Active Galactic Nuclei at hard X-ray energies: absorption, reflection and the unified model

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Thesis

Reference

Active Galactic Nuclei at hard X-ray energies: absorption, reflection and the unified model

RICCI, Claudio

Abstract

Active Galactic Nuclei (AGN) are the most luminous persistent sources in the Universe, and are believed to be powered by accretion onto supermassive black holes. My work is mainly focused on the study of X-ray spectra of AGN, with particular attention to the unification scenarios. AGN are classified according to their optical spectra, and different types of AGN are thought to be intrinsically the same object, just observed from different lines of sight with respect to a molecular torus. In this work, differences between different classes of AGN, unforeseen by the unification model, are found. These differences appear to be due to a larger reflection component in the spectra of more obscured objects, which could be related to the latter having on average more material around than the former.

RICCI, Claudio. Active Galactic Nuclei at hard X-ray energies: absorption, reflection and the unified model. Thèse de doctorat : Univ. Genève, 2011, no. Sc. 4386

URN : urn:nbn:ch:unige-194749

DOI : 10.13097/archive-ouverte/unige:19474

Available at:

http://archive-ouverte.unige.ch/unige:19474

Disclaimer: layout of this document may differ from the published version.

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Département d’astronomie Professeur T. J.-L. Courvoisier

Active Galactic Nuclei at hard X-ray energies:

Absorption, Reflection and the Unified Model

Thèse

présentée à la Faculté des sciences de l’Université de Genève pour obtenir le grade de Docteur ès sciences,

mention Astronomie et Astrophysique

par

Claudio Ricci

de Rome (Italie)

Thèse N^o 4386

Genève

Atelier d’impression ReproMail 2011

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voirAppendix A

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獻給妳, 我的月亮

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They laughed at me as "Prof. Moon,"

As a boy in Spoon River, born with the thirst Of knowing about the stars.

They jeered when I spoke of the lunar mountains, And the thrilling heat and cold,

And the ebon valleys by silver peaks, And Spica quadrillions of miles away, And the littleness of man.

But now that my grave is honored, friends, Let it not be because I taught

The lore of the stars in Knox College, But rather for this: that through the stars I preached the greatness of man,

Who is none the less a part of the scheme of things For the distance of Spica or the Spiral Nebulae;

Nor any the less a part of the question Of what drama means.

Edgar Lee Masters, Spoon River Anthology

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Reśumé

Le travail présenté dans cette thèse a été eﬀectué de 2008 à 2011 au sein de l’ISDC Data Centre for Astrophysics, rattaché à l’observatoire de Genève. La tâche principale de l’ISDC est de collecter, d’analyser et d’archiver les données du satellite d’observation X et gamma INTEGRALde l’Agence Spatiale Européenne (ESA).INTEGRALfût lancé le 17 Octobre 2002 et a pour mission l’observation du ciel dans le domaine des rayons X et gamma, entre 5 keV et 10 MeV, gâce à quatre instruments:

l’imageur en rayons X durs et gamma IBIS, le spectromètre SPI, les deux moniteurs de rayons X JEM-X1 and JEM-X2 et le moniteur optique OMC.

Le thème principal de ma thèse est l’étude des noyaux actifs de galaxies (AGN) dans la bande d’énergie X et gamma. Les AGN sont les sources persistantes les plus lumineuses de l’Univers et sont supposés être le principal responsable du fond cosmique X. Les AGN sont alimentés par l’accrétion de matière par un trou noir supermassif sous forme d’un disque d’accrétion. Les galaxies de Seyfert sont le type d’AGN le plus courant dans l’Univers local et sont classées en type 1 (Seyfert 1s) ou type 2 (Seyfert 2s) selon que leur spectre montre respectivement à la fois des raies larges et fines ou seulement des raies fines. Des objets ayant des caractéristiques intermédiaires sont nommées Seyfert 1.5s. Selon le paradigme généralement reconnu du modèle unifié des AGN, le trou noir supermassif est entouré d’un absorbant anisotrope tel qu’un tore moléculaire, et la seule diﬀérence entre les deux classes est la ligne de visée à l’égard de l’absorbant. Dans ce schéma, la région proche du trou noir supermassif où les raies d’émission larges sont produites (appelée Broad Line Region) est obscurcie par le tore dans les Seyfert 2s puisque nous observons le système par la tranche. On ne peut donc voir que la région de raies étroites plus éloignée du moteur central. Les Seyfert 1s sont plutôt observées des pôles, ce qui permet de voir à la fois les régions de raies étroites et larges. Une classe particulière d’AGN sont les Seyfert 1s à raies étroites (Narrow line Seyfert 1s ou NLS1s), qui sont supposées héberger un trou noir de plus faible masse et qui sont peut-être au début de leur vie.

La bande d’énergie X est le domaine le plus approprié du spectre électromagnétique pour la dé- tection des trous noirs. L’émission de rayons X des AGN est supposée être produite par diffusion Compton inverse des photons UV et optiques générés dans le disque d’accrétion sur une couronne d’électrons très chauds. L’étude de l’émission X des AGN est fondamentale pour comprendre la structure des régions émettrices et absorbantes car elle nous permet d’observer les zones proches de la source centrale. L’absorption affecte les rayons X via deux processus: la diffusion Compton et l’absorption photoélectrique. L’absorption photoélectrique est fortement dépendante de l’énergie mais son influence est négligeable au-delà de 10 keV. La diffusion Compton est indépendante de l’énergie jusqu’à plusieurs centaines de keV mais son effet est significatif seulement si la source est très ab- sorbée. Les spectres X des AGN sont en principe bien représentés par un continu en loi de puissance avec une coupure autour de quelques centaines de keV et une bosse de réflexion aux environs de 30 keV.

Au cours de ma thèse, j’ai travaillé sur de nombreuses questions liées à l’émission de rayons X et gamma des AGN, la plupart du temps en utilisant les données obtenues par l’instrument IBIS à bord du satellite INTEGRALmais aussi par d’autres instruments tels que les télescopes X mous à bord du satelliteXMM-Newton ou le télescope X durs deSwift/BAT.

J’ai collaboré à l’étude du second catalogue d’AGN d’INTEGRAL qui répertorie tous les AGN observés par INTEGRAL à partir de fin 2008 et qui représentait à l’époque la plus grande étude des caractéristiques des AGN en rayons X durs jamais réalisée. J’ai analysé l’émission de rayons X

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mous (0.1–10 keV) d’environ 35 objets pour lesquels aucune information sur la densité de colonne (le paramètre habituellement utilisé pour décrire la quantité d’absorption le long de la ligne de visée) n’était disponible dans la littérature. Cette analyse a été faite en utilisant les données obtenues par Swift/XRT etXMM-Newton/EPIC. Parmi les sources étudiées, trois montraient des caractéristiques particulières dans les X mous. Ces sources sont les suivantes: les galaxies de Seyfert 1s ESO 140-43 et UGC 3142 et la galaxie de Seyfert 2 ESO 383-18. En étudiant de nouvelles données haute qualité de XMM-Newton/EPIC combinées aux données en rayons X durs obtenues parINTEGRALIBIS/ISGRI et Swift/BAT (14–195 keV), nous avons contraint les caractéristiques des nuages absorbants. Nous avons trouvé que les spectres X des sources UGC 3142 et ESO 383-18 ne peuvent pas être décrits par un absorbant unique et homogène. Ils nécessitent tous deux au moins deux absorbants neutres qui couvrent partiellement la source. Le spectre X de la source ESO 140-43 faisait apparaître des signes d’un système d’absorption très complexe composé de trois diﬀérentes couches de matière ionisée avec des états d’ionisation, des densités et des facteurs de couverture diﬀérents. Une des caractéristiques les plus frappantes d’ESO 140-43 est sa variabilité. La source montre une forte variation de flux sur une échelle de temps de six mois entre les deux observations réalisées par XMM-Newtonavec un change- ment remarquable des caractéristiques physiques et du facteur de couverture des trois absorbants.

Sur des délais plus courts, la source présente également une certaine variabilité que nous supposons liée aux changements des facteurs de couverture de l’un des absorbants. Alors que les absorbants dans ESO 383-18 et UGC 3142 sont probablement liés à la présence de bandes de poussière, les absorbants ionisés font certainement partie d’un vent. L’étude des vents dans les galaxies de Seyfert est cruciale pour la compréhension des rétroactions entre l’AGN et la galaxie hôte.

Les rayons X durs représentent une bande d’énergie unique pour tester le modèle unifié des AGN.

Le rayonnement dans cette gamme d’énergie n’est en fait pas affecté par la matière absorbante le long de la ligne de visée, ce qui permet à l’observateur de voir la source centrale. Ainsi, si le modèle unifié des AGN est correct et la seule différence entre les classes est l’absorption, on pourrait s’attendre à voir la même émission X en moyenne des différentes classes d’objets. A partir de toutes les données publiques disponibles réalisées par INTEGRAL IBIS/ISGRI sur les 170 galaxies de Seyfert environ, situées à des redshiftsz <0.2et détectées dans les X durs, nous avons mené l’étude la plus complète du spectre moyen des AGN dans les X durs jamais réalisée. Notre échantillon consistait en 44 Seyfert 1s, 29 Seyfert 1.5s, 78 Seyfert 2s et 14 NLS1s. Afin d’éviter tout effet lié à la présence de gaz obscurcis- sant, nous avons divisé les Seyfert 2s de notre échantillon en deux groupes: Seyfert 2s Compton-thin (NH < 10²⁴cm⁻²) et Seyfert 2s Compton-thick (NH ≥ 10²⁴cm⁻²). Le premier regroupe 68 objets et le second en compte seulement 10. Nous avons produit les spectres moyens des différentes classes d’AGN, incluant toutes les données disponibles pour chacune des sources de l’échantillon. A partir d’une analyse spectrale à la fois dépendante et indépendante du modèle, nous avons trouvé que les spectres moyens des Seyfert 1s et Seyfert 1.5s sont conformes, comme prévu par le modèle unifié. Par ailleurs, le spectre X durs moyen des Seyfert 2s montre une plus forte composante de réflexion que celle des Seyfert 1s et Seyfert 1.5s. Nous avons constaté que cette différence est due à la présence de Seyfert 2s que nous avons définies comme "mildly obscured" (MOB, 10²³ ≤N_H< 10²⁴cm⁻²) alors que les Seyfert 2s qualifiées de "lightly obscured" (LOB, N_H <10²³cm⁻²) ont des caractéristiques plus proches de celles des Seyfert 1s et des Seyfert 1.5s. La réflexion la plus importante que nous ayons trouvé dans les Seyfert 2s ne peut être facilement expliquée par le modèle unifié mais pourrait impliquer un plus grand facteur de couverture de la matière environnante ou l’existence de grumeaux denses dans la ligne de visée. Cependant, une fois que l’on a pris en compte une plus importante réflexion pour les Seyfert 2s MOB, les paramètres de l’émission primaire (l’indice de la loi de puissance

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et l’énergie de coupure) sont compatibles, ce qui confirme l’idée générale que le moteur central est le même pour tous les types d’objets.

Le télescope LAT (Large Area Telescope) à bord du satellite Fermi(une mission de la NASA) est un détecteur de rayons gamma sensible à la bande d’énergie 20 MeV–300 GeV. Le ciel observé dans le domaine des rayons gamma est connu pour être principalement peuplé de blazars et de quelques radio- galaxies très proches, ainsi l’une des découvertes inattendues du Fermi/LAT depuis son lancement en Juin 2008 a été la détection du rayonnement gamma émis par les NLS1s. Motivés par cela, nous avons analysé environ 1.6 année de données et avons cherché l’émission à haute énergie provenant des plus brillantes galaxies de Seyfert. Nous avons trouvé la première preuve de l’émission gamma (avec une détection significative d’environ 8σ) provenant de la Seyfert 2 Compton-thick NGC 1068, et nous avons étudié les nouvelles données de la galaxie à sursaut stellaire NGC 4945 dont la détection a déjà été reportée dans le catalogue de 11 mois. Nous avons étudié et modélisé la distribution spectrale en énergie multi longueur d’onde (du domaine radio aux rayons gamma) de ces deux sources et nous avons conclu que les responsables de l’émission gamma sont diﬀérents pour les deux objets. Dans le cas de NGC 4945, l’émission est dominée par l’activité stellaire, de violents processus de formation d’étoiles qui accélèrent les particules à de très hautes énergies. Les particules interagissent ensuite avec le milieu interstellaire, produisant des photons gamma. Dans le cas de NGC 1068, la luminosité gamma est trop élevée pour être expliquée par l’activité stellaire. Nous avons proposé un modèle dans lequel les photons gamma seraient produits par des processus Compton inverse dans le jet.

INTEGRALa pu apporter d’importantes informations pour la compréhension de la géométrie et des mécanismes physiques des AGN. Les résultats obtenus dans ma thèse, en particulier la découverte d’une plus forte composante de réflexion dans les Seyfert 2s, sont important notamment dans l’optique de la prochaine gènèration de tèlescopes à rayons X (Astro-H, NuSTAR, Athena), qui permettra de mieux comprendre la structure des AGN, et leur contribution au fond cosmique X.

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Acknowledgments

Starting to write this section has proved to be one of the most diﬃcult things during the last months of my PhD. Not because I wouldn’t know who to acknowledge, but rather because words are a re- ductive way to present my thanks to all the people who have helped me during my path.

There are many people I would like to thank, and without whom this work would not have been possible. First of all my I want to thank Prof. Courvoisier for having given me the amazing op- portunity to do my PhD in the ISDC. Thank you Thierry also for all your advices, for having me showed the importance of thinking out of the box, and for having given me the possibility of attending many conferences and schools which have surely contributed to my scientific enrichment. My most sincere thanks also to all the people who followed my progresses (and sometimes my non-progresses) during the years. Volker, you were first person I worked with at ISDC, the first to introduce me to the wonders and mysteries of AGN. Thank you for all your help and your patience, both during my diploma’s thesis and then during the first part of my PhD. Thank you Roland for all your suggestions and help, and for sharing your broad knowledge ofINTEGRALwith me. Thank you Stephane for all your precious help during the last year of my thesis. We started working together relatively late in my PhD, but your help and support in the final stages of my thesis were priceless, I had the luck to learn a lot from you, and without your contribution this manuscript would not have been the same.

Thank you for having your door always open, for all our extremely interesting scientific discussions, and for all your comments on my thesis.

My heartfelt thanks also to all the people of ISDC. First of all to Carla, for being the perfect oﬃcemate, always ready for a chat, to share a laugh, or to hear my complaints during the bad times.

It was great to share the office with you. To Simona, for making me feel at home at ISDC since day 1. Thank you for all your help and for all the parties you and Volker organized, ISDC was a quieter place after you left. Thanks to Laetitia for always being so cheerful, to Dominique and Carlo for the occasional skiing, to Rozenn for her help with the glossary, to Andres for all his crazy stuff, to Nicolas M. for always stopping by our office, to Franca, Marie-Claude and Martine for all the chats and for all the help they provided me during the years. Thanks to all the coffee break/lunch group for all the nice discussions we had over the years: Lucia, Matteo, Pierre, Maria, Nicolas P., Nicolas F., Philippe A., Andrew, Marc A., Ben, Mark G., Reiner and Masa. Thanks also to Andrea, Marc T. and Enrico for all their suggestions and help.

I also want to thank all my friends who made my "outside" life great during the last years. Thanks to the Zio Gio for all the trips and the fun and the crazy stuﬀ, to Chiara, Alice, Giulia, Ale, Leo, Federico, Sara, Luisa, Olivia, Marco and Andrea for all the nights out, we always had a lot of fun, and I’m sure I will miss you a lot after leaving Geneva. Thank you my friends, you made these last years really unforgettable. Thank you to Stefi and Suus for always being such good friends, regardless of the distance. Thanks also to my crazy Israeli friends Benny and Rivay, after we met the first time in India many years ago we have managed to meet every time in diﬀerent places across the world, it was always fun to have you guys around!

Grazie alla mia famiglia per tutto l’aﬀetto ed il supporto dimostratomi durante tutti questi anni.

Grazie per avermi dato la possibilitá di seguire i miei sogni e di avermi sempre incoraggiato, non sarei

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mai riuscito ad arrivare qui senza il vostro aiuto, dal primo giorno di scuola, piu di venti anni fa, sino ad oggi. Grazie papá per essere stato con la tua curiositá sempre un’ispirazione, grazie mamma per avermi sempre spronato a migliorare, e grazie Flavia e Silvia per essermi state sempre vicine, seppur lontane fisicamente. Grazie nonna per tutti gli anni che abbiamo passato insieme e per tutto il tuo aﬀetto.

And finally, I would like to thank the person that in only a few years changed completely my life.

Chin Shin, you know that I would have never been able to be here today without you. Thank you for all your love and your encouragements during these years, your support has been invaluable, and I look forward to our future together.

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Active Galactic Nuclei

Contents

1.1 An historical introduction . . . . 1

1.2 Anatomy of an AGN . . . . 2

1.2.1 The Supermassive Black Hole . . . . 2

1.2.2 The Accretion Disk. . . . 5

1.2.3 The Broad-Line Region . . . . 9

1.2.4 The Narrow-Line Region. . . 10

1.2.5 The Corona . . . 13

1.2.6 The Jets. . . 14

1.3 AGN classification . . . 15

1.3.1 Radio-quiet AGN. . . 16

1.3.2 Radio-loud AGN . . . 17

1.4 The Unified Model . . . 19

1.4.1 Radio-quiet unification. . . 20

1.4.2 Radio-loud unification . . . 20

1.5 A panchromatic view of an AGN . . . 21

1.5.1 Radio emission . . . 21

1.5.2 Infrared emission . . . 21

1.5.3 Optical-UV emission . . . 22

1.5.4 Indicators of bolometric luminosity . . . 22

1.1 An historical introduction

The history of Active Galactic Nuclei (AGN) starts at the beginning of the last century. In 1909 Fathundertook a series of observations aimed at understanding the nature of "spiral nebulae", some of the most enigmatic sources at the time. Scientists were in fact struggling to understand whether these objects were nearby gaseous objects similar to the Orion nebula, or very distant collections of unresolved stars. For most objects Fath found continuous spectra with stellar absorption lines, which indicated emission from unresolved solar-type stars. However, for one object, NGC 1068, he found a peculiar spectrum, showing both bright emission and absorption lines, similar to what had been observed in gaseous nebulae. This result was later confirmed by Slipher(1917) and by Hubble (1926), who found evidence of this behavior also in the spectra of NGC 4051 and NGC 4151. The first systematic study of galaxies showing nuclear emission lines was performed by Seyfert in 1943, who studied the spectra of NGC 1068, NGC 1275, NGC 3516, NGC 4051, NGC 4151 and NGC 7469.

He found that while objects like NGC 1068 had forbidden and permitted lines with roughly similar

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profiles and widths of∼3000 km s⁻¹, objects like NGC 4151 showed narrow forbidden lines and very broad (∼7500 km s⁻¹) hydrogen lines. In the following years the study of AGN developed mostly in the radio, an energy band that had started to be explored by astronomers thanks to the pioneeristic work of Jansky, and to the technological development of radio instruments. As often happens, it was for military reasons that the radio technology was boosted in the early 1940’s. Thus at the end of the World War II several groups of radio engineers turned their eﬀorts to radio astronomy. This led to the discovery of emission from a discrete source in Cygnus, which was then called Cygnus A (Hey et al. 1946), and later from several other sources (Bolton & Stanley 1948). In the following years many more radio sources were discovered and their positions accurately estimated, thanks to surveys like the third Cambridge survey at 159 MHz (Edge et al. 1959). The breakthrough came a few years later from observations of the very bright quasar 3C 273. Thanks to radio observations performed during a lunar occultation, Hazard et al. (1963) were able to define the position of the source with an accuracy of about 1 arcsec. In the same year, based on optical observations Schmidt (1963) discovered the presence of redshifted emission lines, which were almost immediately attributed as being due to Hubble expansion (Greenstein & Schmidt 1964). The luminosity obtained and the cosmological distance excluded the possibility that these peculiar objects were very variable stars.

Few years later Sandage et al.(1965) reported the discovery of a large number of radio quiet objects that resembled quasars. These objects had been found as a characteristic population examining color-color (U−B, B−V) diagrams of stars. They showed in fact a strong "ultraviolet excess", which is now well known to be one of the most important characteristics of AGN. Since then a large number of AGN were discovered and classified according to their main properties. Today we know that many galaxies harbor AGN, radio-quiet objects, as the galaxies discovered by Carl Seyfert (and today called Seyfert galaxies) can usually be found in spiral galaxies, while radio-loud objects (as 3C 273) are usually hosted in elliptical galaxies.

In the following I will first illustrate the several components that constitute an AGN (Sect.1.2), then I will introduce the classification (Sect.1.3) and the unified model (Sect.1.4) for both radio-loud and radio-quiet objects, and I will present how AGN look at diﬀerent wavelengths (Sect.1.5).

1.2 Anatomy of an AGN

AGN are thought to be constituted by several components, all of them very likely intrinsically related to the supermassive black hole lying at their center, and driving their growth and evolution. In the following I will discuss the most important components normally found in AGN: the supermassive black hole (Sect.1.2.1), the accretion disk (Sect.1.2.2), the broad line region (Sect.1.2.3), the narrow line region (Sect.1.2.4), the corona (Sect.1.2.5) and the jets (Sect.1.2.6).

1.2.1 The Supermassive Black Hole

The existence of black holes was predicted byEinsteinin his formulation of General Relativity in 1916.

Already a century before first Mitchell and then Laplace considered the idea of objects so dense that light cannot escape them. Although much theoretical work was done in the years following Einstein’s work, it was not until 1972 that observational evidence for the existence of black holes was found through observations of Cyg X-1 (Bolton 1972,Webster & Murdin 1972).

The presence of supermassive black holes (SMBHs) in the center of AGN was first predicted by Lynden-Bell (1969), and later largely confirmed by several observational evidences. From the short time scales variations (�1 hour) observed in AGN it is possible to obtain an upper limit on the size

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of the central source, which together with mass estimates point towards very high densities, of the order of magnitude of those expected in black holes. A rapidly rotating disk of ionized gas has been detected by HST observations of M87, the velocities measured are in agreement with a disk being Keplerian rotating around a black hole (Ford et al. 1994). Moreover, the luminosities observed in AGN can be easily explained by accretion onto a supermassive black hole (see Sect.1.2.2), and no serious alternative explanation to this mechanism exist, although theories involving starburst have long been debated (e.g.,Terlevich et al. 1992).

The origin of the supermassive black holes hidden at the center of galaxies is still unknown, several scenarios have been proposed (seeDokuchaev et al. 2007 for a review on the topic), amongst which the collapse of supermassive population III stars or the collapse of massive primordial clouds.

Rotating and non-rotating black holes

The role of the surface for a black hole of massM is played by the sphere with a radius corresponding to the Schwarzschild radius (r_g = 2GM/c²). The Schwarzschild radius is the distance from the singularity at which a body cannot escape anymore from the gravitational attraction and is bound to fall inside the black hole. The metric of space-time around a non-rotating black hole is given by theSchwarzschildsolution (1916) to Einstein’s general relativity:

ds² =

�

1−2GM rc²

�

dt²− 1 c²

� dr²

1−^2GM_rc² +r²(dθ²+ sin²θdφ²)

� .

From the metric one can deduce the last stable circular orbit around a non-rotating black hole:

r= 3r_g. So for r <3r_g no stable orbit exist, and forr < r_g particles rapidly fall in the singularity.

The general solution for a rotating black hole was discovered in 1963 by Kerr. The metric for a black hole with an angular momentum J is

ds² =

�

1−2GM r ρc²

�

dt²− 1 c²

�4GM rasin²θ

ρc dtdφ+ ρ

∆dr²

�

−1 c²

�

ρdθ²+ (r²+a²+2GM ra²sin²θ

ρc² ) sin²θdφ²

�

In this formalism a is the angular momentum of the black hole per unit mass (a= (J/M c)) and has the dimension of a distance,∆ =r²−(2GM r/c²) +a² andρ²=r²+a²cos²θ. If the black hole is non-rotating, J =a= 0, and the Kerr metric reduces to the standard Schwarzschild metric. The radius of the horizon for a rotating black hole is:

r₊ = GM

c² + [(GM

c² )²−( J

M c)²]^1/2 (1.1)

The maximum value of the angular momentum a black hole can have is J =GM²/c, and the radius of the latest circular orbit depends on whether the particle is co-rotating or counter-rotating with respect to the black hole. For a maximally rotating black hole this values isr =GM/c² for co-rotating particles, andr = 9GM/c² for counter-rotating ones.

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Figure 1.1: Left panel: correlation between the black hole mass M_BH and the blue luminosity of the galaxy spheroidal component. Right panel: correlation between the central velocity dispersion σ of the host galaxy andM_BH. Filled symbols show ellipticals galaxies, while open symbols show spiral and lenticular galaxies (Ferrarese & Ford 2005).

Measuring black hole masses

One of the most common ways to measure the mass of the black holes (M_BH) in AGN is through the width of their broad emission lines. These lines are emitted from gas orbiting around the central black hole and have normally a velocity broadening of v∼10⁴km s⁻¹. Using the following relation

v²

r = GM_BH r² ,

and considering that the emission region has a size of r ∼ 0.01pc, one obtains a mass of M_BH ∼ 10⁸M_⊙. One of the methods most used to estimate the black hole masses of broad line AGN is the

"Reverberation Mapping" (e.g., Kaspi et al. 2000). This technique exploits the fact that emission lines are lagged with respect to the continuum, from which one can calculate r, the distance of the broad emission line region to the central black hole. Vestergaard & Peterson (2006) estimater from reverberation mapping of the Hβ line and deriveM_BH from

MBH

10⁶M_⊙ = 8.3

�FWHM(H_β) 10³km s⁻¹

�2�

λLλ(5100) 10⁴⁴erg s⁻¹

� .

Due to observational limits, this relation can be used only for objects up to z ∼ 0.75. Above this redshift the line is present in some specific near infrared band. Other emission lines as Mg II and C IV are used to bypass the limits of Hβ.

Several other methods are used to establish the black hole mass. One of the most precise ones is through optical observations of the bulges of the host galaxies (e.g., Magorrian et al. 1998). Water maser emission in circumnuclear disks is often used for obscured AGN. Using the acceleration of masers perpendicular to the line of sight and the distance and speed of masers moving along the line of sight on the outskirts of the disk, it is possible to determine M_BH very accurately (e.g., Moran et al. 1995). Many more indirect methods are also often found in literature. Most of them are based on the empirical relation between the black hole mass and the stellar velocity dispersion σ_s (see

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Fig.1.1 and Sect.5.1). These methods use either direct measurements of σs (Ferrarese & Merritt 2000), or indirect estimates obtained from the width of the [O III] line (Greene & Ho 2005), or from the morphological parameters of the bulge (e.g., O’Dowd et al. 2002). Some estimates use instead the bulge luminosity (Wandel 2002), the K-band stellar magnitude (assuming that it is dominated by the bulge, Novak et al. 2006), the X-ray variability time scales (Gierliński et al. 2008), or the properties of outflowing warm absorber clouds (Morales & Fabian 2002). As we will see in Sect.4, a precise measurements is extremely important to understand the role played by the Eddington ratio (see below for the definition) on the X-ray emission of AGN.

1.2.2 The Accretion Disk

Before the 1970s the importance of accretion in astrophysics was only marginally understood. With the discovery of the first X-ray binaries it was soon realized that the only mechanism that could produce the luminosities observed in many of these sources was accretion onto compact objects.

To have a qualitative idea of the amount of energy produced by accretion processes, one can consider the gravitational potential energy released by the free fall of a mass m on a body of mass M and radius R:

∆E_acc= GM m R

where G is the gravitational constant. To compare it with nuclear processes, the maximum amount of energy per gram released by nuclear fusion of hydrogen is

∆E_nucl= 0.007mc²

where c is the speed of light. Considering for example the values typical for a neutron star (radius of R ∼ 10km and M ∼ M_⊙), the quantity of energy released by accretion would be larger than that of nuclear fusion. The energy released ∆Eacc depends on how compact the accreting object is (∆E_acc ∝M/R), thus one can expect that black holes are the sources that exploit best this source of energy. The luminosity of an accreting black hole can be expressed as

L_acc∝ηmc˙ ²

whereη is the eﬃciency of the conversion of rest mass energy of the accreting matter into radiation.

The maximum binding energy of a particle rotating around a black hole is the amount of energy which has to be lost in order that the material attains the last bound stable circular orbit. In the case of a non-rotating spherically symmetric black hole, the last stable orbit isr_I = 3r_g = 6GM/c², which results in a binding energy that is ∼6% of the mass energy of the particle. In the case of rotating black holes, for corotating particles the last stable orbit is rI = GM/c², while for counter-rotating particles is r_I = 9GM/c². The maximum binding energies of these orbits is (1−1/√

3) of the rest mass energy of the orbiting material in the corotating case and (1−�

25/27) in the counter-rotating case. The corotating case is of the greatest interest because it implies that up to 42.3% of the rest mass energy of the material can be released as it spirals into the black hole through a sequence of almost circular equatorial orbits.

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Eddington luminosity

Accretion is limited by the radiation pressure. At high accretion rates the luminosity increases and so does the outward force exerted by the radiation (F_rad). On a free electron this force is

F_rad = L_accσ_T 4πr²c ,

where σ_T = 6.65×10⁻²⁵cm² is the Thomson cross-section. The radiation pressure balances the gravitational force (F_G=GM m_p/r², wherem_p is the mass of the proton) at the so called Eddington luminosity. For a fully ionized medium in a spherically symmetric geometry such a luminosity is

L_Edd= 4πGM mpc

σ_T �1.3×10³⁸ M M_⊙

ergs s .

The Eddington luminosity can be written, in terms of accretion rate as L_Edd=ηM˙_Eddc². Thus, the accretion rate of a disk radiating at the Eddington luminosity is

M˙_Edd= L_Edd

ηc² = 4πGM m_pc

ηcσ_T ≈1.3×10⁻⁸

� M M_⊙

� �M_⊙ yr

� .

The ratio between the bolometric luminosity of an object (see Sect.1.5.4) and its Eddington limit is called the Eddington ratio:

λ_Edd= L_Bol L_Edd.

Objects withλ_Edd<1 are accreting in the sub-Eddington regime, while those with λ_Edd>1in the super-Eddington one. Although this derivation of the Eddington luminosity is limited by the spherical symmetry approximation, which clearly does not represent the physical situation of an accreting disk, it still provides an useful approximation of the limit luminosity for an accreting system. Another factor aﬀecting the eﬀective Eddington luminosity is the composition of the gas. The discussion above is valid for a fully ionized medium, and it has been shown (e.g., Murray et al. 2005) that L_Edd can be much lower for dusty gas (see also Sect.4.5).

Bondi accretion

Bondi accretion (Bondi 1952) considers spherical accretion of interstellar gas onto a compact object, considering the maximum possible rate of gravitational capture. The Bondi accretion rate is calculated using the distance from the accreting object (here we consider a black hole) at which its gravitational influence equals the thermal velocity of the gas:

M˙_Bondi= 4πR²_Aρ_Ac_s,

where ρ_A is the density of the gas at R_A ∼ GM/c²_s, and cs the sound speed. Assuming that the hot X-ray emitting gas is the supply for the central black hole, it has been possible, thanks to high resolution observations performed byChandra, to resolve the Bondi accretion radius in nearby AGN.

A very recent example is NGC 3115, for whichWong et al.(2011) claim to have resolved the accretion flow within the Bondi radius (Fig.1.2).

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25 10:05:20 15 10 05

41:0042:0043:00-7:44:0045:0046:00

Right ascension

Declination

Figure 1.2: Smoothed 0.3–6.0 keV Chandra image of NGC 3115 (Wong et al. 2011). No point source is detected at the center, but rather a plateau in the diﬀuse X-ray surface brightness.

Shakura-Sunyaev accretion disks

Matter falling onto a compact object from infinity acquires kinetic energy as its gravitational potential energy decreases. In order to conserve the angular momentum the matter cannot fall directly into the compact object, but it is commonly believed to form an accretion disk. The simplest case of accretion disks are the thin accretion disks described byShakura & Sunyaev(1973). In this picture, the matter is supposed to form a geometrically thin and optically thick disk, and follows Keplerian orbits at any radius (v∼(M_BH/R)^1/2). Angular momentum is transported outwards due to viscosity produced by turbulent eﬀects and magnetic instabilities (Frank et al. 1992). The heat energy produced in the disk is then radiated in the form of multi-temperature blackbodies, each of them having a temperature

T(R) =T_∗

�R R_∗

�₋3/4

,

whereR is the radius at which is emitted, R_∗ is the radius of a body of mass M_∗ accreting at a rate M˙, andT_∗ is given by

T_∗ ∝

�M_∗M˙ R³_∗

�¹₄ .

Between frequencies corresponding to the minimum (r_I) and maximum (r_max) distance from the accreting object, the spectrum of the accretion disk is given by:

I(ν)∝ν^1/3.

At frequencies less than that corresponding to the temperature of the disc at r_max, the spectrum tends towards a Rayleigh-Jeans spectrum, I_ν ∝ν², while it shows a cutoﬀ in the inner parts, I_ν ∝ exp(−_kT^hν_in), where Tin is the temperature of the innermost layers of the thin accretion disc. An example of the spectrum expected from an accretion disk (for diﬀerent values of T_in) is shown in

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Fig.1.3.

Figure 1.3: Multicolor blackbody emission from geometrically thin optically thick Shakura-Sunyaev accretion disks with diﬀerent values of the temperature of the inner layerT_in.

The maximum temperature of a thin accretion disk around a Schwarzschild black hole occurs at r= ⁴⁹₃₆r_g, and is:

T_max= 1.4×10⁵

� M˙ 0.1 ˙M_Edd

�1/4� η 0.08

�_−1/4� M 10⁷M_⊙

�_−1/4 K.

For parameters typical of a Seyfert 1 galaxy, one obtains an energy of kT_max = 12eV, where k is Boltzmann’s constant. From this value one can already see that the disk alone cannot account for the X-ray emission observed in AGN. Recently, Davis & Laor(2011) proposed a way to measure the accretion rate in AGN based on the optical luminosityL_opt ≡νL_ν at 4686 Å,

M˙ = 3.5M_⊙yr⁻¹(Lopt,45)^3/2M⁻₈^0.89,

where M₈ = M_BH/10⁸, and L_opt,45 =L_opt/10⁴⁵erg s⁻¹. Using a sample of ∼80 PG Quasars, they also found a possible correlation between the eﬃciency and the black hole mass,s

η�0.089M₈^0.52.

However, it has been argued by Raimundo et al. (2011) that this correlation between η and M_BH might not be intrinsic, but due to a bias in the sample used.

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Alternative accretion scenarios

Besides thin accretion disks, three other solutions to the hydrodynamic equations of viscous rotating flows exist.

• The Shapiro, Lightman, & Eardley solution (1976, SLE) considers an accreting gas forming a two temperature plasma, with the ion temperature being greater than that of the electrons (Ti∼10¹¹K,Te∼10⁸⁻⁹K). The gas is optically thin and would produce a power-law spectrum in the X-rays and in the softγ-rays. This solution is however known to be thermally unstable (Piran 1978), and it is not thought to be realistic.

• The optically thick advection-dominated accretion flow (or optically thick ADAF) solution was found for objects accreting at super-Eddington accretion rates (Katz 1977, Abramowicz et al.

1988). In this scenario, the large optical depth of the inflowing gas traps most of the radiation, advecting it into the central black hole.

• The optically thin ADAF (two-temperature ADAF) solution is found for objects accreting at sub-Eddington accretion rates (Narayan & Yi 1994). In this scenario, the accreting gas has a very low density, and is unable to cool eﬃciently within an accretion time. The viscous energy is therefore stored in the gas as thermal energy instead of being radiated, and is advected into the central object. The gas is optically thin, and has a two-temperature configuration, as for the SLE solution.

An alternative scenario to the accretion disk models, involves accretion of matter in clumps (Courvoisier & Türler 2005,Ishibashi & Courvoisier 2009). The interaction of these clumps at∼100 Schwarzschild radii generates optically thick shocks that produce optical-UV photons, whereas the X-rays are generated in optically thin shocks closer to the central black hole.

1.2.3 The Broad-Line Region

Broad emission lines are one of the dominant characteristics of AGN spectra. These lines are assumed to be Doppler-broadened, and are thought to be produced in a region close to the black hole, normally referred to as broad-line region (BLR). The broad emission lines of AGN are an useful probe of the central engine. In fact the bulk motions in the BLR are regulated by gravity (due to the black hole) and radiation pressure (from the accretion disk), although the influence of outflows (see Chapter5) should probably also be taken into account. The widths of AGN broad lines span over two order of magnitudes, and range from a minimum of ∆v_{F W HM} � 500 km s⁻¹ (not so much larger than the narrow lines) to ∆v_{F W HM} � 10⁴km s⁻¹, with typical values of ∆v_{F W HM} � 5000 km s⁻¹. The strongest lines observed in the typical spectrum of an AGN are the hydrogen Balmer-series lines (Hα, Hβ and Hγ), the hydrogen Lyα, and lines from abundant ions (Mg II, C III], and C IV), for an example see Fig.1.4. Details on these lines are reported in Table1.1. In addition to these lines, there are several which are blended because of their large Doppler widths. A list of these lines is reported in Table1.2.

The broadening of the emission lines is not due to thermal motion. In fact from the lines relative intensities one can calculate that the gas temperature isT ∼10⁴K. The velocity dispersion for a gas with such a temperature would be v ∼(kT /m_p)¹² ∼10 km s⁻¹, which is much less than the typical width observed in AGN. If the broadening is purely thermal, then from the typical widths of broad emission lines, one would have an unrealistic temperature of T � 10⁹K. Most of the broadening is thus attributed to bulk motions of individual line-emitting clouds. The absence of forbidden lines as

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[O III]λ5007 gives a lower limit to the ion density of∼10⁸cm⁻³ (Netzer 1990), although the strong Lyα and C IV are probably produced in environments with higher densities of ∼ 10¹¹cm⁻³. The equivalent width of the Lyα line allows to estimate the fraction of the continuum absorbed by the BLR, and its covering factor (f = Σ/4π), which results to be ∼10%.

The mass of the BLR can be calculated from the luminosity of C IV (MBLR ∼10⁻³L42( C IV)M_⊙) and is negligible compared to that of the central black hole. In fact, even for the most luminous AGN one obtains M_BLR � 10M_⊙. The distance of the BLR from the central black hole can be inferred studying the delay of the broad emission lines with respect to the continuum, and it is found to vary with the luminosity (R_BLR ∝ L^0.7, Kaspi et al. 2000). The average values found are of 0.01-0.1 pc for Seyfert 1s, and up to ∼ 1pc for bright quasars.

Figure 1.4: Ultraviolet spectrum of the Seyfert 1 NGC 5548. The spectrum highlights some of the most important features produced in the BLR. Figure fromPeterson (1997)

1.2.4 The Narrow-Line Region

The narrow-line region (NLR) is the most extensive component of AGN, and the only one which can be resolved by optical observations. Unlike the BLR, in the NLR the electron density is low enough that many forbidden transitions are not collisionally suppressed. Some of the most prominent narrow lines are Lyα λ1216, C IVλ1549, C III]λ1909, [O III]λ4959, [O III]λ5007, [N II]λ6584 and [SII]λ6717. The full width at half maximum for narrow emission lines lies in the range 200�∆v_FWHM�900 km s⁻¹, with most of the lines having values of 350−500 km s⁻¹. Analogously to the BLR, the distance of the NLR to the central engine also scales with luminosity (R_NLR∝L^0.5), and in Seyfert galaxies has sizes of ∼100−300pc, while it might reach diameters of up to few kpc for bright quasars.

The mass of the NLR can be calculated from the luminosity of Hβ(M_NLR�7×10⁵L₄₁( Hβ)/n₃M_⊙, wheren3 is the electron density in units of10³cm⁻³). The mass of the NLR is several orders of magnitude larger than that of the BLR, although the amount of line mission produced is often comparable

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Table 1.1: Typical emission lines observed in AGN, table adapted from Peterson(1997).

Line Relative flux^∗ Equivalent width (Å)

Lyα λ1216 + N Vλ1240 100 75

C IVλ1549 40 35

C III]λ1909 20 20

Mg IIλ2798 20 30

Hγ λ4340 4 30

Hβ λ4861 8 60

∗Normalized in such a way that Lyα+N V=100

for the two regions. This is due to the fact that the emissivity of recombination lines is proportional to the ion density, which makes the BLR more eﬃcient than the NLR.

Figure 1.5: [OIII] ionization cone of NGC 5252 observed by the Wide Field Camera on board of the Hubble Space Telescope (Dadina et al. 2010).

The NLR is normally found to be axisymmetric rather than spherically symmetric. One of the most interesting features of AGN are the "ionization cones" (e.g.,Pogge 1988) which are clearly detected in maps of high-excitation lines as [O III]λ5007(see Fig.1.5). These cones have a [O III]λ5007/Hα flux ratio higher than one, which is characteristics of low-density gas ionized by the AGN continuum.

Outside the cone the ratio is instead lower than the unity, which implies that the gas is mostly ionized by starlight.

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Table 1.2: List of blended broad lines, table adapted fromPeterson (1997).

Feature Contributing Lines

Lyβ + O VIλ1035 Lyβ λ1026; O VIIλλ1032, 1038 Lyα + N V Lyα λ1216; N Vλλ1239, 1243 Si IV + O IV] Si IVλλ1394, 1403; OIV]λ1402 C IVλ1549 C IVλλ1548, 1551

He II + O III] He IIλ1640; OIII]λ1663

C III] + Si III] Al IIIλ1857; Si III]λ1892; C III]λ1909

Small blue bump Balmer continuum (λ <3646 Å); Fe II (many lines) Mg IIλ2798 Mg IIλλ2796, 2803

Fe IIλ4570 + He II Fe II (multiplets 37, 38, & 43); He IIλ4686 Hβ Hβ λ4861; Fe IIλλ4924, 5018 (multiplet 42) Fe IIλλ5190, 5320 Fe II (multiplets 42, 48, 49, and 55)

1.2.5 The Corona

X-ray variability time scales indicate that the X-rays in AGN are produced in a small region located close to the black hole. The possibility that the disk is responsible for the X-ray emission can be discarded on the basis that even very warm disks are not supposed to exceed few hundreds eV (see Sect.1.2.2). It is now widely accepted that the X-ray emission of AGN is produced by Comptonization of optical-UV photons produced in the accretion disk by a corona of hot electrons. Inverse Compton upscatters to higher energies photons when their mean energy (< E >) is smaller than the thermal energy of the electrons (of temperatureT_e): ^<E>_mc2 < ^4kT_mc2^e. Forhν �m_ec² the gain rate of the photon field is given by

dE dt = 4

3σ_TcU_rad(v²/c²)γ²,

where σT is the Thomson cross-section, U_rad the energy density of the photon field, and γ =

�1−v²/c². The spectrum created by this process depends on the Compton parameter y, given by

y= kT_e

mec² max(τ_e, τ_e²),

whereτ_e is the optical depth of the corona. The Compton parameter is basically the product of the average energy gain per scatter and the total number of scatterings. If the electron gas is optically thin, it can be shown (Shapiro et al. 1976) that the spectrum is a power-law with a photon index given by

Γ =−1 2 +

�9 4 +4

y.

As the energy of the photons becomes comparable to the thermal energy of the electrons, the power law emission declines in a cutoﬀ atE_C�3kT_e.

The origin and the geometry (see Fig.1.6) of the hot corona is still debated, several hypothesis

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have been put forward in the last years. In the following I describe three of them: magnetic flares, clumpy discs and aborted jets.

• Magnetic flares. A popular explanation for the heating of the electrons in the corona is related to the presence of flares above the accretion disc, similarly to what is observed in the solar corona. Although the physical details are rather unclear, magnetic flares are a reasonably working hypothesis to model X-ray spectra and variability of AGN (e.g., Haardt et al. 1994, Goosmann et al. 2006). The possible evidence for hot spots corotating with the disk may support this scenario (e.g., Iwasawa et al. 2004). In this scenario the magnetic flares heat the electrons, which then Comptonize the UV photons from the disk, upscattering them into the X-rays.

• Clumpy disks. The hypothesis that the accretion flow breaks up in two phases due to disk instabilities was originally proposed by Guilbert & Rees (1988). In this model one part of the disk is hot and optically thin, while the other is cold and optically thick. The hot phase is responsible for the primary X-ray emission, while the cold phase provides the seed photons for the Comptonization and it reprocesses the primary continuum.

• Aborted jets. The idea that the primary X-ray emission of AGN comes from an aborted jet (a jet that stops in the proximity of the supermassive black hole) was introduced by Henri &

Petrucci(1997). Ghisellini et al.(2004) proposed a lack of momentum to explain the jet failure.

If the matter is launched (at the expense of the angular momentum of the black hole) on the black hole axis in form of blobs, and the velocity of these blobs is lower than the escape velocity, then the blob comes back and collides with the blobs that are being ejected. Kinetic energy is then converted into internal energy, providing the hot electrons necessary to Comptonize the UV photons from the disk.

An alternative model for the production of X-rays has been proposed by Ishibashi & Courvoisier (2009), who pointed out that shocks in a clumpy accretion flow might also produce the observed X-ray luminosities.

1.2.6 The Jets

After the first observations of radio lobes in galaxies in the early 50’s (Jennison & Das Gupta 1953), it was thought that these structures were ejected and completely detached from the host galaxy. It was not until the late 1960’s that it was realized (Rees 1966) that they could actually be powered through jets emitted from the center of galaxies. In the last few decades, in particular thanks to the development of radio interferometry techniques, it has been possible to detect jets in many sources and to resolve even their inner parts (e.g.,Chang et al. 2010). Jets are now known to extend on scales from parsec to hundreds of kpc, are highly collimated and can have continuous appearance or can present knots. Luminous objects sometimes show two jets with very diﬀerent intensities. This, together with the observation of apparent superluminal motions in the jet of some objects (Whitney et al.

1971), is interpreted as an eﬀect of orientation and Doppler boosting. The production, acceleration and collimation of jets is still poorly understood. Some of the most popular models are based on magnetohydrodynamics (e.g., Blandford 2001), and involve the presence of strong electromagnetic fields that convert the rotational kinetic energy of a rotating black hole into an outflow through coupling with diﬀerential rotation.

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Figure 1.6: Some of the possible geometries for the accretion disk and the Comptonizing corona of AGN. The top panel shows aslab or sandwich geometry. The remaining three show photon starved geometries, namely geometries in which the corona is less eﬀectively cooled by soft photons from the disk. The two geometries in the middle panels are often referred to as sphere+disk geometries, while the bottom geometry is often referred to as apatchy corona or pill box(Reynolds & Nowak 2003).

1.3 AGN classification

AGN can be distinguished from normal non-active galaxies in several diﬀerent ways. One of the most commonly used diagnostics are the so called BPT diagrams (Baldwin et al. 1981), which show the ratios of several lines (e.g., [O III]/Hβ vs [N II]/Hα, [O III]/Hβ vs [S II]/Hα, and [O III]/Hβ vs [O I]/Hα, see Fig.1.7). Comparing these ratios to theoretical values, one can easily spot the objects that cannot be photo-ionized by pure stellar emission, but need other mechanisms, such as AGN or shocks. Similarly, one can distinguish AGN from starburst galaxies or pure stellar emission.

Probably the best way to recognize AGN is thanks to X-ray surveys, X-ray emission in AGN being practically ubiquitous. Other techniques involve radio selection, which was also the first method used to find AGN, or the study of IR colors. For a complete review on the subject one can consult Mushotzky (2004).

AGN come in many flavors. A first division between diﬀerent types of AGN is based on the relative strength of their radio emission, which is normally associated to the presence or the absence of a jet.

AGN with a strong radio emission are called radio-loud, while those weak in the radio band are called radio-quiet. The distinction between radio-loud and radio-quiet objects is usually based on the radio loudness parameter, R =F_R/F_o, which is defined as the ratio of monochromatic luminosities at 5 GHz and optical B band at 4400 Å. Kellermann et al. (1989), studying the distribution of the R parameter in the Palomar Bright Quasar Survey, were the first to find evidence of bimodal distribution, which implied the existence of two distinct classes: radio-quiet (R ∼ 0.1−1) and radio-loud (R ∼ 10−1000). The classical boundary between the two populations is set at R ∼ 10 (Kellermann et al. 1994). An alternative way to quantify the radio loudness of a source is to weight the emission in the radio intensity with that in the X-rays. Terashima & Wilson (2003) introduced

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Figure 1.7: BPT diagrams for a large number of IR-selected objects. HII galaxies are in black, Seyfert 2s in blue, and LINERs green (Brightman & Nandra 2011b).

the parameter R_X = νL_ν(5GHz)/L_X, where L_X is the luminosity in the 2−10keV energy band.

The threshold of this parameter, above which the object is considered radio-loud, is commonly set to logR_X =−4.5.

1.3.1 Radio-quiet AGN

Radio-quiet AGN are divided into low luminosity Seyfert galaxies (Sy) and high luminosity quasars.

The division between the two classes is usually set to M_B = −23mag, where M_B is the absolute magnitude in the B band. Seyfert galaxies are the most common AGN in the local universe, and are classified according to their emission lines into Seyfert 1 (showing broad and narrow emission lines), and Seyfert 2s (showing only narrow lines). Many Seyfert galaxies exhibit permitted line profiles with both very broad and relatively narrow components (Osterbrock & Koski 1976, Osterbrock 1981a, Cohen 1983). These objects have been classified as type 1.2, 1.5, 1.8 or 1.9 Seyfert galaxies, depending on the relative contributions of the two components to the total permitted line profiles (see Table1.3).

A parameter often used to distinguish one class from the other is R_s, the ratio of the broad Hβ flux to that of [OIII]λ5007 (e.g., Véron-Cetty & Véron 2010). For Sy1s one has R_s > 5.0, for Sy1.2s 2 < R_s < 5, for Sy1.5s 0.33 < R_s < 2.0 and for Sy1.8s R_s < 0.33. In Sy1.9s and Sy2s Hβ is not detected, thus one cannot use the R_s parameter.

Narrow line Seyfert 1s (NLS1s) are a particular class of Sy1s, characterized in the optical by permitted lines only slightly broader than the forbidden ones. The first of this class to be noticed for the strange properties of its optical spectrum was Mrk 359 (Davidson & Kinman 1978). In the following years many more sources showing similar characteristics were found, which led Osterbrock

& Pogge (1985) to classify them as part of a new class. In the optical, NLS1s have FWHM(Hβ)<

2000 km s⁻¹, they often show a strong emission from FeII multiplets, and they have a weak [ OIII]

emission, with a ratio[ OIII]/Hβ <3. It has been suggested that NLS1 are AGN in their early phase (Grupe et al. 1999), characterized by relatively small black hole masses (e.g.,Grupe & Mathur 2004) and very high accretion rates in terms of Eddington units (e.g., Grupe et al. 2010). In this scheme the small black hole mass would be responsible for the narrow optical lines, the matter in the broad line region (BLR) having lower velocities than in broad line Seyfert 1s (BLS1s).

The optical spectra around the Hβ region of a Sy1, a Sy2 and a NLS1 are shown in Fig.1.9, from which it is possible to see the main diﬀerences amongst these three diﬀerent classes.

Low-ionization nuclear emission-line regions (LINERs) are a peculiar type of AGN (Heckman 1980). Their spectra typically include line emission from weakly ionized or neutral atoms, such as

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Table 1.3: Characteristics of the optical spectra of diﬀerent types of Seyfert galaxies. Table adapted fromVéron-Cetty & Véron(2010).

Class Characteristics

Seyfert 1 Narrow and broad emission lines Seyfert 1.2 Weaker Hβ component than Seyfert 1s

Seyfert 1.5 Strengths of the broad and narrow components in Hβ are comparable

Seyfert 1.8 Broad components are very weak, but detectable in Hβ as well asHα

Seyfert 1.9 Broad component is detected only in the Hα line, and not in the higher order Balmer lines

Seyfert 2 Narrow emission lines only

Figure 1.8: Composite spectrum of Type-1 AGN taken from the SDSS (Vanden Berk et al. 2001).

The continuum is fitted with a broken power law.

O, O⁺, N⁺, and S⁺, while emission from strongly ionized atoms, such as O⁺⁺, Ne⁺⁺, and He⁺, is relatively weak. Spectroscopically, the LINERs might resemble Seyfert 2 galaxies, although some diﬀerences are evident (see also Fig.1.10): the[ OIII]/Hβ flux ratio is much larger in Seyfert 2s, and low-ionization lines are more prominent. LINERs are also classified, analogously to Seyfert galaxies, in type 1 and type 2, depending to the presence or absence of broad-line emission (Ho et al. 1997a).

LINERs are very common: approximately one-third of all nearby galaxies (within approximately 20-40 Mpc) may be classified as LINER.

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4400 4600 4800 5000 5200 5400 5600 0

.5

1 Sy2 (Mrk 1066)

NLS1 (Mrk 42)

Sy1 (NGC 3516)

Hβ He II

He II [O III]

Fe II

Wavelength

Relative Flux

Figure 1.9: Optical spectrum around the Hβ region of the Sy2 Mrk 1066, the Sy1 Mrk 42 and the NLS1 NGC 3516 (Pogge 2000).

Figure 1.10: Optical spectrum of the LINER NGC 1052 (Ho et al. 1993), figure adapted by Peterson (1997). The image shows the most important emission lines, along with some strong absorption lines from the host galaxy.