multi-wavelength analysis of active galactic nuclei.pdf

Multi-Wavelength Analysis of ActiveGalactic Nuclei

A dissertation submitted as partial fulfilment

of the

100-hour certificate course

in

Astronomy & Astrophysics

by

Sameer Patel

M.P. Birla Institute of Fundamental Research

Bangalore, India

December 2014

Declaration

I, Sameer Patel, student of M.P. Birla Institute of Fundamental Research, Bangalore,

hereby declare that the matter embodied in this dissertation has been compiled and

prepared by me on the basis of available literature on the topic titled,

Multi-Wavelength Analysis of Active Galactic Nuclei

as a partial fulfillment of the 100 Hour Certificate Course in Astronomy and Astro-

physics, 2014. This dissertation has not been submitted either partially or fully to any

university or institute for the award of any degree, diploma or fellowship.

Date:

Place:

Signature

Director,

M.P. Birla Institute of Fundamental Research,

Bangalore

i

M.P. Birla Institute of Fundamental Research

Bangalore, India

Abstract

Multi-Wavelength Analysis of Active Galactic Nuclei

by Sameer Patel

This dissertation explores the current research methods and analysis adopted for the

study of Active Galactic Nuclei in all wavelengths of the electromagnetic radiation.

Being the most violent objects that one can see in the present Universe, AGNs have been

attributed to emitting radiation in all wavelengths and still exhibit various unexplained

phenomena, alongside with being the probes to the very early Universe. The unification

of the AGN model is also included for completeness, albeit not confirmed in its entirety.

Acknowledgements

I would never have been able to finish my dissertation without help from friends, and

support from the team at MPBIFR, Bangalore.

I would also like to thank Dr. Babu for constantly reminding us to complete the dis-

sertation timely, and Ms. Komala for guiding me to coast through countless papers

online for reference. I would like to thank Rishi Dua, who as a good friend, was always

willing to help me and give his best suggestions, and Aakash Masand, who helped me

correct typographical errors and grammatical mistakes after painfully proofreading the

final draft.

I would also like to thank my parents. They were always supporting me and encouraging

me with their best wishes.

iii

Contents

Declaration i

Abstract ii

Acknowledgements iii

List of Figures vii

List of Tables ix

Abbreviations x

1 Introduction 1

1.1 The History of AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1

1.2 Active Galactic Nuclei . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1

1.3 The Taxonomy of AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . . 2

1.3.1 Seyferts . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 3

1.3.2 Quasars and QSOs . . . . . . . . . . . . . . . . . . . . . . . . . . . 3

1.3.3 Radio Galaxies . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 5

1.3.3.1 Radio Quiet . . . . . . . . . . . . . . . . . . . . . . . . . 6

1.3.3.2 Radio Loud . . . . . . . . . . . . . . . . . . . . . . . . . . 6

1.3.4 Blazars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 8

1.3.4.1 BL Lacerate Objects . . . . . . . . . . . . . . . . . . . . . 8

1.3.4.2 Optically Violent Variable Quasars . . . . . . . . . . . . . 9

1.3.5 LINERs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 9

2 Non-Thermal Processes 12

2.1 Basic Radiative Transfer . . . . . . . . . . . . . . . . . . . . . . . . . . . . 12

2.2 Synchrotron Radiation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 13

2.2.1 Emission by a Single Electron in a Magnetic Field . . . . . . . . . 13

2.2.2 Synchrotron Emission by a Power Law Distribution of Electrons . 14

2.2.3 Synchrotron Self-Absorption . . . . . . . . . . . . . . . . . . . . . 15

2.2.4 Polarization . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15

2.2.5 Synchrotron Sources in AGNs . . . . . . . . . . . . . . . . . . . . . 16

2.2.6 Faraday Rotation . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17

2.3 Thomson Scattering . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 18

iv

Contents

2.4 Compton Scattering . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19

2.4.1 Comptonization . . . . . . . . . . . . . . . . . . . . . . . . . . . . 20

2.4.2 The Compton Parameter . . . . . . . . . . . . . . . . . . . . . . . 21

2.4.3 Inverse Compton Emission . . . . . . . . . . . . . . . . . . . . . . 22

2.4.4 Synchrotron Self-Compton . . . . . . . . . . . . . . . . . . . . . . . 23

2.5 Annihilation and Pair-Production . . . . . . . . . . . . . . . . . . . . . . . 24

2.6 Bremsstrahlung (Free-Free) Radiation . . . . . . . . . . . . . . . . . . . . 26

3 The IR and Sub-mm Regime 27

3.1 History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 27

3.2 Observations and Detection . . . . . . . . . . . . . . . . . . . . . . . . . . 28

3.3 The Dusty Torus . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 29

3.4 IR Spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 31

3.4.1 The 1 µm Minimum . . . . . . . . . . . . . . . . . . . . . . . . . . 32

3.4.2 IR Continuum Variability . . . . . . . . . . . . . . . . . . . . . . . 33

3.4.3 The Submillimeter Break . . . . . . . . . . . . . . . . . . . . . . . 33

4 The Radio Regime 34

4.1 History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 34

4.2 The “Loudness” of AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . 35

4.3 The Fanaroff-Riley Classification . . . . . . . . . . . . . . . . . . . . . . . 36

4.3.1 Fanaroff-Riley Class I (FR-I) . . . . . . . . . . . . . . . . . . . . . 36

4.3.2 Fanaroff-Riley Class II (FR-II) . . . . . . . . . . . . . . . . . . . . 37

4.4 Radio Lobes and Jets . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 39

4.4.1 The Generation of Jets . . . . . . . . . . . . . . . . . . . . . . . . 40

4.4.2 The Formation of Radio Lobes . . . . . . . . . . . . . . . . . . . . 40

4.4.3 Accelerating the Charged Particles in the Jets . . . . . . . . . . . . 42

4.4.4 Superluminal Velocities . . . . . . . . . . . . . . . . . . . . . . . . 43

5 The Optical-UV Regime 44

5.1 History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 44

5.2 Spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 44

5.2.1 The Optical-UV Continuum and the Accretion Disk . . . . . . . . 45

5.3 Observations in the Optical-UV Region . . . . . . . . . . . . . . . . . . . 47

5.4 Discovery by Optical-UV Properties . . . . . . . . . . . . . . . . . . . . . 51

6 The X-Ray Regime 54

6.1 History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 54

6.2 Probing the Innermost Regions . . . . . . . . . . . . . . . . . . . . . . . . 55

6.3 The X-Ray Spectrum of AGNs . . . . . . . . . . . . . . . . . . . . . . . . 56

6.4 Lineless AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59

6.5 The Central Obscuration . . . . . . . . . . . . . . . . . . . . . . . . . . . 61

6.6 Detection and Observations of AGN in X-Rays . . . . . . . . . . . . . . . 62

6.6.1 X-Ray Observations of AGNs . . . . . . . . . . . . . . . . . . . . . 62

6.6.2 Discovery by X-Ray Properties . . . . . . . . . . . . . . . . . . . . 62

7 The γ-Ray Regime 64

7.1 History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 64

Contents

7.2 Gamma-Ray Loud AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . 67

7.3 γ-Ray Properties of Blazars . . . . . . . . . . . . . . . . . . . . . . . . . . 67

8 The Unified Model of AGNs 70

8.1 The Unification . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70

8.2 Absorbed Versus Unabsorbed AGN . . . . . . . . . . . . . . . . . . . . . . 72

8.3 Radio-Loud Versus Radio-Quiet . . . . . . . . . . . . . . . . . . . . . . . . 78

8.4 Breaking the Unification . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83

List of Figures

1.1 The spectrum of NGC 1275 . . . . . . . . . . . . . . . . . . . . . . . . . . 4

1.2 The visible spectrum of Mrk 1157 . . . . . . . . . . . . . . . . . . . . . . . 4

1.3 The visible spectrum of 3C 273 . . . . . . . . . . . . . . . . . . . . . . . . 5

1.4 The total intensity distribution of 3C 338 . . . . . . . . . . . . . . . . . . 7

1.5 The total intensity distribution of 3C 173P1 . . . . . . . . . . . . . . . . . 7

1.6 The X-ray image of 3C 273’s jet . . . . . . . . . . . . . . . . . . . . . . . 8

1.7 The UV spectrum of NGC 4594 . . . . . . . . . . . . . . . . . . . . . . . . 9

1.8 The spread of emission-line galaxies from the SDSS . . . . . . . . . . . . . 10

1.9 Radio luminosity vs. optical luminosity . . . . . . . . . . . . . . . . . . . 11

2.1 Comparison of a synchrotron source with a blackbody source . . . . . . . 15

3.1 Composite spectrum of Type I AGNs . . . . . . . . . . . . . . . . . . . . . 29

3.2 HST image of NGC 4261 . . . . . . . . . . . . . . . . . . . . . . . . . . . 31

3.3 AGN spectrum continuum . . . . . . . . . . . . . . . . . . . . . . . . . . . 32

4.1 VLA map of 3C 449 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 38

4.2 VLA map of 3C 47 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 39

4.3 Electromagnetic outflows from an accretion disk . . . . . . . . . . . . . . 41

4.4 Contour images of Cygnus A’s jet . . . . . . . . . . . . . . . . . . . . . . . 42

4.5 Superluminal motion of M87’s jet . . . . . . . . . . . . . . . . . . . . . . . 43

5.1 Composite optical-UV spectra of AGNs . . . . . . . . . . . . . . . . . . . 46

5.2 General view of a typical optical-UV SED of AGNs . . . . . . . . . . . . . 46

5.3 Broadband SEDs of AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . 48

5.4 Average optical-UV SED for Type I AGNs . . . . . . . . . . . . . . . . . 48

5.5 Spectrum of LLAGN NGC 5252 . . . . . . . . . . . . . . . . . . . . . . . 49

5.6 Comparison of different broad-line profiles in Type I AGNs . . . . . . . . 50

5.7 u-g color of a large number of SDSS AGNs with various redshifts . . . . . 52

5.8 Discovering AGNs by their broadband colours . . . . . . . . . . . . . . . . 53

6.1 Composite AGN spectrum in extreme UV based on FUSE data . . . . . . 57

6.2 Soft X-ray spectrum of NLS1 Arkelian 564 . . . . . . . . . . . . . . . . . . 58

6.3 Composite spectrum of 15 lineless AGNs with large X-ray-to-optical lu-minosity . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 60

7.1 Multiepoch, multiwavelength spectrum of 3C 279 . . . . . . . . . . . . . . 66

8.1 Schematic representation of unified BL Lac phenomenon . . . . . . . . . . 71

8.2 Schematic representation of the unified AGN model . . . . . . . . . . . . 82

vii

List of Figures

8.3 Anticorrelation between X-ray variability amplitude and black hole mass . 84

List of Tables

2.1 Synchrotron sources in AGNs . . . . . . . . . . . . . . . . . . . . . . . . . 16

8.1 The general unification scheme of AGNs . . . . . . . . . . . . . . . . . . . 83

ix

Abbreviations

AGN Active Galactic Nuclei

SMBH Super Massive Black Hole

QSO Quasi Stellar Objects

IRAS Infrared Astronomical Satellite

NLRG Narrow Line Radio Galaxies

BLRG Broad Line Radio Galaxies

WLRG Weak-Emission Line Radio Galaxies

BLR Broad Line Region

SSRQ Steep Spectrum Radio Quasars

FSRQ Flat Spectrum Radio Quasars

FR-I Fanaroff Riley Type I

FR-II Fanaroff Riley Type II

BL Lac BL Lacertae

OVV Optically Violently Variable Quasars

LINER Low Ionization Nuclear Emission-Line Region

LLAGN Low Luminosity Active Galactic Nuclei

SED Spectral Energy Distribution

SF Star Formation

RIAF Radiatively Inaccurate Accretion Flow

SSC Synchrotron Self Compton

BH Black Hole

NIR Near Infrared

MIR Mid Infrared

FIR Far Infrared

RM Rotation Measure

x

Abbreviations

IC Inverse Compton

UV Ultra-Violet

HST Hubble Space Telescope

MHD Magnetohydrodynamics

FWHM Full Width at Half Maximum

S/N Signal to Noise Ratio

NLS1 Narrow Line Seyfert Type I

SDSS Sloan Digital Sky Survey

BAL Broad Absorption Line

BEL Broad Emission Line

XRB X-Ray Binary

SAS Small Astronomy Satellite

OSO Observing Solar Observatory

HEAO High Energy Astronomy Observatory

2MASS 2 Micron All Sky Survey

XMM X-Ray Multi-Mirror Mission

RGS Reflecting Grating Spectrometer

CCD Charge Coupled Device

ESA European Space Agency

NASA National Aeronautics and Space Agency

COSMOS Cosmic Evolution Survey

EW Equivalent Width

HIG Highly Ionized Gas

BAT Burst Alert Telescope

ROSAT Roentgen Satellite

INTEGRAL International Gamma-Ray Astrophysics Laboratory

CGRO Compton Gamma-Ray Observatory

LAT Large Area Telescope

VLBI Very Long Baseline Interferometry

HESS High Energy Spectroscopic System

MERLIN Multi-Element Radio Linked Interferometer Network

VLA Very Large Array

HBLR Hidden Broad Line Region

Abbreviations

OSSE Oriented Scintillation Spectrometer Experiment

EXOSAT European X-Ray Observatory Satellite

PDS Planetary Data System

HBL High-Frequency Peaked BL Lac Objects

LBL Low-Frequency Peaked BL Lac Objects

RMS Root Mean Squared

Chapter 1

Introduction

1.1 The History of AGNs

Unusual activity in the nuclei of galaxies was first recognised by Minkowski and Humason

(Mount Wilson Observatory), when in 1943 they asked a graduate student Carl Seyfert

to study a class of galaxy with an emission spectrum from the compact bright nucleus.

Most normal galaxies show a continuum with absorption lines, but the emission in the

Seyfert galaxies betrayed the presence of hot tenuous gas. In some cases the emission

lines were broad (Type 1 Seyferts) indicating gas moving with high velocities and in

other objects, the emission lines were narrow (Type 2 Seyferts) indicating that the gas

was moving more slowly. In the 1950s, as radio astronomy became a rapidly developing

science a whole new range of discoveries were made in astronomy. Amongst these were

the Radio Galaxies, which appeared to be elliptical galaxies that were inconspicuous at

optical wavelengths but were shown to have dramatically large, prominent lobes at radio

frequencies, stretching for millions of light years from the main galaxy

1.2 Active Galactic Nuclei

The names “active galaxies” and “active galactic nuclei” (AGNs) are related to the main

feature that distinguishes these objects from inactive (normal or regular) galaxies —the

presence of accreting supermassive black holes (SMBHs) in their centers. As of 2011,

there are approximately a million known sources of this type selected by their color and

1

Chapter 1. Introduction

several hundred thousand by basic spectroscopy and accurate redshifts. It is estimated

that in the local universe, at z ≤ 0.1, about 1 out of 50 galaxies contains a fast-accreting

SMBH, and about 1 in 3 contains a slowly accreting SMBH. Detailed studies of large

samples of AGNs, and the understanding of their connection with inactive galaxies and

their redshift evolution, started in the late 1970s, long after the discovery of the first

quasi-stellar objects in the early 1960s. Although all objects containing active SMBH

are now referred to as AGNs, various other names, relics from the 1960s, 1970s, and

even now, are still being used.

The most powerful active galaxies were discovered with radio telescopes in the 1960’s

and named ‘Quasi-Stellar Radio Sources’, later shortened to QSOs or quasars. Their

huge luminosities (∼ 1042−46 erg s−1) could not be attributed to starlight alone, and the

rapid variability observed (from months down to days) implied that the radiation was

emitted from very small volumes with characteristic linear size of the order of light days.

At the time, it was proving difficult to reconcile these two properties. As more detailed

observations were performed it became clear that AGNs were most likely powered by

accretion of matter onto a central SMBH (>105 M⊙).

It is considerable to add that not all galaxies are active. Our Milky-Way is one of the

numerous galaxies that hosts a SMBH at its galactic center (Schodel et al., 2002), with

MSMBH ∼ 4.6 ± 0.7 × 106M⊙ , but is not considered to be an active galaxy due to

the fact that there is no apparent accretion on to the SMBH. On contrary, the central

regions of an AGN are likely not static, but very dynamic and violent.

1.3 The Taxonomy of AGNs

The observational classification of AGNs is not so clear because of observational limi-

tations, heavy source obscuration (in most cases) and usually varying accretion rate on

many orders of magnitude. Classically, an object is classified as an AGN if :-

• It contains a compact nuclear region emitting significantly beyond what is expected

from stellar processes typical of this type of galaxy.

• It shows the clear signature of a non-stellar continuum emitting process in its

center.

2


• Its spectrum contains strong emission lines with line ratios that are typical of

excitation by a non-stellar radiation field.

• It shows line and/or continuum variations.

1.3.1 Seyferts

Owing the name to Seyfert (1943) who was the first to discover these types, the major-

ity of AGN with visible host galaxies fall under this class, known as Seyfert Galaxies.

Seyfert, in his first observation, had reported a small percentage of galaxies had very

bright nuclei that were the source of broad emission lines produced by atoms in a wide

range of ionization states. These nuclei were nearly stellar in appearance (no powerful

telescopes at that time were available).

Today, these are further divided into two more subcategories :-

• Type I Seyferts: Spectra contain very broad emission lines that include both

allowed lines (H I, He I, He II) and narrower forbidden lines (O [III]). They

generally also have “narrow” allowed lines albeit being comparatively broader than

those exhibited by non-active galaxies. The width of these lines is attributed to

Doppler broadening, indicating that the allowed lines originate from sources with

speeds typically between 1000 and 5000 km s−1

• Type II Seyferts: Spectra contain only narrow lines (both permitted and forbid-

den), with characteristic speeds of about 500 km s−1

1.3.2 Quasars and QSOs

The terms Quasar (Quasi Stellar Radio Source) and QSO (Quasi Stellar Object), often

used interchangeably, are scaled up versions of a Type I Seyfert, where the nucleus has

a luminosity MB < −21.5 + 5 log h0 Schmidt & Green (1983). Maarten Schmidt

recognized that the pattern of the broad emission lines of 3C 273 (the first detected

quasar) was the same as the pattern of the Balmer lines of Hydrogen, but were

severely redshifted to z = 0.158 to unfamiliar wavelengths, thus alluding astronomers

from understanding it.

3


Figure 1.1: The spectrum of NGC 1275. The emission features seen at 5057 A and6629 A are [O III]λ 5007 and Hα, respectively.

(Sabra et al., 2000)

Figure 1.2: The visible spectrum of Mrk 1157, a Seyfert 2 galaxy.(Osterbrock, 1984)

4


In 1963, the Dutch astronomer Maarten Schmidt recognized that the pattern of the

broad lines of 3C 273 was the same as the pattern of the Balmer lines of Hydrogen,

only severely redshifted to z = 0.158, hence alluding astronomers from identifying its

spectrum. The continuous spectrum of a quasar may span nearly 15 orders of

magnitude in frequency, very broad compared with the sharply peaked blackbody

spectrum of a star. Quasars emit an excess of UV light relative to stars and so are

quite blue in appearance. This UV excess is indicated by the “big blue bump” in

(nearly) every quasar spectrum. A quasar’s radio emission may come either from radio

lobes or from a central source in its core.

Figure 1.3: The visible spectrum of 3C 273, a Quasar.(Francis et al., 1991)

1.3.3 Radio Galaxies

These galaxies are very luminous at radio wavelengths, with luminosities up to 1039 W

between 10 MHz and 100 GHz. The observed structure in radio emission is determined

by the interaction between twin jets and the external medium, modified by the effects

of relativistic beaming. These are further subdivided into two categories.

5


1.3.3.1 Radio Quiet

Similar in many aspects to Type I Seyferts, these galaxies show both broad and narrow

lines, the only difference being that they are much more luminous than Type I

Seyferts. They are observed in the absence of relativistic jets, which contribute the

most energies in the radio wavelength spectrum.

• Radio Quiet Type I AGNs: These have relatively low-luminosities and therefore

are seen only nearby, where the host galaxy can be resolved, and the

higher-luminosity radio-quiet quasars, which are typically seen at greater

distances because of their relative rarity locally and thus rarely show an obvious

galaxy surrounding the bright central source.

• Radio Quiet Type II AGNs: These include Seyfert II galaxies at low luminosities,

as well as the narrow-emission-line X-ray galaxies (Mushotzky, 1982). The

high-luminosity counterparts are not clearly identified at this point but likely

candidates are the infrared-luminous IRAS AGN (Hough et al., 1991, Sanders

et al., 1989, Wills et al., 1992), which may show a predominance of Type II

optical spectra.

1.3.3.2 Radio Loud

Usually attributed to AGNs with unipolar/bipolar, relativistic jets beaming out of

their centers, the radio emission from radio-loud active galaxies is synchrotron

emission, as inferred from its very smooth, broad-band nature and strong polarization.

This implies that the radio-emitting plasma contains, at least, electrons with

relativistic speeds (Lorentz factors of ∼ 104) and magnetic fields. However,

synchrotron radiation not being unique to radio wavelengths, if the radio source can

accelerate particles to high enough energies, features which are detected in the radio

may also be seen in the infrared, optical, ultraviolet or even X-ray.

• Radio Loud Type I AGNs: These are called Broad-Line Radio Galaxies (BLRG)

at low luminosities and radio-loud quasars at high luminosities, either Steep

Spectrum Radio Quasars (SSRQ) or Flat Spectrum Radio Quasars (FSRQ)

depending on radio continuum shape.

6


• Radio Loud Type II AGNs: Often called Narrow-Line Radio Galaxies (NLRG),

these include two distinct morphological types: the low-luminosity Fanaroff-Riley

type I (Figure 1.4) radio galaxies (Fanaroff & Riley, 1974), which have

often-symmetric radio jets whose intensity falls away from the nucleus, and the

high-luminosity Fanaroff-Riley type II (Figure 1.5) radio galaxies, which have

more highly collimated jets leading to well-defined lobes with prominent hot

spots.

Figure 1.4: The total intensity distribution of 3C 338, a FR I classified AGN.(Ge & Owen, 1994)

Figure 1.5: The total intensity distribution of 3C 173P1, a FR II classified AGN.(Leahy & Perley, 1991)

7


1.3.4 Blazars

Originally named after what was thought to be an irregular, variable star BL Lacertae,

these are AGNs which are characterized by rapid and large-amplitude flux variability

and significant optical polarization. When compared to quasars with strong emission

lines, blazars have spectra dominated by a featureless non-thermal continuum. The

most well known object in this class is the BL Lacertae. Joining the BL Lac objects in

the blazar classification are the optically violently variable quasars (OVVs), which are

similar to the BL Lacs except that they are typically much more luminous, and their

spectra may display broad emission lines. Blazars are AGNs viewed head on and hence

often have jets associated with them (Figure 1.6)

Figure 1.6: The X-ray image of 3C 273’s jet.(”3C273 Chandra” by Chandra X-ray Observatory - NASA. Licensed under Public

domain via Wikimedia Commons)

1.3.4.1 BL Lacerate Objects

BL Lacertae Objects, or BL Lacs for short, are a subclass of blazars that are

characterized by their rapid time-variability. Their luminosities may change by upto

30% in just 24 hours and by a factor of 100 over a longer time period. BL Lacs are also

distinguished by their strongly polarized power-law continua (30% − 40% linear

polarization) that are nearly devoid of emission lines, suggesting that there are very

powerful EM fields at play. BL Lacs, like quasars, are at cosmological distances. Of all

the BL Lacs that have been resolved, 90% of those appear to reside in elliptical

galaxies.

8


1.3.4.2 Optically Violent Variable Quasars

Almost similar to BL Lacs, OVVs are typically much more luminous and may display

broad emission lines in their spectra. The currently best known example of an OVV is

3C 279.

1.3.5 LINERs

LINERs (Low Ionization Nuclear Emission-line Regions) are types of active galaxies

that have very low luminosities in their nuclei, but with fairly strong emission lines of

low-ionization species, such as the forbidden lines of [O I] and [N II]. The Spectra of

LINERs seem similar to the low-luminosity end of the Seyfert II class, and LINER

signatures are detected in many (most of) spiral galaxies in high-sensivity studies.

These low-ionization lines are also detectable in starburst galaxies and in H II regions

and hence it is sometimes difficult to distinguish between LINERs and starburst

galaxies. In the local universe, they are found in about one-third of all galaxies brighter

Figure 1.7: The UV spectrum of NGC 4594 LINER observed using the HST FOS.(Nicholson et al., 1998)

than B = 15.5 mag. This is larger than the number of local high-ionization AGNs by a

factor of 10 or more. Local high-ionization AGNs and LINERs are present in galaxies

with similar bulge luminosities and sizes, neutral hydrogen gas (H I) contents, optical

colors, and stellar masses. Given a certain galaxy type and stellar mass, LINERs are

9


usually the lowest-luminosity AGNs, with nuclear luminosity that can be smaller than

the luminosity of high-ionization AGNs by 1-5 orders of magnitude. An alternative

name for this class of objects is low-luminosity AGNs (LLAGNs). The strongest

Figure 1.8: The spread of emission-line galaxies from the SDSS on one diagnosticdiagram that uses four strong optical emission lines, Hα, Hβ, [O III] λ5007, and [NII] λ6584, to distinguish galaxies that are dominated by ionization from young stars(green points) from those that are ionized by a typical AGN SED (blue points for high-ionization AGNs and red points for low-ionization AGNs). The AGN and SF groupsare well separated, but the division between the two AGN groups is less clear. Thecurves indicate empirical (solid) and theoretical (dashed) dividing lines between AGNs

and star-forming galaxies.(Groves & Kewley, 2008)

optical emission lines in the spectrum of LINERs include [O III] λ5007, [O II] λ3727,

[O I] λ6300, [N II] λ6584, and hydrogen Balmer lines. All these lines are prominent

also in high-ionization AGNs, but in LINERS, their relative intensities indicate a lower

mean ionization state. For example, the [O III] λ5007/Hβ line ratio in LINERs is 3-5

times smaller than in high-ionization Type-II AGNs. Line diagnostic diagrams are

efficient tools to separate LINERs from high-ionization AGNs. One such example is

shown in Figure 1.8. The exact shape of a LINERs SED is still an open issue. In some

sources, it is well represented by the SED shown in Figure 5.3. Such an SED has a

clear deficit at UV wavelengths compared with the spectrum of high-ionization AGNs.

However, some LINERs show strong UV continua and, occasionally, UV continuum

variations, and it is not entirely clear what fraction of the population they represent.

10


Figure 1.9: (left) Radio luminosity vs. optical (B-band) luminosity for various typesof AGNs. (right) The radio loudness parameter R vs. λ (L/LEdd).

(Sikora et al., 2007)

This is related to the issue of Radiatively Inefficient Accretion Flows (RIAFs) and the

relationship between the mass-accretion rate onto the BH and the emitted radiation.

Point-like X-ray sources have been observed in a large number of LINERs. These

nuclear hard X-ray sources are more luminous than expected for a normal population

of X-ray binaries and must be related to the central source. Many LINERs also

contain compact nuclear radio sources similar to those seen in radio-loud

high-ionization AGNs but with lower luminosity comparable to WLRGs (Figure 1.9).

The UV-to-X-ray luminosity ratio in LINERs is, again, not very well known. In

LINERs with strong UV continua, αox is smaller than in low-redshift, high-ionization

AGNs, consistent with the general trend between αox and Lbol. However, αox is not

known for most LINERs because of the difficulty in measuring the UV continuum.

Like other AGNs, LINERs can be classified into Type-I (broad emission lines) and

Type-II (only narrow lines) sources. The broad lines, when observed, are seen almost

exclusively in Hα and hardly ever in Hβ. This is most likely due to the weakness of the

broad wings of the Balmer lines that are difficult to observe against a strong stellar

continuum. Some, perhaps many, LINERs may belong to the category of real Type-II

AGNs —those AGNs with no BLR. The phenomenon is expected to be more common

among low-luminosity sources and hence to be seen in LINERs. Because of all this, the

classification of LINERs is ambiguous, and the relative number of Type-I and Type-II

objects of this class is uncertain even at very low redshift.

11

Chapter 2

Non-Thermal Processes

Much of the electromagnetic radiation emitted by AGNs is very different from a simple

blackbody emission or a stellar radiation source. The general name adopted here for

such processes is non-stellar emission, but the term non-thermal emission is commonly

used to describe such sources. There are several types of non-stellar radiation

processes.

2.1 Basic Radiative Transfer

Describing the interaction of radiation with matter requires the use of three basic

quantities: the first is the specific intensity Iν , which gives the local flux per unit time,

frequency, area, and solid angle everywhere in the source. The second quantity is the

monochromatic absorption cross section, κν (cm−1), which combines all loss

(absorption and scattering) processes. The third quantity is the volume emission

coefficient, jν , which gives the locally emitted flux per unit volume, time, frequency,

and solid angle. The three are combined into the equation of radiative transfer,

dIνds = −κνIν + jν ,

where ds is a path length interval. The first term on the right in this equation

describes the radiation loss due to absorption, and the second gives the radiation gain

due to local emission processes. One usually defines the optical depth element,

dτν = κνds. Hence,

12

Chapter 2. Non-Thermal Processes

dIνdτν

= −Iν + Sν ,

where Sν = jν/κν is the source function. The formal solution of the equation of

transfer depends on geometry. For a slab of thickness τν in a direction perpendicular

to the slab, it is

Iν(τν) = Iν(0)e−τν +

τν∫

0

e−(τν−t)Sν(t)dt.

For any other direction θ, both τν and dt must be divided by cos θ.

The general equation of radiative transfer is difficult to solve and requires numerical

techniques. However, there are simple cases in which the solution is straightforward.

In particular, the case of a slab and a constant source function that is independent of

τν allows a direct integration and gives the following solution:

Iν = Iν(0)e−τν + Sν(1 − e−τν ).

For an opaque source in full thermodynamic equilibrium (TE), the optical depth is

large, and both Iν and Sν approach the Planck function

Bν(T ) = 2hν3/c2

ehν/kT−1

2.2 Synchrotron Radiation

2.2.1 Emission by a Single Electron in a Magnetic Field

Considering an electron of energy E that is moving in a uniform magnetic field B of

energy density uB = B2/8π, the energy loss rate, −dE/dt, which is also the power

emitted by the electron, P, is given by

P = 2σT cγ2β2uB sin2 α,

where σT is the Thomson cross section,

c is the speed of light,

13


γ = E/mc2 is the Lorentz factor,

β = v/c, and

v is the speed of the electron.

The angular term sin2 α reflects the direction of motion, where α is the pitch angle

between the direction of the motion and the magnetic field. Averaging over isotropic

pitch angles gives

P = (4/3)σT cγ2β2uB.

The radiation emitted by a single electron is beamed in the direction of motion. The

spectral energy distribution (SED) of this radiation is obtained by considering the gyro

frequency of the electrons around the field lines (ωB = eB/γmec) and the mean interval

between pulses (2π/ωB). The calculation of the pulse width is obtained by considering

the relativistic time transformation between the electron frame and the observer frame.

This involves an additional factor of γ2. Thus, the pulse width is proportional to γ−3

or, expressed with the Larmor angular frequency, ωL = eB/mec (which differs from ωB

by a factor of γ), to γ−2. Fourier transforming these expressions gives the mean

emitted spectrum of a single electron, Pγ , which peaks at a frequency near γ2ωL.

2.2.2 Synchrotron Emission by a Power Law Distribution of Electrons

Assuming now a collection of electrons with an energy distribution n(γ)dγ that gives

the number of electrons per unit volume with γ in the range γ − (γ + dγ), the emission

coefficient due to the electrons is obtained by summing Pγ(γ) over all energies:

jν = 14π

∞∫

1

Pν(γ)n(γ)dγ.

There is no general analytical solution to this expression since n(γ) can take various

different forms. However, there are several cases of interest where n(γ) can be

presented as a power law in energy:

n(γ)dγ = n0γ−pdγ.

The additional assumption that all the radiation peaks around a characteristic

frequency, γ2νL, where νL is the Larmor frequency, gives the following solution for jν :

14


4πjν = 23σTn0uBν

−1L

(

ννL

)− p−1

2

.

Figure 2.1: A comparison of a synchrotron source with p = 2.5 (solid line) and a 105

K blackbody source (dotted line).(Netzer, 2013)

2.2.3 Synchrotron Self-Absorption

The source of fast electrons can be opaque to its own radiation. This results in a

significant modification of the emergent spectrum especially at low frequencies, where

the opacity is the largest. It can be shown that in this case,

κν ∝ ν−p+4

2 ,

that is, the largest absorption is at the lowest frequencies. Using the equation of

radiative transfer for a uniform homogeneous medium, we get the solution at the large

optical depth limit, Iν ∝ ν5/2, which describes the synchrotron SED at low energies.

This function drops faster toward low energies than the low-energy drop of a blackbody

spectrum (Iν ∝ ν2). The overall shape of such a source is shown in Figure 2.1.

2.2.4 Polarization

Synchrotron radiation is highly linearly polarized. The intrinsic polarization can reach

70%. However, what is normally observed is a much smaller level of polarization,

15


Source B (G) ν (Hz) γ tcool (yr) E (erg)

Extended radio sources 10−5 109 104 107 1059

Radio jets 10−3 109 103 104 1057

Compact jets 10−1 109 102 101 1054

BH magnetosphere 104 1018 104 10−10 1047

Table 2.1: Synchrotron Sources in AGNs.(Netzer, 2013)

typically 3-15%. This indicates a mixture of the highly polarized synchrotron source

with a strong non-polarized source. For AGNs, especially radio-loud sources, this

polarization is clearly observed. There is also a correlation between high-percentage

polarization and large-amplitude variations. AGNs showing such properties go under

the name blazars. In the NIR-optical-UV spectrum of radio-loud AGNs, the region

around 1 µm shows most of the polarization. The percentage polarization seems to

drop toward shorter wavelengths, in contrast to what is expected from a pure

synchrotron source. This is interpreted as an indication of an additional thermal,

non-polarized source at those wavelengths.

2.2.5 Synchrotron Sources in AGNs

It is thought that most of the non-thermal radio emission in AGNs is due to

synchrotron emission. There are various ways to classify such radio sources using the

slope, (p− 1)/2, and the break frequency below which it is optically thick to its own

radiation. Table 2.1 gives a summary of the properties of several observed and

expected synchrotron sources in AGNs. It includes the typical strength of the

magnetic field, B (in gauss), the Lorentz factor, γ, and the total energy generated in

the source, E, which is obtained by integrating uB over the volume of such sources.

The table also shows the typical cooling time of the source, tcool, which is a

characteristic lifetime defined by

tcool = γmec2

P≃ 5 × 108B−2γ−1sec.

16


2.2.6 Faraday Rotation

Michael Faraday discovered in 1845 that the angle of polarization of an

electromagnetic wave changes when the wave is sent through a medium with a

magnetic field. The so-called Faraday rotation can also affect the synchrotron

emission. Faraday rotation can be understood as the different effect the magnetized

plasma has on the left and right circularly polarized light. Depending on the

orientation with respect to the magnetic field, the components will “see” a different

refractive index. Thus, the phase velocity of the two components will be affected

slightly differently and lead to a shift of their relative phases. This causes the plane of

polarization to rotate, depending on how strong the magnetic field is and what

distance the wave has to travel through the plasma. A similar effect is also observed

with linearly polarized light. Once the linearly polarized synchrotron light is emitted

and travelling towards the observer, it can pass through magnetized material causing

Faraday rotation. This can be the emitting plasma itself, or any magnetized gas along

the line of sight. In astrophysical applications, one can simplify the problem by

considering only free electrons in magnetic fields.

The amount of rotation in the polarization angle depends on the magnetic field

strength and density of the electrons along the line of sight, but also on the frequency

of the electromagnetic wave one observes:

∆θ = λ2RM.

Here, λ is the wavelength of the polarized radiation, and RM is the rotation measure

which is a function of the electron density ne and of the component of the magnetic

field B|| parallel to the line of sight:

∆θ = λ2 e3

2πm2c4

∫

ne(s)B||(s)ds.

Thus, the rotation is larger for low frequencies. This is because the frequency of the

wave is much larger than the gyro-frequency of the electron. The closer the light and

the electron are to a resonant state, and thus the larger the energy transfer from the

wave to the electron. The light from extragalactic sources will not only have to cross

the intergalactic medium, but the interstellar medium of our galaxy as well on its path

17


to the observer. The magnetic field along the line of sight will not be constant, and

importantly, it will not be of the same orientation throughout the path of light. To

determine the net effect of Faraday rotation, it is necessary to measure polarization at

closely spaced frequency interval over many frequencies. Because the rotation affects

the high frequency the least, the best way to get an estimate of the intrinsic

polarization of a synchrotron source is to measure at high frequencies.

2.3 Thomson Scattering

Thomson scattering describes the non-relativistic case of an interaction between an

electromagnetic wave and a free charged particle. The effect was first describe by Sir

Joseph John Thomson, who discovered the electron when studying cathode rays in the

late nineteenth century. The process can be understood as elastic or coherent

scattering, as the photon and the particle will have the same energy after the

interaction as before. For this process of the energy E of the photon has to be much

smaller than the rest energy of the particle:

E = hν ≪ mc2.

Another requirement for Thomson scattering is that the particle must be moving at

non-relativistic speed (v ≪ c). In the classical view of this process, the incoming

photon is absorbed by the particle with charge q, which is set into motion and then

re-emits a photon of the same energy.

Using the classical electron radius r0 = q2/mc2 (Bohr radius), the differential

cross-section of this elastic scattering process can be written as

dσdΩ = 1

2(1 + cos2 θ)r20.

This is symmetric with respect to the angle θ, thus the amount of radiation scattered

in the forward and backward direction is equal. The total cross-section is then given by

σT = 2π

π∫

0

dσdΩ sin θdθ = 8π

3 r20 = 8π3

(

q2

mc2

)2.

18


In the case of electrons, this gives a Thomson cross-section of σT ≃ 6.652 × 10−25 cm2.

The cross-section for a photon scattering on a photon is a factor of

(mp/me)2 ≃ 3.4 × 106 smaller.

Since in the classical view of this process, the electron has no preferred orientation, the

cross-section is independent of the incoming electromagnetic wave. The polarization of

the scattered radiation depends, however, on the polarization of the incoming photon

wave. Unpolarized radiation becomes linearly polarized in the Thomson scattering

process with the degree of polarization being

Π = 1−cos2 θ1+cos2 θ

.

Therefore, polarization of the observed emission can be a sign that the emergent

radiation has been scattered.

Thomson scattering is important in may astrophysical sources. Any photon which will

be produced inside a plasma can be Thomson scattered before escaping in the

direction of the observer. The chance for the single photon to be Thomson scattered

and how many of the photons will be scattered out of or into the line of sight is

quantified in terms of the optical depth τ of the plasma:

τ =∫

σTnedx,

where ne is the electron density, and dx is the differential line element. The mean free

path λT of the photon, that is, the mean distance traveled between scatterings will

thus be λT = (σTne)−1.

2.4 Compton Scattering

The interaction between an electron and a beam of photons is described by the

classical Compton scattering theory. For stationary or slow electrons, one uses energy

and momentum conservation to obtain the relationship between the frequencies of the

coming (ν′

) and scattered (ν) photons. If ~nν and ~nν′ are unit vectors in the directions

of these photons, and cos θ = ~nν · ~nν′ , we get

19


ν = mec2ν′

mec2+hν′(1−cos θ)

.

For non-relativistic electrons, the cross section for this process is given by

dσdΩ = 1

2r2e [1 + cos2 θ],

where re = e2/mec2 is the classical electron radius. Integrating over angles gives the

Thomson cross section, σT . In the high-energy limit, the cross section is replaced by

the Klein-Nishina cross section, σK−N , which is normally expressed using ǫ = hν/mec2.

The approach to the low-energy limit is given roughly by

σK−N ∼ σT (1 − 2ǫ),

and for ǫ ≫ 1,

σK−N ∼ 38σTǫ

[

ln 2ǫ + 12

]

.

2.4.1 Comptonization

The term Comptonization refers to the way photons and electrons reach equilibrium.

The fractional amount of energy lost by the photon in every scattering is

∆νν ≃ − hν

mec2= −ǫ.

Considering a distance r from a point source of monochromatic luminosity Lν in an

optically thin medium where the electron density is Ne, the flux at this location is

Lν/4πr2, and the heating due to Compton scattering is

HCS =

∫

Lν4πr2

NeσT

[

hνmec2

]

dν.

The cooling of the electron gas is the result of inverse Compton scattering. Like

Compton scattering, this process is a collision between a photon and an electron,

except that in this case, the electron has more energy that can be transfered to the

radiation field. In this case, the typical gain in the photon energy is a factor of γ2

larger than the one considered earlier. This factor is obtained by first transforming to

the electron’s rest frame and then back to the laboratory frame. If x is the fraction of

the electron energy kT which is transferred to the photon,

20


⟨

∆νν

⟩

= x kTemec2

,

where Te is the electron temperature. Using this terminology, one can write the cooling

term for the electron gas as

CCS =

∫

Lν4πr2

NeσT

[

xkTemec2

]

dν.

A simple thermodynamical argument suggests that if Compton heating and Compton

cooling are the only heating-cooling processes, and if the radiation field is given by the

Planck function (Lν = Bν), the equilibrium requirement, HCS = CCS , gives x = 4.

Because this is a general relation between a physical process and its inverse, the result

must also hold for any radiation field.

The radiation field in luminous AGNs can be very intense, and the energy density of

the photons normally exceeds the energy density due to electrons. The requirement

HCS = CCS gives, in this case, a Compton equilibrium temperature of

TC = hν4k ,

where the mean frequency, ν, is defined by integrating over the SED of the source,

ν =∫νLνdν∫Lνdν

.

2.4.2 The Compton Parameter

The emitted spectrum of thermal and non-thermal radiation sources that are

embedded in gas with a thermal distribution of velocities is modified due to Compton

and inverse Compton scattering. For high-energy electrons, inverse Compton is the

dominant process, and the resulting collisions will up-scatter the photon energy. The

emergent spectrum is modified, and its spectral shape will depend on the original

shape, the electron temperature, and the Compton depth of the source, which

determines the number of scattering before escape. Considering an initial photon

energy of hνi and the case of thermal electrons with temperature Te such that

hνi ≪ 4kTe, the scattering of such photons by a fast electron will result in energy gain

21


per scattering (inverse Compton scattering). The photon continues to gain energy,

during successive scatterings, as long as hνi ≪ 4kTe. If the final photon energy is hνf ,

and the number of scatterings is N , we get

hνf ≃ hνie

[

N 4kTemec2

]

.

For a medium with Compton depth τe, the mean number of scatterings is roughly

max(τe, τ2e ). Using this, one can define a Compton parameter y,

y = max(τe, τ2e )[

4kTemec2

]

,

such that

hνf ∼ hνiey.

The factor ey is an energy amplification factor. For y > 1, one is in the regime of

unsaturated inverse Comptonization. For y ≫ 1, the process reaches a limit where the

average photon energy equals the electron thermal energy. This is saturated Compton

scattering.

2.4.3 Inverse Compton Emission

An important example is the case of a source whose spectrum is due to scattering of

“soft” photons onto relativistic electrons. Again, considering first the typical energy

following a single scattering and then averaging over the energy distribution of the

photons and electrons, a simple way to estimate the power emitted in the preceding

process is to consider a beam of photons with number density nph and mean energy

before scattering hν0. The energy density of these photons is nphhν0, and the energy

flux of photons incident on a stationary electron is curad = nphhν0c. The mean energy

after scattering, hν, is larger than the mean energy before scattering by a factor of

order γ2. In the rest frame of the electron, the process can be considered as a simple

Thomson scattering with radiated power given by the classical expression

P = σT curad. Thus, the simple Lν ∝ ν(p1)/2 estimate for the laboratory frame emitted

power is P = γ2σT curad. A more accurate derivation of the emitted power must take

22


into account the scattering angle and its transformation between frames. The final

expression in this case is

P = (4/3)σT cγ2β2urad,

which differs from the simple estimate by a factor of order of unity.

The expression for the power emitted due to inverse Compton (IC) scattering is

basically identical to the power emitted by synchrotron radiation, except that the

energy density of the magnetic field, uB, was replaced by the energy density of the

radiation field, urad. Thus, the mean power of the two processes, assuming they take

place in the same volume of space, is simply uB/urad. Also, for the same volume of

space, the energy distribution of the relativistic electrons is given by the same

power-law function used in the synchrotron case, n(γ) ∝ γ−p. Thus, one also gets a

similar dependence of the monochromatic luminosity on the parameter p:

Lν(IC) ∝ ν−(p−1)/2.

2.4.4 Synchrotron Self-Compton

In a compact synchrotron source, the emitted photons can be inverse Compton

scattered by the relativistic electrons that emit the synchrotron radiation. This gives

the photon a big boost in energy. The emergent radiation is synchrotron self-Compton

(SSC) emission. The flux emitted by this process can be calculated by integrating over

the synchrotron radiation spectrum and the electron velocity distribution. To a good

approximation, the resulting spectral index is identical to the spectral index of the

synchrotron source. The synchrotron self-Compton process can repeat itself, in the

same source, by additional scattering of the emergent photons, which results in an

additional boosting, by a factor γ2, to the photons. The natural limit for the process is

when the scattered photon energy extends into the γ-ray and the condition of

hνγ ≪ mec2 (the condition for no Compton recoil of the electron) no longer holds. At

this limit, the resulting radiation density decreases dramatically.

23


2.5 Annihilation and Pair-Production

The observations of γ-ray jets in many AGNs suggest that, under some conditions, the

density of high-energy photons is large enough to result in efficient pair production and

a concentration of both electrons and positrons in some parts of the central source.

Under these conditions, energetic γ-ray photons, with energies much above the rest

energy of the electron, can react with lower-energy photons to create electron-positron

pairs. Short-time-scale variations of the X-ray spectrum, in the lower-luminosity AGN,

indicate extremely small dimensions; γ-ray photons that are associated with the X-ray

source would not be able to escape these regions and would create electron-positron

pairs. Likely locations where such processes take place are in the corona of the central

accretion disk or inside the γ-ray jet.

The process of pair-production and its reverse process (for e− − e+ pair) is given by

e− + e+ γ + γ.

Considering the interaction between a γ-ray photon with frequency νγ , above the rest

mass frequency of the electron, with an X-ray photon of frequency νX below this

frequency and using the notation of unit vectors for the photons, one can write the

threshold frequency for pair production as

vγ =(

mec2

h

)22

νX(1−~nγ ·~nX) .

The γγ cross section is given by

σγγ = 316σT (1 − β2)

[

(3 − β4) ln(

1+β1−β

)

− 2β(2 − β2)]

,

where the value of β for the electron and the positron is measured in the center of

momentum frame. The typical value of σγγ near threshold is ∼ 0.2σT , and it declines

with frequency as ν−1γ .

The size of the radiation source, R, plays an important role in determining the optical

depth of the source and hence the probability of pair-production taking place. This

dependence is usually described by defining a compactness parameter for the γ-ray

source, lγ , using the source size and its luminosity, Lγ . There is an equivalent

24


compactness parameter for the X-ray source, lX . Assuming that the typical γ-ray

photon energy is ǫ = mec2 and that the photon number density is

Nγ =Lγ

4πR2cǫ.

The mean free path of the photons for the pair-production is

λγγ =(

NγσT)−1

and for unit optical depths, R ≈ λ, which gives

LγσT

4πmec3R≈ 1.

This leads to the following expression for the compactness parameter:

lγ =LγσT

4πmec3R,

which is equivalent to the pair production optical depth of the source. In principle, lγ

can be measured from the variability time scale of the γ-ray source. In reality,

however, this is difficult to measure and is occasionally replaced by lX and the X-ray

variability time scale. When lX ≫ hνX/mec2, it will be difficult for the γ-rays to

escape the source without creating pairs.

The rate of the inverse process, pair annihilation, in the non-relativistic limit is

independent of temperature and is roughly 0.4 NeσT c per unit volume, where Ne is the

combined electron-positron density. In a steady state, pair production is balanced by

annihilation,

mec3lγ4πσTR2hνγ

∼= 0.4NeσTc,

where lγ is the compactness parameter for those γ-ray photons for which the source is

optically thick to pair production. This equation can be solved for the mean Thomson

depth in the source, τT . For large τT , the electrons and positrons thermalize because

their interaction time is short compared with the annihilation time. In AGN gas,

where the conditions allow this thermalization, the temperature of the hot, Compton

thick pair plasma can reach 109 K. Such gas can contribute to the observed

high-energy spectrum. It can up-scatter soft (UV) emitted photons and even produce

some free-free electron-positron radiation.

25


2.6 Bremsstrahlung (Free-Free) Radiation

Free-free radiation, formally, is thermal radiation. However in the case of AGNs, the

spectral shape is very different from that of a blackbody. The free-free emissivity due

to ion i of an element of charge Z whose number density is Ni is given by

4πjν = 6.8 × 10−38Z2T−1/2e NeNigff (ν, Te, Z)e−hν/kTe ,

where gff (ν, Te, Z) is the velocity-averaged Gaunt factor, which accounts for

quantum-mechanical effects. This factor is always of the order unity and can change

slightly with frequency, in particular, at X-ray energies gff ∝ ν−0.1. The

Bremsstrahlung radiation extends over a large range of energies and resembles, over

most of this range, a very flat (small spectral index) power law. One can integrate the

free-free emissivity over frequencies to obtain the total energy per unit volume per

second, Cff , where C indicates that this is also the cooling rate due to free-free

emission. The integration gives

Cff = 1.42 × 10−27Z2T1/2e NeNigffNeNi erg s−1 cm−3,

where gff is now the frequency average of the velocity-averaged Gaunt factor. This is

typically in the range 1.1-1.5.

26

Chapter 3

The IR and Sub-mm Regime

3.1 History

The use of IR techniques to measure AGN continua started in the 1970s with the

advent of the first sensitive IR detectors (Low & Kleinmann, 1968). However, the IR

“colours” of Seyfert galaxies are only subtly different than those of normal galaxies

(Kuraszkiewicz et al., 2003), and the equivalent widths of the IR lines are not sufficient

to use as a finding mechanism. Thus, IR color surveys can have a large fraction of

”false” AGN, unless great care is taken.

The first large-scale attempt to find AGN in the IR was based on Infrared

Astronomical Satellite (IRAS) data. de Grijp et al. (1987) showed that AGN had

systematically different 60 µm / 25 µm colours than normal galaxies. An alternative

approach Spinoglio & Malkan (1989) was to obtain optical spectra of every IR-selected

galaxy. This was a follow-up of the idea of Huchra & Burg (1992) to obtain optical

spectra of every optically-selected galaxy, but was not really a survey technique. The

latest use of the IR to find active galaxies is with the Two Micron All Sky Survey

(2MASS; Cutri et al. (2002)). In this survey, ∼ 60% of the objects with

J - K > 2 are found to have the optical properties of AGN. This selection criterion is

bootstrapped by using the near-IR colors of known radio and optically-selected AGN

(Elvis et al., 1994), and thus will tend to find objects with similar properties. The

large space density of these IR-selected objects makes them a major contributor to the

AGN population.

27

Chapter 3. The IR and Sub-mm Regime

The far-IR (FIR) band of thousands of AGNs has been 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. Broadband images with much

improved spatial resolution are now available between 70 and 500 µm. Systematic

surveys have already produced high-quality photometry of hundreds of AGNs and their

host galaxies, up to redshift of 5 and beyond. Lower-sensitivity, high-resolution

spectroscopy over the FIR range is also provided by the Herschel instruments.

3.2 Observations and Detection

Most of the emission in the NIR and MIR bands is due to secondary dust emission.

“Secondary” in this context refers to emission by cold, warm, or hot dust grains that

are heated by the primary AGN radiation source.“Primary” refers to radiation that is

the direct result of the accretion process itself. The temperature of the NIR- and

MIR-emitting dust is between 100 and 2000 K. The dimensions of the dusty structure

emitting this radiation, in intermediate luminosity AGNs, is of order 1 pc. Most of the

thermal FIR emission is thought to be due to colder dust that is being heated by

young stars in large star-forming regions in the host galaxy. In powerful radio sources,

at least part of the FIR emission is due to non-thermal processes much closer to the

center. Broad and narrow emission lines are seen in the NIR-FIR part of the spectrum

of many AGNs. They are thought to originate in the broad- and narrow-line regions.

A very important aspect techniques which use one IR band or a combination of two IR

bands is the ability to detect highly obscured (Compton thick) AGNs. A large fraction

of such objects, especially at high redshift, do not show detectable 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 based on their optical and X-ray

continuum properties. However, their mid-IR (MIR) spectrum is dominated by warm

dust emission, the result of the heating of the central torus by the central source. A

luminosity ratio like L(24 µm)/L(R), where R is the red optical band, will be much

larger in such sources compared with inactive galaxies because the AGN light is

heavily obscured at the R-band. Spectroscopic follow-up of such objects can be used to

look for the unique emission-line spectrum of the AGN. Indeed, systematic searches in

uniformly scanned Spitzer fields reveal a large number of Compton thick AGNs.

28


Figure 3.1 shows a composite 0.3-30 µm spectrum of intermediate-luminosity type I

AGNs. The emission longword of 1 µm is due primarily to secondary radiation from

dust. The dip at 1 µm is due to the decline of the disk-produced continuum on the

short-wavelength side and the rise of the emission due to hot dust on the other side.

Figure 3.1: A composite spectrum of type-I AGNs covering the range 0.340 µm. Theobservations were obtained by several ground-based telescopes and Spitzer and were

normalized to represent a typical intermediate-luminosity source.(Netzer, 2013)

3.3 The Dusty Torus

Dust is the cornerstone of the unification theory of active galactic nuclei (AGNs).

Essentially, all types of AGNs are surrounded by an optically thick dust torus and are

basically the same object but viewed from different lines of sight (Antonucci, 1993,

Urry & Padovani, 1995). The large diversity in the observational properties of AGNs

(eg., optical emission-line widths and X-ray spectral slopes) is simply caused by the

viewing-angle-dependent obscuration of the nucleus: those viewed face-on are

un-obscured (allowing for a direct view of their nuclei) and recognized as Type I

Seyferts, while those viewed edge-on are Type II Seyferts, with most of their central

engine and broad line regions being hidden by the obscuring dust.

29


Apparently, key factors in understanding the structure and nature of AGNs are

determining the geometry of the nuclear obscuring torus around the central engine and

the obscuration (ie., extinction, a combination of absorption and scattering) properties

of the circumnuclear dust. An accurate knowledge of the dust extinction properties is

also required to correct for the dust obscuration in order to recover the intrinsic

optical/UV spectrum of the nucleus from the observed spectrum and to probe the

physical conditions of the dust-enshrouded gas close to the nucleus.

The presence of an obscuring dust torus around the central engine was first indirectly

indicated by the spectropolarimetric detection of broad permitted emission lines

(characteristic of Type I Seyferts) scattered into our line of sight by free electrons

located above or below the dust torus in a number of Type II Seyferts (Heisler et al.,

1997, Tran, 2003) Direct evidence for the presence of a dust torus is provided by IR

observations. The circumnuclear dust absorbs the AGN illumination and re-radiates

the absorbed energy in the IR. The IR emission at wavelengths longward of λ > 1 µm

accounts for at least 50% of the bolometric luminosity of Type II Seyferts. For Type I

Seyferts, ∼ 10% of the bolometric luminosity is emitted in the IR. A near-IR “bump”

(excess emission above the ∼ 2 − 10 µm continuum), generally attributed to hot dust

with temperatures around ∼ 1200-1500 K (near the sublimation temperatures of

silicate and graphite grains), is seen in a few Type I Seyferts (Barvainis, 1987,

Rodrıguez-Ardila & Mazzalay, 2006). Direct imaging at near- and mid-IR wavelengths

has been performed for several AGNs and provides constraints on the size and

structure of the circumnuclear dust torus (Elitzur, 2006). Spectroscopically, the 10 µm

silicate absorption feature and the 3.4 µm aliphatic hydrocarbon absorption feature

are widely seen in heavily obscured Type II Seyferts; in contrast, the 10 µm silicate

emission feature has recently been detected in a number of Type I Seyferts.

To properly interpret the observed IR continuum emission and spectroscopy as well as

the IR images of AGNs, it requires a good understanding of the absorption and

emission properties of the circumnuclear dust. To this end, one needs to know the

composition, size, and morphology of the dust - with this knowledge, one can use Mie

theory (for spherical dust) to calculate the absorption and scattering cross sections of

the dust from X-ray to far-IR wavelengths, and then calculate its UV/optical/near-IR

obscuration as a function of wavelength, and derive the dust thermal equilibrium

temperature (based on the energy balance between absorption and emission) as well as

30


its IR emission spectrum. This will allow us to correct for dust obscuration and

constrain the circumnuclear structure through modeling the observed IR emission and

images. The former is essential for interpreting the obscured UV/optical emission lines

and probing the physical conditions of the central regions; the latter is critical to our

understanding of the growth of the central SMBH.

However, little is known about the dust in the circumnuclear torus of AGNs. Even our

knowledge of the best-studied dust - the Milky Way interstellar dust - is very limited.

Figure 3.2: A HST image of the gas and dust disk in the active galactic nucleus ofNGC 4261.

(”Ngc4261” by Clh288 at en.wikipedia. Licensed under Public domain via WikimediaCommons)

3.4 IR Spectra

The value of spectral index (α) is (almost) constant in the IR region of the spectrum of

an AGN, evident from Figure 3.3. The thermal IR “bump” is due to the emission from

warm (T . 2000 K) dust grains.

31


Figure 3.3: A depiction of the typical features in the continuum observed for manyAGNs.

(Tengstrand et al., 2009)

3.4.1 The 1 µm Minimum

The existence of the IR bump longward of 1 µm has led many authors to conclude that

this emission must be thermal, as the required temperatures are in the right range

(T . 2000 K) for hot dust in the nuclear regions. Sanders et al. (1988) have shown

that a minimum in the SED at ∼ 1 µm is a general feature of AGNs. The hottest dust

has a temperature of ∼ 2000 K; at higher temperatures, dust grains sublimate. This

upper bound of the temperature explains the constancy of the frequency where the

NIR spectrum is the weakest, ie., at the Wien cut-off at a 2000 K blackbody.

One can define a ‘sublimation radius’ as the minimum distance from the AGN at

which grains of a given composition can exist. The dust grains closest to an AGN

probably are graphite rather than silicate, as graphite has a higher sublimation

temperature. The sublimation radius for graphite grains is

r = 1.3L1/2uv46T

−2.81500 pc,

32


where Luv46 is the central source UV luminosity in units of 1046 erg s−1, and T1500 is

the grain sublimation temperature in units of 1500 K (Barvainis, 1987).

3.4.2 IR Continuum Variability

Clear evidence that the hot dust scenario for the origin of the IR continuum has some

merit has been provided by the IR continuum variability characteristics. Unlike

UV/optical variability with little if any time delay, the IR continuum shows the same

variations as the UV/optical continuum, but with a significant time delay. This is

interpreted as a light-travel effect which occurs because of the separation between the

UV/optical and IR continuum-emitting regions; whereas the UV/optical emission

arises in a very compact region, the IR emission arises in dust that is far away from

the central source. The variations occur as the emissivity of the dust changes in

response to the UV/optical continuum that heats it. Within the sublimation radius,

dust is destroyed. Farther out, however, it survives and is heated by the UV/optical

radiation from the central source to approximately the equilibrium blackbody

temperature. The IR continuum arises as this energy is re-radiated by the dust. In the

FIR, the only AGNs that are found to vary are radio-loud sources.

3.4.3 The Submillimeter Break

Observations of the FIR to sub-mm portion of AGN spectra have been made in a

limited number of cases (Chini et al., 1989, Edelson & Malkan, 1987, Hughes et al.,

1993). These observations show that the sub-mm SED decreases rather sharply as one

goes to longer wavelengths, so abruptly that in at least a few cases the spectral index

longward of the sub-mm break must be less than the value of -2.5 expected in the case

of a synchrotron self-absorbed spectrum (ie., Fν ∝ v5/2). At these long wavelengths, a

thermal spectrum can produce a cut-off this sharp because the emitting efficiency of

small grains is a sensitive function of frequency, Qν ∝ νγ , typically with γ ≈ 2 (Draine

& Lee, 1984) so the emitted spectrum can have a very strong frequency dependence,

Fν ∝ ν2+γ .

33

Chapter 4

The Radio Regime

4.1 History

The discovery of radio galaxies preceded the optical discovery of AGNs. It goes back to

the late 1940s and the early 1950s. 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 dimensions 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.

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 correlating 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. The

positional accuracy of optical and radio telescopes is one arcsec or better, and there is

hardly any problem in verifying that the radio and optical emitters are one and the

same source. Most of the early AGN samples were discovered in this way. A

well-known example is the 3C radio sample, which includes some of the most powerful

radio-loud, early-discovered AGNs such as 3C 48 and 3C 273.

34

Chapter 4. The Radio Regime

4.2 The “Loudness” of AGNs

Like optically classified AGNs, there are broad-line radio galaxies (BLRGs), the

equivalent of the Type I sources; narrow-line radio galaxies (NLRGs), the

spectroscopic 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. Hence, usually, the “radio

loudness” parameter, R, is used to separate radio-loud from radio-quiet AGNs. R is a

measure of the ratio of radio (5 GHz) to optical (B-band) monochromatic luminosity,

R = Lν(5 Ghz)

Lν(4400A)= 1.5 × 105 L(5 Ghz)

L(4400A),

where L(5 Ghz) and L(4400 A) represent the value of λLλ at those energies. The

dividing line between radio-loud and radio-quiet AGNs is usually set at R = 10.

Statistics 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 point-like 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,

the spectrum is represented well by a single power law, Fν ∝ ν−αR . 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 variations, 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.

35


4.3 The Fanaroff-Riley Classification

It was first noticed by Fanaroff & Riley (1974) that the relative positions of regions of

high and low surface brightness in the lobes of extragalactic radio sources are

correlated with their radio luminosity. This conclusion was based on a set of 57 radio

galaxies and quasars, from the complete 3CR catalogue, which were clearly resolved at

1.4 GHz or 5 GHz into two or more components. Fanaroff and Riley divided this

sample into two classes using the ratio RFR of the distance between the regions of

highest surface brightness on opposite sides of the central galaxy or quasar, to the

total extent of the source up to the lowest brightness contour in the map. Sources with

RFR < 0.5 were placed in Type I (FR-I) and sources with RFR > 0.5 in Type II

(FR-II). It was found that nearly all sources with luminosity

L(178 Mhz) . 2 × 1025 h−2100 W Hz−1 Sr−1

were of Type I (FR-I) while the brighter sources were nearly all of Type II (FR-II).

The luminosity boundary between them is not very sharp, and there is some overlap in

the luminosities of sources classified as FR-I or FR-II on the basis of their structures.

For a spectral index of α ≃ 1 the dividing luminosity at 5 GHz is

L(5 Ghz) . 7 × 1023 h−2100 W Hz−1 Sr−1

At high frequencies, the luminosity overlap between the two classes can be as much as

two orders of magnitude. Various properties of sources in the two classes are different,

which is indicative of a direct link between luminosity and the way in which energy is

transported from the central region and converted to radio emission in the outer parts.

4.3.1 Fanaroff-Riley Class I (FR-I)

Sources in this class have their low brightness regions further from the central galaxy

or quasar than their high brightness regions (Figure 1.4 and Figure 4.1). The sources

become fainter as one approaches the outer extremities of the lobes and the spectra

here are the steepest, indicating that the radiating particles have aged the most. Jets

are detected in 80% of FR-I galaxies. A jet can begin as one-sided close to the core,

36


but beyond a few kiloparsec it becomes two-sided and continuous, with an opening

angle & 8 that varies along its length. Along the jet the component of the magnetic

field in the plane of the sky is at first parallel to the jet axis, but soon becomes aligned

predominantly perpendicular to the axis.

FR-I sources are associated with bright, large galaxies (D or cD) that have a flatter

light distribution than an average elliptical galaxy and are often located in rich clusters

with extreme X-ray emitting gas (Owen & Laing, 1989, Prestage & Peacock, 1988) As

the galaxy moves through the cluster the gas can sweep back and distort the radio

structure through ram pressure, which explains why narrow-angle-tail or

wide-angle-tail sources, say, appear to be derived from the FR-I class of objects.

A typical FR-I galaxy is shown in Figure 4.1. This is the radio source 3C 449, which is

optically identified with a galaxy of type cDE4 at a redshift of 0.0181, so that 1”

corresponds to 255 h−1100 pc. There are twin jets that are straight for ∼ 30” from the

core, after which they deviate towards the west and terminate into diffuse lobes. These

jets and outer lobes are mirror symmetric about an axis through the core. The jets are

generally smooth in appearance, but higher resolution observations show knots on a

smooth ridge of emission, the southern jet being more knotty than the northern one.

Within ∼ 10” of the nucleus, the surface brightness of the jets is much reduced. The

jets widen at a non-uniform rate close to the core, with the greatest expansion

occurring where the jets are faintest. Beyond ∼ 10” from the nucleus, the opening

angle is constant at 7. The emission from the jets is highly polarized, the average

polarization over the jets being ∼ 30%, and the projected magnetic field is

perpendicular to the jet axis.

4.3.2 Fanaroff-Riley Class II (FR-II)

This class comprises luminous radio sources with hotspots in their lobes at distances

from the center which are such that RFR > 0.5. These sources are called

edge-darkened, which was a terminology when the angular resolution and dynamic

range used in observing the classical sources was not always good enough to reveal the

hotspots as distinct structures. In keeping with the overall high luminosity of this type

of source, the cores and jets in them are also brighter than those in FR-I galaxies in

absolute terms; but relative to the lobes these features are much fainter in FR-II

37


Figure 4.1: VLA map of the FR-I galaxy 3C 449 at 1465 MHz, with angular resolution4.8 × 3.4 arcsec2. The peak flux is 22.2 mJy per beam, with contours drawn at 5%

intervals, beginning with the -5% contour.(Perley et al., 1979).

galaxies. Jets are detected in < 10% of luminous radio galaxies, but in nearly all

quasars. The jets have small opening angles (< 4) and are knotty; the jet magnetic

field is predominantly parallel to the jet axis except in the knots, where the

perpendicular component is dominant. Figure 4.2 shows an example of an FR-II

galaxy, which is a VLA map of the radio quasar 3C 47 made by Bridle et al. (1994).

The most striking feature of the jets in the FR-II class is that they are often one-sided,

as is clearly seen in Figure 4.2. Jet one-sidedness occurs at large (kpc) scales as well as

in the milli-arcsecond jets which are found in compact cores through VLBI

observations. The feature A in the jetted lobe is a hotspot, while feature H on the

38


unjetted side looks like one, but does not qualify for being a hotspot according to the

criteria of Bridle et al. (1994).

Figure 4.2: VLA map of the FR-II quasar 3C 47 made at 4.9 GHz with 1.45 ×1.13 arcsec2 resolution. G is the core, A the jetted hotspot. H does not meet the

hotspot criteria of Bridle et al.(Bridle et al., 1994).

FR-II sources are generally associated with galaxies that appear normal, except that

they have nuclear and extended emission line regions. The galaxies are giant ellipticals,

but not first-ranked cluster galaxies. The environment of FR-II sources does not show

enhanced galaxy clustering over the environment of randomly chosen elliptical galaxies

(Owen & Laing, 1989, Prestage & Peacock, 1988). Owing to the large differences in

the nature of the host galaxies and the environments of the FR-I and FR-II sources, it

is possible that they are intrinsically different types of source not related to each other

through an evolutionary sequence.

4.4 Radio Lobes and Jets

Radio-loud sources usually consist of a radio core, one or two detectable jets, and two

dominant radio lobes. The radio-quiet sources are less luminous at radio wavelengths

by a factor of 103 to 104, consisting of a weak radio core and perhaps a feeble jet. The

39


increased level of activity in radio-loud AGNs is not confined to radio wavelengths,

however; they also tend to be about three times brighter in X-rays than their

radio-quiet counterparts.

4.4.1 The Generation of Jets

The radio lobes are produced by jets if charged particles ejected from the central

nucleus of the AGN at relativistic speeds. These particles are accelerated away from

the nucleus in two opposite directions, powered by the energy of accretion and/or by

the extraction of rotational kinetic energy from the SMBH via the Blandford-Znajek

mechanism (Blandford & Znajek, 1977). The jet must be electrically neutral overall,

but it is not clear whether the ejected material consists of electrons and ions or an

electron-positron plasma. The latter, being less massive, would be more easily

accelerated. The disk’s magnetic field is coupled (“frozen in”) to this flow of charged

particles. The resulting magnetic torques may remove angular momentum from the

disk, which would allow the accreting material to move inward through the disk.

The incredible narrowness and straightness of some jets means that a collimating

process must be at work very near the central engine powering the jet. A thick, hot

accretion disk around the SMBH could provide natural collimation by funneling the

outflowing particles, as shown in Figure 4.3. Because the accreting material retains

some angular momentum as it spirals inward through the disk, it will tend to pile up

at the smallest orbit that is compatible with its angular momentum. Inside this

“centrifugal barrier”, there may be a relatively empty cavity that can act as a nozzle,

directing the accreting gases outward along the walls of the cavity. However, producing

highly relativistic jets, as frequently observed, appear to be difficult to accomplish with

this nozzle mechanism.

Alternatively, magnetohydrodynamic (MHD) effects could play an important role in

accelerating and collimating the relativistic flows.

4.4.2 The Formation of Radio Lobes

As a jet travels outward, its energy primarily resides in the kinetic energy of the

particles. However, the jet encounters resistance as it penetrates the interstellar

40


Figure 4.3: A Sketch of the electromagnetic outflows from the two sides of a rotatingmagnetized accretion disk owing to the unipolar dynamo action.

(CYGAM (CYlindrical GAmma-ray Monitor), Russian Space Research Institute).

medium within the host galaxy and the intergalactic medium beyond. As a result, the

material at the head of the jet is slowed, and a shock front forms there. The

accumulation and deceleration of particles at the shock front cause the directed energy

of the jet to become disordered as the particles “splash back” to form a large lobe in

which the energy may be shared equally by the kinetic and magnetic energy.

The motion of the charged particles and magnetic fields within the lobes of radio-loud

objects contain an enormous amount of energy. For Cygnus A Figure 4.4, the energy of

each lobe is estimated to be approximately 1053 to 1054 J, equivalent to energy

liberated by 107 supernovae!

41


Figure 4.4: Contour images of the Cygnus A radio jet on various scales.(Carilli et al., 1996).

4.4.3 Accelerating the Charged Particles in the Jets

The observations of jets are made possible by inefficiencies in the transport of particles

and energy out to the radio lobes. The spectra of the radio lobes and jets follow a

power law, with a typical spectral index of α ≃ 0.65. The presence of power-law

spectra and a high degree of linear polarization strongly suggest that the energy

emitted by the lobes and jets comes from synchrotron radiation.

The loss of energy by synchrotron radiation is unavoidable, and the relativistic

electrons in jets radiate away their energy after just 10,000 years or so. This implies

that there is not nearly enough time for particles to travel out to the larger radio lobes.

42


This long travel time implies that there must be some mechanism for accelerating

particles in the jets and radio lobes. As one possibility, shock waves may accelerate

charged particles by magnetically squeezing them, reflecting them back and forth

inside the shock. Radiation pressure may also play a role, but is alone not enough to

generate the necessary acceleration.

4.4.4 Superluminal Velocities

Although the standard model of jets and radio lobes requires a steady supply of

charged particles moving at relativistic speeds, evidence for such high velocities is

difficult to obtain. The absence of spectral lines in a power-law spectrum means that

the relativistic velocity of the jet material cannot be measured directly but must be

inferred from indirect evidence. The most compelling argument for relativistic speeds

involves radio observations of material ejected from the cores of several AGNs, with

so-called superluminal velocities. This effect is observed within about 100 pc of the

AGN’s center and probably continues farther out.

Figure 4.5: The apparent superluminal motion of the M87 Jet.

43

Chapter 5

The Optical-UV Regime

5.1 History

The first hint of the violent heritage of today’s galaxies was found by Edward A. Fath

(1908), who was observing the spectra of “spiral nebulae”. Although most showed an

absorption-line spectrum produced by the combined light of the galaxy’s stars, NGC

1068 displayed six bright emission lines. In 1926, Edwin Hubble recorded the emission

lines of this and two other galaxies. Seventeen years later, Carl K. Seyfert reported

that a small percentage of galaxies have very bright nuclei that are the source of broad

emission lines produced by atoms in a wide range of ionization states. These nuclei

were nearly stellar in appearance.

5.2 Spectrum

Figure 3.3 is a rough schematic of the continuum observed for many types of AGNs.

The most notable feature of this SED is its persistence over some 10 orders of

magnitude in frequency. The wide spectrum is markedly different from the thermal

(blackbody) spectrum of a star or the combined spectra of a galaxy of stars, and one

can see that there are many non-thermal radiation processes that are going on at

different stages in the AGNs.

When AGNs were first studied, it was thought that their spectra were quite flat.

Accordingly, a power law of the form : Fν ∝ ν−α was used to describe the

44

Chapter 5. The Optical-UV Regime

monochromatic energy flux, Fν . The spectral index, α, was believed to have a value of

α ≃ 1.

The power received within any frequency interval between ν1 and ν2 is

Linterval ∝

ν2∫

ν1

Fνdν =

ν2∫

ν1

νFνdνν = ln 10

ν2∫

ν1

νFνd log10 ν,

so that equal areas under the graph of νFν vs. log10 ν correspond to equal amounts of

energy. A value of α ≃ 1 reflects the horizontal trend seen at the IR bump of figure 3.3.

The continuous spectra of AGNs are now known to be more complicated, involving a

mix of thermal and non-thermal emission. However Fν ∝ ν−α is still used to

parameterize the continuum. α typically has a value between 0.5 and 2 that usually

increases with increasing frequency, so the curve of log10 νLν (or log10 νFν) vs. log10 ν

in Figure 3.3 is generally concave downward. In fact, the value of α is constant over

only a limited range of frequencies, such as in the IR and visible regions of the

spectrum. The shape and polarization of the visible-UV spectrum indicates that it can

sometimes be decomposed into contributions from thermal sources (blackbody

spectrum, low polarization) and non-thermal sources(power law spectrum, significant

polarization). The thermal component appears as the big blue bump in Figure 3.3,

which can contain an appreciable amount of bolometric luminosity of the source. It is

generally believed that the emission from the big blue bump is due to an optically

thick accretion disk, although some believe that free-free emission may be responsible.

5.2.1 The Optical-UV Continuum and the Accretion Disk

The best-understood disks are thin disks with or without X-ray emitting coronae. The

majority of intermediate- and high-luminosity AGNs found in large surveys are

thought to be powered by such disks. Thin-disk theory suggests that the SED of such

systems contains a broad wavelength band where the spectral slope α (Lν ∝ να) is in

the range 0-0.5. The “classical” thin-disk slope is α = 1/3. X-ray emission from the

hot corona and X-ray reflection from the surface of the disk are additional important

characteristics of such systems. Much observational effort has been devoted to the

measurement of α and the characterization of the part of the continuum showing this

45


Figure 5.1: A Composite Optical-UV Spectra of AGNs(Francis et al., 1991).

Figure 5.2: A General View of the Optical-UV SED of AGNs.

46


slope (the big blue bump). For MBH = 109M⊙ and L/LEdd = 0.1, the peak of this

emission is predicted to be around 1000 A. Even the most sophisticated calculated

spectra are rather limited, and several models show spectra that differ considerably

from the schematic ν1/3 dependence and also, the infrared (λ ≥ 1 µm) part of the SED

is dominated by non-disk emission, in particular, stellar emission by the host galaxy

and thermal emission by warm dust, presumably in the central torus. In radio-loud

AGNs, non-thermal emission can also contribute at this and even shorter wavelength

bands. This obscures the part of the spectrum where the standard thin-disk theory

predicts α = 1/3. One must also consider the (yet hypothetical) possibility that

additional processes, related perhaps to disk winds, are taking place and changing the

observed SED.

5.3 Observations in the Optical-UV Region

Optical images of luminous Type-I AGNs show clear signatures of point-like central

sources with excess emission over the surrounding stellar background of their host

galaxy. The non-stellar 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 luminosity of the nuclear, non-stellar source relative to the host galaxy

luminosity can vary by several orders of magnitude. In particular, many AGNs in the

local universe are much fainter than their hosts, and the stellar emission can dominate

their total light. For example, the V-band luminosity of a high-stellar-mass AGN host

can approach 1044 erg s−1, a luminosity that far exceeds the luminosity of many local

Type-I AGNs. This must be taken into account when evaluating AGN spectra obtained

with large-entrance-aperture instruments. The relative AGN luminosity increases with

decreasing wavelength, and contamination by stellar light is not a major problem at

UV wavelengths. The optical-UV spectra shown in Figure 5.4 and Figure 5.5 represent

typical spectra of high-ionization luminous 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 Type-I spectrum is a composite composed of several thousand

spectra of different redshift AGNs. This is done to illustrate the entire rest wavelength

range of 900-7000 A using only ground-based observations. The data used to obtain

this composite at λ > 5000 A are based on spectra of lower luminosity, low-redshift

47


Figure 5.3: Broadband spectral energy distributions (SEDs) for various types ofAGNs.

(Ho, 2008)

Figure 5.4: The average optical-UV SED of several thousand high-luminosity Type IAGNs

(Vanden Berk et al., 2001)

48


objects, and the SED at those wavelengths is affected by host galaxy contamination.

The Type-II spectrum covers a similar range, but this time, the spectrum is a

combination of a ground-based optical spectrum with a space-borne (HST) UV

spectrum. The striking differences between the high-ionization Type-I and Type-II

Figure 5.5: The spectrum of the low-luminosity, low-redshift type-II AGN NGC 5252.(Tsvetanov et al., 1996)

spectra, which were the reason for the early classification into Seyfert 1 and Seyfert 2

galaxies, are the shape and width of the strongest emission lines. Type-II AGNs show

only narrow emission lines with typical full-width-at-half-maximum (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 − 10000 km s−1 when

interpreted as owing to Doppler motion. The line ratios and line widths of the

49


Figure 5.6: Comparison of different broad-line profiles in a typical Type-I AGN.(Netzer, 2013)

forbidden lines in the spectra of Type-I sources are very similar to those observed in

Type-II spectra and indicate that the basic physics in the narrow line-emitting region

of both classes is the same. The broad emission lines can be used to map the gas

kinematics very close to the central BH and to measure the BH mass. Study of the

spectra of many thousand Type-I AGNs shows a considerable range in optical-UV

continuum slope but little if any correlation between the slope and Lbol. Some of the

observed differences are attributed to a small amount of reddening in the host galaxy

of the AGN or other sources of foreground dust. Although the spectra shown here

clearly illustrate the large differences in emission-line widths between Type-I and

Type-II sources, observational limitations can make it difficult to detect weak broad

emission lines. Slightly obscured or low-luminosity Type-I AGNs are occasionally

classified as Type-II based on their stellar-like continua and narrow emission lines.

This can be the result of reddening of the broad wings of the Hβ line or a relatively

strong stellar continuum, especially in large-aperture, low-spatial-resolution

observations. Higher signal-to-noise (S/N), better-spatial-resolution observations of the

same sources reveal, in some cases, very broad wings in one or more Balmer lines. The

term broad emission lines, which is used to describe the permitted and semi-forbidden

lines in Type-I AGNs, does not necessarily mean similar widths for all lines in all

50


objects. The various broad emission lines show typically different widths, and in

general, the width reflects the level of ionization of the gas, the source luminosity, and

the mass of the central SMBH. Historically, it was found that broad emission lines in

some Type-I AGNs are narrower than narrow emission lines in Type-II sources. A

well-known example is the subgroup of narrow-line Seyfert 1 galaxies (NLS1s). This

class of objects was historically defined as those Seyfert 1 galaxies with

FWHM(Hβ) < 2000 km s−1. In many such sources, FWHM(Hβ) < 1000 km s−1,

similar to the width of Hβ in many Type-II AGNs. Evidently, the distinction between

Type-I and Type-II sources requires other criteria, such as the presence of a non-stellar

continuum; strengths of emission lines typical of Type-I sources such as FeII emission

lines; or the presence of a strong, unobscured X-ray continuum. There are also

differences in the shape and even the velocity of the same line in different objects.

Some examples are shown in Figure 5.6. A similar remark should be made about the

width of the narrow emission lines. For example, the width of the strong [O III] λ 5007

line can depend on the mass of the central SMBH (or, more accurately, the mass of the

bulge). Thus [O III] λ 5007 lines with FWHM≥ 1000 km s−1 are commonly observed

in high-redshift, large-MBH , large-Lbol AGNs.

5.4 Discovery by Optical-UV Properties

As said, typical AGN SEDs are different in several ways from stellar SEDs, in which

they cover a broader energy range and do not resemble a single-temperature blackbody.

This difference provides a simple and efficient way of discovering AGNs using

broadband multicolor photometry. Several color combinations, based on three-band

and five-band photometry, are useful in separating AGNs from stars by their color.

Earlier methods were based on a UVB photometry in large areas of the sky. This

three-band system is useful for discovering low-redshift sources but fails to detect many

high-redshift objects because the spectrum gets effectively red and resembles the colors

of nearby stars. In addition, even the low-redshift AGNs can be confused with the

local population of hot white dwarfs. Some AGNs are intrinsically red, or reddened by

dust, which results in colors that are not very different from those of stars. Moreover,

intrinsically blue, high-redshift AGNs are “effectively red” due to the absorption of

their short-wavelength radiation by intergalactic gas. Figure 5.7 illustrates this and

51


Figure 5.7: The u-g color of a large number of SDSS AGNs with various redshifts.(Abazajian et al., 2009)

shows the u-g colors of a large number of quasars at higher redshifts observed in the

Sloan Digital Sky Survey. The blue AGNs (small values of u-g) are seen at all z < 2.5,

but the color is much redder at higher redshifts. More color combinations and other

techniques are needed to find the high-redshift sources. The more sophisticated

five-band systems overcome most of these difficulties. They use a combination of

several colors and are very efficient in detecting AGNs up to z ≃ 6. The SDSS system,

which produced the largest number of AGN candidates, is a color-color system based

on five photometric bands: u (0.35 µm), g (0.48 µm), r (0.62 µm), i (0.76 µm), and z

(0.91 µm). The system is very efficient for low-redshift AGNs because of the blue color

of such sources. The additional bands help to separate AGNs from white dwarfs. The

five-band system, with its many color combinations, is also very efficient in discovering

high-redshift AGNs. An illustration of the method as adopted by the SDSS survey is

shown in Figure 5.8. Such methods have been shown to produce flux-limited AGN

samples that are complete to a level of 90% and even higher. The total number of

Type-I AGNs discovered in this way, as of 2011, is more than 100,000. Type-I AGNs

can be directly discovered by their spectrum, because of the large contrast between the

strong broad emission and absorption lines and the underlying continuum. This

52


Figure 5.8: Discovering AGNs by their broadband colours. The plots show the AGNlocation in various colour-colour diagrams using the SDSS bands u, g, r, i, z. Blackpoints and contours are stars of different types. Colours specify the redshift of the AGN

(Richards et al., 2004)

method was used in the 1970s and resulted in several large high-redshift samples.

This spectroscopic method is based on objective prism surveys that produce a small,

low dispersion spectrum, instead of a single point, for every object in the field. This

was eventually supplemented by high-resolution follow-up spectroscopy. The method is

most useful in detecting broad-absorption-line (BAL) AGNs. Several such surveys gave

the erroneous impression that BAL AGNs are more common than they really are

because they are difficult to miss in objective prism surveys. The method is very

inefficient in discovering Type-II AGNs, with their relatively weak emission lines and

strong stellar continuum.

53

Chapter 6

The X-Ray Regime

6.1 History

The first X-ray measurements of AGN were made with detectors on-board an Aerobee

rocket in April 1965, which provided evidence for high-energy emission from Cygnus A

and M87 (Byram et al., 1966). A flight in 1969 then provided the first detection of the

radio galaxy Cen A and of the quasar 3C 273, at a significance level of 3.0σ and 3.9σ,

respectively. Rocket flights however did not provide sufficient exposure times to further

advance the field once the brightest sources had been detected. Uhuru, the first X-ray

telescope on an orbiting satellite was then launched in December 1970. The mission

was equipped with two large area proportional counter detector systems, with 840 cm2

effective area each. It performed the first sky survey leading to a catalog of 339 sources

in the 2-6 keV range. The predominant fraction of these sources were found to be

compact, mass-exchanging binaries, commonly known today as X-ray binaries (XRBs).

Examples included Cyg X-3, Her X-1, and Vela X-1 named after their host

constellations and the order of their discovery. These objects provided a general

confirmation of theoretical models that suggested accretion onto a compact object that

can lead to X-ray radiation.

Uhuru also provided the first detection of the Seyfert galaxies NGC 4151 and of NGC

1275, and confirmed the earlier detection of Cen A, Cygnus A, M87, and of 3C 273. In

total, 15 Seyfert galaxies were detected, all of them being Seyfert I or Seyfert 1.5. The

54

Chapter 6. The X-Ray Regime

largest class of extragalactic objects though were the galaxy clusters, of which Uhuru

detected 45.

6.2 Probing the Innermost Regions

During the same decade of the 1960s when the basic AGN paradigm was being

developed, the first cosmic X-ray source (Scorpius X-1), was discovered using a

rocket-borne detector (Giacconi et al., 1967). It quickly became evident that X-ray

emission was characteristic of compact, accretion-powered sources associated with

galactic binaries. With the realization that the deep gravitational wells of massive

black holes were likely the source for the extreme energetics exhibited by quasars, the

generation of X-rays seemed natural. However, the rocket-borne experiments had

limited capabilities and the first significant breakthrough came with the launch of the

Uhuru satellite (also known as SAS-A) in 1970. A catalog of sources detected with

Uhuru ultimately included nearly 340 objects, mostly galactic binaries, but also

including about a dozen AGN (Forman et al., 1978).

The field progressed rapidly during the 1970s, with detectors having larger collecting

area, increased spectral coverage and improved spectral resolution, for example the

Ariel-5 satellite launched in 1975 and the OSO-7 (1974), OSO-8 (1975) and HEAO-1

(1977) (Tucker & Giacconi, 1985). This led to the characterization of AGN as a class

of X-ray sources and to the first detection of the iron Kα line emission from an

extragalactic source. The biggest breakthrough however came later in that decade with

the launch of the Einstein Observatory (formerly, HEAO-2). This was the first true

orbiting X-ray telescope, in that it utilized a concentric array of grazing incidence

mirrors to focus ∼ keV photons onto its focal plane detectors. The resulting images

provided, in addition to vastly improved spatial localization of sources on the sky, a

large leap in sensitivity since the source and celestial background could be effectively

separated.

It had become clear that X-ray emission was a common property of different subclasses

of active galaxies with the X-ray flux comprising a significant fraction (about 5-40%) of

the bolometric emission from such objects (Ward et al., 1988). Rapid variability was

also found to prominent feature of the X-ray emission with kilosecond timescale X-ray

55


flux variations seen in local Seyfert galaxies. This imposed new and increasingly

stringent constraints on the size of the X-ray emission region and strongly supported

the idea that it occurs very close to the active nucleus (Pounds & Turner, 1988). The

origin of X-rays from close to the central black hole means that X-ray data offer a

chance to study the immediate environs of SMBHs and the accretion process that fuels

them. Although the angular scale of the X-ray emission region is too small to image

with current instrumentation, timing analysis and spectroscopy offered methods to

probe these regions indirectly.

Specific spectral signatures were attributed to the characteristics of the gas inflow and

outflow near the central most regions in AGN. The X-ray observations also provided

signatures of reprocessing of radiation in material withing approximative distance of

hundreds of gravitational radii. Features such as the weak, broad emission lines (BELs)

due to low-ionization states of iron as well as other structured deviations from simple

power laws had been identified in the spectra of AGN. George & Fabian (1991) offered

an interpretation of these features in terms of X-irradiation of relatively cold, dense gas

in the vicinity of the central black hole. The emergent spectrum then consists of direct

radiation from the central source plus a scattered or “reflected” spectrum that includes

imprinted photoabsorption, fluorescent emission and Compton scattering from matter

within the surrounding accretion flow. This basic idea has withstood the scrutiny of

improved observational data and has become a tenant of the AGN paradigm.

X-ray observations of AGN are also being applied to address issues of fundamental

black hole physics. The shapes of line profiles have also been applied to models which

in principle allow one to infer an intrinsic property of the central black hole, namely its

intrinsic angular momentum or spin (Brenneman & Reynolds, 2009). The basic idea is

that the asymmetry of a line profile produced in the inner AGN accretion disk depends

in a predictable manner on the black hole’s spin.

6.3 The X-Ray Spectrum of AGNs

In X-rays, the accretion disk surrounding a black hole is believed to produce a thermal

spectrum. The low-energy photons produced from this spectrum are scattered to

higher energies by relativistic electrons, residing for example in the corona above the

56


accretion disk, through inverse Compton process (Section 2.4.3). As the temperature

of the disk and the relativistic electrons energy distribution are limited, the resulting

inverse Compton spectrum has a high-energy cutoff. The spectrum has an

approximate power law shape with a photon index of Γ ≃ 2 extending upto a few

hundred keV. The soft photons involved in the inverse Compton scattering are believed

to originate in the cool thick accretion disk disk with kT < 50 eV, while the relativistic

electron gas has a temperature of about kT ∼ 1000 keV. The photons will be

up-scattered from their initial energy Ei to the energy

Ef = eyEi,

where y is the Compton parameter (Section 2.4.2). In the non-relativistic case with

τ > 0.01, this results in y ≪ 10 and a power law which extends upto the thermal cutoff

in the range Ecut ≃ kT to Ecut ≃ 3kT , determined by the cutoff in the thermal

distribution of the electrons. This spectrum is undergoing reprocessing through

Figure 6.1: Composite AGN spectrum in extreme UV based on FUSE satellite datawith continuum fit (dashed line) and fit to continuum and emission lines (solid line)

(Scott et al., 2004)

57


absorption. In most cases, broad emission lines due to low-ionization stages of iron are

visible in the spectra. Thus, in these objects, the absorber is commonly assumed to

consist of cold (T < 106 K) optically thick material (George & Fabian, 1991). This

reprocessing leads to yet another “bump” in the hard X-ray spectrum (like in the IR),

first observed by the Ginga satellite in the Seyfert I galaxies NGC 7469 and IC 4329A.

The shape and strength of the bump depends on the geometry, chemical composition

and orientation with respect to the line of sight, but has its maximum around 20-30

keV, where the reflection efficiency reaches its maximum. Its measurement is difficult

as the modeled strength of this component is closely linked to the intrinsic absorption

(measured at softer X-rays), spectral slope of the underlying component, and the

high-energy cutoff, which is not well constrained in many cases. The reflecting material

could reside for example in the outer accretion disk, or at the inner edge of the

absorbing gas, or could be located in an outflowing wind. The reflection strength R is

normalized so that R = 1 represents the case of an isotropic source above an infinite

reflecting plane. R is often considered as an estimate of Ω/2π, where Ω is the solid

angle subtended by the reflector as seen from the isotropic X-ray source. When

Figure 6.2: Soft X-ray spectrum of the narrow-line Seyfert I Arakelian 564.(Smith et al., 2008)

studying X-ray spectra in the 2-20 keV band taken by the Ginga satellite, Zdziarski

et al. (1999) observed that the reflection strength R is strongly correlated with the

intrinsic spectral slope as measured in the photon index Γ, as in Figure 6.1. The

correlation is seen not only in AGN, but also in X-ray binaries in their hard state.

This means that sources with steeper X-ray spectra show a larger reflection than those

58


with a flat spectrum. An explanation might be that the photons which, together with

the relativistic electrons in the corona, are the sources for the inverse Compton

component observed in the X-rays, are coming from the cool material which is also

responsible for the reflection component.

Another component observed in AGNs is the soft X-ray excess. A soft (E . 2 keV)

excess over the power law component dominant at higher energies has been found in

the X-ray spectra of many Seyfert galaxies (Saxton et al., 1993). The origin of the soft

excess is still an open issue. In the past, the soft excess had often been associated with

the high-energy tail of the thermal emission of the disk, but it was recently argued that

the temperature of the disk should be nearly constant (kT ≃ 0.1 − 0.2 keV), regardless

of the mass and luminosity of the AGN (Done & Gierlinski, 2004). This result implies

that some other mechanism is at work, as the temperature of the disk should depend

on both the mass of the black hole and the accretion rate.

For bright AGNs, grating spectroscopy can be performed providing substantially

higher spectral resolution that can be obtained with CCDs or proportional counters,

for example R ∼ 103. Figure 6.2 shows an example of a high-resolution soft X-ray

spectrum taken by the RGS instrument on-board ESA’s XMM-Newton satellite.

Absorption and emission lines can be clearly identified, allowing to estimate the

temperature of the absorbing material by identification of the ionized lines. In

addition, the width of the lines give information about the velocity of the particles in

the absorbing clouds, and their displacement with respect to the laboratory wavelength

gives the outflowing (or inflowing) speed.

6.4 Lineless AGNs

Systematic studies of large AGN samples result in the discovery of a subpopulation of

AGNs with extremely weak, sometimes completely undetected emission lines. A

typical upper limit on the EW of the emission lines in such sources is 1 A. The objects

show at least one of the four AGN indicators, usually a non-stellar 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 luminosity,

from very faint objects in the local universe to very luminous AGNs at high redshift.

59


They are referred to in the literature as lineless AGNs, anemic AGNs, dull AGNs, and

other equally original names. Lineless AGNs differ in their optical continuum

Figure 6.3: The composite spectrum of 15 lineless AGNs with large X-ray-to-opticalluminosity. The top panel shows the composite of the 15 sources (top curve) andcompares it with the spectrum of a red galaxy. The bottom panel shows the stellar

subtracted continuum alongside a composite Type-I spectrum.(Trump et al., 2009)

properties from blazars. They do not show a power law continuum; they are mostly

radio quiet; their variability, if any, is of very small amplitude; and the typical

double-peak SED of blazars is not observed. Figure 6.3 shows a composite spectrum of

15 such sources from the COSMOS survey.

The very luminous lineless AGNs are of special interest and may 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 (Equivalent

Width)) compared to other Type-I sources. Broad emission lines in AGNs are known

to show a decrease of line EW with continuum luminosity and/or L/LEdd. The EWs of

the very luminous lineless AGNs (in most cases, only upper limits on the EW) are at

the very end of these distributions. Nevertheless, extremely large L/LEdd is one

possible explanation for the weak broad emission line.

60


A different type of 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 luminosity sources, but for very 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 (RIAF) with inefficient

conversion of gravitational energy to electromagnetic radiation. Such systems can lack

much or all of the (otherwise strong) UV ionizing radiation. This has been proposed as

a possible explanation for the very low luminosity of lineless AGNs such as the ones

shown in Figure 6.3. Regarding the high-luminosity sources, here the Lyman

continuum radiation by the disk depends on the BH mass and accretion rate and can

be extremely weak in disks around very massive BHs. Such systems are likely to show

very luminous continua but no line emission.

6.5 The Central Obscuration

The X-ray opacity of atomic gas is strongly wavelength dependent. For neutral gas

with solar composition, a unit optical depth at 0.3 keV is achieved for hydrogen

column density of about 4.5 × 1020 cm−2. The corresponding column at 5 keV is about

4.5 × 1023 cm−2. At around a column of 1.5 × 1024 cm−2, the gas is Compton thick,

which prevents the transmission of almost all the X-ray radiation above about 8 keV.

For larger columns, all the X-ray radiation is absorbed. X-ray observations provide the

most efficient way for detecting and measuring the line-of-sight-obscuring column in

AGNs. Numerous observations of Type-II sources show a wide column density

distribution with a peak at around 1023 cm−2 and a long tail toward very large

columns. While X-ray absorption does not depend on the dust content of the gas,

much of the obscuring material must be dusty to explain the large opacity at long

wavelengths up to 1 µm and perhaps even more (for solar composition dusty gas with

galactic gas-to-dust ratio τ(5500)A≃ NH/1.5 × 1021 cm−2). Figure 6.3 shows absorbed

X-ray spectra of several Type-II AGNs.

X-ray obscuration is not restricted to Type-II AGNs. In fact, most low-luminosity

Type-I sources show some X-ray absorption along the line of sight with column

densities that range between 1021 and few ×1023 cm−2. Unlike the neutral obscurer in

Type-II sources, in these cases, the gas is highly ionized and, most probably, contains

61


no dust grains. This gas is thought to flow out of the source inside or just outside the

central opening of the torus. This gas is the highly ionized gas (HIG) or the warm

absorber.

6.6 Detection and Observations of AGN in X-Rays

6.6.1 X-Ray Observations of AGNs

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 to 10 keV with a spectral resolution typical of

proportional counters and CCD detectors (∆E ∼ 100 eV). 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 missions improved this situation dramatically by

providing grating spectroscopy of nearby AGNs. The resolution of these instruments,

below about 1 keV, is of order E/∆E ≃ 1000. This has revolutionized X-ray studies of

AGNs and resulted in the identification of hundreds of previously unobserved emission

and absorption lines. Present-day X-ray instruments like Suzaku and Swift/BAT

extend the low-resolution observations to 100 keV and even beyond.

6.6.2 Discovery by X-Ray Properties

Almost all AGNs are strong X-ray emitters. This property can be used to discover

AGNs by conducting deep X-ray surveys. An example of a sample that resulted in the

detection of numerous new AGNs is the ROSAT all-sky survey. 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 1022 cm−2 or larger are

more difficult to detect. Recent (since 1999) deep surveys are those conducted with

Chandra and XMM-Newton. The Chandra surveys are extremely deep because of the

superb resolution of this instrument (about 1”). Unlike the low-energy ROSAT survey,

62


both Chandra and XMM-Newton can observe at higher energies, up to about 10 keV.

However, these missions have only covered a small fraction of the sky. Recent hard

X-ray all-sky surveys include those from Swift Burst Alert Telescope (BAT) and

INTEGRAL missions at energies from 15 to 150 keV. The amount of obscuration at

high energies is much smaller, which results in X-ray discovery of many type-II AGNs.

Needless to say, optical follow-up spectroscopy is needed to confirm those detections.

X-ray observations are not very efficient in discovering very high redshift AGN because

of the limited sensitivity of the X-ray instruments and the sharp drop of X-ray

luminosity of such sources.

63

Chapter 7

The γ-Ray Regime

7.1 History

The impact of gamma-ray astronomy on AGN research did not emerge as rapidly as

did X-ray astronomy, although the fields were initiated more or less concurrently with

1960s rocket flights followed by satellite-borne experiments in the 1970s. The reasons

for this are several-fold. There are fewer gamma-ray photons than lower energy

photons emitted even though the overall energy budget for some AGN may be

dominated by the gamma rays. There are substantial instrumental and celestial

backgrounds at gamma-ray energies that need to be understood and modeled or

subtracted. Gamma-ray detectors tend to be more massive for a given effective

collection area than X-ray detectors and gamma rays cannot be focused. Additionally,

it became apparent that only the radio-loud AGN, which ∼ 5% of the overall

population are prolific emitters of gamma radiation.

In the 1970s, the ESA mission COS-B, along with NASA’s SAS-2, provided the first

detailed views of the Universe in gamma-rays. COS-B, launched in August 1975, was

originally projected to last two years, but it operated sucessfully for nearly seven. It

made the first gamma-ray measurement of an AGN, that being 3C 273 (Swanenburg

et al., 1978). However, it was not ready until 20 years later with the launch of the

Compton Gamma-Ray Observatory (CGRO) that additional gamma-ray detections

were made, starting with the discovery in 1991 of bright gamma-ray emission from 3C

279 (Hartman et al., 1992). New results came quickly after that leading ultimately to

64

Chapter 7. The γ-Ray Regime

the identification of some 70 high-latitude CGRO gamma-ray sources with radio-loud

AGNs. Specifically, BL Lac objects and flat-spectrum radio quasars (FSRQs),

collectively of the blazars subgroup (Section 1.3.4), comprised the entire gamma-ray

sample. It was also clear that the radiative output of the blazars was typically

dominated by gamma-rays. The gamma-ray emission was also found to be variable on

time scales less than a day.

These observations had several immediate implications for physical models. The

emission had to emanate from a compact region. For example, a factor of 2 flux

variation limits, approximately, the size of r of a stationary, isotropic emitter to

r . cδtvar/(1 + z), where δtvar is the variation time scale. The implications from the

early CGRO results, which by this line of reasoning necessitated a very compact

emission region, were problematic in any scenario in which the gamma-ray production

involves such a stationary isotropic source. The problem involved the transparency of a

compact region such as inferred here. If X-rays are produced co-spatially with the

gamma rays, attenuation of the gamma rays due to the process γγ → e+e− for which

the cross-section for attenuation of ∼ 100 MeV gamma rays is in the X-ray range

∼ keV X-ray range. The inferred gamma-ray opacity from the CGRO observations

would exceed unity in many instances. Either the radiating particles were strongly

beamed or the emitting plasma was undergoing bulk relativistic motion. Thus,

beaming was very strongly implied.

Models that had been previously favored to explain the radio-to-optical continua in

these objects, for example Blandford & Konigl (1979), implied that we are viewing

nearly along an axis of a relativistic plasma jet ejected from near the central black

hole, involving non-thermal synchrotron emission. An extension of this scenario

invoking a distinct second spectral component was now clearly required. The basic

idea was that gamma rays emitted by blazars are produced by the same population of

electrons that produced the synchrotron emission via Compton scattering of ambient

low-energy photons. The ambient photon field could be the synchrotron photons

themselves (eg., Maraschi et al. (1992)) or form an external source such as the

accretion disk or broad-line clouds (eg., Dermer et al. (1992)).

Shortly after the CGRO results began to emerge, another major discovery followed

from ground-based Cerenkov gamma-ray telescopes, which measured gamma rays in

65


the ∼ TeV range. Blazars such as Markarian 421 (Punch et al., 1992) and Markarian

501 (Quinn et al., 1996) were detected during a high-amplitude variability episode.

These discoveries established this subclass of AGN as emitters over ∼ 20 decades of

the electromagnetic spectrum. As such, they were a striking sample of the value,

indeed the necessity, inherent in the multiwavelength approach to studying AGN. The

high-energy gamma-ray observations also fit in naturally with the synchrotron plus

Comptonization model scenarios. They also had other potentially significant

implications, not only on the blazar AGN themselves, but on the gamma-ray

transparency of the universe and thus in turn the background radiation fields to the

cosmic star-formation history. In the two decades since these discoveries, gamma-ray

Figure 7.1: The multiepoch, multiwavelength spectrum of the blazar 3C 279,showingthe two characteristic peaks at low and high energies and the long-term variations of

the source.(Bottcher et al., 2007)

studies of AGN have expanded enormously. The Fermi Gamma-Ray Telescope,

launched in 2008, has cataloged approximately 900 gamma-ray AGNs. Advances in

ground-based Cerenkov telescope facilities, as well as in detection and analysis

methodologies, has produced a similar order-of-magnitude increase in the TeV

gamma-ray sample. Multiwavelength campaigns have begun to reveal how jet

formation and propagation may be correlated with the gamma-ray flux variations.

66


7.2 Gamma-Ray Loud AGNs

The group of blazars includes highly variable core-dominated radio-loud sources

showing polarization at radio and optical wavelengths. Many blazars are also powerful

γ-ray emitters, and some of them show one or more of the following properties:-

• Intense, highly variable high-energy emission in the γ-ray part of the spectrum.

• Intense, highly variable radio emission associated with a flat radio spectrum and

occasionally, superluminal motion (Section 4.4)

• Radio, X-ray, and/or γ-ray jet with clear indications for relativistic motion.

• A double-peak SED with a lower-frequency peak at radio-to-X-ray energies and a

high-frequency peak at X-ray-to-γ-ray energies (Figure 7.1).

• Very weak (small EW) broad and/or narrow emission lines indicative of

photoionization by a non-stellar source of radiation on top of a highly variable

continuum.

Blazars can be divided into BL Lac objects (section 1.3.4.1) and fla-spectrum

radio-loud AGNs. The flat radio spectrum blazars are occasionally called flat-spectrum

radio quasars (FSRQs) or OVVs (Section 1.3.4.2). BL Lac objects are often

sub-classified into low-energy-peaked BL Lac objects and high-energy-peaked BL Lac

objects.

7.3 γ-Ray Properties of Blazars

The understanding of the physical mechanism that drives the blazar phenomenon is

strongly coupled to the launch of various advanced X-ray and γ-ray instruments. The

launch of the Fermi Gamma-Ray Telescope in 2008 revolutionized this field by

confirming earlier suggestions that most of the energy in these sources is produced by

relativistic jets and by detecting many more blazars. Most of these discoveries are due

to observations by the Large Area Telescope (LAT), a wide-field-of-view imaging

telescope covering the energy range of ∼ 20 MeV to 300 GeV.

67


The LAT allows blazar variability to be monitored over a wide range of time scales. It

shows that large amplitude variations are very common in most blazars. The

suggestion is that in these sources, most of the non-thermal γ-ray emission arises from

relativistic jets that are narrowly beamed and boosted in a specific direction. The jet

is launched in the vicinity of the central active SMBH, and the angle between the line

of sight and the axis of the jet is typically a few degrees or less. This explains the

superluminal motion often observed in VLBI observations of blazars. There is good

reason to believe that FSRQ blazars are associated with pole-on FRII radio sources

and BL-Lac objects with pole-on FRI sources. Because of this, FRI and FRII sources

are occasionally referred to as the parent population of blazars.

The jet model raises more questions than answers: how is the jet collimated and

confined? What is the composition of the jet close to the launch points and much

further out? What are the details of the conversion between the jets kinetic power and

electromagnetic radiation? Simultaneous multiwavelength observations of blazars are

perhaps the most important tools for answering these and other questions. Today, such

observations can cover a huge energy band, from several centimeters in the radio

through MIR, NIR, optical, UV, X-ray, and all the way up to above 100 GeV.

Ground-based observatories, like HESS, can extend them to even beyond the LAT

energy range.

The fact that most blazars show an SED with two broad peaks (Figure 7.1) is

consistent with the jet model. At lower frequencies, from radio to UV and sometimes

X-rays, the emission is dominated by synchrotron radiation of highenergy electrons in

the jet. The higher-energy part of the SED, from X-rays to γ-rays, is thought to result

from inverse Compton emission (Section 2.4.3).

Detailed studies of blazars by LAT reveal real physical differences inside this

inhomogeneous group of sources. The lower γ-ray luminosity blazars, classified as

BL-Lac objects, have harder γ-ray slope (Γ = 2) compared with flat-spectrum radio

AGNs (Γ = 2.5). A simple power law is not always a good description of the γ-ray

continuum, and in several well-studied cases, the spectrum is better fitted by a broken

power law with a steeper, higher-energy part. There are other differences that relate to

the galaxy type and morphology. Blazars with strong relativistic jets are usually

hosted in elliptical galaxies. However, Fermi found several NLS1s that are also strong

68


γ-ray emitters. These sources are thought to have very large L/LEdd and are hosted in

spiral galaxies with high star formation rates. All this shows that the subclass of

blazars includes objects with very different physical properties that depend on the

central energy source, the BH mass and spin, the exact geometry and inclination, and,

perhaps, the evolutionary phase of the sources.

69

Chapter 8

The Unified Model of AGNs

8.1 The Unification

A fundamental question in AGN research is, whether all these distinct appearances of

the AGN phenomenon can be explained by a common underlying model, or whether

the different classes are intrinsically distinct. It was pointed out earlier on that a

Seyfert galaxy is in fact in most cases, a spiral galaxy to which a faint quasar is added

in the center (Boksenberg et al., 1975, Weedman, 1973). In addition, Kristian (1973)

showed that the fainter quasars indeed appear to have an extended form rather than

being point-like, indicating that they reside in galaxies. Rowan-Robinson (1977) made

an attempt to unify Seyfert galaxies and radio sources. While he correctly assumed

that absorption by dust is important in order to explain the differences in infrared

emission, he did not take into account beaming effects which are an important

ingredient when trying to understand radio-bright AGN. At a 1978 BL Lac conference

in Pittsburgh, the foundations for the beaming unification were outlined (Blandford &

Rees, 1978), a concept which is still believed to be true. In this picture, if an AGN

appears to be a blazar, it is because the emission is beamed along the symmetry axis

of the AGN towards the observer (Figure 8.1). In a next step, Scheuer & Readhead

(1979) proposed that the radio-core dominated quasars could be unified with the

radio-quiet quasars by assuming the former ones are beamed towards the observer,

similar to the case of blazars. This implies that all radio-quiet quasars also host a

relativistic radio jet, but they are only FSRQ when the jet is along the line of sight.

This concept turned out to have a problem though. As Orr & Browne (1982) pointed

70

Chapter 8. The Unified Model of AGNs

Figure 8.1: Schematic representation of a geometrical interpretation of the BL Lacphenomenon.

(Blandford & Rees, 1978)

out, the core-dominated and radio-loud quasars indeed showed extended radio emission

in MERLIN and VLA observations. Therefore, radio-quiet AGN could not simply be

misaligned radio quasars. Later studies explained the differences by two effects:

difference in orientation, and difference in obscuration (Barthel, 1989). A still valid

and rather complete overview of the problem of AGN unification was given by

Antonucci (1993). In the most simplified picture, there are basically two types of

AGN: radio-quiet and radio-loud. For each type, a range of luminosities is observed,

leading for example to the Fanaroff-Riley classes (Section 4.3) as well as to the

distinction between a Seyfert and a quasar. All other observed differences would be

explained by orientation effects. In this scenario all objects which show a quasi-stellar

radio core and blazars would emit beamed radiation towards the observer, with a

closer alignment to the line of sight in the case of the blazars. Radio galaxies, in this

picture, emit their jet at large off-axis angles with respect to the line of sight.

Antonucci (1993) pointed out that the existence of an optically thick torus surrounding

the central regions of an AGN on scales of 1-100 pc would lead to the absence of broad

emission lines in the case of Seyfert II if they were observed edge-on, as their broad-line

region would be hidden, compared to Seyfert I objects, which are mostly observed

face-on. As a narrow-line region lies further away from the central black hole than the

broad-line region, the NLR would still be observable when the BLR is obscured by the

torus. In addition, Antonucci (1993) was also aware of the shortcomings of this simple

model and that it left open the question, what is the intrinsic difference between

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radio-loud and radio-quiet AGN, and why do radio-loud AGN mainly reside in

elliptical galaxies, while radio-quiet AGN are hosted by spiral galaxies.

A subsequent view by Urry & Padovani (1995) explained the unification of the

subgroup of radio-loud AGN. The aim was to study whether the low-luminous FR I

can be the parent population of the BL Lac objects, while FSRQ would be a subset of

the FR II galaxies. Urry & Padovani (1995) pointed out the difference in evolutionary

behaviour between BL Lacs and FSRQ, and considered the suggestion that FSRQ

evolve into BL Lac objects, becoming weak-lined objects by virtue of increased

beaming of the continuum, that is, of a Lorentz factor increasing with cosmic time

(decreasing with redshift; Vagnetti et al. (1991)). A problem became evident though,

as the luminosity functions of the two object groups could not be connected smoothly,

for example, because of very different radio power and line luminosity at comparable

redshifts.

8.2 Absorbed Versus Unabsorbed AGN

If we put aside for the time being the difference between radio-loud versus radio-quiet

AGN, the unified model predicts a distinction between the various types of sources

based solely on the viewing angle. The anisotropy of the AGN population is then

assumed to be caused by the different level of absorption in the line of sight. This

concept leads to certain predictions which can be tested through observations. All the

intrinsic properties, that is, the appearance of the AGN when absorption effects are

not relevant or they have been modeled out, should be similar for all Seyfert types on

one side and all radio emitting sources on the other. On the other hand, one should be

able to observe a consistent set of differences which can be explained by the

absorption, and which should correlate with the optical depth of the absorber.

A stronger test for the unification model is the unification model is the prediction that

the broad-line region lies at smaller radii than the absorbing material, whereas the

narrow-line region should be visible in all Seyfert types as the emitting material resides

further out. This part of the model was based on the observational fact that Seyfert I

galaxies have broad and narrow emission lines whereas Seyfert II have only narrow

lines. The distinction between Type I and Type II should disappear if one finds a way

72


to exclude the influence of absorption in the observation of AGN spectra. One method

to investigate the intrinsic line width in optical spectra is to study the polarized light.

Although the broad-line region (BLR) is hidden by obscuring material in the case of

Type II AGN, the light of the BLR can escape in directions where no material hinders

the view to the central engine and its surroundings. If the BLR’s light then hits, for

example, electrons, it can be scattered into the line of sight of the observer and thus

still reach us. This scattering of photons follows the rules of Thomson scattering

(Section 2.3), and the scattered light can be linearly polarized. Thus, by looking at the

polarized emission only, the broad component in Seyfert II can still be visible, and the

scattering material acts like a mirror which enables us to look behind the absorbing

matter.

A first proof of this concept was given by Antonucci & Miller (1985), who showed that

the Balmer lines in the Seyfert II NGC 1068 are broad when the AGN is observed in

polarized light. They also showed that the non-thermal continuum emission of the

central engine has the same level of polarization as the Balmer and FE II emission

lines, in the case of NGC 1068 about Π ≃ 15%. Subsequent spectropolarimetric

observations of highly polarized Seyfert II galaxies discovered further hidden broad-line

regions (HBLR), for example in Mrk 3, Mrk 348, Mrk 463 E, Mrk 477, Mrk 1210, NGC

7212, NGC 7674, and Was 49b (Miller & Goodrich, 1990, Tran et al., 1992). The

observations confirmed that continuum and broad Balmer lines show the same degree

of polarization, which however can differ substantially from object to object. At the

same time, the narrow forbidden lines do show little or no polarization at all,

confirming that the narrow-line region is observed directly. This discovery was

certainly a strong argument in favour of the unified model. On the other hand, there

are numerous Type II AGN which do not show any broad-line component in polarized

light. In total, only 40% of the Seyfert II galaxies are seen to have an HBLR (Wu

et al., 2011). Jiehao Huang (2002) found that those which have polarized broad lines

are mainly the AGN which have a powerful central engine and thus a high accretion

rate. A similar result was reached by Trump et al. (2011) when studying accretion

rates in a large sample of several hundred AGN with multiwavelength data, ranging

from the infrared upto the X-ray band. They defined the intrinsic luminosity Lint of

these AGN as the sum disk, as measured in the optical and UV band, of the corona on

top of the accretion disk, as detectable in the X-ray band, and of the re-processed

73


emission as visible in the infrared. In this study, a strong dependence of the broad-line

detectability on the accretion rate was evident. Broad emission lines would only be

created in objects with a high Eddington ratio λ = Lint/LEdd > 0.01. The broad-line

region in these objects will be detectable, either directly in the unobscured Seyfert I, or

in the polarized light in the Type II AGN. At low accretion rates with λ < 0.01,

narrow-line AGN are observed which do not show strong absorption. Thus, in these

cases, neither in normal light nor in polarized spectroscopic observation, is a broad-line

detectable. The non-detectability of the broad-line region would not necessarily break

the unification, although it adds another dependency, the strength of the AGN activity.

More recent studies using the Spitzer Space Telescope come to a different conclusion.

Here, the Spitzer mid-infrared data re used as an indicator of the overall AGN power.

The idea is that the circumnuclear dust which is dominating this energy range acts as

a bolometer for the central engine. Although re-processed, the Spitzer data might thus

give a good proxy for the bolometric luminosity of the AGN, hidden or not.

Comparing the MIR luminosity of 46 radio-loud AGNs from a complete sample with

f2.7 Ghz > 2 Jy with [O III] line luminosities, Dicken et al. (2009) find no major

difference between quasars, narrow-line and weak-line radio galaxies, nor between FR-I

and FR-II. Other studies call into question the claim that the [O III] line indeed

represents the bolometric luminosity of the AGN.

Moving further into the infrared, where absorption should play a lesser role, one finds

for example the forbidden narrow line [O IV] at λ = 25.9 µm. Kraemer et al. (2011)

studied a sample of 40 Seyfert galaxies and found that the ratio of [O III]/[O IV] is

lower for the less luminous sources and for the Seyfert II objects. This indicates that

the [O III] luminosity might after all be affected by absorption. On the other hand,

Kraemer et al. (2011) found that in the Seyfert 1.5 galaxy NGC 4151, only one third of

the [O III] emission seems to arise from the inner ∼ 30 pc. Thus, in order for the

Seyfert II galaxies to have the same [O III] profiles as the unobscured Seyfert I, the

absorbing dust must extend out to large radial distances.

Another way to study differences and similarities between different AGN types is to

look at their spectral shape. Absorption also has an influence on the continuum

emission of AGN. The appearance of the underlying continuum will be changed and in

case one can measure the absorption, the intrinsic spectrum can be recovered. The

74


effect on the intrinsic spectrum will obviously be less pronounced for lower absorption,

and should also be less strong when observing at wavelengths less affected by absorbing

material. The optical domain is strongly affected by material in the line of sight. For

example, in the V-band, for a hydrogen column density NH , measured in atoms per

unit area of line of sight, Predehl & Schmitt (1995) showed that the extinction is

Av = NH1.79×1021 cm−2 mag.

Already the galactic hydrogen column density at high-galactic latitude (|b| > 20) is in

the range NH,Gal = 1020 − 1021 cm−2 and thus can lead to an extinction of the optical

flux of Av = 0.5 mag. Through the galactic plane, the observation of extra-galactic

objects in the optical domain is very difficult to impossible. Moving into the infrared

regime, absorption becomes less efficient. In the NIR bands J (λ = 1.3 µm),

H (λ = 1.7 µm), K (λ = 2.2 µm), and L′ (λ = 3.8 µm), the extinction with respect to

the V-band extinction is AJ = 0.276 AV , AH = 0.176 AV , AK = 0.112 AV , and

AL′ = 0.047 AV (Schlegel et al., 1998).

The problem with recovering the intrinsic spectrum in the infrared to optical domain is

that in this energy range, many different components contribute. Not only is the

underlying continuum of the central engine visible, but also the dust and stars in the

bulge of the host galaxy, and the thermal emission of the accretion disk. An energy

range better suited to study the intrinsic spectrum in Type I and Type II objects and

to determine whether it has the same shape in the X-ray and hard X-ray domain. The

X-rays below 10 keV are significantly affected by absorption in the line of sight.

Although some AGN show some contribution of the surrounding starburst activities

and/or some additional excess below 2 keV, the continuum as seen in the ∼ 3 − 10 keV

is dominated by the emission of a central engine (Section 6.5). The only diversion of

the continuum from a single power law, as expected for Comptonization processes, is

then due to the absorption in the line of sight. Therefore, the X-ray range is well

suited for measuring the column density of the absorber.

X-ray data show that most (but not all) AGNs unabsorbed in the X-rays are Seyfert I

type, and most (but not all) AGNs that are absorbed belong to the Seyfert II group

(eg., Awaki et al. (1991)). The distinction between absorbed and unabsorbed appears

around NH = 1022 cm−2 as a dividing line. At energies above 10 keV, absorption will

75


have very little effect on the emission, unless the absorber is Compton-thick, that is

when the column density in the line of sight significantly exceeds a value of

NH = 1.5× 1024 cm−2. When studying a sample of 25 Seyfert II galaxies, Risaliti et al.

(2002) found that 90% show significant variations of their X-ray absorption column

density. This cannot be explained by the simple torus model for the absorber, instead

one has to assume some clumpiness in the absorbing matter. The most prominent of

those “changing look” AGN is the Seyfert 1.8 NGC 1365. In the X-rays, this source

changes from a Compton-thick to a Compton-thin absorber and back on a monthly

time scale, as shown by Risaliti et al. (2005). The component in the X-ray spectrum

which is thought to arise from reflection on the absorbing medium, does not seem to

vary; thus, the assumption that there is indeed a massive, clumpy absorber at some

distance (∼ 1 kpc) from the central engine seems to be valid. Type II sources are not

the only ones that show strong variable column densities in the X-ray. Seyfert I and

Seyfert 1.5, such as the NGC 4151, MCG-6-30-15, and NGC 3227 do as well (Risaliti,

2010), as does the narrow-line Seyfert I galaxy Mrk 766 (Risaliti et al., 2011).

The combination of X-ray results should tell one whether Type I and Type II AGN are

intrinsically the same; in the softer X-rays, we can measure the absorption strength,

and at the hardest X-rays we can, determine the true intrinsic spectral shape. Early

X-ray surveys seemed to indicate that there indeed is a difference in the intrinsic

continuum spectrum of Seyfert I and Seyfert II, in the sense that the spectra of the

observed sources (NH > 1022 cm−2) appeared flatter than those of Type I AGN. This

has been noticed by Zdziarski et al. (1995) based on data taken in the 2-10 and 50-300

keV band by Ginga and CGRO/OSSE, respectively. The same discrepancy between

the spectra of Seyfert I and Seyfert II was later confirmed for example by Gondek

et al. (1996) using combined EXOSAT, Ginga, HEAO-1, and CGRO/OSSE spectra,

and also by INTEGRAL at hard X-rays above 20 keV, where absorption should not

play a role (Beckmann et al., 2006). A study using data from another hard X-ray

experiment covering the 15-200 keV band (BeppoSAX/PDS) using spectra of 45

Seyfert galaxies has come to a similar conclusion, although the spectra of Seyfert II

appeared to be steeper when considering a possible cutoff in the spectra of Seyfert I

galaxies (Deluit & Courvoisier, 2003).

A problem in measuring the true spectral shape is the complex nature of the intrinsic

hard X-ray spectrum. The inverse Compton emission spectrum, which cuts off

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exponentially at around ∼ 100 keV, can also be altered by the reflection from cooler

material leading to a Compton reflection “hump” around 30 keV. Therefore,

high-quality data are necessary in order to disentangle the different components. Data

from the hard X-ray missions BeppoSAX, INTEGRAL, and Swift seem to indicate

now that the underlying continuum has the same spectral slope when all components

are take care of correctly. Analysis of a sample of 105 Seyfert galaxies using the

spectra collected with BeppoSAX in the 2-200 keV band (Dadina, 2007) provided no

evidence of any spectral slope difference when applying more complex model fitting

including a reflection component. The INTEGRAL data show consistent slopes for the

spectra of of unabsorbed/Type I and absorbed/Type II objects already when a simple

power law model is used. When applying more complex models with geometrical

dependencies, the underlying spectral slope seems to be fully consistent over different

inclination angles (Beckmann et al., 2009). Other studies seem to indicate that the

spectral slope is the same, but the reflection component is of different strength when

comparing Seyfert I and Seyfert II galaxies. Ricci et al. (2011) find in an analysis of

hard X-ray spectra of 165 Seyfert galaxies that the strongest reflection is originating

from Seyfert II galaxies with intrinsic absorption of 1023 cm−2 ≤ NH ≤ 1024 cm−2,

whereas objects with more or less absorption do not show this feature strongly. In the

unified model, this is difficult to explain and requires a complex absorption geometry,

in which the objects with a strong reflection component would have to be an absorber

which covers a high fraction of the X-ray source. Clouds of matter of different sizes

could be a possible solution, in the sense that the strongly reflecting sources display

smaller matter clumps in the vicinity of the X-ray source than the other Seyfert

galaxies. The clump size would lead to a larger surface available for reflection,

assuming that the total absorber mass is about the same.

Concerning the accretion activity, the INTEGRAL-selected sources seem to indicate

that the mass of hard X-ray-selected Seyfert galaxies does not depend on the source

type and is on average ∼ 108 ⊙. But at the same time, the average luminosity of Type

I AGNs is higher than that of Seyfert II, and thus also the Eddington rations of

Seyfert I (λ ≃ 0.06) appear higher than those of Seyfert II with λ ≃ 0.02 (Middleton

et al., 2008) in the local Universe. These values have to be treated with caution, as the

black hole masses were determined using different methods. In addition, the X-ray

luminosity was used as a proxy for the bolometric luminosity with

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Lbol = 2 × L3−1000 keV , assuming that the radio to optical branch of the AGN emits as

much as the inverse Compton branch.

Gallo et al. (2010) found that the accretion rate of AGN appears to be a function of

the black hole mass. In their study, the Eddington ratio seems to be anticorrelated

with MBH , another indication for a violation of the unified model. Numerous

investigations have tried to explain the differences between the Seyfert types, which

cannot be covered in the unified model by differences in the geometry or physical

properties of the absorber. Ramos Almeida et al. (2012) conclude based on a small

sample that the absorbing tori in Seyfert II have a larger covering factor, a lower

optical depth, and are more clumpy than those in Seyfert I. In addition, if we assume

that the absorbing medium is not homogeneous, but rather clumpy, observing a Type I

or Type II AGNs is rather given by the probability of the light of the broad-line region

shining through it. In a clumpy absorber model, a small inclination angle object in

which one observes the AGN disk face-on can also appear as a Seyfert II.

Here one also might face a situation where the overall picture agrees with the unified

model, but the model needs further adjustment and dependence on other parameters

than only orientation and radio-loudness.

8.3 Radio-Loud Versus Radio-Quiet

The simple unification scheme which only considers absorption and beaming is not

sufficient to answer the question of why some sources are strong radio emitters, and

some are not. In other words, what makes the central engine produce a jet.

Radio-quiet does not mean that there is no radio emission at all from the AGN, but

that the radio optical flux ratio is low.Also, a radio-bright source is not necessarily

radio-loud, and not every AGN which is radio-quiet has to be a faint radio source.

There is a dichotomy between radio-loud AGNs (broad-line radio galaxies, radio-loud

quasars, FR-I, FR-II) and radio-quiet AGNs (Seyfert galaxies, LINERs). Most Seyfert

galaxies, although being weak radio emitters, do not seem to harbor a jet. High

resolution observations of the radio cores of Seyfert galaxies by Lal et al. (2011) do not

seem to detect any relativistic beaming, which would be a clear indication of the jet.

At the same time, on parsec scales, Seyfert I and Seyfert II appear to be very similar,

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both appearing to have the same compactness. Also, when comparing the inner parsecs

with the extended kpc-scale radio emission, there does not seem to be a difference

between the Type I and Type II objects, as one would expect from the unified scheme.

A first step to unify the radio-loud objects was made when investigating the properties

of the fainter, core-dominated FR-I with the brighter, lobe-dominated FR-II galaxies.

Perley et al. (1980) studied compact radio sources and pointed out that although their

radio properties are different from the FR-II galaxies, the spectra appear consistent if

relativistic beaming effects are considered. Because the jet in FR-I galaxies to the line

of sight, the radio core emission would be enhanced in these cases by relativistic

beaming. Perley et al. (1980) also previously pointed out that the beaming would

roughly explain the fraction of core-dominated sources among the radio galaxies. This

was further investigated by Orr & Browne (1982), using a simple model for the

intrinsic quasar emission consisting of a core which appears relativistically beamed

with a ratio spectral index of αr = 0 and unbeamed radio lobes with spectral index of

αr = −1. The study showed that unification is possible assuming the same average

core Lorentz factor of γ ≃ 5.

A direct correlation of the radio luminosity with black hole mass had been found in

several investigations. This connection has roughly the form of Lr ∝ M2.5BH as found by

Franceschini et al. (1998) and confirmed in several other studies. Thus, following this

result, strong radio emission, and therefore a powerful relativistic jet is a property

connected to the central engines with the highest mass. The same correlation still

seems to hold for radio-quiet objects (Nelson, 2000). Another possibly related trend

which has been found is the correlation of the mass of the central black hole and the

radio-loudness of the AGN. Laor (2000) discovered that nearly all PG quasars with a

black hole mass MBH > 109 M⊙ are radio-loud, while quasars with

MBH < 3 − 108 M⊙ are practically all radio-quiet. This led to the assumption that the

various types of AGN may be largely set by three basic parameters: MBH , Lbol/LEdd,

and inclination angle. It should be pointed out, though, that other studies could not

find a simple relation between the radio-loudness and MBH . One also has to keep in

mind that the radio-loudness depends strongly on whether the whole radio emission of

an object is integrated, or whether only the core flux is measured.

A result which appears to be counterintuitive is the finding that the radio-loudness is

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anticorrelated with the Eddington ratio λEdd. Ho (2002) studied a sample of 80

galaxies including AGN hosting a SMBH of known mass. This study also included

non-active galaxies, and most masses have been determined in the local Universe. First

of all, the study showed that the relation between the mass of the central engine and

its radio luminosity might not be as simple as previously indicated. The Lr ∝ M2.5BH

law might rather describe an upper envelope than presenting a real correlation, as

there are many objects far away but below this line, whereas no object is found above.

This means that the relation appears to be rather as Lr <∝ M2.5BH and would just

indicate a maximum of radiative jet power possible for a given mass of the central

engine. The anticorrelation of radio-loudness and λEdd might reflect the fact that in

objects which are accreting at a low rate, the accretion disk itself is not very

prominent. Therefore, a low λEdd leads to a low thermal disk emission and to a weak

blue bump. As both components dominate the optical/UV band, these objects appear

as “radio-loud”, although they might rather be called “optically quiet”. Another

explanation for the anticorrelation which indicates that the objects with highest

Eddington ratio are the least radio dominated, might be that this rather reflects the

anticorrelation between Eddington ratio and MBH as derived by Gallo et al. (2010).

Thus, the intrinsic mechanism could be a result of smaller black holes accreting more

efficiently. Larger black holes are more likely to produce a significant jet emission; this

would thus appear as an anticorrelation between radio-loudness and Eddington ratio.

Sikora et al. (2007) extended the study of the anticorrelation of Eddington ratio and

radio-loudness to a larger sample of radio-selected AGNs, including broad-line radio

galaxies and FR-I radio galaxies. Including these objects, the scatter of the correlation

of Eddington ratio versus radio-loudness gets much wider. For a given Eddington ratio,

the radio-loudness can fall in a range 5 orders of magnitude wide, and vice versa. The

λ−R∗ anticorrelation is then explained by effects of the spin of the central black hole:

powerful jets can be produced when rotational energy of the central engine can be

extracted via interaction with an external magnetic field, for example from the

accretion disk. This is similar to what one observes from galactic black holes when

they reside in the so-called low/hard state (i.e., low flux and a hard X-ray spectrum).

Sikora et al. (2007) (and many others) use the optical luminosity as a proxy for the

bolometric luminosity, by performing Lbol = 10 × LB. This makes this bolometric

luminosity highly sensitive to the accretion disk’s thermal emission profile, that is, the

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big blue bump, which is different from the bolometric luminosity based on a more

complete model. For example, the presence of inverse Compton emission in the X-ray

domain is a substantial luminosity component. And, absorption can strongly affect the

observed optical flux. Another issue with studying radio-loudness and radio luminosity

of radio galaxies is the question whether only the core luminosity should be used or if

the total radio emission, including the core and the lobes, gives a better estimate of the

overall jet power. The gap between radio-loud and radio-quiet sources appears smaller

when using only the core flux (White et al., 2000), and the core dominated sources

(FR-I) then show a lower radio-loudness than the lobe dominated ones (Rafter et al.,

2011). In a recent study the sample of Sikora et al. (2007) has been investigated using

only core luminosities. Broderick & Fender (2011) find that the λ−R∗ anticorrelation

becomes less pronounced, as one would expect if the black hole spin is indeed the

driving parameter. Instead or in addition to the spin, environmental density or the

black hole mass might again be important here. Broderick & Fender (2011) have

revisited the radio-loud/radio-quiet dichotomy by using a black hole mass normalized

core-only radio luminosity as opposed to the total spatially integrated luminosity.

Their investigation was motivated by the knowledge that the bolometric luminosity, jet

power and black hole mass are interrelated as demonstrated by the ubiquitous

appearance of the black hole fundamental plane. They find that while the bimodal

nature of the AGN population sampled is preserved, the magnitude of the separation is

significantly reduced. Specifically, they find that FR-I and BLRG are on average more

radio-loud than Seyferts and LINER by about 1.6 dex. Recently a new approach has

been applied to try and unify radio-loud and radio-quiet AGN. Garofalo et al. (2010)

consider the relative spin of the central black hole with respect to the accretion disk to

be the crucial factor here. In their scenario, AGNs would start with a black hole which

has a very different, even retrograde, spin with respect to the accretion disk, leading to

strong interaction with the disk and thus strong jets. As the black hole is spun up in

the direction of the accretion disk, the interaction of the rotating black holes with their

magnetospheres becomes less efficient and the jet weakens. Thus, the highest prograde

black hole spins might be discovered in the least active AGN (Garofalo et al., 2010).

This scenario is supported by a theoretical approach of Daly (2009, 2011) determining

the black hole spin that is model-independent, but assumes that spin changes only by

extraction of the reducible black hole mass. This model applied to a small subset of

powerful radio galaxies finds indeed that they harbor low spinning black holes. Further

81


Figure 8.2: Schematic representation of our understanding of the AGN phenomenonin the unified scheme. The type of object one sees depends on the viewing angle,whether or not the AGN produces a significant jet emission, and how powerful thecentral engine is. Radio-loud objects are generally thought to display symmetric jet

emission.

observational support comes from a study of FR-I galaxies, which show low Eddington

ratios (Lbol/LEdd < 0.01) but imply rapidly spinning black holes with j > 0.9 (Wu

et al., 2011). Here,

j = JcGM2

BH

is the dimensionless angular momentum of the black hole and J is its angular

momentum. A case with j close to 1 would represent the case of a Kerr black hole,

while j close to 0 can be treated as a non-rotating Schwarzschild black hole. The

separation of AGNs into sources with and without a jet might not be as clean as

assumed in the past. The gap between radio-loud and radio-quiet appears less

pronounced the more high-quality data of AGN one studies. The dependency of radio

loudness on the Eddington ratio is weak or absent, and the influence of black hole spin,

which is difficult to estimate in the first place, might or might not solve the problem of

which sources produce jets.

82


Type Optical Lines Radio-Quiet Radio-Loud

Type I Broad and narrow linesSeyfert ISeyfert 1.5NLS1

FSRQ, SSRQ, BLRG

Type II Narrow lines onlySeyfert 1.8, 1.9,2LINERs/LLAGN

NLRG, Type II QSO

Type 0 No LinesSgrA∗?Dormant AGN

BL Lac, OVV

Table 8.1: The general unification scheme of AGN, based on the emission lines visiblein the optical domain.

(Tadhunter, 2008)

Notwithstanding all of these issues, an overall unification scheme (Table 8.1) has

evolved over the years which is schematically represented in Figure 8.2. It shows the

radio-loud AGN which are assumed to display a jet in the upper half, and the

radio-quiet objects in the lower part of the figure.

8.4 Breaking the Unification

The basic unified model for AGN predicts differences in appearance only based on

different orientation toward the observer. This causes different absorption effects

intrinsic to the innermost regions of the AGN, as well as geometrical effects regarding

the beaming of the emission. The previous subsections presented several observational

results that do not seem to fit into this scheme.

In many but not all Seyfert II galaxies one can find a hidden broad-line region when

studying the objects in polarized light. The remaining Seyfert II objects, which do not

show any broad-line region, even when observed in polarized light, might simply have

weak BLR emission compared to the underlying continuum, which would make

detection difficult. One explanation might be that the power of the central engine is

not large enough to sufficiently illuminate the BLR. This might also be used as an

explanation, why BL Lac objects, the weakest blazars, do not show any emission lines.

What this comes down to is that one needs to add the total power of the central engine

as a parameter to the unification model in order to make it work. If all Seyfert II cores

83


Figure 8.3: Anticorrelation between the X-ray variability amplitude and the black holemass, with the masses on NGC 4593 and IC 4329A being upper limits. The filled circlesdenote the objects with the black hole masses measured from reverberation mapping.Triangles are based on stellar velocity dispersion measurements, and the open squaredenotes the AGN QPO RE J1034C396. The line indicates the best-fit linear relation.

The intrinsic dispersion of this fit is 0.2 dex.(Zhou et al., 2010)

are absorbed, this should be observable especially in the X-ray domain, where the

emission is supposed to originate from the innermost region around the accreting black

hole. But, there are examples of Type II AGNs, which indeed show no measurable

absorption at soft X-rays, like NGC 3147 and NGC 4698 (Pappa et al., 2001). These

cases might represent the same effect as for the Seyfert II without a detectable

broad-line region even when viewed in polarized light. Also here, there might only be a

weak broad-line region because the intrinsic power of the AGN core is low.

If luminosity is invoked as a parameter in the unified model, one can explain further

effects. For example, in hard X-ray surveys of Seyfert galaxies one observes an

anticorrelation of the fraction of absorbed sources with luminosity. While for an X-ray

luminosity of L20−100 keV = 1042 erg s−1 about 65% of the Seyfert galaxies show an

intrinsic absorption NH > 1022 cm−2, at L20−100 keV = 1045 erg s−1 only 35% are

absorbed (Beckmann et al., 2009). An explanation for this coupling can be the

scenario of a “receding torus” as proposed by Lawrence & Papadakis (1993). The

84


radiation pressure of the central engine pushes the absorbing material out. If one

assumes a simple torus as absorber, one can find a correlation between the luminosity

of the AGN core and the inner radius of the torus of the form Rinner ∝√Lbol.

Assuming that the inner radius is determined by the limit at which the AGN

luminosity can evaporate the dust at a temperature of T = 1000 K, the radius where

this is the blackbody equilibrium temperature is roughly at

Rinner ≃ Lbol 4 × 10−46 erg−1 s pc. For a fixed height of the absorbing torus, this will

lead to a wider opening angle under which the broad-line region is visible with

increasing luminosity. In other words, the fraction of unabsorbed sources we observe

increases with luminosity as observed in the X-rays. Thus, these breaks in unification

can be explained by an additional dependence on luminosity.

Another challenge for the unified model are the differences found in the luminosity

functions of different AGN types. The luminosity function gives a measurement of the

density of sources of a given luminosity per unit volume. In the case of blazars, there

appears to be a difference in the luminosity functions of the bright FSRQ, which show

broad lines in their spectrum, and the fainter, high-frequency peaked BL Lac objects

(HBL). While FSRQ and low frequency peaked BL Lac objects seem to have been

more numerous and/or luminous in the past, that is, they show a positive evolution,

high-frequency BL Lac objects show either no or slightly negative evolution, making

them as numerous and luminous in the local Universe as at redshifts z & 0.3, or even

more abundant now than in the past (e.g., Beckmann et al. (2003)). This presents a

violation of the unification model, which would predict that the distribution in space

does not depend on the AGN type.

As described by Bottcher & Dermer (2002), one way to unify the blazar classes would

be a transformation of FSRQ and LBLs into HBLs as the blazars grow older. In this

model, blazars start as powerful FSRQ with jets of high-energy densities. Strong

cooling limits the electron energies leading to cutoff frequencies for the synchrotron

component at optical wavelengths and for the inverse Compton component in the GeV

energy range. By the time the source of the jet gets less powerful the energy density

within the jet decreases. The cooling efficiency decreases as well resulting in higher

cutoff frequencies for HBLs. The shift of the cutoff frequencies to higher energies is

therefore accompanied by decreasing bolometric luminosities, which is evident from the

decrease of the luminosities in the radio, near IR, and optical bands. Due to the

85


increasing peak frequencies of the synchrotron branch more energy is released in the

X-ray band and the X-ray luminosity increases quite in contrast to the luminosities at

shorter frequencies.

The comparison of Seyfert I and Seyfert II luminosity functions appears to be difficult.

While complete samples of Type I AGN can be compiled relatively easy, Seyfert II

samples often lack completeness or include other narrow-line objects, like LINER or H

II regions. Turning once more to the hard X-ray band in order to achieve complete

samples of AGN in an energy range not affected by absorption, we see that indeed the

luminosity functions of absorbed and unabsorbed sources are similar, although there is

some indication for Seyfert II to dominate at the low luminosity end with

LX < 1043 erg s−1, while Seyfert I are contributing stronger to the high-luminosity

objects. Again, luminosity (or accretion efficiency in terms of Eddington ratio) would

have to be included as a parameter of the unification model in order to explain this

discrepancy.

Variability studies can be used in order to verify the unified model. If all sources are

intrinsically similar, then they should show variability patterns independent of their

source type. Apparently, this is not the case, as the variability seems to be a function

of the mass of the black hole. An anticorrelation of X-ray variability with luminosity in

AGN is observed in the sense that the more luminous AGN are less variable (e.g., Barr

& Mushotzky (1986), Beckmann et al. (2007), Lawrence & Papadakis (1993)). The

same effect has been observed in the UV band and in the optical domain (e.g., de Vries

et al. (2005)), although narrow-line Seyfert I galaxies apparently show the opposite

behavior (e.g., Turner et al. (1999)). Papadakis (2004) explains this correlation as a

connection between luminosity and the mass of the central black hole MBH . This may

be explained if more luminous sources are physically larger in size, so that they are

actually varying more slowly. Alternatively, they may contain more independently

flaring regions and so have a genuinely lower amplitude. The observed correlation

might reflect the anticorrelation of variability and black hole mass. This relation has

been well studied in the 2-10 keV X-ray band. The X-ray variability amplitude, as

measured in the root mean square (RMS) value, is anticorrelated with MBH , as shown

in Figure 8.3. The fit shows a correlation of the form:

log(

MBHM⊙

)

= (4.85 ± 0.20) − (1.05 ± 0.08) log σ2rms.

86

Bibliography

Uttley & Mchardy (2004) explained the anticorrelation of variability and MBH by

suggesting that the X-rays are presumably produced in optically thin material close to

the central black hole, at similar radii (i.e., in Schwarzschild radii, RS) in different

AGNs. As RS = 2GMBH/c2, longer time scales for the variability are expected for the

more massive central engines. These studies might still be affected in part by

absorption. Recent hard X-ray studies of AGN, where absorption should not play a

role, do not seem to detect a difference in variability patterns of Type I and Type II

AGN (Soldi et al., 2009). If this result holds, it would be a strong support for the

unified model.

87

Bibliography

Abazajian, K. N., Adelman-McCarthy, J. K., Agueros, M. A., et al. 2009, ApJS, 182,

543

Antonucci, R. 1993, ARA&A, 31, 473

Antonucci, R. R. J., & Miller, J. S. 1985, ApJ, 297, 621

Awaki, H., Kunieda, H., Tawara, Y., & Koyama, K. 1991, PASJ, 43, L37

Barr, P., & Mushotzky, R. F. 1986, Nature, 320, 421

Barthel, P. D. 1989, ApJ, 336, 606

Barvainis, R. 1987, ApJ, 320, 537

Beckmann, V., Barthelmy, S. D., Courvoisier, T. J.-L., et al. 2007, A&A, 475, 827

Beckmann, V., Engels, D., Bade, N., & Wucknitz, O. 2003, A&A, 401, 927

Beckmann, V., Gehrels, N., Shrader, C. R., & Soldi, S. 2006, ApJ, 638, 642

Beckmann, V., Soldi, S., Ricci, C., et al. 2009, A&A, 505, 417

Blandford, R. D., & Konigl, A. 1979, ApJ, 232, 34

Blandford, R. D., & Rees, M. J. 1978, in BL Lac Objects, ed. A. M. Wolfe, 328–341

Blandford, R. D., & Znajek, R. L. 1977, MNRAS, 179, 433

Boksenberg, A., Shortridge, K., Allen, D. A., et al. 1975, MNRAS, 173, 381

Bottcher, M., & Dermer, C. D. 2002, ApJ, 564, 86

Bottcher, M., Basu, S., Joshi, M., et al. 2007, ApJ, 670, 968

88

Bibliography

Brenneman, L. W., & Reynolds, C. S. 2009, ApJ, 702, 1367

Bridle, A. H., Hough, D. H., Lonsdale, C. J., Burns, J. O., & Laing, R. A. 1994, AJ,

108, 766

Broderick, J. W., & Fender, R. P. 2011, MNRAS, 417, 184

Byram, E. T., Chubb, T. A., & Friedman, H. 1966, Science, 152, 66

Carilli, C. L., Perley, R. A., Bartel, N., & Sorathia, B. 1996, in Astronomical Society

of the Pacific Conference Series, Vol. 100, Energy Transport in Radio Galaxies and

Quasars, ed. P. E. Hardee, A. H. Bridle, & J. A. Zensus, 287

Chini, R., Kreysa, E., & Biermann, P. L. 1989, A&A, 219, 87

Cutri, R. M., Nelson, B. O., Francis, P. J., & Smith, P. S. 2002, in Astronomical

Society of the Pacific Conference Series, Vol. 284, IAU Colloq. 184: AGN Surveys,

ed. R. F. Green, E. Y. Khachikian, & D. B. Sanders, 127

Dadina, M. 2007, A&A, 461, 1209

Daly, R. A. 2009, The Astrophysical Journal Letters, 691, L72

Daly, R. A. 2011, MNRAS, 414, 1253

de Grijp, M. H. K., Lub, J., & Miley, G. K. 1987, A&AS, 70, 95

de Vries, W. H., Becker, R. H., White, R. L., & Loomis, C. 2005, AJ, 129, 615

Deluit, S., & Courvoisier, T. J.-L. 2003, A&A, 399, 77

Dermer, C. D., Schlickeiser, R., & Mastichiadis, A. 1992, A&A, 256, L27

Dicken, D., Tadhunter, C., Axon, D., et al. 2009, ApJ, 694, 268

Done, C., & Gierlinski, M. 2004, in COSPAR Meeting, Vol. 35, 35th COSPAR

Scientific Assembly, ed. J.-P. Paille, 991

Draine, B. T., & Lee, H. M. 1984, ApJ, 285, 89

Edelson, R. A., & Malkan, M. A. 1987, ApJ, 323, 516

Elitzur, M. 2006, New Astronomy Review, 50, 728

Elvis, M., Wilkes, B. J., McDowell, J. C., et al. 1994, ApJS, 95, 1

89

Bibliography

Fanaroff, B. L., & Riley, J. M. 1974, MNRAS, 167, 31P

Forman, W., Jones, C., Cominsky, L., et al. 1978, ApJS, 38, 357

Franceschini, A., Vercellone, S., & Fabian, A. C. 1998, MNRAS, 297, 817

Francis, P. J., Hewett, P. C., Foltz, C. B., et al. 1991, ApJ, 373, 465

Gallo, E., Treu, T., Marshall, P. J., et al. 2010, ApJ, 714, 25

Garofalo, D., Evans, D. A., & Sambruna, R. M. 2010, MNRAS, 406, 975

Ge, J., & Owen, F. N. 1994, AJ, 108, 1523

George, I. M., & Fabian, A. C. 1991, MNRAS, 249, 352

Giacconi, R., Gorenstein, P., Gursky, H., et al. 1967, ApJ, 148, L129

Gondek, D., Zdziarski, A. A., Johnson, W. N., et al. 1996, MNRAS, 282, 646

Groves, B., & Kewley, L. 2008, in Astronomical Society of the Pacific Conference

Series, Vol. 390, Pathways Through an Eclectic Universe, ed. J. H. Knapen, T. J.

Mahoney, & A. Vazdekis, 283

Hartman, R. C., Bertsch, D. L., Fichtel, C. E., et al. 1992, ApJ, 385, L1

Heisler, C. A., Lumsden, S. L., & Bailey, J. A. 1997, Nature, 385, 700

Ho, L. C. 2002, ApJ, 564, 120

Ho, L. C. 2008, ARA&A, 46, 475

Hough, J. H., Brindle, C., Wills, B. J., Wills, D., & Bailey, J. 1991, ApJ, 372, 478

Huchra, J., & Burg, R. 1992, ApJ, 393, 90

Hughes, D. H., Robson, E. I., Dunlop, J. S., & Gear, W. K. 1993, MNRAS, 263, 607

Jiehao Huang, Q. G. 2002, ArXiv Astrophysics e-prints, astro-ph/0207207

Kraemer, S. B., Schmitt, H. R., Crenshaw, D. M., et al. 2011, ApJ, 727, 130

Kristian, J. 1973, ApJ, 179, L61

Kuraszkiewicz, J. K., Wilkes, B. J., Hooper, E. J., et al. 2003, ApJ, 590, 128

90

Bibliography

Lal, D. V., Shastri, P., & Gabuzda, D. C. 2011, ApJ, 731, 68

Laor, A. 2000, ApJ, 543, L111

Lawrence, A., & Papadakis, I. 1993, ApJ, 414, L85

Leahy, J. P., & Perley, R. A. 1991, AJ, 102, 537

Low, J., & Kleinmann, D. E. 1968, AJ, 73, 868

Maraschi, L., Ghisellini, G., & Celotti, A. 1992, ApJ, 397, L5

Middleton, M., Done, C., & Schurch, N. 2008, MNRAS, 383, 1501

Miller, J. S., & Goodrich, R. W. 1990, ApJ, 355, 456

Mushotzky, R. F. 1982, ApJ, 256, 92

Nelson, C. H. 2000, ApJ, 544, L91

Netzer, H. 2013, The Physics and Evolution of Active Galactic Nuclei (New York City,

NY: Cambridge University Press)

Nicholson, K. L., Reichert, G. A., Mason, K. O., et al. 1998, MNRAS, 300, 893

Orr, M. J. L., & Browne, I. W. A. 1982, MNRAS, 200, 1067

Osterbrock, D. E. 1984, QJRAS, 25, 1

Owen, F. N., & Laing, R. A. 1989, MNRAS, 238, 357

Papadakis, I. E. 2004, MNRAS, 348, 207

Pappa, A., Georgantopoulos, I., Stewart, G. C., & Zezas, A. L. 2001, MNRAS, 326, 995

Perley, R. A., Fomalont, E. B., & Johnston, K. J. 1980, AJ, 85, 649

Perley, R. A., Willis, A. G., & Scott, J. S. 1979, Nature, 281, 437

Pounds, K. A., & Turner, T. J. 1988, Mem. Soc. Astron. Italiana, 59, 261

Predehl, P., & Schmitt, J. H. M. M. 1995, A&A, 293, 889

Prestage, R. M., & Peacock, J. A. 1988, MNRAS, 230, 131

Punch, M., Akerlof, C. W., Cawley, M. F., et al. 1992, Nature, 358, 477

91

Bibliography

Quinn, J., Akerlof, C. W., Biller, S., et al. 1996, ApJ, 456, L83

Rafter, S. E., Crenshaw, D. M., & Wiita, P. J. 2011, AJ, 141, 85

Ramos Almeida, C., Levenson, N. A., Alonso-Herrero, A., et al. 2012, Journal of

Physics Conference Series, 372, 012004

Ricci, C., Walter, R., Courvoisier, T. J.-L., & Paltani, S. 2011, A&A, 532, A102

Richards, G. T., Nichol, R. C., Gray, A. G., et al. 2004, ApJS, 155, 257

Risaliti, G. 2010, in IAU Symposium, Vol. 267, IAU Symposium, ed. B. M. Peterson,

R. S. Somerville, & T. Storchi-Bergmann, 299–306

Risaliti, G., Elvis, M., Fabbiano, G., Baldi, A., & Zezas, A. 2005, The Astrophysical

Journal Letters, 623, L93

Risaliti, G., Elvis, M., & Nicastro, F. 2002, ApJ, 571, 234

Risaliti, G., Nardini, E., Salvati, M., et al. 2011, MNRAS, 410, 1027

Rodrıguez-Ardila, A., & Mazzalay, X. 2006, MNRAS, 367, L57

Rowan-Robinson, M. 1977, ApJ, 213, 635

Sabra, B. M., Shields, J. C., & Filippenko, A. V. 2000, ApJ, 545, 157

Sanders, D. B., Phinney, E. S., Neugebauer, G., Soifer, B. T., & Matthews, K. 1989,

ApJ, 347, 29

Sanders, D. B., Soifer, B. T., Elias, J. H., et al. 1988, ApJ, 325, 74

Saxton, R. D., Turner, M. J. L., Williams, O. R., et al. 1993, MNRAS, 262, 63

Scheuer, P. A. G., & Readhead, A. C. S. 1979, Nature, 277, 182

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

Schmidt, M., & Green, R. F. 1983, ApJ, 269, 352

Schodel, R., Ott, T., Genzel, R., et al. 2002, Nature, 419, 694

Scott, J. E., Kriss, G. A., Brotherton, M., et al. 2004, ApJ, 615, 135

Seyfert, C. K. 1943, ApJ, 97, 28

92

Bibliography

Sikora, M., Stawarz, L., & Lasota, J.-P. 2007, ApJ, 658, 815

Smith, R. A. N., Page, M. J., & Branduardi-Raymont, G. 2008, A&A, 490, 103

Soldi, S., Ponti, G., Beckmann, V., & Lubinski, P. 2009, in The Extreme Sky:

Sampling the Universe above 10 keV, 31

Spinoglio, L., & Malkan, M. A. 1989, ApJ, 342, 83

Swanenburg, B. N., Hermsen, W., Bennett, K., et al. 1978, Nature, 275, 298

Tadhunter, C. 2008, New Astronomy Review, 52, 227

Tengstrand, O., Guainazzi, M., Siemiginowska, A., et al. 2009, A&A, 501, 89

Tran, H. D. 2003, ApJ, 583, 632

Tran, H. D., Miller, J. S., & Kay, L. E. 1992, ApJ, 397, 452

Trump, J. R., Impey, C. D., Taniguchi, Y., et al. 2009, ApJ, 706, 797

Trump, J. R., Impey, C. D., Kelly, B. C., et al. 2011, ApJ, 733, 60

Tsvetanov, Z. I., Morse, J. A., Wilson, A. S., & Cecil, G. 1996, ApJ, 458, 172

Tucker, W., & Giacconi, R. 1985, The X-ray universe (Harvard University Press)

Turner, T. J., George, I. M., Nandra, K., & Turcan, D. 1999, ApJ, 524, 667

Urry, C. M., & Padovani, P. 1995, PASP, 107, 803

Uttley, P., & Mchardy, I. M. 2004, Progress of Theoretical Physics Supplement, 155,

170

Vagnetti, F., Cavaliere, A., & Giallongo, E. 1991, ApJ, 368, 366

Vanden Berk, D. E., Richards, G. T., Bauer, A., et al. 2001, AJ, 122, 549

Ward, M. J., Done, C., Fabian, A. C., Tennant, A. F., & Shafer, R. A. 1988, ApJ, 324,

767

Weedman, D. W. 1973, ApJ, 183, 29

White, R. L., Becker, R. H., Gregg, M. D., et al. 2000, ApJS, 126, 133

Wills, B. J., Wills, D., Evans, II, N. J., et al. 1992, ApJ, 400, 96

93

Bibliography

Wu, Y.-Z., Zhang, E.-P., Liang, Y.-C., Zhang, C.-M., & Zhao, Y.-H. 2011, ApJ, 730,

121

Zdziarski, A. A., Johnson, W. N., Done, C., Smith, D., & McNaron-Brown, K. 1995,

ApJ, 438, L63

Zdziarski, A. A., Lubinski, P., & Smith, D. A. 1999, MNRAS, 303, L11

Zhou, X.-L., Zhang, S.-N., Wang, D.-X., & Zhu, L. 2010, ApJ, 710, 16

94

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