Aristotle University of Thessaloniki
Properties of Active Galactic Nuclei (AGN) in the X-Ray cluster population of the XXL Survey
by Eleftheria-Maria Dringa
Master Thesis Supervisors: Elias Koulouridis Manolis Plionis
October 16, 2020 This page intentionally left blank. 0.1 Abstract
In this thesis we investigate the properties of AGN in in the X-Ray cluster population of the XXL Survey. We study the AGN properties in clusters as a function of clus- tercentric distance. To this end, we have used a sample of 165 clusters with uniquely well-defined properties which allowed us to accurately determine their radius and to divide the cluster regions into six bins (r500 −6r500). Whithin this radius we find 166 AGN. For the analysis we use Hyper Suprime-Cam (HSC) photometry and SDSS spectroscopy. Our spectroscopically confirmed AGN cover a redshift range from z ' 0.1 − 1 while our photometric AGN cover a redshift range from z ' 0.01 − 2.5. We study the AGN interactions in cluster center (< r500) and in cluster outskirts (r500 −4r500) with respect to the background (5r500 −6r500) as well as the correlations between parameters such as the morphology of the host galaxies and the spectral types (Type I/Type II). We finally explore the dependence of cluster AGN fraction on redshift. Our results showed a significant excess of merging/interacting AGN in the center and in the outskirts of the cluster compared to the background. Addition- ally we observe a higher frequency of spiral and irregular host AGN galaxies within r500 − 2r500 radius. We conclude that merging/interaction within galaxy clusters is a significant factor in the triggering of AGN that can also affect their morphology.
I 0.2 Περίληψη
Στην εργασία αυτή διερευνούμε τις ιδιότητες των AGN στον πληθυσμό των X-ray σμηνών από το XXL Survey. Μελετούμε τις ιδιότητες των AGN στα σμήνη ως μια συνάρτηση σμηνοκεντρικής απόστασης. Για το σκοπό αυτό, χρησιμοποιήσαμε ένα δε- ίγμα 165 σμηνών με καλά καθορισμένες ιδιότητες που μας επέτρεψε να προσδιορίσουμε με ακρίβεια την ακτίνα τους και να χωρίσουμε τις περιοχές των σμηνών σε έξι bins (r500 − 6r500). Μέσα σε αυτή την ακτίνα βρίσκουμε 166 AGN. Για την ανάλυση χρη- σιμοποιούμε φωτομετρία από το Hyper Suprime-Cam (HSC) και φασματοσκοπία από το SDSS. Τα φασματοσκοπικά επιβεβαιωμένα μας AGN καλύπτουν ένα εύρος ερυ- θρομετάθεσης περίπου z ' 0.1 − 1 ενώ τα φωτομετρικά AGN καλύπτουν ένα εύρος ερυθρομετάθεσης περίπου z ' 0.01 − 2.5. Μελετούμε τις αλληλεπιδράσεις των AGN στο κέντρο, (< r500), και στα άκρα,(r500 − 4r500), των σμηνών σε σχέση με το υπόβα- θρο (5r500 − 6r500), καθώς επίσης και τους συσχετισμούς μεταξύ παραμέτρων, όπως η μορφολογία των host γαλαξιών και οι φασματικοί τύποι (Type I/Type II). Τέλος, εξετάζουμε την εξάρτηση των AGN των σμηνών από την ερυθρομετάθεση. Τα αποτε- λέσματά μας έδειξαν μια σημαντική αύξηση συνένωσης/αλληλεπίδρασης των AGN στο κέντρο και στα άκρα των σμηνών συγκριτικά με το υπόβαθρο. Επιπροσθέτως, παρα- τηρούμε μια αυξημένη συχνότητα σπειροειδών και ακανόνιστων γαλαξιών μεταξύ των ακτίνων r500 − 2r500. Καταλήγουμε στο συμπέρασμα ότι η συνένωση/αλληλεπίδραση στα σμήνη γαλαξιών είναι ένας σημαντικός παράγοντας στην ενεργοποίηση των AGN που μπορεί επίσης να επηρεάσει τη μορφολογία τους.
II Contents
0.1 Abstract...... I
AbstractI 0.2 Περίληψη...... II
Abstract in GreekII
1 Active Galactic Nuclei (AGN)1 1.1 Morphology...... 2 1.1.1 The supermassive black hole...... 2 1.1.2 The central accretion disk...... 3 1.1.3 The broad-line region...... 3 1.1.4 The narrow-line region...... 4 1.1.5 The jets...... 4 1.1.6 The powerful outflows...... 5 1.1.7 The X-ray corona...... 5 1.2 AGN and Host Galaxies...... 6 1.3 AGN activity and galaxy interaction...... 6 1.4 Classification...... 7 1.4.1 How to classify AGN...... 7 1.4.2 Type I and Type II AGNs...... 8 1.4.3 Seyfert Galaxies...... 9 1.4.4 Radio AGNs...... 12 1.4.5 Gamma-Ray AGNs-Blazars...... 14 1.4.6 Lineless AGNs...... 15 1.5 AGN through the Electromagnetic Spectrum...... 15 1.5.1 Optical-UV AGN observations...... 16 1.5.2 Infrared-submillimiter AGN observations...... 16 1.5.3 X-ray AGN observations...... 17 1.5.4 Radio AGN observations...... 17 1.5.5 Gamma-ray AGN observations...... 18 1.5.6 Discovering AGN...... 18 1.6 Unified Model...... 20 1.7 AGN in clusters...... 21 1.7.1 Ram pressure stripping...... 23 1.7.2 Jellyfish Galaxies...... 24 1.7.3 Cluster detection with X-rays...... 24
i 2 Data analysis and processing 25 2.1 XXL: The ultimate XMM extragalactic survey...... 25 2.2 Initial Data...... 26 2.3 Definition of cluster outskirts...... 27 2.4 Definition of Cluster mass...... 29 2.5 Hyper Suprime Cam (HSC)...... 29 2.5.1 The Subaru Telescope...... 30 2.6 Redshift and Mass Distribution...... 31 2.6.1 Cluster distributions...... 32 2.6.2 AGN distributions...... 32 2.7 HSC photometry & SDSS spectroscopy...... 34 2.7.1 1r500...... 35 2.7.2 2r500...... 36 2.7.3 3r500...... 38 2.7.4 4r500...... 39 2.7.5 5r500...... 41 2.7.6 6r500...... 42 2.7.7 Ds9 photometry...... 44
3 Results & Conclusions 45 3.1 Interacting AGN per bin...... 45 3.1.1 Redshift dependence...... 45 3.2 Interacting AGN per bin, mass distinction...... 46 3.2.1 Redshift dependence...... 46 3.3 Spectrum type per bin...... 47 3.3.1 Redshift dependence...... 47 3.4 Morphology type per bin...... 48 3.4.1 Morphology type per Spectrum type/Redshift dependence.. 48 3.5 Conclusions...... 49
Bibliography 51
ii Chapter 1
Active Galactic Nuclei (AGN)
Active galactic nuclei (AGN) are the most energetic persistent objects to the Uni- verse.The names “active galactic nuclei” or “active galaxies” are related to the main feature that distinguishes these objects from inactive (normal or regular) galaxies: the presence of supermassive accreting black holes (BHs) in their centers.
As of 2011, there were approximately a million known sources of this type selected by their color and several hundred thousand by basic spectroscopy and accurate redshifts. It is estimated that in the local universe, at z ≤ 0.1, about 1 out if 50 galaxies contains a fast-accreting supermassive BH, and about 1 in 3 contains a slowly accreting supermassive BH.
Figure 1.1: NGC 4151: Direct Photographs from the AGN, presented in the first article of Carl Seyfert about AGNs in 1943, Nuclear emission in spiral nebulae
Detailed studies of large samples of AGNs and the understanding of their con- nection with inactive galaxies and their redshift evolution started in the late 1970s, long after the discovery of the first quasi-stellar objects (quasars or QSOs) in the early 1960s. Although all objects containing BHs are now referred to as AGNs, various other names, relics from 1960s and 1970s, are still being used. Some of the names that appear occasionally in the literature, such as “Seyfert I galaxies” and “Seyfert II galaxies”, in honor of Carl Seyfert, who observed the first few galaxies
1 of this type in the late 1940s, are the result of an early confusion between different sources that are now known to have similar properties.
As a final point, we ought to define under what exact conditions an extragalac- tic object can be referred as an AGN. The definition of nuclear activity in galaxies can be based either on the physical mechanism involved or on the observational signature of this activity. The physical definition is simple: an extragalactic object is considered to be an AGN if it conyains a massive accreting BH in its center. the observational classification is, in many cases, not so clear because of observational limitations, because of source obscuration, and because the term activity covers many orders of magnitude in accretion rate. The definition adopted here is purely observational- an object is classified as an AGN if at least one of the following is fulfilled:
1. It contains a compact nuclear region emitting significantly beyond what is expected from stellar processes typical of this type of galaxy.
2. It shows the clear signature of nonstellar continuum emitting process in its center.
3. Its spectrum contains strong emission lines with line ratios typical of excitation by a nonstellar radiation field.
4. It shows line and/or continuum variations.
1.1 Morphology
The main components of an AGN are the central accretion disk around the galaxy’s supermassive black hole, the broad-line region (BLR), the obscuring central dust torus and the narrow-line region (NLR). In radio-loud AGNs can be observed another component, which is the pair of jets, originating from the central object. There is also much talk concerning the existence of powerful winds or outflows, originating from the vicinity of the central black hole as a main component as well.
1.1.1 The supermassive black hole It has long been assumed, because no other plausible explanation has emerged, that AGN must be powered by accretion onto massive black holes at the dynamical cen- ters of their host galaxies. These black holes are believed to have masses upwards of a million times the mass of Sun. The AGN are both compact, extremely luminous and often most variable at the shortest observed wavelengths. Thus, a mechanism such as accretion that can provide highly efficient conversion of potential and ki- netic energy to radiation is needed as is a large mass to accomodate high persistent Eddington luminosities.
2 Central supermassive black holes are mow beleived to reside at the center of most galaxies and AGN-like behavior is expected whenever a supply material comes within a critical distance of the central black hole. The most basic characteristic defining a black hole is the presence of an event horizon- a boundary in through which matter and light can fall inward towards the black hole, but can never re-emerge. Any information resulting from an ”event” occuring within the boundary defined by that radius cannot be communicated to an outside observer, thus the origin of the terminology.
1.1.2 The central accretion disk The accretion of matter onto a central massive black hole, leading to the conversion of gravitational potential energy to electromagnetic radiation, has long been postu- lated to power the enormous radiative output of AGNs. Details of the exact nature of this accretion flow remain to be worked out, however, disk accretion scenarios are strongly favored by theoretical arguments.
The accretion disk should be, or at least a large part of an AGN, a compact high-density structure with large column densities and typical dimensions of about 1000 rg. However, its possibles geometries depend heavily on the mass inflow rates, ranging from optically thick, geomatrically thick to thin disks or even a combination of those (two-component disk) for the case of objects with the lowest nass inflow rates. Additionally the disk is a strong optical and UV continuum emitter.
1.1.3 The broad-line region The first region closer to the center, located at about 0.1-1 pc from the supermassive black hole is the broad-line region. It is considered to be a large column-density re- gion (∼ 1023 cm−2) with high-density clouds around the ccreting black hole that are assumed to be able to survive over many dynamical times because they are either confined or extensions of self-gravitating objects, for example stars.
The typical Keplerian velocity of this region ia about 3000 km s−1, a trait re- flected to the broad width of the emission lines in this region’s spectrum. Due to the high density, here, we expect only allowed and semi-forbidden transitions among the broad-lines such as MgII ,CIV ,CIII , and NIV .
A direct method to examine the BLR is via the reverberation mapping (RM). Anal- yses of RN suggests an inhomogeneous BLR over a large range distances which consists of different ”layers” of ionization energies, with higher ionization energy transitions corresponding to smaller radii. Even though the RM can give us valu- able information about the BLR, or even about the black hole in the centre, it is an extremely time-consuming and complex technique.
3 1.1.4 The narrow-line region The next region is the narrow-line region (NLR). Located at about 100 pc (or up to 3 kpc for very luminous AGNs), it has smaller column densities (∼ 1020−21 cm−2) and low density gas concentrated to clouds with ionization levels similar to the BLR. This regions’s typical velocity is about 300 km s−1 therefore this region’s emission lines have the same width of the narrow lines. Unlike the BLR, the gas is likely to contain dust due to a weaker radiation field and also intense forbidden lines are axpected to be observed, such as OI ,OIV , and OIII , since the density is much lower. Unlike the BLR, the NLR morphology is not spherical but appears as two cone-shaped regions, that seems as the radiation of the AGN is not isotropic but direction-dependent.
Between the BLR and the NLR, some pc away from the central BH, we have the obscuring torus which plays an important role in the AGN unification that will be discussed later. The fact that the entire region lies beyond the dust sublimation radius as well as the related opacity at all wavelengths, except the infrared, suggests that it is likely to contain dust and molecular gas. Many observations support the notion that the thick structure, that actively obscures the central area, is mostly flat and shaped as a torus, with a small central opening, that gives us the ability to see the BLR at certain angles.
1.1.5 The jets Another feauture linked especially to radio-loud AGNs is their radio jets. These jets consist of particles, mostly electrons and protons, and have two antiparallel radio lobes that extend well beyond the optical host galaxy. In some cases one of them is not visible but it is believed to be due to relativistic beaming in the jet moving away from us, an effect that leads to a velocity dependent observed flux, in regards to the source luminosity, with respect to the observer.
Figure 1.2: VLBA image of the M87 jet taken at 15GHz
These lobes seem to originate from a central point, which matches the loca- tion of the optical, UV and X-ray continuum source. Keeping in mind that jets appear to have highly relativistic motions, we can safely assume that jets in AGNs must be formed very close to the Schwarzchild radius of the supermassive black hole.
In addition, jets appear to have linear polarization and non thermal emission, traits
4 that are directly linked to synchrotron radiation, explaining their dominance in the radio and suggesting the existence of powerful magnetic fields in the accretion disk.
1.1.6 The powerful outflows Last but not least, another possible feature of all the AGNs could be powerful outflowing winds, or outflows, originating from the accretion disk around the central black hole. Any attempt to calculate the AGN feedback efficiency accurately failed, therefore the scenario of mass loss from the accretion disk in the form of outflowing winds has been proposed.
1.1.7 The X-ray corona X-ray emission is ubiquitous in AGN, and is believed to be produced by Comp- tonization of optical/UV disk photons by a corona of hot electrons located above the SMBH (e.g., Haardt & Maraschi 1991). This Comptonization produces a power- law emission, with a photon index typically of Γ ∼ 1.8−2. The power-law continuum often shows a high-energy roll-over, which is usually located around few hundreds keV. This feature is directly related to the temperature and optical depth of the plasma of hot electrons responsible for the power-law emission.
Observational evidence, such as the presence of blueshifted absorption lines in UV and X-ray spectra, with velocities different than AGN’s systematic one, support the existence of such outflows. Even though many recent AGN studies are applying constraints on their physical and dynamical profile, the physics of such winds and their feedback mechanism is still unknown.
Figure 1.3: Reprocessing of X-ray radiation in AGN
5 1.2 AGN and Host Galaxies
The earliest discovered quasars, in the 1960s, were so bright that their optical images showed no signs of the host galaxies. This resulted in the names quasi- stellar objects (QSOs) and quasars and caused a lot of confusion and even some unusual explanations and models. Interestingly, the much earlier discovery of the first Seyfert galaxies, by Seyfert in 1943, raised no such questions. The luminosity of the central sources in these galaxies was about 2 orders of magnitude below the luminosity of the first discovered quasars, and the galaxy was clearly seen in all cases. The host galaxies of the earlier quasars were soon discovered by ground-based telescopes in sites with good seeing conditions. Faint nebulosities were discovered in all objects with redshift less than about 0.5, where conditions allowed such detection. The launch of the Hubble Space Telescope (HST), in 1990, resulted in superb resolution observations and the extension of such studies to redshifts 2 and 3. Some of the information was detailed enough to enable a systematic study of the morphology, color, and even stellar population of the hosts. Such information is now available for numerous low-redshift AGNs where the study is made easier because of the lower luminosity of low-z AGNs and the vast improvement in the performance of adaptive optics (AO) systems on giant ground-based telescopes. Similar observations of type-II AGNs, where the optical–UV radiation of the central source is completely obscured, allowed us to extend such studies to more sources and to higher redshift. The picture that emerged after several decades of study suggests that the mor- phologies of most AGN host galaxies are similar and even indistinguishable from the morphologies of inactive galaxies at the same redshift. In particular, the color and stellar population of AGN hosts are very similar to those of inactive galaxies of similar mass and morphology.
1.3 AGN activity and galaxy interaction
An interesting idea that was put forward in the early days of the field connected galaxy interaction and merging with the triggering of AGN activity. According to this idea, cold gas located far from the center of a galaxy cannot be brought to the vicinity of the BH because of its high angular momentum. Gravitational interaction with a nearby galaxy distorts the parent galaxy morphology, changes the orbits of the gas and stars, and allows the gas to find its way to the center. Close interactions that lead eventually to mergers were suggested to trigger short episodes of high accretion rate and fast growth of the central BH. The idea of galaxy interaction and merging can be tested observationally. We can count the density of galaxies near an AGN and compare to the mean density in the field. We can also search for galaxy–galaxy encounters and look for distortion and deviations from smooth and ordered morphology. The big advance in understanding AGN hosts is the result of systematic studies of larger AGN samples and a more consistent comparison with samples of inactive galaxies. High-spatial-resolution observations, mostly with the HST, are major tools in such studies. It seems that the vast majority of such hosts do not show indications of strong interaction, and there is no significant difference in the fraction of galaxies with distorted morphology in samples of active and inactive
6 galaxies. The only exception may be the hosts of the most luminous AGNs that show more disturbed morphologies compared with similar mass galaxies with no AGN.
1.4 Classification
1.4.1 How to classify AGN The classification of AGNs into subgroups is based on the history of research in this area. In particular, the discovery and general understanding of quasars, in the early 1960s, preceded the detailed study of the local members of this group, the Seyfert galaxies.
The earlier detailed observations of local, low redshift AGNs provided enough data to define several of the subgroups that are still used today: Seyfert I galaxies, Seyfert II galaxies, radio galaxies, quasars, LINERs, blazars, and so on.
Current AGN classification is based on higher-quality observations of a much larger number of sources, on better understanding of the physics of accretion and the line- emitting processes, and on the realization that many of the observed characteristics depend on the luminosity and inclination of the central source. The main subgroups of today are: type-I radio-quiet AGNs, type-I radio-loud AGNs, type-II radio-quiet AGNs, type-II radio-loud AGNs, LINERs and blazars.
So, after the ”historically” based classification, what parameters must we take un- der consideration in order to subclassify AGNs ? The following are three sets of somewhat different questions that can be posed to subclassify AGNs.
1. The first group is mostly observational:
What is the “power house” of the source? This type of classification refers to the nature of the main energy source. What is the observed SED of the source? This is the observational manifestation of the first quation. What are the properties of the host galaxy? This is related to the idea that the host galaxy properties (mass, gas, dust, evolution) determine the AGN properties. What is the inclination of the source to the observer’s line of sight? The source inclination to the line of sight may cause noticeable observational differences that may affect the classification. What is the phase of activity and the amount of gas supply to the center? This relates to the possibility that the AGN phase of activity may affect its appearance.
2. The second group is based on the properties of the main AGNs components:
7 What is the BH mass? The answer to this question is, in principle, in- dependent of the BH activity (accretion rate) but must be influenced by BH evolution. What is the spin of the BH? For active BHs this will determine the mass- to-energy conversion efficiency and can be used to distinguish different AGNs. What is the accretion rate? This distinguishes AGNs by their total lumi- nosity and/or kinetic power. What is the gas and dust content, and the metallicity, in the nuclear region? This has observational consequences that can be used to classify AGNs by means of optical and IR spectroscopy. What are the properties of the host galaxy? Do different host properties result in different types of AGNs?
3. The third list is a combination of the first two lists:
What is the intensity and EW of the observed emission lines? This distinguishes between weak-line and strong-line objects What is the typical width of the observed emission lines? Do we see only broad lines, only narrow lines, or a combination of both? What is the level of ionization of the line-emitting gas? This can be used to distinguish low-ionization from high-ionization AGNs. How strong is the radio source?This can be used to distinguish radio-quiet from radio-loud AGNs. How strong is the X-ray source?This relative measure can be used to classify objects by their bolometric luminosity. Is there evidence for central obscuration? This is the main way to dis- tinguish type-I and type-II AGNs. What are the variability amplitude and time scale? This distinguishes blazars from other types of AGNs.
1.4.2 Type I and Type II AGNs A basic division that is based on UV, optical and MIR spectroscopy is to characterize an object by the widths and EW of its permitted emission lines, Hα, Hβ, CIV λ1549, MgII λ2798, and so on. The implementation of this method depends on source lu- minosity since very broad emission lines are hard to detect in low-luminosity sources.
Type-I AGNs are those objects with little or no obscuration of the central source of radiation (the central disk or any other source within ∼ 1000rg of the BH and with 45 −1 very broad permitted lines , more than about Lbol/10 erg · s . The dependence on the luminosity stems from a more fundamental dependence on the normalized accretion rate, L/LEdd.
Type-II AGNs are those sources with a completely obscured line of sight to the center at UV, optical and NIR wavelengths and permitted lines with FWHMs that are significantly smaller than the previously mentioned number and are consistent with the velocities of stars in the host galaxy.
8 Most Type-I and Type-II sources show strong forbidden lines at optical and MIR wavelengths. Notable examples are [OIV ]25.9µm,[OIII ]λ5007, [OII ]λ3727, [NII ]λ6584, and so on. In most type-I sources, the forbidden lines are consider- ably more narrow than the permitted lines. In type-II sources, the width and other features of the profile are very similar. An additional difference between the groups is the line EW. In high-luminosity type-I AGNs, the forbidden lines are seen against the AGN continuum, and hence their EWs are considerably smaller than in type-II sources, where the lines are seen against the (fainter) stellar continuum.
The spectral differences between type-I and type-II AGNs can be explained with a simple model that involves central obscuration. The obscurer has axisymmetric structure, such as a thick disk with height and radius of order unity. It is normally referred to as the central torus, and its exact structure is a major area of research. Models involving central tori of different properties are quite successful in explain- ing many AGN properties, including the NIR-MIR SED and the relative numbers of type-I and type-II sources in the local universe which are normally reffered to as ”unified models or unification schemes”.
1.4.3 Seyfert Galaxies Seyfert galaxies are the most common class of AGN we observe in the local Universe. Owing to their relative proximity we have the best available spectra and images for them, allowing us to study in detail the physical processes at work.
Optical classification Seyfert galaxies were the first AGN identified as such, firstly observed by Carl Seyfert in 1943.
What distinguishes Seyfert galaxies from nonactive galaxies in direct images is the bright, central point-like core. With the advent of larger telescopes and improved instrumentation, it became straightforward to obtain spectra of the unresolved core and from the host galaxy, and it was discovered that the highly ionized emission lines are indded emitted from the central core, while the hostr had typical galac- tic spectra. These host galaxy spectra are interpreted as a synthesis of the stellar populations and thus they display photospheric absorption lines imprinted on their continuum. Emission lines from low-ionized gas may be present depending on the type of the host galaxy.
The identification as a Seyfert galaxy is nowadays based on the spectral signa- ture of the AGN core. In case this spectrum shows highly ionized emission lines, it qualifies as a Seyfert. Studying the spectra of Seyfert galaxies, Khachikian and Weedman found two generally distinct types of optical spectra. While all spectra had narrow and unresolved emission lines from highly ionized material, only some of the Seyferts also exhibited broad lines. They seperated the Seyfert galaxies into classes according to their relative width of narrow (forbidden) lines and Balmer lines
9 (Khachikian and Weedman 1974). In the Seyfert class 1 the Balmer lines , mainly Hα, Hβ, Hγ, would appear broader than the forbidden lines, like the [OII] and [OIII] lines, and the [NII], [NeIII] and [NeIV] lines. In type 2 objects, both the
Figure 1.4: Optical spectrum of the narrow-line Seyfert 1 galaxy HS 1747+6837B forbidden and the Balmer lines show the same narrow width. A closer inspection of the spectra shows that actually Seyfert 1 galaxies have a broad and a narrow component in the Balmer lines. Another indication of a type 1 AGN is the presence of a strong Fe[II] line at 4570 A˚. In the continuum emission, one observes a super-
Figure 1.5: Optical spectrum of the Seyfert 2 galaxy HE 0201-3029 (z=0.036) position of the host galaxy and the AGN core. In Seyfert 2 galaxies the AGN core is usually less dominant with respect to the surrounding galaxy than in the Seyfert 1 objects. This is one reason why it is generally more difficult to find Seyfert 2 galaxies based on on their optical spectra: often longer exposure times are needed in order to seperate the weak emission line spectrum superposed on the synthetic stellar spectrum of the galaxy. On the contrary, Seyfert 1 galaxies often exhibit a strong continuum which cannot be attributed to the surrounding galaxy. This con- tinuum emission appears to be featureless, that is, without the characteristic stellar absorption lines. These featureless spectra often exhibit nonthermal continua which
10 further distinguish them from stellar spectra. This can be a somewhat misleading term, if we recall that the continuum emission in the AGN core can be understood as the signature of the accretion disk, and thus a superposition of many black body spectra of a continuus range of temperatures.
In spite of the clear distinction between Seyfert 1 and Seyfert 2, intermidiate and borderline objects are observed. These objects are sorted into subclasses Seyfert 1.2, 1.5, 1.8 and 1.9 according to their Balmer line characteristics, following the scheme introduced by Osterbrock (1977). The subtypes are based on the relative width of the Hβ line.
X-Ray classification Up to now we have distinguished Seyfert types based on the optical spectrum, seper- ating them into type 1 and type 2. A similar distinction is made in the X-rays based on the intrinsic absorption measurement in the soft (E << 5keV ) X-ray band. Here, intrinsic absorption indicates matter which is close to the central engine of the AGN, whereas other absorption of line of sight photons, for example in our own Galaxy, has been accounted for. The absorption is measured as a column density of 2 hydrogen NH in the line of sight in atoms per cm . The unification scheme for AGN, which will be described later, assumes that the differences between type 1 and type 2 Seyferts result from the amount of absorbing material close to the central engine.
The most strongly absorbed sources are often referred to as type 2, and those with lower absorption are labeled type 1. As a dividing line a hydrogen column density 22 −2 of NH = 10 cm has been chosen. Most, but not all, Seyfert galaxies which have 22 −2 an inferred intrinsic absorption of NH < 10 cm are classified spectroscopically as Seyfert 1 or Seyfert 1.2, while most, but again not all, Seyfert galaxies with 22 −2 NH > 10 cm are optical Seyfert 1.8, 1.9, or 2. It is likely that the transition between absorbed and unabsorbed sources is smooth, and that not all type 1 AGN exhibit low absorption (Awaki et al.,1991). In addition, Seyfert 2 galaxies (like NGC 3147 and NGC 4698) show no intrinsic absorption (Pappa et al., 2001). A catalog of hard X-ray-selected AGN contains typically around 10 percent of objects in which the optical classification does not match the one based on intrinsic absorption mea- surements, as seen for example in the AGN catalogs based on data from the hard X-ray satellites Swift (Tueller et al., 2008) and INTEGRAL (Bechmann et al., 2009). Thus it is important to know for each study whether the classification of an object has been based on X-ray or optical measurements. In addition, in some objects a variation of the intrinsic absorption with time has been observed, which causes objects to be unabsorbed in one X-ray measurement, while they are of absorbed type in another.
As a final remark, there is a subset of Seyfert-1 galaxies, the so-called narrow emis- sion line Seyfert 1 galaxies (NLS1) are strong X-ray emitters, but while their Hα lines are broad, their Hβ line is narrow, similarly as in Seyfert 1.9 (Osterbrock and Pogge, 1985)
11 Figure 1.6: The Perseus cluster as seen by the Chandra X-ray telescope. The center of the cluster is dominated by the emission of the Seyfert and radio galaxy NGC 1275.
1.4.4 Radio AGNs When the central region of a quasar is hidden but the object produces bright radio jets and large radio luminosities, the existence of an AGN is assumed. A scheme developed by Fanaroff and Riley in 1974 classifies radio galaxies according to the extended radio structure and to whether they are edge brightened (FRII sources) or edge darkened (FRI sources).
In spatially resolved FRI sources, the radio spots are seperated by less than half the overall size of the source, whereas in FRII sources, those points (e.g., the bright- est points in the radio lobes) are seperated by more than half the total extent of the source. FRI radio sources are much more numerous and less luminous than FRII radio sources. The dividing luminosity between the two groups is close to the break in the radio luminosity function. The division in luminosity is seen very clearly when plotting radio power versus the optical magnitude of the galaxy. However, the location of the dividing line itself is luminosity dependent such that it is at larger L(5GHz) in galaxies that are more luminous in the optical and NIR bands. An- other division is based on the luminosity and ionization of the narrow emission lines. Low-power radio galaxies with FRI radio structure are invariably weak-line galaxies with typical LINER emission-line ratios. Most radio galaxies show a synchrotron −αR power-law continuum, Lv ∝ v , extending over a large frequency range with clear indications for synchrotron self-absorption at low frequencies. Core-dominated sources show a ”flat” spectrum (flat- spectrum radio-loud AGNs) with αR < 0.5. Lobe-dominated, weak- core sources usually show a much steeper spectrum with αR > 0.5 (steep-spectrum radio-loud AGNs). It seems that much of the difference in the measured value of αR is due to the inclination of the central radio jet to the
12 Figure 1.7: The z=0.458 FRII radio galaxy 3C 200
Figure 1.8: The FRI radio galaxy M84 which shows a jet-dominated radio structure line of sight.
Speaking about radio jets, another way to subdivide radio-loud AGNs is based on the radio jets properties. Radio jets are easily found and mapped by modern radio techniques, and their basic properties like the degree of collimation and the radio SED, are known for many sources. Here, again, there is a dichotomy between jets that are strongly collimated and show indications of relativistic motions and less collimated, less extended jets. Strongly collimated jets are associated with both
13 FRI and FRII sources, but there are clear differences between them. Evidence for relativistic motion in such jets comes from radio variability, beaming and superlu- minal motion of clumps or blobs very close (a fraction of a pc) to the central BH. The appearance of the jet depends strongly on its viewing angle.
1.4.5 Gamma-Ray AGNs-Blazars The group of blazars includes highly variable core-dominated radio-loud sources showing polarization at radio and optical wavelengths. Many blazars are also power- ful gamma-ray emitters, and some of them show indications of superluminal motion. Specifically, a blazar is defined as an AGN when shows one or more of the following properties:
1. Intense, highly variable high-energy emission in the gamma-ray part of the spectrum.
2. Intense, highly variable radio emission associated with a flat radio spectrum and occassionally superluminal motion.
3. Radio, X-ray, and/or gamma-ray jet with clear indications for relativistic mo- tion.
4. A double-peak SED with a lower-frequency peak at radio-to-X-ray energies and a high-frequency peak at X-ray-to-gamma-ray energies.
5. Very weak (small EW) broad and/or narrow emission lines indicative of pho- toionization by a nonstellar source of radiation on top of a highly variable continuum.
Figure 1.9: Slit spectrum of the blazar RXJ 1211.9+2242
Blazars can be divided into BL Lacertae (BL-Lac) objects and flat- spectrum radio- loud AGNs. The flat radio spectrum blazars are occasionally called flat-spectrum radio quasars (FSRQs) or optically violently variable QSOs (OVVs).
14 1.4.6 Lineless AGNs Systematic studies of large AGN samples result in the discovery of a subpopulation of AGNs with extremely weak, usually totally undetected emission lines. A typical upper limit on the EW of the emission lines in such sources is 1A˚. The objects show at least one of the four AGN indicators, usually a nonstellar continuum with, occasionally flux variations. A clear indication for the active BH is an observed point X-ray source in many of the sources. The objects cover a large range in lumi- nosity, from vary faint objects in the local universe to very luminous AGNs at high redshift. They are referred to in the literature as ”lineless AGNs”, ”anemic AGNs”, ”dull AGNs”.
Lineless AGNs differ in their optical continuum properties from blazars. They do not show a power-law continuum, they are mostly radio-quiet, their variability is of very small amplitude and the typical double-peak SED of blazars is not observed.
The very luminous lineless AGNs are of special interest and many have a unique role in AGN evolution. These are high-redshift sources with extremely weak broad emission lines that are 1 or 2 orders of magnitude fainter (in term of line EW) com- pared to other type-I sources.
An explanation for the weak emission lines is related to the properties of the central accretion disk in such objects. This applies to very low as well as very high luminos- ity sources, but for different reasons. A very low accretion rate through the central disk can result in heating of the central part and the onset of radiation-inefficient advection-dominated accretion flow with inefficient conversion of gravitational en- ergy to Electromagnetic radiation. Such systems can lack much or all of the UV ionizing radiation. This has been proposed as a posible explanation for the very low luminosity of lineless AGNs.
1.5 AGN through the Electromagnetic Spectrum
The characteristic spectral signatures of AGNs are easily distinguished in several wavelength bands. It is customary to refer the spectral energy distribution (SED) and describe it in terms of the monochromatic luminosity per unit frequency −1 −1 −1 −1 (Lvergs Hz ), per unit energy (LEergs Hz ), or per unit wavelength −1 −1 −1 −1 (Lλergs A ). The equivalent monochromatic fluxes (Fv, FE, or Fλ erg s A ) contain an additional unit of cm−2 and are used to describe the observed properties. The conversion between frequency and wavelength is obtained from single energy conservation considerations:
Lvdv = Lλdλ (1.1)
The SED of many AGNs can be described, over a limited energy range, as
−a Lv ∝ v or (1.2) −β Lλ ∝ λ (1.3)
15 where α is the frequency spectral index, β is the wavelength spectral index, and β = 2 − α. For example, the observed 1200-6000 Angstrom continuum of many −0.5 −1.5 luminus AGNs is described, adequately, by Fv ∝ v or Fλ ∝ λ . This single power-law approximation clearly fails for wavelengths below 1200 Angstrom or above ∼6000 Angstrom.
1.5.1 Optical-UV AGN observations Optical images of luminus type-I AGNs show clear signatures of pointlike central sources with excess emission over the surrounding stellar background of their host galaxy. The nonstellar origin of these sources is determined by their SED shape and by the absence of strong stellar absorption lines. Type-II AGNs do not show such excess.
The optical-UV spectra shown in Figures 1 and 2 represent typical spectra of high- ionization luminus type-I and type-II AGNs. the added ”high-ionization” is needed to distinguish such sources from low-ionization type-I and type-II sources.
The striking differences between the high-ionization type-I and type-II spectra, which were the reason for the early classification into Seyfert I and Seyfert II galaxies, are the shape and width of the strongest emission lines. Type-II AGNs show only narrow emission lines with typical FWHM of 400 − 800 km s−1. In type-I spectra, all the permitted line profiles, and a few semi-forbidden line profiles, indicate large gas velocities, up to 5000 − 10.000 km s−1 when interpreted as owing to Doppler motion.
1.5.2 Infrared-submillimiter AGN observations The earlier infrared (IR) observations of AGNs provided broadband photometry in the near-IR (NIR) bands, J(∼ 1.2µm), H(∼ 1.6µm),K(∼ 2.2µm), and L(∼ 3.5µm). Advances in building IR detectors and launch IR-dedicated experiments allowed the extension of the measurements in wavelength and spectral resolution. Intermidiate- resolution J,H and K spectroscopy of AGNs is now a standard procedure. Ground- based N-band (∼ 10µm) imaging is commonly performed, and mid-IR (MIR) spec- troscopy, mostly by ISO and Spitzer, have provided good-quality spectra of hundreds of sources, some at redshift as high as 3.
The far-IR (FIR) band of thousands of AGNs has beedn observed by IRAS, with limited spatial resolution, and by Spitzer, with much improved resolution. The 2009 launch of Herschel is the most recent development in this area.
Submillimeter observations of a small number of AGNs are also available covering the range from 0.4 to about 1.2 µm.
16 1.5.3 X-ray AGN observations X-ray images of AGNs are usually not very interesting; a point source at all X-ray energies in type-I sources and a point source in hard X-rays only in type-II AGNs. Low resolution X-ray spectra of AGNs are available since the late 1970s. They cover the energy range from about 0.5 keV with a spectral resolution typical of propor- tional counters and CCD detectors. Using optical band terminology, these can be described as broad- or intermediate-band photometry. The situation is somewhat improved at higher energies, close to the strong 6.4 keV iron Kα line, where the resolution approaches that of low-dispersion optical spectroscopy. The Chandra and XMM-Newton mission, launched in 1999, improved the situation dramatically by providing grating spectroscopy of nearby AGNs. This has revolutionized X-ray stud- ies of AGNs and resulted in the identification of hundreds of previously unobserved emission and absorption lines.
The soft X-ray spectrum of many type-I AGNs is dominated by a plethora of narrow absorption lines superimposed on a strong X-ray continuum. This must represent material along the line of sight to the source. Narrow emission lines are often as- sociated with the strongest absorption lines. X-ray spectroscopy of type-II AGNs, those with an obscured soft X-ray continuum, shows the narrow emission lines more clearly because of the attenuated central continuum.
1.5.4 Radio AGN observations The discovery of radio galaxies preceded the optical discovery of AGNs. It goes back to the late 1940s and the early 1950s (except, of course, for the famous paper by sSeyfert from 1943). Many of these sources were later shown to have optical-UV spectra that are very similar to the various types of optically discovered AGNs. The main features of many such sources are single- or double-lobe structures with dimen- sions that can exceed those of the parent galaxy by a large factor and strong radio cores and/or radio jets in some sources that coincide in position with the nucleus of the optical galaxy.
Like optically classified AGNs, there are broad-line radio galaxies (BLRGs), the equivalent of the type-I sources; narrow-line radio galaxies (NLRGs), the spectro- scopic equivalent of type-II AGNs; and even weak-line radio galaxies (WLRGs), the equivalent of LINERs.
While most AGNs show some radio emission, there seems to be a clear dichotomy in this property. It is therefore customary to define the ”radio loudness” parameter, R, which is used to seperate radio-loud from radio-quiet AGNs. R is a measure of the ratio of radio (5 GHz) to optical (B-band) monochromatic luminosity, L (5GHz) L(5GHz) R = v = 1.36 × 105 , (1.4) Lv(4400A) L(4400A) where L(5 GHz) and L(4400 A) represent the value of λLλ at those energies. The di- viding line between radio-loud and radio-quiet AGNs is usually set at R=10. Statis-
17 tics of a large number of AGNs show that about 10 percent of the sources are radio-loud, with some indication that the ratio is decreasing with redshift.
Much of the radio emission in radio-loud AGNs originates in a pointlike radio core. The spectrum of such core-dominated radio sources suggests emission by a self- absorbed synchrotron source. Except for the self-absorption low-frequency part, −αR the spectrum is represented well by a single power law, Fv ∝ v . Sources with αR < 0.5 are usually referred to as ”flat-spectrum radio sources”, and those with αR > 0.5 are ”steep spectrum radio sources”. There is a clear connection between the radio structure and the radio spectrum of such sources. Steep-spectrum radio sources show lobe dominated radio morphology and are also less variables. Flat- spectrum sources have in general higher luminosity cores, larger amplitude varia- tions, and weak or undetected lobes. This dichotomy is interpreted as a dependence on the viewing angle to the core. In steep-spectrum sources, one is looking away from the direction of the nuclear radio jet, and the radio emission is more or less isotropic. In flat-spectrum sources, we are looking at a small angle into the core. The intensity is boosted due to the relativistic motion of the radio-emitting particles, and the variations are amplified. In many cases, there is evidence for superluminal motion in such sources.
1.5.5 Gamma-ray AGN observations Observations at energies above 100 keV show that most AGNs are weak high-energy emitters. However, this is not the case for a small fraction (less than 10 percent) of the population that are strong γ-ray emitters. All these AGNs are also power- ful, core-dominated, radio-loud sources. They are highly variable, at all wavelength bands, and it is thought that their γ-ray emission is highly collimated. If correct, the apparent high-energy luminosity is large, but the isotropic γ-ray emission is not a large fraction of Lbol.
Because of the low spatal resolution of present-day γ-ray instruments, all AGNs detected at such energies appear as point sources. The most advanced γ-ray obser- vatory as of 2011, the Fermi Gamma-Ray Space Telescope, allows us to probe their extremely high-energy emission up to about 300 GeV.
1.5.6 Discovering AGN There are various methods to discover AGNs, some almost as old as the subject itself. The most recent methods have resulted in the largest number of new sources and the most uniform samples. The most important techniques, in terms of the number of newly detected sources, are as follows.
• Discovery by Optical-UV properties
It is a fact that typical AGN SEDs are different in several ways from stel- lar SEDs.They cover a broader energy range and do not resemble a single- temperature blackbody. This difference provides a simple and efficient way of
18 discovering AGNs using broadband multicolor photometry. Several color com- binations,based on 3-band and 5-band photometry, are usuful in seperating AGNs from stars by their color. Type-I AGNs can be directly discovered by their spectrum, becaue of the large contrast between the strong broad emis- sion and absorption lines and the underlying continuum. This method resulted in several large high-redshift samples, however is very inefficient in discover- ing type-II AGNs, with their relatively weak emission lines and strong stellar continuum.
• Discovery by Radio properties
About 10 percent of all AGNs are core-dominated radio-loud sources. This provides an additional way to identify AGNs in deep radio surveys by corre- lating their radio and optical positions. Stars are extremely weak radio sources, and hence an optical point source that is also a strong radio source is likely to be a radio-loud AGN.
• Discovery by X-ray properties
Almost all AGN are strong X-ray emitters. This property can be used to discover AGNs by conducting deep X-ray surveys. The most sensitive deepest X-ray surveys tend to pick bright soft X-ray sources with strong 0.5-2 keV emission. Type-II AGN with obscuring column densities of 1022cm−2 or larger are more difficult to detect. Additionally, X-ray observations are not very efficient in discovering very high redshift AGN because of the limited sensitiv- ity of the X-ray instruments and the sharp drop of X-ray luminosity of such sources.
• Discovery by IR properties
Several recent IR surveys have been used to search for AGNs using their unique IR properties. This rewuires the use of at least two IR bands or a combination of one IR band with X-ray or optical observations. A very important aspect of such techniques is the ability to detect highly obscured AGNs. Specifically, a large fraction of such objects, especially at high redshift, do not show de- tectable X-ray emission, and being type-II sources, their optical spectrum is completely dominated by the host galaxy. Such sources would not be classified as AGNs base on their optical and X-ray properties if it weren’t for their IR -or mid-IR- properties.
• Discovery by Continuum Variability
This is an independent method based on the fact that optical variability is very common in type-I AGNs. The few available variability studies of large AGN samples show that their typical structure function is very different from the structure function of variable stars. Thus, the method can clearly distin- guish the two types of sources. Also, it does not suffer from the limitation of the color selection methods in various redshift bands where AGN colours
19 are similar to those of some stars, for example, at 2.5 < z < 3. The method requires at least two visits per field and follow-up spectroscopy. It is expected that this way of discovering type-I AGNs will become dominant when the Large Synopsis Survey Telescope (LSST) becomes operational around 2017.
1.6 Unified Model
A major observational and theoretical challenge is to construct a general picture that connects the various subgroups of AGNs. Such a scheme starts from the central supermassive accreting black hole and the central accretion disk and adds other features, related to the gas and dust in the system and the radio jets. All this is an attempt to make the ”AGN unified model”.
Figure 1.10: A side view of AGNs showing the main ingredients of a unified model
The simplest addition of a torus-like obscurer to the two central components, the black hole and the accretion disk, already explains most of the observed differ- ences between radio-quiet type-I and type-II AGNs. Such a structure introduces a viewing angle parameter that determines what AGN components will be seen from a given line of sight. This can account for the different observed properties of MIR, NIR, optical and UV emission lines in type-I and type-II sources, the luminosity and variability of the optical -UV continuum, the different amount of obscuration of the central X-ray source and the shape of the MIR continuum.
Adding radio-loud and gamma-ray-loud objects to this model requires an additional component and several adjustments. the additional component is arelativistic jet emanating from the vicinity of the BH. The complications arise from the fact that FRI and FRII jets are rather different in their physical properties and cannot be considered as a single component in this model. Reasonably, we can assume that in both cases, the jet’s direction at its launch point is along the spin direction of the BH and perpendicular to the plane of the accretion disk. This results in inclination- angle dependence of the jet luminosity and variability.
In a generalized unified model, blazars are those FRI and FRII objects seen at
20 a very small angle to the inner jet direction. Other radio-loud AGNs are seen at a somewhat larger angle. BLRGs are seen at somewhat larger arngles, about 60 degree, and NLRGs are the geometrical equivalent of radio-quiet type-II AGNs.
By stating that all types of AGN are the same, the Unified Model can make predic- tions, that can in fact be tested by observations. One of them, for instance, is related to the morphology of AGN itself questioning whether the BLR is in fact closer to the supermassive black hole than the NLR. Even though the BLR is hidden in the case of Type 2 AGN, we could see polarized light originating from that area.
However, extensive observations of highly polarized light have shown non thermal continuum emission from the central engine, as it had been assumed, with polariza- tion similar to the Balmer and FeII lines, validating the Unified Model, even though the degree of the polarization seems to correlate with the object’s accretion rate. The Unified Model, also, seems to be validated by the fact that many classifications in different wavelengths, for the same objects, agree on the general type of the AGN. In cocnclusion, the accuracy of the Unified Model is an open topic which certainly needs further investigation.
1.7 AGN in clusters
Clusters of galaxies are the largest gravitationally bound systems in the Universe. Their mass can exceed 10 15 M , that is 3 orders of magnitude more than the mass of our own Galaxy with M MW ’ 10 12 M (Battaglia et al., 2005). The main baryonic mass content of galaxy clusters is not in the galaxies themselves but in between, in the intracluster medium (ICM), which is bound by the deep gravitational potential of the cluster. Most of the inferred cluster mass exists as dark matter though: only about 5% of a cluster’s mass resides in the galaxies, and 10% in the hot ICM, 85% being made of dark matter. Thus, galaxy clusters are by far the largest constituent of the mass in the Universe. While only 0.4% of the total mass-energy content of the Universe is attributable to visible matter in stars and in the interstellar medium, the hot intracluster medium accounts for 4%, while the dark matter adds 26% to the budget. The remaining 70% is currently attributed to dark energy. A number of experiments including ground- and space-based telescopes are currently planned with the goal understanding the nature of dark matter and dark energy. So there is hope that we will eventually shed some light on this dark side of the Universe.
Galaxy clusters are observed in a variety of size and morphology. They can con- tain 50–1000 galaxies, although ∼ 50 gravitationally bound galaxies would be bet- ter called a galaxy group, like the Local Group, which is formed by the Milky Way, M31, the Triangulum Galaxy, NGC 3109, and some more ∼ 30 or so satellite galax- ies. Many, but not all galaxy clusters have one bright and massive galaxy in their center, which seems to dominate the core of the cluster. These galaxies are called the central dominant (cD) galaxies and are usually red and of early-type, with low star-formation rate, like other massive elliptical galaxies. In some cases these cD galaxies are AGN. The Virgo Cluster, for example, hosts in its center M87, which
21 appears to be a misaligned blazar. M87 has a prominent jet which has been ob- served throughout the electromagnetic spectrum. Another example of an AGN be- ing a cluster cD galaxy is NGC 1275 located at the center of the Perseus cluster. The nearest cluster, the Coma Cluster does not have a single cD, but is instead apparently dominated by two nonactive galaxies, NGC 4874 and NGC 4889. The latter is a giant elliptical galaxy that has recently been identified to host one of the most massive (but inactive) black holes known in the Universe, with a mass of 10 MBH w 10 Msun (McConnell et al., 2011).
There are several ways to classify galaxy clusters. One is the richness, which is the number of galaxies belonging to the cluster. In order not to include too many back- ground objects into the richness parameter, one defines richness as the number of galaxies which are not more than 2 mag fainter than the brightest member galaxy of the cluster. Another classification aspect can be the morphology of the cluster. A regular cluster is spherical symmetric and usually contains a large number of galax- ies. Regular clusters show a high central concentration and host mainly elliptical galaxies. One example for a regular and rich cluster is the Coma Cluster. Irregular clusters of galaxies appear in nonsymmetric morphology and have younger galaxy populations. About half of the galaxies in irregular clusters are spiral (late-type) galaxies. The Virgo Cluster is an irregular cluster.
Identifying a galaxy cluster as such is not always an easy task in the optical domain. Early studies by Abell and Zwicky in the 1950s relied on searching for agglomerates of at least 50 galaxies which were similar in size and luminosity, in order to ensure that they were at the same redshift distance. An easier way to find galaxy clusters is to search for extended, high-latitude X-ray sources. The intraclus- ter medium is comprised of hot gas with typical temperatures of T = 107 − 108K and as such the clusters emit bremsstrahlung radiation. The size of a cluster can be defined assuming that the cluster mass is in relaxed equilibrium and thus the kinetic energy is balanced by the gravitational potential.
Also galaxy clusters, like their constituent galaxies, can undergo merging events. With respect to mergers, clusters can be sorted into three categories. Merging clusters are the equivalent phenomenon to galaxy mergers. Also in the process of two merging clusters, there is a significant amount of gas mixing. Bow shocks in the ICM can heat up the cluster, as seen for example in the Ophiuchus Cluster. A spectacular example for a galaxy merger is the Bullet Cluster. This sys- tem is undergoing a particular high-velocity merging, with a relative speed on the order of 4500 km/s. Observations in the X-rays show that the shock of the merg- ing gas not only heats it up, but also slows down its bulk velocity (Markevitch et al., 2002). On the contrary, the individual galaxies are largely unaffected and remain collisionless. The same appears to be true for the dark matter in the cluster, which does not seem to be swept up like the ICM, but continues with the galaxy flow. Thus, the merging of galaxy clusters seems to be an efficient way to separate bary- onic from dark matter, at least for the duration of the process. This would indicate a low cross-section for dark matter.
22 Post merger clusters do not show anymore direct evidence for a recent merger, like bow shocks or disturbed morphology. But they remain at a high temperature as a result of the mixing of the ICM, and the cooling flows in their centers are still relatively weak. Finally, relaxed clusters have been undisturbed for a long time and their dynamical configuration satisfies the virial theorem. They are cooling down through radiation (mainly in the X-rays), which leaves the center of the cluster cooler than its sur- roundings. This causes matter to flow into the direction of the center of the cluster, the so-called cooling flow (Fabian, 1994). The in-falling matter forms new galaxies and causes a starburst in the center. An example for such a relaxed cluster with a cooling flow is the Perseus Cluster. As galaxies are the hosts of AGN, our main questions are: how many of the AGN reside in clusters of galaxy? What is the feedback effect of the AGN on the cluster and vice versa? Because of the AGN heating, the cold baryon content is significantly reduced in the central galaxy and thus inhibiting its star formation. In analogy to the case of AGN–host galaxy feedback, this AGN–cluster feedback can affect the entire inner core of a massive clusters. The heating and cooling at the cluster core are kept in balance by the AGN. When the gas cools down in the central ICM, more matter gets accreted onto the supermassive black hole, leading to enhanced activity causing reheating of the ICM. This in turn slows down the ac- cretion processes, causing the AGN to fall dormant again. And thus a new feedback cycle can begin. A feedback cycle like this is directly related to a phenomenon called ”Ram pressure stripping”.
1.7.1 Ram pressure stripping Galaxy clusters are permeated by hot, X-ray emitting gas known as the intra-cluster medium. As individual galaxies move within such clusters, they experience this intra-cluster gas as a ‘wind’ – much like the wind experienced by a moving bicyclist, even on a still day. ‘Ram pressure stripping’ occurs if this wind is strong enough to overcome the gravitational potential of the galaxy to remove the gas contained within it. The result of ram pressure stripping is a galaxy which contains very little cold gas. This effectively halts star formation in the galaxy, supporting the belief that ram pressure stripping could be one of the processes responsible for the morphology density relation. Gunn & Gott compare the ram pressure a galaxy moving within a cluster would experience due to the intracluster medium (just the gas in the cluster but outside of the galaxy) with the force available to hold the gas in a typical spiral galaxy. They find that, in the central regions of clusters, the pressure exceeds the force holding the gas in, so that spiral galaxies passing through these regions will be stripped of the gas they have between their stars. This agreed with observations that had been done even as early as the 1950s. Previous work had ascribed this absence of spiral galaxies to collisions between galaxies, but the density of galaxies in clusters is low enough that they will not collide with each other often.
23 1.7.2 Jellyfish Galaxies A jellyfish galaxy is a type of galaxy found in galaxy clusters. They are characterised by ram pressure stripping of gas from the affected galaxy by the intracluster medium, triggering starbursts along a tail of gas.
Figure 1.11: Jellyfish galaxies, Fumaglli+14 (MUSE spectroscopy)
1.7.3 Cluster detection with X-rays The expansion and structure formation history of the Universe is imprinted on the spatial distribution and number density of its largest collapsed entities, galaxy clus- ters. This makes galaxy clusters powerful probes for constraining cosmological pa- rameters such as the dark energy equation of state (e.g., Vikhlinin et al. 2009; Allen et al. 2011). Among others, X-ray observations of galaxy clusters are of particular interest because they trace the bulk of the baryonic component, the hot intracluster medium (ICM). (Florian K¨afer,Nicolas Clerc et al. 2020) The gas in the cluster’s gravitational field emits X-rays in high temperature, (107 − 108K) so they can be easily detected as extended sources with X-ray telescopes. Their detection has two main advantages. Candidate clusters are not affected by projection phenomena and we can relatively easily measure their temperature and their luminosity. As a result we can find their mass and their radius (size) and finally make important assump- tions about the extragalactic physical mechanisms and the universe evolution, as we speak about large scale structures.
Figure 1.12: Cluster X-ray detection
24 Chapter 2
Data analysis and processing
Our aim in this thesis is to research the properties of AGN in X-ray galaxy clusters. Specifically we study:
1. AGN interactions (tidal tails, neighbours)
2. AGN morphology (host galaxy)
3. AGN spectra (obscured, unobscured)
Why do we choose AGN?
AGN have a strong relation with the galaxy evolution, they co-evolt. In addition we find them in high redshifts because they are very luminous objects so we can go deeper in the Universe with AGN observations. Last but not least they correlate with the large-scale structure, namely the clusters which are the largest gravitationally bound systems in the Universe. Clusters can be easily detected in X-rays so as the AGN. We use XXL survey at XMM-Newton in the X-ray band. For the analysis we use Hyper Suprime-Cam (HSC) images and photometry and spectroscopy from SDSS and other spectra. The analysis is for the northern sky because HSC images are only from the northern field.
2.1 XXL: The ultimate XMM extragalactic sur- vey
Over the past two decades, the Chandra and XMM-Newton observatories have sup- ported numerous studies of X-ray-selected clusters of galaxies, active galactic nuclei (AGNs), and the X-ray background. Here, the XXL Survey is the largest XMM programme totaling some 6.9 Ms to date and involving an international consortium of roughly 100 members. The XXL Survey covers two extragalactic areas of 25 deg2 each at a point-source sensitivity of ∼ 5x10−15ergs−1cm−2 in the [0.52] keV band (completeness limit). The survey’s main goals are to provide constraints on the dark energy equation of state from
25 the space-time distribution of clusters of galaxies and to serve as a pathfinder for future, wide-area X-ray missions. The survey also includes cluster studies, AGN evolution, and large-scale structure, which are being conducted with the support of approximately 30 follow-up programmes.
Figure 2.1: XXL cluster detection-North
Every circle has approximately the size of the moon and corresponds to each detection of the field.
Figure 2.2: XMM-XXL/Rosat
2.2 Initial Data
For our calculations we used the current values for cosmological parameters and a flat Universe Ωm = 0.3, ΩΛ = 0.7,H0 = 73.4 km/s/Mpc (2.1)
26 Our dataset consists of a catalogue with 166 (82 spectroscopic and 84 photo- metric) AGN in clusters in the northern XXL survey at XMM-Newton in the X-ray band. For the analysis we use Hyper Suprime-Cam (HSC) images and photometry and spectroscopy from SDSS and other spectra.
• r500: 11 sources
• 2r500: 20 sources
• 3r500: 18 sources
• 4r500: 38 sources
• 5r500: 37 sources
• 6r500: 42 sources
What we define as r500 ? We use overdensity radii to define region in which properties are measured. Overdensity Radii • A radius within which the mean density is ∆ times the ρc at the cluster’s redshift Use overdensity radii to define region in which properties are measured • R500 means ∆ = 500, a radius measured out to in typical X-ray observations A radius within which the mean density is Δ times the critical
density (ρc) at the cluster's redshift Clusters are centrally concentrated so larger Δ correspond to smaller radii ρ Write radii as RΔ
● e.g. R200 means Δ=200 N.B. here ρ is the total mass density (not just gas) R
Overdensity radii allow fair comparisonFigure of properties 2.3: Density-Radius of clusters relation of different sizes, key part of self-similar model
2.3 Definition of cluster outskirts
A plethora of physical effects is believed to be acting in the outskirts of galaxy clus- ters, which ebbed away long ago in more central regions. This includes, e.g., break- down of equilibrium states like hydrostatic equilibrium (e.g., Nagai et al. 2007b), thermal equi- librium and equipartition (e.g., Fox & Loeb 1997), and ionization equilibrium (e.g., Wong et al. 2011). It is also in the outskirts, where structure formation effects should be widespread. Moreover, the primary processes of intra- cluster medium (ICM) enrichment with heavy elements (e.g., Schindler & Diaferio
27 2008) may be identified by determining the metal abundance up to the cluster out- skirts.
Where are the cluster outskirts?
r500 < clusteroutskirts < 3r200 (2.2) where r500 is the observational limit for X-ray temperature measurements, while 3r200 captures most of the interesting physics before clearly enetering the regime of the warm-hot ICM.
Figure 2.4: Simulated galaxy cluster. The white circles indicate r500, r200, rvir, and 3r200 moving outwards, respectively (Roncarelli et al. 2006)
A theoretical recipe that can be used to define a cluster “border,” “boundary,” or at least a “characteristic” radius is the spherical collapse model (e.g., Amendola & Tsujikawa 2010). Based on this very idealistic model, a virial radius, rvir, sepa- rating the virialized cluster region from the outer “infall” region, can be obtained by requiring the mean total mass density of a cluster, < ρtot > , to fulfill
3Mtot(< rvir) vir < ρtot > (< rvir) ≡ 3 = ∆c (z)ρc(z) (2.3) 4πrvir
where ρc(z) is the critical density of the Universe at redshift z. The virial over- vir density ∆c (z) is a function of cosmology and redshift (Kitayama & Sunto 1996).
• Another possibility to define a virial radius is to use the region whithin which the condition of virial equilibrium (2Ekin = −Epot) is satisfied. (Reiprich et al. 2013)
• We use a fixed value for ∆c in both observations and simulations 3Mtot( 28 • r500 ' 0.65r200 (Navarro et al. 1997) • Measurements limited to r500 explore only ' 10% of the total cluster volume • Due to their high particle backgrounds XMM-Newton is basically limited to . r500 for robust gas temperature measurements 2.4 Definition of Cluster mass The total mass of galaxy clusters can be determined by measuring ICM properties, like density, temperature, and pressure. Under the assumption that the ICM is in hydrostatic equilibrium with the gravi- tational potential, the integrated total mass profile, Mtot(< r), is given by 1 dP GMtot(< r) = − 2 (2.4) ρgas dr r where P is the gas pressure,, ρgas its density, and G the gravitational constant. Applying the ideal gas equation P = kBρgasT gas/µmp results in kBTgasr dlnρgas dlnTgas Mtot(< r) = − ( + ) (2.5) Gµmp dlnr dlnr where µ ≈ 0.6 is the mean particle weight in units of the proton mass, mp, and kB is Boltzmann’s constant. So, the total mass within a given radius depends on the gas temperature at this radius, as well as the temperature gradient, and the gas density gradient. There is no dependence on the absolute value of the gas density, only on its gradient. 2.5 Hyper Suprime Cam (HSC) The Hyper Suprime-Cam (HSC) is a gigantic mosaic CCD camera, which is attached at the prime focus of 8.2m Subaru Telescope. The HSC uses 104 main science CCDs, which covers 1.5-deg field-of-view in diameter with a pixel scale of 0.17 arcseconds, as well as 4 CCDs for auto guider, and 8 CCDs for focus monitoring. Figure 2.5: Hyper Suprime Cam (HSC)-XXL 29 Figure 2.6: HSC/Subaru telescope 2.5.1 The Subaru Telescope Subaru Telescope is located on the summit of Maunakea, a dormant volcano on the Big Island of Hawaii. The summit of Maunakea is an isolated peak that protrudes above most of the Earth’s weather systems. The air pressure up on Maunakea is only two-third of what it is at the sea level. Clouds typically form below the summit where an inversion layer keeps the clouds from rising to the summit. Because Hawaii is isolated from any other land mass, trade winds blow smoothly over the islands, and there are few cities to pollute its dark skies. The summit of Maunakea is one of the best astronomical observing sites in the world. With 13 telescopes from 11 countries in operation, Maunakea has more telescopes from more countries than any other observing site. Besides Subaru, there are three other 8-10 m class telescopes on Maunakea, the Gemini North telescope and the two Keck telescopes. Maunakea is an irreplaceable natural and cultural resource. As a science reserve, development is carefully managed to balance the needs of preservation as well as exploration. 30 Figure 2.7: Clusters-Subaru Gallery 2.6 Redshift and Mass Distribution The AGN of our dataset in the sky. -3.0 -4.0 -5.0 Declination -6.0 -7.0 30.0 32.0 34.0 36.0 38.0 Right Ascension Figure 2.8: XXL survey and AGN distribution 31 2.6.1 Cluster distributions 30 30 25 25 20 20 15 15 Clusters Clusters 10 10 5 5 0 0 0.00 0.25 0.50 0.75 1.00 1.25 1.50 1.75 2.00 0 10 20 30 40 50 60 Redshift z M [×1013M ] (a) Redshift (b) Mass Figure 2.9: Redshift and Mass distributions of the clusters Obviously, we see that the most of the clusters exist in a redshift range between 0.2 and 0.6. However ∼ 25 clusters have a higher redshift (0.7-2). Concerning the mass more than 60% of these clusters are massive. (A massive cluster has 14 Mcluster > 10 Msun 2.6.2 AGN distributions Here are the redshift distributions for the AGN of our dataset for all radii (bins). 2 5 4 3 1 2 Number of galaxies Number of galaxies 1 0 0 0.15 0.20 0.25 0.30 0.35 0.40 0.45 0.50 0.00 0.25 0.50 0.75 1.00 1.25 1.50 1.75 Redshift z Redshift z (a) 1st bin (b) 2nd bin 32 5 7 6 4 5 3 4 3 2 Number of galaxies Number of galaxies 2 1 1 0 0 0.2 0.4 0.6 0.8 1.0 0.2 0.4 0.6 0.8 1.0 1.2 Redshift z Redshift z (a) 3d bin (b) 4th bin 8 7 10 6 8 5 6 4 3 4 Number of galaxies Number of galaxies 2 2 1 0 0 0.2 0.4 0.6 0.8 1.0 0.0 0.5 1.0 1.5 2.0 Redshift z Redshift z (a) 5th bin (b) 6th bin And here are the mass distributions for the AGN of our dataset for all radii (bins). 4 4 3 3 2 2 Number of AGNs Number of AGNs 1 1 0 0 10 15 20 25 30 35 5 10 15 20 25 30 35 40 M (×1013 M ) M (×1013 M ) (a) 1st bin (b) 2nd bin 33 5 14 12 4 10 3 8 2 6 Number of AGNs Number of AGNs 4 1 2 0 0 0 10 20 30 40 50 60 10 20 30 40 50 60 M (×1013 M ) M (×1013 M ) (a) 3d bin (b) 4th bin 12 12 10 10 8 8 6 6 Number of AGNs 4 Number of AGNs 4 2 2 0 0 10 20 30 40 50 60 0 10 20 30 40 50 60 M (×1013 M ) M (×1013 M ) (a) 5th bin (b) 6th bin 2.7 HSC photometry & SDSS spectroscopy To begin with, we want to examine all the AGN from our sample in HSC and see their morphology and their environment, specifically we are looking for merging phenomena or interactive neighbours. We divide our dataset into 3 morphological categories. Elliptical (E), Spiral (S) and Irregular (Irr). For the elliptical and the spiral the definition is obvious. As irregulars we count the galaxies that have disturbed morphology, tidal tails or substructures. Concerning the interacting phenomena we take into account as Interacting AGN the cases when there are close galaxies with the same redshift in a radius ∼ 100kpc or when the morphology is disturbed, an evidence that shows that probably this AGN had a merging incident. Most of our spectra become from SDSS. We divide them in Narrow-line and Broad- line spectra. Moreover we edit the fits image from the HSC in DS9 in order to have an even more clear image. Here we present some of our AGN (from all bins) as we show them in HSC. 34 2.7.1 1r500 (a) HSC (b) SDSS Figure 2.16: AGN from Cluster 82 (a) HSC (b) SDSS Figure 2.17: AGN from Cluster 183 35 (a) HSC (b) SDSS Figure 2.18: AGN from Cluster 117 2.7.2 2r500 (a) HSC (b) SDSS Figure 2.19: AGN from Cluster 159 36 (a) HSC (b) SDSS Figure 2.20: AGN from Cluster 23 (a) HSC (b) SDSS Figure 2.21: AGN from Cluster 60 37 2.7.3 3r500 (a) HSC (b) SDSS Figure 2.22: AGN from Cluster 152, z=0.21 (a) HSC (b) SDSS Figure 2.23: AGN from Cluster 83, z=0.42 38 (a) HSC (b) SDSS Figure 2.24: AGN from Cluster 25, z=0.26 2.7.4 4r500 (a) HSC (b) SDSS Figure 2.25: AGN from Cluster 1, z=0.61 39 (a) HSC (b) SDSS Figure 2.26: AGN from Cluster 83, z=0.51 (a) HSC (b) SDSS Figure 2.27: AGN from Cluster 25, z=0.18 40 2.7.5 5r500 (a) HSC (b) SDSS Figure 2.28: AGN from Cluster 96, z=0.51 (a) HSC (b) SDSS Figure 2.29: AGN from Cluster 140, z=0.28 41 (a) HSC (b) SDSS Figure 2.30: AGN from Cluster 111, z=0.3 2.7.6 6r500 (a) HSC (b) SDSS Figure 2.31: AGN from Cluster 96, z=0.51 42 (a) HSC (b) SDSS Figure 2.32: AGN from Cluster 78, z=0.95 (a) HSC (b) SDSS Figure 2.33: AGN from Cluster 91, z=0.18 43 2.7.7 Ds9 photometry (a) Elliptical host galaxy (b) Spiral host galaxy (c) Merging host galaxy (d) Irregular host galaxy Figure 2.34: Characteristic AGN, analyzed with ds9 44 Chapter 3 Results & Conclusions In order to further investigate the possible existence of a correlation between the AGN in clusters and their properties we make some plots to see if there are indeed any differences between the AGN in the clusters outskirts (r500-4r500) and the background (5r500 and 6r500). 3.1 Interacting AGN per bin 0.8 0.7 0.6 0.5 0.4 0.3 0.2 Fraction of interacting AGN 0.1 0.0 1 2 3 4 5 6 R/R500 Figure 3.1: Fraction of interacting AGN per radius r/r500 √ The errorbars correspond to 1σ deviation ( N/N). 3.1.1 Redshift dependence We wanted to see if there are any differences including the distance and the timescale, in other words the redshift factor. For that reason we divided our sample in two 45 smaller datasets. The first one includes the AGN with z 6 0.5 and the second one includes the AGN with z > 0.5. 0.5 0.5 0.4 0.4 0.3 0.3 0.2 0.2 Fraction of interacting AGN 0.1 Fraction of interacting AGN 0.1 0.0 0.0 1 2 3 4 5 6 1 2 3 4 5 6 R/R500 R/R500 (a) z 6 0.5 (b) z > 0.5 3.2 Interacting AGN per bin, mass distinction We consider that when the mass is greater than 1014 we have a massive cluster. 1.0 M < 1014M M 1014M 0.8 0.6 0.4 Fraction of AGN 0.2 0.0 1 2 3 4 5 6 R/R500 Figure 3.3: Fraction of interacting AGN per bin with mass distinction √ The errorbars correspond to 1σ deviation ( N/N). 3.2.1 Redshift dependence We wanted to see if there are any differences including the distance and the timescale, in other words the redshift factor. For that reason we divided our sample in two 46 smaller datasets. The first one includes the AGN with z 6 0.5 and the second one includes the AGN with z > 0.5. 0.8 0.8 M < 1014M M < 1014M 14 14 0.7 M 10 M 0.7 M 10 M 0.6 0.6 0.5 0.5 0.4 0.4 0.3 0.3 Fraction of AGN Fraction of AGN 0.2 0.2 0.1 0.1 0.0 0.0 1 2 3 4 5 6 1 2 3 4 5 6 R/R500 R/R500 (a) z 6 0.5 (b) z > 0.5 3.3 Spectrum type per bin 1.0 Narrow Broad 0.8 0.6 0.4 Fraction of AGN 0.2 0.0 1 2 3 4 5 6 R/R500 Figure 3.5: Spectrum type distinction √ The errorbars correspond to 1σ deviation ( N/N). 3.3.1 Redshift dependence We wanted to see if there are any differences including the distance and the timescale, in other words the redshift factor. For that reason we divided our sample in two smaller datasets. The first one includes the AGN with z 6 0.5 and the second one includes the AGN with z > 0.5. 47 8 Narrow Narrow Broad 4.0 Broad 7 3.5 6 3.0 5 2.5 4 2.0 3 1.5 Fraction of AGN Fraction of AGN 2 1.0 1 0.5 0 0.0 1 2 3 4 5 6 1 2 3 4 5 6 R/R500 R/R500 (a) z 6 0.5 (b) z > 0.5 3.4 Morphology type per bin Spiral 70 Elliptical Irregular 60 50 40 AGN 30 20 10 0 1 2 3 4 5 & 6 Bins Figure 3.7: Morphology type distinction √ The errorbars correspond to 1σ deviation ( N/N). 3.4.1 Morphology type per Spectrum type/Redshift depen- dence We wanted to see if there are any differences including the distance and the timescale, in other words the redshift factor. For that reason we divided our sample in two smaller datasets. The first one includes the AGN with z 6 0.5 and the second one includes the AGN with z > 0.5. 48 Spiral 16 Spiral Elliptical Elliptical 8 Irregular 14 Irregular 12 6 10 AGN AGN 8 4 6 4 2 2 0 0 Narrow Broad Narrow Broad Spectrum type Spectrum type (a) z 6 0.5 (b) z > 0.5 3.5 Conclusions We studied the AGN properties in clusters as a function of clustercentric distance. To this end, we have used a sample of 81 clusters and 166 AGN up to z ' 2.5 with uniquely well-defined properties for this redshift, which allowed us to accurately determine their radius and to divide the cluster regions into six bins. Specifically we saw that: • The interactions are related to the AGN creation in clusters, according to previous studies. • Despite the fact that in the centre of the clusters, especially in massive clusters (probably because of the ram pressure stripping), there are only few AGN, these AGN present an important merging rate. We notice the same behaviour in the outskirts too, in both cases compared with the background (5r500, 6r500). 14 • Interacting AGN present similar mass distribution (∼ 60%M = 10 Msun & 14 ∼ 40%M < 10 Msun) in all bins, except from bin 2. The 2nd bin, the first 14 “outskirts” region, we notice an inversion (∼ 60%M < 10 Msun & ∼ 40%M = 14 10 Msun)which needs further insight. • The spectra show a relative uniformity in all radii with a dominance of the Narrow line objects. • In terms of the morphology concerning the AGN host galaxy in clusters we see the same results both in the centre and in the outskirts in comparison with the background. Namely most of the hosts are elliptical, less are spiral and there are few irregulars. Nevertheless, we have to point out that despite the fact that this ratio exists in all bins, we observe a significant increase in of spirals and irregulars in 1st and 2nd bin. We conclude that merging/interaction within galaxy clusters is a significant factor in the triggering of AGN that can also affect their morphology. 49 Theoretically, non-axisymmetric perturbations can cause mass inflow during galaxy interactions and merging, and can lead to AGN triggering (Koulouridis et al. 2013; Koulouridis 2014; Ellison et al. 2011; Villforth et al. 2012; Hopkins et al. 2014). Therefore, the detected AGN excess can be explained by a high rate of galaxy merging (e.g. Ehlert et al. 2015) caused by the particular conditions in the cluster outskirts. In more detail, according to the cold dark matter (CDM) paradigm of hierarchical structure formation, many galaxies experience high-density environments before they become cluster members, either as members of smaller groups or by forming within large-scale filaments. As our future plan, we would like to improve our confidence level from 1σ to 3σ and shed some light on • the causes that could suppress the AGN interactions in galaxy clusters • SED fitting of all the sources • X-ray fitting 50 Bibliography [1] Netzer,H., The Physics and Evolution of Active Galactic Nuclei, 2013. [2] Beckmann, V., Shrader, C., Active Galactic Nuclei, 2012. [3] Schneider,P., Extragalactic Astronomy and Cosmology: An Introduction, 2006. [4] Koulouridis et al. The AGN dependence on cluster mass, 2020 [5] E. Koulouridis and I. Bartalucci High density of active galactic nuclei in the outskirts of distant galaxy clusters, 2018 [6] Reiprich et al. Outskirts of Galaxy Cluters , 2013 [7] Umetsu et al. Weak-lensing Analysis of X-Ray selected XXL Galaxy Groups and Clusters with Subaru HSC Data, 2020 [8] Ehlert et al. X-ray bright active galactic nuclei in massive galaxy clusters III: New insights into the triggering mechanisms of cluster AGN, 2018 [9] Kafer et al. Toward the low-scatter selection of X-ray clusters , 2020 [10] Roncarelli et al. Simulated X-ray galaxy clusters at the virial radius: Slopes of the gas density, temperature and surface brightness profiles, 2006 [11] Amendola et al. Dark energy: theory and observations, 2010 [12] Martini et al. An Unexpectedly High Fraction of Active Galactic Nuclei in Red Cluster Galaxies, 2002 [13] Martini et al. The distribution of active galactic nuclei in clusters of galaxies, 2007 [14] Martini et al. The evolution of active galactic nuclei in clusters of galaxies to redshift 1.3, 2009 [15] E.Koulouridis, M.Plionis, V.Chavushyan, D.Dultzin-Hacyan, Y.Krongold, C.Goudis Local and Large Scale Environment of Seyfert Galaxies, 2013. [16] Stefano Bianchi, Roberto Maiolino, Guido Risaliti AGN Obscuration and the Unified Model, 2012. 51