxies could seem not reasonable, however in this moment we have not yet a good alternative, moreover no simple analytic functions describing the geometrical distribution of emitting stars and absorbing dust, both on small and large scales, are yet available. The radiative transfer models of Witt & Gordon (2000) have however shown that the FIR to UV flux ratio, being mostly independent of the geometry, of the star formation history (the two radiations are produced by similar stellar populations) and of the adopted extinction law, is a robust estimator of the dust extinction at UV wavelengths. From the value of the TIR/FUV

179 180 A. The extinction correction

(measured or obtained with the recipes presented in Chapter 7 we can thus estimate A(FUV), following Buat et al. (2005):

A(F UV ) = 0.0333 y3 + 0.3522 y2 + 1.1960 y + 0.4967 [mag] (A.1) − ∗ ∗ ∗ where y is log(T IR/F UV ). A(λ) can be derived from A(FUV) once an extinction law and a geometry for the dust and star distribution are assumed. We adopt the sandwitch model, where a thin layer of dust of thickness ζis embedded in a thick layer of stars: 1 ζ(λ) A(λ) = 2.5 log − 1 + e−τ(λ)·sec(i) + − ·   2   ζ(λ) + 1 e−τ(λ)·sec(i) [mag] (A.2) τ(λ) sec(i) · −  ·  where the dust to stars scale height ratio ζ(λ) depends on λ (in units of A)˚ as:

ζ(λ) = 1.0867 5.501 10−5 λ. (A.3) − × · This has been calibrated adopting the average between the optically thin and optically thick cases with λ dependent dust to star scale height ratios given by Boselli & Gavazzi (1994). Observations of some edge-on nearby galaxies show that it is still unclear whether ζ depends or not on λ (Xilouris et al. 1999). As shown in Gavazzi et al. (2002a), however, similar values of Ai(λ)are obtained in the case of a sandwitch model and of the extreme case of a slab model (ζ = 1), meaning that the high uncertainty on ζ is not reflected on A(λ). In the case of the FUV band ( λ 1530 A),˚ ζ = 1, and Eq. A.2 reduces to a simple slab model. In this case τ(UV) ∼can be derived by inverting Eq. A.2:

τ(UV) = [1/sec(i)] 0.0259 + 1.2002 A (FUV) + 1.5543 A (FUV)2 + · × i × i 0.7409 A (FUV)3 + 0.2246 A (FUV)4 (A.4) − × i × i  using the galactic extinction law k(λ) (Savage & Mathis 1979), we than derive:

τ(λ) = τ(UV) k(λ)/k(UV) (A.5) ·

and we compute the complete set of Ai(λ) using Eq. A.2. Appendix B

Estimate of the < 912A˚ flux from Hα + [NII]

The stellar radiation field with λ <912 Aionizes˚ the gas, which re-emits, via recom- bination lines. If the gas is optically thick in the Lyman continuum, the number of photons in a specific recombination line is directly proportional to the number of star photons in the Lyman continuum. In the case of Hβ this number is given by equation (5.23) in Osterbrock (1989). For Hα we have:

∞ Lν dν = LHα C (B.1) Z hν · ν0

where: eff o αHβ (H , T ) FHα 1/C = hνHβ o (B.2) αB(H , T ) FHβ Assuming T =10000K and the Osterbrock case B: eff o −14 3 −1 αHβ (H , T ) = 3 10 (cm sec ) o × −13 3 −1 αB(H , T ) = 2.59 10 (cm sec ) F × and Hα = 2.87 FHβ From the Hα luminosity it is thus possible to recover the number of ionizing photons, which can be compared with the similar quantity derived from the on the model spectrum. A conservative estimate of the uncertainty on the derived < 912 A˚ flux is 1 mag.

181

Bibliography

Abadi, M. G., Moore, B., & Bower, R. G. 1999, MNRAS, 308, 947

Abazajian, K., Adelman-McCarthy, J. K., Agueros,¨ M. A., et al. 2005, AJ, 129, 1755

Abraham, R. G., Tanvir, N. R., Santiago, B. X., et al. 1996a, MNRAS, 279, L47

Abraham, R. G., van den Bergh, S., Glazebrook, K., et al. 1996b, ApJS, 107, 1

Adami, C., Biviano, A., & Mazure, A. 1998, A&A, 331, 439

Alton, P. B., Trewhella, M., Davies, J. I., et al. 1998, A&A, 335, 807

Alton, P. B., Xilouris, E. M., Bianchi, S., Davies, J., & Kylafis, N. 2000, A&A, 356, 795

Andreon, S. 1999, A&A, 351, 65

Arosio, I. 2005, Undergraduate Thesis, Universit´a degli Studi di Milano-Bicocca

Balogh, M., Eke, V., Miller, C., et al. 2004, MNRAS, 348, 1355

Barnes, J. E. & Hernquist, L. 1996, ApJ, 471, 115

Bechtold, J., Forman, W., Jones, C., et al. 1983, ApJ, 265, 26

Beers, T. C., Flynn, K., & Gebhardt, K. 1990, AJ, 100, 32

Beers, T. C., Gebhardt, K., Forman, W., Huchra, J. P., & Jones, C. 1991, AJ, 102, 1581

Bekki, K. 1999, ApJ, 510, L15

Bekki, K. & Couch, W. J. 2003, ApJ, 596, L13

Bekki, K., Couch, W. J., & Shioya, Y. 2002, ApJ, 577, 651

Bell, E. F. 2002, ApJ, 577, 150

183 184 BIBLIOGRAPHY

Bell, E. F. & Kennicutt, R. C. 2001, ApJ, 548, 681

Bernardi, M., Alonso, M. V., da Costa, L. N., et al. 2002, AJ, 123, 2990

Bernardi, M., Sheth, R. K., Annis, J., et al. 2003, AJ, 125, 1817

Bertin, E. & Arnouts, S. 1996, A&AS, 117, 393

Binggeli, B., Sandage, A., & Tammann, G. A. 1985, AJ, 90, 1681

Bird, C. M. 1994, AJ, 107, 1637

Bird, C. M. & Beers, T. C. 1993, AJ, 105, 1596

Biviano, A., Katgert, P., Mazure, A., et al. 1997, A&A, 321, 84

Bliton, M., Rizza, E., Burns, J. O., Owen, F. N., & Ledlow, M. J. 1998, MNRAS, 301, 609

Boissier, S., Boselli, A., Buat, V., Donas, J., & Milliard, B. 2004, A&A, 424, 465

Boissier, S., Gil de Paz, A., Madore, B. F., et al. 2005, ApJ, 619, L83

Boissier, S. & Prantzos, N. 2000, MNRAS, 312, 398

Boissier, S., Prantzos, N., Boselli, A., & Gavazzi, G. 2003, MNRAS, 346, 1215

Boselli, A., Boissier, S., Cortese, L., et al. 2005a, ApJ, 623, L13

Boselli, A., Cortese, L., Deharveng, J. M., et al. 2005b, ApJ, 629, L29

Boselli, A. & Gavazzi, G. 1994, A&A, 283, 12

Boselli, A. & Gavazzi, G. 2002, A&A, 386, 124

Boselli, A. & Gavazzi, G. 2006, PASP, submitted

Boselli, A., Gavazzi, G., Donas, J., & Scodeggio, M. 2001, AJ, 121, 753

Boselli, A., Gavazzi, G., Lequeux, J., et al. 1997a, A&A, 327, 522

Boselli, A., Gavazzi, G., & Sanvito, G. 2003a, A&A, 402, 37

Boselli, A., Iglesias-P´aramo, J., V´ılchez, J. M., & Gavazzi, G. 2002a, A&A, 386, 134

Boselli, A., Lequeux, J., & Gavazzi, G. 2002b, A&A, 384, 33 BIBLIOGRAPHY 185

Boselli, A., Sauvage, M., Lequeux, J., Donati, A., & Gavazzi, G. 2003b, A&A, 406, 867

Boselli, A., Tuffs, R. J., Gavazzi, G., Hippelein, H., & Pierini, D. 1997b, A&AS, 121, 507

Bouchet, P., Lequeux, J., Maurice, E., Prevot, L., & Prevot-Burnichon, M. L. 1985, A&A, 149, 330

Bower, R. G. & Balogh, M. L. 2004, in Clusters of Galaxies: Probes of Cosmological Structure and Galaxy Evolution, 326

Bravo-Alfaro, H., Cayatte, V., van Gorkom, J. H., & Balkowski, C. 2000, AJ, 119, 580

Bravo-Alfaro, H., Cayatte, V., van Gorkom, J. H., & Balkowski, C. 2001, A&A, 379, 347

Bressan, A., Chiosi, C., & Fagotto, F. 1994, ApJS, 94, 63

Brown, T. M., Bowers, C. W., Kimble, R. A., Sweigart, A. V., & Ferguson, H. C. 2000, ApJ, 532, 308

Bruzual, G. & Charlot, S. 2003, MNRAS, 344, 1000

Buat, V. 1992, A&A, 264, 444

Buat, V., Boselli, A., Gavazzi, G., & Bonfanti, C. 2002, A&A, 383, 801

Buat, V., Donas, J., Milliard, B., & Xu, C. 1999, A&A, 352, 371

Buat, V., Iglesias-P´aramo, J., Seibert, M., et al. 2005, ApJ, 619, L51

Buat, V. & Xu, C. 1996, A&A, 306, 61

Burgarella, D., Buat, V., & Iglesias-P´aramo, J. 2005, MNRAS, 360, 1413

Burstein, D., Bertola, F., Buson, L. M., Faber, S. M., & Lauer, T. R. 1988, ApJ, 328, 440

Burstein, D. & Heiles, C. 1982, AJ, 87, 1165

Butcher, H. & Oemler, A. 1978, ApJ, 226, 559

Butcher, H. & Oemler, A. 1984, ApJ, 285, 426

Byrd, G. & Valtonen, M. 1990, ApJ, 350, 89 186 BIBLIOGRAPHY

Caldwell, N., Rose, J. A., & Concannon, K. D. 2003, AJ, 125, 2891

Calzetti, D. 1997, AJ, 113, 162

Calzetti, D. 2001, PASP, 113, 1449

Calzetti, D., Bohlin, R. C., Kinney, A. L., Storchi-Bergmann, T., & Heckman, T. M. 1995, ApJ, 443, 136

Calzetti, D., Kennicutt, R. C., Bianchi, L., et al. 2005, ApJ, in press, astro- ph/0507427

Calzetti, D., Kinney, A. L., & Storchi-Bergmann, T. 1994, ApJ, 429, 582

Caplan, J. & Deharveng, L. 1986, A&A, 155, 297

Cardelli, J. A., Clayton, G. C., & Mathis, J. S. 1989, ApJ, 345, 245

Cayatte, V., Kotanyi, C., Balkowski, C., & van Gorkom, J. H. 1994, AJ, 107, 1003

Cayatte, V., van Gorkom, J. H., Balkowski, C., & Kotanyi, C. 1990, AJ, 100, 604

Charlot, S. & Fall, S. M. 2000, ApJ, 539, 718

Charlot, S. & Longhetti, M. 2001, MNRAS, 323, 887

Chemin, L., Cayatte, V., Balkowski, C., et al. 2005, A&A, 436, 469

Chung, A., van Gorkom, J. H. Kenney, J. D. P., & Vollmer, B. 2005, astro-ph/0507592

Churazov, E., Forman, W., Jones, C., & B¨ohringer, H. 2003, ApJ, 590, 225

Cohen, M., Sasseen, T. P., & Bowyer, S. 1994, ApJ, 427, 848

Colless, M. 1989, MNRAS, 237, 799

Colless, M., Dalton, G., Maddox, S., et al. 2001, MNRAS, 328, 1039

Combes, F., Dupraz, C., Casoli, F., & Pagani, L. 1988, A&A, 203, L9

Conselice, C. J., Blackburne, J. A., & Papovich, C. 2005, ApJ, 620, 564

Contursi, A., Boselli, A., Gavazzi, G., et al. 2001, A&A, 365, 11

Cortese, L., Boselli, A., Buat, V., et al. 2006, ApJ, in press

Cortese, L., Boselli, A., Gavazzi, G., et al. 2005, ApJ, 623, L17 BIBLIOGRAPHY 187

Cortese, L., Gavazzi, G., Boselli, A., Iglesias-Paramo, J., & Carrasco, L. 2004, A&A, 425, 429

Cortese, L., Gavazzi, G., Boselli, A., et al. 2003a, A&A, 410, L25

Cortese, L., Gavazzi, G., Iglesias-Paramo, J., Boselli, A., & Carrasco, L. 2003b, A&A, 401, 471

Couch, W. J., Barger, A. J., Smail, I., Ellis, R. S., & Sharples, R. M. 1998, ApJ, 497, 188

Cowie, L. L. & Songaila, A. 1977, Nature, 266, 501

Dale, D. A., Helou, G., Contursi, A., Silbermann, N. A., & Kolhatkar, S. 2001, ApJ, 549, 215

De Propris, R., Colless, M., Driver, S. P., et al. 2003, MNRAS, 342, 725 de Rijcke, S., Michielsen, D., Dejonghe, H., Zeilinger, W. W., & Hau, G. K. T. 2005, A&A, 438, 491

Deharveng, J.-M., Boselli, A., & Donas, J. 2002, A&A, 393, 843

Deharveng, J.-M., Sasseen, T. P., Buat, V., et al. 1994, A&A, 289, 715

Devereux, N. A. & Young, J. S. 1990, ApJ, 359, 42

Dickey, J. M. & Gavazzi, G. 1991, ApJ, 373, 347

Dickinson, M., Papovich, C., Ferguson, H. C., & Budav´ari, T. 2003, ApJ, 587, 25

Domainko, W., Mair, M., Kapferer, W., et al. 2005, astro-ph/0507605

Donas, J., Milliard, B., & Laget, M. 1991, A&A, 252, 487

Donas, J., Milliard, B., & Laget, M. 1995, A&A, 303, 661

Donati, A. 2004, Undergraduate Thesis, Universit´a degli Studi di Milano

Donnelly, R. H., Forman, W., Jones, C., et al. 2001, ApJ, 562, 254

Donnelly, R. H., Markevitch, M., Forman, W., et al. 1998, ApJ, 500, 138

Draine, B. T. 2003, ApJ, 598, 1017

Dressler, A. 1980, ApJ, 236, 351 188 BIBLIOGRAPHY

Dressler, A. 2004, in Clusters of Galaxies: Probes of Cosmological Structure and Galaxy Evolution, 207–+

Dressler, A., Oemler, A. J., Couch, W. J., et al. 1997, ApJ, 490, 577

Dressler, A. & Sandage, A. 1983, ApJ, 265, 664

Dressler, A. & Shectman, S. A. 1988, AJ, 95, 985

Driver, S. P., Fernandez-Soto, A., Couch, W. J., et al. 1998, ApJ, 496, L93+

Duc, P.-A., Bournaud, F., & Masset, F. 2004, A&A, 427, 803

Duc, P.-A., Brinks, E., Springel, V., et al. 2000, AJ, 120, 1238

Duc, P.-A. & Mirabel, I. F. 1998, A&A, 333, 813

Duc, P.-A. & Mirabel, I. F. 1999, in IAU Symp. 186: Galaxy Interactions at Low and High Redshift, 61

Dutil, Y. & Roy, J. 1999, ApJ, 516, 62

Enßlin, T. A., Biermann, P. L., Klein, U., & Kohle, S. 1998, A&A, 332, 395

Enßlin, T. A. & Bruggen,¨ M. 2002, MNRAS, 331, 1011

Fabricant, D., Franx, M., & van Dokkum, P. 2000, ApJ, 539, 577

Ferguson, H. C. 1994, in Dwarf Galaxies, 475

Ferguson, H. C. & Binggeli, B. 1994, A&A Rev., 6, 67

Ferguson, H. C. & Sandage, A. 1991, AJ, 101, 765

Forbes, D. A., Beasley, M. A., Bekki, K., Brodie, J. P., & Strader, J. 2003, Science, 301, 1217

Forman, W., Churazov, E., David, L., et al. 2003, astro-ph/0301476

Franzetti, P. 2005, Ph.D. Thesis, Universit´a degli Studi di Milano-Bicocca

Frei, Z., Guhathakurta, P., Gunn, J. E., & Tyson, J. A. 1996, AJ, 111, 174

Fujita, Y. 2004, PASJ, 56, 29

Fujita, Y. & Nagashima, M. 1999, ApJ, 516, 619

Gavazzi, G. 1978, A&A, 69, 355 BIBLIOGRAPHY 189

Gavazzi, G., Bonfanti, C., Sanvito, G., Boselli, A., & Scodeggio, M. 2002a, ApJ, 576, 135

Gavazzi, G., Boselli, A., Cortese, L., et al. 2006, A&A, in press

Gavazzi, G., Boselli, A., Donati, A., Franzetti, P., & Scodeggio, M. 2003a, A&A, 400, 451

Gavazzi, G., Boselli, A., Mayer, L., et al. 2001a, ApJ, 563, L23

Gavazzi, G., Boselli, A., Pedotti, P., Gallazzi, A., & Carrasco, L. 2002b, A&A, 386, 114

Gavazzi, G., Boselli, A., Pedotti, P., Gallazzi, A., & Carrasco, L. 2002c, A&A, 396, 449

Gavazzi, G., Boselli, A., Scodeggio, M., Pierini, D., & Belsole, E. 1999a, MNRAS, 304, 595

Gavazzi, G., Carrasco, L., & Galli, R. 1999b, A&AS, 136, 227

Gavazzi, G., Catinella, B., Carrasco, L., Boselli, A., & Contursi, A. 1998, AJ, 115, 1745

Gavazzi, G., Contursi, A., Carrasco, L., et al. 1995, A&A, 304, 325

Gavazzi, G., Cortese, L., Boselli, A., et al. 2003b, ApJ, 597, 210

Gavazzi, G., Donati, A., Cucciati, O., et al. 2005, A&A, 430, 411

Gavazzi, G., Franzetti, P., Scodeggio, M., Boselli, A., & Pierini, D. 2000, A&A, 361, 863

Gavazzi, G. & Jaffe, W. 1987, A&A, 186, L1

Gavazzi, G., Marcelin, M., Boselli, A., et al. 2001b, A&A, 377, 745

Gavazzi, G., Perola, G. C., & Jaffe, W. 1981, A&A, 103, 35

Gavazzi, G., Pierini, D., & Boselli, A. 1996, A&A, 312, 397

Gavazzi, G. & Trinchieri, G. 1983, ApJ, 270, 410

Gavazzi, G., Zaccardo, A., Sanvito, G., Boselli, A., & Bonfanti, C. 2004, A&A, 417, 499

Gavazzi, G., Zibetti, S., Boselli, A., et al. 2001c, A&A, 372, 29 190 BIBLIOGRAPHY

Geller, M. J. & Huchra, J. P. 1989, Science, 246, 897

Ghigna, S., Moore, B., Governato, F., et al. 1998, MNRAS, 300, 146

Gil de Paz, A. & Madore, B. F. 2005, ApJS, 156, 345

Girardi, M., Giuricin, G., Mardirossian, F., Mezzetti, M., & Boschin, W. 1998, ApJ, 505, 74

Glazebrook, K., Blake, C., Economou, F., Lilly, S., & Colless, M. 1999, MNRAS, 306, 843

Gnedin, O. Y. 2003, ApJ, 582, 141

Godwin, J. G., Metcalfe, N., & Peach, J. V. 1983, MNRAS, 202, 113

Godwin, J. G. & Peach, J. V. 1982, MNRAS, 200, 733

Golev, V. & Prugniel, P. 1998, A&AS, 132, 255

Gomez, P. L., Loken, C., Roettiger, K., & Burns, J. O. 2002, ApJ, 569, 122

Gordon, K. D., Clayton, G. C., Witt, A. N., & Misselt, K. A. 2000, ApJ, 533, 236

Goto, T. 2005, MNRAS, 359, 1415

Gottl¨ober, S., Klypin, A., & Kravtsov, A. V. 2001, ApJ, 546, 223

Grebel, E. K. 2000, in ESA SP-445: Star Formation from the Small to the Large Scale, 87

Grebenev, S. A., Forman, W., Jones, C., & Murray, S. 1995, ApJ, 445, 607

Greggio, L. & Renzini, A. 1990, ApJ, 364, 35

Gunn, J. E. & Gott, J. R. I. 1972, ApJ, 176, 1

Hawkins, E., Maddox, S., Cole, S., et al. 2003, MNRAS, 346, 78

Hayakawa, A., Furusho, T., Yamasaki, N. Y., Ishida, M., & Ohashi, T. 2004, PASJ, 56, 743

Haynes, M. P. & Giovanelli, R. 1984, AJ, 89, 758

Haynes, M. P., Jore, K. P., Barrett, E. A., Broeils, A. H., & Murray, B. M. 2000, AJ, 120, 703 BIBLIOGRAPHY 191

Heckman, T. M., Robert, C., Leitherer, C., Garnett, D. R., & van der Rydt, F. 1998, ApJ, 503, 646

Heisler, J., Tremaine, S., & Bahcall, J. N. 1985, ApJ, 298, 8

Helfer, T. T., Thornley, M. D., Regan, M. W., et al. 2003, ApJS, 145, 259

Helou, G., Khan, I. R., Malek, L., & Boehmer, L. 1988, ApJS, 68, 151

Hibbard, J. E., Bianchi, L., Thilker, D. A., et al. 2005, ApJ, 619, L87

Hibbard, J. E. & van Gorkom, J. H. 1996, AJ, 111, 655

Hinz, J. L., Rieke, G. H., & Caldwell, N. 2003, AJ, 126, 2622

Hinz, J. L., Rix, H.-W., & Bernstein, G. M. 2001, AJ, 121, 683

Hoopes, C. G., Heckman, T. M., Strickland, D. K., et al. 2005, ApJ, 619, L99

Hubble, E. & Humason, M. L. 1931, ApJ, 74, 43

Hubble, E. P. 1925, The Observatory, 48, 139

Hummel, E. & Saikia, D. J. 1991, A&A, 249, 43

Hunsberger, S. D., Charlton, J. C., & Zaritsky, D. 1998, ApJ, 505, 536

Ibata, R., Irwin, M., Lewis, G., Ferguson, A. M. N., & Tanvir, N. 2001, Nature, 412, 49

Iglesias-P´aramo, J., Boselli, A., Cortese, L., V´ılchez, J. M., & Gavazzi, G. 2002, A&A, 384, 383

Iglesias-P´aramo, J., Boselli, A., Gavazzi, G., Cortese, L., & V´ılchez, J. M. 2003, A&A, 397, 421

Iglesias-P´aramo, J., Boselli, A., Gavazzi, G., & Zaccardo, A. 2004, A&A, 421, 887

Ilbert, O., Tresse, L., Zucca, E., et al. 2004, astro-ph/0409134

Inoue, A. K. 2003, PASJ, 55, 901

Iono, D., Yun, M. S., & Mihos, J. C. 2004, ApJ, 616, 199

Isobe, T., Feigelson, E. D., Akritas, M. G., & Babu, G. J. 1990, ApJ, 364, 104

Issa, M. R., MacLaren, I., & Wolfendale, A. W. 1990, A&A, 236, 237 192 BIBLIOGRAPHY

Jarrett, T. H., Chester, T., Cutri, R., Schneider, S. E., & Huchra, J. P. 2003, AJ, 125, 525

Johnston, K. V., Sackett, P. D., & Bullock, J. S. 2001, ApJ, 557, 137

Jones, C., Mandel, E., Schwarz, J., et al. 1979, ApJ, 234, L21

Jore, K. P., Broeils, A. H., & Haynes, M. P. 1996, AJ, 112, 438

Katz, N. & White, S. D. M. 1993, ApJ, 412, 455

Kauffmann, G. 1995, MNRAS, 274, 153

Kauffmann, G., Heckman, T. M., Tremonti, C., et al. 2003a, MNRAS, 346, 1055

Kauffmann, G., Heckman, T. M., White, S. D. M., et al. 2003b, MNRAS, 341, 33

Kauffmann, G., White, S. D. M., & Guiderdoni, B. 1993, MNRAS, 264, 201

Kenney, J. D. & Young, J. S. 1988, ApJS, 66, 261

Kenney, J. D. P., Rubin, V. C., Planesas, P., & Young, J. S. 1995, ApJ, 438, 135

Kenney, J. D. P. & Yale, E. E. 2002, ApJ, 567, 865

Kennicutt, R. C. 1983, ApJ, 272, 54

Kennicutt, R. C. 1992, ApJ, 388, 310

Kennicutt, R. C. 1998, ARA&A, 36, 189

Kennicutt, R. C., Tamblyn, P., & Congdon, C. E. 1994, ApJ, 435, 22

Kewley, L. J. & Dopita, M. A. 2002, ApJS, 142, 35

Kewley, L. J., Geller, M. J., Jansen, R. A., & Dopita, M. A. 2002, AJ, 124, 3135

Kewley, L. J., Jansen, R. A., & Geller, M. J. 2005, PASP, 117, 227

Klypin, A., Kravtsov, A. V., Valenzuela, O., & Prada, F. 1999, ApJ, 522, 82

Kodaira, K., Watanabe, T., Onaka, T., & Tanaka, W. 1990, ApJ, 363, 422

Kodama, T., Yamada, T., Akiyama, M., et al. 2004, MNRAS, 350, 1005

Kogut, A., Spergel, D. N., Barnes, C., et al. 2003, ApJS, 148, 161

Kojima, M. & Noguchi, M. 1997, ApJ, 481, 132 BIBLIOGRAPHY 193

Kong, X., Charlot, S., Brinchmann, J., & Fall, S. M. 2004, MNRAS, 349, 769

Koopmann, R. A. & Kenney, J. D. P. 2004a, ApJ, 613, 866

Koopmann, R. A. & Kenney, J. D. P. 2004b, ApJ, 613, 851

Kotanyi, C., van Gorkom, J. H., & Ekers, R. D. 1983, ApJ, 273, L7

Kroupa, P. 1998, MNRAS, 300, 200

Kroupa, P., Tout, C. A., & Gilmore, G. 1993, MNRAS, 262, 545

Kuntschner, H. & Davies, R. L. 1998, MNRAS, 295, L29

Labb´e, I., Rudnick, G., Franx, M., et al. 2003, ApJ, 591, L95

Laird, E. S., Nandra, K., Adelberger, K. L., Steidel, C. C., & Reddy, N. A. 2005, MNRAS, 359, 47

Lampton, M., Deharveng, J. M., & Bowyer, S. 1990, in IAU Symp. 139, 449

Larson, R. B., Tinsley, B. M., & Caldwell, C. N. 1980a, ApJ, 237, 692

Larson, R. B., Tinsley, B. M., & Caldwell, C. N. 1980b, ApJ, 237, 692

Leitherer, C. & Heckman, T. M. 1995, ApJS, 96, 9

Leitherer, C., Schaerer, D., Goldader, J. D., et al. 1999, ApJS, 123, 3

Lilly, S. J., Le Fevre, O., Hammer, F., & Crampton, D. 1996, ApJ, 460, L1+

Longhetti, M., Bressan, A., Chiosi, C., & Rampazzo, R. 1999, A&A, 345, 419

Machacek, M. E., Jones, C., & Forman, W. R. 2004, ApJ, 610, 183

Madau, P., Ferrara, A., & Rees, M. J. 2001, ApJ, 555, 92

Madau, P., Pozzetti, L., & Dickinson, M. 1998, ApJ, 498, 106

Madgwick, D. S., Lahav, O., Baldry, I. K., et al. 2002, MNRAS, 333, 133

Maloney, P. R., Hollenbach, D. J., & Tielens, A. G. G. M. 1996, ApJ, 466, 561

Martin, D. C., Fanson, J., Schiminovich, D., et al. 2005, ApJ, 619, L1

Mathis, J. S. 1990, ARA&A, 28, 37 194 BIBLIOGRAPHY

Mayer, L., Mastropietro, C., Wadsley, J., Stadel, J., & Moore, B. 2005, MNRAS sub- mitted, astro-ph/0504277

McGaugh, S. S. 1991, ApJ, 380, 140

Mendes de Oliveira, C., Cypriano, E. S., Sodr´e, L., & Balkowski, C. 2004, ApJ, 605, L17

Merritt, D. 1983, ApJ, 264, 24

Merritt, D. 1984, ApJ, 276, 26

Meurer, G. R., Heckman, T. M., & Calzetti, D. 1999, ApJ, 521, 64

Meurer, G. R., Heckman, T. M., Leitherer, C., et al. 1995, AJ, 110, 2665

Mihos, J. C. 2004a, in Clusters of Galaxies: Probes of Cosmological Structure and Galaxy Evolution, 278

Mihos, J. C. 2004b, in IAU Symposium, 390

Mihos, J. C., Richstone, D. O., & Bothun, G. D. 1991, ApJ, 377, 72

Miller, R. H. 1986, A&A, 167, 41

Milliard, B., Donas, J., & Laget, M. 1991, Advances in Space Research, 11, 135

Milliard, B., Donas, J., Laget, M., Armand, C., & Vuillemin, A. 1992, A&A, 257, 24

Mobasher, B., Colless, M., Carter, D., et al. 2003, ApJ, 587, 605

Moore, B., Ghigna, S., Governato, F., et al. 1999, ApJ, 524, L19

Moore, B., Katz, N., Lake, G., Dressler, A., & Oemler, A. 1996, Nature, 379, 613

Moore, B., Lake, G., & Katz, N. 1998, ApJ, 495, 139

Mori, M., Ferrara, A., & Madau, P. 2002, ApJ, 571, 40

Morrissey, P., Schiminovich, D., Barlow, T. A., et al. 2005, ApJ, 619, L7

Mulchaey, J. S. 2000, ARA&A, 38, 289

Nagashima, M. & Yoshii, Y. 2004, ApJ, 610, 23

Neff, S. G., Thilker, D. A., Seibert, M., et al. 2005, ApJ, 619, L91

Neistein, E., Maoz, D., Rix, H., & Tonry, J. L. 1999, AJ, 117, 2666 BIBLIOGRAPHY 195

Nulsen, P. E. J. 1982, MNRAS, 198, 1007

O’Connell, R. W. 1999, ARA&A, 37, 603

Okamoto, T. & Habe, A. 1999, ApJ, 516, 591

Oke, J. B. 1974, ApJS, 27, 21

O’Neil, K. 2004, AJ, 128, 2080

Oosterloo, T. & van Gorkom, J. 2005, A&A, 437, L19

Osterbrock, D. E. 1989, Astrophysics of gaseous nebulae and active galactic nu- clei (Research supported by the University of California, John Simon Guggenheim Memorial Foundation, University of Minnesota, et al. Mill Valley, CA, University Science Books, 1989, 422 p.)

Ostriker, J. P. 1980, Comments on Astrophysics, 8, 177

Panuzzo, P., Bressan, A., Granato, G. L., Silva, L., & Danese, L. 2003, A&A, 409, 99

Pinkney, J., Roettiger, K., Burns, J. O., & Bird, C. M. 1996, ApJS, 104, 1

Poggianti, B. 2004a, Baryons in Dark Matter Halos

Poggianti, B. M. 2004b, in Clusters of Galaxies: Probes of Cosmological Structure and Galaxy Evolution, 246

Poggianti, B. M. & Barbaro, G. 1997, A&A, 325, 1025

Poggianti, B. M., Bridges, T. J., Komiyama, Y., et al. 2004, ApJ, 601, 197

Poggianti, B. M., Bridges, T. J., Mobasher, B., et al. 2001, ApJ, 562, 689

Poggianti, B. M., Smail, I., Dressler, A., et al. 1999, ApJ, 518, 576

Quilis, V., Moore, B., & Bower, R. 2000, Science, 288, 1617

Rich, R. M., Salim, S., Brinchmann, J., et al. 2005, ApJ, 619, L107

Rines, K., Geller, M. J., Kurtz, M. J., & Diaferio, A. 2003, AJ, 126, 2152

Roediger, E. & Hensler, G. 2005, A&A, 433, 875

Rubin, V. C., Waterman, A. H., & Kenney, J. D. P. 1999, AJ, 118, 236

Sakai, S., Kennicutt, R. C., van der Hulst, J. M., & Moss, C. 2002, ApJ, 578, 842 196 BIBLIOGRAPHY

Salpeter, E. E. 1955, ApJ, 121, 161

Savage, B. D. & Mathis, J. S. 1979, ARA&A, 17, 73

Saviane, I., Hibbard, J. E., & Rich, R. M. 2004, AJ, 127, 660

Schechter, P. 1976, ApJ, 203, 297

Schindler, S. & Mueller, E. 1993, A&A, 272, 137

Schlegel, D. J., Finkbeiner, D. P., & Davis, M. 1998, ApJ, 500, 525

Scodeggio, M., Gavazzi, G., Franzetti, P., et al. 2002, A&A, 384, 812

Seibert, M., Martin, D. C., Heckman, T. M., et al. 2005, ApJ, 619, L55

Shang, Z., Brinks, E., Zheng, Z., et al. 1998, ApJ, 504, L23

Skillman, E. D., Kennicutt, R. C., & Hodge, P. W. 1989, ApJ, 347, 875

Smail, I., Ellis, R. S., Dressler, A., et al. 1997, ApJ, 479, 70

Smail, I., Kuntschner, H., Kodama, T., et al. 2001, MNRAS, 323, 839

Smith, A. M. & Cornett, R. H. 1982, ApJ, 261, 1

Smith, G. P., Treu, T., Ellis, R. S., Moran, S. M., & Dressler, A. 2005, ApJ, 620, 78

Sodre, L. J., Capelato, H. V., Steiner, J. E., & Mazure, A. 1989, AJ, 97, 1279

Somerville, R. S. 2002, ApJ, 572, L23

Springel, V. & Hernquist, L. 2003, MNRAS, 339, 312

Springel, V., White, S. D. M., Jenkins, A., et al. 2005, Nature, 435, 629

Steidel, C. C., Adelberger, K. L., Giavalisco, M., Dickinson, M., & Pettini, M. 1999, ApJ, 519, 1

Stein, P. 1997, A&A, 317, 670

Struble, M. F. & Rood, H. J. 1999, ApJS, 125, 35

Sullivan, M., Treyer, M. A., Ellis, R. S., et al. 2000, MNRAS, 312, 442

Sun, M. & Murray, S. S. 2002, ApJ, 576, 708

Tantalo, R., Chiosi, C., Bressan, A., & Fagotto, F. 1996, A&A, 311, 361 BIBLIOGRAPHY 197

Terlevich, A. I., Kuntschner, H., Bower, R. G., Caldwell, N., & Sharples, R. M. 1999, MNRAS, 310, 445

Thuan, T. X. & Sauvage, M. 1992, A&AS, 92, 749

Tonry, J. & Davis, M. 1979, AJ, 84, 1511

Tonry, J. L., Schmidt, B. P., Barris, B., et al. 2003, ApJ, 594, 1

Toomre, A. & Toomre, J. 1972, ApJ, 178, 623

Tornatore, L., Borgani, S., Matteucci, F., Recchi, S., & Tozzi, P. 2004, MNRAS, 349, L19

Trager, S. C., Faber, S. M., Worthey, G., & Gonz´alez, J. J. 2000, AJ, 120, 165

Tremonti, C. A., Heckman, T. M., Kauffmann, G., et al. 2004, ApJ, 613, 898

Trentham, N., Sampson, L., & Banerji, M. 2005, MNRAS, 357, 783

Treu, T., Ellis, R. S., Kneib, J.-P., et al. 2003, ApJ, 591, 53

Treu, T., Ellis, R. S., Liao, T. X., & van Dokkum, P. G. 2005, ApJ, 622, L5

Treyer, M., Wyder, T. K., Schiminovich, D., et al. 2005, ApJ, 619, L19

Treyer, M. A., Ellis, R. S., Milliard, B., Donas, J., & Bridges, T. J. 1998, MNRAS, 300, 303

Trimble, V. 1995, PASP, 107, 1133

Trujillo, I., Rudnick, G., Rix, H.-W., et al. 2004, ApJ, 604, 521

V´azquez, G. A. & Leitherer, C. 2005, ApJ, 621, 695

van den Bergh, S. 1976, ApJ, 206, 883

van Dokkum, P. G., Franx, M., Kelson, D. D., et al. 1998, ApJ, 500, 714

van Gorkom, J. H. 2004, in Clusters of Galaxies: Probes of Cosmological Structure and Galaxy Evolution, 306

van Zee, L., Salzer, J. J., Haynes, M. P., O’Donoghue, A. A., & Balonek, T. J. 1998, AJ, 116, 2805 van Zee, L., Skillman, E. D., & Haynes, M. P. 2004, AJ, 128, 121

Vogt, N. P., Haynes, M. P., Herter, T., & Giovanelli, R. 2004, AJ, 127, 3273 198 BIBLIOGRAPHY

Vollmer, B., Balkowski, C., Cayatte, V., van Driel, W., & Huchtmeier, W. 2004a, A&A, 419, 35

Vollmer, B., Beck, R., Kenney, J. D. P., & van Gorkom, J. H. 2004b, AJ, 127, 3375

Vollmer, B., Braine, J., Combes, F., & Sofue, Y. 2005, A&A in press, astro- ph/0507252

Vollmer, B., Cayatte, V., Balkowski, C., & Duschl, W. J. 2001, ApJ, 561, 708

Walker, I. R., Mihos, J. C., & Hernquist, L. 1996, ApJ, 460, 121

Wang, B. & Heckman, T. M. 1996, ApJ, 457, 645

Wehner, E. H. & Gallagher, J. S. 2005, ApJ, 618, L21

Weilbacher, P. M., Duc, P.-A., Fritze v. Alvensleben, U., Martin, P., & Fricke, K. J. 2000, A&A, 358, 819

West, M. J. & Blakeslee, J. P. 2000, ApJ, 543, L27

West, M. J. & Bothun, G. D. 1990, ApJ, 350, 36

West, M. J., Villumsen, J. V., & Dekel, A. 1991, ApJ, 369, 287

White, S. D. M. & Rees, M. J. 1978, MNRAS, 183, 341

Whitmore, B. C., Gilmore, D. M., & Jones, C. 1993, ApJ, 407, 489

Willman, B., Governato, F., Wadsley, J., & Quinn, T. 2004, MNRAS, 355, 159

Witt, A. N. & Gordon, K. D. 2000, ApJ, 528, 799

Wyder, T. K., Treyer, M. A., Milliard, B., et al. 2005, ApJ, 619, L15

Xilouris, E. M., Byun, Y. I., Kylafis, N. D., Paleologou, E. V., & Papamastorakis, J. 1999, A&A, 344, 868

Xu, C. & Buat, V. 1995, A&A, 293, L65

Xu, C. K., Donas, J., Arnouts, S., et al. 2005, ApJ, 619, L11

Yahil, A. & Vidal, N. V. 1977, ApJ, 214, 347

Yi, S., Demarque, P., & Oemler, A. J. 1998, ApJ, 492, 480

Yi, S. K., Yoon, S.-J., Kaviraj, S., et al. 2005, ApJ, 619, L111 BIBLIOGRAPHY 199

Yoshida, M., Ohyama, Y., Iye, M., et al. 2004, AJ, 127, 90

Yoshii, Y. & Arimoto, N. 1987, A&A, 188, 13

Zabludoff, A. I., Geller, M. J., Huchra, J. P., & Ramella, M. 1993, AJ, 106, 1301

Zaritsky, D., Kennicutt, R. C., & Huchra, J. P. 1994, ApJ, 420, 87

Zibetti, S., Gavazzi, G., Scodeggio, M., Franzetti, P., & Boselli, A. 2002, ApJ, 579, 261

Zwicky, F., Herzog, E., & Wild, P. 1961, Catalogue of galaxies and of clusters of galaxies (Pasadena: California Institute of Technology (CIT))

List of Figures

1.1 An example of the heterogeneous population of galaxies that inhabit our Universe. Mosaic of RGB (g,r,i) images adapted from Frei et al. (1996) ...... 6

2.1 Cross section of the instrument portion of GALEX. The optical path is outlined in blue. Overall dimensions of the view shown are 1.5 m 1 m (adapted from Morrissey et al. 2005)...... 16 2.2 The transmittance profile for the NUV and FUV GALEX filters. Dif- ferent galaxy spectral energy distributions are superposed...... 19 2.3 Example of GALEX image. GALEX NGS observation of NGC4631. In the color table, red-green (gold) is used for NUV, and blue for FUV. 19

3.1 The UV luminosity functions for the four analyzed data sets...... 27 3.2 The composite UV luminosity function of 3 nearby clusters. The solid line represents the best Schechter fit to the data for MUV 16.5. . . 28 3.3 The UV bi-variate composite luminosity functions of nearb≤y−clusters. Red (UV B > 2) and blue (UV B < 2) galaxies are indicated with − − empty and filled circles respectively...... 29 3.4 The cluster and the field UV luminosity functions. The composite cluster LF is given with filled circles. The solid line indicates the best Schechter fit of the field LF of Sullivan et al. (2000). The normalization is such that the two LFs match at M 19.25...... 31 UV ∼ − 4.1 The GALEX observation of Abell1367. ROSAT X-ray contour are superposed in black. The tick rectangular region indicates the region covered by the optical catalogues used for the star/galaxy discrimination. 34 4.2 Comparison between FOCA (upper image) and GALEX (lower image) observation of the center of Abell1367. It emerges clearly the strong improvement in resolution and sensitiveness of new GALEX data. . . 35 4.3 Left: The comparison between FOCA and GALEX NUV (left) and FUV (right) magnitudes of galaxies in Abell1367. The continuum line shows the best linear fit to the data...... 36

201 202 LIST OF FIGURES

4.4 The redshift completeness per bin of UV magnitude in Abell 1367. . . 37 4.5 The GALEX NUV (left) and FUV (right) LF for Abell 1367. Open dots are obtained using the subtraction of field counts obtained by Xu et al. (2005); filled dots are obtained using the completeness corrected method. The solid line represents the best Schechter fit. The dotted line shows the composite nearby clusters 2000 A˚ LF by (Cortese et al. 2003a). The dashed line represents the GALEX local field LF (Wyder et al. 2005), normalized in order to match the cluster LF at M AB ∼ 17.80...... 38 − 4.6 The NUV (left) and FUV (right) bi-variate LF of A1367. Star-forming and quiescent galaxies are indicated with empty triangles and filled squares respectively. The dashed line represents the GALEX local field LF (Wyder et al. 2005), normalized as in Fig.4.5 ...... 38 4.7 The FUV-NUV color magnitude relation for confirmed members of A1367. Symbols are as in Fig.4.6 ...... 41 4.8 The optical (r0-band) distribution for star forming (blue histogram) and quiescent (red histogram) galaxies in our sample...... 41

5.1 Cumulative redshift distribution for galaxies in the studied region. . . 46 5.2 Velocity histogram and stripe density plot for the members of Abell 1367. Arrows mark the location of the most significant weighted gaps in the velocity distribution...... 47 5.3 Local deviations from the global kinematics for galaxies in Abell 1367 as measured by the Dressler & Shectman (1988) test. Galaxies are marked with open circles whose radius scales with their local deviation δ from the global kinematics. The ROSAT X-ray contours are shown with dotted lines...... 51 5.4 Palomar DSS image of the central region ( 1.3 square degrees) of ∼ Abell 1367 studied in this Chapter. The iso-density contours for the 146 confirmed cluster members are superposed. The lowest iso-density contour correspond to 3σ above the mean density in the field (left). The ROSAT X-ray contours are superposed in red (right). The straight line indicates the position of the abrupt gas temperature gradient detected by ASCA (Donnelly et al. 1998), used to divide our sample into two subclusters: the North-West and the South-East...... 52 5.5 The LOS velocity field (left) and the velocity dispersion field (right) for the whole region studied in this Chapter. The LOS velocity and the velocity dispersion are computed using the 10 nearest neighbors to each pixel, whose size is 36 arcsec2. The iso-density contours for the 146 confirmed cluster members are superposed in black...... 53 LIST OF FIGURES 203

5.6 A 3D sketch of Abell 1367 summarizing the various sub-components described in Section 5.5. The cluster is viewed from its near side, as suggested by the eyeball indicating the observer’s position...... 54 5.7 Blow-up of the NW substructure of Abell 1367. The arrows indicate the direction of radio head tails associated with 97-079 and 97-073 and the orientation of the NAT radio galaxy 97-095. The dashed region shows the distribution of the diffuse cluster radio relic (Gavazzi 1978). The iso-density contours for the 146 confirmed cluster members are superposed...... 55 5.8 The LOS velocity distribution for galaxies in the NW (upper) and in the SE (lower) subclusters...... 56 5.9 The velocity dispersion radial profile of the NW (upper) and the SE (lower) subclusters...... 57 5.10 The distribution of galaxies belonging to the South-East subcluster. Triangles indicate galaxies with LOS velocity > 7500 km s−1, circles galaxies with LOS velocity < 5800 km s−1 and squares galaxies with LOS velocity comprises in the range 5800 km s−1 < V < 7500 km s−1. The ROSAT X-ray contours are shown...... 58 5.11 The LOS velocity distribution for emission line (upper) and non emis- sion line galaxies (lower) in the whole cluster sample...... 59 5.12 Projected density map of non emission line (left) and emission line (right) galaxies in Abell 1367. The iso-density contours of the 146 confirmed cluster members are superposed...... 60 5.13 The bound and unbound orbit regions in the (Vrel, α) plane. The bound-incoming solutions (BIa and BIb), the bound-outgoing solu- tions (BO) and the unbound-outgoing (UO) solutions are indicated with solid lines. The dotted lines show the dividing line between bound and unbound regions. The vertical solid lines represent the observed Vrel and the dashed regions their associated 1σ uncertainty...... 63

6.1 The near-UV (left column) and far-UV (right column) to optical and near-IR color magnitude relations. Colors are in the AB magnitude system. Open circles are for ellipticals, filled circles for dwarfs, crosses for lenticulars (S0-S0a). Galaxies redder than the dashed line are un- detectable by the present survey (at the NGS limit). Largest 1σ errors for luminous and dwarf systems are given...... 72 6.2 The relationship between the UV color index (F UV NUV ) and a) the total H band luminosity, b) the B-H color index, −c) the logarithm of the central velocity dispersion and d) the Mg2 index. Symbols are as in fig. 6.1. Labeled points indicate objects having unusual radio or optical properties (see Sect. 3)...... 74 204 LIST OF FIGURES

6.3 The relationship between the UV color index (F UV NUV ) and the total H band luminosity. Symbols are as in fig. 6.1. The− optical spectra available for dwarf ellipticals are presented...... 77 6.4 The relationship between the UV color index (F UV NUV ) and the − total H band luminosity. Symbols are as in fig. 6.1. The optical spectra available for ellipticals are presented...... 78 6.5 The relationship between the UV color index (F UV NUV ) and the total H band luminosity. Symbols are as in fig. 6.1. The− optical spectra available for lenticulars are presented...... 79

7.1 Ratio of the total infrared to far ultraviolet luminosity as a function of the ultraviolet spectral slope (lower x-axis) and the FUV-NUV color (upper x-axis). Open circles indicates our secondary sample while filled circles represent the primary sample. The dashed line represents the best linear fit to starburst IRX-UV relation. The solid line indicates the best bisector linear fit for our primary sample. The stars indicate the sample of IUE starbursts. Mean error bars for the plotted data are shown in the lower right corner, in this and subsequent figures. The residuals from the best linear fit for normal galaxies are shown in the bottom panel...... 88 7.2 Relation between the birthrate parameter computed from the Hα emis- sion, and the distance from the L /L β relation for starbursts. T IR F UV − The solid line represents the best linear fit...... 90 7.3 The relation between the ultraviolet spectral slope β and the Hα at- tenuation obtained from the Balmer decrement. Symbols are as in Fig.7.1. Solid line represents the best linear fit to our primary sample (equation 7.14) while the dashed line indicate the best-fit for starburst galaxies obtained by Calzetti et al. (1994) (equation 7.13). Arrows in- dicate galaxies for which the value of A(Hα) is a lower limit of the real value (i.e. Hβ undetected). The residuals from the best linear fit for normal galaxies are shown in the bottom panel...... 93

7.4 Relation between gas metallicity and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig.7.1. The solid lines show the best linear fit for our primary sample. The residuals from the best linear fits for normal galaxies are shown in the upper panels...... 95

7.5 Relation between the galaxy size and the LT IR/LF UV ratio for starburst (left panel) and normal galaxies (right panel). Symbols are as in Fig. 7.1. Mean values and uncertainties in bins of 0.30 log(Diameter) are given...... 95 LIST OF FIGURES 205

7.6 Relation between the gas to dust ratio and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample...... 97

7.7 Relation between the H-band luminosity and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panel...... 98

7.8 Relation between the TIR+FUV luminosity and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1...... 98

7.9 Relation between the mean H-band surface brightness (µe) and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panel.100

7.10 Relation between the star formation rate density and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample.The residuals from the best linear fit for normal galaxies are shown in the upper panel. . 100

7.11 Relation between the Hα and far ultraviolet luminosity and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. Hα luminosity is corrected for dust attenuation using the Balmer decrement, while the FUV flux is uncorrected. The solid lines show the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panels...... 102 7.12 Relation between the observed Hα and far ultraviolet luminosity and the LT IR/LF UV ratio (left) or β (right). Symbols are as in Fig. 7.1. Hα luminosity is the observed value not corrected for dust attenuation. The solid lines show the best linear fit for our primary sample.The residuals from the best linear fit for normal galaxies are shown in the upper panels...... 102

8.1 The combined NUV and FUV image of NGC 4438. The regions de- scribed in sect. 3 of the text are labeled 1 to 7. The horizontal line is 10 kpc long (assuming a distance of 17 Mpc)...... 109 8.2 The Hα+[NII] contours (red, in arbitrary scale, in between 8 10−17 and 6 10−16 erg cm−2 s−2 arcsec−2, with σ= 5 10−17 erg cm−2 s−2 arcsec−2, from Boselli & Gavazzi (2002)) are superposed to the NUV gray-level image of NGC 4438...... 111 8.3 Chandra image of NGC 4438 in the 0.3-2 keV energy band with Hα contours superposed. Adapted from Machacek et al. (2004) ...... 111 206 LIST OF FIGURES

8.4 The RGB (FUV=blue, NUV=green, B=red) color map of NGC 4438 and NGC 4435. The SED of each region defined in Fig. 1 are given in the lower plot of each frame. Crosses indicate the observed data, arrows upper limits (in mJy), the red dashed line the evolved population fit as determined by the model of Boissier & Prantzos (2000), the dotted blue line the starburst SED (from Starburst 99) and the dashed green line the combined fitting model. The burst luminosity contribution (for the age corresponding to the minimum χ2) in the band FUV, B and K is also given. The upper panel gives the variation of the reduced χ2 parameter (black continuum line, in logarithmic scale) and of the burst mass fraction (red dotted line) as a function of the age of the burst (in Myr). The lower panel of region 4 gives the integrated 3500 to 7000 A,˚ R=1000 spectrum of the main body of the galaxy (black continuum line) compared to the fitted model (red dashed line). . . . 115

9.1 The radial profile of observed (open symbols) and extinction-corrected (filled symbols) H-band surface brightness (left) and of the rotational velocity (center) used to constrain the model without interaction (rep- resented by the black solid line). The total gas radial profile (right) predicted by the unperturbed model (solid black line) is compared to the observed one (green filled circles), obtained by summing the HI component (red line) to the molecular one (blue and light blue) and correcting for Helium contribution ( 1.4), and to the model including × the interaction (black dashed line)...... 120

9.2 Ram pressure stripping intensity (in arbitrary units) as a function of time (Eq.9.1). Adapted from Vollmer et al. (2001)...... 121

9.3 The radial profile of the observed (empty green circles) and extinction- corrected (filled green circles) total gas, Hα, FUV (1530 A),˚ NUV (2310 A),˚ B and i surface brightness. The yellow shaded area marks the range in between the observed (bottom side) and extinction-corrected (top side) surface brightness profiles. Surface brightnesses are compared to the model predictions without interaction (black solid line) or with in- teraction for several 0 and t0 parameters. Equal maximum efficiency −2 −1 (0=1.2 M kpc yr ) and different age: t0=100 Myr, red continuum line (the adopted model); t0=500 Myr, grey long dashed line, t0=1.5 Gyr, dashed magenta line. Equal age (t0=100 Myr) and different max- −2 −1 imum efficiency: 0=3 M kpc yr , blue dotted line; 0=1/3 M kpc−2 yr−1, orange dotted line...... 122 LIST OF FIGURES 207

9.4 The observed and model surface brightness (a), color (b) radial profiles of NGC 4569. In the model profiles the continuum lines are for models with gas removal, dashed lines for unperturbed models. c) the reduced χ2 as a function of the look-back time of the ram-pressure event for a −2 −1 few efficiencies 0 (M kpc yr ). Models were computed each 100 Myr for 0 < t0 < 500 Myr, 200 Myr for 1.5 < t0 < 0.5 Gyr and 1 Gyr −2 −1 for 6.5 < t0 < 1.5 Gyr; and each 0.2 M kpc yr efficiencies between 0.4 and 1.6 (only the more relevant are shown here). d) the variation of the effective surface brightness (mean surface brightness within Re, the radius containing half of the total light) and radius due to differential variation of the star formation history of NGC 4569. Open triangles are for the unperturbed model, the other symbols for different ages of the interaction (100 Myr, 1.5 and 5.5 Gyr)...... 125 9.5 The RGB (continuum subtracted Hα =blue, FUV=green, NUV=red) color map of NGC 4569 ...... 126

10.1 The four Arecibo HI pointings obtained in the region of the BIG group, superposed to the r0 band image. The size of each circle correspond to the telescope beam...... 129 10.2 GALEX NUV image of the Blue Infalling group (BIG)...... 130 10.3 High-contrast Hα+[NII] band frame of the BIG group...... 132 10.4 Upper panel: The position and the width (rectangular areas on the right) of the three slits obtained for CGCG97-125. The slits are super- posed to the Hα + [NII] net image. Lower Panel: The three different rotations curves obtained for CGCG97-125. Letters indicate the dif- ferent regions as labeled in the upper panel...... 136 10.5 The low resolution 2D spectrum obtained at ESO/3.6 for the knots DW3d (left) and DW3e (right), shows a significant difference ( 500km s−1) in the velocity of the two knots...... ∼. . . . . 137 10.6 Stellar shells are seen around galaxy 97-125 in the r0 band image of BIG. No continuum emission is detected from the low brightness trails (except K2)...... 139 10.7 Extended low brightness trails appear in the Hα+[NII] NET frame of BIG...... 140 10.8 HI column density distribution in BIG. Contours are 0.5, 1.0, 2.0, 3.0, 4.0, 5.0, and 6.020cm−2. Adapted from Sakai et al. (2002) ...... 144 10.9 HI position-velocity diagram centered on CGCG 97-125. Adapted from Sakai et al. (2002) ...... 144 10.10The HI spectra obtained for each pointing...... 146 208 LIST OF FIGURES

10.11Comparison between the combined HI spectrum obtained from the four different Arecibo pointings, and the single pointing on the NW trail. It appears clearly the presence of a low velocity component not associated to the bright galaxies in BIG...... 147 10.12The relation between Metallicity and B-band Luminosity (with linear best-fit) for galaxies in nearby clusters (empty circles, adapted from Gavazzi et al. 2004). The triangles mark the mean metallicity obtained for the individual knots of BIG...... 151 10.13Comparison between the drift-scan integrated (blue) and nuclear (red) spectrum of CGCG97-125...... 153 10.14The SED of CGCG97-125, corrected for internal extinction. Nuclear and drift-scan integrated spectra are shown in green. Black circles in- dicate photometric observations and their relative uncertainties. Best fitting models for the nuclear spectrum (red) and for the starburst com- ponent (blue) are given. The resulting best fitting SED for CGCG97- 125 is presented in black...... 155 10.15The star formation history of CGCG97-125 as obtained from the SED fitting procedure...... 155 10.16The 2D high resolution spectrum (left) and the optical rotation curve (right) of CGCG97-120 ...... 156 10.17B-R color map of BIG (Blue = B; Red = R)...... 160 10.18The observed smoothed (step 3) one dimensional spectra. The object identification and telescope are labeled on each panel...... 161 10.18Continue...... 162 10.18Continue...... 163

11.1 The distribution of the individual HαE.W. measurements in the Virgo cluster along the Hubble sequence (small dots) and of the median EW(Hα) in bins of Hubble type. Error bars are drawn at the 25th and 75th percentile of the distribution.Filled symbols represent HI-def< 0.4 (unperturbed) objects and open symbols HI-def> 0.4 (HI deficient) galaxies...... 170 11.2 The star formation rate as a function of density, comparing groups of galaxies with clusters. The upper and lower horizontal dashed lines show the 75% percentile and the median of the equivalent widths. The hashed region shows the relation for the complete sample, while the solid line shows the relation for systems with 500 km s−1 < σ < 1000 km s−1 (left) and σ < 500 km s−1 (right). The dependence on local density is identical irrespective of the velocity dispersion of the whole system. Figure taken from Bower & Balogh (2004)...... 171 LIST OF FIGURES 209

11.3 The ratio of the isophotal Hα and r0 radii as a function of the HI deficiency for galaxies in the Virgo cluster...... 172 11.4 The clustercentric radial distribution of the individual EW(Hα) mea- surements in the Virgo cluster. High and low (B-band) luminosity galaxies are given with open and filled dots respectively. Median in th th bins of 0.5 R/RV ir are given. Error bars mark the 25 and 75 per- centile of the distribution...... 173

List of Tables

2.1 Selected Performance Parameters (Morrissey et al. 2005) ...... 17

3.1 Integral redshift completeness in bin of 0.5 magnitudes...... 26 3.2 The completeness-corrected differential number of galaxies per bin of magnitude ...... 30

4.1 Best Fitting Parameters...... 37

5.1 The spectrograph characteristics ...... 45 5.2 1D substructure indicators for the whole cluster sample ...... 48 5.3 The most significant weighted gaps detected in the velocity distribution of the whole cluster sample...... 49 5.4 3D substructure indicators for our sample ...... 50 5.5 Mass estimate for Abell 1367 ...... 61 5.6 Two-body model parameters ...... 65 5.7 The 119 new redshift measurements ...... 67 5.7 Continue ...... 68

6.1 Main relations for early type galaxies ...... 73

7.1 Linear realtions useful to estimate the LT IR/LF UV ratio (log(LT IR/LF UV ) = a x + b)...... 105 × 10.1 Redshifts of the galaxies in the BIG group...... 133 10.2 Line fluxes, corrected for internal extinction, of the galaxies in the BIG group...... 135 10.3 Properties of galaxies in BIG...... 138 10.4 Metallicities of the galaxies in the BIG group...... 150 10.5 Best-fitting parameters for the nuclear and starburst component of CGCG97125...... 153

211