Universidad Nacional Autónoma de México
Posgrado en Astrof́ısica
Instituto de Astronomı́a
X-ray analysis of Seyfert 1 galaxies with optical
polarization:
a test for unification models
TESIS
que para optar por el grado de:
Doctora en Ciencias (Astrof́ısica)
Presenta:
Miriam Eugenia Gudiño Yáñez
Tutoras:
Dra. Anna Lia Longinotti[1]
Dra. Elena Jiménez-Bailón[2]
Miembros del Comité Tutor:
Dr. Takamitsu Miyaji[1]
Dr. Tomás Verdugo[1]
Dra. Alenka Negrete[1]
[1]Instituto de Astronomı́a, UNAM
[2]Quasar Science Resource S.L. European Space Astronomy Centre (ESAC)
Ciudad Universitaria, CDMX, México. Febrero, 2025.
 
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Acknowledgments
First and foremost, I would like to express my gratitude to my advisors, Dra. Anna Lia
Longinotti and Dra. Elena Jiménez-Bailón, for their invaluable guidance and feedback through-
out this doctorate.
I am profoundly grateful to my friends and family. This milestone would not have been
possible without their support and inspiration. With every accomplishment in my life, I carry
the cheering, encouragement, and wisdom of Karla, Rodrigo, and Jorge, thank you for always
being there.
To my incredible partner, Leo, I want to specially acknowledge your steadfast presence
through thick and thin. Thank you for your unwavering support and love.
Lastly, to my most joyous buddies —Lucas, Félix, Julieta, and Andrés— you keep my mind
constantly curious and imaginative. You are the most wondrous and amazing little Viking
crew. And Isa, you are such an impressive lady, you are the big cousin the crew needs to keep
growing into their amazing selves. I love you all oh so very much. Always remember this quote
by J.R.R. Tolkien: ”Not all those who wander are lost.” I deeply believe that a mind that
wanders will always be surrounded by great discoveries.
Thank you.
3
Contents
1 Active Galactic Nuclei 13
1.1 Observational History . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 13
1.2 The main physical structure of AGN . . . . . . . . . . . . . . . . . . . . . . . . 16
1.2.1 The central region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 16
1.2.2 Line emitting regions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 18
1.2.2.1 The Broad Line Region . . . . . . . . . . . . . . . . . . . . . . 18
1.2.2.2 The Narrow Line Region . . . . . . . . . . . . . . . . . . . . . . 19
1.2.3 The molecular torus . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 20
1.2.4 The radio Jet . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 21
1.3 AGN Spectral Energy Distribution . . . . . . . . . . . . . . . . . . . . . . . . . 22
1.3.1 Radio and Infrared . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 24
1.3.2 Optical and UV . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 24
1.3.3 X-ray spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 24
1.3.4 γ−ray emission . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 26
1.4 AGN classification and the Unified Model . . . . . . . . . . . . . . . . . . . . . 27
1.4.1 AGN classification . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 27
1.4.1.1 Types 1 and 2: classification based on optical emission lines . . 28
1.4.1.2 X-ray classification: obscured and unobscured . . . . . . . . . . 29
1.4.2 The Unified Model . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 30
1.4.2.1 Modifications and challenges to the Unified Mode . . . . . . . . 31
2 Optical polarization in Seyfert galaxies 35
2.1 Optical polarization . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 35
2.2 Polarized Seyferts galaxies . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 38
2.2.1 Towards a unification scheme for polarized Seyfert galaxies . . . . . . . . 39
2.2.1.1 Characteristics of polarized type 1 Syeferts . . . . . . . . . . . . 40
2.2.1.2 Polarization in the context of the Unified Model, interpreting
polarization position angle . . . . . . . . . . . . . . . . . . . . . 42
2.3 X-ray analysis, testing the unification scheme . . . . . . . . . . . . . . . . . . . 43
2.3.1 Cold Absorber . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 44
2.3.2 Warm Absorber . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 46
2.3.3 Hypothesis: Is X-ray absorption related to the polarization regions? . . . 48
2.4 Objectives . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 50
3 X-ray analysis of optically polarized Seyferts 51
3.1 Sample selection of polarized Seyfert . . . . . . . . . . . . . . . . . . . . . . . . 51
3.1.1 Weak polarization . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 53
3.2 XMM-Newton . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 53
3.2.1 The European Photon Imaging Camera (EPIC) . . . . . . . . . . . . . . 53
3.3 Processing the XMM-Newton data . . . . . . . . . . . . . . . . . . . . . . . . . . 55
4
CONTENTS 5
3.3.1 Data processing . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 57
3.4 Spectral Models . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 58
3.4.1 Galactic neutral absorption . . . . . . . . . . . . . . . . . . . . . . . . . 59
3.4.2 Power law continuum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59
3.4.3 Fe emission lines . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59
3.4.4 Soft Excess . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59
3.4.5 Absorption . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 62
3.4.5.1 Cold Absorption . . . . . . . . . . . . . . . . . . . . . . . . . . 62
3.4.5.2 Warm Absorption . . . . . . . . . . . . . . . . . . . . . . . . . 63
3.5 Methodology of the X-ray Spectral Analysis . . . . . . . . . . . . . . . . . . . . 64
3.6 Statistical tests . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 65
3.6.1 F-test . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 65
3.6.2 Akaike Information Criterion . . . . . . . . . . . . . . . . . . . . . . . . 65
3.7 Sherpa & PyXspec: fitting packages . . . . . . . . . . . . . . . . . . . . . . . . . 66
3.7.1 Sherpa . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 66
3.7.1.1 Results obtained with Sherpa . . . . . . . . . . . . . . . . . . . 68
3.7.2 PyXspec . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 69
3.7.3 Ionized reflection . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 72
4 Results and discussion 75
4.1 Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 75
4.1.1 Polar-polarized sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . 79
4.1.2 Equatorial-polarized sources . . . . . . . . . . . . . . . . . . . . . . . . . 79
4.1.3 Weakly-polarized sources . . . . . . . . . . . . . . . . . . . . . . . . . . . 80
4.2 Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 80
4.2.1 Considerations on the sample . . . . . . . . . . . . . . . . . . . . . . . . 80
4.2.2 The X-ray spectrum of optically polarized Seyfert 1: a first order descrip-
tion. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 82
4.2.3 The X-ray absorption in optically polarized Seyfert 1 . . . . . . . . . . . 83
4.2.4 Interpreting the unified scheme by Smith through X-ray absorption . . . 85
4.3 Notes on individual sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 89
5 Conclusions and future work 95
5.1 Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95
5.2 Future work . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 96
A Spectral Analysis 107
B Spatial Analysis 112
List of Figures
1.1 The physical structure of AGN . . . . . . . . . . . . . . . . . . . . . . . . . . . 16
1.2 Optical spectra of NGC 4151 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19
1.3 Ionization cones in NGC 5728 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 20
1.4 Resolved image of the torus . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 21
1.5 The radio emission of M87. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 22
1.6 SED of an AGN . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 23
1.7 X-ray spectrum of an AGN . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 25
1.8 Two different spectral shapes for ionized and neutral reflection . . . . . . . . . . 26
1.9 Example of soft excess in Was 45 . . . . . . . . . . . . . . . . . . . . . . . . . . 27
1.10 Radio loud sources types FRII and FRI . . . . . . . . . . . . . . . . . . . . . . . 28
1.11 AGN Types 1 and 2 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 29
1.12 X-ray absorption in the AGN spectra . . . . . . . . . . . . . . . . . . . . . . . . 30
1.13 X-ray AGN classification: absorbed and unabsorbed sources. . . . . . . . . . . . 30
1.14 The AGN Unified Model . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 32
2.1 Light Polarization . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 36
2.2 Polarization by scattering . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 37
2.3 Polarization of Hα line of Akn 120 . . . . . . . . . . . . . . . . . . . . . . . . . 41
2.4 The scattering regions responsible for polarization in Seyfert galaxies . . . . . . 42
2.5 Cold absorption upon a power law continuum . . . . . . . . . . . . . . . . . . . 45
2.6 Cold absorbers in NGC 1365 and ESO 323-G077 . . . . . . . . . . . . . . . . . . 45
2.7 Modeling the warm absorber . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 46
2.8 High-resolution spectrum of NGC 3783 . . . . . . . . . . . . . . . . . . . . . . . 47
2.9 Ionized regions in AGN . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 48
2.10 An X-ray test to the unification scheme of polarized Seyferts . . . . . . . . . . . 49
2.11 The geometry of ESO 323-G077, a polar-polarized source . . . . . . . . . . . . . 49
3.1 XMM Newton Observatory . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 54
3.2 EPIC-pn focal plane . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 54
3.3 Out of time events . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 55
3.4 Epatplot for Pile-up . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 56
3.5 EPIC-pn Background of Was 45 . . . . . . . . . . . . . . . . . . . . . . . . . . . 57
3.6 Source and background regions for spectrum extraction . . . . . . . . . . . . . . 58
3.7 Extracted spectrum of Was 45 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59
3.8 Soft excess theoretical models . . . . . . . . . . . . . . . . . . . . . . . . . . . . 61
3.9 The relxill model for ionized reflection . . . . . . . . . . . . . . . . . . . . . . . 62
3.10 Models for cold absorption . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 63
3.11 Models for warm absorption . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 64
3.12 Fitting process flow chart . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 67
3.13 Power law index for different soft excess models analyzed with Sherpa . . . . . . 69
3.14 Absorption test done with Sherpa . . . . . . . . . . . . . . . . . . . . . . . . . . 70
6
LIST OF FIGURES 7
3.15 Four models of soft excess tested with PyXspec . . . . . . . . . . . . . . . . . . 71
3.16 Example of PyXspec code . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 71
3.17 Was 45 fitted with relxill . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 73
4.1 Spectral counts . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 81
4.2 X-ray Luminosity . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 81
4.3 Incidence of absorption . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 84
4.4 Column density distribution . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 86
4.5 Distribution of column density and corresponding ionization parameter of the
Warm Absorbers . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 86
4.6 Equivalent width vs Luminosity . . . . . . . . . . . . . . . . . . . . . . . . . . . 88
4.7 The column density values: our results vs reported in literature . . . . . . . . . 89
4.8 The ionization parameter: our results vs reported in literature . . . . . . . . . . 90
A.1 Polar polarized sources - spectral fits . . . . . . . . . . . . . . . . . . . . . . . . 108
A.2 Equatorial Polarized sources - spectral fits . . . . . . . . . . . . . . . . . . . . . 110
A.3 Weakly Polarized sources - spectral fits . . . . . . . . . . . . . . . . . . . . . . . 111
B.1 Polar polarized sources - images . . . . . . . . . . . . . . . . . . . . . . . . . . . 113
B.2 Equatorial polarized sources - images . . . . . . . . . . . . . . . . . . . . . . . . 114
B.3 Weakly polarized sources - images . . . . . . . . . . . . . . . . . . . . . . . . . . 115
List of Tables
3.1 Sample: X-ray data of polarized Seyfert 1 galaxies . . . . . . . . . . . . . . . . . 52
3.2 Soft excess testing for Mrk 704 with Sherpa . . . . . . . . . . . . . . . . . . . . 68
3.3 Absorption test for Mrk 704 with Sherpa . . . . . . . . . . . . . . . . . . . . . . 68
3.4 Relxill fits . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 74
4.1 Results from the hard band fits . . . . . . . . . . . . . . . . . . . . . . . . . . . 76
4.2 Results from the full energy band . . . . . . . . . . . . . . . . . . . . . . . . . . 77
4.3 Results from the Absorption test . . . . . . . . . . . . . . . . . . . . . . . . . . 78
4.4 Incidence of absorption . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83
8
Abstract
This project focuses on the study of the X-ray emission of Active Galactic Nuclei (AGN) and its
possible relation to optical polarization properties. A galaxy hosting an active nucleus exhibits
a highly luminous central region and continuous emission across the entire electromagnetic
spectrum. Examples of these properties are the broad and narrow optical emission lines asso-
ciated to a prominent ionizing continuum and a strong X-ray emission that characterizes all
AGN. These and other characteristics point to a great variety of physical processes that take
place in various regions of AGN, such as the central core, Broad Line Region (BLR), Narrow
Line Region (NLR), etc. The presence or absence of some of these properties in some AGN
spectra form the basis for various classification criteria and are often attributed to orienta-
tion effects. For example, a region of obscuring material is sometimes responsible for blocking
specific emitting regions from the line of sight, thus explaining some of the observed differences.
AGN are classified into two main types based on their optical emission lines: type 1 AGN
exhibit both broad and narrow lines, while type 2 AGN show only narrow lines. The detection
of polarized broad emission lines in type 2 galaxies, specifically NGC 1068, revealed that the
broad-line region (BLR) is present but obscured from direct view (Antonucci & Miller 1985).
This discovery led to the Unified Model (Antonucci 1993; Urry & Padovani 1995), which posits
that an optically thick torus surrounds the BLR and central region, obscuring the AGN’s direct
emission depending on its orientation. In this model, the observed spectral features, including
the presence or absence of broad lines, are determined by the orientation of the torus relative
to our line of sight.
Optical polarization of AGN may also offer insights into the geometry of the circumnuclear
region, as it reveals light that is otherwise hidden from detection. We focus our work on Seyfert
(Sy) galaxies with known optical polarization and consider the orientation effects related to it.
In general, the polarization position angle (PA) of type 2 Sy is found to be perpendicular to
the principal axis of the system, i.e., the rotation axis of the accretion disk. This is the polar-
polarization region, PL-pol, and corresponds to the well-established AGN ionization cones that
trace the kpc-scale region from where the narrow optical lines typical of AGN arise. A second
scattering region is postulated to account for the measurements on polarized type 1 Sy, where
the polarization PA is parallel to the principal axis of the system. This scattering region is
described as co-planar to the accretion disk, in the equatorial plane of the torus, and is referred
to as the equatorial-polarization region, EQ-pol. Interestingly, Smith et al. report a list of type
1 sources with polarization properties similar to those of type 2, they are polar-polarized. In
accordance to the Unified Model, Smith et al. (2002, 2004, 2005) proposed that both scatter-
ing regions are present in all Seyfert sources and that the observed optical polarization is also
orientation-dependent. This unified scheme relates these observations with the orientation of
the torus: a source with no obscuration effect by the torus shows the spectrum of a type 1 Sy
with EQ-pol; a source where the torus is obscuring the central region corresponds to a type 2
Sy with PL-pol; and, lastly, a type 1 Sy with PL-pol represents an intermediate type between
a type 2 Sy with PL-pol and a type 1 Sy with EQ-pol.
9
10 Abstract
Moving to the X-ray band, we can apply a different criterion to differentiate between AGN
types. X-rays trace material along the line of sight through absorption, corresponding to a
dichotomy between obscured and unobscured sources, consistent with this orientation-based
framework. The presence of circumnuclear gas associated with X-ray absorption and scattering
is a well established feature of AGN.
This thesis is addressed to explore if a physical association is viable between the X-ray
absorption and the material responsible for the optical polarization. The X-ray regime is espe-
cially appropriate for assessing the role of the torus absorption effect, thus providing a useful
tool to test the presence or absence of the absorbing material co-spatial with the outer layers
of the torus on the line of sight. The study of X-ray absorption can provide an independent
test to the polarization unification scheme proposed by Smith et al. by comparing the X-ray
absorption properties, if present, to the known optical polarization characteristics. As a first
approach, we would expect a type 1 PL-pol source to show the imprints of the outer layers
of the torus in the form of absorption by mildly ionized gas, which is often observed in type
1 Sy galaxies. On the other hand, we would expect a type 1 EQ-pol source to show less ab-
sorption or to be unabsorbed. In order to test this hypothesis, we carried out a systematic
analysis of the X-ray spectra of 25 type 1 Sy sources with optical polarization reported by
Smith et al. and classified into three sub-samples: 11 type 1 Sy with PL-pol, 8 with EQ-pol
and 6 with weak polarization but described as “candidates for null-polarization” by Smith et al.
Our analysis consists of building a nested model, starting with a simple base and progres-
sively incorporating additional elements. The choice of each model component is based on the
known characteristic features of the typical X-ray spectra of a Seyfert source. As we include
components in the model, we can statistically test whether the newly added element produces
an improved fit. The resulting model is then used to test the presence of absorption.
The absorption test consists of separately fitting a neutral absorption and an ionized ab-
sorption component and determining which best approximates the data. We aim to determine
whether the source is unabsorbed, affected by an ionized absorber (warm absorption), or af-
fected by a neutral absorber (cold absorption). The results of this test are then compared
to their known polarization, with the ultimate aim to establish a relation between the X-ray
absorber and the polarization scattering regions. We find that 100% of the PL-pol subsample
show some form of absorption in comparison with 75% of the EQ-pol, posing an interesting dis-
cussion regarding the presence of absorbing material and its relationship with the polarization
regions. Moreover, our premise is supported by the finding that only 33% of the sources with
weak or null polarization exhibit presence of absorption.
Abstract 11
The structure of this document is as follows:
• Chapter 1: Active Galactic Nuclei. This chapter introduces Active Galactic Nuclei,
describing their structure, the main features of their emission, as well as some of the
classification criteria and the Unified Model.
• Chapter 2: Optical polarization in Seyfert Galaxies. This chapter reports on
studies of polarized Seyfert galaxies, primarily by Smith et al., describing their unification
scheme. We then present our proposed test of this unification scheme through the study
of X-ray absorption, establishing our hypotheses and objectives.
• Chapter 3: X-ray analysis of optically polarized Seyfert galaxies. This chapter
presents and describes our X-ray sample and analysis. We begin by introducing our
sample and providing a description of the XMM-Newton X-ray observatory and the data
processing. This is followed by the X-ray analysis, starting by detailing the models
chosen for the fitting process. Next, we describe our three-step methodology and the
corresponding statistical tests. Finally, we outline the procedures we employed using
two distinct fitting approaches: a command-line-based method with Sherpa and a script
developed for PyXspec.
• Chapter 4: Results and discussion. We present, organize, and describe the results
of our analysis, followed by a discussion of various aspects of our results.
• Chapter 5: Conclusions and future work. The final chapter is dedicated to our
conclusions and suggestions for future work.
Chapter 1
Active Galactic Nuclei
The light of a galaxy comprises the integrated light of its stars, with its spectrum characterized
by absorption lines from stellar atmospheres and emission lines from regions of hot gas. Some
galaxies exhibit highly energetic nuclear emission across the entire electromagnetic spectrum.
These are active galaxies, and the term Active Galactic Nuclei (AGN) refers to the highly
energetic and compact region found in the center of active galaxies. While various distinctions
can be identified based on their observational properties, AGN are typically characterized by
the following key attributes, as outlined in Netzer (2013):
• An AGN is the compact central region, with emission that is “beyond what is expected
from stellar processes” in its host galaxy and bolometric luminosities that can go up to
1048 erg s−1, (Ginzburg & Ozernoy 1977).
• The central region of the host galaxy emits a unique non-stellar continuum spanning the
entire electromagnetic spectrum.
• The spectrum shows strong optical emission lines from excitation by a non-stellar radia-
tion field.
• The emission lines and/or the continuum exhibit variability on timescales ⩽ 1 day,
(Ginzburg & Ozernoy 1977).
AGN are the most powerful sources of the Universe and play an important role in galaxy
evolution, with energy outputs that may affect the star formation activity and evolution of their
host galaxy, (Heckman & Best 2014; Hopkins & Elvis 2010). The study of AGN is of great
relevance to Extragalactic Astrophysics and their study has significantly improved throughout
decades of observations, analysis, and technological advances. In this chapter, we relate a brief
observational history (Shields 1999), followed by a description of the main structures that form
an AGN, all derived from the analysis and interpretation of observational data. We end this
chapter by introducing the Unified Model, developed in the 1990s to explain the AGN config-
uration and the observed differences as an orientation effect.
1.1 Observational History
The origins of AGN research can be traced back to early studies of what was then known as
“spiral nebulae”, so named for their distinctive spiral structures. In the 1920s, astronomers de-
bated whether these objects were independent extragalactic systems or part of the Milky Way.
12
1.1. OBSERVATIONAL HISTORY 13
However, the foundation for their identification and initial spectral studies was laid years earlier.
In 1908, Edward Fath set out to study these objects, conducting a series of observations at
Lick Observatory. He reported that these nebulae consisted of a collection of solar-type stars,
with the typical dark (absorption) lines from stellar atmospheres. Notably, Fath found the pe-
culiar spectrum in NGC 1068 (M77) and determined that it was a composite spectrum of both
bright (emission) and dark (absorption) lines, (Fath 1909). These findings were later confirmed
by Slipher in 1917, (Slipher 1917). In 1924-25, Edwin Hubble published the calculations for the
distances of M33 and M31 using the period-luminosity relation of Cepheids stars, confirming
that both sources are, in fact, extragalactic (Hubble 1926).
In the following decades, many milestones contributed to the understanding of these spiral
nebulae, now called galaxies. In 1943, Carl Seyfert conducted a systematic examination of six
galaxies: NGC 1068, NGC 1275, NGC 3516, NGC 4051, NGC 4151, and NGC 7469. Seyfert
reported the nuclear broad emission lines present in the spectra and identified them as a par-
ticular type of galaxy, (Seyfert 1943). Consequently, galaxies exhibiting broad emission lines
became known as “Seyfert Galaxies”.
The detection of extraterrestrial radio signals also played a very important role. In the late
40s and early 50s, the partnership of optical and radio astronomy set out to detect discrete
radio sources, determine their location, and identify their optical counterparts. One example
is the work on 3C 295, the 295th source from the 3rd Cambridge Catalogue of Radio Sources
(3C), published in 1959. Minkowski (1960) was able to identify the optical counterpart of 3C
295 and determined the redshift of the source at z = 0.46, making it one of the furthest sources
known at the time. Many publications on the identification of sources and their redshift mea-
surements followed.
A particular breakthrough was the work of Schmidt on 3C 273, a source whose accurate
position had been obtained a year earlier by Hazard et al. (1963). Schmidt published the
spectral analysis revealing emission lines at unfamiliar wavelengths (Schmidt 1963). The work
by Schmidt and others concluded that these “quasi-stellar objects” were highly luminous and
compact sources at great distances. Credit for the term “quasar” is attributed to Hong-Yee
Chiu, who coined it in a publication in Physics Today in May 1964.
The 1960s and the following decades brought significant developments in X-ray astronomy.
In 1962, the American Science and Engineering (AS&E) launched the first rocket with an X-ray
detector on board. Using this detector, Giacconi et al. (1964) reported a bright X-ray cosmic
background, peaking in the Scorpius constellation. In December of 1970, NASA launched
Uhuru, the first satellite commissioned for X-ray astronomy that made the first X-ray detec-
tions from Seyfert galaxies: NGC 1275 and NGC 4151, (Gursky et al. 1971). The following
decades saw many X-ray missions such as Ariel V (1974) or the ıOrbiting Solar Observatories
(OSO) consisting of eight different satellites launched from 1962 to 1975. Data collected from
these missions unveiled one of the most remarkable features of AGN: their variability. The
detected fluctuations in the flux observed in objects such as Cen A (NGC 5128) over short
periods, as brief as six days in the case of Cen A, led to the conclusion that the emission arises
from a compact nuclear region, 1015−1016 cm, (Fabian et al. 1976). This finding highlighted the
deep connection between X-rays and the nature of the central engine driving AGN phenomena.
The availability of data, extending from radio to γ − rays, has significantly improved since
the 1980s and onward. In particular, in X-ray Astronomy, the development of space observato-
14 CHAPTER 1. ACTIVE GALACTIC NUCLEI
ries has made it possible to study the X-ray emission of AGN. The ROSAT satellite, developed
by the Deutsches Zentrum fur Luft und Raumfahrt (DLR), operational from 1990 to 1999,
provided the first X-ray all-sky survey, (Trümper 1982; Voges 1993). The Advanced Satellite
for Cosmology and Astrophysics (ASCA) was also of great significance in the area of X-ray
astronomy. This mission by the Japan Aerospace Exploration Agency (JAXA), had a life span
from 1993 to 2001. The sensitivity of the instruments of ASCA allowed the first detailed spec-
trum of AGN, (Reynolds 1997; Reeves & Turner 2000). These missions also paved the way for
successors like XMM-Newton by the European Space Agency (ESA) and the Chandra X-ray
Observatory by NASA, two ongoing missions that have provided high resolution imaging and
spectroscopy since 1999. Another prominent X-ray observatory was Suzaku, a JAXA mission
active from 2005 until 2015 that featured both X-ray imaging and spectroscopy capabilities
which allowed for sensitive observations across a broad energy range, from soft to hard X-rays
(Mitsuda et al. 2007; Serlemitsos et al. 2007; Kelley et al. 2007; Takahashi et al. 2007). The
mission provided valuable insights into the nature of AGN, confirming the presence of relativis-
tically broadened iron lines (e.g. Reeves et al. 2006; Miller 2007) and probing the structure of
the inner regions of accretion disks around supermassive black holes (SMBH) (e.g. Reynolds
2015). Its high sensitivity to faint sources also helped in studying heavily obscured AGN, which
are otherwise challenging to detect in other wavebands.
Since the pioneers Chandra and XMM-Newton, new missions have continued to advance the
field of X-ray astronomy. The Swift X-Ray Telescope launched in 2004 and has since provided an
extensive compilation of high-energy sources, (Burrows et al. 2005). The Burst Alert Telescope
catalog (BAT), which is updated periodically, includes data on gamma-ray bursts (GRBs),
AGN, galaxy clusters, etc (e.g. Koss et al. 2017; Ricci et al. 2017; Shimizu et al. 2018). The
Nuclear Spectroscopic Telescope Array (NuSTAR) (Harrison et al. 2013), launched by NASA in
2012, was the first satellite capable of focusing high-energy X-rays (above 10 keV), making it
possible to observe the hard X-ray spectrum of AGN with unprecedented detail (e.g Brandt &
Alexander 2015; Fabian et al. 2015). Another important development is the eROSITA instru-
ment aboard the Spectrum-Roentgen-Gamma (SRG) mission, launched in 2019 (Predehl et al.
2021). Operated by the Max Planck Institute for Extraterrestrial Physics, eROSITA is con-
ducting an all-sky survey with an X-ray sensitivity far exceeding that of ROSAT significantly
increasing the number of known X-ray sources, including AGN (Brunner et al. 2022; Merloni
et al. 2024). These missions have not only unveiled new aspects of AGN but also provided
critical insights into the physics driving their extreme behaviors.
Future missions are also set to push the boundaries of X-ray astronomy. The X-Ray Imaging
and Spectroscopy Mission (XRISM), a joint project by JAXA and NASA, launched in 2023.
The Advanced Telescope for High-Energy Astrophysics (Athena), a highly anticipated mission
developed by ESA is set to launch in the 2030s and it will offer a combination of high spatial
resolution and spectral sensitivity, enabling detailed studies of AGN evolution and the role of
SMBH in galaxy formation.
Over a century of research has shaped our understanding of AGN, starting from their iden-
tification to detailed studies on the physical processes driving their behavior. Today, advanced
instruments allow us to probe AGN across all electromagnetic wavelengths. Moreover, future
missions promise to expand our knowledge even further. By observing AGN across different
wavelengths and redshifts, astronomers can trace the history of black hole activity and their
role in shaping galaxy evolution, (Harrison 2017). Powered by accretion onto a SMBH, AGN
emit vast amounts of energy that can regulate star formation (Page et al. 2012) and drive
feedback processes (Schawinski et al. 2007), influencing the properties of their host galaxies.
1.2. THE MAIN PHYSICAL STRUCTURE OF AGN 15
These energetic nuclei also serve as natural laboratories for exploring extreme physics (Rani
et al. 2019), such as strong gravitational fields and high-energy radiation.
1.2 The main physical structure of AGN
Active Galactic Nuclei (AGN) are defined as such by their shared characteristics that suggest
a common underlying physical structure. The description of AGN structure derives from the
observational properties detected throughout several decades since their discovery. The main
components are the central SMBH surrounded by an emitting accretion disk. A distribution of
hot electrons is found immediately above the accretion disk, the X-ray corona, followed by a gas
region, extending from 0.1-1 pc, that gets ionized and excited by the accretion disk, producing
broad emission lines and is thus called the Broad Line Region (BLR). A toroidal circumnuclear
distribution of neutral gas and dust, located on pc-scale, that can sometimes block our line of
sight to the innermost region (see Sect. 1.4). Lastly, there is a further extended Narrow Line
Region (NLR), extending on kpc-scale. This structure, located at the center of a spiral galaxy,
is depicted in Fig. 1.1.
Figure 1.1: The physical structure of an AGN. The central region consists of a SMBH sur-
rounded by an emitting accretion disk at scales of ∼ 0.01 pc, shown in blue. This central
region is surrounded by the BLR, ∼ 0.1 pc. In the BLR block, we indicate a small region
located immediately above the accretion disk, the X-ray Corona, denoted with a small red star.
Further away from the BLR, we find the dusty torus, 1 − 100pc. Lastly, we have the NLR, a
diffuse gas region extending up to 103 pc. All of these are located in the very center of a spiral
galaxy. Image adapted from Malygin et al. (2019).
1.2.1 The central region
The source of primary emission consists in a system formed of a SMBH surrounded by an
emitting accretion disk (Salpeter 1964; Lynden-Bell 1969), a configuration supported by AGN
high luminosities and their short-term variability.
16 CHAPTER 1. ACTIVE GALACTIC NUCLEI
A well established phenomenon in AGN is their optical and UV variability, characterized
by fluctuations in brightness over timescales as short as ≲ 60 min. This short-term variabil-
ity indicates that the radiation arises from a compact region, whose size is estimated to be
∼ 0.02 ly ≈ 1016 cm (Smith & Hoffleit 1963; Fabian et al. 1976; Ginzburg & Ozernoy 1977).
This consideration led to the postulate that a compact object of high mass is located in the
nuclear region of AGN. The presence of a central SMBH is consistent with this scenario and in
turn, provides a mechanism for the high luminosities due to its large gravitational potential.
We begin by defining the Schwarzschild radius, rS, which is determined by the mass of the
SMBH. This radius corresponds to the distance at which the escape velocity equals the speed
of light, marking a point known as the event horizon, a boundary beyond which matter is in-
evitably accreted into the SMBH. Eq. 1.1 defines the Schwarzchild radius given by the SMBH
mass, MBH , where G is the gravitational constant1, and c is the speed of light2.
rS =
2GMBH
c2
(1.1)
As matter is pulled inwards by the gravitational force of the SMBH, it gains kinetic energy
and loses angular momentum, forming a rotating accretion disk. Gravitational potential en-
ergy gets converted into radiation, serving as the source of the primary emission, i.e., the main
engine. Various models have been proposed for the accretion mechanism and emission process.
In particular, Shakura & Sunyaev (1973) describes a geometrically thin and optically thick disk
where viscosity due to turbulence transports angular momentum outwards, heating the disk.
The resulting radiation is in the form of multiple black bodies of different temperatures, T .
The luminosity of the radiated energy is given by Eq. 1.2, where the left hand side corresponds
to the release of gravitational energy, with Ṁ corresponding to the rate of accreted mass. The
right hand side represents the black body radiation3, σT 4, per unit area, πr2.
L =
GMBHṀ
2r
= 2πr2σT 4 (1.2)
By balancing the radiation pressure with the gravitational force of the SMBH (Frad, FG, on
the left and right-hand sides of Eq. 1.3 respectively), we obtain the maximum luminosity that
can be emitted by an object before Frad from the emitted photons overcomes the gravitation
pull, limiting the accretion power. This is the Eddington Luminosity, LEdd, dependent on the
MBH . Eq. 1.4 is the Eddington Luminosity of a compact object of mass MBH ; where σT
corresponds to the Thomson cross section4 and mp is the mass of the proton5.
LσT
4πr2c
=
GMBHmp
r2
(1.3)
LEdd =
4πGMBHmpc
σT
(1.4)
1G= 6.67× 10−11 Nm2kg−2
2c = 299, 792, 458 m s−1
3σ = 5.67× 10−18 W m−2K−4 is the Stefan-Boltzmann constant
4σT = 6.65×10−29 m2: Thomson cross section is the probability that a photon interacts with a free electron.
5mp = 1.67× 10−27 kg
1.2. THE MAIN PHYSICAL STRUCTURE OF AGN 17
Evidently, there is a tight relation between the observed luminosity of the AGN and the
SMBH mass. For example, Shields (1978) provides an estimation where for a SMBH mass of
108M⊙, the emission from the accretion disk would reach a luminosity of L ∼ 1046 erg s−1. The
temperature corresponding to these parameters is of the order to 105 K, or ∼0.01 keV, an actual
upper end to the disk temperature and corresponding to the UV region of the electromagnetic
spectrum (see Sect. 1.3).
Moreover, the prominent X-ray continuum in AGN spectra, detailed in Sect. 1.3.3, also
presents short-term variability, similar to the variability of the Optical/UV emission. This sug-
gests that an X-ray emitting region comes from a compact source, such as an accreting SMBH.
However, X-rays are not consistent with the thermal emission from the accretion disk, thus a
hot medium is proposed to be located in the immediate proximity of the accretion disk, this
is called the X-ray corona and is described as a plasma of relativistic electrons. The nature
and structure of the X-ray corona are still a subject of research, bearing a strong dependence
on the data quality across the X-ray band. Possible explanations for the formation of this
plasma include the strong magnetic fields of the accretion disk and intense radiation pressure
from the inner accretion disk, (e.g Ursini et al. 2020; Hinkle & Mushotzky 2021; Cackett et al.
2021; Gronkiewicz et al. 2023). Sect. 1.3.3 describes the emission resulting from the interplay
between the emission from the X-ray corona and the surrounding material.
1.2.2 Line emitting regions
Optical emission lines are a key observational characteristic of AGN and are a crucial feature in
distinguishing them from normal galaxies. These lines originate in regions ionized by radiation
from the central engine, with photoionization from bound-free transitions being the dominant
emission process. There are two main types of AGN emission lines, distinguished mainly by
their width, and emitted in two physically different gas regions, the Broad and Narrow Line
Regions, BLR and NLR respectively.
1.2.2.1 The Broad Line Region
The broad emission lines, observed in the optical and UV spectra of AGN, have widths in
ranges between 102 − 104 km s−1. Their width is a consequence of the gas motion, positioning
the BLR in the vicinity of the accretion disk, ∼ 0.1 − 1 pc. The width of these lines can be
used to estimate the mass of the black hole with the relation in Eq.1.5; where for a distance of
r = 0.1 pc and a velocity broadening of 104 km s−1, v, the resulting SMBH mass is of 108M⊙,
(Netzer 2013; Peterson 1997).
v2
r
=
GMBH
r2
(1.5)
Moreover, broad lines exhibit variability, which further corroborates the proximity to the
central region. Fig. 1.2 shows the optical spectrum of NGC 4151, (Shapovalova et al. 2008), de-
picting two different epochs, 1996 and 2005, with a maximum and minimum of flux respectively.
The flux variation, by a factor of 6, is evident in the broad lines such as Hα λ6563 and Hβ λ4861.
Broad lines also show a myriad of excitation and ionization states, and only from permit-
ted atomic transitions. The absence of broad lines from forbidden transitions indicates that
the gas emitting region has high density, as these forbidden transitions would get collisionally
18 CHAPTER 1. ACTIVE GALACTIC NUCLEI
Figure 1.2: The optical spectrum of NGC 4151 in two different epochs. The 1996 observation
shows a higher flux in comparison with the measurement from 2005. This is an example of
the variability of the broad emission lines of AGN. Lines corresponding to emission lines Hα at
λ6563 and Hβ at λ4863, with blue markers, correspond to broad emission lines, whilst the O
III λλ4959 5007 are narrow emission lines, marked in red. Credit: Shapovalova et al. (2008).
suppressed. The estimation of the gas density ranges between 108 − 1011 cm−3 (Netzer 1990).
1.2.2.2 The Narrow Line Region
In Fig. 1.2, we see the narrow line that corresponding to [O III] λ5007. This line arises from
a forbidden transition and does not show the flux variability featured by the broad lines in
the same object. These differences tell us about a second line emitting gas region of differ-
ent characteristics than the BLR, the NLR. Typical widths of these lines are in a range of
∼ 300− 1000 km s−1.
The constant flux in the narrow lines indicates that the NLR is further away from the central
engine. The estimated spatial scale of the NLR is of ∼ 100 to 1000 pc, (Netzer 2013; Beck-
mann & Shrader 2012). Additionally, the presence of lines from forbidden and semi-forbidden
transitions, e.g. [O III] λ5007, implies that this gas region is less dense than the BLR. The
density estimation for the NLR is of 103 − 106 cm−3.
Due to this larger extension, it is possible to obtain images of the NLR, a well-established
double cone structure depicted in Fig. 1.3. These are bi-conical regions of cylindrical symmetry,
extending co-axial to the rotation of the accretion disk, (Osterbrock 1993; Wilson & Tsvetanov
1994). We will refer to this as the principal axis of the system: the rotation axis of the accretion
disk aligned with the ionization cones.
1.2. THE MAIN PHYSICAL STRUCTURE OF AGN 19
Figure 1.3: The ionization cones of NGC 5728. Observations were conducted using the Hubble
Space Telescope. Left image depicts the source filtered and continuum subtracted for Hα +
[NII] lines at λλ 6548, 6583. The right image corresponds to the contours of the same filtered
observation at various percentages of peak values. Credit: Wilson et al. (1993).
1.2.3 The molecular torus
The IR emission, along with the obscuration of optical and UV radiation, provides evidence for
the presence of neutral and optically thick material surrounding the AGN. This distribution of
material must be geometrically thick to effectively obscure the central region (Krolik & Begel-
man 1988). Specifically, this region is characterized as a circumnuclear concentration of neutral
gas and dust, forming a toroidal structure situated beyond the BLR, with an inner radius of
∼1 parsec and an outer radius extending up to 100 parsecs. Raban et al. (2009) were able to
resolve the nuclear region of the Seyfert 2 galaxy NGC 1068 using MIDI, the mid-infrared inter-
ferometer at the Very Large Telescope (VLT) of the European Southern Observatories (ESO),
in Paranal, Chile. They were able to identify two different mid-infrared regions, the first one
1.35 pc long and 0.45 pc thick in full width at half-maximum (FWHM) with a temperature of
800 K and a second component of 3 × 4 pc in FWHM and a temperature of 300K, besides evi-
dence that the gas distribution is not homogeneous but clumpy. The Atacama Large Millimeter
Array (ALMA) has provided data with which the torus of NGC 1068 has been successfully re-
solved, mapping the emission of the different molecules and imaging the distribution of this
gas, tracing its dynamics and estimating parameters such as scale and mass (Garćıa-Burillo
et al. 2016; Imanishi et al. 2018; Garćıa-Burillo et al. 2019). Fig. 1.4 shows two images of the
torus resolved from the ALMA observatory data.
One possible formation mechanism for this molecular torus involves outflows and radiation
feedback from the accretion disk. In this scenario, material is expelled from the innermost re-
gions of the disk, where intense gravitational and thermal forces prevail. This expelled material
interacts with infalling gas, leading to cooling and accumulation in the so-called ”molecular
torus” (Wada 2012). Further ongoing research into the geometry, physical properties, and
origins of this gas distribution, (e.g López-Gonzaga et al. 2016; Hönig 2019) is crucial for un-
derstanding its relationship to different types of AGN (see Sect. 1.4).
20 CHAPTER 1. ACTIVE GALACTIC NUCLEI
Figure 1.4: Resolved image of the torus of NGC 1068 through observations from the ALMA
observatory. The blue contours on the left panel correspond to the HCO molecule superimposed
on the CO scale. The right panel shows a close-up view of the inner part of the torus, with
radius ≈ 8 pc, from higher-resolution data. The blue line marks the direction of a large-scale
ionized outflow of the galaxy, and the red line shows the main rotation axis of the torus. Yellow
ellipses represent the beam sizes of ALMA observations. Credit: Garćıa-Burillo et al. (2019).
1.2.4 The radio Jet
Finally, some AGN show strong radio emission corresponding to synchrotron radiation emitted
from relativistic particles in a strong magnetic field. These structures can extend up to scales
of ∼ 103 kpc and are highly collimated streams of relativistic plasma ejected from the central
region, (Rees 1978).
The jets are powered by the accretion onto the SMBH. As matter spirals inward, the grav-
itational energy released during accretion heats the surrounding material and generates strong
magnetic fields. These magnetic fields extract rotational energy from either the black hole or
the accretion disk itself, accelerating charged particles to relativistic speeds and ejecting them
into forming these radio emitting streams (Blandford et al. 2019; Hada 2019; Marscher et al.
2008, etc).
An example corresponding to the radio structures of M87 (NGC 4486) is shown in Fig. 1.5.
The top left image shows a composite of optical and infrared, covering a scale of > 5000 ly ∼
1500 pc and showing the luminous AGN and the emerging jet. The top right panel corresponds
to a scale of 0.5 ly ∼ 0.15 pc, depicting the emission at 43 GHz, obtained by the East-Asian
VLBI Network collaboration (EAVN) (An et al. 2018). This image shows the origin of the
collimated jet at the central region of the AGN. The bottom panel shows the shadow of the
central SMBH in M87, at a scale of 0.01 ly, obtained by the Event Horizon Telescope (EHT)
(EHT Collaboration et al. 2019). The image shows the radio emission of a bright ring that
corresponds to radio photons lensed by the black hole. From these images, the EHT collabo-
ration was able to determine a BH mass of ∼ 6.5 × 109 M⊙ and peak brightness temperature
of T ∼ 109 K, consistent with synchrotron emission, (Hada 2019). This is the very first direct
image of a SMBH located at the center of an AGN, an image produced by combining the data
from 8 radio telescopes (at the time of the M87 publication) and provided the first direct ob-
1.3. AGN SPECTRAL ENERGY DISTRIBUTION 21
servational evidence of the central region of an AGN.
Figure 1.5: Top left is an optical/IR HST image of M87, credit: HST STScI/AURA. The top
right image is a 43 GHz obtained by the EAVN collaboration. The bottom image corresponds
to the first direct image of a SMBH, located at the center of the M87 galaxy, showing the
shadow of the black hole (EHT Collaboration et al. 2019). Image credit: Hada (2019).
1.3 AGN Spectral Energy Distribution: emission through
the electromagnetic spectrum
The Spectral Energy Distribution (SED) of an AGN is a crucial tool for inferring its structure,
as the detected emission across the electromagnetic spectrum originates from distinct regions
within the AGN. The multi-wavelength components, from radio to gamma-rays, provide in-
formation on the various physical processes such as synchrotron or thermal radiation and,
in consequence, provide clues about the spatial distribution and condition of various regions.
Thus, the SED serves as a map, connecting the observed spectrum to the physical structure of
the AGN.
Customarily, the SED is referred to in terms of monochromatic luminosity per unit fre-
quency, Lν erg s−1 Hz−1. Fig. 1.6, depicts the main emission components that constitute the
SED of AGN, (Harrison 2014). This section is dedicated to briefly summarizing the observed
features of the AGN SED from lowest to highest photon energy, from radio waves to gamma
rays.
22 CHAPTER 1. ACTIVE GALACTIC NUCLEI
Figure 1.6: The SED of Active Galactic Nuclei. The colored lines correspond to the distinct
components contributing to the integrated SED, depicted by the solid black curve. The light
grey line corresponds to the spectrum of a non-active, star forming galaxy. Credit: Harrison
(2014).
1.3. AGN SPECTRAL ENERGY DISTRIBUTION 23
1.3.1 Radio and Infrared
On the lower energy end of the electromagnetic spectrum, we find the radio band. The radio
emission of AGN primarily arises from synchrotron radiation, produced by relativistic electrons
spiraling around magnetic fields. As seen by the two orange lines in Fig. 1.6, there can be a
significantly different radio emission in AGN. The line of higher flux, log νFν = −2, corresponds
to the emission associated with large-scale jets of relativistic matter (Blandford et al. 2019).
These sources are known as “radio-loud” AGN and represent around 10% of the known AGN
population. When present, jets can be indicators of the AGN orientation. The lower flux group,
log νFν = −6, are known as “radio-quiet” sources. Their radio emission is associated with the
compact core, linked to the immediate vicinity of the SMBH.
The next section of the electromagnetic spectrum is the infrared emission (IR), divided into
far (FIR), mid (MIR), and near-infrared (NIR). The measured IR continuum is depicted by a
red dotted line in Fig. 1.6. This emission is thermal radiation from dust particles primarily
attributed to the circumnuclear region surrounding the central SMBH. This dust absorbs high-
energy ultraviolet and optical radiation emitted by the accretion disk and re-emits it in the
infrared range. The IR continuum peaks at around ∼ 20− 50 µm, in the MIR range (Sanders
1999). The grey line indicates contributing emission from star-formation activity in the host
galaxy, being the dominant IR source at NIR ranges. Approaching the FIR and sub-mm wave-
lengths, the processes of emission can be once again dominated by phenomena intrinsic to the
AGN.
1.3.2 Optical and UV
The main component in the optical and UV bands is a continuum emission called the Big
Blue Bump (BBB), for its characteristical shape, depicted by a dotted blue line in Fig. 1.6.
This feature is a product of the thermal emission of the accretion disk, peaking in the UV, at
∼ 1100−1300Å corresponding to a temperature on the order of ∼ 104 K, (Elvis et al. 1994). A
good portion of this feature falls in the Far-UV, from ∼ 2000 Å, and gets significantly affected
by Galactic absorption. The cut-off energy of the BBB is below 0.6 keV ≈ 20 Å, approaching
the soft X-ray range.
A significant portion of the optical and UV emission originates from the accretion disk, and
it is accompanied by prominent broad and narrow emission lines, generated by ionized gas in
the BLR and NLR, respectively.
1.3.3 X-ray spectrum
The focus of this project, and a prominent characteristic of the SED of AGN, is their X-ray
emission. Therefore, we describe its properties in greater detail. Figure 1.7 shows a typical
X-ray spectrum of AGN, represented by the black solid line. The dotted red, blue, and green
lines correspond to different components of the spectrum.
X-ray primary emission originates from the corona, located in the proximity of the accretion
disk. Haardt & Maraschi (1993), proposed a two-phase model in which the hot electron corona
is located in an optically thin layer, where photons from a colder, underlying, optically-thick
accretion disk undergo inverse Compton mechanism. Thus, the primary mechanism driving
X-ray emission is Compton up-scattering of thermal UV/optical photons from the accretion
24 CHAPTER 1. ACTIVE GALACTIC NUCLEI
Figure 1.7: The X-ray spectrum of an AGN. The solid black line depicts the integrated spec-
trum, and dotted colored lines correspond to different components. The primary emission is a
power law continuum (red), with an energy cut-off at ∼ 100 keV. There is a reflection com-
ponent (blue) comprised of the reflection hump that peaks between 20-40 keV and the Fe Kα
emission line at 6.4 keV. At lower energies, < 2 keV, we see emission that exceeds that of the
main continuum, the soft excess. Credit: Ricci (2011).
disk, producing a the primary continuum, shown with a red dotted line in Fig. 1.7. This
emission is power law-shaped, characterized by a photon index of Γ = α + 1; where α is the
power law spectral index. Γ is an indicator of the slope of the spectrum and it typically ranges
from 1.5-2.1 (Corral et al. 2011; Singh et al. 2011; She et al. 2017).
The primary X-ray emission interacts with the surrounding material and gets reprocessed
through various mechanisms such as photoelectric absorption and Compton scattering. This
gives rise to different reflection features of the X-ray spectrum (Falocco et al. 2014; Ricci et al.
2011; Victoria-Ceballos et al. 2023), such as the Compton-hump and the Fe Kα line. The state
of the reflecting material distinguishes between neutral or ionized reflection.
Neutral reflection occurs in the molecular torus. The dotted blue line in Fig.1.7 depicts
this neutral reflection component. Compton scattering of the primary emission produces the
Compton hump (“reflection hump” in Fig.1.7), peaking at around 20-40 keV. This feature is
particularly prominent in sources with high column density, NH ≥ 1.5× 1024 cm−2 (Ricci et al.
2011). The other typical feature produced via Compton reflection is the Fe Kα emission line
at 6.4 keV (Jimenez-Bailon et al. 2005), shown as a narrow Gaussian-shaped dotted line in
Fig.1.7. This is a fluorescent line, emitted due to an electron transitioning from the L-shell
(n=2) into the K-shell (n=1) after a K-shell electron gets ejected in an event of photoelectric
absorption.
If the reflecting medium is ionized, Compton reflection arises with different spectral char-
acteristics. This situation can occur when Compton reflection takes place in the accretion disk
itself, rather than onto the molecular torus. This produces different characteristics in the spec-
trum, mainly affecting the line position and the continuum shape due to the electron density
1.3. AGN SPECTRAL ENERGY DISTRIBUTION 25
Figure 1.8: The spectrum of NGC 4593 fitted with a component of ionized reflection, shown
in red in the top panel and with a neutral reflection model depicted in blue in the bottom
panel. In both cases, the black line corresponds to the overall model including the power law
continuum and the reflection component. Image adapted from Victoria-Ceballos et al. (2023).
and temperature of the ionized gas. Fig.1.8 depicts these two different spectral shapes from
ionized and neutral reflection.
At lower energies, < 1 keV, many AGN exhibit an excess flux over the main continuum
component, the so-called soft-excess, shown with a green dotted line in Fig. 1.7. The true
nature of this emission is still a subject of debate. It is sometimes interpreted as thermal radia-
tion from the accretion disk or Comptonization of thermal disk photons or as a consequence of
ionized reflection, (Done et al. 2007; Turner & Miller 2009; Boissay et al. 2016; Petrucci et al.
2018). In Fig. 1.9 we show an example of the soft excess in one of our sources, Was 45. In this
case, the soft excess is modeled by a combination of thermal emission fitted with a blackbody
spectrum and an ionized reflection component with the same model from the top panel in Fig.
1.8.
1.3.4 γ−ray emission
Lastly, we briefly mention the high-energy frontier of the electromagnetic spectrum, > 100 keV.
Gamma-ray emission in AGN is primarily associated with the presence of jets. The gamma
rays are produced through various high-energy processes involving the relativistic protons and
electrons, that constitute jets and interact with seed photons from the nuclear emission or the
synchrotron radiation from the jet itself. As it is associated to the jets, this emission is found
in radio-loud AGN. The detection of high energy photons poses a challenge for modern astro-
26 CHAPTER 1. ACTIVE GALACTIC NUCLEI
Figure 1.9: Soft excess found in Was 45 modeled by a combination of thermal emission and
ionized reflection. The dotted lines correspond to different components added to the main
power law, marked with “Fe” for the Fe Kα line, “BB” for the blackbody spectrum, and “Ref”
for the ionized reflection.
physics, due to the need for specialized detectors with significant advancements made possible
by observatories such as the Fermi Gamma-ray Space Telescope (Massaro et al. 2016).
1.4 AGN classification and the Unified Model
AGN are characterized by a realm of very diverse properties. While their spectral differences
provide the means to classify them, their shared characteristics have made it possible to pos-
tulate a model that unifies their diversity in a common framework.
1.4.1 AGN classification
AGN can be classified as high and low luminosity, with Lbol ∼ 1045 erg s−1 as the threshold.
Moreover, high luminosity AGN are typically distant objects, with a redshift of z=0.2 as the
typical separation between distant and nearby objects. The radio flux is another classification
criterion. In Fig. 1.6 we see two orange lines indicating two different fluxes, logν Fν = −6
and -3.5, corresponding to “radio-quiet” and “radio-loud” AGN, respectively. In the radio-loud
regime, there are two main subtypes: FRI and FRII (Fanaroff & Riley 1974). Type FRI cor-
responds to sources where radio emission is primarily concentrated close to the central active
galactic nucleus, with a prominent jet that extends outward but does not exhibit strong struc-
ture or significant collimation at large distances. Type FRII have a significant radio emission on
large scales, displaying bright radio lobes that are located far from the central engine, typically
at several hundred kpc away. Fig. 1.10 shows an example of two radio-loud sources, with 3C
405 showing the structures of a type FRII and 3C 31 corresponding to a type FRI. We do not
delve any deeper into the radio emission and classification given that we will be working with
radio-quiet sources.
In other energy bands, different aspects can provide the means to distinguish between differ-
ent types of AGN. We will describe the two classification schemes that are particularly relevant
to our project: the optical classification into types 1 and 2 and the X-ray classification into
1.4. AGN CLASSIFICATION AND THE UNIFIED MODEL 27
Figure 1.10: Two radio maps from the Very Large Array (VLA) of two radio loud sources with
different classifications. On the left, the FRII source 3C 405 (Cygnus A), and 3C 31, an FRI
source, on the right. Credit: NRAO/AUI.
absorbed and unabsorbed.
1.4.1.1 Types 1 and 2: classification based on optical emission lines
Types 1 and 2 were first identified prior to the formulation of the Unified Model, with mentions
of this classification tracing back to the first analysis of Seyfert sources (Khachikian &Weedman
1974; Weedman 1977). Figure 1.11 provides an observational comparison between these two
types: the top spectrum corresponds to a Type 1 AGN, showing broad Hα and Hβ lines, while
the bottom spectrum corresponds to a Type 2 AGN, displaying only narrow lines (Netzer 1990).
These two main types can be further divided into subtypes by specific characteristics of
the lines, such as the comparison between narrow and broad components and strengths of the
Balmer lines, (Osterbrock 1981). We provide a brief description of these subtypes in the fol-
lowing table, adapted from the Véron-Cetty & Véron (2010) catalog by Ricci (2011).
AGN types 1 and 2 and intermediate.
Type Characteristics
1 Both broad and narrow lines
1.2 Hβ is broad but weaker than in type 1
1.5 The strength of the broad and narrow components of Hβ are comparable
1.8 Broad component in Hα and Hβ is present but very weak
1.9 Broad component is only present in Hα
2 No broad lines
28 CHAPTER 1. ACTIVE GALACTIC NUCLEI
Figure 1.11: Optical spectra of two AGN with emission lines of different strength and width.
The top spectrum corresponds to a type 1, with broad Hα and Hβ (from the BLR) and two
narrow [O III] lines (from the NLR). The bottom spectrum corresponds to a type 2 source,
showing only the narrow component of Hα and Hβ. Credit: Morgan (2002).
1.4.1.2 X-ray classification: obscured and unobscured
In Sects. 1.2 & 1.3.3, we described the X-ray spectrum and its origins, with the primary
component being a power law continuum emitted from the X-ray corona. Soft X-rays (E< 2
keV) are particularly susceptible to absorption by intervening material, which can significantly
alter the shape of the X-ray spectrum. Thus, X-ray absorption can impact the shape of the
primary spectrum, and serves as the crucial classification criterion for AGN in the X-ray regime.
X-ray absorption is mainly characterized by the measured hydrogen column density of the
absorbing material, NH, in units of atoms cm−2. The value of this parameter serves as a
threshold to distinguish between unabsorbed, with no detectable NH intrinsic to the AGN and
absorbed AGN. As the column density increases, the power law component gets progressively
absorbed and the reflection component starts being the dominant component in the spectrum.
A limit value is NH = 1.5 × 1024 cm−2, beyond which the power law gets suppressed and the
reflection component becomes dominant, making the AGN opaque at energies below 10 keV,
with the transmitted power law component getting completely absorbed. These are called
“Compton-Thick” sources. Fig.1.12 shows how the shape of the X-ray spectrum gets affected
by the column density, beginning with an unabsorbed case with a spectrum dominated by the
power law continuum and up to the Compton Thick case with log NH ≥ 24.5. In Fig.1.13
we show a comparison between two real spectra. The left-side spectrum corresponds to the
unabsorbed NGC 7469, where we can see the power law primary continuum. The right-side
spectrum is that of the absorbed Mrk 348, showing the power law continuum significantly ab-
sorbed at energies below 2 keV.
1.4. AGN CLASSIFICATION AND THE UNIFIED MODEL 29
Figure 1.12: The X-ray spectrum is affected by absorption of different column densities. From
a completely unabsorbed source, we see how the power law continuum gets absorbed as the
column density gets higher. For values of NH ≥ 1.5 × 1024 cm−2, the primary continuum gets
completely suppressed at energies < 10 keV. Credit: Gilli et al. (2007).
Figure 1.13: X-ray spectra of two AGN that can be easily distinguished by the absorption of
the primary emission. The left spectrum corresponds to the unabsorbed NGC 7469, a type
1 source. The right spectrum corresponds to Mrk 348, a type 2 source, showing how the
continuum emission is absorbed below 2 keV. Credit: Singh et al. (2011).
According to the measured NH, X-ray absorption can vary significantly, ranging from un-
absorbed cases to heavily absorbed, i.e. Compton-Thick sources. X-ray absorption not only
distinguishes between different AGN types but also offers important information on the ma-
terial surrounding the central region of the AGN. We will discuss more details about X-ray
absorption in Ch. 2, Sect. 2.3.
1.4.2 The Unified Model
The motivation for developing a unifying model for AGN can be traced back to observations
of polarized light in the optical band. NGC 1068 is considered an archetypal type 2 source, yet
30 CHAPTER 1. ACTIVE GALACTIC NUCLEI
Antonucci & Miller (1985) reported the presence of broad emission lines in its polarized spec-
trum. This finding was interpreted as evidence that the broad-line region (BLR) exists but is
obscured from direct view. Furthermore, optical polarization studies of a sample of radio-loud
galaxies (Antonucci 1983) revealed polarization perpendicular to the radio source, suggesting
that photons emitted parallel to the radio source are being scattered. These studies lead to the
conclusion that there is neutral, obscuring material distributed in a toroidal shape around the
central region. This material should cover the BLR while leaving the narrow-line region (NLR)
unobscured; therefore, the proposed scale for this gas region is between 1 − 100 pc (Krolik &
Begelman 1988). This gas region, referred to as the torus, was introduced into the structure
of AGN as a fundamental component of the Unified Model, a framework that has since been
supported by infrared observations (Garćıa-Burillo et al. 2016, 2019).
The Unified Model posits that the fundamental structure is consistent across all AGN types,
but the orientation of the obscuring torus determines the features detected in their spectra. Ac-
cording to Antonucci (1993), when observed edge-on, the torus completely obscures the central
engine and the broad-line region (BLR), resulting in the classification of the object as a type 2
AGN. In contrast, when a source is viewed face-on, the observer has a direct line of sight to the
central region, revealing both broad and narrow emission lines characteristic of a type 1 AGN.
This classification scheme is grounded in the optical emission lines discussed in Sect. 1.4.1.1.
Two years after the initial proposal of the Unified Model, Urry & Padovani (1995) expanded
the framework to address differences observed at radio frequencies, focusing on the presence or
absence of jets and their orientation relative to the observer’s line of sight.
A schematic illustration of the Unified Model is depicted in Fig.1.14, where we can dis-
tinguish the various types of AGN. The diagram distinguishes two main classification criteria
that divide AGN into four main groups: high/low power divides AGN according to their lumi-
nosity, threshold value of Lbol ∼ 1045 erg s−1, radio-loud/radio-quiet separates them according
to the flux of their radio emission, with types FRI and FRII indicated in the radio-loud sec-
tion. Many publications provide insights in AGN radio emission (such as Wilson & Colbert
1994; Xu et al. 1999; Cirasuolo et al. 2003; Kellermann et al. 2016, etc). Along the diagram,
there are different lines of sight marked, indicating the orientation scheme of the Unified Model.
We direct our attention to the Seyfert section in the lower left-hand quadrant of the dia-
gram: low-luminosity, radio-quiet sources mainly classified on whether the line of sight passes
through the obscuring torus or looks directly into the central region. In Fig.1.14, we see lines of
sight marked with BL or NL (for radio-loud sources). In the Seyfert sources quadrant, they are
analogously marked as Seyfert 1 and Seyfert 2 (Sy 1 and Sy 2 from now on). This distinction
is due to the orientation of the torus, marked as “dusty absorber” in the diagram. In Sy 1, the
line-of-sight looks directly into the central region, observing both the BLR and the NLR. In Sy
2, the line-of-sight is aligned with the torus, with the BLR obscured and the spectrum showing
only narrow emission lines.
1.4.2.1 Modifications and challenges to the Unified Mode
Many observations have provided evidence that supports the Unified Model. For instance, the
detection of hard X-rays (> 2 keV) in Sy 2, (e.g. Maiolino et al. 1998): while in a Sy 2, the
BLR may be completely covered by the torus, the hard X-ray spectrum is not absorbed by its
column density. This finding aligns with the predictions of the Unified Model, suggesting that
the intrinsic properties of Sy 1 and Sy 2 galaxies are fundamentally the same and the differences
1.4. AGN CLASSIFICATION AND THE UNIFIED MODEL 31
Figure 1.14: The Unified Model for AGN. This scheme divides AGN into main classes: high and
low luminosity, radio-quiet, and radio-loud. The scheme explains that all the main differences
in AGN types can be accounted for by introducing an obscuring torus and relating the line of
sight to its orientation. Credit: Beckmann & Shrader (2012).
32 CHAPTER 1. ACTIVE GALACTIC NUCLEI
arise from the orientation of our line of sight with respect to the obscuring torus. However,
advancements in detector technology and detection techniques have made high-resolution data
across all bands of the electromagnetic spectrum available, posing new challenges for the classi-
cal unification scheme. These improvements have revealed complexities in AGN behavior that
cannot be easily reconciled with the classical Unified Model.
A review by Bianchi et al. (2012) provides important insights into the complexity of obscu-
ration in AGN, leaving the idea of a homogeneous torus behind. In general, X-ray obscuration
and variability studies reveal the presence of absorbing clouds of different column densities,
located at various scales, and possibly resulting from a combination of neutral and ionized ab-
sorbers. Thus, the modeling of the obscuring torus has proven to be more complex than what
was originally postulated. MIR interferometric studies have confirmed the presence of sub-
parsec and parsec scale dust distributions, indicating two-component structures with varying
temperatures (Garćıa-Burillo et al. 2016, 2019). The scatter of the silicate absorption feature
with X-ray gaseous column density suggests a clumpy structure of the dusty absorbers, (e.g.
Nikutta et al. 2009).
Since the early 2000s and with the improvement in IR detectors, a model where the gas of the
torus is not distributed in a homogeneous toroidal shape but rather in dense clouds or clumps
of molecular material, which can vary significantly in density and size (Nenkova et al. 2002;
Elitzur & Shlosman 2006; Nenkova et al. 2008a,b). Moreover, the scale of the absorbers has also
been diversified. Some studies have reported that column density variability is associated with
high-velocity clouds, > 103 km s−1 that produce eclipse-like events in X-ray spectra and are
therefore associated with BLR scale (e.g. Risaliti et al. 2005; Beuchert et al. 2015). These new
findings have emerged with the growing richness of X-ray archives, which allowed multi-epoch
spectroscopy to be performed in a larger number of AGN. The resulting not-so-uniform view
of the obscuration process in AGN has challenged the idea of a static, homogenous absorber,
as initially postulated by the unification model.
In their work, Ramos-Almeida & Ricci (2017), provide an IR/X-ray review of the known
properties of the circumnuclear environment around the SMBH and its accretion disk. They
find this material to be clumpy and dynamic, with its covering factor influenced by the accre-
tion itself. The IR analysis shows the obscuring material as a bridge between BLR and NLR
and is made of an equatorial disk and a polar component. Moreover, Spinoglio & Fernández-
Ontiveros (2019) suggest that the orientation effect does not sufficiently explain the type 1 vs
type 2 dichotomy, as transitions between types have been detected, and that the role of the
host galaxy may play an important role.
An interesting remaining issue is the findings of “true” Sy 2 sources. Tran (2001, 2003) and
further work have found Sy 2 sources with no hints of BLR, suggesting weak emission from the
main engine, incapable of producing a BLR. Another finding that poses an open question for
the unifying scheme is AGN lacking evidence of the NLR, originally described as “ubiquitous”,
(e.g. Armus et al. 2007; Bianchi 2009; Panessa et al. 2009). The Netzer (2015) review also pro-
vides insights into the true distribution and properties of the torus, with newer models involving
disk winds and hydrodynamic simulations that connect the large-scale galactic disk to the in-
ner accretion flow. It also highlights the issue surrounding true type 2 AGN. Changing-look
AGN also challenge the classical unified model by revealing that AGN can undergo significant
transformations in their observed properties over relatively short timescales, (Hon et al. 2020;
Ricci & Trakhtenbrot 2023).
1.4. AGN CLASSIFICATION AND THE UNIFIED MODEL 33
In conclusion, while the Unified Model has provided a solid foundation for the study of
AGN, the continuous improvement in detectors and modeling techniques is constantly fueling
newer research. From the detection of hard X-rays in Sy 2 galaxies to the structure of the ob-
scuring torus revealed by X-ray variability and MIR interferometric studies, our understanding
of AGN continues to evolve. This project aims to add to this evolution by comparing measured
optical polarization with an analysis of the X-ray spectra.
Chapter 2
Optical polarization in Seyfert galaxies
Studies of polarized light provide crucial insights into the central regions of AGN, revealing
emissions from otherwise obscured structures. As described in Sect. 1.4.2, polarization re-
search laid the foundation for the AGN Unified Model by confirming the presence of a BLR
in type 2 sources (Antonucci 1984; Antonucci & Miller 1985). Polarization studies are partic-
ularly valuable for examining the innermost and unresolved pc-scale regions within the AGN
and the the surrounding circumnuclear material (Axon et al. 2008). These studies also enhance
our understanding of the unification scheme due to their sensitivity to AGN geometry. For
instance, Marin (2014, 2016) used polarimetric studies to estimate AGN nuclear inclination.
Additionally, spectropolarimetry has allowed for key measurements related to AGN evolution,
including magnetic fields associated with the accretion disk (Piotrovich et al. 2017), constraints
on the size of the BLR (Songsheng & Wang 2018), and estimates of the supermassive black
hole (SMBH) mass (Baldi et al. 2016; Popović et al. 2018; Piotrovich et al. 2023).
In this chapter, we detail the work of Smith et al. (2002, 2004, 2005), (from here on Smith
et al.). They provide a significant study on the polarization properties of a sample of 46 Seyfert
type 1 sources and establish a unification scheme that explains that has been further supported
by work such as (Marin 2014, 2016; Baldi et al. 2016, etc.). This unification scheme provides
the basis of our study, therefore we dedicate this chapter to describing some pivotal work on
polarized Seyfert galaxies and, specifically, the measurements by Smith et al. We will then
present the main goal of this project based on a comparison between the polarization unified
scheme and a study on X-ray absorption.
2.1 Optical polarization
Light as an electromagnetic wave is characterized by mutually perpendicular electric and mag-
netic fields, with both fields oscillating orthogonal to the direction of propagation. When the
direction of propagation varies randomly in all possible planes, it is said that the light is unpo-
larized.
Polarization of light occurs when the electric field aligns along a specific axis, typically
within a singular plane, resulting in a uniform wave orientation. The orientation of polarized
light denotes its polarization type. Linear polarization has an E with constant amplitude,
propagating within a single plane. Circular polarization happens when E vectors oscillate in
a circular pattern, rotating in time at the frequency of the radiated energy. Lastly, elliptical
polarization is a combination of linear and circular polarization, and, in general, any form of
polarization can be considered a form of partially elliptically polarized light. Fig. 2.1 illus-
34
2.1. OPTICAL POLARIZATION 35
Figure 2.1: Process of light polarization via transmission through a filter, aligning the
electric field (E) in a specific orientation. The image depicts two examples of polariza-
tion according to the form of light propagation, linear and left-handed circular polariza-
tion. Credit: Wikipedia contributors. Polarizer. In Wikipedia, The Free Encyclopedia:
https://en.wikipedia.org/wiki/Polarizer. Author: Dave3457.
trates light becoming polarized after passing through a filter that transmits the electric field in
a particular direction, with the first filter producing linear polarization and the quarter-wave
plate (λ/4) causing the linearly polarized light to become circularly polarized.
A form of characterizing polarized light is through the four Stokes parameters I, Q, U
and V, (Stokes 1851). Consider an electric field, E represented by two orthogonal vectors as
given by Eq. 2.1. The propagation time is indicated by t, and z corresponds to the direction.
Amplitudes are denoted with Ex, Ey, and ϵx, ϵy are their respective phase, with k denoting the
wave number and w the angular frequency:
E(z, t) = Ex cos(κz − ωt+ ϵx)̂i+ Ey sin(κz − ωt+ ϵy)ĵ (2.1)
Thus, the four Stokes parameters are defined in terms of the electric field by Equations 2.2,
2.3, 2.4 and 2.5:
I = E2
x + E2
y (2.2)
Q = E2
x − E2
y (2.3)
U = 2ExEy cos ϵ (2.4)
V = 2ExEy sin ϵ (2.5)
Linear polarization is described by the Q and U parameters. The Q parameter corresponds
to polarization along the horizontal and vertical directions, and the U parameter corresponds
to linear polarization at ±45circ relative to the horizontal and vertical axes. The V parameter
describes the elliptical/circular polarization and, lastly, the parameter I represents the total
intensity, a sum of unpolarized and polarized light: I = Iu + Ip. For Iu = 0, i.e. fully polarized
36 CHAPTER 2. OPTICAL POLARIZATION IN SEYFERT GALAXIES
Figure 2.2: Polarization by scattering, where light interaction with matter induces partial
or complete polarization through dispersion. Credit: Byju’s. Polarisation by Scattering:
https://byjus.com/physics/polarisation-by-scattering/.
light, I = Ip = (Q2+U2+V 2)
1
2 . In the case of partially polarized light, the degree of polarization
is expressed by Eq.2.6 and represents the fraction of polarized light. The percentage of polarized
light relative to the total light is given by Eq. 2.7. The polarization position angle, Eq. 2.8,
represents the direction in which the electric field of the polarized light oscillates.
p =
√
(Q2 + U2 + V 2)
I
(2.6)
p(%) =
Ip
I
× 100 (2.7)
θ =
1
2
tan−1
(
U
Q
)
(2.8)
Phenomena such as reflection, transmission, and scattering can cause light to be polarized.
In particular, polarization by scattering occurs when unpolarized light interacts with matter
along its propagation path, getting dispersed and becoming partially or fully polarized, depend-
ing on the scattering angle and the scattering particles. A simple illustration of polarization
by scattering is shown in Fig. 2.2, where an incident light ray, marked with E, encounters
a scattering particle. This interaction causes the electric field to get orientated in a specific
direction. The observer can thus detect light that is fully polarized, partially polarized, or
unpolarized depending on the orientation of the observed ray.
In general, Seyfert galaxies exhibit some percentage of their light linearly polarized due to
scattering of the light from the central region by dust or electrons in the environment of the
AGN (Martin et al. 1983; Thompson & Martin 1988). Scattering processes can take place, for
example, in the BLR, where free electrons and/or ionized outflowing gas can produce signifi-
cant polarization. In the molecular torus, the scattering by dust leads to wavelength-dependent
polarization.
Polarization by scattering varies depending on the size and type of particle encountered.
Particularly, Rayleigh scattering occurs when light interacts with particles much smaller than
its wavelength. This form of scattering is strongly wavelength-dependent, being more effective
at shorter wavelengths. The intensity of the scattered light I ∝ λ−4, indicating how much light
2.2. POLARIZED SEYFERTS GALAXIES 37
is being scattered in a particular direction. Thus, Rayleigh scattering is more effective at shorter
wavelengths, leading to a higher degree of polarization in the ultraviolet and blue regions of the
spectrum, where scattering dominates, causing a decrease in polarization at redder wavelengths.
On the other hand, Thomson scattering occurs when light interacts with free electrons,
regardless of the wavelength of the incident light. The incoming electromagnetic wave induces
oscillations in the electron, causing it to re-emit light in multiple directions. When observed
at a 90-degree angle from the incident direction, the scattered light is fully polarized (Fig.
2.2). The probability of Thomson scattering is characterized by the Thomson cross-section a
constant that determines the likelihood of scattering for each interaction with a free electron.
The value of this constant is expressed in Eq. 2.9, where e is the charge of the electron and me
its mass, ϵ0 is the permittivity of free space, and c is the speed of light. Thomson scattering is
especially significant in ionized gas such as the plasma typically found in AGN environments.
σT =
8π
3
(
e2
4πϵ0mec2
)2
≈ 6.65×−29 m2 (2.9)
Spectropolarimetry allows the detection and measurement of the polarization properties of
light across a range of wavelengths, determining the orientation and intensity of the emitted
light. Most notorious, studies on polarized optical emission from Seyfert galaxies provided the
foundation for the Unified Model.
2.2 Polarized Seyferts galaxies
Seyfert galaxies are distributed in the low end of AGN properties: low redshift, radio quiet,
and moderate luminosity, Lbol = 1043 − 1046 erg/s. In the optical band, Seyfert galaxies are
mainly classified as Sy 1 and Sy 2 (types 1 and 2), with the difference being that broad emission
lines are only present in Sy 1, and both types exhibit narrow emission lines in their spectra.
Interestingly, spectropolarimetric studies of Sy 2 revealed the presence of polarized broad emis-
sion lines (e.g. Antonucci & Miller 1985; Miller & Goodrich 1990). This finding served as the
foundation of the Unified Model. Polarization studies of AGN have taken place since the 1960s,
and, in general, the majority of Seyfert galaxies show some form of linear optical polarization,
typically low, p < 0.2− 0.5%, (Martin et al. 1983; Thompson & Martin 1988).
In particular, the work on NGC 1068 has had significant relevance to the studies of polarized
AGN. Located at a distance of ∼ 13.5 Mpc, NGC 1068 (M77), is the brightest type 2 source
in the local Universe, and is considered the archetype of a Sy 2 galaxy. As such, it has been
extensively studied since the early 20th century, (Slipher 1917); with Walker (1966) being one
of the first reports of polarized light found in AGN. In later years, Angel et al. (1976) determine
that the detected polarization pattern was evidence of a distribution of dust clouds, motivated
by the findings that infrared emission is arising from an optically thick dust cloud, (Becklin
et al. 1973). Angel et al. (1976) concluded that the detected polarization of NGC 1068 is due to
scattering. Subsequently, Miller & Antonucci (1983) found a very highly and constantly polar-
ized continuum, with p ∼ 16%, and suggested that the polarization mechanism could be either
synchrotron radiation or scattering by an asymmetrical distribution of electrons. In further
work, Antonucci & Miller (1985) found that NGC 1068 shows a highly polarized component
of the Balmer and Fe ii emission lines. As a Sy 2, these broad components are not detected
in unpolarized light. Moreover, the polarization of the broad lines is different in strength and
38 CHAPTER 2. OPTICAL POLARIZATION IN SEYFERT GALAXIES
position angle than that of polarization that characterizes the narrow lines, indicating that the
lines arise from two different regions with different physical conditions. It is also reported that
polarization position angle is perpendicular to the symmetry axis, determined by the projected
radio source.
The findings of Sy 1-like broad emission lines in the archetype of Sy 2 served as the founda-
tion of the Unified Model, described in Sect. 1.4.2. Antonucci (1993) postulated a model where
the continuum source and the BLR are enveloped, and therefore obscured, by a neutral and
optically thick toroidal gas distribution. This gas and dust region, called the torus, is co-axial
to the rotation of the accretion disk, the principal axis of the system. Light from the central
region escapes through a bi-conical region that extends along this axis and corresponds to the
NLR. Electrons located above and below the torus, in this bi-conical region, are responsible for
scattering light from the central region, and polarizing it. Given that the E field is orthogonal
to the scattering plane, the polarization position angle results perpendicular to the principal
axis. Thus, this model explains that both BLR and NLR are present in all Seyfert sources and,
while the emission from the NLR is ubiquitous, the emission from the BLR cannot always be
detected in direct light due to the orientation of the torus.
Previous studies (Antonucci 1983, 1984) found that Sy 1 galaxies typically exhibit polariza-
tion with a position angle parallel to the principal axis of the system, whereas in Sy 2 galaxies,
the position angle is perpendicular. This indicates that two distinct scattering regions, of dif-
ferent orientations, contribute to the observed polarization: a polar-scattering region and an
equatorial-scattering region, explained in Sect. 2.2.1.2. Subsequent works (Brindle et al. 1990;
Miller & Goodrich 1990; Goodrich & Miller 1994) have provided additional support for these
findings.
Building on this, Tran (2001, 2003) provide a spectropolarimetric study of Sy 2 sources in
millimetric bands. They corroborate the findings of polarized Sy 1 regions hidden within the
obscuring material of Sy 2 sources. Intriguingly, they identified several “true Sy 2” galaxies
in which this polarized BLR appears to be completely absent. As mentioned in Sect. 1.4.2.1,
this poses a challenge to the classical Unified Model, as it suggests that there are type 2 AGN
whose differentiation from type 1 goes beyond an orientation effect.
In this context, we now turn to the work of Smith et al., who developed a model to ex-
plain the polarization observed in different Seyfert galaxies within the framework of the Unified
Model. Their study provides a unifying perspective on the polarization properties of these
systems, providing valuable insights into the scattering regions and the geometry of AGN.
2.2.1 Towards a unification scheme for polarized Seyfert galaxies
We provide a review of the work by Smith et al. as it serves as the foundation of this project.
This analysis of polarized Seyfert galaxies is distributed in three articles by Smith et al. and
reports the characteristics of the polarization detected in a sample of 46 Sy 1 galaxies. Data
were obtained primarily between 1996 and 1999 using the 4.2m William Herschel Telescope
(WHT) and the 3.9m Anglo–Australian Telescope (AAT), with additional observations taken
in 1991 and 2002.
The observations were taken at equal exposures for each of the half-wave plate angles of
0, 22.5, 45, 67◦. Each exposure produces two mutually perpendicular spectra of opposite po-
2.2. POLARIZED SEYFERTS GALAXIES 39
larization. The Stokes parameters can be derived from the intensity ratio of these two spectra.
The first two exposures yield the Q Stokes parameter, the next two exposures the U parame-
ter, (Tinbergen & Rutten 1997). In practice, the software Starlink is used, with particular
routines such as ccd2pol in tsp (The FIGARO Time Series/Polarimetry Package) to process
and reduce the data and obtain the intensities of the orthogonally polarized spectra and, from
there, calculate the I, Q, U Stokes parameters and the average of degree of polarization, p,
and polarization position angle, θ, for both the continuum and the Hα line.
Determining the characteristic values of polarization intrinsic to the AGN must take into
consideration possible contamination by foreground polarization, i.e. interstellar polarization
from the host galaxy and polarization produced from the Galaxy. In order to determine Galactic
polarization, Smith et al. refer to a relationship between E(B-V), extinction, and polarization
percentage established by Serkowski et al. (1975). Smith et al. were able to separate polariza-
tion intrinsic to the AGN from interstellar polarization because the latter changes weakly with
wavelength, whilst AGN polarization shows changes in degree of polarization across a short
range of wavelengths.
2.2.1.1 Characteristics of polarized type 1 Syeferts
The reported general characteristics of the polarization of the sample of Sy 1 are the following:
• The average detected polarization ranges from 0.5− 5%.
• All sources exhibit changes in both the degree of polarization and the polarization PA as
a function of the line of sight velocity across the Hα line profile. These variations can
indicate a kinematic structure of the Hα line, revealing different regions and velocities
within the scattering material that contribute to the observed polarization.
• Some sources exhibit rotation from blue to red in the PA associated to the Hα line.
This describes a change in the direction of polarized light as function of wavelength and
suggests that different parts of the BLR or other scattering regions are contributing to the
polarized light at different velocities. A rotation in PA can indicate a geometry such that
orientation of the scattering material varies with velocity, like the presence of a rotating
disk, for example.
• The polarization of the Hα line is characterized by a dip of polarization over the line
core, i.e. central peak emission, flanked by an increase over the line wings, explained by
thermal velocity broadening.
• For some objects of the sample, the average polarization PA is approximately parallel
to the principal axis of the system. Interestingly, some of the objects, all Sy 1, have an
average polarization PA perpendicular to the axis.
The polarization of Akn 120 is shown in Fig. 2.3. The first image, right, corresponds to the
plots produced for the entire polarized flux, specifically around the Hα wavelength. The plots
show that the core of the broad Hα line has similar polarization to that of the red wing of the
line, while the peak in degree of polarization of the blue wing differs. The plots on the left
correspond to the polarization of the Hα line, with subtracted continuum. The first panel in
both panels shows a polarization PA rotation across the line. We see that the PA deeps below
that of the continuum in the blue wing and and exhibits an increment in θ towards the red
wing, (Smith et al. 2002).
40 CHAPTER 2. OPTICAL POLARIZATION IN SEYFERT GALAXIES
Figure 2.3: The polarization of Akn 120. The right set of plots correspond to the spectropo-
larimetric data of Akn 120, and the left show the polarization of the Hα line after subtracting
the continuum. For both cases, the first panel depicts the polarization PA, followed by the
polarized flux density, the percentage of polarization and the total flux density. Credit: Smith
et al. (2002).
Polarization PA changes are interpreted as a function of velocity shifts, with scattering
varying according to the line of sight, and constraining the geometry of the scattering region,
placing it close to the BLR. The BLR is confined to a rotating disk, with a co-planar disk of
scattering material surrounding it.
The decrease in polarization at the core of the Hα line is primarily due to the combination
of both direct and scattered light in the observed spectrum. In the core of the line, the direct
light, which originates from the central regions of the AGN and travels without being scattered,
is narrower in profile and more intense compared to the scattered light. This direct component
dilutes the polarization because it dominates the core region, reducing the overall polarization
signal. In contrast, the light scattered at the BLR or scattering regions further away, usually
shows broader profiles due to the range of velocities contributing to the scattering. When the
direct light overlaps with this scattered light, the polarization signal gets reduced because the
unpolarized direct light adds to the total flux without contributing to the polarization. This
effect is particularly pronounced in the line core, where the direct emission is strongest.
In comparison, Sy 2 show very small changes in the polarization PA and local peaks in the
degree of polarization associated with the broad lines. This is explained as an effect of the
presence of a diluting unpolarized continuum emitted by the accretion disk. Similar to the
effects on the polarization of the Hα line, the contribution of unpolarized light to the overall
spectrum also causes depolarization. This diluting component is also responsible for a reported
increase in the degree of polarization at shorter wavelengths. The increase in the degree of
polarization at shorter wavelengths in Seyfert galaxies can be attributed to more efficient scat-
tering processes at these wavelengths, causing a decrease in polarization at redder wavelengths.
2.2. POLARIZED SEYFERTS GALAXIES 41
2.2.1.2 Polarization in the context of the Unified Model, interpreting polarization
position angle
Smith et al. were able to determine the average polarization PA parallel to the principal axis
in 11 Sy 1 sources, in agreement with previous studies such as Antonucci (1983, 1984); Brindle
et al. (1990), which also report on Sy 2 sources. Interestingly, Smith et al. find 12 Sy 1 sources
that exhibit polarization characteristics similar to Sy 2 sources, mainly that the average po-
larization PA is perpendicular to the principal axis. These results led to the postulation of a
unification scheme for polarized Sy, shown in Fig.2.4. This scheme is based on the presence of
two scattering regions: the polar-scattering region and the equatorial-scattering region. The
inclination i is measured from the principal axis towards the line-of-sight, with αP and αE
corresponding to the half-opening angles of the polar and equatorial scattering regions respec-
tively. The unification scheme proposed by Smith et al. proposes that both scattering regions
are present in all Seyfert galaxies and the observed polarization is due to an orientation effect.
Figure 2.4: Scattering regions responsible for the polarization of Seyfert galaxies. The polar
scattering region along the principal axis with half opening angle αP and the equatorial scat-
tering region, co-planar to the BLR and of αE half opening angle. The inclination is taken from
the principal axis to the line-of-sight. Credit: Smith et al. (2004).
In the polar scattering scenario, polarization is due to scattering by material located in the
ionization cones formed by radiation from the central region that extend along the principal
axis, corresponding to the NLR, see Sect. 1.2.2.2. Light scattered in this region is character-
ized by an average polarization PA perpendicular to the axis. This scenario corresponds to the
typical dominant polarization observed in Sy 2 sources.
42 CHAPTER 2. OPTICAL POLARIZATION IN SEYFERT GALAXIES
The second scattering region, the equatorial-scattering scenario, is described as a distri-
bution of material within the obscuring torus and along its equatorial plane, forming a disk
co-planar to the BLR. Light polarized in this region will have a PA parallel to the principal
axis, as measured in several Sy 1 sources.
In accordance with the Unified Model, (Antonucci 1993), a Sy 2 source is edge-on, with the
line-of-sight aligned with the obscuring torus. Emission from the central engine and the BLR
escapes through the ionization cones, and material in the NLR scatters the emission, which will
show as PL-pol. On the other hand, Sy 1 sources are looked at face-on, i.e. the line-of-sight
sees directly into the central engine and the BLR. Both polar and equatorial regions are there-
fore visible in Sy 1 sources, with a tendency to be dominated by equatorial scattering, EQ-pol.
Finally, Sy 1 exhibiting PL-pol represent an intermediate case, with the orientation paradigm
positioning the line of sight passing through the upper layers of the torus. According to this
scheme, we can expect four cases for this unification scheme for polarized Seyferts:
1. Null polarization, i = 0◦ - The line-of-sight is aligned to the main axis. Both scattering
regions exhibit circular symmetry, and the polarization vectors cancel each other out,
resulting in null polarization. This scenario would correspond to a classical Sy 1 source,
where there would be no obscuration of the central emission.
2. EQ-pol, 0 < i < 45◦ - There is little to no obscuration towards the central region. The two
scattering regions are present, orthogonal to each other. Polarization from the equatorial
scattering region is the stronger component. This effect of dominant polarization can be
due to photons from the central region not getting blocked or diluted by the geometry
of the system, while the polar-scattered light from the NLR tends to be less dominant
because the scattering angles are not optimal for high polarization. This scenario also
corresponds to a Sy 1 source but is looked at through a range of inclinations between the
axis and before intercepting the torus.
3. PL-pol, i ≈ 45◦ - This represents an intermediate case between EQ-pol Sy 1 and PL-pol
Sy 2. These cases are interpreted as emission from the central region passing through
the upper layers of the torus, where the light scattered by the equatorial region gets
attenuated, producing the polar scattered light to dominate the spectrum.
4. PL-pol, i > 45◦ - The line-of-sight corresponds to an edge-on view of the source, with the
light from the central region getting obscured by the torus. This case corresponds to the
typical case of Sy 2, with only polar polarization present in the spectrum.
2.3 X-ray analysis: a test for the optical polarization
unification scheme
For this project, we propose an independent test of the unification scheme presented by Smith
et al. through X-ray analysis, specifically, by studying the presence of X-ray absorption.
When present, X-ray absorption is a key process that shapes the observed X-ray spectra, as
X-ray emission carries the signatures of the material along the line of sight. Interactions such
as photoelectric absorption and Compton scattering attenuate the intensity of X-ray photons.
Photoelectric absorption is effective at lower energies, becoming irrelevant for energies > 10
keV, and for column densities between NH ∼ 1021−1024 cm−2 (Ricci 2011). On the other hand,
2.3. X-RAY ANALYSIS, TESTING THE UNIFICATION SCHEME 43
Compton scattering, where the interaction of X-ray photons with free electrons can result in
the loss of energy of the X-ray photons, is energy independent and becomes effective at higher
column densities, NH ≥ 1024 cm−2.
Absorption is attributed to the presence of material along the line of sight, which can be
in a neutral or ionized state. The properties of the absorber, such as its ionization state, col-
umn density, and chemical composition, can significantly affect the X-ray spectrum, leading to
features like absorption edges, absorption lines, and complex spectral shapes. In the following
sections, we will explore the so-called “cold absorber” and “warm absorber”, due to neutral
and ionized material respectively.
2.3.1 Cold Absorber
The cold absorber refers to absorbing material in a neutral atomic state. X-ray photons emitted
from the inner region interact with atoms from the surrounding mediums, ionizing or exciting
the material through photoelectric absorption. This effect depends on the energy of the incident
photon and the atomic number of the encountered atom, with the photoelectric cross section
expressed as a function of the atomic number, Z, and the photon energy, E, as described in
Eq. 2.10, with β typically ranging between 2.5-3 for X-rays (Longair & Longair 1992). As the
photon energy increases, the efficiency of photoelectric absorption decreases.
σph(E,Z) ∝ ZnE−β (2.10)
The photoelectric absorption also depends on the column density of the medium, with
the cut-off energy, ECO, for the emitted power law dropping below certain energies. For
1022 cm−2 < NH < 1023cm−2 → ECO < 2keV, for values of 1023 cm−2 < NH < 1024cm−2
ECO ranges between 2 and 10 keV, and for values above NH > 1024 the ECO = 10 keV.
At higher column densities, the dominant effect is Compton scattering, resulting in two main
regimes mentioned in Sect. 1.4.1.2: Compton Thin and Compton Thick. The limit value of
1.5× 1024 cm−2, corresponds to an optical depth, τ = 1 for Compton scattering which depends
on the electronic density of the medium, Ne, and the Thomson cross section σT (Longair &
Longair 1992):
τ = NeσT → Ne =
1
σT
≈ 1.5× 1024 (2.11)
As the column density increases, the power law continuum gets absorbed, with Compton
thick sources showing a power law suppressed below < 10 keV, and the spectrum being domi-
nated by the reflection component (Risaliti 2002; Maiolino & Risaliti 2007). The FeKα emission
line becomes more prominent, reaching equivalent widths1 of up to 1 keV. For NH ∼ 1025 cm−2,
the continuum gets completely absorbed at energies < 10 keV. In Fig. 2.5, we show a plot of
an unabsorbed power law model with an index of Γ = 1.8 represented with a blue line. This
plotted model is a result of the power law convolved with the response of the EPIC-pn detector.
Effects from instruments and background noise cause a deviation from a straight line. The rest
of the colored lines represent the same power law model affected by cold absorption of different
column densities starting with NH = 1021 cm−2 and up to 1023 cm−2, where we see the power
1Equivalent Width (EW) is a measure of the strength of the line relative to the continuum, representing the
width of a rectangle with the same area as the line.
44 CHAPTER 2. OPTICAL POLARIZATION IN SEYFERT GALAXIES
law strongly suppressed at energies < 2 keV.
The location of the cold absorber is not always associated with the torus. For instance,
constraints on the location of this absorber have placed it in scales up to 100 pc (Guainazzi
et al. 2005), identifying Sy 2 whose X-ray spectra are affected by absorption but are, overall,
in the Compton-Thin regime. On the other hand, studies of multi-epoch column density vari-
ations in some AGN indicate that there can also be an inner cold absorber, in the scale of a
few pc, located in the proximity of the BLR (Risaliti et al. 2005; Elvis et al. 2004; Maiolino
& Risaliti 2007). Furthermore, these studies indicate that variations can go up to changing a
source from Compton-Thin to Compton-Thick, with obscuration events blocking the primary
X-ray emission. Fig. 2.6 illustrates two cases of obscuration events that significantly change
the shape of the spectrum. The first case, on the right-hand side, corresponds to NGC 1365
(Risaliti et al. 2005). This source showed a significant variation between observations 2 and 3,
plotted in red and blue respectively, where the spectrum in blue is dominated by the primary
emission, and the spectrum in red is dominated by reflection. The left image shows the plot
of ESO 323-G077 (a source analyzed in this thesis), where obscuration drastically changes the
shape of the spectrum to a reflection dominated emission, shown in red.
Figure 2.5: A power law model, Γ = 1.8, as it gets affected by cold absorption of different
column densities. We see how the power law gets suppressed at higher energies as the column
density gets higher.
Figure 2.6: Cold absorption in NGC 1365, right, and ESO 323-G077, left. Both sources show
variation due to obscuration events that shape the spectrum from the one dominated by the
primary emission to the ones dominated by reflection, depicted in red for both cases. Credit:
Image adapted from the studies of Risaliti et al. (2005); Miniutti et al. (2014).
2.3. X-RAY ANALYSIS, TESTING THE UNIFICATION SCHEME 45
2.3.2 Warm Absorber
X-ray absorption due to ionized material was first detected by Halpern (1984), who analyzed
the spectrum of an AGN using data from Einstein (HEAO-2). More detailed analysis of these
“warm absorbers” was possible in the following decades thanks to missions such as the Ad-
vanced Satellite for Cosmology and Astrophysics (ASCA), in operations from 1993 until 2001.
Using data from ASCA, Reynolds & Fabian (1995) presented the modeling and analysis of this
ionized absorber, including constraining its location to be within or just outside the BLR. In a
study of a sample of 24 Sy 1 galaxies, Reynolds (1997) found that ∼ 50% of the sources exhibit
the presence of warm absorbers.
The warm absorber is characterized by its column density and ionization parameter depend-
ing on the X-ray luminosity Lx, the electron density n of the absorbing gas, and the distance
from the primary source R, Eq. 2.12. For a column density of NH ∼ 1022 cm−2 and ionization
parameter of ξ = 35 erg cm s−1, the estimated gas temperature is T = 105 K. Among the first
modelings of warm absorbers, Fig. 2.7 shows an ASCA observation of MCG-6-30-15, (Fabian
et al. 1994). The image shows the absorbed spectrum on the right, with the residuals on the
left showing the OVII and OVIII edges at ∼ 740, 870 eV respectively. The parameters of the
absorber found for this source are: logNH = 22.13, log ξ = 44.7.
ξ =
Lx
nR2
(2.12)
Figure 2.7: The spectrum of MCG-6-30-15 before the modeling of the warm absorber. The
image on the right corresponds to the spectrum fitted by a power law, which only produces
a good fit at higher energies, (< 3 keV). The soft band shows clear signs of absorption, as
displayed in the residuals on the left, where the edges of OVII and OVIII are identified. Credit:
Fabian et al. (1994).
It is thanks to the ongoing Chandra and XMM-Newton missions, both launched in 1999,
that astronomers were able to obtain the first high resolution X-ray data of AGN and conduct
more in-depth studies of the warm absorption phenomena. Kaastra et al. (2000) published the
first high-resolution X-ray spectrum of a Seyfert galaxy, NGC 5548. The analysis revealed the
presence of blue-shifted absorption lines corresponding to highly ionized material, which they
modeled as an outflowing shell of ionized gas. Krongold et al. (2003), presented a tool that
allows for detailed modeling of the ionized absorber, the PHASE ionization code. The analysis
of NGC 3783 with PHASE revealed over 100 absorption features that suggested a two-phase
ionized absorber consistent with one outflow at ∼ 750 km s−1. We present an example of
the high-resolution spectrum of NGC 3783, as taken by the Chandra Medium Energy Grating
46 CHAPTER 2. OPTICAL POLARIZATION IN SEYFERT GALAXIES
(MEG) and analyzed with PHASE in Fig. 2.8.
Figure 2.8: High-resolution spectrum of NGC 3783 taken by the MEG instrument on board
the Chandra X-ray Observatory. The absorption lines predicted by PHASE are marked at the
top of each plot. Credit: Krongold et al. (2003).
The location of the warm absorbers in AGN spans a wide range of spatial scales. Blustin
et al. (2005) argues that warm absorption is not an essential process of AGN as it is not
ubiquitous in samples of Seyfert sources. Moreover, their detailed analysis based on the use
of photoionization codes and grating spectra, reports on the presence of ionized material dis-
tributed in multiple layers of different column densities and ionization parameters. Blustin
et al. (2005) estimates the location of the absorber through the outflow velocity, assuming that
it needs to be equal or larger than the escape velocity associated with the SMBH. This is
expressed in Eq. 2.13, which gives the minimum distance of the absorber to the SMBH, with
G being the gravitational constant and R the distance to the SMBH. The absorber is assumed
geometrically thin, with depth ∆r so ∆r/R ≤ 1 and the column density, NH , is expressed
as a function of the material density n(R) and volume filling factor Cv: NH ∼ n(R)Cv∆r.
This expression, combined with the ionization parameter from Eq. 2.12 allows to estimate the
maximum distance to the SMBH, Eq. 2.14. The estimates of these distances mostly range in
the order of a few pc for most sources reported by Blustin et al. (2005).
Vout =
√
2GMSMBH
R
→ R ≥
2GM
V 2
out
(2.13)
R ≤
LionCv(R)
ξNH
(2.14)
Furthermore, Laha et al. (2014, 2016) provides an updated review of these ionized re-
gions, finding warm absorbers in ∼ 65% of their sample. The work of Laha et al. (2021)
2.3. X-RAY ANALYSIS, TESTING THE UNIFICATION SCHEME 47
Figure 2.9: Laha et al. (2021) establishes the parameters of the warm absorber as outflows of
velocities that range between 100-2000 km/s and places the material in a scale from 0.1 pc -
1kpc from the central engine.
reports the general characteristics of warm absorbers: a range of ionization parameters of
log ξ = −1 to 3 erg cms−1, column densities of log NH = 21 to 22.5 cm−2 and outflow velocity
ranging from 100-2000 km/s. The suggested scale for the location of the ionized material ranges
from 0.1 pc up to 1 kpc from the central engine. Fig. 2.9 presents a brief description of different
ionized regions found in AGN, responsible for different contributions to the X-ray spectrum,
and distinguishing them from one another according to their measured parameters.
2.3.3 Hypothesis: Is X-ray absorption related to the polarization
regions?
Through X-ray analysis, our goal is to compare the regions responsible for the detected po-
larization to the presence of X-ray absorption. The proposed absorption test is based on the
premise stated by the polarization unified model: both polar and equatorial scattering regions
are present in all Seyfert sources, with the inclination towards the line of sight determining
which one dominates the optical polarization. In turn, these scattering regions would show as
absorption in the X-ray band, as shown in Fig. 2.10. Relating the scattering regions and the
four polarization scenarios described in Sect. 2.2.1 to X-ray absorption poses the following cases:
1. Null polarization i ≈ 0◦ - The line-of-sight looks directly into the central region of the
source, which would appear completely unabsorbed in the X-ray.
2. Equatorial polarization 0◦ < i < 45◦ - EQ-pol Sy 1. The line of sight passes through the
NLR. We expect this type of source to be unabsorbed or, if absorption is present, it would
be by ionized material, as is reported by Reynolds (1997); Laha et al. (2014, 2016).
48 CHAPTER 2. OPTICAL POLARIZATION IN SEYFERT GALAXIES
3. Polar polarization where i ∼ 45◦ - PL-pol Sy 1. The line of sight passes through the upper
layers of the torus. These layers have lower optical depth and get ionized by the central
region, showing as warm absorption in the X-ray regime. An example of this proposed
geometry is shown in Fig. 2.11, corresponding to the spectrum of ESO 323-G077 in the
left panel of Fig.2.6. This geometry proposed by Jiménez-Bailón et al. (2008) describes
how a PL-pol Sy 1 source is seen passing through the polar scatterer which dominates the
optical polarization and shows as a “soft scattered power law” in the X-ray spectrum.
4. Polar polarization where i > 45◦ - PL-pol Sy 2. The line-of-sight passes through the
optically deep obscuring torus, showing as cold absorption in the X-ray.
Figure 2.10: The proposed test to the unification scheme by Smith establishes a comparison
between the measured polarization and the presence of absorption in the X-ray spectrum.
Figure 2.11: The geometry of ESO 323-G077, a polar-polarized source. Line of sight passes
through the material responsible for the polar-polarization in the optical band, showing as
absorption in the soft X-ray band. Credit: (Miniutti et al. 2014).
2.4. OBJECTIVES 49
2.4 Objectives
The objective of this study is to conduct an independent test of the unification scheme pro-
posed for polarized Seyfert galaxies by analyzing their X-ray spectra. Specifically, we aim to
investigate the presence of absorption in these sources, seeking to establish connections between
the absorbing material to the gas and dust regions responsible for the detected polarization.
The spectral analysis focuses on examining the general characteristics of our sample based
on the main components of a typical X-ray spectrum of an AGN, as described in Sect. 1.3.3.
These components provide a base model, which is used to test for the presence of absorption.
Our goal is to classify each source as being affected by either a cold or warm absorber or as
unabsorbed. Results will be systematically organized based on polarization classifications, en-
abling a comparison between polarization sub-samples. This analysis is based on the hypothesis
illustrated in Fig. 2.10 as a framework for the relationship between X-ray absorption properties
and the AGN orientation scheme.
Chapter 3
X-ray analysis of optically polarized
Seyfert galaxies
In Ch. 2, we presented the polarization study and the unification scheme based on their po-
larization position angle. In this chapter, we present the X-ray data available from 25 sources
out of the original 46 sources in the study by Smith et. al. We provide a description of the
XMM-Newton satellite and the EPIC-pn camera, the instrument from which we obtained our
data. We then outline the data reduction and analysis, from the extraction of spectral products
to spectral fitting. This chapter includes a detailed description of the models used to fit the
spectra. The results will be reported in Ch. 4.
3.1 Sample selection of polarized Seyfert
Smith et al. reported 23 Sy 1 galaxies with intrinsic polarization: 11 classified as EQ-pol and
12 as PL-pol (see Fig. 2.10). From these 23 sources, we obtained all public EPIC-pn data for
11 PL-pol and 8 EQ-pol galaxies from the XMM-Newton Science Archive1. Additionally, we
included 6 sources with weak polarization, which Smith described as “null-polarization candi-
dates” (see Sect. 3.1.1). In total, our X-ray sample consists of 25 sources, organized according
to the polarization position angle: PL-pol, EQ-pol and Weak pol., as presented in Table 3.1.
The coordinates (RA, Dec) in J2000 and the redshift of each source are taken from the
NASA/IPAC Extragalactic Database (NED)2. We indicate their classification according to
their optical spectra, explained in Sect. 1.4.1.1, which was also obtained from NED. The col-
umn density NGal
H corresponds to gas in the Galaxy distributed along the line of sight of each
source, taken from the High Energy Astrophysics Science Archive Center (HEASARC) server3,
(Kalberla et al. 2005). For each observation, we provide the XMM-Newton Science Archive
observation ID number and the exposure time. In cases where multiple observations were avail-
able, we selected the one with the longest available exposure time.
1XMM-Newton archive: https://nxsa.esac.esa.int/nxsa-web/#home
2NED Database: http://ned.ipac.caltech.edu
3HEASARC server: https://heasarc.gsfc.nasa.gov/cgi-bin/Tools/w3nh/w3nh.pl
50
3.1. SAMPLE SELECTION OF POLARIZED SEYFERT 51
(1) (2) (3) (4) (5) (6) (7) (8)
Object RA Dec z NGal
H Type Obs ID Exp.
[×1020 cm−2] [ks]
Polar Polarization
Mrk 1218 129.546 +24.8953 0.02862 3.1 Sy 1.8 0302260201 13.9
Mrk 704 139.608 +16.3053 0.02923 2.7 Sy 1.2 0502091601 98.2
Mrk 1239 148.080 -1.6121 0.01993 4.1 NLSy1 0891070101 105.0
NGC 3227 155.877 +19.8651 0.00386 1.9 Sy 1.5 0782520601 107.9
WAS 45 181.181 +31.1772 0.02500 1.4 Sy 1.9 0601780601 39.5
Mrk 766 184.611 +29.8129 0.01293 1.9 NLSy1/Sy 1.5 0109141301 129.9
Mrk 231 194.059 +56.8737 0.04217 0.9 Sy 1 0770580501 26.5
NGC 4593 189.914 -5.3443 0.00831 1.7 Sy 1 0784740101 142.1
ESO 323-G077 196.609 -40.4147 0.01501 7.7 Sy 1.2 0694170101 132.6
IRAS 15091-2107 227.999 -21.3171 0.04461 7.9 NLSy1 0300240201 23.0
Fairall 51 281.225 -62.3648 0.01418 6.3 Sy 1 0300240401 26.9
Equatorial Polarization
I Zw1 13.396 +12.6934 0.06117 4.6 NLSy1 0743050301 141.2
Akn 120 79.048 -0.1498 0.03271 9.9 Sy 1 0721600401 133.3
NGC 3783 174.757 -37.7387 0.00973 0.1 Sy 1.5 0112210201 137.8
Mrk 841 226.005 +10.4378 0.03642 2.0 Sy 1.5 0882130401 132.0
Mrk 876 243.488 +65.7193 0.12109 2.4 Sy 1 0102040601 12.8
KUV 18217+6419 275.489 +64.3434 0.29705 3.5 Sy 1.2 0506210101 14.3
Mrk 509 311.041 -10.7235 0.03440 3.9 Sy 1.5 0306090201 85.9
Mrk 304 334.301 +14.2391 0.06576 4.9 Sy 1 0103660301 47.3
Weak polarization
Mrk 896 311.587 -2.813 0.02642 3.1 Sy 1 0112600501 11.2
Mrk 926 346.181 -8.686 0.04702 2.9 Sy 1.5 0790640101 59.0
NGC 7213 332.318 -47.167 0.00584 1.1 Sy 1.5 0605800301 132.5
NGC 7469 345.82 +8.874 0.01627 4.5 Sy 1.2 0760350801 101.6
NGC 7603 349.736 +0.244 0.02876 3.4 Sy 1.5 0305600601 16.8
PG 1211+143 183.574 +14.054 0.0809 2.7 Sy 1 0745110201 104.0
Table 3.1: X-ray data of polarized Seyfert 1 galaxies. The table is organized according to
the polarization characteristics of the sources. Columns 2 and 3 report coordinates (RA, Dec)
in J2000, column 4 indicates the redshift, column 5 corresponds to the column density of the
Galaxy from HEASARC, column 6 indicates AGN classification taken from NED. Columns 7
and 8 are the XMM-Newton observation ID and the exposure time.
52 CHAPTER 3. X-RAY ANALYSIS OF OPTICALLY POLARIZED SEYFERTS
3.1.1 Weak polarization
Within the sources studied by Smith et al., several are reported to be significantly contaminated
by interstellar polarization, or data are too noisy to determine the polarization characteristics.
Particularly, seven sources are deemed “intrinsically weakly polarized”, with an average polar-
ization of p ≤ 0.3%, while the data has an uncertainty of 0.1 − 0.3%. Smith et al. argue that
these sources could be consistent with the null polarization scenario, i.e., i ≈ 0◦. However,
Smith et al. also mentioned the possibility of very different scattering column densities in both
the equatorial and polar regions and from source to source, leading to a wide range of varia-
tion in the degree of polarization, p. In their own words: “... an independent constraint on
orientation is needed to discriminate between such effects and pole-on orientation as the cause
of their low polarization”, Smith et al. (2004).
Given their status as “null-polarization candidates”, we decided to include these sources
with weak polarization in our X-ray analysis and use them as a comparison group for the other
two sub-samples with well-defined polarization. Thus, we can compare three different polariza-
tion scenarios using the results of our X-ray analysis.
3.2 XMM-Newton
The X-ray Multi-Mirror Mission (XMM-Newton) was developed by the European Space Agency
(ESA) and launched on December 10th, 1999. It is one of the main missions of Horizon2000, a
long-term program by ESA that began in 1983 with the purpose of space exploration (Jansen
et al. 2001). XMM-Newton follows a highly elliptical orbit with an inclination of 40 degrees, a
perigee of ∼ 6,000, and an apogee of about ∼ 115,000 km. This orbital design minimizes the
impact of radiation from Earth on the instruments of the spacecraft and enables longer exposure
times, expanding its observational capabilities. The payload of XMM-Newton consists of three
co-aligned Wolter telescopes of 7.5 m focal length each. Invented by Hans Wolter in 1952, this
type of telescope consists of nested mirrors using grazing-incidence optics so the incident X-rays
graze the mirrors at shallow angles and get reflected towards the detector. The left image of
Fig. 3.1 shows the mirrors in parabolic and hyperbolic arrays so the X-rays can be reflected
at angles of ≤ 1 degree. At the end of the focal length is the Focal Plane Assembly (FPA),
which contains the focal-plane instruments: two Reflection Grating Spectrometers (RGS) and
three imaging cameras, the EPIC-pn, and EPIC-MOS 1 and 2 (Barré et al. 1999; Lumb et al.
2012). The configuration of XMM-Newton and the location of the instruments are depicted in
the right image in Fig. 3.1.
3.2.1 The European Photon Imaging Camera (EPIC)
The European Photon Imaging Camera (EPIC) consists of three different X-ray CCD cameras
located behind the X-ray telescopes. The EPIC cameras operate in photon counting mode,
providing imaging over a field of view (FOV) of 30 arcmin, detecting high-energy photons
nominally ranging from 0.15 up to 15 keV with a spectral resolution of E/∆E ∼ 20 − 50 and
an angular resolution of 6 arcsec. Two of the cameras are equipped with metal oxide semi-
conductor (MOS) CCDs. The third camera, EPIC-pn, consists of a single chip of very high
resistivity, allowing for efficient detection of over 90% of incoming X-rays between 0.5 and 10
keV, (XMM-Newton Users Handbook 2024).
3.2. XMM-NEWTON 53
Figure 3.1: Schematic of the observatory system of XMM-Newton. The image on the left
shows the parabolic and hyperbolic array of the Wolter telescopes on board the spacecraft.
Credit: ESA/AOES Medialab. The image on the right is a schematic of the XMM-Newton
observatory system, with the Mirror Support Platform (MSP) at one end and the FPA at
the other, carrying the observing instruments, including the EPIC-pn camera, marked in blue.
The Service Module (SVM) carries the spacecraft subsystems such as the sun shields and the
antennae. Credit: Lumb et al. (2012).
Figure 3.2: The EPIC-pn CCD focal plane. The left image is a simple schematic of the 12 CCD
array of the camera. Credit: XMM-Newton Users Handbook. The right image corresponds to
the EPIC-pn Full Frame mode exposure of Was 45, a Seyfert source part of our sample encircled
in green.
This study is carried out with the CCD data obtained by EPIC-pn, the resolution and effi-
ciency over its range of energy, 0.5−10 keV, (Strüder et al. 2001; Turner et al. 2001). A scheme
of the CCD configuration of EPIC-pn is shown in Fig. 3.2. The left image is a simple sketch
of an array of 12 CCDs that comprise the EPIC-pn configuration, with a 30 arcmin diameter.
The right image corresponds to the EPIC-pn exposure of Was 45, one of the sources in our
sample.
54 CHAPTER 3. X-RAY ANALYSIS OF OPTICALLY POLARIZED SEYFERTS
Figure 3.3: Example of the exposure of Fairall 51, showing OoT in the left image. The right
panel corresponds to the exposure after OoT have been subtracted.
3.3 Processing the XMM-Newton data
1. Out of Time events.
Out-of-Time (OoT) events occur due to the registering of photons during the readout of
the CCD, alongside the actual integration interval, affecting the region of the source and
the region between the source and the readout node. Their effect broadens the spectral
features and appears as a strip of wrongly reconstructed events. An example of this strip
present in an observation and the same observation after the OoT correction is shown in
Fig. 3.3. To correct for these OoT events, a two-step process is employed: first, an OoT
event list is generated using the epchain tool, which simulates the behavior of these out-
of-time photons. This process assigns new positions along the RAWY (vertical position
of detected X-ray events within the CCD) axis based on a random shift and applies the
necessary corrections for charge transfer inefficiency. Next, the OoT event list is used to
create a clean image or spectrum by subtracting the expected contribution of OoT events
from the observation data4.
2. Pile-up.
Pile-up happens when multiple X-ray photons reach the detector within a short amount
of time, particularly in sources with high count rates, hitting the same pixel or adjacent
pixels and getting registered as a single event. This causes the hardening of the spectrum
as two lower energy photons get registered as a single photon of higher energy. Before
proceeding with the pile-up correction, the amount of pile-up must be evaluated. The
SAS task epatplot produces a plot that depicts the observed versus expected pattern dis-
tribution as well as the observed-to-model fractions for single and double events within a
certain energy range. An example of a source with a slight but non-negligible pile-up is
shown in Fig.3.4. The graph shows the expected pattern of single and double events in
a solid line, with the observed pattern superimposed. A single and double ratio of 1.0 (1
sigma errors are given) indicates the absence of pile-up, in comparison with the s< 1.0
and d> 1.0 shown in the graph (lines red and blue respectively). Once determined that
pile-up is present, the epproc and epchain tasks offer the option of pile-up correction using
different methods. These tasks identify regions of the CCD where pile-up is likely to occur
based on the count rate and the spatial distribution of events. The correction can involve
excluding pile-up pixels or quantifying pile-up by analyzing the shape of the count rate
4OoT correction thread: https://www.cosmos.esa.int/web/xmm-newton/sas-thread-epic-oot
3.3. PROCESSING THE XMM-NEWTON DATA 55
Figure 3.4: The plot produced by the epatplot task, to evaluate the presence of pile-up corre-
sponding to the source ESO 323-G 077. The solid line corresponds to the expected distribution
pattern for the different events, with the observed distribution superimposed. The single and
double events ratio indicate pile-up is present in this observation.
distribution and determining the fraction of events that have been affected5. The pile-up
found in our sample is minimum, as we can see in the example of ESO 323-G077.
3. EPIC background.
In general, there can be two sources of X-ray background: the cosmic X-ray background
(CXB) and the instrumental background6. The latter includes the detector noise, only
relevant at low energies, ∼ 200 eV, and the interaction with high energy particles that
introduces two sources of background. The first is due to soft protons, Ep < 100 keV, that
give rise to flaring periods in the light curves > 10 keV when photons are focused on the
CCD. There is also background due to particles of E> 100 MeV interacting with the metal
components in the satellite, characterized by a flat spectrum that contains lines due to
fluorescence from the metals in the structure. The Science Operation Center (SOC) has
analyzed these background components and developed the corresponding tasks to filter
it and produce the clean events files used for the extraction of the spectrum7. We show
an example of the background spectrum in Fig. 3.5, where the extracted background is
plotted in red superimposed to the spectrum of the source, in black.
5Pile-up correction thread: https://www.cosmos.esa.int/web/xmm-newton/sas-thread-epatplot
6Background Components: https://www.cosmos.esa.int/web/xmm-newton/epic-background-components
7Filtering events file for flaring particles: https://www.cosmos.esa.int/web/xmm-newton/sas-thread-epic-
filterbackground
56 CHAPTER 3. X-RAY ANALYSIS OF OPTICALLY POLARIZED SEYFERTS
Figure 3.5: The 0.5-10 keV spectrum of Was 45 is shown in black. Data in red corresponds to
the background spectrum.
3.3.1 Data processing
The public data from the XMM-Newton mission are available on the XMM-Newton Science
Archive website. We obtain the Calibration Files, containing calibration data and parameters
necessary for processing the raw data which allow us to obtain a calibrated event file that
includes X-ray detected events such as energy, arrival time, and position. The data of each
source corresponds to the Observation Data Files (ODF). The ODF files contain the scientific
raw, unprocessed data from each observation, i.e. the events list, as well as data containing
the telemetry that provides information on the status of the satellite, instrument settings, and
pointing directions during the observation. From each ODF, we obtain the event list, which
includes X-ray detected events such as energy, arrival time, and position.
We processed the ODF files and used the most updated Current Calibration Files (CCF),
at the moment of downloading the observations, to produce calibrated event lists from which
to extract our spectral products. We used the Science Analysis System (SAS) software, SAS
v.19.1 (Gabriel 2017), that contains the series of tasks for processing and extracting the im-
ages and spectra corresponding to each source, following a standard procedure detailed by the
XMM-Newton Science Operations Center (SOC).
The EPIC-pn data were reprocessed with the epproc task and were filtered for high back-
ground events, producing the clean events file. Using the SAO Image DS9 display, (Joye 2006),
we selected a circular region that encloses the source, centered at the peak of X-ray emission.
The source radius varied from 30 up to 68 arcsec, depending on the target and the observation
mode of the camera, Full Frame, Large Window, or Small Window. For cases where the ob-
servation was taken in the Full Frame or Extended Full Frame mode and the extraction region
radius is around 30 arcseconds. For observations in the Small Window mode, the source is cen-
tered within the smaller frame, typically resulting in a proportionally larger extraction region,
an example of this frame is shown in the top image of Fig. 3.6. This explains the significant
differences in the sizes of the extraction regions across different observation modes. Using the
evselect task, we can extract the spectrum that corresponds to the source region, whose center
coordinates and radius were selected by opening the image in DS9. Using the same procedure,
we selected the background region defined by a circular region with no contribution from any
source in the CCD, preferably selected on the same chip of the source to minimize background
fluctuations. The radius of the background regions varies from 50 up to 96 arcsec, depending
on the target. An example of the region selection for the source Was 45 is shown in Fig. 3.6.
With the tasks rmfgen and arfgen, we generate the response matrices: Ancillary Response Files
3.4. SPECTRAL MODELS 57
Figure 3.6: Top image is the EPIC-pn observation of Mrk 704 in Small Window mode, already
corrected for OoT, with a source extraction region of ∼ 45 arcsec. The bottom images corre-
spond to Was 45 shown in Fig. 3.2, in Full Frame mode. This close up shows two green circles,
the smaller one that corresponds to the source with a size of ∼ 40 arcsec and the bigger one,
∼ 85 arcsec, that was selected as the background.
(ARF) provide information about the effective area of the telescope and instrument response,
and Response Matrix Files (RMF) that describe the energy response of the instrument. The
resulting source spectra were binned with the task specgroup in order to obtain a minimum of
25 counts per energy bin, allowing us to use the χ2 statistics. As an example, the extracted
0.5-10 keV spectrum of Was 45 is shown in Fig. 3.7.
We processed and applied the corresponding corrections to the data and extracted the pn
spectrum of our 25 sources. We can now move on to the analysis, which we describe in the
following sections of this chapter.
3.4 Spectral Models
The spectral analysis is structured to start by fitting the main spectral components that are
common to the vast majority of AGN. The aim is to obtain a baseline model and report on the
general characteristics of each source. Once the baseline model is well established, we conduct
the absorption test to determine whether the source is unabsorbed, affected by cold absorption,
or influenced by warm absorption. The models employed in our analysis are available in the
Xspec Spectral Fitting Package v 12.10.1n. We provide a detailed description of these models
in this section. Each subsection begins with the XSPEC name of the model in italics.
58 CHAPTER 3. X-RAY ANALYSIS OF OPTICALLY POLARIZED SEYFERTS
Figure 3.7: The 0.5-10 keV extracted spectrum of the EPIC-pn observation of Was 45 with an
exposure time of 39.5 ks.
3.4.1 Galactic neutral absorption
TBabs. To model the absorption from our Galaxy, we chose the Tuebingen-Boulder Interstellar
Medium (ISM) absorption model. This model accounts for the photoelectric absorption by
neutral ISM gas from our Galaxy along the line of sight, with abundances from Wilms et al.
(2000). We used Galactic Hydrogen column values, NH , from the “Column Density Tool” from
the HEASARC server with measurements from the Bekhti et al. (2016) survey. These values
are reported in Table 3.1 for each source, which are kept fixed throughout the fitting process.
3.4.2 Power law continuum
zpowerlw. The primary emission is modeled as a power law continuum characterized by the
photon index Γ, which indicates the spectral slope. The model is defined by the formula:
A(E) = K[E(1 = z)]−α; where α is the photon index and K is the normalization of the model
in units of [photons/keV/cm2/s] at 1 keV. These are left as free parameters. In Figures 3.8,
3.17, 3.10 and 3.11, this component is plotted with Γ = 1.8, a typical value for Sy 1 sources,
(Corral et al. 2011; Singh et al. 2011).
3.4.3 Fe emission lines
zgauss. As reported in Sect. 1.3.3, the most prominent emission line in AGN X-ray spectra is
the Fe Kα at 6.4 keV. This feature is modeled using a Gaussian line profile with the line energy
fixed at 6.4 keV and its width set to 0.05 keV. The normalization, corresponding to the line
flux, is left as a free parameter and is expressed in units of photons/cm2/s. To further improve
the fit, we also tested for the presence of emission lines from ionized iron: Fe xxv at 6.697 keV
and Fe xxvi at 6.966 keV, (Bianchi et al. 2005). These additional lines were also modeled with
the Gaussian line profile, fixed energies at their rest-frame values, and a line width of 0.05 keV.
3.4.4 Soft Excess
This excess emission above the extrapolated power-law component is thought to arise from
either the inner accretion disk or reprocessed radiation, (Sobolewska & Done 2007; Done et al.
2007). Given that the nature of the soft excess emission is still debated, we considered five
different models to fit this feature.
3.4. SPECTRAL MODELS 59
1. Black body emission: zbbody. This model represents a redshifted blackbody spectrum,
commonly used to approximate thermal emission in AGN spectra, (Corral et al. 2011;
Page et al. 2003; Piconcelli et al. 2005; Krongold et al. 2003). It can account for thermal
emission from the accretion disk, which may contribute to the observed soft excess in X-
rays. This model yields the black body temperature in units of [keV]. The normalization
of this model is related to the size and the distance of the emitting region. The first panel
of Fig.3.8 depicts the contribution of a blackbody spectrum of temperature of 0.1 keV
added to a power law with a slope of Γ = 1.8. The red line in the figure corresponds to
the contribution of the blackbody model.
2. Comptonization: compbb. This model describes a blackbody spectrum that undergoes
inverse Compton scattering, (Nishimura et al. 1986; Noda et al. 2013). The Comptoniza-
tion parameter accounts for the product of the electron temperature and the optical depth
of the plasma. In this case, the electron temperature is fixed at 50 keV. The model yields
two key parameters: the blackbody temperature (kT, in keV) and the optical depth of the
plasma (τ). A lower optical depth implies fewer scattering events, resulting in a spectrum
that is closer to the original blackbody emission, while a higher optical depth leads to
more significant Compton scattering and a harder spectrum. The contribution of adding
a Comptonized blackbody to a power law, Γ = 1.8, is depicted in the second panel of Fig.
3.8.
3. Circumnuclear ionized gas: apec. Astrophysical Plasma Emission Code. This model
represents the emission spectrum from collisionally ionized gas. It can be interpreted as
heated gas in the circumnuclear material around AGNs, contributing to the soft excess
emission via thermal bremsstrahlung and line emission from ionized metals (e.g. Buhari-
walla et al. 2023). The model yields a temperature for the plasma and typically assumes
solar abundances for the metals. It is based on the AtomDB atomic database, which
contains atomic data for collisional and radiative processes for elements ranging from hy-
drogen to zinc. The spectrum corresponding to this model is depicted in the third panel
of Fig. 3.8.
4. A second power law: zpowerlw. With this model, we introduce a second power law
component with a different photon index from the one fitted to the hard band. The full
continuum is thus modeled by two power-law components with different slopes, where the
second photon index typically results in a steeper power law. An example of a second
power law modeling the soft excess is shown in the bottom image of Fig. 3.8.
5. Ionized reflection: relxill. We explore the possibility of the soft excess being due
to Compton reflection of the power law onto the accretion disk. This model describes
reflection from the ionized disk, allowing for the inclusion of relativistic effects due to the
strong gravitational field near the black hole, (Garćıa et al. 2014). The model combines
the xillver reflection code, (Garćıa & Kallman 2010; Garćıa et al. 2011, 2013), with the
relline ray tracing code, (Dauser et al. 2010, 2013). As a result, relxill can effectively
reproduce the soft excess observed below 2 keV by modeling the reflected emission that
arises from the disk as it is illuminated by the X-ray corona. The power law emission
accounted for in the relxill model is linked to the power law component already fitted
to the spectra. We fixed the reflection fraction to -1.0, indicating that the model will
only account for the reflected emission from the accretion disk. The free parameter is
the ionization parameter of the accretion disk, responsible for the contribution at lower
energies. Fig. 3.17 depicts the main power law component (black dotted line) with the
relxill model added (red line). The resulting model is depicted with a black solid line.
60 CHAPTER 3. X-RAY ANALYSIS OF OPTICALLY POLARIZED SEYFERTS
Figure 3.8: Comparison of Soft excess theoretical models. For all the graphs, the black solid
line corresponds to the main continuum zpowerlw, of Γ = 1.8 (black dotted line), plus the soft
excess component. The different models in red represent the soft excess component as indicated
at the top of each graph. For the first three models, the temperature has been set to 1 keV,
with the compbb modeled for an optical depth of 1. For the last model, the second power law
component has been set to Γ = 2.5.
3.4. SPECTRAL MODELS 61
Figure 3.9: Relxill theoretical model. The contribution of ionized reflection modeled with the
relxill, depicted with a red line. The model is added to a power law, depicted with a black
dotted line. The resulting model is shown with a black solid line.
3.4.5 Absorption
As described in Sect 2.3, absorption is characterized by the physical conditions of the absorbing
gas, mainly its column density and ionization state, thus we distinguish two main conditions:
neutral and ionized absorbers. XSPEC provides many options to fit both of these types of
absorption, in particular, we selected two different models for each.
3.4.5.1 Cold Absorption
1. zphabs. This model represents photoelectric absorption by neutral material at a specified
redshift, fixed to the redshift of the source in this particular case, effectively simulating
the effects of neutral absorption intrinsic to the AGN. It assumes solar abundances and
yields the Hydrogen column, NH in units of [1022 cm−2]. The effect of neutral absorption
on a power law continuum is depicted in the top panel of Fig. 3.10.
2. zTBabs. This is an updated version of the Tuebingen-Boulder photoelectric absorption
model that places the absorbing material at the redshift of the source. This model uses
a cross-section obtained from X-ray absorption, considering three contributions to the
absorption by three different forms of the ISM: dust/grains, atomic gas, and molecular
gas. For the gas-phase, it sums the photoionization cross sections of the different elements,
for the molecule-phase it considers the cross section of the Hydrogen molecule and, lastly,
it takes into consideration shielding by the grain-phase, although the contribution is
very small. From this model, we will be reporting the Hydrogen column as well, NH in
units of [1022 cm−2]. Fig. 3.11 depicts a comparison between the previous model and
this one, depicted in the bottom panel. In both cases, the neutral absorber is set to
NH = 1022 atoms cm−2.
62 CHAPTER 3. X-RAY ANALYSIS OF OPTICALLY POLARIZED SEYFERTS
Figure 3.10: Cold absorption models. The depicted model corresponds to a power law of
Γ = 1.8 affected by the cold absorber by a neutral absorber of NH = 1022 cm−2. The main
difference between these two models is the cross-section of the photoelectric absorption, where
zTBabs includes three components of the ISM.
3.4.5.2 Warm Absorption
1. absori. This model, developed by Magdziarz & Zdziarski (1995) and based on the work
of Done et al. (1992), describes the absorption of X-rays by ionized gas. The model uses
ionization and recombination rates to simulate the effects of ionized absorption on the
X-ray spectrum. The temperature of the absorber is fixed at 105 K and the Fe abundance
is set to solar values. The model provides two main parameters: the ionization parameter
that indicates the degree of ionization of the gas, and the column density in units of
10−22 cm−2. The top panel of Fig. 3.11 corresponds to the absori model multiplied by a
power law model of Γ = 1.8.
2. zxipcf. Reeves et al. (2008). This model is described by four parameters: the redshift of
the source, the column density in units of 1022 cm−2, the covering fraction, f , and the
ionization parameter, log ξ. The ionization parameter is defined: ξ = L/nr2; where L is
the X-ray luminosity of the source, n is the electronic density of the ionized gas and r2
is the squared distance of the ionized cloud to the source of ionizing radiation. In our
analysis, the covering factor f is fixed to 1.0, where 1 − f represents the portion of the
source that is seen directly. The bottom panel of Fig. 3.11 corresponds to the power law
affected by zxipcf.
3.5. METHODOLOGY OF THE X-RAY SPECTRAL ANALYSIS 63
Figure 3.11: Models for warm absorption. Both figures show a power law of Γ = 1.8 affected by
a warm absorber of NH = 1022 cm−2 and an ionization parameter of log ξ = 2.5. The top panel
shows the effect of the absori model and the bottom panel corresponds to the zxipcf model, with
a covering factor set to 1.0. The main difference is that, while absori models a more uniform
absorber with fixed temperature, zxipcf provides a more detailed approach.
3.5 Methodology of the X-ray Spectral Analysis
To characterize the X-ray emission using the models listed in the previous section, we developed
a systematic three-step analysis. This approach builds a nested model, starting with a basic
fit and progressively adding components to account for additional spectral features. Such a
structure allows us to evaluate whether each newly added component statistically improves the
overall fit (see Sect. 3.6). We also consider the value of the χ2
ν = χ2/DOF , as an indicator
of the goodness of fit, where an ideal case of χ2
ν = 1.0 indicates that the spectrum is well
approximated by the chosen model within the statistical errors. The analysis is divided into
three main steps to sequentially add each of the components used for fitting our sample.
Step 1: Hard-band (2-10 keV). We start with the hard band. As discussed in Sect.
1.3.3, absorption effects are generally not expected at energies > 2 keV. Our first step is to fit
the main continuum component and determine its photon index. Next, we search for emission
lines, which can often be identified by examining the residuals from the power-law fit. In par-
ticular, we test for the presence of emission lines in the Fe K band at fixed rest frame energies
and widths. Lastly, we tested for possible effects of neutral absorption which, as mentioned,
would not significantly affect the spectral shape at these energies, unless in the case of high
column densities. The model resulting from this step serves as our “baseline model”.
64 CHAPTER 3. X-RAY ANALYSIS OF OPTICALLY POLARIZED SEYFERTS
Step 2: Full energy band (0.5-10 keV). Our analysis, particularly the absorption test,
requires extending the energy band to include the soft X-rays (0.5-2 keV). The next step is
to fit the “baseline model” to this energy range. We use this fit to verify the presence of soft
excess by testing it against models 1, 2, 3, and 4 in Sect. 3.4.4. The inclusion of the ionized
reflection is also part of this second step, with the relxill working as an additional contribution
to the soft excess, fitted in conjunction with the previously tested models. We will be testing
the presence of absorption using the resulting model from this step .
Step 3: The Absorption Test. Having fitted the continuum (power law + soft excess),
we can now move on to test the absorption effects. We separately fit the neutral and ionized
absorbers to the existing model, obtained from Step 2. For the corresponding statistical test,
see Sect. 3.6, indicates if the spectrum is affected by absorption and, if so, which model is
preferred by the data. We will classify each source into one of three categories: (1) unabsorbed,
(2) affected by cold absorption, or (3) affected by warm absorption.
3.6 Statistical tests
As we build a nested model, we must determine whether each new component significantly
improves the fit of the data. The different tested components are either multiplicative or
additive, this determines which statistical test to use to determine if the additional component
is significant. The final model, after all the steps enlisted in Sec.3.5, is of the form:
GalacticAbsorption*IntrinsicAbsorption*(PowerLaw+EmissionLines+SoftExcess)
We used the F-test for additive components and the AIC for multiplicative components. In
our analysis only absorption components are multiplicative.
3.6.1 F-test
The F-test is a statistical test used to compare the fit statistics obtained by applying different
models to the same data set. The test assesses whether the newly added component results
in an improved fit. It uses the ratio of the χ2 and the degrees of freedom (D.O.F) from the
models that are to be compared; where the subindex 2 indicates the newer, more complex
model: (χ2/D.O.F )1
(χ2/D.O.F )2
. The significance of the F-test result is assessed by calculating the p-value,
the probability of observing the F-test statistic that is as large, or larger than the one obtained
from the data if the simpler model is the correct model. Thus, a p-value< 0.05 indicates that
the more complex model significantly improves the fit, resulting in an F-test of > 95%, (Pro-
tassov et al. 2002).
3.6.2 Akaike Information Criterion
For multiplicative components, such as the absorption models, we used the Akaike Information
Criterion (AIC, Akaike 1974). We first calculate the AIC value, Eq. 3.1, where k corresponds
to the number of free parameters in each model. The comparison between both models is given
by the factor, FAIC, resulting from Eq. 3.2, where AICNEW refers to the model with the new
component, and AIC0 to the previous model. The inverse, 1/FAIC, indicates the improvement
3.7. SHERPA & PYXSPEC: FITTING PACKAGES 65
of the new and more complex model, (e.g. Krongold et al. 2021).
AIC = 2k + χ2 (3.1)
FAIC = exp((AICNEW − AIC0)/2) (3.2)
As mentioned, the statistical testing allows us to determine whether the contribution of
each component is of statistical significance to the general model, and so, the results presented
in Ch. 4 will include the results of both the F-test and the AIC.
3.7 Sherpa & PyXspec: fitting packages
For the fitting process, we employed two Python-based tools for X-ray analysis: Sherpa and
PyXspec, both of which provide access to the extensive library of spectral models available in
XSPEC. Sherpa is a command-line modeling and fitting application included in the Chandra
Interactive Analysis of Observations (CIAO) software, developed by the Chandra X-ray Center,
(Doe et al. 2007; Refsdal et al. 2009, 2011). PyXspec serves as a Python interface to the XSPEC
spectral fitting package, allowing for a more streamlined and automated approach to spectral
analysis. The script was developed in a Python v 3.7.4. environment with the PyXspec library,
(Gordon & Arnaud 2021). The flowchart in Fig. 3.12 shows the fitting process, following the
three main steps previously described.
3.7.1 Sherpa
For our initial approach, we utilized Sherpa, fitting each source interactively. This hands-on
process allowed for a better understanding of the fitting process and visualized the contribution
that each component had to the overall model. This approach provided a solid foundation for
understanding the model-building process and refining the methodology for subsequent analy-
ses.
At this stage, in the hard band, we only tested for the presence of Fe Kα after fitting the
power law continuum. It was also during this stage that we tested the first three soft excess
models listed in Sect. 3.4.4: zbbody, compbb, and apec. We detail an example of this first
testing for the case of Mrk 704, reporting the resulting parameters in Table 3.2. First, we could
conclude that the apec model does not result in a reasonable fit for a type 1 Sy: the slope of the
power law is steeper and the emission gas temperature is too high to consider it a contribution
to the soft X-ray emission. The zbbody and compbb models yield similar model parameters, with
the power law index and the soft excess temperature resulting in very close values. Moreover,
the optical depth of compbb results in 2.5× 10−9, such a low value indicates that this value is
negligible and that the spectrum resembles that of a blackbody. For these reasons, we conclude
that the zbbody model is better fitted for the case of Mrk 704. However, this fit results in a
χ2
ν = 2.8, a high value that indicates this model still needs to account for further aspects of the
spectrum.
66 CHAPTER 3. X-RAY ANALYSIS OF OPTICALLY POLARIZED SEYFERTS
Hard band:
phabs*zpowerlw
Fitting Process
+zgauss 
phabs*(zpowerlw+zgauss)
phabs*(zpowerlw+zgauss+zgauss)
phabs*(zpowerlw+zgauss+zgauss+zgauss)
phabs*(zpowerlw)
Testing for cold absorption
*zphabs
phabs*zphabs*(zpowerlw+zgauss)
phabs*(zpowerlw+zgauss)
Soft Excess Testing
+zbbody +compbb+apec
Baseline
Baseline + zbbody
Baseline
Baseline + apec
Baseline
Baseline + compbb
Absorption Test 
*zphabs o
*zTBabs 
*absori o
*zxipcf
ColdAbs*Baseline+SoftExcess
Source not affected by cold absorption
WarmAbs*Baseline+SoftExcess
Source not affected by warm absorption
Preferred Model:
Source is Unabsorbed
Source affected by cold absorption
Source affected by warm absorption
Baseline Model
Extending the baseline model to 
0.5 - 10 keV
+zpowerlw
Baseline
Baseline + zpowerlw
One Soft Excess
Figure 3.12: Flow chart for the fitting process. Following the three steps from Sect.3.5, the
first section corresponds to the hard band (in purple), fitting the main continuum and testing
for emission lines and possible neutral absorption. Then the fit is extended to the full energy
band and we proceed to produce fits with different soft excess models (in green). Last, the
resulting model from the soft excess testing is used to conduct our absorption test (in blue),
where possible final results are listed in the final blue square.
3.7. SHERPA & PYXSPEC: FITTING PACKAGES 67
Model Γ kT χ2 DOF
apec 4.04± 0.09 8.0± 0.2 824.15 164
zbbody 1.841± 0.005 0.0803± 0.0010 462.16 165
compbb 1.841± 0.004 0.0779+0.0006
−0.0002 462.64 164
Table 3.2: Soft excess testing for Mrk 704 using Sherpa. We show the resulting photon index
and temperature in keV of each of the three models, as well as their corresponding statistics.
Continuing with the third step, we conduct the absorption test to the resulting model of
zpowerlw+zgauss+zbbody. We used absori for the ionized absorber and ztbabs for the neutral
one. Continuing with the Mrk 704 example, Table 3.3 reports the results of the absorption test.
In this case, we see that the preferred model corresponds to absori*(zpowerlw+zgauss+zbbody),
with a χ2
ν = 2.2 in comparison with the χ2
ν = 2.8 that results from the fit with a cold absorber.
Moreover, when inspecting the resulting parameters of the fits, we see that the cold absorber
yields a column density where the parameter value is of the same order of magnitude as the
parameter error, indicating a weak constraint on the column density.
Model Γ keV nH ξ χ2 DOF
absori 1.909± 0.009 0.090± 0.002 0.088± 0.009 4.84+1.71
−1.15 359.17 163
zphabs 1.846+0.008
−0.009 0.0808+0.0010
−0.0011 0.005± 0.006 - 461.56 164
Table 3.3: Absorption test for Mrk 704 using Sherpa. We show the resulting photon index and
temperature in keV, as well as the ionization state of the warm absorber, ξ and the NH in units
of 1022 cm−2 of each absorber with their corresponding statistics.
In conclusion, the fitting process of Mrk 704 done with Sherpa results in the source being
affected by warm absorption, which is consistent with our premise that a PL-pol type 1 source
would show the presence of the outer, ionized layers of the torus in the line-of-sight. However,
the poor fit of χ2
ν = 2.2 is still an indicator that the source requires further analysis. Following
the process of this example of Mrk 704, we fitted all 11 sources from the PL-pol subsample. We
chose to also fit the 6 weakly polarized sources in order to have a point of comparison for the
results.
3.7.1.1 Results obtained with Sherpa
We show the resulting Γ obtained by testing three soft excess models in Fig. 3.13, in order to
visualize the variation of the main continuum component according to the nature of the second
continuum component, i.e. the soft excess. The typical range of Γ values is shown with a red
horizontal stripe (Corral et al. 2011; She et al. 2017). The values of the photon index show
significant variation for fits with different models for the soft excess component, with only Mrk
926 and NGC 7213 resulting in the expected values of Γ for the three soft excess models. We
would expect that, if the continuum is well fitted, the value of Γ would be confined within a
narrow range of values. Mostly, the fits yield power law indexes above the expected values of
Γ, suggesting that the continuum is not sufficiently well fitted.
68 CHAPTER 3. X-RAY ANALYSIS OF OPTICALLY POLARIZED SEYFERTS
IR
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x
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Model
APEC
BB
CBB
1.5    2.1
Figure 3.13: Power law index for different soft excess models tested with Sherpa. The red
shaded stripe indicates a range of 1.5 < Γ < 2.1, considering the typical range of values for the
photon index of Sy 1 galaxies.
Moreover, as explained with the example of Mrk 704, in order to determine which of the soft
excess models produces the best result, we need to review the rest of the parameters involved,
such as the temperature of the emitting gas. This process not always clearly concluded a single
preferred model.
Another problem of this fitting process is that many of the fits do not produce results where
the statistics of the model can reliably tell us whether the spectrum is affected by absorption.
We show a graph of the values of the χ2
ν of the model resulting from the absorption test for each
source in Fig. 3.14. The different colors correspond to the preferred soft excess model, and the
different shapes correspond to the preferred absorption models: with squares indicating sources
affected by cold absorption and triangles indicating the presence of warm absorption. We can
see that only six sources result in a χ2
ν < 1.5. These high values of χ2
ν are far from ideal, as we
require our conclusions on the presence of absorption to be supported by the model statistics.
This poses a main issue for our first set of results. At this point of the project, we chose to
change the approach of the fitting process which we describe in the next section.
3.7.2 PyXspec
The choice of switching to a Python script within the XSPEC fitting package, PyXspec, was
motivated by the intention of producing a systematic analysis as well as the aim of finding
a more solid statistical solution. The command line interface in Sherpa is not ideally suited
for the systematic fitting of a given sample, as this process often requires source by source
interaction with model parameters to adjust values and achieve the best fit for each source.
To reduce the need for manual adjustments, we created Python scripts that implement all the
3.7. SHERPA & PYXSPEC: FITTING PACKAGES 69
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1
2
3
4
5
2
2 Values by Source - Fits with Sherpa
Absorption Type
absori
zphabs
Soft Excess Model
apec
zbbody
compbb
Figure 3.14: Resulting χ2
ν for fits obtained in the absorption test done with Sherpa.
models used in Sherpa, following the three-step approach detailed in Sect. 3.5.
In this process, we added a fourth soft excess test: a second power law component of an
independent slope from that of the continuum. Similar to Fig. 3.13, we show the resulting Γ
for the four soft excess models tested with their corresponding PyXspec scripts in Fig. 3.15.
There is more agreement with the expected photon index in comparison to those reported in
Fig. 3.13. For example, we can see that the values of Γ of the source IRAS 15091-2107 are
closer to the 1.5 < Γ < 2.1 range in this set of results obtained with the PyXspec script than
to the ones obtained with Sherpa.
The development of a PyXspec script allows us to delimit each model parameter within
the expected range for these sources, as in the case of the blackbody temperature shown in
Fig. 3.16, and fit all the chosen models without the need to keep interactively adjusting the
parameters, resulting in a more systematic approach to the fitting process.
Having all the necessary scripts, including the corresponding statistical testing for each
step, (see Sect.3.6), we ran the same procedure for each source using the models described in
Sect. 3.4 up to the absorption test. At this point, we decided to select one model for the soft
excess. Having used 4 different soft excess models resulted in a difficult analysis that deviated
attention from the absorption test, requiring a full discussion per source for the preferred soft
excess before the absorption could be tested. Indeed, the interpretation of the soft excess is
a very complex and debated topic within the AGN community, which is out of the scope of
this thesis. We began by discarding the second zpowerlw being a model that is not statisti-
cally preferred over the other three, being the only one rejected by the F-test for some of the
sources. The second discarded model was apec, given that it does not homogeneously fit the
sample, meaning that the model yields parameters that do not make sense in the context of
these sources. Lastly, regarding the compbb model, some of the sources result in a negligible
optical depth, which indicates that the zbbody model is the one, out of the four, more suited
to homogeneously fit our sample. This is the reason why, in Table 4.2 on Ch. 4, we report our
70 CHAPTER 3. X-RAY ANALYSIS OF OPTICALLY POLARIZED SEYFERTS
IR
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Model
BB
CBB
APEC
2ndPwLw
1.5    2.1
Figure 3.15: Values of Γ for four models of soft excess tested with PyXspec scripts. The red
horizontal region corresponds to typical values for type 1 sources.
PyXspec script example
Figure 3.16: Example of Python code for scripting an entire model in order to avoid manually
adjusting parameter to parameter. This case corresponds to the case where a blackbody spec-
trum is added to account for the soft excess. Model is indicated in “modelname”. The script
lists the parameters, with “nH” and z corresponding to Galactic column density and source’s
redshift respectively. The script fixes the energy, 6.4 keV, and width 0.05 keV, of the Fe Kα
line. The blackbody temperature is not fixed but is indicated to start at 0.1 keV, having a
minimum of 0.01 keV and a maximum of 1.0 keV and varying in 0.001 keV intervals.
3.7. SHERPA & PYXSPEC: FITTING PACKAGES 71
final results obtained only using zbbody as a soft excess component.
We also introduced the testing for additional ionized Fe emission lines. When inspecting
the fits produced for the hard band, we realized that the overall results for the fitting of the
continuum could be further improved by accounting for additional emission lines corresponding
to Fe xxv at 6.697 keV and Fe xxvi at 6.966 keV, that are often present in type 1 AGN (e.g.
Mrk 590 & Mrk 205 Longinotti et al. 2007; Reeves et al. 2001, respectively).
Another significant change when switching to PyXspec is the choice of absorption models.
Concerning the neutral absorption, changing to ztbabs provides a more detailed and realistic
treatment of X-ray absorption in the ISM by considering the additional contributions from
molecular gas and dust components, while zphabs offers a simpler approximation based solely
on neutral gas absorption. We also change the absori model for zxipcf, as the former is an
obsolete model that was applied in the 90s to model AGN ionized absorption. With the advent
of Chandra and XMM-Newton and high-resolution spectroscopy, our knowledge of ionized gas
in AGN has significantly increased, resulting in the development of finer photoionization codes
and more complete atomic databases. Considering this, zxipcf provides a more refined descrip-
tion and a more realistic description of the absorber. We also tested a two-layered ionized
absorber by adding a second zxipcf to the model, also fixing the covering factor.
Lastly, one note is to mention that at the time of changing from Sherpa to PyXspec, a
better data set for MRK 1239 became available. In 2020, the observation with ID 0891070201,
initially included in the sample and analyzed with Sherpa, was not optimal as it was not cen-
tered on the target and had very short exposure, less than 10 ks. In December 2021, a new
data set became available, obs. ID 0891070101, with exposure > 100ks and centered on the
source. We replaced the previous observation, therefore results reported in Ch. 4 are derived
from analyzing this newer observation.
All these changes are reported in the paper by Gudiño et al. (2024), where we present the
analysis done with PyXspec for the sources with known polarization, i.e PL-pol and EQ-pol.
The complete set of results is listed in Ch. 4 and will be the main point of our further discussion.
3.7.3 Ionized reflection
Lastly, motivated by the fact that these sources have been previously studied by different groups
and are known to have complex spectra, (see Sect.4.3), we chose to conduct one last test by
including ionized reflection considering the possibility that it could significantly improve the
overall fits (Victoria-Ceballos et al. 2023). This test was applied once the results of the ab-
sorption test were already published and for this reason, we decided to apply it to a reduced
number of sources. We included sources from all the three sub-samples, in order to cover the
three scenarios of polar, equatorial, and weak polarization. We present the results of these tests
in Table 3.4.
We decided to compare the fits between zbbody only and zbbody+relxill. We can see that,
while the relxill contribution is significant in cases such as NGC 4593 or Was 45, the model is
not homogeneously favored as a good description for the soft excess. The example of Was 45 is
shown in Fig. 3.17, where we have marked the contribution of each additive component: zgauss
at 6.4 keV, the zbbody, and, in the bottom panel, relxill. Therefore, we are confident that the
inclusion of ionized reflection does not introduce a significant change in the soft excess model-
72 CHAPTER 3. X-RAY ANALYSIS OF OPTICALLY POLARIZED SEYFERTS
ing, and in turn, to the absorption parameters and the overall conclusion we derived. However,
the inclusion of this model poses an interesting approach that is likely to be considered in the
future development of this work.
Figure 3.17: Was 45 fitted with a zbbody, top, and then a zbbody+relxill, bottom. This source
does show a significant improvement by adding ionized reflection to its modeling. Parameters
of both fits are reported in Table 3.4.
3.7. SHERPA & PYXSPEC: FITTING PACKAGES 73
Model Γ Temp. log ξ χ2 DOF Ftest
Mrk 1218
zbbody 0.8414+0.10
−0.12 0.841+0.11
−0.09 106.12 90
+relxill 0.72X 0.84+0.11
−0.08 2.99X 104.46 88 50%
Mrk 231
zbbody 0.46±0.08 0.170±0.013 133.18 79
+relxill 0.69+0.22
−0.09 0.169+0.012
−0.014 0.3 +1.8
−0.3 127.36 77 82.10%
NGC 4593
zbbody 1.687±0.004 0.075±0.002 440.66 166
+relxill 1.643+0.012
−0.008 0.0603+0.006
−0.004 3.28+0.03
−0.06 370.91 164 99.99%
IRAS 15091-2107
zbbody 1.31+0.06
−0.07 0.5±0.03 164.48 127
+relxill 1.50±0.02 1.0E−4X 4.7X 349.92 125 χ2
BB < χ2
Rel
Mrk 704
zbbody 1.843 ±0.008 0.077±0.002 367.25 163
+relxill 1.770+0.014
−0.023 0.063+0.004
−0.006 3.29+0.02
−0.06 277.52 161 99.99%
Fairall 51
zbbody 0.999±0.007 0.0499±0.0013 6609.08 162
+relxill 1.047±0.003 0.010X 2.9899+0.0012
−0.0015 7458.54 160 χ2
BB < χ2
Rel
Was 45
zbbody 0.14±0.03 0.093±0.008 848.58 123
+relxill 1.86±0.03 0.135±0.010 1.104E−7X 319.09 121 99.99%
IZw 1
zbbody 2.33±0.03 0.118+0.007
+0.008 294.20 133
+relxill 2.31+0.04
−0.05 0.111+0.010
−0.012 3.62X 297.33 131 χ2
BB < χ2
Rel
Mrk 841
zbbody 1.46±0.02 0.104+0.003
−0.002 361.71 155
+relxill 1.28+0.04
−0.05 0.089+0.006
−0.010 3.45+0.04
−0.05 331.64 153 pvalue > 0.05
Mrk 926
zbbody 1.744±0.008 0.144+0.004
−0.005 267.30 167
+relxill 1.82±0.02 0.158+0.029
−0.013 1.0+0.3
−0.2 206.90 165 99.99%
NGC 7213
zbbody 1.695±0.010 0.173±0.006 209.21 163
+relxill 1.69±0.02 0.175 ±0.006 4.309X 208.90 161 11.3%
NGC 7603
zbbody 1.99±0.02 0.139±0.009 200.61 151
+relxill 2.02+0.02
−0.03 0.25±0.03 1.58+0.02
−0.03 153.71 149 99.99%
Table 3.4: Results of fits comparing the blackbody on its own and with the relxill added. The
table reports the resulting parameters: Γ, blackbody temperature [keV], the relxill ionization
parameter and the corresponding statistic. X denotes parameters for which the error could not
be calculated.
Chapter 4
Results and discussion
4.1 Results
As outlined in Sect. 3.5, our analysis consisted of three main steps aimed at drawing general
conclusions regarding the X-ray properties of our sample and conducting the absorption test.
This section reports the results of each of the steps presented in Tables 4.1, 4.2, and 4.3 and
described in Sections 4.1.1,4.1.2, and 4.1.3. All three tables are organized according to the
polarization characteristics –PL-pol, EQ-pol, and Weak Polarization–.
In the three tables, the second column corresponds to the fitted model. In Table 4.1, it
corresponds to the resulting baseline model. For Table 4.2, the baseline model does not include
any contribution from neutral absorption, regardless of the result in the hard band. This is
due to the fact that, when extending to the full energy band and comparing between fits with
and without cold absorption, the results indicated that the neutral absorber was not significant
once the model was extended to the full energy band. We removed it and tested the soft excess
against a simple baseline model consisting of the power law and the significantly detected Fe
emission lines.
Table 4.3 presents the results of the absorption test. Highlighted in bold, we indicate the
preferred model for each source: ColdAbs, WarmAbs, 2WarmAbs (indicating two warm ab-
sorbers), or “Unabs” for unabsorbed sources. Sources with the mark (*) and no model in
bold indicate that the AIC results are inconclusive for either of the absorption scenarios. The
X-ray luminosity, reported in column 9, corresponds to that of the preferred model. The last
three columns correspond to the AIC test: the AIC factor of the model without any absorber
(AICBB), the AIC value of each of the models with absorbers, and the factor that indicates the
fit improvement between the AICBB and AIC see Sect. 3.6. Finally, a parameter marked with
X indicates that the value falls to a minimum and error cannot be estimated.
74
4.1. RESULTS 75
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10)
Source Model NH Γ Γ Fekα EW Fexxv Fexxvi χ2
ν
[cm−2] Norm Norm [keV] Norm Norm
Polar Polarization
Mrk 1218 PwLw - 1.37 ±0.10 0.4+0.3
−0.3 - - - - 0.94
IRAS 15091-2107 PwLw+Fe - 1.72 ±0.06 2.3 ±0.2 1.3 ±0.6 0.08+0.03
−0.02 - - 1.10
NGC 4593 PwLw+Fe+Fe2 - 1.74 +0.04
−0.07 4.64 +0.2
−0.14 2.8 ±0.2 0.16± 0.02 - 5.5 +4.7
−4.3 1.12
Mrk 231 NH(PwLw) 2.4 +1.8
−1.7 1.0 ±0.3 6.0 +4.0
−2.0 - - - - 1.13
Fairall 51 NH(PwLw+Fe) 2.3 ±0.2 1.89 ±0.04 8.8 ±0.7 2.6 ±0.5 0.10+0.02
−0.03 - - 1.33
Mrk 704 PwLw+Fe - 1.82 ±0.02 2.69 ±0.06 1.1 ±0.2 0.12 ±0.02 - - 1.37
NGC 3227 NH(PwLw+Fe) 0.29 ±0.09 1.58 ±0.02 6.2 ±0.2 3.2 ±0.3 0.10±0.01 - - 1.49
Was 45 NH(PwLw+Fe) 7.2+0.9
−1.0 1.56 ±0.13 0.58 +0.16
−0.12 1.0 ±0.2 0.3± 0.3 - - 1.66
Mrk 766 PwLw+Fe+Fe2+Fe3 - 2.11±0.01 7.76±0.11 1.0±0.2 0.06± 0.03 0.9±0.2 0.6±0.2 2.10
ESO 323-G077 NH(PwLw+Fe+Fe2) 5.5±0.4 0.59±0.04 0.011+0.009
−0.008 1.94±0.08 0.5± 0.5 - 6.3±0.3 5.54
Mrk 1239 NH(PwLw+Fe) 57.7+2.8
−2.7 4.3±0.17 186+78.9
−53.6 0.017±0.002 0.2+0.2
−0.3 - - 6.0
Equatorial Polarization
Mrk 876 PwLw - 1.62 ±0.11 1.7 ±0.3 - - - - 0.86
NGC 3783 NH(PwLw+Fe+Fe2) 0.6 ±0.2 1.50 ±0.03 4.9 +0.3
−0.2 4.5 ±0.4 0.14+0.0.03
−0.02 - 1.3 ±0.4 1.65
Mrk 841 NH(PwLw+Fe) 0.7 ±0.4 1.52 ±0.06 1.5 ±0.2 1.0 ±0.2 0.10+0.03
−0.04 - - 1.75
I Zw1 PwLw - 2.33 ±0.05 2.7 ±0.2 - - - - 1.85
Mrk 509 PwLw+Fe - 1.855±0.010 8.41+0.12
−0.11 1.7±0.3 0.06+0.02
−0.03 - - 2.30
Akn 120 PwLw+Fe+Fe2+Fe3 - 1.932±0.008 9.88+0.11
−0.10 2.2±0.2 0.08±0.02 0.7±0.2 0.8±0.2 2.50
KUV 18217+6419 PwLw - 1.14±0.03 5.8122±0.0003 - - - - 2.53
Mrk 304 NH*PwLw 6.0±0.8 1.74±0.11 1.1+0.3
−0.2 - - - - 3.0
Weak Polarization
Mrk 896 NH(PwLw+Fe) 0.9+1.1
−0.0 2.4± 0.2 2.3± 1.0 0.6± 0.3 0.20±0.11 - - 1.03
NGC 7603 PwLw+Fe - 1.89± 0.03 5.4± 0.2 1.3± 0.5 0.08+0.04
+0.14 - - 1.09
Mrk 926 PwLw - 1.71± 0.02 10.9± 0.4 - - - - 1.12
NGC 7213 PwLw+Fe+Fe2 - 1.704± 0.014 2.50± 0.05 1.6± 0.2 0.14+0.11
−0.17 - 0.6± 0.2 1.40
NGC 7469 PwLw+Fe - 1.832± 0.011 7.94± 0.11 2.8± 0.3 0.10+0.08
−0.13 - - 1.59
PG 1211+143 PwLw+Fe+Fe2 - 1.83±0.02 0.996±0.002 0.268±0.006 0.07+0.04
−0.11 - 0.79+0.16
−0.14 1.40
Table 4.1: Hard band analysis. The resulting baseline model from the hard band fits.
Columns 3 to 8 report the value of the model parameters, with NH expressed in units of
×1022 cm−2 and photon index normalization in units of ×10−3[photons/keV/cm2/s]. When an
emission line is considered significant, we report the line normalization in columns 6, 7 and 8,
in units of ×10−5[photons/cm2/s]. The model statistics in column 9.
76 CHAPTER 4. RESULTS AND DISCUSSION
(1) (2) (3) (4) (5) (6) (7) (8)
Source Model Γ Γ kT BB χ2
ν F-test
Norm [keV] Norm
Polar Polarization
Mrk 1218 PwLw 0.95±0.03 0.21±0.08 - - 2.24
PwLw+BB 0.84+0.10
−0.11 0.13±0.02 0.84 +0.11
−0.09 0.1±0.2 1.18 99.99%
Mrk 231 PwLw 0.96±0.07 0.046±0.003 - - 4.30
PwLw+BB 0.46 ±0.08 0.023 ± 0.003 0.169 +0.014
−0.012 0.0116 ± 0.0012 1.68 100%
IRAS 15091-2107 PwLw+Fe 1.437±0.011 1.38±0.02 - - 6.47
PwLw+Fe+BB 1.31±0.03 0.93±0.04 0.52±0.02 0.24±0.02 1.98 100%
Mrk 704 PwLw+Fe 1.988±0.006 3.10918±0.011 - - 25.81
PwLw+Fe+BB 1.833±0.007 2.746±0.017 0.073±0.002 2.0 2.29 100%
NGC 4593 PwLw+Fe+Fe2 1.777±0.003 4.774±0.011 - - 25.32
PwLw+Fe+Fe2+BB 1.688±0.004 4.42±0.02 0.0740.002 1.28+0.14
−0.12 2.60 100%
Was 45 PwLw+Fe 0.23±0.03 0.042±0.002 - - 9.24
PwLw+Fe+BB 0.14±0.03 0.037±0.002 0.092±0.008 0.03±0.02 6.9 99.99%
ESO 323-G077 PwLw+Fe+Fe2 0.33±0.02 0.0493±0.0012 - - 32.5
PwLw+Fe+Fe2+BB 0.33±0.02 0.0262 ±0.0007 0.151 ±0.004 0.0212±0.0008 14.13 100%
Mrk 766 PwLw+Fe+Fe2+Fe3 2.247±0.003 8.342±0.014 - - 92.02
PwLw+Fe+Fe2+Fe3+BB 2.0565±0.004 7.33±0.02 0.0757±0.0009 3.17+0.13
−0.12 13.45 100%
NGC 3227 PwLw+Fe 1.376±0.003 4.502±0.010 - - 40.28
PwLw+Fe+BB 1.411±0.005 4.11±0.02 0.87±0.02 0.49±0.02 27.53 99.99%
Fairall 51 PwLw+Fe 1.055±0.007 2.12±0.02 - - 56.97
PwLw+Fe+BB 1.24±0.04 14.18±0.03 1.15±0.03 1.54+0.10
−0.09 37.57 99.99%
Mrk 1239 PwLw+Fe 3.38±0.04 0.14 ±0.02 - - 70.38
PwLw+Fe+BB -0.14±0.03 0.0133±0.0007 0.064±0.002 0.156±0.015 31.41 100%
Equatorial Polarization
Mrk 876 PwLw 1.99±0.04 2.550.04 - - 2.55
PwLw+BB 1.73 ± 0.03 1.64 ± 0.06 0.102 ± 0.006 0.48+0.6
−0.05 1.09 100%
I Zw1 PwLw 2.548±0.013 3.18±0.02 - - 4.50
PwLw+BB 2.34±0.03 2.75±0.06 0.109±0.007 0.426579+0.05
−0.04 2.25 100%
Mrk 841 PwLw+Fe 1.819±0.012 1.878±0.012 - - 24.81
PwLw+Fe+BB 1.46±0.02 1.29±0.02 0.101±0.003 0.60±0.04 2.36 100%
KUV 18217+6419 PwLw 1.547±0.010 10.07±0.08 - - 18.50
PwLw+BB 1.14±0.02 5.83±0.02 0.21095±0.005 1.45067±0.04 2.63 100%
Akn 120 PwLw+Fe+Fe2+Fe3 2.11289±0.003 11.95±0.02 - - 21.43
PwLw+Fe+Fe2+Fe3+BB 1.98113±0.006 10.56 ±0.06 0.138±0.002 0.77±0.02 4.20 100%
Mrk 509 PwLw+Fe 2.108±0.003 10.94±0.02 - - 33.20
PwLw+Fe+BB 1.955±0.006 9.64±0.05 0.102±0.002 1.38+0.06
−0.05 4.75 100%
Mrk 304 NH*PwLw 0.54±0.02 0.112±0.003 - - 12.16
NH*PwLw+BB 0.46±0.02 0.997±0.003 0.066 ±0.004 0.23±0.06 8.43 99.99%
NGC 3783 PwLw+Fe+Fe2 1.284±0.005 3.37±0.02 - - 40.38
PwLw+Fe+Fe2+BB 1.199+0.005
−0.006 3.05±0.02 0.0668+0.0017
−0.0011 2.9±0.3 14.80 100%
Weak Polarization
Mrk 896 PwLw+Fe 2.22 ± 0.02 1.75 ± 0.02 - - 1.42
PwLw+Fe+BB 2.18+0.03
−0.04 1.68 ± 0.04 0.07 ± 0.02 0.6+0.3
−0.4 1.11 99.99%
Mrk 926 PwLw 1.855±0.004 15.98±0.04 - - 2.65
PwLw+BB 1.73 ± 0.02 11.6 ± 0.2 0.156 ± 0.009 0.69 ± 0.07 1.19 100%
NGC 7213 PwLw+Fe+Fe2 1.789±0.003 2.78±0.06 - - 3.2
PwLw+Fe+Fe2+BB 1.696 ± 0.010 2.50 ± 0.03 0.170 ± 0.006 0.11 ± 0.010 1.30 100%
NGC 7469 PwLw+Fe 2.023±0.012 9.61±0.13 - - 24.0
PwLw+Fe+BB 1.885±0.006 8.56+0.04
−0.05 0.103±0.002 1.03±0.04 2.71 100%
NGC 7603 PwLw+Fe 2.140±0.010 7.13±0.04 - - 3.13
PwLw+Fe+BB 2.00 ± 0.02 6.27 ± 0.14 0.128 ± 0.008 0.59 ± 0.06 1.36 100%
PG 1211+143 PwLw+Fe 2.170±0.007 1.28±0.05 - - 34.6
PwLw+Fe+BB 1.827±0.009 0.979±0.008 0.0889±0.0014 0.68+0.04
−0.03 2.4 100%
Table 4.2: Full energy band. For each source, the first row shows the baseline model, and
the second row corresponds to the model with a black body (BB) as soft excess. The columns
3 to 6 report model parameters, with BB normalization in units of ×10−4[L39/[D10(1 + z)]2].
Columns 7 and 8 correspond to the fit statistics and the F-test resulting from comparing the
two models.
4.1. RESULTS 77
(1) (2) (3) (4) (5) (6) (7) (8) (9) (10) (11) (12) (13)
Source Model nH log ξ NH2 log ξ2 Γ kT [keV] Lum χ2
ν AICBB AIC 1/FAIC
Polar Polarization
Mrk 1218 ColdAbs 0.47+0.14
−0.12 - - - 1.37±0.09 0.062±0.012 0.56 ± 0.03 1.17 151.81 113.66 1.9 × 108
WarmAbs 0.15 ± 0.10 1.2 +0.3
−1.6 - - 1.0 +1.0
−0.3 0.80 +0.80
−0.10 1.21 118.18 2.01 × 107
Mrk 231 ColdAbs 0.6 ± 0.2 - - - 0.65 ± 0.10 0.106 +0.015
−0.012 0.324±0.012 1.29 140.41 110.78 2.7 × 106
WarmAbs 0.05 0.4 - - 0.48 +0.10
−0.07 0.168 ±0.013 1.82 152.41 0.002
NGC 4593 ColdAbs 0.060 ±0.010 - - - 1.735 ±0.009 0.0782 +0.0013
−0.0014 1.94 449.37 334.59 8.4 × 1024
WarmAbs 0.14±0.02 2.36±0.08 - - 1.727 ±0.007 0.079 ±0.003 0.437±0.001 1.40 253.78 3.0 × 1042
Fairall 51 ColdAbs 0.0X - - - 1.24±0.04 1.15±0.03 37.80 6095.58 6097.58 0.37
WarmAbs 1.31±0.03 0.57+0.06
−0.05 - - 1.81±0.03 0.111±0.003 3.35 549.52 > 1090
2WarmAbs 2.50 +1.03
−0.84 2.5 +0.5
−0.2 1.0 +0.4
−0.2 0.46 +0.05
−0.04 1.86 ±0.09 0.113 +0.006
−0.004 1.268 ±0.007 1.60 269.84 > 1090
IRAS 15091-2107 ColdAbs 0.08 ± 0.03 - - - 1.51 +0.06
−0.07 0.51 +0.04
−0.03 1.83 301.57 278.51 1.02 × 105
WarmAbs 0.12 +0.05
−0.12 -0.46 +0.14
−0.23 - - 1.63 +0.10
−0.07 0.45 +0.07
−0.11 4.562±0.05 1.63 251.32 8.2 × 1010
Mrk 704 ColdAbs 0.052±0.014 - - - 1.877±0.014 0.0738±0.0014 2.05 382.46 344.67 1.6 × 108
WarmAbs 0.13 +0.03
−0.04 2.15 +0.14
−0.15 - - 1.882 +0.012
−0.013 0.076 ± 0.003 1.84 318.36 8.3 × 1013
2WarmAbs 0.061 +0.002
−0.009 0.28 +0.46
−0.13 0.23 +0.02
−0.14 2.44 +0.60
−0.11 1.9 ±0.9 0.113 +0.113
−0.003 2.919±0.010 1.73 293.35 2.2 × 1019
Was 45 ColdAbs 1.6±0.3 - - - 0.69+0.09
−0.08 0.052+0.006
−0.005 4.54 858.41 565.55 > 1090
WarmAbs 2.74+0.07
−0.09 0.276921+0.005
−0.003 - - 1.17+0.05
−0.06 0.095±0.003 5.96 737.15 2.0 × 1062
2WarmAbs 0.24 +0.05
−0.04 2.04+0.15
−0.10 2.5±0.09 0.278±0.004 1.30+0.07
−0.05 0.096+0.003
−0.004 0.325±0.006 3.20 398.63 4.4 × 1087
Mrk 766 ColdAbs 0.085 ±0.008 - - - 2.138±0.005 0.0758±0.0005 11.72 2220.43 1926.75 5.9 × 1063
WarmAbs 0.39+0.05
−0.03 0.86±0.05 - - 2.175+0.018
−0.009 0.120 +0.004
−0.003 0.128±0.002 2.87 480.05 > 1090
NGC 3227 ColdAbs 0.0X - - - 1.400±0.003 0.838±0.012 27.71 4635.78 4640.26 0.11
WarmAbs 0.36±0.02 1.56±0.03 - - 1.509±0.006 0.63±0.03 6.12 1023.29 > 1090
2WarmAbs 0.233±0.006 0.84±0.02 0.166+0.037
−0.014 2.730+0.102
−0.007 1.596+0.003
−0.004 0.15±0.02 0.144±0.001 3.39 573.97 > 1090
ESO 323-G077 ColdAbs 0.97±0.05 - - - 0.6±0.5 0.085±0.002 15.84 1259.63 2436.98 X
WarmAbs 0.30±0.02 -1.35+0.05
−0.15 - - 0.14±0.03 0.116±0.004 0.230±0.001 6.64 1027.53 2.5 × 1050
Mrk 1239 (*) ColdAbs 0.02+0.03
−0.02 - - - -0.15±0.03 0.145+0.004
−0.005 31.06 10495.39 4390.27 > 1090
WarmAbs 0.05X 5.96±0.04 - - -0.16±0.03 0.147 ±0.002 0.159±0.003 31.45 4415.32 > 1090
Equatorial Polarization
Mrk 876 Unabs - - - - 1.73 ± 0.03 0.102 ± 0.006 12.9±0.4 1.09 106.95 - -
ColdAbs 0.0X - - - 1.73+0.06
−0.05 0.110+0.014
−0.015 1.10 108.96 0.37
WarmAbs 0.04±0.04 0.2+0.2
−1.4 - - 1.73+0.06
−0.05 0.13+0.02
−0.03 1.19 117.70 0.005
1Zw 1 ColdAbs 0.0X - - - 2.80±0.03 1.0+0.004
−1.0 3.02 307.83 409.14 X
WarmAbs 0.09+0.04
−0.09 -0.2+0.6
−0.2 - - 2.3+0.3
−0.4 0.12+0.011
−0.005 8.65±0.03 1.91 261.57 1.11 × 1010
Mrk 841 ColdAbs 0.0X - - - 1.45±0.02 0.102±0.003 2.38 375.48 378.51 0.22
WarmAbs 0.12+0.04
−0.06 1.3+0.2
−0.5 - - 1.47±0.02 0.120+0.004
−0.007 0.825±0.007 2.15 342.26 1.6 × 107
KUV 18217+6419 ColdAbs 0.0X - - - 1.15±0.02 0.199+0.004
−0.005 2.65 445.24 447.61 0.3
WarmAbs 1.7+1.7
−1.2 3.33+0.09
−0.16 - - 1.15±0.02 0.199+0.004
−0.005 1426.3±9.9 2.47 417.12 1.3 × 106
Akn 120 ColdAbs 0.0X - - - 1.981±0.006 0.138±0.002 3.66 617.45 619.45 0.4
WarmAbs 0.30 +0.05
−0.05 0.27+0.02
−0.03 - - 1.98±0.006 0.1364+0.0009
−0.0013 11.42±0.02 3.49 587.54 3.1 × 106
Mrk 509 Unabs - - - - 1.955±0.006 0.102±0.002 12.5±0.2 4.75 802.98 - -
ColdAbs 0.0X - - - 1.955±0.005 0.102±0.002 4.78 805.77 0.3
WarmAbs 0.014+0.002
−0.05 4.6+0.13
−0.12 - - 1.952±0.007 0.122±0.002 6.29 1052.53 X
Mrk 304 ColdAbs 0.96±0.05 - - - 0.85±0.03 0.0463±0.0015 3.335±0.007 6.03 1146.04 817.48 2.2 × 1071
WarmAbs 0.05±0.03 0.02±0.16 - - 0.51±0.02 0.078±0.003 7.92 1064.92 4.1 × 1017
NGC 3783 (*) ColdAbs 0.375±0.015 - - - 1.452±0.011 0.0654+0.0007
−0.0006 2.82 2453.57 477.19 > 1090
WarmAbs 0.250±0.008 -0.440.04 - - 1.418±0.006 0.090±0.002 0.688±0.002 7.13 1178.32 > 1090
Weak Polarization
Mrk 896 Unabs - - - - 2.18 ± 0.04 0.07 ± 0.02 1.059±0.014 1.11 123.78 - -
ColdAbs 0.0X - - - 2.20+0.08
−0.07 0.99±0.03 3.07 324.99 0.0
WarmAbs 0.02X 5.6+3.6
−5.6 - - 2.18+0.04
−0.03 0.07 ± 0.02 1.13 127.66 0.14
Mrk 926 Unabs - - - - 1.73 ± 0.02 0.156 ± 0.009 34.210±0.04 1.19 193.17
ColdAbs 0.0X - - - 1.737+0.012
−0.020 0.154 ± 0.010 1.20 195.37 0.33
WarmAbs 0.04+0.04
−0.82 3.450.00.0 - - 1.73 ± 0.02 0.158+0.014
−0.004 1.20 197.02 0.15
NGC 7213 Unabs - - - - 1.696 ± 0.010 0.170 ± 0.006 0.1238±0.0001 1.30 224.19
ColdAbs 0.01X - - - 1.783±0.005 0.121438+0.004
−0.003 6.14 1008.43 0.0
WarmAbs 0.05+0.05
−0.82 −0.96 ± 0.13 - - 1.721 ± 0.011 0.143 ± 0.004 1.23 213.88 173.13
NGC 7469 ColdAbs 0.0X - - - 2.09±0.02 0.99±0.002 32.24 462.33 5363.91 0.0
WarmAbs 0.05+0.0003
−0.05 2.72+0.14
−0.19 - - 1.891±0.007 0.111±0.002 2.618±0.003 2.33 396.53 1.9 × 1014
NGC 7603 Unabs - - - - 2.01 ± 0.02 0.128 ± 0.008 5.51±0.03 1.36 215.64
ColdAbs 0.0X - - - 1.99 ± 0.02 0.129+0.007
−0.009 1.37 217.77 0.35
WarmAbs 0.05X 4.95+0.23
−0.19 - - 2.00 ± 0.02 0.139+0.004
−0.020 1.39 220.75 0.08
PG 1211+143 ColdAbs 0.0X - - - 1.833+0.010
−0.008 0.0885±0.0013 2.38 377.03 379.03 0.37
WarmAbs 0.16 ± 0.05 1.54+0.12
−0.16 - - 1.87 ± 0.02 0.106 ± 0.004 8.58±0.04 1.60 256.63 1.4 × 1026
Table 4.3: Absorption test. Each absorber is reported with its respective NH in units of
×1022 cm−2 and log ξ for warm absorbers. Photon index and BB temperature are reported
in columns 7 and 8, followed by the X-ray luminosity in units of ×1043 erg s−1 and model
statistics. The last three columns correspond to the AIC test.
78 CHAPTER 4. RESULTS AND DISCUSSION
4.1.1 Polar-polarized sources
We start by describing the hard band fits reported in Table 4.1. The first component corre-
sponds to the power law continuum. The photon index ranges from the flatter 0.6 to the steeper
4.3. In particular, for 8 sources the photon index is in a range of 1.4 < Γ < 2.1, (< Γ >= 1.7).
Also in the hard band, we tested the presence and significance of emission lines corresponding
to Fe Kα, Fexxv, and Fexxvi. We found a significant Fe Kα line for 9 of the sources, with
only Mrk 1218 and Mrk 231 not showing any emission line. Additionally, NGC 4593 and ESO
323-G077 show the Fexxvi line, and Mrk 766 shows both Fexxv and Fexxvi lines.
Extending to the full energy range, we find that soft excess is ubiquitous, see Table 4.2.
The significance of this component is indicated by an F-test > 99% for all sources. For the
black body component, we find temperatures of kT < 1 keV for 9 out of the 11 sources. Fairall
51 yields a higher temperature, however, it decreases once we add absorption components, as
reported in Table 4.3.
Regarding the absorption test, according to the AIC, we determine the presence of absorp-
tion in all the PL-pol sources, see Table 4.3. Mrk 1218 and Mrk 231 favor the cold absorption
scenario, resulting in column densities on the order of 1021 cm−2. For NGC 4593, IRAS 15091-
2107, Mrk 766, and ESO 323-G077, the AIC indicates that warm absorption is the best scenario.
The fit is further improved by adding a second warm absorption component for the cases of
Fairall 51, Mrk 704, Was 45, and NGC 3227. For the warm absorbers, all the column den-
sities are within the order of NH ∼ 1020 − 1022 cm−2. We obtain a large range of values in
the ionization parameters, from −1.4 < log ξ < 5.9. In particular, for Mrk 1239 the results
are inconclusive, with the AIC value indicating that the fit improves by adding an absorption
component, but none of the two scenarios (cold vs warm) can be statistically favored.
4.1.2 Equatorial-polarized sources
In the hard band of the EQ-pol sources, we find one source with a flat spectrum corresponding
to a power law index of ∼ 1.1 (KUV 18217+6419), the remaining 7 sources have a photon index
in the range of 1.5 < Γ < 2.3. The Fe Kα line is significantly detected in 4 sources: NGC 3783,
Mrk 841, Mrk 509, and Akn 120. NGC 3783 also shows the Fexxvi line and Akn 120 shows
both the Fexxv and Fexxvi lines.
In the full energy range, we again obtain a result that supports that the soft excess is ubiq-
uitous, with F-tests > 99% for all the sources. The black body temperature is found to be
lower than 1 keV in all sources.
As to the absorption test, we do not find that this component is ubiquitous, as in the case of
the PL-pol sample. We can determine that 6 out of 8 sources favor the presence of absorption.
Mrk 876 and Mrk 509 are better fitted without any absorber, according to the small value of
the AIC test. Mrk 304 favors the cold absorber scenario, with column density of the order of
NH ∼ 1021 cm−2. I Zw1, Mrk 841, KUV 18217+6419 and Akn 120 favor the scenario of warm
absorption, yielding column densities of the orders of NH ∼ 1020 − 1021 cm−2 and a range of
ionization parameters of 3.3 > log ξ > −0.2. None of the fits improved by the addition of a
second warm absorption component. NGC 3783, similar to the case of Mrk 1239, results in an
AIC that indicates fit improvement by the addition of an absorption component, however, the
test does not indicate whether the favorable scenario involves either cold or warm absorption.
4.2. DISCUSSION 79
4.1.3 Weakly-polarized sources
The hard band fits for the weakly polarized yield a smaller range of photon index values. For 5
sources, the photon index is between 1.7 < Γ < 1.9, < Γ >= 1.79, and one source has a steeper
power law of index 2.4 (Mrk 896). We find the Fe Kα present in the spectra of 5 sources, with
NGC 7213 also showing the Fexxvi line. Only Mrk 926 does not show any emission lines.
Extending to the full energy range, we once again find the soft excess to be ubiquitous,
with temperatures below 1keV except for the case of NGC 7469. Just as is the case for ESO
323-G077 and Fairall 51, the BB temperature of NGC 7469 gets adjusted to a lower value once
we do the absorption test.
The results of the absorption test indicate that only two sources show a significant improve-
ment when absorption is added to the model. For both sources, NGC 7469 and PG 1211+143,
the favored model is the warm absorption one. Their resulting parameters are in a range simi-
lar to those of the previous sub-samples – PL-pol and EQ-pol –, with NH ∼ 1020 − 1021 cm−2
and log ξ of −0.96, 1.54, respectively. No source in this sub-sample favors the cold absorption
scenario.
4.2 Discussion
To organize the discussion of our results, we first consider the subsamples and the possible bias
due to the data selection. We then examine the key spectral features observed, including the
power-law continuum, Fe emission lines, and soft excesses. This is followed by a discussion on
the absorption test, examining the incidence of absorption and the parameters of the found
absorbers. We will compare our findings among sub-samples, and explore their correlations
with optical polarization and the unification scheme. Additionally, we compare our findings
with those from previous studies, offering both a complementary perspective and a compara-
tive analysis of these sources, Sect. 4.3. When available, we cite X-ray analysis of the same
XMM-Newton observation as the one used in our work.
4.2.1 Considerations on the sample
We considered the possible bias due to the data selected for this analysis. Our selection crite-
rion was to choose the longest exposure time for each source. Since the exposure times of the
spectra in our sample do not provide a uniform S/N, this choice may introduce some bias as
the three sub-samples contain sources with a higher number of counts, which makes it easier to
find deviations from the continuum and absorption features are more easily detected. Figure
4.1 is a histogram of the number of counts for each source, color-divided by sub-samples. We
can see that there is no significant distinction between the subsamples regarding the quality of
the observation in terms of the number of X-ray counts, which, if present, might have hindered
the modeling of the spectrum.
Regarding the X-ray luminosity of the sources, which may be intuitively associated with
a higher frequency of the absorber due to the photoionization process, we show in Figure 4.2
a histogram of the X-ray luminosity calculated for the model resulting from the absorption
test. This plot shows that the EQ-pol sources are generally more luminous than the other two
subsamples. The difference in X-ray Luminosity could be an indicator of a viewing angle effect
80 CHAPTER 4. RESULTS AND DISCUSSION
and/or structural differences in the emission and scattering regions around these sources. We
find no bias due to either counts or luminosities among our subsamples.
0.0 0.2 0.4 0.6 0.8 1.0 1.2 1.4
Counts 1e6
0
1
2
3
4
5
6
7
Nu
m
be
r o
f s
ou
rc
es
Spectral data counts
PL-pol
EQ-pol
W-pol
Figure 4.1: The spectral counts from 0.5 to 10 keV in the EPIC pn spectra.
1042 1043 1044 1045 1046
X-ray Luminosity [erg/s]
0
2
4
6
8
10
Nu
m
be
r o
f s
ou
rc
es
X-ray Luminosity
PL-pol
EQ-pol
W -pol
Figure 4.2: X-ray Luminosity distribution of our sample, (0.5-10 keV). The value plotted for
each source corresponds to the value reported in Table 4.3, the luminosity of the preferred
model after the absorption test.
4.2. DISCUSSION 81
4.2.2 The X-ray spectrum of optically polarized Seyfert 1: a first
order description.
Regarding the common X-ray properties of our sample as derived from the spatial analysis and
the fitting of the main spectral components, we highlight the following results. All the sources
of our sample are consistent with a point-like source regarding their X-ray emission described
in Sect. 1.3.3.
The power law continuum
In the hard band, the main continuum component is well fitted with a power law for the
majority of the sample, with photon index ranging between 1.4 < Γ < 2.3 for 21 of the 25
sources, with a mean value of Γ = 1.79. This is considered to be typical values of Γ for Sy 1
galaxies from studies of large samples of Seyfert galaxies, e.g. (Piconcelli et al. 2005; Cappi
et al. 2006; Singh et al. 2011; Corral et al. 2011; She et al. 2017). Deviations from this range
of values of the power law index can be due to a reflection component that has not been con-
sidered, like the case of ESO 323-G077. Another example is Mrk 231 which results in a flatter
slope, Γ = 1.0, with previous analysis reporting this source to have a scattered power law also
affected by reflection.
The incidence of the Fe Kα emission line
We detect the Fe Kα line in 18 sources, ∼ 72% of our sample. This narrow component of
the line is associated with Compton reflection taking place off the molecular torus, (Nandra
2006; Ricci et al. 2014). We focused exclusively on fitting the narrow component of the Fe
Kα line, without incorporating relativistic broadening effects, which are typically associated
with emission originating near the accretion disk. The narrow Fe Kα line produced by neutral
reflection is almost ubiquitous in type 1 AGN spectra, offering a solid probe of the torus region.
In contrast, the broadened Fe Kα, is not always present in AGN spectra and its detection
requires a much higher number of counts as shown by De La Calle Pérez et al. (2010), making
it less suitable for our systematic approach. Nandra et al. (2007) reports the presence of the
narrow component of the line in ∼ 80% of a Sy 1 sample that includes 7 of our sources, with
a consistent result between the two analyses of the same sources (i.e the analysis reported by
Nandra et al. (2007) and the one reported herein).
The soft excess, a phenomenological approach
Extending to the soft energy band, we find the soft excess present in 100% of our sources.
This is an interesting result in itself when considering studies such as Singh et al. (2011) and
Corral et al. (2011), where the incidence of soft excess is on 85% of their samples (17/20 and
35/41 respectively). Notably, in the case of Corral et al. (2011), it is mentioned that the esti-
mate should be considered a lower limit given the specifics of their analysis, hinting at a larger
incidence of soft excess. After testing four different soft excess models, we decided to stick
with a black body model to characterize the soft excess, (Sect. 3.7 (Ricci et al. 2017). All the
sources yield temperatures below 0.5 keV, with a resulting average temperature of 0.13 keV.
These are consistent with the temperatures found by Corral et al. (2011).
The inclusion of ionized reflection
As a part of the ongoing investigation on the nature of the soft excess, it has been suggested
82 CHAPTER 4. RESULTS AND DISCUSSION
that the Compton reflection off the inner accretion disk could provide a significant contribution
to the emission < 2keV. Nardini et al. (2011), for example, find that a good approximation
for the Suzaku spectrum of Akn 120 is by a combination of ionized reflection, taking place
at the accretion disk, and neutral reflection off the torus. Waddell & Gallo (2020) suggests
that blurred reflection could provide the description behind the soft excess found in NLSy1.
Moreover, Garćıa et al. (2019) study whether the soft excess in Mrk 509 can be fitted by either
a Comptonizing corona or a relativistically blurred ionized reflection from the accretion disk,
concluding that while the reflection component provides a good description of the data, further
work is needed to truly interpret the parameters involved. Including ionized reflection as a
contributor to the soft excess definitively poses an interesting idea. However, our particular
test with ionized reflection proved to deviate from the original approach of testing common
components in the whole sample. While the inclusion of ionized reflection proved significant
in sources like NGC 4593, it was also rejected by the F-test in sources like Fairall 51 or IRAS
15091-2107, see Table 3.4 in Sect. 3.7.3. Thus the model is not favored homogeneously by
our sample. We therefore conclude that further tests with this component are better suited to
future work in some sources, as explained in Sect. 5.2.
To conclude with the comments on the common characteristics of Sy 1 galaxies, we see
that while these provide a reasonable first approximation for our sample, our sources exhibit
greater spectral complexity, and require a more detailed model to achieve a statistically good
description of the spectrum (with χ2
ν ≈ 1.0). However, this complexity deviates from a purely
systematic fitting approach, as certain features in the spectra suggest the presence of additional
processes not included in this study. One example is the continuum shape of Mrk 1239 which
is affected by Compton reflection (Buhariwalla et al. 2020), highlighting the need for further
refinement to capture the full range of spectral characteristics in our sources. In Sect. 4.3, we
provide examples of these details as reported in previous publications on these sources.
4.2.3 The X-ray absorption in optically polarized Seyfert 1
The incidence of absorption in a sample of 25 Seyferts
Subsample CAbs WAbs 2 WAbs Unabs N/A Total sources
PL-pol 2 4 4 0 1 11
EQ-pol 1 4 0 2 1 8
WK-pol 0 2 0 4 0 6
Total sources 3 10 4 6 2 25
Table 4.4: Incidence of absorption: the results of the absorption test. Table indicates
how many sources from each subsample result in each of the possible absorption scenarios
reported in Table 4.3. CAbs = cold absorption, WAbs = warm absorption, 2WAbs = 2 warm
absorbers, Uabs = source is unabsorbed, N/A = test results inconclusive in regards to the
absorption scenario.
The incidence of absorption varies among subsamples, with the PL-pol being the most af-
fected by absorption and the Weak pol. having the highest fraction of unabsorbed sources. Let
us consider the actual statistics of the presence of absorption, as shown in Table 4.4 as well as in
Fig. 4.3. Overall, absorption is found in 68% of our sample (17/25). Specifically, 56% (14/25)
are affected by warm absorption and 3 sources favor the cold absorption scenario. Unabsorbed
4.2. DISCUSSION 83
sources correspond to 24% of our sample, (6/25). For 2 sources, Mrk 1239 and NGC 3783, we
are unable to determine which scenario is favored. The incidence of absorption is depicted in
Fig. 4.3, where we indicate the absorption scenario favored by each of the 25 sources: warm
absorption, two warm absorbers, cold absorption, unabsorbed, or an undetermined scenario.
We used different colors for each of the three subsamples: PL-pol in blue, EQ-pol in green, and
Weak pol. in purple.
Mrk 
12
18
Mrk 
23
1
NGC 45
93
Fai
ral
l 5
1
IRA
S 1
50
91
-21
07
Mrk 
70
4
Was 
45
Mrk 
76
6
NGC 32
27
ES
O 32
3-G
07
7
Mrk 
12
39
Mrk 
87
6
1Z
w 1
Mrk 
84
1
KU
V 18
21
7+
64
19
Akn
 12
0
Mrk 
50
9
Mrk 
30
4
NGC 37
83
Mrk 
89
6
Mrk 
92
6
NGC 72
13
NGC 74
69
NGC 76
03
PG
 12
11
+14
3
Source
ColdAbs
WarmAbs
2WarmAbs
Unabsorbed
Undetermined
Ab
so
rp
tio
n 
Sc
en
ar
io
Incidence of Absorption
Sub-sample
PL
EQ
WK
Figure 4.3: Incidence of absorption according to the different scenarios tested. Three
sources with cold absorption, 10 with warm absorption, and 4 sources with two warm absorbers.
There are 6 sources that resulted unabsorbed, and 2 sources undetermined. We tested a total
of 25 sources, with each color corresponding to a polarization subsample: 11 polar-polarized in
blue, 8 with equatorial polarization in green, and 6 with weak polarization in purple.
To compare the absorption incidence between subsamples, we applied the Fisher’s Exact
Test (Fisher 1922), which is particularly suitable for small sample sizes. For sources with an
undetermined absorption scenario in our analysis, we considered results from individual X-ray
analysis on the specific source reported in the literature. These are Mrk 1239 and NGC 3783,
both affected by warm absorption according to detailed literature analyses (Buhariwalla et al.
2020, 2023; Blustin et al. 2002; Mao et al. 2019).
The resulting p-value from conducting the Fisher’s Exact Test comparing EQ-pol and PL-pol
is 0.164, which is considerably higher than the standard threshold of 0.05 used for 95% confi-
dence. This means that while we have an incidence of absorption of 100% in the PL-pol vs 63%
in the EQ-pol, there is not a significant statistical difference between these two sub-samples.
In contrast, when comparing the PL-pol sample with Weak pol. sources, the resulting p-value
is 0.006, indicating a statistical difference regarding the incidence of absorption up to 99%
confidence. This finding suggests that the gas responsible for X-ray absorption is associated
with the region where the scattering that causes optical polarization takes place. Moreover, we
can interpret the presence of absorbing gas as an orientation effect, with Weak pol. being seen
84 CHAPTER 4. RESULTS AND DISCUSSION
completely face on, without the effects of any absorber, and the orientation of the scattering
regions is such that polarization cancels out.
The absorption parameters
We examine the parameters of the identified absorbers before interpreting them within the
context of the unification scheme. Starting with the column density of both cold and warm
absorbers, we find values ranging within NH ∼ 1020 − 1022 cm−2. Merloni et al. (2014) reports
on the incidence of obscuration among AGN, proposing a threshold of NH = 1021.5 cm−2 as
a separation between X-ray obscured and unobscured sources. This threshold is marked by a
red line in Fig.4.4, which shows the column density distribution for all absorbed sources. In
this figure, sources affected by warm absorption are marked with circles, while those with cold
absorption are marked with triangles. This separation allows us to distinguish/classify among
Compton-thin sources that show the effects of absorption but their spectra are still dominated
by the primary emission.
The criterion of Merloni et al. (2014), refers specifically to neutral/cold absorption. Thus,
we find two PL-pol sources corresponding to this “obscured” classification, Mrk 231 and Mrk
1218. For Mrk 231 it has been suggested that this obscuration is responsible for the flat power
law slope, Γ ≈ 0.6 (Braito et al. 2004, Sect. 4.3). This scenario could also be consistent with
the other PL-pol with cold absorption, Mrk 1218.
Particularly, for the warm absorbers, we see that the value of the parameters, (both column
density and ionization parameters) are consistent with the range reported by Laha et al. (2021):
NH ∼ 1021−1022.5 cm−2 and −1.0 < log ξ < 3.0. An exception is observed for KUV 18217+6419
and the second absorber of NGC 3227, where the ionization parameter is log ξ > 3. The fit
of both sources can be significantly improved by considering a more complex model (see Sect.
4.3). Fig. 4.5 corresponds to the distribution of NH values with the color map corresponding
to log ξ. No specific correlation is observed indicating that a higher NH tends to have a lower
ionization parameter, or vice versa.
4.2.4 Interpreting the unified scheme by Smith through X-ray ab-
sorption
The unified scheme by Smith et al. postulates that in EQ-pol Sy 1, the line of sight passes
through the NLR and central region, with polarization from a scattering region co-planar to
the accretion disk dominating the optical polarization spectrum. In contrast, a PL-pol Sy 1 is
observed through the outer layers of the torus, causing polarization from the bi-conical polar
region, the NLR, to dominate the optical polarization spectrum. Finally, Smith proposed that
sources of weak polarization can be considered Null-polarization candidates, with the line of
sight aligned to the principal axis, looking directly into the central region, (see Fig. 2.10).
In terms of X-ray absorption, a completely unabsorbed source would correspond to the
Null-polarization description. Conversely, EQ-pol can be considered classical type 1 sources.
While the incidence of warm absorption in type 1 sources has been largely documented (e.g.,
Reynolds & Fabian 1995; Laha et al. 2014, 2016), a higher incidence of absorption would be
expected in PL-pol sources given they are seen through the outer layers of the torus, which is
likely associated to ionized X-ray absorption. This scenario is supported by the results obtained
by Blustin et al. (2005), who argues that warm absorbers in Seyfert galaxies are more likely
4.2. DISCUSSION 85
0.5 0.1 0.5 1.0 1.5 2.0 2.5 3.0
NH (1022 atoms/cm2)NGC 32
27
_2NGC 74
69Mrk 
70
41Z
w 1Mrk 
84
1
IRA
S 1
50
91
-21
07NGC 45
93
PG
 12
11
+14
3
Mrk 
70
4_2
Was 
45
ES
O 32
3-G
07
7Akn
 12
0Mrk 
76
6
NGC 32
27
Fai
ral
l 5
1_2
KU
V 18
21
7+
64
19Fai
ral
l 5
1Was 
45
_2Mrk 
12
18Mrk 
23
1Mrk 
30
4
So
ur
ce
Distribution of NH values
Sub-sample
PL
EQ
WK
Figure 4.4: The column density distribution for the sources that are affected by absorption,
with circles indicating sources with preferred warm absorption and triangles for those with
preferred cold absorption. According to their optical polarization, blue corresponds to polar
sources, green to equatorial, and purple to sources with weak polarization. Labels marked with
a 2 correspond to the second warm absorber found in some cases. The red line corresponds
to a value of NH = 1021.5 cm−2, a threshold between obscured and unobscured AGN, a value
suggested by Merloni et al. (2014).
0.0 0.5 1.0 1.5 2.0 2.5
NH (1022 atoms/cm2)
NGC 3227_2
NGC 7469
Mrk 704
1Zw 1
Mrk 841
IRAS 15091-2107
NGC 4593
PG 1211+143
Mrk 704_2
Was 45
ESO 323-G077
Akn 120
Mrk 766
Mrk 1218
Mrk 231
NGC 3227
Mrk 304
Fairall 51_2
KUV 18217+6419
Fairall 51
Was 45_2
So
ur
ce
Distribution of NH/Ionization Parameter
1
0
1
2
3
4
5
lo
g 
Figure 4.5: Distribution of column density of the Warm Absorbers. The color map indicates
the corresponding ionization parameter for each absorber. Sources labeled with a 2 correspond
to sources with a second warm absorber.
86 CHAPTER 4. RESULTS AND DISCUSSION
to originate from outflows associated with the dusty torus, making an association to PL-pol
sources more plausible.
The incidence of absorption among PL-pol sources is 100%, (considering the reported result
of Mrk 1239 from Buhariwalla et al. 2020, 2023). In 82% of these sources, we confirm the
presence of at least one warm absorber, as supported by the AIC reported in Table 4.3; specif-
ically the value of 1/FAIC , Sect. 3.6. Studies on large samples of Sy 1, (Reynolds & Fabian
1995; Laha et al. 2021) report the presence of warm absorbers in 50-65% of the sources. In
comparison, the higher incidence of warm absorption in our sample seems very consistent with
the hypothesis that the gas producing the X-ray absorption and the material responsible for
the optical polarization are related.
On the other hand, 63% of the sources with EQ-pol are affected by warm absorption, with
only one favoring the cold absorption scenario. This percentage of incidence of warm absorp-
tion is consistent with the reports from (Laha et al. 2014, 2016), where warm absorbers are
significantly found in 65% of a sample of 26 Seyfert galaxies selected from the Catalog of AGN
in XMM Archive (CAIXA) with the criteria of being unobscured sources, NH < 2× 1022 cm−2.
Furthermore, in a review on ionized outflows found un AGN, Laha et al. (2021) establishes that
these ionized clouds can be located from 0.1pc-1kpc, which corresponds to scales that go from
the BLR up to the NLR. Moreover, the resulting unabsorbed sources can lend support to the
premise of absorption due to material that is not ubiquitous in AGN, while the torus is always
present but not aligned with our line of sight. Overall, we see no deviation from the expected
spectra of Seyfert 1 sources.
The work by Bianchi et al. (2007) reports on the “Iwasawa–Taniguchi (IT) effect”, an X-ray
phenomenon analogous to the optical Baldwin effect where there is a decrease in the equiva-
lent width (EW) of the Fe Kα emission line with increasing X-ray luminosity, indicating that
the line becomes less prominent relative to the continuum at higher luminosities (Iwasawa &
Taniguchi 1993). The equivalent width (EW) is a measure of the strength of an emission line
relative to the underlying continuum; thus, it serves as an indicator of the prominence of the
Fe Kα line. In the context of the IT effect, the decrease in EW with increasing luminosity can
be attributed to the torus geometry. Higher X-ray luminosities correlate with a larger torus
opening angle, thus its inner wall is smaller. This results in the radiation from the central re-
gion being reflected off of a smaller surface. This geometry decreases the torus covering factor,
allowing more direct X-ray continuum emission to reach the observer. As a consequence, the Fe
Kα line, which is predominantly produced by reflection off the torus, becomes less prominent
relative to the continuum, thus a lower EW.
In the context of polarization, this effect would be evident in the EQ-pol sources, which
exhibit increased scattered light from the torus, showing a larger torus opening angle and lead-
ing to a more prominent polarized signal from the equatorial scattering region. Our sample
conservatively shows this relationship with the EW of the Fe Kα line (see Table 4.1), as shown
in the distribution on Fig. 4.6. Moreover, we observe that both the EQ-pol and the Weak pol.
sources generally have larger luminosities than the PL-pol subsample, hinting to a similar effect
in the polarization domain, see Fig. 4.2.
Finally, regarding the sample of Weak pol., Smith et al. argue that the weak polarization
can be an indicator that the sources have null polarization intrinsic to the AGN. This scenario
is described as the line of sight aligned with the principal axis, corresponding to an unabsorbed
source in X-rays. We find results consistent with this, as warm absorption is only confirmed in
4.2. DISCUSSION 87
Figure 4.6: Equivalent width of the Fe Kα emission line vs the X-ray Luminosity.
∼ 33% of this subsample, with no source showing signs of cold absorption and ∼ 67% of them
resulting as unabsorbed.
Interestingly, Merloni et al. (2014) finds that around 30% of sources that are optically un-
obscured show X-ray absorption. This result lends support to the interpretation by Smith,
as they argue that the line of sight is grazing the edge of the torus so the cloud distribution
is less dense but the covering fraction for the X-ray corona, given its smaller size, is much
larger than the BLR. In the work cited in Sect. 4.3, the notes on individual sources report
the distance estimation between the ionized absorber and the ionizing source for some of them.
Determining the distance of each absorber requires a highly detailed analysis, based on several
assumptions since this parameter cannot be directly measured. In general, we can see that
most of the absorbers are estimated to the scales of the torus. However, the interpretation of
the outer layers of the torus by Smith corresponds to an outdated torus model. If we consider
a non-homogeneous torus with a clumpy structure of clouds, (Nenkova et al. 2002; Elitzur &
Shlosman 2006; Garćıa-Bernete et al. 2022), we could revisit the interpretation for type 1 PL-
pol sources.
88 CHAPTER 4. RESULTS AND DISCUSSION
4.3 Notes on individual sources
In this section, we briefly present properties taken from the available literature on each source
of our sample. Comparing our results on the incidence of absorption to the previous reports on
each source. These results obtained from literature publications, generally report an analysis of
higher complexity as they are derived by a dedicated analysis of individual sources that often
include the use of photionization models and grating spectra. Interestingly, we find agreements
for 14/25 of the sources regarding the presence and type of absorber affecting the X-ray emission.
We show a comparison of the reported parameters and the ones obtained from our analysis
in Figures 4.7 and 4.8. The sources depicted in these graphs are those for which we find at least
one absorber in both our analysis and in the literature. We do not consider sources resulting
in cold absorption, such as Mrk 231 which is in agreement with results from previous publica-
tions. We also omit sources for which our resulting model does not coincide with the one in
the literature. For example, Akn 120 that results affected by warm absorption in our analysis
but unabsorbed in the literature. Lastly, we did not include Mrk 1239 and NGC 3783, because
our analysis results in an undetermined absorption scenario. Overall, we find mixed results
regarding this comparison, even if for all the sources the model corresponds to a warm ab-
sorber. Sources such as NGC 7469, IRAS 15091-2107, and NGC 4593 show agreement between
previous reports and our results, while sources like ESO 323-G077 and NGC 3227 show large
discrepancies in the obtained parameters. These discrepancies occur for sources where our re-
sulting χ2
ν is larger, suggesting that the simpler model approach is not sufficient to analyze these
sources with complex spectra and affected by processes that were not considered in this project.
In general, while the incidence of absorption is mostly in agreement with previous publica-
tions, the characterization of the absorbers probes to need more detailed work. This can, in
turn, support our interpretation regarding absorbing clouds and their correlation to scattering
regions, which opens the door for future work using this project as the basis, as described in
Sect.5.2.
Figure 4.7: The column density values. The results from our analysis, in blue, compared to
those reported in the literature, in yellow.
4.3. NOTES ON INDIVIDUAL SOURCES 89
Figure 4.8: The ionization parameter. The results from our analysis are in blue compared to
the results reported in the literature in yellow.
Polar-polarized sources
Mrk 1218
Hernández-Garćıa et al. (2017) concludes that this source is best fitted by a power law with
a cold absorber of NH ∼ 9× 1020 cm−2 in the soft energy range. This result is consistent with
our analysis.
Mrk 231
This is a well studied source and, among the many publications, Braito et al. (2004) ar-
gues that the resulting flat photon index suggests a heavily absorbed spectrum consisting of
a scattered power law component and a reflected component, and even a reflection dominated
scenario. We do not consider any reflection components, but our hard band results are consis-
tent with a continuum well fitted by a power law with a flat slope and a column density of the
order of ∼ 1022cm−2.
NGC 4593
Our results for this source are consistent with previous analyses, such as Brenneman et al.
(2007); Ebrero et al. (2013); Ursini et al. (2016). In particular, Ebrero et al. (2013), worked
with grating spectra and found 4 warm absorbers of different ionization states with at least
one of high-ionization (log ξ ∼ 2.5) and determined its distance from the source of around a
few pc. This result is also consistent with the work of Ursini et al. (2016), where they report
two different warm absorbers, with the high-ionization one consistent with the column density
and ionization parameter that we found; they determined this component to be at a distance
of ≤ 3 pc from the central region; with a low ionization absorber being located at 1.5 kpc.
Fairall 51
From our test, this source shows significant improvement by adding a second warm absorber.
In fact, the work Svoboda et al. (2015) presents a detailed study of this source with data ob-
90 CHAPTER 4. RESULTS AND DISCUSSION
tained from Suzaku. Their results include the report of complex absorption that consists of
up to 3 different highly ionized absorbers plus a neutral column of NH ∼ 4 × 1022 cm−2 that
affects the hard band. An improvement in the fit is reported by allowing the covering factor
of the zxipcf model to vary. This study also considers the effects of reflection, thus producing
a much better fit via a much more complex model for this variable source. With an average
taken from the parameters resulting from two observations, the variable absorber of ionization
parameter of ξ ≈ 15erg cm−2s−1 is estimated to be located at 0.05 pc from the central region.
From the luminosity at 5100 Å of the source, the estimated size of the BLR is of 0.035 pc.
IRAS 15091-2107
Jimenez-Bailon et al. (2007a) reports a cold and a warm absorber in this source. We find our
best fitted model to be the one with a warm absorption since we do not consider a combination
of both warm and cold absorption.
Mrk 704
Our best fit for this source includes 2 warm absorbers. This result is consistent with what is
reported in Laha et al. (2011) & Matt et al. (2011), where the absorption is interpreted as the
line of sight passing close to the torus. The parameters of the absorbers are highly variable, as
reported by Matt et al. (2011), where the second observation analyzed is the same as the one
reported in this project. As an example of the complexity that these sources show, produced a
best fit model including a second soft excess element and partial covering by a cold absorber.
Was 45
Our results favor the presence of at least 2 warm absorbers. However, there are features
in the spectra that are not accounted for by our model. This source is yet to be studied in
depth, having a particularly complex spectrum whose analysis would definitively benefit from
the RGS data.
Mrk 766
We find a power law continuum on the steeper end of the typical Sy range, Γ ∼ 2.2, as
well as three different Fe emission lines and the presence of a warm absorber. This result is
consistent with the studies made by Miller et al. (2006) & Turner et al. (2007), particularly
with the reports of a warm absorber with log ξ ∼ 1 and NH ∼ 1021 cm−2, without considering
any reflection or partial covering. In a variability study, Risaliti et al. (2011) also finds warm
absorption and estimated a lower limit for the location of the absorbing clouds, corresponding
to the BLR.
NGC 3227
In our analysis, the results favor the model with 2 warm absorbers. This result is consistent
with the work by Markowitz et al. (2009), where the best fit model finds 2 warm absorbers,
besides a strong soft excess also absorbed by cold gas. Moreover, by estimating a maximum
distance, they place the ionized absorber outside the BLR. Newman et al. (2021) also finds 2
warm absorbers, as well as a partially covered power law and a reflection component.
ESO 323-G077
4.3. NOTES ON INDIVIDUAL SOURCES 91
Our model does not accurately fit the continuum, possibly due to the omission of a reflection
component and a soft excess that is modeled by more than one element. Jiménez-Bailón et al.
(2008); Miniutti et al. (2014) present a more detailed analysis of the complex spectrum of this
source. Both analyses find 2 warm absorbers as well as cold absorber that affect the power law
in the hard band. Miniutti et al. (2014) presents the idea that the variability observed in the
warm absorber can be due to a clumpy torus or clouds in the BLR and proposes this to be a
source of intermediate type between 1 and 2, being observed at an angle of ∼ 45◦.
Mrk 1239
Buhariwalla et al. (2020, 2023) report a very detailed analysis on this NLSy1. The hard
band requires a relativistic blurred reflection component besides the power law which is par-
tially covered by ionized material. This source has starburst activity, which affects and even
dominates the soft band. Since we do not account for these features, our model fails to fit the
continuum and the absorption test results are inconclusive.
Equatorial-polarized sources
Mrk 876
We report this source as unabsorbed. This result is consistent with work reported in
Bottacini (2022), using NuSTAR data, and Bottacini et al. (2014) with XMM-Newton and
Swift. This unabsorbed scenario corresponds to a classical Sy 1 and suggests that the polariza-
tion is due to regions associated with the NLR.
IZw 1
Gallo et al. (2007); Costantini et al. (2007) report on the spectral details and variability
between two observations. They find evidence of absorption of two different ionization states.
Silva et al. (2018) worked on the same observation as we did and found a variable multi-phase
ionized absorber by two gas components. The resulting parameters of the first absorber are
consistent with the absorber we report. The outflow velocity estimated for the low ionization
component is of ∼ 1900 km/s and for the high ionization is ∼ 2500 km/s. By assuming that the
outflow velocity is greater or equal to the escape velocity, they estimate the absorbers distance to
the central source at ∼ 0.07, 0.04 pc respectively, placing them in the scale of the accretion disk.
Mrk 841
Our best fit for this source corresponds to the warm absorption scenario. Longinotti et al.
(2010) find a two-phase warm absorber of log ξ1,2 = 1.5 − 2.2, 3, with the lower ionization one
corresponding to a NH ∼ 1021 cm−2. This result is obtained with high resolution data, which
can provide more detailed insights on the characteristics of the absorber. In this work, they
estimate a lower limit for the density of the absorbing gas > 103cm−3, which corresponds to the
distance of the absorbing cloud to the central source of a few tens of pc, placing the absorber
in the scale of the torus.
KUV 18217+6419
The best fit from our analysis indicates the presence of a warm absorber. This source shows
92 CHAPTER 4. RESULTS AND DISCUSSION
interesting features in the hard band, Jimenez-Bailon et al. (2007b), but we found no studies
on absorption in the soft band.
Akn 120
This source is referred to as a “bare-nucleus” AGN, and thus, there are no reports of this
source being affected by absorption. The spectrum is characterized by a strong soft excess
emission, as studied by Matt et al. (2014) and Porquet et al. (2018).
Mrk 509
This source has been largely studied, in particular in a campaign led by Kaastra et al.
(2011), where we can find reports on the many components of the spectrum. In particular,
Detmers et al. (2011) finds multiple absorption systems, with three different velocity compo-
nents. This is not consistent with our results, where this source appears unabsorbed, according
to the AIC. This is an example of how an oversimplified model like ours does not recover the
spectral complexity of the source.
Mrk 304
In our work, the AIC favors the cold absorption scenario. Our result differs from previous
studies of the same observation. We find a cold absorber of ∼ 1021 cm−2 and our baseline model
does not reproduce the convex shape of the continuum. Piconcelli et al. (2004) find that the
convex shape of the spectra corresponds to a heavily obscured source with column density up
to ∼ 1023 cm−2, and their final fit includes a multi-phase ionized absorber.
NGC 3783
According to the AIC, we are unable to determine whether the warm or cold absorber yield
a better fit for this source. There are many previous publications on this AGN including reports
on obscuring events. Based on RGS data, Blustin et al. (2002), reported the detection of a
two-phase warm absorber. This result is confirmed by Mao et al. (2019).
Weakly-polarized sources
Mrk 896
The work of Page et al. (2003) reports a model without the effects of absorption and a
power law slope of Γ = 2.03. These results are consistent with ours, with our analysis resulting
in a photon index Γ ≈ 2.2, with a χ2
ν = 1.11, also favoring the unabsorbed scenario.
Mrk 926
This source results as unabsorbed in our analysis, with a χ2
ν = 1.19. Previous work on its
spectra does not report significant absorption and, in fact, it is referred to as a “bare” AGN.
In a multi-epoch study by Chalise et al. (2022), no absorption is reported, consistent with our
results. The photon index varies among epochs, with the data fromXMM-Newton combined
with data from NuSTAR yielding a Γ = 1.88. This combined data takes into consideration the
reflection component found at higher energies.
4.3. NOTES ON INDIVIDUAL SOURCES 93
NGC 7213
Consistent with our results, Starling et al. (2005); Emmanoulopoulos et al. (2013) find this
source unaffected by absorption. In particular, Emmanoulopoulos et al. (2013) finds the best fit
to include a third Fe emission line, while we were able to fit the Gaussian profile corresponding
to Fe Kα and FeXXVI. Their best fit produces a power law index of Γ ≈ 1.9, while ours is
of Γ ≈ 1.7. Their approach includes two thermal components in addition to the power law
continuum; whilst we only added the BB for the soft excess.
NGC 7469
This is one of the two sources from this subsample that favor the warm absorption scenario.
This source has had an extensive multiwavelength analysis, with reports on detailed aspects of
its X-ray emission e.g. Behar et al. (2017); Peretz et al. (2018); Middei et al. (2018). These
reports find significant flux variability and yet a constant power law of index Γ = 1.78. In these
results, a reflection component is included, as well as soft excess modeled with a Comptoniza-
tion component. The ionized outflows found show multiphased warm absorption of ionization
parameters log ξ = 0.4, 1.6. This is among the results of characterizing absorption using RGS
data, (Middei et al. 2018). Behar et al. (2017) reports no variability of the column densities
and interprets this as the outflow originating far from the ionizing source, from 12 up to 31 pc.
NGC 7603
We did not find a publication of an in-depth analysis of this source. In our analysis, this
source favors the unabsorbed scenario, consistent with the description of possible null polariza-
tion due to the line of sight aligned with the principal axis.
PG 1211+143
This source has been extensively studied as it is associated with Ultra Fast Outflows (UFO)
(Laha et al. 2021). The work by Pounds et al. (2016a,b), among others, reports in great detail
an analysis of the complex spectrum of this source, finding highly ionized and high velocity
outflows. Our results report a warm absorber, which is far more simplistic than the actual
spectrum when considering the high-resolution RGS data.
Chapter 5
Conclusions and future work
5.1 Conclusions
We presented the results of a systematic analysis of the X-ray spectra of 25 polarized Seyfert
galaxies: 11 with PL-pol, 8 with EQ-pol, and 6 with Weak pol.. Our analysis consisted of fitting
the main components of a typical Seyfert spectrum and testing the response when absorption
was added to the model. The conclusions of this work are listed in the following bullet points.
• In regards to the general spectral shape, we find that the power law continuum and the
Fe Kα provide a good first approximation for a Sy 1 spectrum in the hard band, with
the mean value for the photon index Γ = 1.79 and the Fe Kα present in ∼ 72% of the
sample. Moreover, at lower energies, all sources require at least one additional component
to account for the soft excess, which we chose to model using a blackbody. Furthermore,
many of our sources display significant spectral complexity, with some requiring multiple
components to explain the soft X-ray emission adequately. For instance, in the literature,
sources like ESO 323-G077 have been successfully modeled by including two components
of collisionally ionized gas (Miniutti et al. 2014). This suggests the need for more complex
spectral modeling to achieve a more robust best fit model, which could then improve the
accuracy of the absorption test.
• In regards to the incidence of absorption, we find that 68% of our sample is affected by
an absorber; with 100% of PL-pol and 63% of the EQ-pol sources affected by absorption.
The difference among subsamples seems to indicate an intrinsic diversity of the scattering
medium in the two groups of sources. This is further supported by the low incidence of
absorption in the weakly polarized sources, ∼ 33%, considered by Smith as candidates
for null polarization. This suggests that optical polarization and X-ray absorption can
be related to the same material, lending support to the unification model proposed by
Smith.
• While we observe a distinction between subsamples, another interesting result emerges
when examining the entire sample of polarized Sy 1. The incidence of absorption among
the 25 Sy 1 is of ∼ 68%, with ∼ 56% confirmed to be affected by warm absorption.
The fraction of warm absorbers found in Sy 1 can reach up to 65% (Laha et al. 2021),
suggesting no significant difference in the general incidence of warm absorption. Further-
more, warm absorption is significantly more common than cold absorption, as would be
expected from type 1 sources, indicating the presence of ionized gas but not the optically
thick obscuring torus in the line of sight.
94
5.2. FUTURE WORK 95
• At first approximation, our work provides a promising test for the use of X-ray absorption
as a tool for explaining the properties of the observed polarization and their interpretation
in the context of the AGN unification model. It is desirable that further work can include
more recent polarization measurements and a larger sample of X-ray sources with known
polarization, see Sect.5.2. It is important to note that, while the model presented by
Smith can be regarded as a first approximation to a unification scheme, the availability
and quality of X-ray data and therefore our understanding of the X-ray emission of AGN
has significantly improved since Smith’s unification scheme was first proposed (2002-2005).
This poses the interesting option of taking this systematic analysis further by considering
a multi-layered absorber, a less homogeneous torus, and a more detailed characterization
of the absorbing gas.
• The X-ray absorption test could be also improved by considering a more complex modeling
of the absorbers. Considering several absorbing layers and/or a combination of neutral
and ionized absorbers can provide further insights into the physical conditions of the
absorbing material. This analysis can significantly benefit from the inclusion of high
resolution data, see Sect. 5.2.
5.2 Future work
As our analysis centered on low resolution data and a first approximation model, we find these
results very useful to set the basis of future work that explores the correlation between the
regions responsible for polarization by scattering and the ones that produce X-ray absorption.
We propose four main aspects that can enrich the results so far obtained for this project:
• A more complex model. As a first approximation model, we only included the main
components from 0.5-10 keV of the AGN X-ray spectrum. A promising idea is to further
explore the effects of ionized reflection, as some of the sources did show improvement in
their soft excess fitting by the addition of the relxill model (Victoria-Ceballos et al. 2023).
With the inclusion of data > 10 keV, provided by the NuSTAR satellite, we could improve
the fitting of ionized reflection and also consider neutral reflection, which we know can
affect the shape of the continuum. Overall, the inclusion of reflection components offers
the possibility of improving the baseline model.
• High-resolution data. The studies of high-resolution data, such as RGS in the case of
XMM-Newton and the High and Low Energy Transmission Gratings (HETG and LETG)
onboard Chandra, allow for the characterization of the absorbers. Modeling a more de-
tailed absorption scenario can further improve the statistics of the general fits but, more
importantly, provides the means for a deeper understanding of the ionized gas on the line
of sight and can help determine their location in regards to the central region.
• A larger sample. Our incidence of absorption does hint towards a relationship between
absorbing gas and scattering regions. This result can be supported by the study of a
larger sample:
X-ray regime: conducting the X-ray analysis and corresponding absorption test to
more/different sources with known optical polarization. This would provide a better
statistical foundation to confirm the observed trends in the relationship between X-ray
absorption and scattering regions. In particular, Afanasiev et al. (2019) provides polar-
ization characteristics of a sample of Sy 1 whose observations were taken between 2011
96 CHAPTER 5. CONCLUSIONS AND FUTURE WORK
and 2014. They find that the equatorial scattering scenario is dominant for a sample of
30 type 1 Sy. Of these 30 sources, 10 are already analyzed in this thesis. Out of the
remaining 20, 19 have public RGS and EPIC-pn data available. The X-ray absorption
analysis of these spectra can provide more insights into the relation between absorption
and scattering region.
Optical polarization: We did not find any other large sample of Sy 1 sources with
PL-pol. However, Smith reported 14 sources with “undetermined” polarization, i.e. they
were not classified in any subsample at the time of the studies by Smith et al. These
sources are not considered in this thesis, however, an X-ray analysis of their spectra pos-
sibly accompanied by a new study on their optical polarization can enrich our results.
Considering the lack of confirmed PL-pol type 1 sources, this study could benefit from
spectropolarimetric studies on nearby, bright Sy 1 sources. Determining the optical po-
larization of a larger sample of Sy 1 can lead to supporting our results.
Incorporating sources with both high-quality X-ray data and well-characterized polar-
ization measurements would allow for more robust modeling of the absorption and its
connection to the scattering regions.
• Considering variabilty:
1. We did not consider X-ray variability in our analysis, as our study relies on single ob-
servations per source, which are not simultaneous with the optical band observations.
However, incorporating X-ray variability studies can offer valuable insights into the
properties of the absorbers. X-ray variability, such as changes in the column density
or ionization state, can reveal important information about the characteristics of the
absorbing material, such as its location. For instance, rapid variability might indi-
cate absorbers located closer to the central engine (Krongold et al. 2007), possibly
within the BLR (Risaliti et al. 2005), whereas slower changes could be attributed to
larger scale structures like the torus (Longinotti et al. 2009). Additionally, monitor-
ing variability can help distinguish between multiple layers of absorption or uncover
transient absorption events. This deeper characterization would provide a better
understanding of the relationship between absorption, polarization, and the physical
conditions in AGN.
2. Optical polarization variability is also an interesting approach. For example, the
study by Afanasiev et al. (2019) reports as EQ-pol four sources that are reported as
PL-pol in the study by Smith et al. Thus, while the optical polarization variability
appears less common than the X-ray counterpart, it can be an aspect to be studied,
especially if having a larger sample of Seyfert sources with known polarization.
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Appendix A
Spectral Analysis
All the spectra are fitted in a range of 0.5-10 keV. First column corresponds to the baseline
model: power law + significant Fe emission lines. Second column are fits with the addition
of soft excess as black body. Third column is the resulting model for absorbed and “undeter-
mined” sources and empty for unabsorbed ones.
Polar Polarized sources
Cold Absorption
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
Mrk 1218
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
(d
at
a-
m
od
el
)/e
rro
r
PwLw
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
Mrk 1218
1 2 3 4 5 6 7 8 910
Energy (keV)
2.5
0.0
2.5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+BB
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
Mrk 1218
1 2 3 4 5 6 7 8 910
Energy (keV)
4
2
0
2
(d
at
a-
m
od
el
)/e
rro
r
ColdAbs*(PwLw+BB)
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
Mrk 231
1 2 3 4 5 6 7 8 9 10
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
Mrk 231
1 2 3 4 5 6 7 8 9 10
Energy (keV)
2.5
0.0
2.5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+BB
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
Mrk 231
1 2 3 4 5 6 7 8 9 10
Energy (keV)
2
0
2
(d
at
a-
m
od
el
)/e
rro
r
ColdAbs*(PwLw+BB)
Warm Absorption
2 4 6 8 10
10 1
100
101
co
un
ts
 s
1  k
eV
1
NGC 4593
1 2 3 4 5 6 7 8 910
Energy (keV)
0
20
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2
2 4 6 8 10
10 1
100
101
co
un
ts
 s
1  k
eV
1
NGC 4593
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2+BB
2 4 6 8 10
10 1
100
101
co
un
ts
 s
1  k
eV
1
NGC 4593
1 2 3 4 5 6 7 8 910
Energy (keV)
2.5
0.0
2.5
(d
at
a-
m
od
el
)/e
rro
r
WarmAbs*(PwLw+Fe+Fe2+BB)
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
IRAS 15091-2107
1 2 3 4 5 6 7 8 910
Energy (keV)
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
IRAS 15091-2107
1 2 3 4 5 6 7 8 910
Energy (keV)
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+BB
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
IRAS 15091-2107
1 2 3 4 5 6 7 8 910
Energy (keV)
2.5
0.0
2.5
(d
at
a-
m
od
el
)/e
rro
r
WarmAbs*(PwLw+Fe+BB)
106
107
2 4 6 8 10
10 1
101
co
un
ts
 s
1  k
eV
1
Mrk 766
1 2 3 4 5 6 7 8 910
Energy (keV)
25
0
25
50
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2+Fe3
2 4 6 8 10
10 1
101
co
un
ts
 s
1  k
eV
1
Mrk 766
1 2 3 4 5 6 7 8 910
Energy (keV)
10
0
10
20
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2+Fe3+BB
2 4 6 8 10
10 1
101
co
un
ts
 s
1  k
eV
1
Mrk 766
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
WarmAbs*(PwLw+Fe+Fe2+Fe3+BB)
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
ESO 323-G077
1 2 3 4 5 6 7 8 910
Energy (keV)
10
0
10
20
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
ESO 323-G077
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2+BB
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
ESO 323-G077
1 2 3 4 5 6 7 8 910
Energy (keV)
10
0
10
(d
at
a-
m
od
el
)/e
rro
r
WarmAbs*(PwLw+Fe+Fe2+BB)
2 Warm Absorbers
2 4 6 8 10
10 1
100
co
un
ts
 s
1  k
eV
1
Fairall 51
1 2 3 4 5 6 7 8 910
Energy (keV)
20
0
20
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe
2 4 6 8 10
10 1
100
co
un
ts
 s
1  k
eV
1
Fairall 51
1 2 3 4 5 6 7 8 910
Energy (keV)
0
20
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+BB
2 4 6 8 10
10 1
100
co
un
ts
 s
1  k
eV
1
Fairall 51
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
2WarmAbs*(PwLw+Fe+BB)
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
Mrk 704
1 2 3 4 5 6 7 8 910
Energy (keV)
0
20
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
Mrk 704
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+BB
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
Mrk 704
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
(d
at
a-
m
od
el
)/e
rro
r
2WarmAbs*(PwLw+Fe+BB)
2 4 6 8 10
10 1
100
co
un
ts
 s
1  k
eV
1
NGC 3227
1 2 3 4 5 6 7 8 910
Energy (keV)
20
0
20
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe
2 4 6 8 10
10 1
100
co
un
ts
 s
1  k
eV
1
NGC 3227
1 2 3 4 5 6 7 8 910
Energy (keV)
20
0
20
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+BB
2 4 6 8 10
10 1
100
co
un
ts
 s
1  k
eV
1
NGC 3227
1 2 3 4 5 6 7 8 910
Energy (keV)
10
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
2WarmAbs*(PwLw+Fe+BB)
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
WAS 45
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
10
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
WAS 45
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+BB
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
WAS 45
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
2WarmAbs*(PwLw+Fe+BB)
Undetermined
2 4 6 8 10
10 4
10 2
co
un
ts
 s
1  k
eV
1
Mrk 1239
1 2 3 4 5 6 7 8 910
Energy (keV)
0
10
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe
2 4 6 8 10
10 5
10 3
10 1
co
un
ts
 s
1  k
eV
1
Mrk 1239
1 2 3 4 5 6 7 8 910
Energy (keV)
0
10
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+BB
2 4 6 8 10
10 3
10 2
10 1
co
un
ts
 s
1  k
eV
1
Mrk 1239
1 2 3 4 5 6 7 8 910
Energy (keV)
0
10
(d
at
a-
m
od
el
)/e
rro
r
WarmAbs*(PwLw+Fe+BB)
Figure A.1
108 APPENDIX A. SPECTRAL ANALYSIS
Equatorial Polarized sources
Cold Absorption
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
Mrk 304
1 2 3 4 5 6 7 8 910
Energy (keV)
0
10
(d
at
a-
m
od
el
)/e
rro
r
PwLw
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
Mrk 304
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+BB
2 4 6 8 10
10 2
10 1
co
un
ts
 s
1  k
eV
1
Mrk 304
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
ColdAbs*(PwLw+BB)
Warm Absorption
2 4 6 8 10
10 2
100
co
un
ts
 s
1  k
eV
1
IZw 1
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
10
(d
at
a-
m
od
el
)/e
rro
r
PwLw
2 4 6 8 10
10 2
100
co
un
ts
 s
1  k
eV
1
IZw 1
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+BB
2 4 6 8 10
10 2
100
co
un
ts
 s
1  k
eV
1
IZw 1
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
(d
at
a-
m
od
el
)/e
rro
r
WarmAbs*(PwLw+BB)
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
Mrk 841
1 2 3 4 5 6 7 8 910
Energy (keV)
10
0
10
20
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
Mrk 841
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+BB
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
Mrk 841
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
(d
at
a-
m
od
el
)/e
rro
r
WarmAbs*(PwLw+Fe+BB)
2 4 6 8 10
10 1
100
101
co
un
ts
 s
1  k
eV
1
KUV 18217+6419
1 2 3 4 5 6 7 8 910
Energy (keV)
10
0
10
(d
at
a-
m
od
el
)/e
rro
r
PwLw
2 4 6 8 10
10 1
100
101
co
un
ts
 s
1  k
eV
1
KUV 18217+6419
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+BB
2 4 6 8 10
10 1
100
101
co
un
ts
 s
1  k
eV
1
KUV 18217+6419
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
WarmAbs*(PwLw+BB)
2 4 6 8 10
10 1
100
101
co
un
ts
 s
1  k
eV
1
Akn 120
1 2 3 4 5 6 7 8 910
Energy (keV)
10
0
10
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2+Fe3
2 4 6 8 10
10 1
100
101
co
un
ts
 s
1  k
eV
1
Akn 120
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2+Fe3+BB
2 4 6 8 10
10 1
100
101
co
un
ts
 s
1  k
eV
1
Akn 120
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
WarmAbs*(PwLw+Fe+Fe2+Fe3+BB)
Unabsorbed
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
Mrk 876
1 2 3 4 5 6 7 8 910
Energy (keV)
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
Mrk 876
1 2 3 4 5 6 7 8 910
Energy (keV)
2
0
2
(d
at
a-
m
od
el
)/e
rro
r
PwLw+BB
109
2 4 6 8 10
10 1
100
101
co
un
ts
 s
1  k
eV
1
Mrk 509
1 2 3 4 5 6 7 8 910
Energy (keV)
0
20
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe
2 4 6 8 10
10 1
100
101
co
un
ts
 s
1  k
eV
1
Mrk 509
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+BB
Undetermined
2 4 6 8 10
10 1
100
co
un
ts
 s
1  k
eV
1
NGC 3783
1 2 3 4 5 6 7 8 910
Energy (keV)
20
0
20
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2
2 4 6 8 10
10 1
100
co
un
ts
 s
1  k
eV
1
NGC 3783
1 2 3 4 5 6 7 8 910
Energy (keV)
10
0
10
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2+BB
2 4 6 8 10
10 1
100
co
un
ts
 s
1  k
eV
1
NGC 3783
1 2 3 4 5 6 7 8 910
Energy (keV)
10
0
10
(d
at
a-
m
od
el
)/e
rro
r
WarmAbs*(PwLw+Fe+Fe2+BB)
Figure A.2
110 APPENDIX A. SPECTRAL ANALYSIS
Weakly-Polarized
Warm Absorption
100 101
10 1
100
101
co
un
ts
 s
1  k
eV
1
NGC 7469
100 101
Energy (keV)
0
20
(d
at
a-
m
od
el
)/e
rro
r
AbsGal*(PwLw+Fe)
100 101
10 1
100
101
co
un
ts
 s
1  k
eV
1
NGC 7469
100 101
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
AbsGal*(PwLw+Fe+BB)
100 101
10 1
100
101
co
un
ts
 s
1  k
eV
1
NGC 7469
100 101
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
AbsGal*WarmAbs*(PwLw+Fe+BB)
100 101
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
PG 1211+143
100 101
Energy (keV)
0
20
(d
at
a-
m
od
el
)/e
rro
r
AbsGal*(PwLw+Fe)
100 101
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
PG 1211+143
100 101
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
AbsGal*(PwLw+Fe+BB)
100 101
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
PG 1211+143
100 101
Energy (keV)
0
5
(d
at
a-
m
od
el
)/e
rro
r
AbsGal*WarmAbs*(PwLw+Fe+BB)
Unabsorbed
2 4 6 8 10
10 2
100
co
un
ts
 s
1  k
eV
1
Mrk 896
1 2 3 4 5 6 7 8 910
Energy (keV)
2.5
0.0
2.5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe
2 4 6 8 10
10 2
100
co
un
ts
 s
1  k
eV
1
Mrk 896
1 2 3 4 5 6 7 8 910
Energy (keV)
2
0
2
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+BB
100 101
10 1
100
101
co
un
ts
 s
1  k
eV
1
Mrk 926
100 101
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
AbsGal*(PwLw)
100 101
10 1
100
101
co
un
ts
 s
1  k
eV
1
Mrk 926
100 101
Energy (keV)
2.5
0.0
2.5
(d
at
a-
m
od
el
)/e
rro
r
AbsGal*(PwLw+BB)
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
NGC 7213
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2
2 4 6 8 10
10 2
10 1
100
co
un
ts
 s
1  k
eV
1
NGC 7213
1 2 3 4 5 6 7 8 910
Energy (keV)
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+Fe2+BB
2 4 6 8 10
10 1
101
co
un
ts
 s
1  k
eV
1
NGC 7603
1 2 3 4 5 6 7 8 910
Energy (keV)
5
0
5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe
2 4 6 8 10
10 1
101
co
un
ts
 s
1  k
eV
1
NGC 7603
1 2 3 4 5 6 7 8 910
Energy (keV)
5.0
2.5
0.0
2.5
(d
at
a-
m
od
el
)/e
rro
r
PwLw+Fe+BB
Figure A.3
Appendix B
Spatial Analysis
The following images are organized in the following order, for each source, from top to bottom
and left to right: soft band (0.5-2 keV), hard band (2-10 keV), X-ray in the full energy band
(0.5-10 keV) with optical contours superimposed and the optical image from the DSS with the
X-ray contours superimposed. In all cases, the red x marks the point of maximum X-ray counts
and the cyan x marks the maximum in the optical.
Polar Polarized sources
111
    
  
Mrk 704 
Was 45 Mrk 766 
NGC 3227 ESO 323-G077 
 Mrk 1239
112 APPENDIX B. SPATIAL ANALYSIS
Figure B.1
    Mrk 841 KUV 18217+6419 
ELSA Gl 
Akn 120 Mrk 509 
Mrk 304 NGC 3783 
  
  
113
Equatorial Polarized sources
Figure B.2
114 APPENDIX B. SPATIAL ANALYSIS
Weakly Polarized sources
Figure B.3