institute of natural and applied sciences university … · 2019. 5. 10. · naklar, galaksi-otesi...

INSTITUTE OF NATURAL AND APPLIED SCIENCES

UNIVERSITY OF CUKUROVA

MSc THESIS

Husne DERELI

OBSERVATION OF MEDIUM-HIGH REDSHIFT ACTIVE GALACTICNUCLEI AND THEIR IMPACT ON COSMOLOGY

DEPARTMENT OF PHYSICS

ADANA, 2010

INSTITUTE OF NATURAL AND APPLIED SCIENCESUNIVERSITY OF CUKUROVA


By Husne DERELI

A THESIS OF MASTER OF SCIENCEDEPARTMENT OF PHYSICS

We certify that the thesis titled above was reviewed and approved for the award of degreeof the Master of Science by the board of jury on ...........................

Signature........................Prof.Dr. Aysun AKYUZSUPERVISOR

Signature...............................Prof.Dr.Mehmet Emin OZELMEMBER

Signature...........................................Assist.Prof.Dr. Nuri EMRAHOGLUMEMBER

This MS.c. Thesis is prepared in Department of Physics of Institute of Natural andApplied Sciences of Cukurova UniversityRegistration Number:

Prof.Dr. Ilhami YEGINGILDirector

The Institute of Natural and Applied Sciences

Note: The usage of the presented specific declarations, tables, figures and photographs eitherin this thesis or in any other reference without citation is subject to “The Law of Arts andIntellectual Products” numbered 5846 of Turkish Republic.

ABSTRACT

MSc THESIS


Husne DERELI

CUKUROVA UNIVERSITYINSTITUTE OF NATURAL AND APPLIED SCIENCES

DEPARTMENT OF PHYSICS

Supervisor: Prof.Dr. Aysun AKYUZYear: 2010, Pages: 88Jury: Prof.Dr. Aysun AKYUZ

Prof.Dr.Mehmet Emin OZELAssist.Prof.Dr. Nuri EMRAHOGLU

Active Galactic Nuclei (AGNs) are the most powerful and long-lived objectsin the Universe. They have very large bolometric luminosities up to 1048 ergs/s. Theyprovide a very attractive way to probe cosmology to high redshifts. It is thought thatthe mechanism to produce the observed properties of AGNs might be the accretion ofmatter onto a compact object. The compact object could be a black hole that has amass of order 106 - 109 M⊙. AGNs are classified into two general classes called radioquiet and radio loud ones. One of the subclasses in the radio loud AGNs, the blazars,largely dominate the high-energy extragalactic sky.

Fermi-LAT data from two blazars, 3C 454.3 and B2 1520+31, were analyzedin the present thesis, main goal being to model the AGN spectra precisely at the LATenergy range of 20 MeV - 300 GeV. Data were also exploited to build a phenomeno-logical model of medium-high redshift AGNs using the maximum likelihood methoddeveloped for the LAT. Target AGNs show (at > 100 MeV) a bent spectrum, mainly asignature of the acceleration mechanism that produces a jet. This features is probable aresult of the interaction of the emitted higher energies (> 10 GeV) γ-rays with the Ex-tragalactic Background Light (EBL). It is anticipated that observation of high energyemission from AGNs may shed light both on the acceleration processes taking placein AGNs and on the nature of EBL itself. Present work also provides a contribution tothe list of spectra of AGN sources compiled of by LAT.

Key Words: AGN: Active Galactic Nuclei, Blazar, 3C 454.3, B2 1520+31, Fermi-LAT

I

OZ

YUKSEK LISANS TEZI

ORTA-YUKSEK KIZILA KAYMA DEGERLI AKTIF GALAKTIKCEKIRDEKLERIN GOZLEMI VE KOZMOLOJIYE ETKILERI

Husne DERELI

CUKUROVA UNIVERSITESIFEN BILIMLERI ENSTITUSU

FIZIK ANABILIM DALI

Danısman: Prof.Dr. Aysun AKYUZYıl: 2010, Sayfa: 88Juri: Prof.Dr. Aysun AKYUZ

Prof.Dr.Mehmet Emin OZELYrd.Doc.Dr. Nuri EMRAHOGLU

Aktif Galaktik Cekirdekler (AGC’ler) evrende oldukca guclu ısıma yapan veuzun omurlu nesnelerdir ve 1048 ergs/s kadar varan cok buyuk bolometrik ısımaguclerine sahiptirler. Yuksek kızıla kaymalara kadar uzanan kozmolojik calısmalaraonemli katkılar saglamaktadırlar. AGC’lerin gozlenen ozelliklerini olusturan mekaniz-manın, sıkı bir cisim uzerine yıgılan maddeden kaynaklandıgı dusunulmektedir. Busıkı cismin kutlesi 106 - 109 M⊙ civarında hesaplanmaktadır. AGC’ler, genellikleradyo ısıması yapan ve yapmayan kaynaklar olarak iki genel sınıfa ayrılırlar. Radyoısıması yapan AGC’ler sınıfından bir bolumu ’blazar’lar olarak adlandırılır. Bu kay-naklar, galaksi-otesi yuksek enerjili ısımalarının kaynagı olarak one cıkarlar.

Bu tezde, iki blazarın (3C 454.3 ve B2 1520+31) Fermi-LAT verileri analizedilmistir. Calısmanın temel amacı, LAT enerji aralıgında AGC’lerin tayflarının enyuksek olasılık (maximum likelihood) yontemi ile modellenmesidir. Orta ve yuksekkızıla kayma degerli AGC’lerin gorungusel (fenomolojik) bir modeli icin LAT ver-ileri kullanılmıstır. AGC’ler (> 10 GeV) ivmelenme mekanizmasının urunu olanjetin varlıgının isareti olarak bukumu olan bir tayf sunmaktadır. Fenomolojik mod-elin anlasılması hem AGC’lerin hem de galaksi-otesi ardalan ısımasının anlasılmasınayardımcı olabilir. Yaptıgımız calısmanın, LAT tarafından belirlenen AGC kay-naklarının tayf listesine onemli bir katkı saglaması beklenmektedir.

Anahtar Kelimeler: AGC: Aktif Gokada Cekirdekler, Blazar, 3C 454.3, B21520+31, Fermi-LAT

II

ACKNOWLEDGEMENTS

I would like to thanks many people for their helps and support for this work.

First, I would like to express my gratitude to my supervisor, Prof. Dr. Ay-

sun Akyuz, whose expertise, understanding, and patience, added considerably to my

graduate experience. I am grateful for her guidance throughout my graduate study her

frequent advice and valuable comments. My special thanks are also to Prof. Dr. M.

Emin Ozel for his helpful comments and correlations on my thesis.

I will also acknowledge my external supervisor, Dr. Denis Bastieri at University

of Padua (in Padua, Italy) for his suggestions, supervision, teaching as well as provision

of basic materials for this work. My special thanks goes to Dr. Riccardo Rando, for the

motivation, he has provided. I am further grateful to him for his help during my study

in Padova. I also thank to the Fermi LAT Collaboration allowing me into cooperation.

I would also like to thank the other members of Fermi group at Padova, Miss Sara

Buson, Mr. Luigi Tibaldo, Ms. Svenja Carrigan, Mr. Gabriele Navarro, and Mr.

Matteo Balbo for the assistance they provided at all levels of my work. I would also

like to thank my friends, Mr. Sezgin Aydın, Miss Meltem Degerlier, and Mr. Luca

Silvestrin and others in Padova for their friendship and help.

My special thanks go also 11th COSPAR (in India) meeting for providing com-

fortable conditions in for my study.

I would also like to thank my family, for the support they provided me through

out this study education life and this thesis. I must also acknowledge my friend, Miss

Sevinc Mantar, without whose love, encouragement and editing assistance, I would not

have finished this thesis.

Finally, I would like thank to my colleagues and friends Miss Eda Sonbas, Mr.

Ilham Nasıroglu, Miss Sukriye Cihangir, Mr. Hasan Avdan, Mr. Abdullah Iskender,

Miss. Emine Gurpınar, the name of other my friends that i can not write and Mr. Hakkı

Gorgulu in UZAYMER for their helps during my works.

III

CONTENTS PAGE

ABSTRACT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . I

OZ . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . II

ACKNOWLEDGEMENTS . . . . . . . . . . . . . . . . . . . . . . . . . . . III

CONTENTS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . IV

LIST OF TABLES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . VI

LIST OF FIGURES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . VII

LIST OF SYMBOLS AND ABBREVIATIONS . . . . . . . . . . . . . . . . . X

1 INTRODUCTION . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 11.1 Types of Active Galactic Nuclei . . . . . . . . . . . . . . . . . . . . 2

2 PREVIOUS WORKS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 82.1 Properties of AGNs . . . . . . . . . . . . . . . . . . . . . . . . . . . 92.2 Radio-Quiet and Radio-Loud AGNs . . . . . . . . . . . . . . . . . . 92.3 Classification of AGNs: Type2, Type1 and Type0 Objects . . . . . . . 10

2.3.1 Unification Model for AGNs . . . . . . . . . . . . . . . . . . 112.4 Blazars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15

2.4.1 Blazars Properties . . . . . . . . . . . . . . . . . . . . . . . 162.4.2 Blazar Classification . . . . . . . . . . . . . . . . . . . . . . 172.4.3 Radio Galaxies and Blazars . . . . . . . . . . . . . . . . . . 20

2.5 Impact of Active Galactic Nuclei on Cosmology . . . . . . . . . . . . 25

3 BRIEF OVERVIEW AND THE METHOD OF ANALYSIS . . . . . . . . 273.1 Gamma-ray Production Processes . . . . . . . . . . . . . . . . . . . 273.2 Gamma-ray Detection Techniques . . . . . . . . . . . . . . . . . . . 313.3 Gamma-ray Telescopes . . . . . . . . . . . . . . . . . . . . . . . . . 33

3.3.1 Compton Telescopes . . . . . . . . . . . . . . . . . . . . . . 363.3.2 Pair-tracking Telescopes . . . . . . . . . . . . . . . . . . . . 38

3.4 Fermi Gamma-ray Space Telescope (FGST) . . . . . . . . . . . . . . 393.4.1 The Large Area Telescope (LAT) . . . . . . . . . . . . . . . 403.4.2 The Gamma-ray Burst Monitor (GBM) . . . . . . . . . . . . 42

3.5 The LAT Mission . . . . . . . . . . . . . . . . . . . . . . . . . . . . 43

4 INTERPRETATION OF OBSERVATIONS AND RESULTS . . . . . . . . 454.1 The Maximum Likelihood Method . . . . . . . . . . . . . . . . . . . 454.2 The Unbinned Analysis . . . . . . . . . . . . . . . . . . . . . . . . . 46

4.2.1 Diffuse γ-ray Emission . . . . . . . . . . . . . . . . . . . . . 484.3 Analysis of LAT Data for 3C 454.3 . . . . . . . . . . . . . . . . . . . 49

IV

4.4 Analysis of LAT Data for B2 1520+31 . . . . . . . . . . . . . . . . . 59

5 CONCLUSIONS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 64

REFERENCES . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 66

RESUME . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70

1 APPENDIX-A . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 711.1 Fermi Acceleration Mechanism . . . . . . . . . . . . . . . . . . . . . 71

V

LIST OF TABLES PAGE

Table 1.1. Type of galaxies under discussion with their defining properties . . 7

Table 2.1. The Unified Model for AGNs. Focusing on UV/optical properties

(emission line widths) and on radio properties (quiet/loud) it is pos-

sible to classify AGNs population as illustrated below. The observa-

tion angle and the black hole spin are probably the main causes of

this partition. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 13

Table 3.1. LAT science requirements compared with EGRET performances . . 43

Table 4.1. The broken power - law model parameters over 6 months data for

3C 454.3 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51

Table 4.2. The simple power - law model parameters for Galactic diffuse,

isotropic diffuse and 3 point sources. I is integral flux (I > 100

MeV). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51

Table 4.3. Flux and time values of 3C 454.3 as a result of likelihood fit . . . . 53

Table 4.4. Flux and time values of 3C 454.3 as a result of ASP . . . . . . . . . 53

Table 4.5. A broken power - law model parameters for the high flux state of 3C

454.3. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 55

Table 4.6. A broken power - law model parameters for the medium flux state

of 3C 454.3. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 55

Table 4.7. A simple power - law model parameters for the state of low flux data

of 3C 454.3. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 57

Table 4.8. A broken power - law model parameters for B2 1520+31 . . . . . . 60

Table 4.9. A simple power - law model parameters for B2 1520+31 . . . . . . 60

VI

LIST OF FIGURES PAGE

Figure 1.1. To illustrate the optical classification (Padovani and Urry, 1995) . . 5

Figure 2.1. A diagram of an active galaxy, with its primary components . . . . 8

Figure 2.2. The unification model for AGNs (not to the scale) (Padovani and

Urry, 1995) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 12

Figure 2.3. Fermi-LAT all-sky gamma-ray image, 3C 454.3 is indicated at the

lower left quarter . . . . . . . . . . . . . . . . . . . . . . . . . . . 23

Figure 3.1. Schematic view of synchrotron radiation . . . . . . . . . . . . . . . 28

Figure 3.2. Schematic view of bremsstrahlung . . . . . . . . . . . . . . . . . . 28

Figure 3.3. Schematic view of inverse Compton scattering . . . . . . . . . . . 30

Figure 3.4. The high energy γ-rays horizon. The shaded region is the large opti-

cal depth zone: photons at these energies from given sources at their

redshifts are significantly attenuated (Diehl, 2001) . . . . . . . . . 32

Figure 3.5. The Earth’s atmospheric transparency for electromagnetic radiation

at different energies (Diehl, 2001). The lower energy domain of γ-

rays, satellite and balloon experiments are required, while in the high

energy domain, ground based instruments are used. . . . . . . . . . 32

Figure 3.6. Schematic view of the CGRO observatory . . . . . . . . . . . . . . 34

Figure 3.7. Schematic view of a Compton telescope working principle. . . . . . 37

Figure 3.8. The EGRET telescope (Esposito,1999) . . . . . . . . . . . . . . . 39

Figure 3.9. Artist’s concept of the Fermi Gamma-ray Space Telescope (FGST) . 40

Figure 3.10.A schematic display of the LAT . . . . . . . . . . . . . . . . . . . 41

Figure 3.11.Track production in a LAT tower . . . . . . . . . . . . . . . . . . . 42

Figure 4.1. 68% containment radius versus energy at normal incidence for LAT

(left panel) and at ± 60o off-axis (right panel). The vertical line

represents our energy selection limit at 200MeV. . . . . . . . . . . 46

VII

Figure 4.2. Spectrum of γ-ray flux as a function of energy for 3C 454.3. Here,

six months of LAT data under the likelihood method, for the broken

power law model. . . . . . . . . . . . . . . . . . . . . . . . . . . . 52

Figure 4.3. (a) The light curve of real data analysis results by ’likelihood’

method. The photons (>100 MeV) are used for this analysis. Each

point indicates 2 week’s data of data. (b) The same light curve is ob-

tained by the Automated Science Processing (ASP) methods. Each

point corresponds 1 week of data . . . . . . . . . . . . . . . . . . . 54

Figure 4.4. Three flux level states for 3C454.3 are defined. Upper part is the

’high flux’, middle part is the ’medium flux’, below part is the ’low

flux’ states. The portion near 250 MET is analyzed separately. . . . 54

Figure 4.5. Spectrum for the high flux state of 3C 454.3. Break is at 4.05 GeV

and spectral differ by ∆γ = γ1 − γ2 = 1.16 . . . . . . . . . . . . . . 55

Figure 4.6. Spectrum for the medium flux state of 3C 454.3. Breaking point is

at Eg≃2.53 GeV and indices before and after the break differs by

about 0.55 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 56

Figure 4.7. Spectrum for the low flux state of 3C 454.3. . . . . . . . . . . . . . 56

Figure 4.8. This plot shows the spectra for is high, medium and low flux states

together for 3C 454.3. Break energy moves to a higher energy as we

move from medium to high flux states. For the low state, no break is

observed. ) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 57

Figure 4.9. The plot is Spectral Energy Distribution for high, medium and low

flux states as defined in the text. . . . . . . . . . . . . . . . . . . . 58

Figure 4.10.Energy spectra from the real (red circles) and simulated (green trian-

gles) data for B2 1520+31 resulting from summing the power - law

distributions with parameters flux and photon index, as measured in

weekly bins. The red dashed line represents the Power Law fits and

blue dashed line represents Broken Power Law fits. . . . . . . . . . 61

VIII

Figure 4.11.Time variation of the flux obtained with the likelihood analysis for

B2 1520+31 in the 100 MeV - 300 GeV band. Red points are real

and green points are simulated data. The LAT was operated in the

survey mode throughout these observations except during the pe-

riod 2008 August 4 - 2010 January 25, when it was operated in the

pointed mode. Each point represents one week’s data. The error bars

are statistical only. . . . . . . . . . . . . . . . . . . . . . . . . . . 62

Figure 4.12.The variation of spectral index for B2 1520+31 against flux. These

values are obtained by the likelihood analysis for B2 1520+31 in the

100 MeV - 300 GeV band. . . . . . . . . . . . . . . . . . . . . . . 63

Figure 4.13.The variation of the flux with aperture photometry analysis for B2

1520+31 in the 100 MeV - 300 GeV band. The LAT was operated

in the survey mode throughout these observations except during the

period 2008 August 4 - 2010 January 25, when it was operated in the

pointing mode. Each point is for one week’s duration of data. The

error bars are statistical only. . . . . . . . . . . . . . . . . . . . . . 63

Figure 1.1. (a) To illustrate the collisions between a particle of mass m and a

cloud of mass M. (b) To illustrate the collisions between a particle

and equal numbers of clouds moving in opposite directions in one

dimension (Longair 1983) . . . . . . . . . . . . . . . . . . . . . . 72

IX

LIST OF SYMBOLS AND ABBREVIATIONS

ACD Anticoincidence Detector

AGILE Astro-rivelatore Gamma a Immagini LEggero

AGN Active Galactic Nucleus

BATSE Burst and Transient Source Experiment

BGO Bismuth Germanate

BL Lac BL Lacertae

BLRG Broad-Line Radio Galaxies

BPL Broken Power-Low

CAL CALorimeter

CR Cosmic Ray

DAQ Data AcQuisition electronics

DGE Diffuse Galactic -ray Emission

EBL Extragalactic Background Light

EGRET Energetic Gamma Ray Experiment Telescope

FGST Fermi Gamma-ray Space Telescope

FoV Field of View

FSRQs Flat Spectrum Radio Quasars

FWHM Full Width at Half Maximum

FRI Fanaroff-Riley type I

FRII Fanaroff-Riley type II

GBM GLAST Burst Monitor

GRBs Gamma-Ray Bursts

GRID Gamma-Ray Imaging Detector

HBLs High-Frequency Peaked BL Lac

HESS High Energy Stereoscopic System

IACT Imaging Atmospheric Cherenkov Technique

ICS Inverse Compton Scattering

INTEGRAL INTErnational Gamma Ray Astrophysics Laboratory

IR InfraRed

X

IRAS InfraRed Astronomical Satellite

IRF Instrument Response Function

ISM InterStellar Medium

LAT Large Area Telescope

LBLs Low Frequency Peaked BL Lacs

MAGIC Major Atmospheric Gamma-ray Imaging Cherenkov

NASA National Aeronautics and Space Administration

NLRG Narrow-Line Radio Galaxies

OSSE Oriented Scintillation-Spectrometer Experiment

QSO Quasi-Stellar Objects

OVVs Optically Violently Variables

PMT PhotoMultiplier Tube

PL Power Low

PSF Point Spread Function

ROI Region Of Interest

SEDs Spectral Energy Distributions

SSC Synchrotron Self Compton

SSD Silicon Strip Detector

TKR TracKeR

TS Test Statistic

UV UltraViolet

VLBI Very Long Baseline Interferometry

XI

1. INTRODUCTION Husne DERELI

1. INTRODUCTION

Galaxies play an important role in explaining the formation and evolution of

the Universe. They normally emit light at all wavelengths. Generally, galaxies are

large systems of stars and interstellar matter. Typically, they contain in the range of

∼106 to 1012 stars. Their masses are ∼106 to 1012 times the Sun’s mass (M⊙). A

typical galaxy consists of a disk, a central bulge, and a galactic hole that is a roughly

spherical distribution of stars and globular clusters surrounding the disk. Usually, they

are separated by millions of light years distance.

Galaxies are classified into three main types according to their morphology as

spiral, elliptical and irregular. Beyond this, astronomers have also elaborated new,

more complex criteria of classification from their appearances. Different types of

galaxies have a rather broad energy distribution. Some of them show significant emis-

sion in the full electromagnetic range from radio wavelengths to the X-ray and gamma-

ray range. The latter types of emissions originate mainly from a very small, central

region of an active galaxy, which is called the Active Galactic Nucleus (AGNs). Such

active galaxies form many different types of AGN families which may differ in their

spectral properties, their luminosities and their ratio of nuclear luminosities to that of

the total stellar emissions.

The central engine of an AGN is a strong energy source which may be highly

variable and very bright compared to the rest of the galaxy. Most models of active

galaxies are built on the possibility of a supermassive black hole at the center of it. The

dense central galaxy provides material which accretes onto the black hole releasing a

large amount of gravitational energy. Part of the energy in this hot plasma is emitted

as X-rays and gamma rays.

AGNs are also known to be the most powerful, long-lived objects in the Uni-

verse. In many ways, AGNs are special laboratories for extreme physics. They can be

observed at significant redshifts where the Universe was only a fraction of the age it is

now. If a typical AGN can be used as a standard candle, they may provide a very at-

tractive way to probe cosmos to high redshifts. The main mysteries with AGNs are that

1


they produce very high luminosities in a very small volume, probably the size of solar

system, through physical processes other than the nuclear fusion that powers stars.

AGNs also have very large bolometric luminosities up to 1048 ergs s−1, cor-

responding to about 1000 times the luminosity of a normal galaxy. Their sizes are

estimated to be of the order of light days or less (R ∼ 1013 - 1016 cm ∼ 1-1000 AU)

(Maraschi, Tavecchio 2003). Taking into account the observed variability, the causal-

ity relation implies that most of their bolometric luminosity is produced from a region

that is extremely compact. The only mechanism to produce the observed properties

of AGNs is the accretion of matter onto a compact object. Their mass must be of the

order of 106 - 109 M⊙. This range of masses definitely excludes a neutron star as ac-

creting object; to obtain the luminosities typically observed, the presence of a central

supermassive black hole seems to be required (Rees, 1984a; Rockefeller, 2005).

In this view, a massive black hole accretes matter from the inner regions of

the host galaxy; because the infalling matter possesses angular momentum, the flow

is organized in a disk structure (the so-called accretion disk) where the matter pulled

toward the black hole by gravity and loses angular momentum through viscous or

turbulent torques. Its gravitational energy is then converted into radiation; the energy

that has to be dissipated in order to reach the disk’s inner boundary is the maximum

energy released by the accretion disk.

1.1. Types of Active Galactic Nuclei

There are mainly 3 types of active galaxies: Seyfert galaxies, quasars, and

blazars. Which name is given depends on the angle between the observer and the

object, also on the mass of the object and finally on how much mass the black hole

accretes.

The activities seen in the AGNs are caused by the gaseous matter falling into,

and interacting with the supermassive central objects. Sometimes, the spectra of these

nuclei indicate enormous gas masses in rapid motion. These galaxies with such a

nucleus are called Seyfert galaxies.

2


The first strong and broad emission lines were discovered in the galaxy NGC

1068 in 1908 by Edward A. Fath. However, only the systematic analysis by Carl

Seyfert in 1943 drew the focus of astronomers to this new class of galaxies. The cores

of these Seyfert galaxies had extremely high surface brightness, and the spectrum of

their central region was dominated by emission lines of very high excitation. Some of

these lines were extremely broad.

There are two types of for Seyfert galaxies. As was first published by

Khachikian and Weedman (1974), one can distinguish two distinct subclasses of

Seyfert galaxies, depending on the presence or absence of broad bases on the per-

mitted emission lines in their spectra. Seyfert 1 galaxies have both very broad and

narrower emission lines, where ’narrow’ still means several hundred km/s and thus a

significantly larger width than characteristic velocities (like rotational velocities) found

in normal galaxies. Seyfert 2 galaxies, however, show only one set of emission lines

which are comparatively narrow and originate from a low-density ionized gas (elec-

tron density 103 to 106 electrons cm−3) with velocity widths corresponding to several

100 km/s, which is somewhat broader than the emission lines from non-active galac-

tic nuclei. These lines are frequently referred to as ’narrow lines’ and occur for both

permitted and forbidden spectral lines. Seyfert 1 galaxies, in addition, show a set of

’broad lines’ corresponding to velocities up to 1000 km/s, occurring only for the per-

mitted lines, which indicates higher densities (109 electrons cm−3). To summarize,

Seyfert 1 show broad hydrogen emission lines and narrow forbidden lines. Seyfert 2

show only the narrower lines which are narrow hydrogen emission lines and narrow

forbidden lines.

By measuring their redshifts, it is found that Seyferts are much closer to us

than quasars or blazars. Their luminosity is considerably lower than that of quasars.

On optical images they are identified as spiral galaxies.

While Some AGNs are faint or quiet, others bright or loud in the radio; the

latter are called radio galaxies. A famous example is the radio galaxy M87. Radio

galaxies are elliptical galaxies with an active nucleus. They were the first sources that

were identified with optical counterparts in the early radio surveys. Two characteristic

3


radio galaxies are Cygnus A and Centaurus A. Similar to Seyfert galaxies, radio galax-

ies are also distinguished by their emission lines: broad-line radio galaxies (BLRG)

and narrow-line radio galaxies (NLRG). In principle, two types of radio galaxies can

also be considered as radio-loud Seyfert 1 and radio-loud Seyfert 2 galaxies, but with

different morphologies for their host galaxy. A smooth transition between BLRG and

quasars also seems to exist. They are again distinguished by their optical luminosities

similar to Seyfert galaxies.

Quasars are active galaxies which are very far away from us. Some of the

quasars we have seen so far are 12 billion light-years away. These objects were first

revealed as radio sources with no corresponding visible counterpart. They were as-

sociated with point-like objects of very small angular size, comparable in size with

stars. Nevertheless, their large luminosity and high redshifts were too high to be ex-

plained by whole stars in a galaxy. Quasars have even more exotic nuclei, which are

extremely compact and extremely bright, outshining their whole parent galaxy. They

are so rare and the nearest is so remote that the brightest of them, 3C273, about 2 bil-

lion lightyears away in the constellation Virgo, is only of magnitude 13.7, and none of

them is in Messier’s or even in the NGC or IC catalogs of Dreyer, (1888, 1895).

Most of these sources are nearly invisible in the radio domain of the spectrum;

such sources are called ’radio-quiet’. In particular, they have a blue optical energy

distribution, strong and broad emission lines, and in general a high redshift. Hence,

apart from their radio properties, these sources appear to be like quasars. Therefore,

they were called radio-quiet quasars, or ’quasi-stellar objects’ (QSOs). Today this ter-

minology is no longer very common, because the clear separation since a clear sources

with and without radio emission is not considered to be valid any more. Radio-quiet

quasars also show radio emission if they are observed at sufficiently high sensitivity.

In modern terminology, the expression QSO encompasses both the quasars and the

radio-quiet QSOs. About 10 times more radio-quiet QSOs than quasars are thought

to exist. The QSOs are the most luminous AGNs. Their core luminosity can be as

high as a thousand times that of luminous galaxy. Therefore, they outshine their host

galaxies and appear point-like on optical images. Host galaxies of QSOs of lower lu-

4


minosity were identified and some are spatially resolved by Hubble Space Telescope

(HST) (Schneider, 2006).

Most scientists believe that, even though these types look quite different to us,

they are really all the same thing viewed from different directions. Figure 1.1 illustrates

the unified model, according to which the classification of an active galactic nucleus

depends on the inclination of its symmetry axis to observer’s line of sight.

Figure 1.1. To illustrate the optical classification (Padovani and Urry, 1995)

In most QSO’s, a strong jet of relativistic particles emanates perpendicular to

the plane of the accretion disc. If we look towards the central engine along the jet axis

(<10o), i.e., basically directly into the jet, we observe a blazar. If we look along the

jet axis ∼10◦ - 20◦ we observe a quasar or a Seyfert type 1 galaxy with a flat radio

spectrum. A steep spectrum quasar or a Seyfert type 2 galaxy is observed at offset

angles of the order of ∼30◦. A typical radio galaxy, showing two oppositely aligned

jets, is observed at viewing angles perpendicular to the jet axis.

One of the special classes of AGNs is called ’blazars’. They are very bright in

the radio band. This result is from looking directly down into a jet emitting in syn-

chrotron radiation. Blazars are also a class of powerful but highly variable γ-ray emit-

ters. Unified model of AGNs describes them as super-massive black holes surrounded

by accretion disks characterized by out-flowing jets. Although they represent only

5


a small percentage of the overall AGNs population, they largely dominate the high-

energy extragalactic sky. The reason is that, most of the non-thermal power, which

arises from relativistic jets, narrowly beamed and boosted in the forward direction, is

emitted in the γ-ray band. The general properties of all types of normal and active

galaxies are shown in Table 1.1. We will concentrate on the study of blazars, which

will be explained in detail in the following chapters.

6


Table 1.1. Type of galaxies under discussion with their defining properties

Normal

galaxy

Radio

galaxy

Seyfert

galaxy

Quasar Blazar

Example Milky Way M87,

Cygnus A

NGC 4151 3C273 BL Lac,

3C279

Galaxy

type

spiral elliptical,

irregular

spiral irregular elliptical

L/L⊙ <104 106-108 108-1011 10111014 1011-1014

MBH /M⊙ 3x106 3x109 106-109 106-109 106-109

Radio

emission

properties

weak core, jets,

lobes

only≈5%

radio-loud

only≈5%

radio-loud

strong,

short-time

variable

Optical

/NIR

fully ab-

sorbed

old stars,

continuum

broad emis-

sion lines

broad emis-

sion lines

weak or no

lines

X-ray

emission

weak strong strong strong strong

Gamma

emission

weak weak medium strong strong

Variability unknown months to

years

hours to

months

weeks to

years

hours to

years

7

2. PREVIOUS WORKS Husne DERELI

2. PREVIOUS WORKS

Active galaxies are studied at all wavelengths. Since they can change their

behavior on short timescales, it is useful to study them simultaneously at all energies.

Figure 2.1. A diagram of an active galaxy, with its primary components

Recent studies resulted a consensus model with an approximate structure for

AGNs. The common picture is illustrated in Figure 2.1 (Holt et al. 1992; Padovani and

Urry, 1995). At the center is a supermassive black hole whose gravitational potential

energy is the ultimate source of the AGN luminosity. If the black hole is spinning,

energy may be extracted electromagnetically from the black hole itself. Matter pulled

toward the black hole loses angular momentum through viscous and turbulent pro-

cesses in an accretion disk. The disk glows brightly at ultraviolet (UV), and perhaps at

soft X-ray wavelengths. But the detailed physics is somewhat hidden because of their

strongly anisotropic radiation patterns.

8


2.1. Properties of AGNs

AGNs refer to the central region of a galaxy that is observed to be particu-

larly energetic, in many cases outshining the rest of the galaxy.They are mostly distin-

guished by the following characteristics :

• a compact (less than 0.1 pc) and bright nucleus. It has luminosity up to

1047ergs−1 overcoming that of the whole galaxy,

• presence of broad or narrow emission lines in the optical spectra produced by

non-stellar processes in the disk and surrounding environment,

• high variability of the electromagnetic emission, on time scales from minutes to

years,

• symmetrically opposite jets propagating from the central core possibly showing

superluminal motions,

• continuum non-thermal emission in several wavelength, from radio to γ-ray

band,

On the other hand, it is not hundred percent certain that all AGNs are very

compact, non-stellar and quite massive objects.

Despite these common features and basically simple origin for the primary en-

ergetic output, the spectral energy distributions (SEDs) of AGNs are extremely com-

plex. Depending on their spectral properties, their luminosity, and the selection criteria,

AGNs have been classified into a large numbers of classes and subclasses. Two impor-

tant classification criteria are based on the characteristic of the importance of the radio

emission and the optical emission lines.

2.2. Radio-Quiet and Radio-Loud AGNs

Historically, the first important division on AGNs was made on the basis of the

relative importance of the radio emission with respected to the optical one. Keller-

9


mann (1989) found that the radio-to-optical ratio R, of AGNs are defined as R ∼ 100

- 1000. AGNs populating these two peaks are called ’radio-quiet’ and ’radio-loud’,

respectively; the boundary line is drawn at R = 10 (Cirasuolo, 2003). The radio-loud

ones represent only 10% - 20% of the AGN population. In most of these sources, radio

observations have revealed directed outflows on scales from 1017 cm (∼0.1 ℓy) to sev-

eral times 1024 cm (∼ 106 ℓy). This outflows often were having apparent superluminal

velocities. The continuum from the infrared (IR) to the optical band of radio-loud and

radio-quiet AGNs is rather similar, suggesting that two classes have similar thermal

components (Urry and Padovani, 1995). The total spectrum of radio-loud sources is

flatter, and basically could be described as a combination of a radio-quiet emission with

a non-thermal component extending from radio to γ-rays. For a special class of radio-

loud AGNs, known as ’blazars’, this non-thermal radiation dominates the observed

continuum. In the ν vs ν Fν plane, it shows a typical ’double humped’ shape which

is characterized by an extreme variability, in particular at higher energies (Sambruna,

1996).

2.3. Classification of AGNs: Type2, Type1 and Type0 Objects

Two principal systems of optical emission lines were identified in the spectra of

AGNs (Netzer, 1990). The first were the broad emission lines mainly, from Lyα , [Mg

II] and [C IV], with typical Full Width at Half Maximum (FWHM), values & 2000

km s−1). It is also showed and variability on time scale of days to weeks to months,

depending on the ionization state. The second type of system was composed of narrow

(FWHM . 1500 km s−1) and time-constant lines. Although it is well known that they

are far too simple, the application of photoionization models could provide constraints

on the physical state of the emitting gas. In particular, for the first component, the

measure of the line ratios provides a measure of its density usually in the range n ∼

108 - 1012 cm−3. Its variability time scale implies that the typical size for this region

of gas is lower than ∼ 2 x 1018 cm (about 2 light years). In the second component, the

presence of forbidden emission lines indicates that they are produced in low density

10


regions, n ∼ 103 - 106 cm−3. Objects with broad permitted lines and narrow forbidden

lines have been called ’Type 1’ AGNs; their spectra is dominated by a UV excess,

or big blue bump. A second important component, the IR bump, emerges in the ∼

1-300 µm region. On the other hand, the optical spectrum of a ’Type 2’ AGNs show

only narrow lines and their continuum emission does not show a strong UV excess,

while the IR component is strongly enhanced. In this classification, a special case was

pointed out: the Broad Absorption Line QSOs designated as a optical spectra having

strong absorption lines. This special case, ’Type 0’, was represented by the class of

BL Lacertae objects (BL Lacs), characterized by the particular weakness or absence

(equivalent width (EW) < 5 A) of emission lines, and the presence of a strong non-

thermal continuum.

2.3.1. Unification Model for AGNs

A schematic diagram of the current understanding for AGNs is shown in Fig-

ure 2.2. These is a luminous accretion disk surrounding the central black hole. Strong

optical and UV emission lines are produced in clouds of gas moving rapidly in the

gravitational potential of the black hole. These are also called ’broad-line clouds’

(dark blobs in Figure 2.2). The optical and ultraviolet radiation from the object were

obscured along some lines of sight by a torus or warped disk of gas and dust, which is

well outside the accretion disk and broad-line region. Beyond the torus, slower mov-

ing clouds of gas produce emission lines with narrower widths (grey blobs in Figure

2.2). For a 108M⊙ black hole, the gravitational radius is ∼ 10 −5 pc (about 2 AU),

the accretion disk emits mostly from the regions ∼ (3 - 100) x 10−5 pc (6-200 AU),

the broad-line clouds are located within ∼ 0.1 pc (∼0.3 ℓy). The narrow-line region

extends approximately from 1 to a few 103 pc, and radio jets have been detected on

scales from 0.1 to several times 100 kpc.

A thick dense and dusty torus (or warped disk) obscures a big portion of the

broad - line region from transverse lines of sight; some continuum and broad - line

emission can be scattered into those lines of sight by hot electrons that pervade the

11


Figure 2.2. The unification model for AGNs (not to the scale) (Padovani and Urry, 1995)

region. A hot corona above the accretion disk may also play a role in producing the

hard X-ray continuum.

In the Figure 2.2, clouds located further out (up to few kpc; e.g. Bennert, 2002)

have lower densities (n ∼ 103 −106cm−3) and smaller velocities than the narrow-line

region with typical FWHM < 1000 - 2000 km s−1. Narrow emission lines and also

some forbidden lines owing to the relatively low electron densities (ne ∼ 1010cm−3),

are permitted. The strongest of the forbidden transitions are from ionized oxygen and

neon.

We do not observe broad emission lines in some AGNs, but we almost always

observe narrow emission lines from all AGNs and this was, the existence of a thick

dusty torus has been postulated. This feature is assumed to be located outside the

accretion disk and to obscure the broad-line region at certain orientations of the AGNs

with respect to our line of sight. Strong evidence that such a torus indeed exists comes

from direct observations of broad emission lines in the polarized scattered light of

numerous narrow-line AGNs (Antonucci and Miller, 1985).

Outflows of energetic particles are deduced to occur along the poles of the

disk or torus, escaping and forming collimated radio-emitting jets and sometimes giant

12


Table 2.1. The Unified Model for AGNs. Focusing on UV/optical properties (emission

line widths) and on radio properties (quiet/loud) it is possible to classify AGNs population as

illustrated below. The observation angle and the black hole spin are probably the main causes

of this partition.

AGNType 2

(Narrow Line)

Type 1

(Broad Line)

Type 0

(unusual spectra)

Black

Radio

quiet

Seyfert 2 Seyfert 1

QSO

Hole

Spin?

⇓Radio

NLRG

FR I

FR IIBLRG Blazars

BL Lacs

(FSRQ)

loud

SSRQ

FSRQ

decreasing angle to the line of sight ⇒

radio sources when the host galaxy is an elliptical. However, when the host is a gas-

rich spiral forming only very weak radio sources are observed. The plasma in the jets,

on the smallest scales, streams outward with very high velocities. This beams beaming

the radiation relativistically in the forward direction.

Table 2.1 shows the principal classes of AGNs (Lawrence, 1987), organized ac-

cording to their radio-loudness and their optical spectra, i.e., whether they have broad

emission lines (Type 1), only narrow lines (Type 2), or weak or unusual line emission

(Type 0). Roughly 15 - 20% of AGNs are radio-loud, meaning they have ratios of

radio (5 GHz) to optical (B-band) flux F5/FB & 10, (Kellermann, 1989), although this

13


fraction increases with optical and X-ray luminosities, reaching for example ∼ 50%

at MB . -24.5. With few exceptions, the optical and ultraviolet emission-line spectra

and the IR to soft X-ray continuum of most radio-loud and radio-quiet AGNs are very

similar and so must be produced in more or less the same way (Sanders, 1989). The

characteristic of radio-loudness itself may be related in some way to host galaxy type

or to black hole spin, which might enable the formation of powerful relativistic jets.

The radio-loud Type 1 AGNs are called Broad-Line Radio Galaxies (BLRG) at low lu-

minosities and radio-loud quasars (RLQs) at high luminosity, [either Steep Spectrum

Radio Quasars (SSRQ) or Flat Spectrum Radio Quasars (FSRQ)] depending on radio

continuum shape, with the dividing line set at αr = 0.5 where the radio spectrum is

measured at a few GHz. Radio-quiet Type 2 AGNs include Seyfert 2 galaxies at low

luminosities, as well as the narrow-emission-line X-ray galaxies. The high-luminosity

counterparts are not clearly identified at this point but likely candidates are the infrared-

luminous InfraRed Astronomical Satellite (IRAS). AGNs, which may show a predom-

inance of Type 2 optical spectra. Radio-loud Type 2 AGNs, often called Narrow-Line

Radio Galaxies (NLRG), include two distinct morphological types: the low-luminosity

Fanaroff-Riley type I radio galaxies, which have often-symmetric radio jets whose in-

tensity falls away from the nucleus, and the high-luminosity Fanaroff-Riley type II

radio galaxies, which have more highly collimated jets leading to well-defined lobes

with prominent hot spots, (Fanaroff and Riley 1974).

There are no known radio-quiet BL Lacs. A subset of Type 1 quasars, including

those defined as Optically Violently Variable (OVV) quasars, Highly Polarized Quasars

(HPQ), Core-Dominated Quasars (CDQ) or FSRQ, are probably also found at a small

angle to the line of sight. Their continuum emission strongly resembles that of BL Lac

objects (apart from the presence of a blue ’bump’ in a few cases) and, like BL Lac

objects, they are characterized by very rapid variability, with very high and variable

polarization, high brightness temperatures, often in excess of the Compton limit T

∼ 1012 K (Quirrenbach et al. 1992), and superluminal velocities of compact radio

cores. Although the names OVV, HPQ, CDQ, and FSRQ reflect different empirical

definitions, evidence is accumulating that they are all more or less the same thing -

14


that is, the majority of flat-spectrum radio quasars tend to show rapid variability, high

polarization, and radio structures dominated by compact radio cores. Hereafter we

refer to all of them simply as FSRQ. BL Lacs and FSRQ together are called blazars.

Even though the FSRQ have strong broad emission lines like Type 1 objects, they

are noted in the Type 0 column in Table 1, because they have the same blazar-like

continuum emission as BL Lac objects (Padovani and Urry, 1995).

OVVs are one subclass of QSOs, characterized by the very strong and rapid

variability of its optical radiation. The flux of OVV’s can vary by significant fractions

on time-scales of days. Besides this strong variability, OVVs also stand out because of

their relatively high polarization of optical light (typically a few percent) whereas the

polarization of normal QSOs is lower (below ∼1%). OVVs are usually strong radio

emitters. Their radiation also varies in other wavelength regions besides the optical,

with shorter time-scales and larger amplitudes as one moves to higher frequencies

(Schneider, 2006).

2.4. Blazars

Blazars are the most enigmatic class of active galactic nuclei (AGNs) which

are characterized by an extreme variability. This was explained by Blandford and Rees

(1978) in terms of highly relativistic motions of emitting particles. Subsequently, with

the introduction of the unification scenario of AGNs, blazars were interpreted as radio-

loud sources with a relativistic jet that points toward us (Massaro, 2009).

Observationally, blazars are characterized by their large amplitude chaotic vari-

ability measured in all accessible spectral bands, from radio up to TeV energies. The

variability often manifests itself as a very high flux state that lasts for months to years,

with more rapid, smaller amplitude flares superimposed on those high states. Opti-

cal and radio data show a high degree of polarization, and the radio data reveal the

presence of strong emission components that arise from extremely compact, spatially

variable structures with a core-jet morphology, often associated with apparent superlu-

minal expansion (Abdo et al., 2009).

15


2.4.1. Blazars Properties

Blazars are defined as any radio loud AGN whose jet is pointing toward us. The

source is then recognized as a blazar if it shows the signatures of a beamed emission.

In particularly, blazars have;

• compact, core-dominated radio emission with flat radio spectral index. The first

seems really a peculiar feature of the blazar phenomenon, since searches for

radio-quiet counterparts have been unsuccessful;

• a very wide broad-band non-thermal spectrum, extending from radio up to GeV

and TeV energies and are, relatively bright and luminous at any observed fre-

quency;

• strong and rapid variability at all bands, with large amplitudes, particularly in

the optical, UV and X-ray and at higher energies (MeV and GeV);

• superluminal motion of radio compact regions (called ”blobs”), as seen in Very

Long Baseline Interferometry (VLBI) images;

• one sided jets, i.e. the path of blob emission extends only to one side (a jet on

the other side is also expected to exist, but its emission is heavily dumped by the

relativistic effects, pointing in the opposite direction);

• high brightness temperatures (Tb ∼ 1011−1018K), close to or above the Compton

limit (Tb ≈ 1012K);

• stong and highly variable gamma-ray emission.

The first two properties have been fundamental in the ’hunt’ for blazars, the

most powerful tool in selecting them from large survey samples, where the number of

candidates can be very large. The cross correlation between radio and X-ray bands

hance, of these catalogues are at present one of the most efficient ways of selecting

new samples. Examples are the Deep X-ray Radio Blazar Survey, DXRBS, and the

Sedentary Survey of High Energy Peaked BL Lacs (Giommi, 1999).

16


Another characteristic that is often found is their high polarization (>1-2%).

They show a wide range of polarization levels in the optical band, varying from nearly

zero (<1%, compatible with the levels induced by dust) up to 50-60%. Because there

are sources which present all the signatures of the blazar phenomena but do not show

a polarization in the optical band. Therefore, polarization is not considered any longer

among the defining properties, even if the majority of blazars do show polarized optical

fluxes up to levels >3% (Kuhr and Schmidt, 1990).

2.4.2. Blazar Classification

Blazars are presently divided into two main classes: BL Lac objects and Flat

Spectrum Radio Quasars (FSRQs).

Historically, the first member of the BL Lac class to be discovered as BL Lac-

ertae (BL Lac). This was a compact and highly variable radio source that had been

first identified with a star by Hoffmeister in 1929 (Gasparrini, 2009). In 1968, Schmitt

noticed that a variable radio source was located at the same position of BL Lacertae

(Schmitt, 1968). The radio source VRO 42.22.01 had been detected at the Vermillon

River Observatory (in Illinois, USA) by MacLeod in 1965. BL Lacertae was not a

periodic variable, but rather its intensity varied irregularly with no apparent pattern of

its brightening or dimming. When the spectrum of this variable star was taken, it was

discovered that in the optical, it was featureless; there were no emission lines as from

quasars, and no absorption lines as found in most stars. Strittmatter and several others

identified four objects closely similar to BL Lac by 1972 (Strittmatter, 1972); therefore

this class was named as ’BL Lac Objects’. Later in 1974, Gunn and Oke, determined

that BL Lacertae was actually located in a normal elliptical galaxy (Oke and Gunn,

1974). By blocking the light from the central region of the source, light from the sur-

rounding area showed absorption lines that permitted an estimate of its redshift z ≃

0.07. This corresponds to a distance at about 420 Mpc (Hubble constant H0= 100 h0

km s−1 M pc, h0 = 0.7), indicating that the core of BL Lac shines with a luminosity L

≃ 1046 erg s−1. The discovery that some radio-loud quasars show a continuum simi-

17


lar to BL Lac Objects, but with the occurrence of broad spectral lines, was important

for a complete definition of the class of new sources. Moore and Stockman gave an

important contribution in 1981 (Moore and Stockman, 1981) when they performed a

polarization survey in which they discovered 17 high polarization quasars (HPQs) and

discussed their link to BL Lacs (Moore and Stockman, 1981). So the division of this

new class of Active Galactic Nuclei in two main subgroups, now referred to as BL Lacs

and Flat Spectrum Radio Quasars (FSRQ), was firmly established (Massaro, 2009).

Many BL Lacs turned out to have temporarily weak emission lines especially

when in a faint state, and this started to blur their distinction from quasars. Therefore,

the first survey, the Einstein Medium Sensitivity Survey (EMSS) at X-ray frequencies,

introduced a well-defined limit on their emission line strength. The EMSS, begin to

sample serendipitous X-ray sources from pointed observations, covered an area of ∼

700 square degrees with (0.3 - 3.5 keV) X-ray flux limits down to a few 10−13 erg

cm−2 s−1 and yielded 44 BL Lacs (Gasparrini, 2009).

BL Lacs are a subclass of AGNs with very strongly varying radiation, and they

do not have strong emission and absorption lines. The optical radiation of BL Lacs is

highly polarized. Since no emission lines are observed in the spectra of BL Lacs, the

determination of their redshift is often difficult and sometimes impossible. In some

cases, absorption lines are detected in the spectrum which are presumed to derive from

the host galaxy of the AGN and are then identified with the redshift of the BL Lac.

BL Lac objects are constituted with featureless optical spectra and FSRQs in

which, typically, there are prominent spectral lines. Both classes show high time vari-

ability over the whole electromagnetic spectrum, from radio waves to TeV energies,

coupled with high polarization detected in the radio and optical bands. Another dif-

ference among these two classes is that BL Lacs do not exhibit apparent cosmological

evolution and are observed at redshifts z < 1, while FSRQs are observed up to z ≃ 5.

SED is characterized by two broad bumps, interpreted as two emission components.

The former component typically peaks from the IR to the X-ray band, and the sec-

ond one in the γ-rays up to TeV energies. A possible classification criterium for BL

Lacs in terms of the SED peak energy position of the first component was proposed

18


by Giommi and Padovani (1995). They named high-frequency peaked BL Lac (HBLs)

objects those in which the synchrotron peak is between the UV band and the X-rays,

and low-frequency peaked BL Lacs (LBLs), that show the first bump in the IR-optical

range. Another BL Lac subclass was introduced subsequently to distinguish the so

called intermediate BL Lac (IBLs) objects with the transition between the first and the

second component in the keV range.

The discovery of intense medium-energy γ radiation from over 60 Blazars with

the Energetic Gamma Ray Experiment Telescope (EGRET) instrument on board the

Compton Observatory (Hartman, 2001) showed that non-thermal γ-ray production is

an important dissipation mechanism of their jets. This scenario was enriched by the

discovery of the TeV emission of Mrk 421 that was the first extragalactic source de-

tected at these energies in the range by the Whipple (very high energy photon) tele-

scopes (Punch, 1992, Petry, 1996). Subsequently, with the advent of other Atmo-

spheric Cerenkov Telescopes (CAT), like High Energy Stereoscopic System (HESS)

and MAGIC (Major Atmospheric Gamma-ray Imaging Cherenkov Telescope), about

twenty BL Lacs have been detected, and recently one FSRQ was also discovered as a

TeV emitter.

Finally, the most recent results of the Auger atmospheric fluorescence system

seem to indicate Blazars as the sources of the highest energy cosmic rays (CRs). Usu-

ally, SEDs of a BL Lacs are interpreted in terms of Synchrotron Self Compton (SSC)

models in which synchrotron photons, emitted by a population of electrons accelerated

in the relativistic jets, are scattered into the second component via inverse Compton

(IC) scattering by the same electrons. The SEDs of FSRQs, typically, require other

spectral components, as for example soft seed photons produced in regions external

to their jets, in order to account for their high energy emission. These emission mod-

els are generally named as External Compton (EC) radiation. An important issue that

raised new investigations on Blazars in the recent years concerns their spectral shape

and its evolution. About 20 years ago Landau (1986), studying a sample of LBLs over

a very broad frequency range, noticed that the SEDs of BL Lacs appear to be curved

and that the best description could be given in terms of a parabolic fit on a double-

19


log plot. This spectral shape is also known as log-normal distribution. More recently,

the log-parabolic model was also used to describe the X-ray spectra of HBLs such as

Mrk 421 and Mrk 501 (Tanihata, 2004, Massaro, 2004). In particular, Massaro (2004)

first attempted an interpretation and showed that this distribution can be understood in

terms of statistical acceleration mechanisms. The idea that the log-parabolic spectral

shape is not only a good and simple empirical model to fit the SEDs of Blazars, but

that can provide important clues to understand the physical conditions and acceleration

mechanisms in their jets. On the other hand, the fact that this distribution, in principle,

can be obtained as a solution of the diffusion equation for relativistic particles, can be

traced back to the early works on the physics of radio sources in the classic paper by

Kardashev (1962) (Massaro, 2009).

2.4.3. Radio Galaxies and Blazars

BL Lacs and Flat Spectrum Radio Quasars are strong radio sources character-

ized by their distinct optical spectra. Whereas in BL Lacs we observed have no or very

weak emission lines and their continuum emission is usually fitted by a power-law or

by a logarithmic parabola, FSRQs exhibit both strong narrow and broad emission lines.

The current unification scheme attempts to combine radio galaxies with BL Lacs and

quasars, assuming that the latter are radio-galaxies in which the jet points in a direc-

tion very close to the line of sight which means that their actual member would be

much higher. Relativistic effects amplify the non-thermal continuum produced in the

jet, producing all the peculiar characteristic observed in these sources. In particular

BL Lacs object appear to be the beamed versions of FRI radio-galaxies and quasars of

FRII.

Radio-loud AGNs are generally found to reside in luminous ellipticals

(McLure, 1999, Urry, 2000) which supports the unification of blazars and radio galax-

ies in general but does not provide a test for the BL Lac/FRI and quasar/FRII as-

sociations in particular. Studies of environmental properties of radio loud AGN are

somewhat inconclusive. Quasars and FRII are found to reside in clusters of similar

20


richness (Wold, 2000). On the other hand, Wurtz (1997) found for a relatively large

sample (45 sources) of BL Lacs that, their environments are more similar to those of

quasars and FRIIs rather than FRIs. However, it is not clear if the environments of

FRIs and FRIIs differ at all. Prestage and Peacock (1988) found that, for a sample of

∼ 200 radio sources with redshift z < 0.25, FRI radio galaxies laid in richer clusters

than FRII radio galaxies. But at higher redshift (z ≃ 0.5) their environments were

found to be similar (Hill and Lilly, 1991). More recent studies, albeit for much smaller

samples of radio-loud AGN, conclude that they have similar cluster environments also

at low redshifts, z ≃ 0.2 (McLure and Dunlop, 2001). In any case a common result of

these studies is that the cluster properties of all types of radio-loud AGNs span a large

range.

Quasars are found to have extended radio powers and morphologies typical

of FRII radio galaxies (Fernini, 1997). BL Lacs, however, can have extended radio

powers typical of both FRIs and FRIIs (Cassaro, 1999, Rector and Stocke, 2001).

Regarding their narrow emission lines, these are relatively weak or absent in FRI radio

galaxies as observed for BL Lacs (one of their defining criteria). On the other hand,

quasars have (by definition) strong narrow emission lines, and these can be both weak

and strong in FRII radio galaxies (Tadhunter, 1998). Therefore our current view that

BL Lacs are solely beamed FRI radio galaxies appears problematic (Gasparrini, 2009).

Applying this condition to the luminosity functions as a well defined sample,

Urry and Padovani (1995) were able to calculate the beaming properties of the blazar

populations: they found that the Lorentz factor Γ , both in FRI and FRII, lies in range

5 < Γ < 40. This result is in agreement with the values found with measures of the

superluminal motion (Vermeulen and Cohen, 1994).

Urry and Padovani (1995) have successfully applied this model also to the ob-

served radio luminosity functions of quasars (FSRQ and SSRQ) and FRII radio galax-

ies from the 2 Jy sample. A similar test of the BL Lacs/FRI unification scheme is more

subtle. BL Lacs not only are much rarer than quasars and, therefore, complete samples

suffer from small number statistics, but their almost featureless spectra make a redshift

determination often difficult. Nevertheless, comparison of luminosity function of BL

21


Lacs and FRI radio galaxies at radio, optical and X-ray frequencies showed (within the

errors) good agreement with the beaming hypothesis (Padovani and Urry, 1990, 1991;

Urry, 1991).

Many Blazars have been detected in γ-rays by EGRET of CGRO and by Whip-

ple and other Imaging Atmospheric Cherenkov Technique (IACTs) at very high en-

ergies. The most intriguing results are their short time variability and the inferred

huge γ-ray luminosity. These facts are explained assuming that we are viewing almost

along the axis of a relativistically out-flowing plasma jet. Seyfert and radio galaxies

were detected by OSSE and BATSE detectors of CGRO at energies between 50 and

150 keV. The only radio-galaxy detected at high energies by previous generation satel-

lites or IACTs was Centaurus A, detected at MeV energies by COMPTEL and above

100 MeV by EGRET.

After successfully launched, the Fermi-GLAST observatory ushered in a new

era of observational astronomy in the energetic γ-ray band. The Fermi-Large Area

Detector (LAT) is designed to address several different scientific objects which will

enlarge our knowledge of the γ-ray sky. The LAT is an imaging, wide field of view

high-energy pair conversion telescope with energy range from ∼ 20 MeV to 300 GeV,

and surveys the whole sky every three hours (Atwood, 2009). It aims at a detailed

study of different astrophysical topics such as: galactic sources, which are pulsars, su-

pernova remnants, X-rays binaries and sources of the solar system, diffuse sources and

molecular clouds, extra-galactic sources such as galaxy clusters and AGNs. Moreover,

the progress in several areas requires multi-wavelength observations with both ground

and space-based telescopes. In particular, with regard to the higher energy range cov-

ered by the LAT data, an important and crucial topic is the calibration of the Cherenkov

telescopes (such as MAGIC, HESS and the forthcoming CTA (Cherenkov Telescope

Array)) in order to merge the information collected in the GeV-TeV energy range by

different instrument and to extend our understanding of the TeV energy range (Tibaldo,

2007). LAT’s predecessor EGRET had indicated that the most prominent extragalactic

γ-ray sources are blazars, a subclass of active galactic nuclei.

In particular, I have studied, using LAT observations, two of these quasars 3C

22


Figure 2.3. Fermi-LAT all-sky gamma-ray image, 3C 454.3 is indicated at the lower left

quarter

454.3 and B2 1520+31 in this thesis. Both sources are an optical violent variable

(OVV) like 3C 279 that is one of the brightest blazars discovered to emit in gamma-

ray band by EGRET (Giuliani, 2008). 3C 454.3 (PKS 2251+158) is a quasar/blazar

located off the galactic plane. This object has been undergoing pronounced long-

term outbursts since 2000, and was remarkably active in 2005, when it reached the

largest apparent optical luminosity ever recorded from an astrophysical source apart

from GRBs (Fuhrmann, 2006, Villata, 2006). It lies some 7.1 billion light-years away

in the constellation Pegasus and is currently (2010) undergoing a flaring episode that

makes it particularly bright, especially in the gamma-ray part of the spectrum. 3C

454.3 has a redshift z=0.859 and is a well known Flat Spectrum Radio Quasar (FSRQ)

(Vercellone, et al., 2008). It was detected above 100 MeV several times in the γ-ray

band by the EGRET telescope, with an average photon index of Γ = 2.2 (Hartmann,

1999). In 2005, it underwent a very active phase in optical and X-ray bands, triggering

intensive observations in the radio, optical and X-ray bands (Villata, 2006, Giommi,

2006, Pian, 2006). During the summer of 2007, 3C454.3 was active again, reaching a

level of the optical emission comparable to that of 2005.

Several observations of 3C 454.3, in the optical, X-ray and γ-ray bands were

carried out: Kungliga Vetenskapsakademin (KVA), optical-UV: Swift / UVOT, X-

ray : Swift / XRT (Giommi, 2006), GeV band: AGILE / GRID, in the soft γ- ray

23


bands: by OSSE (McNaron-Brown 1995) and COMPTEL (Zhang, 2005). The AGILE

(Astro-rivelatore Gamma a Immagini LEggero) satellite (Tavani et al. 2008), while still

continuing its science verification phase, detected an intense emission from 3C 454.3

(Vercellone et al. 2008a). Triggered by these observations, the MAGIC observed 3C

454.3 in July and August 2007. Another γ-ray active phase was recorded by AGILE

in November-December 2007 (Vercellone, 2008b, Vercellone, 2009), which triggered

further observations with MAGIC during that period. No signal in 2009 was detected

by MAGIC and the other IACTs.

An excellent correlation is found between the IR, optical, UV and gamma-ray

light curves, with a time lag of less than one day for 3C 454.3. This source shows a

very strong, correlated variability between the peak of the synchrotron component (at

IR, optical, and UV wavelengths) and the peak of the gamma-ray component. On the

contrary, no such correlation is seen between X-rays and any other band (Bonning et

al. 2008).

AGILE performed the most intensive and long term campaign on 3C 454.3 dur-

ing May-June 2008, resulting in a continuous 50-days long monitoring, collecting also

data with WEBT (Whole Earth Blazar Telescope), REM (Rapid Eye Mount), Swift

and RXTE (Rossi X-Ray Timing Explorer ). This long observation showed the highly

variable nature of 3C 454.3, not only on short but also on longer time scales (Marisaldi

et al., 2009).

As expected, 3C 454.3 was detected easily by Fermi (Tosti, 2008). Owing to

its high flux state, it was possible for the LAT to measure its variability properties on

time scales less than a day. LAT observations indicate that the most prominent extra-

galactic energetic γ-ray sources are blazars, a sub-class of active galactic nuclei whose

overall flux is dominated by emission from a relativistically inner (≤ pc) jet. An early

LAT observation of 3C 454.3 has highlighted the capabilities of the instrument. This

observation drove also further constrains on the emission mechanisms and structure of

the object.

The LAT also observed an increasing gamma-ray flux from a source position-

ally consistent with B2 1520+31 at (RA: 15 22 09.99, Dec: +31 44 14.4), (Beasley et

24


al. 2002), since 20 April 2009. This object is known to be an OVV in Flat Spectrum

Radio Quasar (FSRQ) with a redshift of 1.487 (Emmerd et al. 2005). Preliminary

analysis indicates that on 20 April 2009, the source was in a high state with a gamma-

ray flux (E>100MeV) of 1.0 +/-0.3 x 10-6 ph cm−2 s−1 (errors statistical only) on a

time scale of a day and reached a value of 1.9 +/- 0.7 x10-6 ph cm−2 s−1 (errors sta-

tistical) in the 6-hour interval starting at 06:00 UT of the same day. During this period

the source had a flux around 4 times greater than the average flux reported in the LAT

Bright AGN Source List on first three months (Abdo et al. 2009). Because Fermi op-

erates in an all-sky scanning mode, regular gamma-ray monitoring of this source will

continue, in the future.

2.5. Impact of Active Galactic Nuclei on Cosmology

The past few decades have led to a gigantic development in the field of cosmol-

ogy. Early sky surveys aiming at collecting numerous galaxies revealed the first indi-

cations of cellular structures in the distribution of local galaxies. Due to their restricted

luminosities, galaxies can only be observed in the nearby Universe and more powerful

sources are needed in order to study the large-scale structures in the far Universe. Ac-

tive Galactic Nuclei (AGNs) generate the required luminosities and are useful sources

mapping structures in the deep Universe. Especially, the lack of samples at intermedi-

ate distances, between the local Universe and the epoch when the number density of

AGNs peaked, has prevented us from making inferences on structure evolution.

A more through understanding of active galaxies will allow astronomers to

make measurements of the Extragalactic Background Light (EBL). The EBL is mainly

made up of infrared light produced by stars and hot dust, and the intensity of the

EBL contains information on the rate of star formation when the Universe was much

younger than it is today. Blazars can be used to measure the intensity of the EBL

by estimating how much of the high energy gamma-ray emission is lost between the

source and the Earth due to collisions with EBL photons (Tibaldo, 2007).

Photons above 10 GeV can probe the era of galaxy formation (z ≥ 1.0) through

25


absorption by near UV, optical, and near IR extragalactic background light (EBL). The

EBL at IR to UV wavelengths is the accumulated radiation from structure and star for-

mation. Main contributors of its subsequent evolution in the universe are the starlight

in the optical to UV band, and IR radiation from dust reprocessed starlight (Primack

et al. 2001; Hauser and Dwek 2001). Since direct measurements of EBL suffer from

large systematic uncertainties (due to contamination by the bright foreground, e.g., in-

terplanetary dust, stars and gas in the Milky Way, etc.), the indirect probe provided by

the absorption of high energy γ-rays from blazars via pair production (γ + γ → e+ +

e−), during their propagation in the EBL fields, can be a powerful tool for probing the

EBL density, and its subsequent evolution.

26

3. BRIEF OVERVIEW AND THE METHOD OF ANALYSIS Husne DERELI

3. BRIEF OVERVIEW AND THE METHOD OF ANALYSIS

3.1. Gamma-ray Production Processes

γ-rays are high energy electromagnetic radiation produced by the decay of ex-

cited nuclei. In astrophysics we call γ-rays the electromagnetic radiation at energies

above 0.5 MeV.

In the astrophysical context many processes can produce γ-rays, both thermal

and nonthermal (Longair, 2004).

Blackbody radiationAlso called the thermal radiation, it is produced by a large population of elec-

tromagnetically interacting particles and fields in equilibrium forming a black body.

The spectrum, characterized by the temperature T , follows the Planck distribution

where the intensity I at the frequency ν is given by

I(ν) =8πhν3

c21

e−hν/kBT −1(3.1)

The radiation spectrum has a peak at a wavelength λmax

ε ≃ λmaxT (3.2)

with the constant ε = 2.898x10−3 m K, so the radiation peak is in the optical region for

T ∼ 6000 K (about the Sun’s surface temperature). For γ-rays of 1 MeV, T is 2x109 K.

The most typical γ-ray sources are in fact powered not by thermal, but by nonthermal

processes such as synchrotron radiation and others.

Synchrotron radiationWhen electrons encounter a magnetic field, they spiral along the field lines in

a helical path. This means that their direction is constantly changing, and hence they

are accelerating and therefore they will emit radiation as shown in Figure 3.1. This

radiation is called synchrotron radiation.

We could define νg the gyration frequency of a charged particle of mass m and

charge q moving in a magnetic field B, with a pitch angle θ between the particle speed

and the magnetic field direction:

27


Figure 3.1. Schematic view of synchrotron radiation

νg =qB

2πmsinθ (3.3)

The accelerated particle will emit photons with a peak at the frequency

νs =32

γ2νg (3.4)

where γ is the particle Lorentz factor. An electron with an energy of 1 GeV in the

interstellar magnetic field (∼ 1 µG) radiates synchrotron photons in the radio band.

Synchrotron radiation may occur in the UV or X-ray region, and can reach γ-ray ener-

gies in extreme cases such as on the surface of neutron stars with B & 1010 G.

BremsstrahlungWhen a charged particle is decelerated by an electric field, it emits pho-

tons as shown in Figure 3.2. This phenomenon is called the breaking radiation, or

’bremsstrahlung’. For example, an electron passing near (interacting with) a nucleus

changes its trajectory and emits radiation. We must distinguish between thermal and

nonthermal bremsstrahlungs, that is the emission from charged particles in thermal

equilibrium and emission from accelerated particles through a nonthermal process re-

spectively.

Figure 3.2. Schematic view of bremsstrahlung

Bremsstrahlung from energetic ions colliding with ambient electrons or nuclei

is always negligible, so the most important process for γ-ray production at low energies

28


is the interaction of energetic electrons with low energy nuclei, and at higher energies,

interaction of electrons with both high energy nuclei and electrons.

The dominant luminous component in a cluster of galaxies is the emission from

the hot (T∼107 to 108 K) intracluster medium. This emission is characterized by ther-

mal bremsstrahlung. Thermal bremsstrahlung radiation occurs when the particles pop-

ulating the emitting plasma are at a uniform temperature and are distributed according

to the Maxwell - Boltzmann distribution given by

f(υ) = 4π( m

2ßkT

)3/2υ2 exp

(−mυ2

2kT

)(3.5)

where speed, υ , is defined as

υ =√

υ2x +υ2

y +υ2z (3.6)

The bulk emission from such a gas is in thermal bremsstrahlung. The energy

emitted per cubic centimeter per second also called the ’free-free emission’ source

function can be written in the compact form

ε f f = 1.4x10−27T 1/2neniZ2gB (3.7)

with cgs units [erg cm−3 s−1]. Here ’ff’ stands for ’free-free’ interaction, 1.4x10−27

cm−2 s−1 is the condensed form of the physical and geometrical constants associated

with integrated power per unit area per unit frequency, ne and ni are the electron and

ion densities, respectively. Z is the number of protons of the bending ion, gB is the

frequency averaged Gaunt factor and is of the order unity, and T is the global X-ray

temperature in the medium determined from the spectral cut-off frequency.

~ν = kT (3.8)

above which exponentially small amount of photons are created because the energy

required for creation of such a photon is available only by electrons in the tail of the

Maxwell distribution.

29


This process is also known as ’bremsstrahlung cooling’ since the plasma is

optically thin to photons at these energies and the energy radiated is emitted freely into

the universe.

Inverse Compton Scattering (ICS)When a low energy photon interact of with a high energy particle (usually, an

electron), this phenomenon is called Inverse Compton Scattering. In inverse Compton

scattering, a high energy electron transfers both energy and momentum to a lower

energy scattering photon is shown in Figure 3.3.

Figure 3.3. Schematic view of inverse Compton scattering

The interaction results in the production of a γ-ray, with a typical energy:

Eγ ≃ 1.3(

Ee

TeV

)2( Eph

2 ·10−4 eV

)GeV (3.9)

Inverse Compton Scattering (ICS) plays an important role in the galaxies in

regions of high photon density.

Relativistic electrons can easily produce low energy photons by synchrotron

radiation. Then these photons interact, via ICS with electrons emitting them, in the so

called Synchrotron Self Compton (SSC) mechanism. In this case, synchrotron photon

frequency is νs ∝ B ·E2e , while the resulted γ-photon will have a frequency νIC ≈ νs ·

E2e ∝ B ·E4

e .

Nuclear transitionsEnergy levels of nuclei have typical spacings of ∼ 1 MeV in magnitude. Ra-

dioactive decay of a resultant nucleus or energetic interactions can produce an excited

state X∗, which generates a γ-ray through the decay

X∗ → X+ γ (3.10)

30


The most important lines for γ-ray astronomy are by 12C at 4.438 MeV, by 16O

at 6.129 MeV and by 26Mg at 1.809 MeV. The cross section for excitation into these

levels is maximized in resonances, so nuclear lines trace low energy cosmic rays (CRs)

and tell us about nucleosynthesis in different regions.

Decays and annihilationThe most important decay for γ-ray production is the neutral pion (π0) decay.

Pions are generated by strong interaction during collision of high energy CRs with

ambient gas or nuclei. The neutral pion π0 decays rapidly into two γ-rays, with an

energy distribution peaked at about 70 MeV in the pion rest frame. The π0 decay

bump is offsetted and broadened by the momentum distribution of the high energy

collisions producing pions.

Particle-antiparticle annihilation also produces γ-rays. The lightest pair is the

electron-positron one, which gives two or more photons with a total energy of 1.022

MeV in their center of mass frame. In the same way, hadronic and maybe some exotic

particle-antiparticle pairs can annihilate and generate γ-rays.

Effect of gamma ray propagationLow energy γ-rays cross a long path of interstellar space without hardly any

interactions. Over large distances γ-rays can be absorbed through interaction with

low energy photons, such as cosmic microwave background, IR radiation or starlight,

producing e+ e− pairs. Therefore the horizon for a γ-ray of energy Eγ is defined by the

pair production process, which is possible only above a threshold energy of

Eth =2m2

ec4

[1− cosϕ ](1+ z)2Eγ≃(

1+ z4

)−2

·(

30 GeVEγ

)eV (3.11)

if ϕ is the photons scattering angle and z is the source redshift. In Figure 3.4, the γ-ray

horizon is shown for different energies and redshifts.

In addition, the attenuation of very high energy γ-rays is possible in regions

of high density galactic interstellar radiation field (ISRF) such as the galactic center

(Tibaldo, 2007). Relative amount of attenuation can be estimated by comparing the

energy density of the medium with that of interstellar medium.

3.2. Gamma-ray Detection Techniques

γ-rays interact with the upper atmosphere of the Earth and produce electromag-

netic showers. Therefore we can divide γ-ray astrophysics in different domains. The

31


Figure 3.4. The high energy γ-rays horizon. The shaded region is the large optical depth

zone: photons at these energies from given sources at their redshifts are significantly attenuated

(Diehl, 2001)

space-based γ-ray astrophysics ranges from 500 keV to about 300 GeV, where high

energy photons are detected directly from space with satellites or balloon experiments.

At energies & 100 GeV however, the ground-based γ-astrophysics detects are more ef-

fective, due to the shower development in the Earth’s atmosphere. Then higher energy

detectors are placed at the Earth’s surface like is shown in Figure 3.5.

Figure 3.5. The Earth’s atmospheric transparency for electromagnetic radiation at different

energies (Diehl, 2001). The lower energy domain of γ-rays, satellite and balloon experiments

are required, while in the high energy domain, ground based instruments are used.

Photons above some eV’s can be detected only after interaction through pho-

toelectric absorption, Compton scattering or pair production, according to the energy

range. In these three processes, charged particles are produced, and then these particles

32


detected. A γ-ray detector has to be dense enough to have a high interaction probabil-

ity to create a charged particle. Also it has to convert a significant fraction of charged

particle’s energy into a measurable form.

Observations of celestial γ-rays in space are complicated by the low fluxes and

the high background levels, from both charged and neutral particles. Charged particles

are mainly electrons, protons and nuclei. They form a background which needs to be

effectively rejected. This can be achieved by surrounding the detector with a thin scin-

tillator. In addition, there is an external neutral background, from chargeless particles

like pions and neutrons produced in interactions between CRs and the atmosphere: it

can be dealt with by using cuts in direction by using effective absorbers and/or other

arrangements. The internal neutral background is due to instrument’s materials acti-

vation from CRs and atmospheric neutrons and decay of natural radioactive elements.

These can be minimized by lowering the mass of detector structure.

As was explained above, γ-rays entering the Earth atmosphere produce elec-

tromagnetic and particle showers. Their detection is mainly carried out using the

Cherenkov light produced by electrons and positrons in the shower: a charged particle

moving through a transparent medium faster than the speed of light in that medium

emits photons. There is a threshold energy for this process (known as Cherenkov radi-

ation) given by

Eth = mc2 ·(

n√n2 −1

)(3.12)

if n is the medium refraction index. Photons are emitted in a very short time (∼ ps) in

a cone along the particle direction with an aperture angle θ given by

cosθ =1

βn(3.13)

3.3. Gamma-ray Telescopes

A gamma ray telescope aims at not only to detect photons but also to mea-

sure their direction, energy and arrival time. Early detection attempts of γ-rays in the

primary cosmic radiation were performed with balloon and rocket experiments in the

1940s and early 1950s. Also, there was a person form Turkey who is Hakkı Ogelman

33


between experimenters into about 1960 (Ozel, 1978).Measurements with early instru-

ments were often limited by low counting statistics or systematic uncertainties. De-

pending on the energy range impressive progresses were made with new imaging tech-

niques (Lichti and Georgii, 2001).

Imaging via collimators are used in γ-ray in gamma ray telescopes below

some MeV.

For X-rays, imaging is usually achieved by using collimators which define the

instrument field of view. These collimators are arrays of tubes with absorbing walls,

so only X-rays with a path parallel to the tube can reach the detector. Because of the

high penetration power of γ-rays, these collimators can not be used at γ-ray energies:

too much material would be needed and too much background radiation would be

produced.

In later years, actively collimated γ-ray telescopes, i.e. telescopes collimated

with an active shield, like a plastic scintillator, have been developed and successfully

used: Third Orbiting Solar Observatory (OSO-3), the γ-ray experiment on board of the

Solar-Maximum Mission (SMM), the Gamma Ray Imaging Spectrometer (GRIS) and

the Oriented Scintillation-Spectrometer Experiment (OSSE) on the Compton Gamma

Ray Observatory (CGRO). We will concentrate or the latter for a successful applica-

tion.

The Oriented Scintillation-Spectrometer Experiment, OSSE, was one of

the four instruments on the CGRO the of National Aeronautics and Space Adminis-

tration (NASA) is shown in Figure 3.6. This large observatory was operated from

1991 to 2000, giving the first complete γ-ray sky survey, at several γ-ray energy bands.

Figure 3.6. Schematic view of the CGRO observatory.

OSSE was a low energy γ-ray detector and was consisted of four identical de-

34


tectors, that could be independently rotated within 192◦. The main element of each

detector was a phosphor-sandwich consisting of a NaI(Tl) crystal with a diameter of

33 cm and a 7.6 cm thick CsI(Na) crystal at the rear side. A passive tungsten collimator

was placed on the front side, defining a field of view of 3.8◦×11.4◦. This detector was

surrounded by an annular shield of NaI(Tl) crystals. The different scintillation decay-

time constants of NaI and CsI were used to distinguish γ-rays from above and below.

The spectral resolution was controlled via the PhotoMultiplier Tubes (PMTs) voltage,

using a calibration source of 60Co. OSSE had an effective area of 400 cm2 around 1

MeV, with an energy resolution of about 5-10% (Lichti and Georgii, 2001).

Imaging via modulation techniques using in a γ-ray telescope with a wide

field of view can reach a good angular resolution modulating the signal from the source

to the detector. The best modulation technique uses arrays of opaque and transparent

elements arranged in a regular pattern, called coded masks. The most important coded

mask telescopes have been SIGMA and INTEGRAL.

The French γ-ray telescope SIGMA was on board, in 1989 the Russian mis-

sion GRANAT, which began its mission in 1989. SIGMA gave the first survey in the

transition region between hard X-rays and soft γ-rays (35 keV - 1 MeV). It mainly

observed the galactic center, detecting about 30 sources and also discovering the so

called galactic microquasars, objects that show a jet structure emanating from a com-

pact radio cone. The mask was made of tungsten and a shield of CsI scintillator was

used as an aperture-defining device (Lichti and Georgii, 2001).

The INTErnational Gamma Ray Astrophysics Laboratory (INTEGRAL),launched in 2002 by ESA, carried out transports two instruments based on SIGMA

type designs: IBIS and SPI, both operating from 15 keV to 10 MeV. There were two

additional instruments, JEM-X, an X-ray monitor which works from 3 keV to 35 keV,

and OMC, an optical telescope observing between 500 and 850 nm.

The Imager on Board of the INTEGRAL Satellite (IBIS), has an array of 53×53

opaque elements which allows an angular resolution of 12 arcsec. It consists of two

planes: the upper layer, made up of 16348 CdTe pixels (for a total area of 2621 cm2),

is used to measure between 15 keV and 400 keV with an energy resolution of 7%;

the lower layer, made of 4096 seperate CsI scintillators (total area 3318 cm2), for the

detection of γ-rays from 200 keV to 10 MeV with a typical energy resolution of 6%.

The aperture was defined by an active bismuth germanate (BGO) shield around the

detector and a passive tungsten collimator between the mask and the BGO shield.

The SPectrometer on Integral, SPI, has a coded mask like IBIS, but its main

detector consisted of an array of 19 germanium (Ge) crystals. This detector allows

35


an energy resolution of 0.2% between 20 keV and 8 MeV. The whole detector is sur-

rounded by a BGO active shield (Lichti and Georgii, 2001).

3.3.1. Compton Telescopes

The absorption probability for γ-rays in matter reaches a minimum in the range

from 1 to 5 MeV, so the imaging techniques described above do not work well at these

energies. Here the dominant interaction mechanism is Compton scattering, in which, a

photon interacts with an electron, producing a photon with a final energy Ef, depending

on the scattering angle φ, according to the formula

Ef =mec2Eγ

Eγ(1− cosφ)+mec2 (3.14)

Compton telescopes are used up to ∼ 20 MeV, although pair production be-

comes dominant over 5 MeV, due to the better performances in term of identification

of photons, angular and energy resolution.

A Compton telescope consists of two detector planes, the scatter and the ab-

sorption planes, separated by a ∼ 2 m distance. According to the Klein-Nishina cross

section formula (in the limit Eγ/mec2 ≫ 1)

σKN = r20 ·

πmec2

Eγ·[

ln(

2Eγ

mec2 +12

)](3.15)

(with the classical electron radius r0 = e2/4πε0mec2 = 2.8 · 10−15 m) the Compton

interaction probability is proportional to the electron density, i.e. to the atomic number

Z, while the photoelectric effect and pair production probability is proportional to Zn

with n ≥ 2. In the scatter detector Compton process has to be favored, so a low-

Z material must be used. The absorption detector has to absorb the energy of the

scattered photon, therefore it must be built using a high-Z material.

In an ideal case, the measurement of the energy losses in the scatter detector E1

and in the absorption detector E2 would allow the evaluation of the incoming photon

energy and scattering angle, in accordance with:

Eγ = E1 +E2 (3.16)

36


φ = arccos[

1−mec2 ·(

1E2

− 1E1 +E2

)](3.17)

To estimate photon directions one needs to know the interaction positions, ei-

ther using detector planes made of small modules or applying the Anger-camera tech-

nique (the pulse heights of PMTs viewing a scintillator allow to know the interaction

position). An obvious disadvantage of Compton telescopes is that the infalling direc-

tion of γ-rays is not uniquely identified: the only thing one can reconstruct is the event

circle, whose opening is given by the scattering angle φ in Figure 3.7.

Figure 3.7. Schematic view of a Compton telescope working principle.

Breakthrough performance of Compton telescopes were achieved with COMP-

TEL on CGRO (Lichti and Georgii, 2001).

COMPTEL was the first successful COMPton TELescope on the CGRO. It

detected photons from 700 keV to 30 MeV. Its main components were:

• a scatter detector array of liquid organic scintillator;

• an absorption detector array of NaI(Tl) crystals;

• anticoincidence shields made up of plastic scintillator;

• tagged 60Co sources for instrument calibration.

37


The Anger-camera technique allowed to measure photon direction within an event cir-

cle of ∼ 0.76◦ radius. The COMPTEL effective area ranged from 10 cm2 to 50 cm2

depending on photons energies and, the energy resolution varied from 3% to 15%. The

point source sensitivity for a two week observation was about 10−5 cm−2 s−1 (Lichti

and Georgii, 2001).

3.3.2. Pair-tracking Telescopes

For γ-rays with energy & 20 MeV, pair-tracking telescopes are used, because

the most important interaction process is pair production. A pair-tracking telescope

usually has the following parts:

• an anticoincidence shield;

• a conversion device (because the conversion probability is proportional to Z2, it

has to consist of a high Z material like tungsten (W) or Tantalum (Ta));

• a pair-tracking device (in past missions, a spark chamber);

• a time of flight measurement system or a Cherenkov detector, to help the antico-

incidence shield in rejecting the background;

• a calorimeter (CAL) for the absorption and measurement of electromagnetic en-

ergy.

The first pair-tracking telescope featuring a good signal to noise ratio was Sec-

ond Small Astronomy Satellite (SAS-2), launched in 1972. This telescope was a co-

operation with NASA, Italian Space Agency and Middle East Technical University.

It functioned properly for 6 months and it detected the first galactic γ-ray sources in

the 100 MeV range (Crab, Vela and Geminga pulsars) and diffuse galactic emission

(Fichtel et al., 1975; Ozel, 1978; Akyuz, 1993). The other important experiments in

the energy range were ESA’s European laboratories colleboration experiment COS-B

(Bignami et al., 1975) launched in 1975. The EGRET, on board of the CGRO, and

more recent Italian telescope AGILE (Tibaldo, 2007). We will give short information

on the latter experiments.

The Energetic Gamma Ray Experiment Telescope, EGRET, was one of the

four CGRO experiment, and was sensitive to γ-rays in the energy range from 20 MeV

to 30 GeV is shown in Figure 3.8.

38


Figure 3.8. The EGRET telescope (Esposito,1999)

Central unit of EGRET instrument was a multilevel wire-grid spark chamber

with tantalum conversion layers. It had a trigger telescope consisting of plastic scin-

tillator sheets in the lower part of the spark chamber. A time of flight measurement

discriminated between upward and downward moving charged particles. The CAL

was made of NaI(Tl) crystals. The field of view was about 0.5 sr, with an energy

resolution of 0.5◦ at 10 GeV. EGRET had an energy resolution of 20-25% FWHM,

an effective area of 1000 cm2 on axis. The sensitivity limit at 3σ corresponded to a

minimum flux of 10−7 cm−2 s−1 (Esposito,1999).

Italian gamma ray telescope Astro-rivelatore Gamma a Immagini LEggero, AGILE,

was launched in April 2007. Its main instrument was the Gamma-Ray Imaging De-

tector (GRID), operating between 30 MeV and 50 GeV. It consisted of a plastic

scintillator anticoincidence system, a silicon-tungsten (Si-W) tracker and a CsI

calorimeter. In contrast with previous generation instruments, it did not require gas

operations and high voltages. The new tracking technique applied to allowed a good

angular resolution (15′ for intense sources), an unprecedentedly large field of view

of about 2.5 sr and an efficiency comparable to that of EGRET (Lichti and Georgii,

2001). This complete our short survey of previous gamma-ray experiments.

3.4. Fermi Gamma-ray Space Telescope (FGST)

The Gamma-ray Large Area Space Telescope (GLAST) was launched on June

11, 2008 into a 565 km orbit with an inclination of 25.6 degrees with Earth’s equator

and then renamed the Fermi Gamma Ray Space Telescope after starting its scientific

mission on 11 August, 2008. The Fermi mission has a expected lifetime requirement

39


of 5 years and a possible extension goal of 10 years.

Figure 3.9. Artist’s concept of the Fermi Gamma-ray Space Telescope (FGST)

The Fermi payload has two science instruments, the Large Area Telescope

(LAT) and the Gamma-ray Burst Monitor (GBM). Its overall appearance is presented

in Figure 3.9. Fermi-LAT is a pair-conversion gamma-ray telescope sensitive to pho-

ton energies from 20 MeV to 300 GeV. It is a new-generation detector which provides

an unprecedented sensitivity in the γ-ray detection. I will explore the energy range

of entire MeV-GeV bands, including a part of the electromagnetic spectrum still not

covered by any other instrument. In the so called standard sky survey mode, the LAT

monitors all regions of the sky every 3 hours, leading to a highly uniform exposure

on longer timescales. Full details of the instrument, onboard and ground data process-

ing, can be found in Atwood (2009). The second instrument, GBM, aims to detected

γ-ray-bursts (GRB) with their position and measure their energy spectrum. It is made

of 12 Sodium Iodide (NaI) and 2 Bismuth Germanate (BGO) scintillation detectors.

GBM covers the lower part of the interested energy range, from few keV to about 1

MeV by NaI, and from 150 keV to 30 MeV by BGO. Overlapping with the LAT, this

detector gives a good coverage of the energy range to study GRB spectrum and solar

flares (Atwood 2009).

3.4.1. The Large Area Telescope (LAT)

The Large Area Telescope, (LAT), is the main instrument of the Fermi obser-

vatory. A schematic view of the LAT instrument with its inner structure is presented in

Figure 3.10.

The LAT is composed by a segmented anticoincidence detector (ACD) that

40


Figure 3.10. A schematic display of the LAT

surrounds the whole instrument to reject the charged-particle background. Inside the

ACD, 4 × 4 = 16 towers each measuring 43.25 cm × 43.25 cm × 84 cm are posi-

tioned. A tracker module (TKR) is located in each tower on top of the corresponding

calorimeter (CAL) module while on the bottom the Tower Electronics Modules are

placed with the Data Acquisition electronics (DAQ).

The TKR module consists of 18 (x,y)-pairs of silicon strip detector planes, with

a spacing of 228 µm between strips. The first 12 pairs are covered by a conversion

tungsten plate of 0.03 radiation lengths. The following 4 planes are covered by a 0.18

radiation lengths converter, while the last 2 tracker planes have no converter. The CAL

consists of 1536 CsI(Tl) crystals in 8 layers, for a total depth of 8.5 radiation lengths

(Rando, 2004).

Each TKR (Figure. 3.10) is composed by 18 trays one above each other. For the

backbone structure of the trays carbon has been chosen because of its large radiation

length, high stiffness modulus to density ratio, good thermal conductivity, and thermal

stability.

All trays are of similar construction, every one consists of a x− y couple of

Silicon Strip Detector (SSD) planes and, depending on its position, a tungsten (W) foil

of variable thickness. Tungsten was chosen for its high Z, in order to improve photon

conversion in electron-positron couples as the conversion probability is proportional to

Z2. As mentioned, the first 12 couples of SSD (counting from top) have a W conversion

foil of 0.03 radiation lengths (r.l.), the following 4 couples have a W foil of 0.18 r.l.

and the last 2 couples have no converters. The trays on the top have thin converters to

41


Figure 3.11. Track production in a LAT tower

optimize the Point Spread Function (PSF) at low energy. On the other hand, the trays

with a higher W thickness allow to maximize the effective area, and thus to increase

the conversion probability (even if, in this way, angular resolution is degraded due to

the Coulomb multiple scattering). Finally the last 2 trays, with no W foils, maintain a

good precision in the determination of the CAL entering point is shown in Figure 3.11.

The total TKR depth is about 1.5 radiation lengths (Buson, 2009).

Each SSD plane contains 16 units (4×4): four adjacent ladders, each one made

up by four square SSDs bonded edge to edge. Each SSD sensor has 384 strips on a

single side, with a pitch (i.e. distance between centers of adjacent strips) of 228 µm.

The TKR contributes to the first-level trigger for the LAT. Each detector layer

generates a logical signal OR of all of its 1536 channels, and the coincidence of suc-

cessive layers (typically 3 x− y planes) provides a trigger request that will be used by

subsequent subsystems. Finally, to reconstruct the track from the SSD hits, an iterative

procedure.

LAT performaces are compared with that of EGRET of CGRO

(ref:NASA/GLAST), in Table 3.1.

3.4.2. The Gamma-ray Burst Monitor (GBM)

The GLAST Burst Monitor includes 12 NaI scintillation detectors and 2 BGO

scintillation detectors. The NaI detectors cover the lower part of the energy range,

from a few keV to about 1 MeV and provide burst triggers and locations. The BGO

detectors cover the energy range of 150 keV to 30 MeV, providing a good overlap with

42


Table 3.1. LAT science requirements compared with EGRET performances

Quantity LAT EGRETEnergy range 20 MeV - 300 GeV 20 MeV - 30 GeVPeak effective area >8000 cm2 1500 cm2

Field of view >1.5 sr 0.5 srAngular resolution <3.5◦ (100 MeV) 5.8◦ (100 MeV)Energy resolution <10% 10%Deadtime per event <100 µs 100 msSource location determination <0.5′ 15′

Point source sensitivity < 6 ·10−9 cm−2 s−1 ∼ 10−7 cm−2 s−1

the NaI at the lower end, and with the LAT at the high end. GBM generates on-board

trigger for 250 GRBs per year. The reason for this limitation is that The GBM flight

software specifies different repoint criteria depending on whether or not the burst is

already within the LAT FoV, defined as within 60 of the instrument zenith (+Z) axis.

The primary science goal of the GBM is the joint analysis of spectra and time histories

of GRBs observed by both the GBM and the LAT. Secondary objectives are to provide

near-real time burst locations on-board to permit repointing of the spacecraft to obtain

LAT observations of delayed emission from bursts, and to disseminate burst locations

rapidly to the community of ground-based observers. Also, one of the goals of GBM

is to provide information to allow reorienting the Fermi observatory to position strong

bursts near the center of the LAT field of view (FoV) for extended observations. The

GBM detectors have been calibrated from 10 keV to 17.5 MeV using various gamma

sources (Meegan, 2009).

3.5. The LAT Mission

In order to understand of the high energy gamma ray sky from the previous

observatories SAS-2 (Fichtel et al., 1975) and COS-B (Bignami et al., 1975) missions

led to the EGRET instrument (Thompson et al., 1993). Fermi follows the successful

launch of AGILE by the Italian Space Agency in April 2007 (Tavani et al., 2008). LAT

offers several new opportunities for determining the nature of high energy sources and

advancing knowledge in astronomy, astrophysics, and particle physics, such as

• yield an extensive catalog of several thousand high-energy sources obtained from

an all-sky survey;

43


• measure spectra from 20 MeV to more than 50 GeV for several hundred sources;

• localize point sources to 0.3 - 2 arc minutes;

• map and obtain spectra of extended sources such as SNRs, molecular clouds,

and nearby galaxies;

• measure the diffuse isotropic γ-ray background up to TeV energies;

• permit rapid notification of high-energy γ-ray bursts and transients and facilitate

monitoring of variable sources (Band, 2009);

• explore the discovery space for dark matter (Atwood, 2009).

LAT enables detailed studies of time-resolved broad-band gamma ray spectra

of a broad range of sources, including active galaxies. As active galactic nuclei (AGN),

form an important component of point source in the high energy γ-ray sky, are discov-

ered by EGRET on the Compton Observatory (Hartman et al. 1992; Fichtel et al.

1994),

The first list of such AGNs are detected by the Fermi-LAT, the

LAT Bright AGN Sample (LBAS) (Abdo et al. 2009) includes bright, high-galactic

latitude (|b|>10o) AGNs detected by high significance (Test Statistic TS > 100) dur-

ing the first three months of scientific operation. This sample comprises 58 Flat Spec-

trum Radio Quasars (FSRQs) (among them our sources 3C 454.3 and B2 1520+31) 42

BLLac-type objects (BLLacs), two radio galaxies and four quasars of unknown type.

44

4. INTERPRETATION OF OBSERVATIONS AND RESULTS Husne DERELI

4. INTERPRETATION OF OBSERVATIONS AND RESULTS

Main goal of this thesis is to report on the results of analysis of FGST data

for the quasars 3C 454.3 and B2 1520+31. To obtain the light curves of these two

sources, the standard LAT analysis software, ScienceTools v.9.r12 is used, and a max-

imum likelihood fit of the model parameters is performed. The source model includes

a point source, a component for the Galactic diffuse emission (derived using the GAL-

PROP code; Strong et al. 2004), an isotropic component that represent the extragalactic

diffuse emission and the residual instrument background, in combination.

Maximum likelihood analysis is the official analysis method of the Fermi col-

laboration for their data. This method has already been used in previous experiments,

e.g. by COMPTEL and EGRET. Therefore, a short description of the method will be

given in the next section.

4.1. The Maximum Likelihood Method

To derive information from the measured events, e.g. the flux or spectral index

of a source, we use the maximum likelihood method. This method estimates the values

of the set of parameters maximizing the likelihood that the chosen source model fits

the collected data.

In the best case, we can bin events in energy E, direction p and arrival time t in

such a way that each bin contains only one single photon in an unbinned analysis.

Maximum likelihood method allows also the comparison of different models

for which the likelihood ratio test is found to be a powerful tool. According to Wilks’

theorem (Buson, 2009), if we consider a model M with m parameters and a second

model M0 with a subset of h < m parameters, the test statistic given by

TS = 2(lnL− lnL0) (4.1)

is distributed asymptotically as a χ2 with k = m−h degrees of freedom. This compar-

ison can be applied only within the likelihood analysis, so that the result we will have

in which of the models tested with the likelihood method is more consistent with the

data.

There are two different methods that use the maximum likelihood analysis. The

first one is ’the unbinned’ analysis which use LAT data obtained by detection of γ-rays

from a point-source. The second is ’the binned’ analysis which is designed to analyze

45


data in a large region of the sky, not being particularly interested in achieving a very

high precision in source position. In this work the unbinned analysis is used, since 3C

454.3 and B2 1520+31 are extragalactic point sources, with well known positions.

4.2. The Unbinned Analysis

Unbinned analysis consist of the following steps:

• spatial and temporal selection of events: A first selection is made on the energy

of the detected γ-ray. Only the events in the energy range between 100 MeV -

300 GeV are considered. Their systematic uncertainties are not well understood

to date. Regarding the region of extension to account for the analysis, a useful

limit can be inferred by the Point Spread Function, (PSF) of the LAT. Referring

to Figure 4.1 where the LAT containment radius is displayed versus energy at

normal incidence and at 60o incidence, we see that, for the lower energy range

selected, 68% of photons are contained within ≤ 3 degrees, in a Gaussian

approximation, and we have 99% of photons, in about 8 degrees. Therefore

in fitting, we can restrict the analysis to a region of interest (ROI) of θ ≤ 10o

degrees ensuring not an excessive loss of events.

Gamma-event selection

Figure 4.1. 68% containment radius versus energy at normal incidence (left panel) and at ±60o off-axis (right panel). The vertical line represents our energy selection limit at 200MeV.

We make a cut to remove albedo gammas (i.e. photons produced by interaction

of γ-rays with the Earth’s atmosphere) from analysis. This is performed by re-

moving from analysis the time intervals when the analyzed source angle with

respect to the detector zenith was > 105o (and thus in the range of the Earth

albedo).

46


• we create a count map of the studied region to have a first overview of the other

possible sources within the extraction region.

• we further create an exposure map which is used to calculate the number of

events per unit time expected for each bin. The exposure E is calculated as

the integral of the total response over the entire ROI extension and is defined as

(Buson, 2009)

E =∫

ROIdt dE ′ d p′ T (t)A(E ′, p′) (4.2)

in bins of time, logarithmic energy and cosine of incidence angle (The grid used

for the exposure calculation is much wider than the one used to calculate the

likelihood). In each bin, one can obtain the expected number of photons from

the source by multiplying the model flux with the corresponding exposure:

Nexp = EdΦdE

(4.3)

Generating the exposure map requires two steps: First, calculation of the life-

times as a function of energy and cosine of inclination angle (the routine for the

LAT data, expCube, provides the information about how long a single position of

the sky has been observed). With this quantity we can compute the map (called

expMap in the routine) choosing an acceptance cone with a radius larger than the

ROI for the event selection (ROI is chosen as a circle 10◦ wide centered in the

point source coordinates). An acceptance cone larger than the ROI is necessary

to ensure that photons located outside the ROI, but still coming from the source,

can be accounted of, in calculating the size of the instrument PSF.

• We developed a source model, taking into account the radiation emitted in a re-

gion larger than the ROI to avoid systematics effects. This model has to contain

all the possible sources in the extraction region that have to be fit: the galac-

tic diffuse emission, the isotropic components and point-source objects located

nearby.

47


• Then we fit the model parameters to the detected counts: For this first, we obtain

the expected distribution of observed photons by convolving the source model

with the Instrument Response Functions (IRFs). Then we calculate the likeli-

hood logarithm and maximize it to find the most probable set of parameters with

their covariance matrix.

These are the guide lines to perform the unbinned analysis in accordance to

with the provided Science Tools v.9.r12 as part of the LAT software. The analysis

presented here uses the post-launch IRFs P6 V3 DIFFUSE while the optimizer used

to find the parameters estimation is MINUIT (Buson, 2009).

4.2.1. Diffuse γ-ray Emission

To analyze galactic and extragalactic sources, a crucial point is to provide an

accurate determination of the diffuse emission which dominates the γ-ray sky. The

diffuse γ-ray emission consists of two components: the galactic diffuse emission and

the isotropic component. In the likelihood analysis an important step consists of de-

veloping a model which accounts for the radiation emitted in the selected ROI.

• Diffuse Galactic γ-ray Emission (DGE) comes mainly from the galactic plane

and is produced by interactions of high energy cosmic rays (CRs) with the inter-

stellar gas and the interstellar radiation fields. Energetic particles involved are

primarily protons and electrons: the protons interact via π0-production with the

interstellar gas, while the electrons interact via bremsstrahlung with the interstel-

lar medium (ISM) and via inverse Compton scattering (ICS) with the radiation

field.

Therefore, determining the DGE requires a model of CR propagation and must

account for the distribution of the target gas and the interstellar radiation field.

Such models are based on the theory of particle transport and interactions in the

ISM and can exploit data provided by different observations. An important con-

tribution to the development of these models has been given by EGRET (Hunter

1997, Strong, 2004).

The DGE model used to analyze LAT data has been developed in the pre-launch

period and it is now constantly updated and improved (Porter, 2008). The model

is based a numerical method and the corresponding computer code to calculate

galactic CR propagation and γ-ray production, the GALPROP run used for our

work, (Abdo, 2009).

48


• Isotropic Component (isotropic background) is difficult to disentangle from the

intense galactic diffuse foreground because it is relatively weak and has a contin-

uum spectrum with no distinguishing features. Moreover, its modeling is more

complex with respect to DGE, since all the components which may contribute

are still not well determined. Besides, determination of the isotropic compo-

nent depends on the adopted model for the DGE spectrum, which itself is not

yet firmly established. However, for the source analysis described here we do

not care to separate the various contributions and therefore we use an isotropic

spectrum inferred from LAT data themselves, assuming the galactic model pre-

viously described.

4.3. Analysis of LAT Data for 3C 454.3

3C 454.3 source has been undergoing pronounced long-term outbursts since

2000. The data (on RA:343.49, Dec:16.14) from the LAT, covering 2008 September

1 - 2009 April 26, indicate strong, highly variable γ-ray emission with an average flux

of 1.2 x 10−6 photon cm−2 s−1, for energies > 100 MeV.

The LAT also observed an increasing gamma-ray flux from a source position-

ally consistent with B2 1520+31 on RA: 230.54, Dec: +31.73, since 20 April 2009.

This object is known to be a Flat Spectrum Radio Quasar (FSRQ) with a redshift of

1.487 (Cutini 2009).

Various spectral functions are available to model the spectra of the sources

within the ROI. Regarding the diffuse components, we assume that they are uniform

at high galactic latitudes (as such is the exposure of the instrument) and thus, for the

galactic component, the spectrum is fitted with a constant function and the spatial part

is modeled with the GALPROP code 54 77Xvarh7S (Strong et al. 2004a, b) and the

corresponding isotropic component which is used as the code provided by the Diffuse

Fermi Group up to date (Buson, 2009).

The same model is used for all the point-source objects within the ROI, i.e. a

power law function described by the following equation:

dNdE

=N(γ + 1)Eγ

Eγ+1max −Eγ+1

min

(4.4)

where γ is the spectral index and dN/dE is the differential flux (for point sources it is

49


expressed in units of cm−2 s−1 MeV−1 while for extended sources it is considered per

unit solid angle (steradian, or sr) and so expressed in units of cm−2 s−1 MeV−1sr−1).

The integrated flux N is treated as a free parameter, together with the index γ , to be

evaluated over the fixed energy range Emax −Emin (that in our analysis is the selected

energy range). The errors are calculated in the minimization procedure and they do not

account for the systematics.

In the particular case of the interested source, we made a first attempt to fit it

with a power law function. The fit results show a problem with the spectral shape,

highlighted in the behavior of residuals at high energies. Therefore, accounting for the

spectral shape, we decide to try to fit it with a broken power law function (Band et al.,

1993) , i.e. a power law with a spectral index γ1 with Emin = 200 MeV and Eb is the

broken point energy, and another different spectral index γ2 from Eb = 2.461GeV to

Emax = 180 GeV. A broken power law function is therefore described by the following

equation:

dNdE

= N ·[∫ Eb

Emin

(EEb

)γ1

dE +∫ Emax

Eb

(EEb

)γ2

dE]−1

·

(

EEb

)γ1

if E ≤ Eb(EEb

)γ2

if E ≥ Eb

(4.5)

In this function, in addition to the flux and index parameters, the break-point energy

Eb is treated as a free parameter.

Using likelihood analysis results we plot the spectrum of flux as a function of

energy is shown in Figure 4.2. The data in this plot was accumulated over 6 months. It

enables us to discriminate between different spectral models. The spectra is exhibiting

a strong departure from a pure power law (PL), and a broken power-law (BPL) model

is favored as the best fit for this source. Estimated parameters of the fit with the model

of broken power law (eq. 4.5), are presented in Table 4.1.

The other estimated parameters that are fitted with simple power law are

the Galactic diffuse, isotropic diffuse emissions and the 3 point sources (MID0735,

MID0743, MID0752) in our ROI. This result shows a bit high of the Galactic Diffuse

value and a bit low of the Isotropic Diffuse value for a total ’Galactic’ + ’Isotropic’

reaching at about 2. Because we are working at high galactic latitude, where the galac-

tic and isotropic component have about the same intensity. These two source models

are scaled to unity (in units, i.e. ph/cm2/s/sr, for the totaly LAT diffuse models, our

parameter is a ratio, so it’s a dimensional), so, ideally, we would get as a result a scale

50


Table 4.1. The broken power-law model parameters over 6 months data for 3C 454.3

Energy band 200 MeV < E < 180 GeVFlux (E>100) (1.36 +/- 0.032) E−6 ph/cm2/sγ1 -2.43 +/- 0.022γ2 -3.56 +/- 0.14Eb 2.461 +/- 0.142 GeV

factor of 1 for each. But there is considerable correlation, away from the galactic plane

where the characteristic gas distribution allows us to easily identify the galactic diffuse

emission from the isotropic (galactic is clumpy; isotropic is isotropic all spread). At

high galactic latitudes these two cannot be distinguished very well, so, in the fit, one of

them goes up and the other goes down, due to the high correlation. The total intensity

is obviously constant, so the sum should always be close to 2. These values are shown

in Table 4.2.

Table 4.2. The simple power - law model parameters for Galactic diffuse, isotropic diffuseand 3 point sources. I is integral flux (I > 100 MeV).

Galactic Diffuse Value 1.476 +/- 0.054Isotropic Diffuse Normalization 0.600 +/- 0.072MID0735 I(100) = 1.482 10-7 γ= -2.63MID0743 I(100) = 0.216 10-7 γ= -1.79MID0752 I(100) = 0.456 10-7 γ= -2.87

Despite the good statistics, no curvature is apparent in the energy range below

the break in Figure 4.2 (Adbo et al., 2010). Flux above 100 MeV is (1.40 +/- 0.03) 10−6

ph cm−2 s−1 and energy break is at 2.5 +/- 0.14 GeV. Points in the plots are calculated

with Fabio’s SED macro method. Fabio’s method is used to split ROI events in energy

bins, fit each bin separately and obtain the flux of source in that bin from the fit. These

results are plotted as points in a graph, Figure 4.2.

To see if there is time variability, we divided the total time interval in 18 pieces

each lasting about 2 weeks (1200000s) and fitted each interval. Now, diffuse compo-

nents and additional point sources are kept fixed. The source 3C 454.3 is then kept

free to change. As a result of likelihood fit, values of flux and time are shown in Table

51


Figure 4.2. Spectrum of flux as a function of energy for 3C 454.3. Here six months of data

sample is analyzed with the likelihood method by using broken power law model.

4.3. After the standard analysis methods, we calculated an average flux for 3C 454.3

which indicates strong, highly variable γ - ray emission with an average flux of ∼1.2 x 10−6 photon cm−2 s−1, for > 100 MeV. After that, we plot the light curve and

we compared this with Automated Science Processing (ASP) system’s plot (flux versus

Mission Elapsed Time, MET). The results are shown in Figure 4.3.

Afterwards, to study the evolution of the spectrum shape in time, we define

3 states in the light curve of 3C 454.3 as ’high flux’, ’medium flux’ and ’low flux’

intervals. These states are shown in Figure 4.4.

We have chosen 6 points for high flux state, 7 points for medium flux state, 5

points for low flux state, and fitted data for the three intervals. For each states graphs

of spectrum were drawn as in the Figures 4.5 - 4.7. After fitting BPL model, the

parameters obtained for the three intervals are shown in the Tables 4.5 - 4.7.

52


Table 4.3. Flux and time values of 3C 454.3 as a result of likelihood fit

Tstart (s) Tstop (s) Time (average) Flux(x10−4 phcm−2s−1)240800000 242000000 241400000 0,033643242000000 243200000 242600000 0,027751243200000 244400000 243800000 0,027797244400000 245600000 245000000 0,022871245600000 246800000 246200000 0,014883246800000 248000000 247400000 0,013451248000000 249200000 248600000 0,014523249200000 250400000 249800000 0,018250250400000 251600000 251000000 0,008723251600000 252800000 252200000 0,006122252800000 254000000 253400000 0,007541254000000 255200000 254600000 0,004231255200000 256400000 255800000 0,001817256400000 257600000 257000000 0,001238257600000 258800000 258200000 0,001606258800000 260000000 259400000 0,001005260000000 261200000 260600000 0,001727261200000 262400000 261800000 0,007944

Table 4.4. Flux and time values of 3C 454.3 as a result of ASP

Time (s) Flux(phcm−2s−1) Time (s) Flux(phcm−2s−1)261748800 2,96E-007 251467200 7,79E-007261144000 2,65E-007 250862400 6,86E-007260539200 1,93E-007 250257600 1,14E-006259934400 1,37E-007 249652800 1,32E-006259329600 1,64E-007 249048000 1,50E-006258724800 1,46E-007 248443200 1,31E-006258120000 2,21E-007 247838400 9,59E-007257515200 2,27E-007 247233600 1,13E-006256910400 1,79E-007 246628800 1,36E-006256305600 2,02E-007 246024000 1,15E-006255096000 3,65E-007 245419200 1,32E-006254491200 4,75E-007 244814400 2,28E-006253886400 5,27E-007 244209600 2,04E-006253281600 8,26E-007 243604800 2,85E-006252676800 6,40E-007 243000000 3,42E-006252072000 5,18E-007 242395200 1,97E-006

53


Figure 4.3. (a) The light curve of real data analysis results by ’likelihood’ method. The

photons (>100 MeV) are used for this analysis. Each point indicates 2 week’s data of data.

(b) The same light curve by the Automated Science Processing (ASP) methods. Each point

corresponds 1 week of data

Figure 4.4. Three flux level states for 3C454.3 are defined. Upper part is the ’high flux’,

middle part is the ’medium flux’, below part is the ’low flux’ states. The portion near 250 MET

is analyzed separately.

54


Figure 4.5. Spectrum for the high flux state of 3C 454.3. Break is at 4.05 GeV and spectral

indices differ by ∆γ = γ1 − γ2 = 1.16

Table 4.5. A broken power - law model parameters for the high flux state of 3C 454.3.

Fit with a Broken Power-Law ModelEnergy band 200 MeV < E < 180 GeVFlux (E>100) (2.569 +/- 0.0668) E−6 ph/cm2/sγ1 -2.401 +/- 0.0025γ2 -3.417 +/- 0.146Eb 4.051 +/- 0.184 GeV

Table 4.6. A broken power - law model parameters for condition of the medium flux data of3C 454.3.

Fit with a Broken Power-Law ModelEnergy band 200 MeV < E < 180 GeVFlux (E>100) (1.05 +/- 0.00005) E−6 ph/cm2/sγ1 -2.451 +/- 0.0044γ2 -3.899 +/- 0.385Eb 2.531 +/- 0.316 GeV

55


Figure 4.6. Spectrum for the medium flux state of 3C 454.3. Breaking point is at Eg≃2.53

GeV and indices before and after the break differs by about 0.55

Figure 4.7. Spectrum for the low flux state of 3C 454.3.

56


Table 4.7. A simple power - law model parameters for condition of low flux data of 3C454.3.

Fit with a Simple Power-Law ModelEnergy band 200 MeV < E < 180 GeVFlux (E>100) (0.279 +/- 0.034) E−6 ph/cm2/sγ -2.61 +/- 0.10

We applied a simple power law for low flux data sample because the energy

break from the fit results greater than maximum energy. A simple power law model

parameters are shown Table 4.7.

We summarize the spectrum evolution for each state of 3C 454.3 in Figure 4.8.

That is superimposed to see the variation. Red curve is for high, blue curve is for

medium and the grey curve is for the low flux states. Break energy seems to move

towards the lower energies as flux level is degreasing and the break as flux gets from

high for medium flux state does not exist in the low flux state. Probably we do not have

enough data to observe of the break.

Figure 4.8. This plot shows the spectra for is high, medium and low flux states together for

3C 454.3. Break energy moves to a higher energy as we move from medium to high flux states.

For the low state, no break is observed.

As a result we used broken power - law for high and medium flux states and

57


simple power law for low flux state. SED of each states for (≥200 MeV) are shown in

the Figure 4.9.

The Figure 4.8 gives no information about errors, but the Figure 4.9 is more

complete as all the fit results are shown, with 1-sigma uncertainty level in a band, so we

have an idea of the statistical uncertainty from the fit (parameters errors and parameter

correlations are considered). However, break point energies differ a bit from Figure

4.8, especially at high energies, moving from Eb=4.0 GeV to Eb=3.0 GeV. It changes

unnoticeably in the medium flux level. Also, the error is changed at high flux states as

seen in the Figure 4.9, the reason for this change is not understood.

Figure 4.9. The plot is Spectral Energy Distribution for high, medium and low flux states as

defined in the text.

58


4.4. Analysis of LAT Data for B2 1520+31

A similar analysis using the same software was carried out for a second source

B2 1520+31. Two methods were used for the analysis of this source data. First one

is the same likelihood method. Spectral analysis using in five different energy bands

and their light curves are obtained with this method. Second method was, the ’aperture

photometry’, also used to obtain the light curves for this source. The likelihood analy-

sis is a more rigorous approach and offers the potential to reach a greater sensitivity. It

also leads to a more accurate flux measurement, as background can be modeled. Be-

sides, more detailed source models can also be applied. In contrast, the aperture pho-

tometry is also useful, while it is computationally less demanding. It requires fewer

analysis steps and provides a model independent measure of the flux.

520 days (17.3 months) of data (from 2008 August 4 to 2010 January 25) se-

lected in energy range from 100 MeV to 300 GeV were used for spectral analysis. The

spectrum was fitted with a constant function plus a spatial part. Data were modeled

with the GALPROP code 54 77Xvarh7S for the galactic diffuse and isotropic com-

ponents. Proper fitting of the source analysis requires simultaneous modeling of the

background contributions. These contributions are incorporated (in the XML exten-

sion) into model file in addition to the source models.

The likelihood analysis was performed with the standard analysis tool gtlike,

which is part of the Fermi-LAT ScienceTools software package (version v.9.r12).

The first set of instrument response functions (IRFs) tuned with the flight data,

P6 V3 DIFFUSE, was used in the analysis. In contrast to the preflight IRFs, these IRFs

take into account for corrections for pile-up effects. This correction being higher for

lower energy photons, the measured photon index of a given source is about 0.1 higher

(i.e. the spectrum is softer) with this IRF set as compared to the P6 V1 DIFFUSE one

used previously in Abdo et al. (2009e). Photons were selected in circular regions of

interest (ROI), 20o in radius, centered at the positions of the sources of interest. The

isotropic background (the sum of residual instrumental background and extragalactic

diffuse gamma-ray background) was modeled with a simple power-law. The GAL-

PROP model (Strong et al. 2004a,b), version gll iem v01.fit, was used for the galactic

diffuse emission, with both flux and spectral photon index left free in the fit. All point

sources with statistical source identification parameter TS>25 in the 6-month source

list, lying within the ROI and a surrounding 10o - wide annulus, were modeled in the

fit with single power-law distributions were used as ROI + 5o (Abdo, et al., 2010).

Using likelihood analysis results spectrum of flux of B2 1520+31 as a function

59


of energy is given in Figure 4.10. In this work, the accumulated data over 17 months

10 days were fitted with simple PL using five different energy ranges: 100-300 MeV,

300-1000 MeV, 1-3 GeV, 3-10 GeV, 10-100 GeV. The results compared with Fermi-

LAT catalog (version gll psc v02.fit) results are shown in Figure 4.10. The data in

all energy range were also fitted with a simple power law (PL) model and a broken

power-law (BPL) model. These results were later compared with each other as shown

in Figure 4.10. Log(likelihood) values are used to calculate the probability-value with

the help of test statistics formula (Abdo, et al., 2009). Then we decided to obtain the

best fit of simple PL model for this source. Model parameters are shown in Table 4.8

and Table 4.9.

Table 4.8. A broken power - law model parameters for B2 1520+31

Fit with a Broken Power-Law ModelEnergy band 100 MeV < E < 300 GeVFlux (E>100) (4.307 +/- 0.082) E−7 ph/cm2/sγ1 -2.278 +/- 0.024γ2 -2.604 +/- 0.044Eb 1.21 +/- 0.107 GeVTS 14444.3-log(likelihood) 2334872.716

Table 4.9. A simple power - law model parameters for B2 1520+31

Fit with a Simple Power-Law ModelEnergy band 100 MeV < E < 300 GeVFlux (E>100) (0.187 +/- 0.0069) E−7 ph/cm2/sγ -2.384 +/- 0.013TS 14666.6-log(likelihood) 23333614.152

To get the light curve (flux versus time) of B2 1520+31, we divided the total

time interval into 78 pieces each lasting about 1 week (60000s) and we applied the fit

for each interval. Now, diffuse component and additional point sources are kept fixed,

60


Figure 4.10. Energy spectra from the real (red circles) and simulated (green triangles) data

for B2 1520+31 resulting from summing the power - law distributions with parameters flux and

photon index, as measured in weekly bins. The red dashed line represents the Power Law fits

and blue dashed line represents Broken Power Law fits.

61


as B2 1520+31 flux is kept free. After the standard likelihood analysis methods we

obtained an average flux and a spectral index value per week for B2 1520+31. These

results are compared with catalog results. The light curve of B2 1520+31 is shown in

Figure 4.11 and a plot of spectral index versus flux is shown in Figure 4.12.

Figure 4.11. Time variation of the flux obtained with the likelihood analysis for B2 1520+31

in the 100 MeV - 300 GeV band. Red points are real and green points are simulated data.

The LAT was operated in the survey mode throughout these observations except during the

period 2008 August 4 - 2010 January 25, when it was operated in the pointed mode. Each

point represents one week’s data. The error bars are statistical only.

As seen in Figure 4.12, there is no clear correlation between flux and hardness

values is observed, so we cannot say e.g. that a brighter sources is harder, as one could

expect from the AGN unified model.

From the aperture photometry analysis methods, we obtained average flux

value per week for B2 1520+31. We plot its light curve as given in Figure 4.13. Since

we were only interested in the source region, we chose a ROI as 2o. Because the

light curves obtained from aperture photometry procedure, background value are not

subtracted.

62


Figure 4.12. The variation of spectral index for B2 1520+31 against flux. These values are

obtained by the likelihood analysis for B2 1520+31 in the 100 MeV - 300 GeV band.

Figure 4.13. The variation of the flux with aperture photometry analysis for B2 1520+31 in

the 100 MeV - 300 GeV band. The LAT was operated in the survey mode throughout these

observations except during the period 2008 August 4 - 2010 January 25, when it was operated

in the pointing mode. Each point is for one week’s duration of data. The error bars are statistical

only.

63

5. CONCLUSIONS Husne DERELI

5. CONCLUSIONS

This thesis reports one of first attempts to provide a description of variability of

3C 454.3 by the LAT instrument onboard the FGST. The source was in a flaring/active

state since 2000 and, it was easily detected and showed rapid variability which is de-

scribed as symmetric flares with rise and fall time of ∼ 3.5 days. This source was

expected to play an important role in the LAT observations, because it was reported in

the first three months LAT Bright AGN Source List. We also presented a description

of variability of B2 1520+31 that had a flux around 4 times greater than the average

flux reported in the LAT Bright AGN Source List.

The spectral analysis of both of these sources were performed with the stan-

dard procedures and methods provided by the Fermi Science Tools (v.9.r12). FGST

performs a maximum likelihood fit of the model parameters for all source, including

for 3C 454.3 and B2 1520+31. The diffuse components (Galactic and extragalactic)

are also fit with the data, hence taking into account also the residual instrumental back-

ground. The Galactic diffuse emission is derived using the GALPROP code, while the

extragalactic diffuse emission was modeled with an isotropic power law component.

All blazars spectra measured by EGRET were represented with pure PLs. How-

ever, FGST has revealed that the spectra of some low-energy peaked blazars display

a strong departure from a pure-PL behavior, with a BPL function as the best model,

according to its improved sensitivity. We found, for example, that the γ - ray spectrum

of 3C 454.3 is not a simple power-law with the index ∼ 2.3; instead, steepens toward

higher energies. A good, (but not unique), description of its spectrum is a broken

power-law with photon indices of ∼ 2.4 and ∼ 3.5, below and above a break at ∼ 2.4

GeV, respectively.

The observed break might be due to photon-photon absorption to pair produc-

tion; however, this would require a space region responsible for production of γ - ray

flux to be sufficiently close to the accretion disk/black hole system to produce the spec-

tral signatures of the reprocessed γ - rays in the X - ray photon energy range. These

are not observed yet.

We found that the γ - ray spectrum of second source, B2 1520+31 is a simple

power-law, that is, it does not steepen toward higher energies. If we force a BPL

spectrum, its spectrum to this source photon indices of ∼ 2.2 and ∼ 2.6, below and

above at a break ∼ 1.2 GeV, respectively is obtained. We conclude that simple PL

gives a better fit than BPL for B2 1520+31.

We observed a significant break for 3C 454.3 around 2.4 GeV and not so signif-

64


icant break for B2 1520+31 around 1.2 GeV. Break were ascribed to mirroring a similar

feature in the underlying emitting electron energy distribution; the Klein-Nishina effect

was not ruled out, though the importance of photon-photon pair production requires the

gamma-ray emission region to be close to the supermassive black hole. Clearly, under-

standing the details of the spectral break is important for understanding the structure

and location of the dissipation region of jets in active galaxies.

In a large statistical sample, for 19 of the 22 brightest LBAS FSRQs, a like-

lihood ratio test (LRT, Mattox et al. 1996) rejects the hypothesis that spectrum is a

PL (null hypothesis) against the one that the spectrum is a BPL, at a confidence level

greater than 97% (Abdo et al., 2010). The fact that spectra for most FSRQs are best

modeled by a broken power law with a break in the 1-10 GeV range is quite unex-

pected. This break becomes a distinctive feature of these sources.

We also calculated that average flux for 3C 454.3 indicates a strong, highly

variable γ - ray emission with an average flux of 1.2 x 10−6 photon cm-2 s-1, for

energies > 100 MeV. Concerning its light curve, sometimes 3C 454.3 shows some

anomalous bumps. These bumps are not yet fully understood. Moreover, nobody even

knows how it behaves during those peaks. They thought that the integrated flux rose

along with index, but from our earlier analysis, it seems more likely that, there are

at least two different behaviors, which could be interpreted as a sort of warming and

cooling phases. The spectra of such bumps can tell as about their nature.

We now know that Blazars constitute the most numerous classes of extragalac-

tic γ-ray sources. Fermi-LAT observations seen to be very fundamental to explain the

real nature of blazars. In addition, the components of the VHE class are of crucial

importance in the comprehension of the blazar physics and in the understanding of

extragalactic γ-ray interactions. In short, considering their effects on cosmology; the

mystery of a supermassive black hole in the center of AGNs are still unexplained. Our

observations by LAT and other similar experiments will help to improve our under-

standing of how supermassive black holes power AGNs. The comprehension of the

emission process related to the physics of blazars is particularly important to under-

stand the nature of these objects that represent one of the most energetic emission of

universe.

65

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69

RESUME

She was born in Malatya. She graduated from primary and secondary schools in

Malatya. After that, She enrolled in the high school of Hekimhan Lisesi and graduated

in 2001. She enrolled in the Physics Department of Cukurova University and She

graduated in 2006. She continued to study for my Masters degree in High Energy

Astrophysics, at the Institute of Natural and Applied Sciences in Cukurova University.

She has been worked at Galileo Galilei Physics department in University of Padova as

an Erasmus student.

70

1. APPENDIX-A Husne DERELI

1. APPENDIX-A

The purpose of this appendix is to give a formulation of Fermi acceleration

mechanism. A complete review of this subject can be found in Longair, 1983.

1.1. Fermi Acceleration Mechanism

The Fermi acceleration can be view from different approach. In Fermi’s origi-

nal picture, he imagined charged particles being reflected from ’magnet mirrors’ asso-

ciated with the Galactic magnetic field. These mirrors are in random motion (to first

order). Fermi asked what the energies of such particles would be if they remained in

between such clouds for a certain time T.

In a simplified version of the problem, we consider only the one-dimensional

case. There are as many mirrors or clouds moving towards the particle as away from

it for an external observer. Consequently, the particle makes ’head on’ and ’following’

collisions with these clouds. First, we work out the change of energy of the particle in

a single collision. Let us do this for the relativistic speeds.

The situation is shown in Figure (a). We suppose the cloud is infinitely massive

(compared with particle) so that its velocity is not changed in the collisions. The center

of momentum frame is therefore that of the cloud moving at velocity V. The energy of

the particle in this frame is

E′= γV (E +V p) (1.1)

where

γV = (1− V 2

c2 )−1/2 (1.2)

The relativistic three-momentum of the particle will be:

p′= γV (p+

V Ec2 ) (1.3)

In the collision, the particle’s energy is conserved, E′be f ore = E

′a f ter = E

′, and its

71


Figure 1.1. (a) To illustrate the collisions between a particle of mass m and a cloud of mass

M. (b) To illustrate the collisions between a particle and equal numbers of clouds moving in

opposite directions in one dimension (Longair 1983)

momentum vector reversed, p′ →−p

′. Therefore, transforming back to the observer’s

particles frame we find

E′′= γv(E

′+V p

′) (1.4)

After a bit of manipulation of this results, we can express it in the form

E′′= E +2γ2

v EVc(Vc+

υc) (1.5)

i.e.

∆E = 2γ2v E

Vc(Vc+

υc) (1.6)

Now, if instead it was a following collisions, energy would have been lost

∆E =−2γ2v E

Vc(υc− V

c) (1.7)

However, we notice that there is greater probability of head-on than following

collisions. From Figure (b), we see that the frequency of encounters is just proportional

72


to the relative velocity of the particle and cloud, i.e. V + υ for head-on collisions

and υ −V for following collisions. Therefore the probability of a head-on collision12((V + υ)/υ) and of a following collision 1

2((υ −V )/υ). Therefore, the net mean

energy gain per collision is

∆E =−12(υ +V

υ)2γ2

v EVc(Vc+

υc)+

12(υ −V

υ)2γ2

v EVc(υc− V

c)E (1.8)

This simplifies down to

∆EE

= 4γ2v (

Vc)2 (1.9)

if V ≪ c, this is ∆E/E = 4(V/c)2. Therefore, the rate of gain of energy is

dEdt

= 4M(Vc)2E = αE (1.10)

where M is the number of collisions per second. Notice that this represents

exponential of the energy of the particle.

Now, we assume that the particle stays in the accelerating region for a charac-

teristic time τ . Then we write down the diffusion equation for particle acceleration and

find the solution for N(E) in equilibrium, i.e.

dNdt

= D∇2N +∂

∂E[b(E)N(E)]− N

τ+Q(E) (1.11)

We are interested in the steady-state solution and, hence, dN/dt=0. We are not

interested in diffusion and hence D∇2N = 0 and we assume there are no sources, Q(E)

= 0. The energy loss term is b(E) = - dE/dt which in our case is −αE. Therefore

equation (1.11) reduces to

− ∂∂E

[αEN(E)]− N(E)τ

= 0 (1.12)

Differentiating with respect to E and rearranging, we get

dN(E)dE

=−(1+1

ατN(E)

E(1.13)

73


Therefore

N(E) = constant ·E−(1+α−1τ−1) (1.14)

In conclusion, we have managed to derive a power law energy spectrum of

”clouds”. This is general idea of Fermi for clouds, but now we know that power spec-

trum of photon is generated by bremsstrahlung or supernova remnants and supernova

bumps, so photons show power law spectrum.

Now, we can derive simple power law of ”proton and electron” or ”photons”

for Fermi-LAT from this formula (1.14)

dNdE

= A ·Eγ (1.15)

where A is a constant, γ is a index∫ Emax

Emin

dN(E)dE

dE = N (1.16)

Then

∫ Emax

Emin

A ·EγdE = N (1.17)

A =N∫ Emax

EminEγdE

=N

1γ+1Eγ+1|Emax

Emin

=N · (γ +1)

Eγ+1max −Eγ+1

min

(1.18)

dNdE

1Eγ =

N · (γ +1)

Eγ+1max −Eγ+1

min

(1.19)

After that we obtained equation (4.4) that is simple power law of Fermi - LAT;

dNdE

=N · (γ +1) ·Eγ

Eγ+1max −Eγ+1

min

(1.20)

A =N∫ Eb

Emin( E

Eb)γ1dE +

∫ EmaxEb

( EEb)γ2dE

(1.21)

Finally, we obtained equation (4.5) that is broken power law of Fermi - LAT;

74


dNdE

= N ·[∫ Eb

Emin

(EEb

)γ1

dE +∫ Emax

Eb

(EEb

)γ2

dE]−1

·

(

EEb

)γ1

if E ≤ Eb(EEb

)γ2

if E ≥ Eb

(1.22)

75

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