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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
Chapter 1. Introduction
• 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
Chapter 1. Introduction
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
Chapter 1. Introduction
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
Chapter 1. Introduction
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
Chapter 1. Introduction
• 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
Chapter 1. Introduction
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
Chapter 1. Introduction
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
Chapter 1. Introduction
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
Chapter 1. Introduction
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
Chapter 2. Non-Thermal Processes
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
Chapter 2. Non-Thermal Processes
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
Chapter 2. Non-Thermal Processes
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
Chapter 2. Non-Thermal Processes
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
Chapter 2. Non-Thermal Processes
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
Chapter 2. Non-Thermal Processes
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
Chapter 2. Non-Thermal Processes
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
Chapter 2. Non-Thermal Processes
⟨
∆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
Chapter 2. Non-Thermal Processes
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
Chapter 2. Non-Thermal Processes
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
Chapter 2. Non-Thermal Processes
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
Chapter 2. Non-Thermal Processes
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
Chapter 2. Non-Thermal Processes
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
Chapter 3. The IR and Sub-mm Regime
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
Chapter 3. The IR and Sub-mm Regime
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
Chapter 3. The IR and Sub-mm Regime
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
Chapter 3. The IR and Sub-mm Regime
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
Chapter 3. The IR and Sub-mm Regime
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
Chapter 4. The Radio Regime
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
Chapter 4. The Radio Regime
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
Chapter 4. The Radio Regime
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
Chapter 4. The Radio Regime
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
Chapter 4. The Radio Regime
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
Chapter 4. The Radio Regime
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
Chapter 4. The Radio Regime
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
Chapter 4. The Radio Regime
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
Chapter 5. The Optical-UV Regime
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
Chapter 5. The Optical-UV Regime
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
Chapter 5. The Optical-UV Regime
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
Chapter 5. The Optical-UV Regime
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
Chapter 5. The Optical-UV Regime
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
Chapter 5. The Optical-UV Regime
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
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Chapter 5. The Optical-UV Regime
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
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Chapter 5. The Optical-UV Regime
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
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Chapter 6. The X-Ray Regime
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
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Chapter 6. The X-Ray Regime
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)
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Chapter 6. The X-Ray Regime
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
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Chapter 6. The X-Ray Regime
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
Chapter 6. The X-Ray Regime
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.
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Chapter 6. The X-Ray Regime
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
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Chapter 6. The X-Ray Regime
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,
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Chapter 6. The X-Ray Regime
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.
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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
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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
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Chapter 7. The γ-Ray Regime
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.
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Chapter 7. The γ-Ray Regime
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.
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Chapter 7. The γ-Ray Regime
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
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Chapter 7. The γ-Ray Regime
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.
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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
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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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Chapter 8. The Unified Model of AGNs
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
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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
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Chapter 8. The Unified Model of AGNs
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
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Chapter 8. The Unified Model of AGNs
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
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Chapter 8. The Unified Model of AGNs
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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Chapter 8. The Unified Model of AGNs
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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Chapter 8. The Unified Model of AGNs
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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Chapter 8. The Unified Model of AGNs
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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Chapter 8. The Unified Model of AGNs
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
80
Chapter 8. The Unified Model of AGNs
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
Chapter 8. The Unified Model of AGNs
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
Chapter 8. The Unified Model of AGNs
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
Chapter 8. The Unified Model of AGNs
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
Chapter 8. The Unified Model of AGNs
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
Chapter 8. The Unified Model of AGNs
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
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