Study of the Galactic Diﬀuse X-Ray Emission
with the Suzaku Satellite
Hideki Uchiyama
Department of Physics, Graduate School of Science, Kyoto University
Kitashirakawa Oiwake-cho, Sakyo-ku, Kyoto, 606-8502, Japan

2

Abstract
The Milky Way galaxy emits diﬀuse X-rays (the Galactic Diﬀuse X-ray Emission, GDXE) fromthe center, ridge and bulge. The GDXE accompanies the strong Fe XXV Kα (6.7 keV), Fe XXVIKα (6.97 keV) and Fe I Kα (6.4 keV) lines. The origin of the GDXE and the three Fe Kα lines
has been an open question though supernova remnants (SNRs) or cataclysmic variables (CVs)are suggested as the probable candidates of the origin. The spatial distributions and propertiesof the Fe Kα lines in the GDXE are the key to revealing the origin.
In this thesis, we show the results of the Suzaku observation of the region from the Galactic
center (GC) to the ridge and the bulge (−3◦ < l < 2◦ , −2◦ < b < 1◦ ). We also report theanalysis results of a SNR located in the GC region, Sgr A East, and a newly found X-ray point
source, Suzaku J1740.5−3014, in the data set.
The intensity proﬁles of the highly ionized Fe Kα lines and the continuum in the GDXE are
described with a phenomenological two-exponential component model. The e-folding scale of theintensity proﬁle along the Galactic longitude changes around |l| = 1◦ –2◦ . The transition of theline intensity ratio of Fe XXVI Kα to Fe XXV Kα also occurs in |l| = 1◦ –2◦ , which corresponds
to the border between the nuclear bulge (NB) and the Galactic disk (GD). These facts suggestthat the GDXE contains two components which accompany the NB and the GD, respectively.
The scale height of the intensity proﬁle along the Galactic latitude changes around |b| = 0◦ .5–
1◦ . The line intensity ratio of Fe XXVI Kα to Fe XXV Kα is not diﬀerent in −2◦ < b < 1◦ ,signiﬁcantly.
The equivalent widths of the Fe I Kα and Fe XXV Kα lines are anti-correlated in the GD
regions (|l| = 1◦ –3◦ , |b| < 0◦ .1). From this relation, we obtained the equivalent widths of the Fe′ and EW ′ , respectively. The
I Kα and Fe XXV Kα lines to their associated continua, EW6.4
6.7
′
EW6.4
is signiﬁcantly smaller than those reported previously in the inner NB regions (Sgr A′ is almost same among the every region. It suggests that the origin of
and C) though the EW6.7
the Fe I Kα line might be diﬀerent between the NB and the GD.
To reveal whether the GDXE is the superposition of point sources or truly diﬀuse plasma,
we compared the intensity proﬁle of the GDXE with those of near-infrared observations and
X-ray point sources. However, conclusive results are not obtained.
In the analysis of Sgr A East, we detected the non-thermal component for the ﬁrst time anddetermined the temperature accurately. Sgr A East has two-temperature plasma of kT ∼ 6 and∼ 1.2 keV. The temperature of kT ∼ 6 keV is peculiarly hot as a SNR and close to that of the
GDXE. However, the total line intensity ratio of Fe XXVI Kα to Fe XXV Kα is 0.05, and too
low to explain that (0.2–0.4) of the GDXE.
From Suzaku J1740.5−3014, we detected three Fe Kα lines and the coherent and clear pulsation of 432.1 s. These features are typical for an intermediate polar (IP), one of the categories ofcataclysmic variables. The interstellar absorption indicates that Suzaku J1740.5−3014 is located
not in the GC region but in the distance of ∼ 2 kpc from the Sun. The luminosity of ∼ 1033 erg
s−1 is also typical for an IP.
We compared the GDXE spectra with those of Sgr A East, Suzaku J1740.5−3014, otherCVs, other SNRs and point sources in the GC region. The spectra of these sources are diﬀerent
from that of the GDXE.

Contents
1 Introduction

12 Review
2.1 Structure of the Milky Way Galaxy and the Galactic Center Region . . . . . . .
2.2 Discovery of the Galactic Diﬀuse X-ray Emission and its Fe Kα Lines . . . . . .
2.3 Candidates for the Origin of the Fe XXV Kα and Fe XXVI Kα Line Emissions .
2.3.1 Truly Diﬀuse Plasma Origin . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.2 Faint Point Sources Origin . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.3 Charge Exchange Origin . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.4 Candidates for the Origin of the Fe I Kα Line Emission . . . . . . . . . . . . . .
2.4.1 Truly Diﬀuse Interstellar Gas Origin . . . . . . . . . . . . . . . . . . . . .
2.4.2 Faint Point Sources Origin . . . . . . . . . . . . . . . . . . . . . . . . . .
2.5 Recent Observation of the GDXE . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.5.1 GDXE in the Sgr A Region (|l| < 0.3◦ ) . . . . . . . . . . . . . . . . . . . .
2.5.2 GDXE in the Inner Region of the NB (0◦ .3 < |l| < 2◦ ) . . . . . . . . . . .
2.5.3 GDXE in the Galactic Ridge (|l| > 8◦ ) . . . . . . . . . . . . . . . . . . . .
2.5.4 GDXE in the Galactic Bulge (|b| > 1◦ ) . . . . . . . . . . . . . . . . . . . .
2.6 Review of Sgr A East . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.7 Summary of the Recent Study and Subjects Tackled in This thesis . . . . . . . .

3
3
5
6
6
7
10
10
10
10
11
11
12
13
14
16
16

3 Instrument
3.1 Overview of the Suzaku Satellite . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2 X-Ray Telescope (XRT) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3 X-Ray Imaging Spectrometer (XIS) . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.1 Introduction of X-ray Charge Coupled Device Cameras . . . . . . . . . .
3.3.2 Overview of the XIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.3 Charge Injection Capability . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.4 Performance of the XIS . . . . . . . . . . . . . . . . . . . . . . . . . . . .

19
19
20
23
23
24
26
29

4 Galactic Diﬀuse X-ray Emission
4.1 Observation & Data Reduction . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.1 Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.2 Data Reduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.3 Response Function of the XIS . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.4 Background . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.5 Exclusion of Bright Sources . . . . . . . . . . . . . . . . . . . . . . . . . .

35
35
35
35
37
37
37

i

ii

CONTENTS

4.2

4.3

Spatial Distribution of the GDXE . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.1 Intensity Proﬁles along the Galactic Longitude . . . . . . . . . . . . . . .
4.2.2 Intensity Proﬁles along the Galactic Latitude . . . . . . . . . . . . . . . .
4.2.3 Two-Dimensional Model for the Intensity Proﬁle . . . . . . . . . . . . . .
Detailed Spectral Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

40
40
43
53
56

5 Sgr A East69
5.1 Observation & Data Reduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . 69
5.2 Spectral Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70
5.2.1 Emission Lines . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70
5.2.2 Sulfur Line Diagnostics . . . . . . . . . . . . . . . . . . . . . . . . . . . . 71
5.2.3 Iron and Nickel Line Diagnostics . . . . . . . . . . . . . . . . . . . . . . . 72
5.2.4 Thin Thermal Plasmas and Power-Law Component . . . . . . . . . . . . 76
5.3 Short Discussion for Sgr A East . . . . . . . . . . . . . . . . . . . . . . . . . . . . 76
5.3.1 Thermal Plasma . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 76
5.3.2 Hard Tail . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 76
5.3.3 Chemical Compositions . . . . . . . . . . . . . . . . . . . . . . . . . . . . 79
6 Suzaku J1740.5–3014
81
6.1 Observation & Data Reduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . 81
6.2 Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 82
6.2.1 Imaging . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 82
6.2.2 Spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83
6.2.3 Timing . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 86
6.3 Short Discussion for Suzaku J1740.5-3014 . . . . . . . . . . . . . . . . . . . . . . 87
6.3.1 Identify of Suzaku J1740.5-3014 by Comparison with Previous X-ray Observations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 87
6.3.2 Suzaku J1740.5-3014 as a New Intermediate Polar Candidate . . . . . . . 88
7 Discussion
91
7.1 Spatial Distributions of the Fe Kα Line Emissions of the GDXE . . . . . . . . . 91
7.1.1 Uncertainty of the Intensity Measurement . . . . . . . . . . . . . . . . . . 91
7.1.2 Comparisons with Stellar Mass Distribution Observed in the Infrared . . 91
7.1.3 Comparisons with Distribution of X-ray Point Sources Observed with
Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95
7.2 Spectra of the GDXE . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95
7.2.1 Line Intensity Ratios of Fe XXVI Kα to Fe XXV Kα . . . . . . . . . . . . 95
7.2.2 Relation between the Equivalent Widths of the Fe I Kα and Fe XXV KαLine Emissions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 100
7.2.3 Center Energy of the Fe XXV Kα Line Emission . . . . . . . . . . . . . . 102
7.3 Comparison between the Intensities of the GDXE and X-ray Point Sources Detected with Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 103
7.4 Properties of Sgr A East and Other GC SNR and Their Contribution to the GDXE103
7.5 Properties of Suzaku J1740.5-3014 and Other Point Sources and Their Contribution to the GDXE . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 105

CONTENTS

iii

8 Conclusion

107

A X-ray Emission from Thin-Thermal Plasma109

B CTI Correction and In-Orbit Calibration for SCI
113
B.1 Property of the CTI in the SCI Mode and the “Sawtooth” CTI model . . . . . . 113
B.2 CTI Measurement in Orbit . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 115
B.2.1 Calibration Data . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 115
B.2.2 Determination of the CTI Parameters . . . . . . . . . . . . . . . . . . . . 118
B.3 Energy Scale Uniformity and Resolution in the SCI Mode . . . . . . . . . . . . . 121

Chapter 1

Introduction
Since the beginning of the astronomy, people had observed the sky in the visible light mainly.
Visible light is emitted from objects at a temperature of a few thousands K like the Sun. The
temperature of the X-ray emitting object is about 107 K. Before 1960s, most of scientists did not
imagine that such a hot object commonly exists in the universe. Ricardo Giacconi and Bruno
Benedetto Rossi, however, launched an Aerobee rocket carrying Geiger counters and discovered
the ﬁrst extrasolar X-ray source, Sco X-1 [23]. It was the beginning of the X-ray astronomy.
Bruno Benedetto Rossi said;”I must admit that my main motivation for pressing forward was a deep seated faith in theboundless resourcefulness of nature, which so often leaves the most daring imagination of man
far behind.”
As he said, ”unimaginable” X-ray emissions have been discovered from a lot of celestial
objects, for example, X-ray binaries like Sco X-1, of course, young stellar objects [116], late typestars [67], clusters of the galaxies [13], and comets [57] etc. Our Galaxy (the Milky Way galaxy)is one of the most unimaginable X-ray sources.
The Galactic Diﬀuse X-ray Emission (GDXE) extends over more than 100◦ along the Galactic
plane but only ∼ 2◦ across it (e.g. [137, 142]) and thus it has been called as the Galactic “ridge”
X-ray emission. One of the most interesting features of the GDXE is its strong Fe Kα line
emissions discovered with the Tenma satellite [44]. With the Ginga satellite, it was revealed
that the Fe Kα emission concentrates sharply at the Galactic center (GC) region [47]. With the
ASCA satellite, the Fe Kα line feature is resolved to three lines at the energy of 6.4, 6.7 and6.97 keV [48].
The 6.7 and 6.97 keV lines correspond to the Fe XXV Kα and XXVI Kα line emissions. These
Fe ions imply the existence of the optically thin high-temperature plasma (kT = 5–10 keV or
T = several 107 K). If there is truly diﬀuse hot plasma gas which ﬁlls the Galactic plane, thehuge energy input is required to sustain it. The most probable energy suppliers are supernovae
but the required supernova rate is about one per 10 years [45], which is one order lager than the
recently estimated rate [17, 65]. In addition, the temperature of the GDXE is higher than thoseof known young supernova remnants (SNRs). Unknown energy suppliers and plasma heatingprocess might exist in our Galaxy. An alternatively explanation of the GDXE, proposed soonafter its discovery (e.g. [44]), is the superposition of unresolved faint point sources. The mostprobable candidates of such faint point sources are now thought to be cataclysmic variables(CVs), which are binaries of a white dwarf and a main-sequence star. CVs have three Fe Kαlines and thus resemble the spectrum of the GDXE. However, the detail spectral analysis shows
1

2

CHAPTER 1. INTRODUCTION

the diﬃculty of this scenario. The equivalent widths of Fe Kα lines of CVs in the vicinity of the
Sun are typically 100–200 eV [20]. The equivalent width of Fe XXV Kα of the GDXE, however,is more than 400 eV [55, 143]. It means that the integration of the spectra of known CVs cannot
explain the spectrum of the GDXE. The GDXE might be composed of unknown type of faintpoint sources with larger equivalent widths of Fe Kα lines than those of well-known CVs.
The origin of the GDXE has been still an open question, which might link to unknown
activities of our Galaxy or new X-ray sources. The recent studies, especially spectral analysis
using the Suzaku data, have shown that the properties of the GDXE in the GC region (|l| < 1◦ )1[51, 55, 83] and the ridge (|l| > 8◦ ) [143] are clearly diﬀerent. The intermediate region between
them, however, has not been studied well. This region is important to know the transition of theproperty of the GDXE from the GC to the ridge. To know the entire structure of the GDXE,
this region is one of the last lacked pieces. In addition, the edge of the nuclear bulge (NB) is
located around |l| = 1◦ –3◦ [68]. The NB is a stellar complex in which recent star formations or
young ionizing stars are present diﬀerently from the old Galactic bulge. If the GDXE is related
to the NB, the property of the GDXE might change around |l| = 1◦ –3◦ . In this region, however,neither of the spatial proﬁle and spectra of the GDXE has been known well. The X-ray point
sources around this region also have not been studied well.
In this thesis, we analyzed the data of the Suzaku observation of the regions from the GC
to the ridge and the bulge (−3◦ < l < 2◦ , −2◦ < b < 1◦ ). The main aim of this thesis is toreveal the properties of the GDXE from the GC to the ridge and the bulge and to make clear
the diﬀerence of the GDXE of each region, especially the Fe Kα line emissions. Based on theseresults, we will tackle the origin of the GDXE. This thesis is organized as below. In chapter 2,
we reviewed the past and recent study of the GDXE, especially its Fe Kα lines. In chapter 4, theSuzaku data of GDXE of the GC to ridge region (−3◦ < l < 2◦ , |b| < 1◦ ) and the bulge region(b = −2◦ ) were analyzed. The intensity proﬁle of the Fe Kα line emissions and the results ofdetail analysis of the GDXE spectra are shown in this chapter. In the following two sections, weanalyzed peculiar sources as samples of the possible origin of the GDXE. In chapter 5, a peculiar
hot SNR in the GC, Sgr A East was studied. In chapter 6, a newly found CV candidate, Suzaku
J1740.5–3014 was studied as a sample of X-ray point sources in the region between the GC andridge. In chapter 7, we made discussions. Chapter 8 is the conclusion of this thesis.

1

The notations l and b denote the coordinates of the Galactic longitude and latitude.

Chapter 2

Review
The instruction for the structure of the Milky Way galaxy is in section 2.1. The candidate
origins of the Fe Kα lines in the GDXE are reviewed in section 2.3 and 2.4. Recent studies of
the GDXE with Suzaku, Chandra and RXTE are reviewed in section 2.5. Review of Sgr A East
is in section 2.6. The summary of the recent studies is tabled in section 2.7.

2.1Structure of the Milky Way Galaxy and the Galactic Center
RegionHere we review the structure of the Milky Way galaxy and the Galactic center (GC) region
brieﬂy according to mainly Mezger et al. [67]. The Milky Way galaxy is composed of the
Galactic disk (GD), the Galactic bulge (GB), the nuclear bulge (NB) and the Galactic halo as
shown in ﬁgure 2.1.
The GD (3 kpc < R < 14 kpc where R is the radius from the dynamical center.) is the main
component of the Milky Way galaxy. The GB (0.3 kpc < R < 3 kpc) is the ellipsoidal region in
the inner region. The GB consists of mainly old evolved stars.
The GC is thought to be at a distance of 7.5–8.5 kpc from the Sun. In this thesis, we
adopt 8 kpc as the distance to the GC (see [66] and references therein). The typical interstellar
absorption toward the GC region is NH = 6 × 1022 cm−2 [96] and thus the visible light from the
GC region is absorbed. We can observe the GC region only in the limited band; below λ ≥ 2µm (radio and near-infrared) and above E ≥ 2 keV (X- and γ-ray).
The large and dense stellar complex of the innermost region (R ≤ 300 pc or |l| ≤ 2◦ ) is called
the nuclear bulge (NB). The NB shows ample evidence of recent and ongoing star formation(e.g. populous clusters of massive stars, very young IR clusters, molecular clouds that are able
to form new massive clusters, and massive main-sequence stars etc. See van Loon et al. [136]and references therein.) and thus this region is distinct from the old GB. An unusually dense
molecular cloud complex exists in the central few hundred pc (|l| ≤ 1◦ ) of the Milky Way galaxy
and it is often referred to as the Central Molecular Zone (CMZ). The CMZ contains roughly 10%of our Galaxy’s molecular gas. The Sgr A–D clouds are well-known massive molecular clouds in
the CMZ. The CMZ is thought to be associated with the NB.
The central region (R ≤ 5 pc) is known as the Sgr A complex which was found as a very strong
radio source in the dawn of the radio astronomy [37]. Now the Sgr A complex is decomposed totwo components, Sgr A West and Sgr A East (ﬁgure 2.2). Sgr A West is a thermal radio source
3

4

CHAPTER 2. REVIEW

Figure 2.1: Zoom picture of the Galaxy: a: Schematic cross section of the Galaxy. b: False
color COBE/DIRBE image of the Galactic Bulge c: Schematic representation of Galactic andNuclear Bulge. d: Schematic representation of the Sgr A Radio and Giant Molecular Cloud
Complex. [68]

2.2. DISCOVERY OF THE GALACTIC DIFFUSE X-RAY EMISSION AND ITS FE Kα LINES5
which has a characteristic spiral structure. A compact radio source, Sgr A∗ , is located in thecenter of Sgr A west. Observations of the orbital motions of the stars near Sgr A∗ reveal thatthe mass of Sgr A∗ is ∼ 4 × 106 M⊙ (e.g. [25]) and thus Sgr A∗ is thought to be the nuclear
massive black hole and the dynamical center of the Milky Way galaxy. The X-ray from Sgr A∗
is detected with Chandra and the ﬂux is 1.3 × 10−13 erg cm−2 s−1 in the 2–10 keV band (ﬁgure2.3) [3]. The X-ray luminosity of Sgr A∗ is 2 × 1033 erg s−1 , which is much lower than those ofother active galactic nuclei (∼ 1042−49 erg s−1 ). The spirals of Sgr A West are thought to be
the ﬂows of gas falling into Sgr A∗ [67]. Sgr A East is a non-thermal radio shell embedding Sgr
A West. Now it is thought to be a supernovae remnant. More detail review of Sgr A East is
shown in section 2.6.

Figure 2.2: VLA radio continuum image of the GC showing the shell-likestructure of non-thermal Sgr A East
(light blue and green) and the spiralshaped structure of thermal Sgr A
West (red) at λ= 6 cm [144].

2.2

Figure 2.3: Chandra image of Sgr A West in the
0.5–7 keV band with contours of VLA 6 cm [3].
The central radio and X-ray point source is SgrA∗ .

Discovery of the Galactic Diﬀuse X-ray Emission and its Fe
Kα Lines

The large extended X-ray emission from our Galaxy – the Galactic Diﬀuse X-ray Emission,
hereafter GDXE — has been known since the dawn of the X-ray astronomy. For example, in
1970, just 8 years later after the ﬁrst X-ray observation of Giacconi et al., Cooke et al. [14]observed with the Skylark rocket and reported a clear excess of X-ray ﬂux on the Galactic planewhere no strong X-ray point sources present. With early X-ray astronomy satellites such as
UHURU, ARIEL-V, HEAO-1 and EXOSAT, it was reveal that the GDXE extends over morethan 100◦ along the Galactic plane but only ∼ 2◦ across it (see [137] and and references therein).
Thus it has been called as the Galactic “ridge” X-ray emission.
One of the most interesting features of the GDXE is its strong Fe Kα line emissions. Japanese
X-ray satellites have contributed the discovery and the study of these lines. With the Gas
Scintillation Proportional Counter (GSPC) on the Tenma satellite, the Fe Kα line feature was

6

CHAPTER 2. REVIEW

discovered [44] (ﬁgure 2.4). With the Ginga satellite, it was revealed that the Fe Kα emission
extends over 100◦ along the plane and concentrates sharply at the center region [47] (ﬁgure 2.5).
Due to the high energy resolution of the CCD camera (Solid state Imaging Spectrometer; SIS)
on the ASCA satellite, the Fe Kα line feature in the GC region is resolved to three lines at theenergy of 6.4, 6.7 and 6.97 keV [48] (ﬁgure 2.6). The 6.7 and 6.97 keV lines correspond to the Fe
XXV Kα and XXVI Kα line emissions. These Fe ions imply the existence of hot plasma whose
temperature is kT = 5–10 keV or T = several 107 K. The 6.7 and 6.97 keV lines correspond to
the Fe XXV Kα and XXVI Kα line emissions. These Fe ions imply the existence of extremely
hot plasma whose temperature is kT = 5–10 keV or T = several 107 K. The 6.4 keV line, on
the other hand, corresponds to the Fe I (or low-ionzied Fe) Kα line emission. It indicates thatcold matter should be excited by some high-energy process to radiate X-rays.

Figure 2.4: X-ray spectrum of the Galactic
ridge obtained with theTenma satellite [44].

2.3

Figure 2.5: The proﬁle of the X-ray emission in the 2–18 keVenergy range (a) and that of the iron-line emission (b) on the
Galactic plane scanned with the Ginga satellite. The line center energy is plotted in (c). [47]Candidates for the Origin of the Fe XXV Kα and Fe XXVIKα Line Emissions

Thin thermal plasma is the most probable origin of the Fe XXV Kα and Fe XXVI Kα line
emissions. We show the outline of X-ray emission from optically thin thermal plasma, especiallycollisional ionization equilibrium (CIE) plasma in appendix A. In the following sections, we show
candidate emitters of the highly ionized Fe Kα lines in the case of the truly diﬀuse plasma origin(section 2.3.1) or the unresolved point sources origin (section 2.3.2). Charge exchange origin isreviewed in section 2.3.3.

2.3.1

Truly Diﬀuse Plasma Origin

One possible origin of the highly ionized Fe Kα lines in the GDXE is truly diﬀuse hot plasma
ﬁlling the Milky Way galaxy. The most probable candidate to supply such diﬀuse hot plasma aresupernova explosions. The ejected material from a supernova explosion expands supersonicallyand generate shock waves, which heat interstellar matter and the ejected material. As a result,2.3. CANDIDATES FOR THE ORIGIN OF THE FE XXV Kα AND FE XXVI Kα LINE EMISSIONS7

Figure 2.6: X-ray spectrum of the Galactic center obtained with the ASCA satellite [48].
the plasma of a supernovae remnant (SNR) are made. The integration of the SNR plasma might
be the origin of the GDXE [45].
SNRs have, however, two defects as the origin of the GDXE. First, the typical temperatureof the SNRs is lower than that of the GDXE (kT = 5–10 keV). Even in the case of young
SNRs whose age t is less than 1000 yr, the temperature kT is lower than 4 keV. The observed
temperatures of famous young SNRs are summarized in table 2.1. The second problem is theenergy budget. If truly diﬀuse hot plasma gas ﬁlls the Galactic plane (V ∼ 1066 cm3 ), its totalluminosity and energy are about 2 ×1038 erg s−1 and 1056 η 1/2 erg (η is a ﬁlling factor) [41]. The
total energy is extremely large. The sound speed of such high temperature plasma (vs ∼ 1400
km s−1 for kT = 6 keV) is higher than the escape velocity from the Galactic plane (according to[11], vescape ∼ 700 km s−1 at the region on the Galactic plane of 3 kpc from the center). The hot
plasma will escape from the Galactic plane (thickness ∼ 1 kpc) within ∼ 106 years. The energy
input of 1050 erg yr−1 is required to sustain this plasma. The most probable energy suppliers
are supernovae whose energy is ∼ 1051 erg. The required supernova rate is about one per 10
years [45], which is one order lager than the recently estimated rate [17, 65]. The luminosity of
the GDXE cannot be explained, either. The X-ray luminosity (LX ) of the SNRs is LX = 1035−36
erg s−1 (see table 2.1 and [45, 131]). To explain the luminosity ∼ 1038 erg s−1 of the GDXE by
young SNRs (t ∼ 1000 yr), one supernova must occur every 10 years in the Galaxy [45] again. If
the origin of the GDXE is the truly diﬀuse hot plasma, unknown plasma heating process might
exist in the Milky Way galaxy.

2.3.2

Faint Point Sources Origin

Some kinds of X-ray point sources emit the highly ionized Fe Kα lines. The superposition ofsuch point sources might be origin of the GDXE. For example, binaries of Wolf-Rayet and OB
star have the spectra of thin-thermal plasma at kT = 0.1–6 keV and emit the Fe XXV Kα line
(e.g. [29]). Wolf-Rayet binaries are, however, so luminous (LX > 1033 erg s−1 ) that they are

8

CHAPTER 2. REVIEW

Table 2.1: Examples of young supernova remnants.

Cas A
KeplerTycho
SN1006kT (keV)
3.8
4.5
2.9
1.9

LX (erg s−1 )1 × 1036
3 × 1035
2 × 1035
1 × 1035

Thermal energy (erg)
5 × 1049
4 × 1049
1 × 1049
2 × 1049

Age t (yr)
∼ 300
∼ 300
∼ 400
∼ 1000Ref.
[132][9, 27][132][38, 46]

resolved easily. In addition, they are not numerous and thus cannot be the origin of the GDXE.
The candidate source should be faint, numerous and hot.
Cataclysmic variables (CVs) are the possible origin of the GDXE. They are binaries of awhite dwarf and an ordinary star and the most numerous accretion-powered X-ray sources.
Their local space density is 3 ×10−5 pc−3 [114] or 1.2 × 10−5 M⊙ −1 [115].
Systems with non-magnetized white dwarfs, which comprise 80% of CVs, have luminosities
between 1029.5 –1032 erg s−1 and spectra that can be described with kT =1–25 keV plasma froman accretion shock (ﬁgure 2.7).
Magnetized cataclysmic variables (mCVs), which are referred to as polars and intermediate
polars (IPs) depending on whether or not the rotational period of the white dwarf is synchronized
to its orbital period. The magnetic ﬁeld of the polars is so strong (107−7.8 G) that the rotationaland orbital periods of are synchronized and accretion disks are not formed. IPs are thought
to form accretion disks (ﬁgure 2.8). In mCVs, the matter spilled over Roche Lobe of the staraccretes onto the magnetic pole and becomes hot by the free-fall shock. Thus, they have spectrathat can be described with kT =1–25 keV plasma and often have X-ray pulses synchronizing
to their rotational periods. IPs have hard continuum components (Γ ∼ 0) due to the selfabsorption. The typical luminosities of the polars and the IPs are LX = 1030−33 and 1031−33 erg
s−1 .
The line intensity ratio of Fe XXVI Kα to Fe XXV Kα of the CVs is typically ∼ 0.5 andhigher than that of the GDXE (0.3–0.4). The equivalent widths of the Fe XXV Kα and XXVIKα lines are, however, typically ∼ 200 and ∼ 100 eV and too small to explain those of the
GDXE, ∼ 400 and ∼ 150 eV (ﬁgure 2.9).

Figure 2.7: Schematic view of a non-magnetic CV [33]. Matter transferred through the accretion
disk onto a white dwarf releases its Keplerian motion energy and become plasma by the shock.

2.3. CANDIDATES FOR THE ORIGIN OF THE FE XXV Kα AND FE XXVI Kα LINE EMISSIONS9

Figure 2.8: Schematic view of magnetic cataclysmic variables. Left: Polar [15]. Right: Intermediate Polar (IP) [93].

Figure 2.9: Relation between the ionization temperature (horizontal axis) and the summed
equivalent width of the Fe XXV Kα and Fe XXVI Kα lines of mCVs [20]. The circle andrectangle marks show polars and IPs respectively. The ionization temperature is determinedfrom the line intensity ratio of Fe XXVI Kα to Fe XXV Kα. The typical values of the GDXE,
the ioniztion temperature ∼ 6 keV and the summed equivalent width ∼ 550 eV, are shown by
the red dashed lines.

10

2.3.3

CHAPTER 2. REVIEW

Charge Exchange Origin

The charge exchange scenario is that bare or hydrogen-like Fe ions in cosmic-ray capture aelectron in the interstellar matter and emit the Fe XXVI Kα and Fe XXV Kα lines. The
scenario was proposed mainly to explain the line broadening observed with ASCA [122]. The
spectrum of the GDXE in the Sgr A region revealed with Suzaku, however, prefers the thermal
plasma origin to the charge exchange origin (see section 2.5.1).

2.4

Candidates for the Origin of the Fe I Kα Line Emission

The Fe I Kα line is emitted from the cold matter when a K-shell electron in a neutral Fe atom
is given the energy higher than the K-edge energy (7.1 keV) and ionized.
The origin of the Fe I Kα line in the GDXE can be truly diﬀuse source or integrated point
sources. In the case of the diﬀuse scenario, the cold interstellar diﬀuse gas is ionized by hard
X-ray photons (photoionization) or electrons in the cosmic-ray (collisional ionization). In section
2.4.1, we show the diﬀerence between the photoionization and the electron-collisional ionization.
In the case of the point sources scenario, the most probable candidates are CVs. We review theFe I Kα emission of CVs in section 2.4.2.

2.4.1

Truly Diﬀuse Interstellar Gas Origin

The continuum component is emitted by the Thomson scattering of the incident X-ray photon
in the photoionization case, and by the bremsstrahlung of the incident electrons in the collisional
ionization case. The equivalent width is diﬀerent between the photoionization and the collisional
ionization. In the case of the photoionization, the calculated equivalent width of the Fe I Kα lineis ∼ 700 eV assuming Γ ∼ 2 for the incident photons [87]. In the case of the electron-collisional
ionization, the calculated equivalent width of the Fe I Kα line is ∼ 300 eV assuming that theenergy distribution of the incident electrons dN/dE is proportional to E −2 [124].
Tatischeﬀ [124] pointed out that, in the electron-collisional ionization scenario, the incident
electrons cannot penetrate the interstellar gas thicker than NH = 4 × 1021 cm−2 and thus theintrinsic absorption of the bremsstrahlung should be less than this value.
In the Galactic center region, clumpy Fe I Kα line emissions (hereafter ”6.4 keV-clumps”),
which are correlated to the distribution of the molecular clouds in the CMZ, have been found[48, 50, 85]. The spectral parameters of the 6.4 keV-clumps are summarized in table 2.2. Becauseof the large equivalent widths of 1–2 keV and large absorptions of NH = 1023−24 cm−2 , the Fe I
Kα line emissions of the 6.4-keV clumps are thought to be the photoionization origin and calledas X-ray reﬂection nebulae (XRNe). The past ﬂare of Sgr A∗ is suggested as the external X-ray
radiation source [35, 53, 54, 77, 108] for the 6.4-keV clumps in the CMZ. The origin of the Fe I
Kα in the GDXE might be thinly extended XRNe but it is still an open question.

2.4.2

Faint Point Sources Origin

CVs emit the Fe I Kα lines. The origin is thought to be the photoionization of the cold matterin the accretion disk or the surface of the white dwarf irradiated by the hot plasma emission(see ﬁgure 2.7 and 2.8). The typical equivalent width of the Fe I Kα line of CVs is ∼ 100 eV.
It is small because the continuum includes not only the Thomson scatter component but also

2.5. RECENT OBSERVATION OF THE GDXE

11

Table 2.2: Properties of the 6.4-keV Clumps in the GC Region.

Sgr A source 1
Sgr A source 2
Sgr B2 cloud
M 0.74–0.9
M 0.51–0.10
M 359.43–0.07
M 359.47–0.15

EW6.4 ∗ (keV)
1.2 ± 0.1
1.1 ± 0.1
1.13+0.05−0.02
1.6+0.4−0.3
1.4 ± 0.3
2.2+0.3−0.4
2.0+0.2−0.1

Γ1.83 ± 0.03
1.86+0.03−0.02
3.2+0.9−0.6
1.4+0.4−0.7
1.8+0.4−0.3
1.7+0.1−0.2
1.6+0.3−0.1

NH (1022 cm−2 )
23+4
−3
23+2
−4
96+25
−8
40+14
−11
15+2
−1
9+5
−4
8+4
−1

Ref.
[55]
[55]
[53]
[53]
[85]
[84]
[84]

∗ The equivalent width of the Fe I Kα line.

direct component. Note that the equivalent width strongly depends on the geometry of the CV
structure.

2.5

Recent Observation of the GDXE

Recently, the study of the GDXE has been improved very much, especially with Suzaku andChandra. Due to the good energy resolution of the Suzaku satellite, it has been revealed that the
GDXE spectra are diﬀerent between the Galactic center (GC) and the ridge [51, 143]. With thegood spatial resolution of the Chandra satellite, some fraction of the GDXE has been resolved
into point sources. Here we review the results of the recent GDXE observations.

2.5.1

GDXE in the Sgr A Region (|l| < 0.3◦ )The spectrum of the GDXE in |l| < 0◦ .3, the Sgr A region, is studied well by Koyama et al. [50]with the Suzaku XIS (ﬁgure 2.10). The center energy of the Fe XXV Kα line is determined to be
6680+3−6 eV with the careful calibration. This observed value supports not the charge exchange
scenario (6666 eV) but the CIE plasma scenario (6680 eV) as the origin of the Fe XXV Kα line
emission. The line intensity ratio of Fe XXVI Kα to Fe XXV Kα is determined to be 0.33 ± 0.02.
Using the intensity ratios among highly ionized Fe and Ni Kα and Kβ lines, the temperature of
the plasma is determined to be kT = 6.5 ± 0.1 keV. The line ratios severely constraint the shapeof the bremsstrahlung continuum of the thermal component. The non-thermal hard tail with
Γ = 1.4+0.5−0.7 was successfully detected for the ﬁrst time by subtracting the thermal continuum
from the observed one. The line intensity ratio or the temperature along the longitude is almost
constant (0.3–0.4, kT =6–7 keV) in the region of −0◦ .4 < l < 0◦ .2. The intensity proﬁle of the
Fe XXV Kα line disagrees with that of X-ray point sources detected with Chandra (ﬁgure 2.11).
This result indicates that the major fraction of the Fe XXV Kα line emission is not due to theintegration of the unresolved point sources but truly diﬀuse.
Koyama et al. [55] analyzed the GDXE spectra in the Sgr A region (|l| < 0◦ .3) and foundthe anti-correlation between the equivalent widths of the Fe I Kα and Fe XXV Kα lines (ﬁgure2.12). The relation between the equivalent widths of the Fe I Kα (EW6.4 ) and Fe XXV Kα(EW6.7 ) lines can be ﬁtted with the equation EW6.7 + 0.50 (±0.06) × EW6.4 = 0.62 (±0.07)
keV. Using this relation, they phenomenologically decomposed the GDXE in the Sgr A into two

12

CHAPTER 2. REVIEW

Figure 2.10: The GDXE spectrum in
the Sgr A region with the Suzaku XIS.
The models of the thermal plasma of
kT = 6.5 keV and the non-thermal
hard tail are also shown [50].

Figure 2.11: The intensity proﬁle in the Sgr A
region of Fe XXV Kα are given by the squares,
while the integrated point sources in the 4.7–8
keV band are plotted by the crosses [50].

components; the Fe I Kα line emission with its associated continuum (6.4-component) and Fe
XXV Kα line emission with its associated continuum (6.7-component). From the line intensity
ratio of Fe I Kα to Fe XXV Kα in the Sgr A region (∼ 1 : 1), Koyama et al. determined thatthe ﬂux ratio of the 6.4-component to the 6.7-component is 1:2. Combined with the result [50]that the ﬂux ratio of the non-thermal component to the thermal component is 1:1, Koyama et
al. determined that the lower limit of the contribution of the point sources is 1/6 [55].
Using the Chandra data, Muno et al. [78] found that the integrated spectrum of the point
sources in the Sgr A region (R < 9′ ) has three Fe Kα lines like the GDXE (ﬁgure 2.13). The line
intensity ratio of Fe XXVI Kα to Fe XXV Kα is ∼ 0.6, similar to those of CVs near the Sun.
The point sources in the Sgr A region also have hard continuum emissions (Γ ∼ 0.9) and thetypical luminosities of 1031−33 erg s−1 . Muno et al. [78] suggest that most of the point sourcesin the Sgr A region are IPs. The equivalent width of the Fe XXV Kα line of the point sourcesin the Sgr A region is ∼ 400 eV and larger than that of a typical CV, but close to that of the
GDXE.
Revnivtsev et al. [103] observed the region of about 10 arcmin2 at (l, b) = (−0◦ .09, −0◦ .03) for
918 ks with Chandra and resolved 40% of the GDXE into X-ray point sources at the sensitivity
limit of 1.4 × 10−15 erg cm−2 s−1 (2–10 keV).

2.5.2

GDXE in the Inner Region of the NB (0◦ .3 < |l| < 2◦ )The GDXEs in the Sgr B (0◦ .3 < l < 0◦ .8) and C (−0◦ .8 < l < −0◦ .6) regions have been studied
in detail by Nobukawa et al. [85] and Nakajima et al. [84] with the Suzaku data. Nobukawa et al.
[85] reported that the line intensity ratio of Fe XXVI Kα to Fe XXV Kα of the Sgr B region is 0.3–
0.4, same as that of the Sgr A region. The line intensity proﬁle along the longitude also connects
to that of the Sgr A region smoothly. Nakajima et al. [84] analyzed the Suzaku data of the Sgr
C region in the same way as Koyama et al. [55] and found the anti-correlation of the equivalent
widths between the Fe I Kα and Fe XXV Kα lines. The obtained relation is diﬀerent from thatof the Sgr A region and given by the equation EW6.7 + 0.22 (±0.12) × EW6.4 = 0.53 (±0.06)
keV (ﬁgure 2.14).

2.5. RECENT OBSERVATION OF THE GDXE

Figure 2.12: Relation between the equivalent widths of the Fe I Kα and Fe XXV
Kα lines in the Sgr A region. The solid line
shows the relation of EW6.7 +0.50 (±0.06)×EW6.4 = 0.62 (±0.07) [55].

13

Figure 2.13: The integrated spectrum of
the point sources detected with Chandra in
the Sgr A region [80].

In both of the Sgr B and C regions, giant molecular clouds (e.g [129]) and 6.4-keV clumps
were observed (see table 2.2). Inui et al. [35] and Koyama et al. [54] found the time variabilityof the line intensity of the Fe I Kα line in the Sgr B2 cloud and suggested that the time validityof the Fe I Kα line might reﬂect that of the past Sgr A∗ ﬂare.

Figure 2.14: Relation between the equivalent widths of the Fe I Kα and Fe XXV Kα lines in the
Sgr C region. The solid line shows the relation of EW6.7 + 0.22 (±0.12) × EW6.4 = 0.53 (±0.06).
The dashed line shows the relation in the Sgr A region [84].

2.5.3

GDXE in the Galactic Ridge (|l| > 8◦ )Using the Suzaku data, Ebisawa et al. [19] resolved the Fe Kα complex into the three linesand detected the Fe I Kα line from the ridge spectrum (l = 28◦ .5) clearly for the ﬁrst time.
Yamauchi et al. [143] studied the Fe Kα lines of the GDXE in the ridge region (|l| > 8◦ )

14

CHAPTER 2. REVIEW

systematically. They found that the line intensity ratio of Fe XXVI Kα to Fe XXV Kα in theridge is 0.22 ± 0.06, signiﬁcantly smaller than that in the GC region (|l| < 1◦ ), 0.38 ± 0.02 (ﬁgure2.15). These line ratios correspond to the plasma temperatures of kT ∼ 7 keV (0.38) and ∼ 5.5
keV (0.22). Yamauchi et al. [143] suggested that the diﬀerence in the line intensity ratios mightcome from the spatial diﬀerence of the plasma in the diﬀuse scenario or the candidate sources in
the GC region have a higher temperature plasma emission than those in the ridge systematicallyin the point sources scenario.
Using the data of the RXTE satellite [10], Revnivtsev et al. studied the intensity proﬁle of
the GDXE in the 3–20 keV band in |l| = 1–100◦ and compared it with the near-infrared (NIR)
observation [101]. The intensities of the GDXE and the NIR are correlated. Revnivtsev etal. calculated the X-ray luminosity to the stellar mass ratio of the GDXE in the ridge, LX /M= (3.5 ± 0.5) × 1027 erg s−1 M⊙ −1 , which is close to that of the point source in the vicinityof the Sun, LX,local /M = (5.3 ± 1.5) × 1027 erg s−1 M⊙ −1 and suggested that the origin of the
GDXE is the integration of X-ray point sources of the known class like the solar neighborhood.
Revnivtsev et al. also studied the line intensity proﬁle of the Fe Kα complex of the GDXE in|l| = 1–100◦ and showed that it agrees with that of the NIR (the left panel of ﬁgure 2.16).
Revnivtsev et al. did not resolve the three Fe Kα lines because they used a proportional counter(PCA) on RXTE, whose energy resolution is ∼ 18% at 6 keV [102].
Using the Chandra data of the 200 ks exposure, Ebisawa et al. [18] and Revnivtsev etal. [103] resolved 10 or 19% of the GDXE at l = 28◦ .5 into point sources at the sensitivity limitof 3 ×10−15 erg cm−2 s−1 (2–10 keV) or 1.2 ×10−15 erg cm−2 s−1 (1–7 keV).

Figure 2.15: Proﬁle of the line intensity ratio of Fe XXVI Kα to Fe XXV Kα along thelongitude [143].

2.5.4

GDXE in the Galactic Bulge (|b| > 1◦ )Using the RXTE data, Revnivtsev et al. revealed that the intensity proﬁles of the Fe Kα
complex (the right panel of ﬁgure 2.16) [102] and the continuum [101] (3–20 keV) in the bulge
match that of the NIR similarly to the ridge.
Using the Chandra data with 1 Ms exposure, Revnivtsev et al. resolved 85% of the GDXE
in the 2.6 arcmin-radius region at (l, b) = (0◦ .113, −1◦ .424) into point sources at the sensitivity
limit of 10−16 erg cm−2 s−1 in the 0.5–7 keV band (ﬁgure 2.17) [105].

2.5. RECENT OBSERVATION OF THE GDXE

15

Figure 2.16: The intensity proﬁle of the Fe Kα complex along the longitude (left) and thelatitude (right) observed with the RXTE. The dashed line shows the intensity proﬁle of the NIR(COBE/DIRBE, 4.9 µm) [102].

Figure 2.17: GDXE spectrum and its resolved fraction. (a) Spectra collected by Chandra.
Black data points show the spectrum of the total emission from the 2.6 arcmin-radius region;
the collective spectrum of all detected point sources is presented in blue and the spectrumof the remaining unresolved emission in the current observations is in red. (b) Fraction of
the X-ray emission resolved by Chandra into point sources as a function of X-ray photon
energy [105].

16

2.6

CHAPTER 2. REVIEW

Review of Sgr A East

Sgr A East is a non-thermal radio-shell structure ﬁlled with thermal X-ray emission in the
GC region. In X-rays, Chandra (Maeda et al. [63]) found a center ﬁlled morphology, and thespectrum consisted of strong Kα lines from highly ionized S, Ar, Ca, and Fe. The spectrum was
ﬁtted with a simple isothermal plasma model of temperature ∼2 keV, an absorption column ofNH ∼ 1023 cm−2 and a luminosity of ∼ 8 × 1034 ergs s−1 in the 2–10 keV band. From the radio
and X-ray morphologies, Sgr A East is classiﬁed to be a “mixed morphology” SNR [63].
XMM-Newton (Sakanao et al. [112]) found a two-temperature structure of 0.94+0.06−0.09 keV
+0.5
and 4.2−0.4 keV. The age of Sgr A East is estimated to be 4000—10000 yr [63, 112] so thetemperature of kT ∼ 4 keV is hot as SNRs. With follow-up deep Chandra observation, Park et
al. [92] analyzed the spectra spatially divided in the Center, Plume and North regions of Sgr A
East. The Center and North spectra require two-temperature structures composed of kT ∼ 1+19
keV and kT = 5.3+0.7−0.6 keV (Center) or kT = 11−4 keV (North). Such hot temperatures are close
to that of the GDXE. Sgr A East is located in the GC region, where extreme physical conditions
are manifested. These are a high stellar density, dense stellar gas, a strong magnetic ﬁeld and thepresence of a massive black hole. Hot SNRs like Sgr A East might be made in such an extreme
environment and contribute to the strong GDXE in the GC region. Chandra found many clumpy
and ﬁlamentary structures in Sgr A East. Some of them must be non-thermal X-ray ﬁlaments,
which would be related to the extreme conditions. The existence of the non-thermal component,
however, brings the uncertainty to the previous temperature determinations, which were mainly
based on the shape of the thermal bremsstrahlung continuum.

2.7

Summary of the Recent Study and Subjects Tackled in This
thesis

The recent results we have reviewed above are summarized in table 2.3. This table shows theintermediate region between the GC and ridge regions has not been studied well. In this thesis,
we study the subjects shown in green.

2.7. SUMMARY OF THE RECENT STUDY AND SUBJECTS TACKLED IN THIS THESIS17

Table 2.3: Summary of the recent studies of the GDXE. The results written in red (blue) support
whether diﬀuse (point sources) origin. The issues studied in this thesis are shown in green.

Sgr A region|l|<0°.3

Between GC &Ridge |l|=0°.3–3°

Ridge|l|>8°

Comparison
between the
profiles of the
GDXE andpoint sources• The profile of
Fe XXV K! isdifferent fromthat of X-ray
point source
detected with
Chandra.

•Not studied.

•The profile of the
GDXE of |l|>1°
matches that of the
NIR observation.

Equivalent
widths (EW)of Fe K!

•The EWs of the
GDXE are toolarge as known
CVs.
•The EWs of thepoint sources in
the GC are aslarge as those of
the GDXE.

!The EWs of the
GDXE has notbeen not studied
well.
!The spectra ofpoint sources inthis region havenot been studied
well.

!The EWs of the
GDXE are too large
as known CVs.
•The EWs of known
CVs near the Sun are
too small as the
GDXE.

EW of FeXXV K! v.s.
EW of Fe I K!

!They are anticorrelated.

!Except for the
Sgr C region,other regions are
not studied.

!Not studied.

Intensity ratioof Fe XXVI K! to Fe XXV
K!

!0.33
•It is too large as
well-known
SNRs. Sgr A
East might be
hot ?

•Not studied.

!0.22

!Not
studied.

Fraction of
the GDXE
resolved into
point sources.

•40%!The flux of the
GDXE itself has
not been studied.

•19%

•85%

Bulge
|b|>1°

Chapter 3

Instrument
3.1

Overview of the Suzaku Satellite

Suzaku (ASTRO-E2) [69] is the ﬁfth Japanese X-ray astronomy satellite following Hakucho,
Tenma, Ginga and ASCA (ﬁgure 3.1). It was launched by Japan Aerospace Exploration Agency(JAXA) with the M-V launch vehicle from Uchinoura Space Center (USC) on 2005 July 10.
The length is 6.5 m after an extensible optical bench (EOB) deployment. The total mass is1700 kg. The attitude is controlled by momentum wheels and magnetic torquers to make theangle between the Sun and the solar panel less than 30◦ . The area of the sky accessible at a time
is a belt within which the Sun angle is between 65◦ and 115◦ because the pointing direction islimited by the power constraint of the solar paddle, which is perpendicular to the optical axesof the instruments.

Figure 3.1: Schematic view of the Suzaku satellite in orbit [69].
The orbit of Suzaku is a near circular orbit at 570 km altitude with an inclination angleof 31◦ . Compared with the high-altitude ellipsoidal orbits of Chandra [138] and XMM-Newton[36], this low-altitude orbit has a low and timely stable non-Xray background (NXB). Thus
Suzaku is good at detecting faint and diﬀuse X-ray emissions. On the other hand, the observing
eﬃciency is lower than those of Chandra and XMM-Newton. Most targets are occulted by the
Earth for about one third of each orbit. Observation is also interrupted by passages of the South
Atlantic Anomaly, where the inner Van Allen radiation belt makes its closest approach to the
Earth’s surface and thus the intensity of cosmic-rays is extremely high. The observing eﬃciency
is about 40%.
19

20

CHAPTER 3. INSTRUMENT

XRT−S
Sun shade
XRT−I(4 units)

EOB
Star
tracker

Side panel(8 panels in total)

Gyro−
scopes
XRS

Base
panel

HXD
XIS
(4 units)

Figure 3.2: Side view of Suzaku with the internal structures [69].
Suzaku is equipped with ﬁve sets of X-Ray Telescopes (XRT [117]). At the focal plane,
there are four X-ray sensitive CCD cameras (the X-ray Imaging Spectrometer; XIS [50]) andan X-ray micro-calorimeter (the X-Ray Spectrometer; XRS [43]). The Hard X-ray Detector(HXD [126]), a non-imaging collimated hard X-ray scintillating instrument, is also installed.
The optical axes of these instruments are co-aligned. Thus Suzaku can observe a celestial object
with these instruments at the same time and obtain a broadband spectrum which extends from
0.2 to a few hundred keV. The arrangement of these instruments is shown in ﬁgure 3.2.
Before the start of astronomical observations, XRS was out of function because a thermal
short between liquid helium and solid neon tanks resulted in the liquid evaporation. The dataof the HXD are not used in this thesis. The details of the XRT and the XIS are show in thefollowing sections.

3.2

X-Ray Telescope (XRT)The XRT [117] is a X-ray telescope with grazing-incidence reﬂective optics consisting of tightly
nested, thin-foil conical mirror shells (ﬁgure 3.3). There are ﬁve sets of the XRT onboard theSuzaku. One is for the XRS (XRT-S) and the others are for the XIS (XRT-I).
Refractive indices of matter to X-rays is less than 1 and thus X-rays cannot be focused bya refracting telescope. X-rays are full-reﬂected if the incident grazing angle is shallow enough(≤ 1◦ ), and thus X-rays are focused using this full-reﬂection. An approximate Wolter-I optics
is used in the XRT. The shapes of the primary and secondary reﬂectors, which are a paraboloid
and hyperboloid respectively in the original Wolter type I optics, are both approximated bycones (ﬁgure 3.4). This approximation ultimately limits the imaging capability of the mirror
shell but contributes to make the weight of XRT reduced. The focal length of the XRT is 4.75 m.
The mirror shell is made of aluminum and a gold layer adhered to it. About 170 mirror
shells mounted concentrically in one XRT module. Due to the tightly nested thin-foil structure,3.2. X-RAY TELESCOPE (XRT)

21

Figure 3.3: Picture of the XRT [117].

Figure 3.4: Schematic view of an X-ray telescope composed of nested mirror shells with a
Wolter-I optics [72].

22

CHAPTER 3. INSTRUMENT

the XRT provides a large aperture eﬃciency with a moderate imaging capability in the energy
range of 0.2–12 keV despite a total weight of only 20 kg. For example, the eﬀective area of one
XRT is as large as that of the HRMA, which is the X-ray telescope onboard Chandra. The
weight of the HRMA is, however, ∼ 1500 kg. The total on-axis eﬀective area of the four XRTs
is shown in ﬁgure 3.5 compared with the other X-ray missions.
The eﬀective area for oﬀ-axis incident X-rays is smaller than that for on-axis X-rays because
oﬀ-axis X-rays are partially obscured by neighbor mirror shells. This ”vignetting” eﬀect isprominent at high energy. A maximum incident angle for full-reﬂection is smaller for higher
energy X-rays than for lower ones. High energy X-rays are reﬂected only at inner mirror shells,
where the cone angles is small and the mirror shells are tightly packed (see ﬁgure 3.4), andhighly subject to shadowing. The oﬀ-axis angle dependences of the eﬀective area for the soft
(3–6 keV) and the hard (8–10 keV) bands are shown in ﬁgure 3.6.

Figure 3.5: Total eﬀective area of thefour XRT-Is compared with these of
XMM-Newton and Chandra. Transmissions of the thermal shield and theoptical blocking ﬁlter, and the quantum eﬃciency of the CCD are all taken
into account [117].

Figure 3.6: Oﬀ-axis angle dependenceof the eﬀective area normalized at the
optical axes in the soft and hard bands
shown with black and red points respectively (vignetting). Data are obtained with the Crab nebula observations. The solid lines show the simulated models [117].

Half power diameter (HPD) is the diameter of a circle which contains 50% of ﬂux of a point
source. HPD is a standard value to evaluate the imaging capability. The HPDs of the XRTs
are 1′ .8–2′ .3. Point spread function (PSF) is an intensity proﬁle of a point source as a functionof radius from the point source. Enclosed energy function (EEF) is a integration of PSF andshows the ratio of the ﬂux of the enclosed circle to the total ﬂux of a point source. PSF andEEF of one of the XRTs are shown in ﬁgure 3.7 and 3.8.
Each XRT is made of four sections (“quadrants”, see ﬁgure 3.3). Since the spaces between
the quadrants are dead areas, point source images appear to be a cross.
Tightly-nested structures increases the possibility of reﬂection other than the normal double
reﬂection within the telescope structure. Because of these unexpected incidences, X-ray of brightsources outside the ﬁeld of view reaches the focal plane. This is called stray light and it has
been a large problem for nested X-ray telescopes (e.g. ASCA/XRT). In the case of Suzaku XRT,3.3. X-RAY IMAGING SPECTROMETER (XIS)

Figure 3.7: Point spread function of
the XRT of XIS 0. This is obtainedwith the observation of SS Cyg, which
is a bright point-like source [117].

23

Figure 3.8: Enclosed energy function
XRT of XIS 0. This is obtained with
the integration of the PSF in ﬁgure 3.7[117].

front-mounted pre-collimators are installed and decrease the stray light by one order comparedwith that of ASCA [72].

3.3

X-Ray Imaging Spectrometer (XIS)

3.3.1

Introduction of X-ray Charge Coupled Device Cameras

The XIS consists of MOS-type charge coupled devices (CCDs). An X-ray CCD is a silicon
semiconductor detector with two-dimensional imaging capability. Each pixel consists of Metal(polysilicon), Oxide (SiO2 ) as an insulator and p-type Semiconductor (silicon). This is called a
MOS structure (ﬁgure 3.9).

Metal
Oxide
Semiconductor

Figure 3.9: Schematic view of a MOS structure.
A brief explanation of the principle of an X-ray CCD is the following. In a p-type semiconductor, holes are the major carriers. A positive voltage applied to the metal electrodes repels
holes in the semiconductor and leaves behind a depleted layer where an electric ﬁled exists.
When an X-ray photon enters into a pixel, photoelectrons are created and then electron-hole

24

CHAPTER 3. INSTRUMENT

pairs are generated. The number of the generated electrons is proportional to the energy of theincident photon. The average energy required to form an electron-hole pair is 3.65 eV in −90◦ C.
The holes and electrons are separated by the electronic ﬁeld in the depletion layer before they
recombine. The electrons are collected into the potential wells beneath the electrodes as signal
charges. After the exposure period, the signal charges are transferred to the readout node byclocking the voltages of the electrodes.
The CCDs of the XIS employee the three-phase transfer type electrode structure. As shownin ﬁgure 3.10, in this type, three electrodes are deposited on each pixel. By applying the threephase clock voltage to the electrodes, the potentials beneath the electrodes change periodicallyand the signal charges are transferred.
The CCD of the XIS is classiﬁed as a frame transfer (FT) CCD in the scheme of the chargetransfer and readout. A FT CCD consists of an imaging area and a frame-store region shown inﬁgure 3.11. The imaging area is exposed but the frame-stored region is shielded from light. Afterthe exposure period, the signal charges are transferred in a short time from the imaging area tothe frame-store region. The signal charges in the frame-store region are transferred verticallyto the horizontal shift register by each row. The charges in the horizontal shift register aretransferred to the readout node by each pixel. This process is repeated until the charges in all
the pixels of the frame-store are read out. During this readout of the charges, new exposure
starts in the imaging area. The signal charges for the next exposure time can be read out
continuously.
In this way, amounts of charges in each pixel are measured and an image and a spectrum aretaken with an X-ray CCD camera. An X-ray CCD camera has also time resolution dependingon the exposure time.
1 pixelClock voltage

1

Electrode

Imaging area

Potential

Frame−store region
Time

Horizontal shift register

ReadoutnodeElectrons

Figure 3.10:[121].

3.3.2

Three phase transfer Figure 3.11:[121].

2

Frame transfer CCD

Overview of the XIS

The XIS is an X-ray CCD camera system which can take image, spectral and photometrical datain the 0.2–12 keV band. The XIS has four sensors (XIS 0, 1, 2, and 3). Each sensor consists of
one CCD chip and a camera body. A photograph and the cross section of one of the XIS sensors(CCD + camera body) are shown in ﬁgure 3.12 and 3.13.

3.3. X-RAY IMAGING SPECTROMETER (XIS)

25

Electric
valve

Hood
105 mmBaffleWall calibration
sourceOptical blocking
filter (OBF)
Cover shield
Thermo-electric
CCD
cooler (TEC)
Heat sink

Base
70 mm

Door calibration
sourceBonnet
125 mmDoor

Flexible print

Figure 3.13: Cross section of the XIS
Figure 3.12: Picture of one of the XIS
sensor [50].
sensors [50].

The CCD chip of the XIS is a MOS-type three-phase FT CCD. A photograph and theschematic view of the CCD is shown in ﬁgure 3.14 and 3.15. The CCD chips of the XIS arefabricated by the Lincoln Laboratory of MIT and basically same as those of ACIS/Chandra [12]except for a charge injection structure. The details of the charge injection structure are foundin the following section 3.3.3.
The upper half of the CCD shown in ﬁgure 3.15 is an imaging area, and the lower half is aframe-store region. Both regions have 1024 × 1024 pixels. The signal charges are transferreddown in ﬁgure 3.15. The pixel sizes of the imaging area and the frame-store region are 24 µm ×
24 µm and 21 µm × 13.5 µm respectively. In the imaging area, one pixel and the entire ﬁeldof view correspond to 1′′ .04 × 1′′ .04 and 17′ .8 × 17′ .8, respectively, combined with the XRT
whose focal length is 4.75 m. Note that the spatial resolution of the XIS for celestial objects isdetermined by not the pixel size of the XIS but the imaging capability of the XRT.
Each CCD chip are divided into four segments (from A to D) and each segment has one
readout node. One of the CCD chips (XIS 1) is a back-illuminated type (BI) and the others
are a front-illuminated type (FI). The front side of a CCD has a structure of electrode andinsulator (“gate” structure) made of poly-Si and SiO2 layers. The thicknesses of the poly-Si and
SiO2 layers are ∼ 0.28 µm and ∼ 0.44 µm in the case of the XIS. The FI CCD is less sensitive
than the BI CCD to soft X-rays since the gate structure absorbs soft incident X-rays. On theother hand, the depletion layer of the FI CCD is thicker (∼ 65 µm) than that of the BI CCD(∼ 42 µm) and thus the FI CCD is more sensitive to hard X-rays than the BI CCD.
For the in-orbit calibration, calibration sources of 55 Fe are equipped (“Wall calibration
source” in ﬁgure 3.13). They emit the Mn I Kα (5.895 keV) and Kβ (6.492 keV) [8, 56] lines tothe upper edge of the segment A and D constantly (see ﬁgure 3.15).
The CCDs of the XIS are operated at a temperature of −90◦ C cooled by radiators and26

CHAPTER 3. INSTRUMENT

Figure 3.14: Picture of the CCD installed in the base [50].

Figure 3.15: Schematic view of the
CCD [50].

Pertier coolers. Each XIS has an optical blocking ﬁlter (OBF) to block optical and UV photons.
The OBF is made of 100 nm polyimide coated with Al of a 120 nm thickness.
XIS 2 suddenly showed an anomaly on November 9, 2006, and it has not been operated sincethen. Although there is no direct evidence, the micro-meteoroid impact might have caused theanomaly 1 .

3.3.3

Charge Injection Capability

A charge injection (CI) structure is installed adjacent to the topmost row of the imaging area.
The CI structure allows us to inject a commandable amount of charge in a nearly arbitrary spatial
pattern. The XIS is equipped with the CI structure for the ﬁrst time among the satellite-onboard
CCD cameras. In the case of the XIS, the CI has been used in two ways to mitigate the eﬀectof the radiation damage in orbit; checker-ﬂag CI (CFCI) and spaced-row CI (SCI) techniques.
The CI of the XIS is important to analyze the data used in this thesis because the diﬀerence ofthe observation dates are up to three and half years. Generally, systematic errors from a timedegradation of a CCD become serious problems for such data sets. In the case of the XIS data
used in this thesis, however, such systematic errors are small due to these CI techniques.
The essential function of the CI is described in ﬁgure 3.16 and 3.17. A serial CI registerof 1024 pixels long is attached next to the upper edge of the imaging area. An input gate isequipped left of the CI register (ﬁgure 3.16). Pulling down the potential for electrons at the
input gate and the next electrode (S3 in ﬁgure 3.17), the potential well is ﬁlled with charges
with an amount of Q. Then, pulling up the potential, the charges are spilled. The amount ofcharges is controlled by the oﬀset voltage between the input gate and the next electrode (S3).
It can inject charges from ∼ 50 e− to ∼ 4000 e− per pixel; the equivalent X-ray energies range
from ∼300 eV to ∼15 keV. The deposited charges are transferred horizontally through the serial1

See http://www.astro.isas.jaxa.jp/suzaku/proposal/ao3/suzaku td/ and similar micro-meteoroid impact
events of XMM-Newton are reported by Strüder et al. [120].

3.3. X-RAY IMAGING SPECTROMETER (XIS)27

CI resister, vertically through transferred into the imaging area, and to an arbitrary position of
the CCD chip. Prigozhin et al. [97] have reported details concerning the CI structure.

Figure 3.16: Schematic view of the CI structure [83].
Figure 3.17: Schematic
view for the CI at the
input gate [83].

X-ray CCDs in orbit suﬀer from radiation damage. The damage causes the increase of thecharge transfer ineﬃciency (CTI) which is deﬁned as the ratio of lost charges to transferredcharges in one pixel by one transfer. The main origin of the CTI thought to be the increaseof defects in the crystal lattice generated by cosmic-rays in orbit. These defects form charge
traps that can capture signal electrons. The CTI results in the degradation of the determination
accuracy of the energy and the energy resolution for two reasons: (1) the pulse height strongly
depends on the position of an X-ray event, since the X-ray event loses more electric charges as
the number of transfer increases, and (2) the loss of charge is a stochastic process, and thus
there is a ﬂuctuation in the amount of lost charge.
Checker-Flag Charge Injection
One strategy to mitigate the eﬀect of the CTI is to measure the CTI of each column precisely andcorrect the lost charges. Here the ”column” is deﬁned as the vertical column of the pixels along
the ACTY in ﬁgure 3.15. The checker-ﬂag CI (CFCI) technique makes the column-to-column
CTI measurement possible. Here the outline of the CFCI technique is described referring ﬁgure
3.18. First, a “test” charge packet is injected into the top CCD row. Then, after a gap of a few
rows, ﬁve continuous packets are injected with the same amount of charge of the test packet.
The former four packets are called “sacriﬁcial” charge packets, while the last one is called a“reference” charge packet. The test packet loses its charge by charge traps. On the other hand,
the reference packet does not suﬀer from any charge loss, because the traps are already ﬁlled by
the preceding sacriﬁcial charges. Thus, we can measure the CTI of each column by comparing
the pulse heights of the reference charge and the test charge.
Using the obtained CTI values, the column-to-column CTI correction is applied for the data

28

CHAPTER 3. INSTRUMENT

Figure 3.18: Schematic view of how to measure CTI with the CFCI technique [83].
observed between the launch and September 2006. The energy resolution of the FI CCD at
5.9 keV was ∼140 eV (FWHM) at August 2005, and had degraded to ∼200 eV at August 2006
without the column-to-column CTI correction. The column-to-column CTI correction improves
the energy resolution to ∼170 eV (FWHM) at August 2006. Nakajima et al. [83] and Ozawa etal. [89] reported the details of the CFCI in orbit and the column-to-column correction.
However, we cannot correct the ﬂuctuation by the readout-to-readout in a speciﬁc column,
which still degrades energy resolution even after the column-to-column CTI correction in principle.

Spaced-row Charge Injection
Spaced-row charge injection (SCI) technique can reduce the CTI actively and improve the energy
resolution. In the SCI mode, a charge packet is injected into CCD rows periodically during theobservation. The injected charge ﬁlls the radiation-induced traps as a sacriﬁcial charge packet,
and thus prevents some of the traps from capturing signal charges produced by X-rays. Results
based on ground experiments using the SCI technique with radiation-damaged CCDs have beenreported by Tomida et al. [128] and Bautz et al. [5], but no in-orbit experiment had been
done. The Suzaku XIS operated with the SCI technique in orbit for the ﬁrst time from August
2006. In the SCI mode of the XIS, a charge is injected into every 54th row. The amounts of
the injected charge into each pixel are equivalent to the X-ray energy of ∼6 keV and of ∼2 keV,for the FI and BI, respectively. The energy resolution was improved from ∼200 eV to ∼140 eV
at 5.9 keV [6]. The SCI has been applied in a normal observation since October 2006 (theSCI mode). In the SCI mode, since the CTI depends on the distance from a charge injectedrow, the pulse heights show a sawtooth distribution depending on the transfer number shown inﬁgure 3.19. This non-uniformity of the pulse heights degrade the determination accuracy of theenergy and the energy resolution. Thus we developed a new CTI correction method for the SCI
mode, which is called a “sawtooth” CTI correction. This method corrects the non-uniformity
like shown in ﬁgure 3.19. All the XIS data of the SCI mode are processed with the sawtoothcorrection. The details of the sawtooth correction and CTI measurement method for the SCI
mode are described in appendix B and [134].

3.3. X-RAY IMAGING SPECTROMETER (XIS)

29

1640

Pulse height (ch)

1630
1620
1610
1600

Readout
Corrected

1590
800

830

860
890
920
i (transfer number)

950

980

Figure 3.19: Relation between pulse height of the MnI Kα line from the onboard calibration
source and transfer number i: This is the data of February 2008. Black and red marks are data
before and after our new CTI correction, respectively [134].

3.3.4

Performance of the XISDetermination Accuracy of the Energy Scale
Figure 3.20 shows the time histories of the center energy of the Mn I Kα line obtained withthe onboard calibration sources. These data are processed with the column-to-column CTI
correction for the SCI-oﬀ data or the saw-tooth CTI correction for the SCI-on data. The center
energy of the Mn I Kα line had been constantly 5.895 keV until January 2009 but that of the
BI decreased systematically by 20 eV. The reason is not clear now. The uncertainty of theenergy-scale determination is less than ∼ 20 eV (0.3%) at 5.9 keV except for these BI data.
$SU

7\SLFDOHUURU

MnI Kα peak [eV]6HS

Time since 2005-7-10 [day]Figure 3.20: Time histories of the center energy of the Mn I Kα line obtained with the onboard
calibration sources. Left: SCI-oﬀ mode [89]. Right: SCI-on mode [145].

Figure 3.21 shows the time histories of the center energy of the O VII Kα line (0.570 keV)30

CHAPTER 3. INSTRUMENT

obtained from 1E 0102-7219, which is a well-studied supernovae remnant (SNR) and used asa celestial calibration source for X-ray astronomical satellites recently [94]. The uncertainty ofenergy-scale determination of the XIS is less than ∼ 10 eV (1 – 2%) at 0.5 – 1 keV.

Corrected pulse height (ch)OVIII Kα peak [eV]

170

165

~ 4 eV

160
400

500

600
700
800
Time since 2005-7-10 (day)

900

1000

Time since 2005-7-10 [day]Figure 3.21: Time histories of the center energy of the O VII Kα of 1E 0102.2-7219. Left:
SCI-oﬀ mode [89]. Right: SCI-on mode, where the center energy is described in the unit ofpulse height channel and the result of XIS 1 is shown [134].

There is no good celestial calibration source for the 1.5 – 3 keV band, the uncertainty islargest in this band. Recent studies of SNRs using the XIS show the energy-scale determinationin the 1.5 – 3 keV can have the systematic error of 20–30 eV (1 – 2%).
Responce and Energy Resolution
The pulse height distribution of the signals for a monochromatic incident X-ray energy is calleda response function of the detector. The schematic response function of the XIS is shown inﬁgure 3.22. The main-peak component is made by competely-gathered signal charges. The
sub-peak, triangle and constant components come from events in which some of signal charges
are lost. The Si escape and Si Kα lines happen in the cases that a ﬂuorescent Si X-ray photon
generated by a photoelectric absorption escapes or is absorbed in other pixels. The detail ofphysical interpretations of the response function are reported by Matsumoto et al. [64].
The energy resolution of the XIS is deﬁned as the width of the main peak. The time historyof the widths of Mn I Kα from the onboard calibration sources is shown in ﬁgure 3.23. The
width of the Mn I Kα is almost 0 eV2 , and thus ﬁgure 3.23 shows the time histories of the energy
resolution at 5.9 keV. The energy resolution of the BI has degraded more rapidly than that of
the FI in the SCI mode.
The time history of the energy resolution at 0.57 keV (O VII Kα) is shown in ﬁgure 3.24.
This ﬁgure is obtained with the observations of 1E 0102.2-7219, whose emitting narrow lines
were resolved with the grating spectrometer of XMM-Newton (see [98] and [94] in detail).
The calibration data base for the XIS includes these time history of the energy resolution
properly.
2

Strictly speaking, since this line is a blend of Mn K I α1 5.888 keV and α2 5.899 keV, the width of the merged
Mn I Kα is a few eV. See figure 42 of [43] for examples

3.3. X-RAY IMAGING SPECTROMETER (XIS)

31

Figure 3.22: Schematic response function of the XIS. 1. main peak, 2. sub peak, 3. triangle
component, 4. Si escape, 5. Si Kα line, and 6. constant component. The physical interpretationof the each component is shown in[64].

$SU
7\SLFDOHUURU

(a)
180

Mn I Kα FWHM (eV)MnI Kα FWHM [eV]6HS

170

160

XIS 1(BI)
Oct.
2006

Mar.
2008

150
XIS 0(FI)
140
400

500

600
700
800
Time since 2005-7-10 (day)

900

1000

Time since 2005-7-10 [day]Figure 3.23: Time histories of the width of the Mn I Kα line obtained with the onboard calibration sources. Left: SCI-oﬀ mode [89]. Right: SCI-on [134].

32

CHAPTER 3. INSTRUMENTO VIII Kα FWHM (eV)70

XIS 1(BI)60

XIS 0(FI)
50

400

500

600
700
800
Time since 2005-7-10 (day)

900

1000

Figure 3.24: Time history of the line widths of the O VII Kα obtained with the onboard
calibration sources in the SCI-on mode [134].

Quantum Eﬃciency and Eﬀective Area
Conversion eﬃciency of an incident X-ray photon to an detected X-ray event is called quantum
eﬃciency (QE). QE of the XIS is shown in ﬁgure 3.25. The diﬀerence of the QE between the FI
and the BI comes from the illumination type and the thickness of the depletion layers.

Figure 3.25: Quantum eﬃciency of the XIS as a fuction of incident energy [50].
The QE in the low energy band has degraded by unexpected contamination (due to outgassing from the satellite) accumulating on the OBF [50, 73]. The contaminant is mainlycomposed of carbon and oxygen. Although the chemical composition of the contaminant is still
uncertain, it is assumed to be C/O=6 in number. The time evolution of the contamination
has been measured using regular observations of 1E 0102-072. The eﬀect of contamination is3.3. X-RAY IMAGING SPECTROMETER (XIS)33

included in the eﬀective area of the XRT properly. Note that, in this thesis, Fe Kα lines (∼ 7keV) are main targets and the contamination eﬀects almost nothing to the ∼ 7-keV band.
When the XIS data are analyzed, the observed pulse height distribution is ﬁtted with models
convolved with this response function including the quantum eﬃciency of the XIS and theeﬀective area of the XRT.
Non-X-ray BackgroundCosmic-rays in orbit causes non-X-ray background (NXB). Figure 3.26 shows the spectra of
NXB accumulated during night-earth observations in which no celestial X-ray is incident.

Figure 3.26: Spectra of the NXB of the FI (XIS 0) and the BI (XIS 1) [125].
The ﬂux of the NXB depends on the geomagnetic cut of rigidity (COR). A COR is theminimum energy of the cosmic-ray which can penetrate the magnetic ﬁeld of the Earth and astandard value to evaluate the strength of the cosmic-ray of orbits. The ﬂux of the NXB also
depends on the position on the CCD chip. It increases along ActY (ﬁgure 3.27) because pixelsof the larger ActY take the longer time before the readout in the frame-store region and suﬀer
the more cosmic-ray radiation [140]. Tawa et al. [125] modeled the NXB spectra as a functionof the COR and the position on the CCD. The model reconstruct the NXB spectra with theaccuracy of 5% for any observation. In this thesis, the reconstructed spectra are used.
Due to the near-earth orbit, the NXB of the XIS is lower than X-ray CCD cameras on othersatellites, especially in the 7–10 keV band (ﬁgure 3.28). The XIS is the best instrument to detectdiﬀuse faint X-ray emissions like the Galactic diﬀuse X-ray emission studied in this thesis.

34

CHAPTER 3. INSTRUMENT

Figure 3.27: Dependence of the NXB on ACTY in the case of XIS 2 [140].

Figure 3.28: XIS background counting rate as a function of energy. The background rate wasnormalized with the eﬀective area and the ﬁeld of view, which is a good measure of the sensitivity
determined by the background for spatially extended sources. The background spectra of ASCA,
Chandra, and XMM-Newton are shown for comparisons [69].

Chapter 4

Galactic Diﬀuse X-ray Emission
In this chapter, we show the result of the Suzaku observation of the region from the GC to
the ridge and the bulge (−3◦ < l < 2◦ , −2◦ < b < 1◦ ). In section 4.1, we describe about
the observations and the data reduction. In section 4.2, we give the spatial distribution of the
GDXE in the 5–10 keV band, especially the Fe Kα line emissions. In section 4.3, we show thedetailed spectral analysis of the GDXE. Discussions based on the results given here will be madein chapter 7.

4.1

Observation & Data Reduction

4.1.1

Observation

The Suzaku data of the GC region used in this thesis were obtained from October 2005 to March
2009. The data in b = −2◦ were also analyzed. The total eﬀective exposure is about 3.2 Ms.
The logs of the observations are summarized in table 4.1.
In these observation, the XIS operated in the normal clocking full window mode 1 . In thismode, the time resolution of the XIS is 8.0 s. 71 % of the total exposure was observed without
XIS2. The SCI technique was used in 75 % of the total exposure. We did not use the BI data
after January 2009 because the energy-scale uncertainty is large as described in section 3.3.4.

4.1.2

Data Reduction

The observation data were downloaded from DARTS/ISAS2 . In order to apply the latest CTI
correction, we reprocessed the data using the calibration data base (CALDB) released at 200810-203 and xispi in the HEADAS software package version 6.6.14 . This CALDB includes thetime history of the CTI. Hereafter, Typewiter means the name of the software in the HEADAS
package.
After the reprocessing, we cut the epoch of Earth elevation angles less than 5◦ (to avoid
targets hiding the edge of the Earth), the day Earth elevation angle less than 20◦ (to avoid
1

In the normal clocking full window mode, the FT CCDs are operated in the way described in section 3.3.1.
The details of the clocking and window modes are shown in Koyama et al. [50].
2<http://www.darts.isas.jaxa.jp/astro/suzaku/>.
3<http://www.astro.isas.jaxa.jp/suzaku/caldb/history/>.
4
See <http://heasarc.gsfc.nasa.gov/ FTP/software/ftools/release/archive/Release Notes 6.6.1>.
The
HEADAS package is released by Goddard Space Flight Center/NASA and generally used for the data analysis of high energy astronomy. Tools released by the Suzaku team are also included in this package.

35

36

OBSID

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

Pointing Direction
α2000 (◦ ) δ2000 (◦ ) l(◦ )b(◦ )100027010
100027020
100037010
100037030
100037040
100037060
500018010
500005010
100048010
501040010
501040020
501008010
501009010
501010010
501049010
501050010
501051010
501052010
501053010
501057010
501039010
501046010
501047010
501048010
501054010
501055010
501056010
501058010
501059010
501060010
502022010
102013010
502020010
502002010
502003010
502004010
502005010
502006010
502007010
502008010
502009010
502010010
502011010
502016010
502017010
502018010
502051010
503007010
503008010
503009010
503010010
503011010
503012010
503013010
503014010
503015010
503016010
503017010
503018010
503072010
403009010

266.51
266.30266.30266.20266.51266.88266.16266.77266.51
266.69266.69266.50266.19266.26265.38265.91265.70265.50265.30266.03267.02265.98265.78265.58266.63266.43266.23266.98267.09267.29266.81266.51
267.19267.16266.96267.48267.29266.18265.86265.66267.48266.38266.57266.23266.47266.06266.94266.44266.78266.45266.04265.95266.30265.67265.18265.03264.87264.72264.56265.99266.48-28.92-29.17-29.17-29.35-28.93-28.44-29.47-28.63-28.93-28.38-28.38-29.17-28.91-30.37-29.75-29.65-29.93-30.21-30.50-30.11-28.35-28.90-29.19-29.47-29.25-29.54-29.82-27.72-27.94-27.65-28.88-28.93-28.13-29.14-29.42-29.31-29.60-28.62-28.45-28.73-27.37-28.34-28.05-30.11-30.09-30.24-28.15-28.57-29.13-29.34-29.55-29.83-29.94-30.07-30.75-30.96-31.17-31.38-31.60-29.21-28.78

0.05
-0.25
-0.25
-0.45
0.06
0.64
-0.57
0.43
0.06
0.61
0.61
-0.15
-0.07
-1.29
-1.17
-0.83
-1.17
-1.50
-1.83
-1.17
0.78
-0.17
-0.50
-0.83
-0.17
-0.50
-0.83
1.30
1.17
1.50
0.23
0.06
1.05
0.17
-0.17
0.17
-0.17
0.17
0.17
-0.17
1.83
0.50
0.83
-1.08
-0.95
-1.27
0.91
0.33
0.00
-0.33
-0.70
-0.97
-0.91
-1.30
-2.10
-2.35
-2.60
-2.85
-3.10
-0.42
0.17

-0.07
-0.05
-0.05
-0.07
-0.07
-0.10
-0.09
-0.12
-0.07
0.07
0.07
-0.19
0.18
-0.64
0.33
0.00
-0.00
0.00
-0.00
-0.33
-0.16
0.33
0.33
0.33
-0.33
-0.33
-0.33
0.20
0.00
0.00
-0.27
-0.07
-0.17
-0.67
-0.67
-1.00
-1.00
0.33
0.67
0.67
-0.00
0.33
0.33
-0.48
-0.65
-0.43
0.01
0.17
-0.38
-0.24
-0.05
-0.13
-0.44
-0.05
-0.05
-0.05
-0.05
-0.05
-0.05
0.17
0.02

502059010

268.38-29.96

-0.00

-2.00

Table 4.1: Observation data list
Observation
Observation
Start (UT)
End (UT)
Galactic center 5◦ × 2◦ region
2005-09-23T07:18:25 2005-09-24T11:05:19
2005-09-24T14:17:17 2005-09-25T17:27:19
2005-09-29T04:35:41 2005-09-30T04:29:19
2005-09-30T06:06:32 2005-09-30T07:41:20
2005-09-30T07:43:01 2005-10-01T06:21:24
2005-10-10T12:28:01 2005-10-12T07:05:23
2006-02-20T12:45:25 2006-02-23T10:50:14
2006-03-27T23:00:22 2006-03-29T18:12:15
2006-09-08T02:23:24 2006-09-09T09:06:15
2006-09-21T17:29:01 2006-09-23T06:52:20
2006-09-24T05:03:12 2006-09-25T07:00:24
2006-09-26T14:18:16 2006-09-29T21:25:14
2006-09-29T21:26:07 2006-10-01T06:55:19
2006-10-07T02:16:52 2006-10-08T10:19:19
2006-10-08T10:22:40 2006-10-09T02:19:24
2006-10-09T02:20:25 2006-10-09T13:39:24
2006-10-09T13:40:09 2006-10-10T06:44:24
2006-10-10T06:45:09 2006-10-10T21:18:14
2006-10-10T21:18:59 2006-10-11T10:06:14
2006-10-11T10:07:27 2006-10-12T03:28:14
2007-03-03T12:20:20 2007-03-05T12:39:25
2007-03-10T15:03:10 2007-03-11T03:55:14
2007-03-11T03:55:59 2007-03-11T19:04:14
2007-03-11T19:04:59 2007-03-12T08:09:14
2007-03-12T08:11:07 2007-03-12T23:58:24
2007-03-12T23:59:09 2007-03-13T15:40:19
2007-03-13T15:41:12 2007-03-14T05:00:24
2007-03-14T05:02:29 2007-03-15T18:55:14
2007-03-15T18:55:51 2007-03-17T05:06:19
2007-03-17T05:07:04 2007-03-18T20:58:14
2007-08-31T12:33:33 2007-09-03T19:00:25
2007-09-03T19:01:10 2007-09-05T05:20:20
2007-09-06T00:26:47 2007-09-09T16:10:19
2007-10-09T16:40:54 2007-10-10T03:40:24
2007-10-10T03:41:13 2007-10-10T15:20:24
2007-10-10T15:21:17 2007-10-11T01:00:24
2007-10-11T01:01:17 2007-10-11T11:32:20
2007-10-11T11:34:01 2007-10-11T23:07:14
2007-10-11T23:09:15 2007-10-12T09:52:14
2007-10-12T09:52:59 2007-10-12T21:50:19
2007-10-12T21:52:24 2007-10-13T07:30:19
2007-10-13T07:32:00 2007-10-13T18:50:24
2007-10-13T18:51:09 2007-10-14T05:30:24
2008-03-02T18:08:00 2008-03-04T17:40:19
2008-03-06T13:26:36 2008-03-08T16:00:24
2008-03-08T16:02:17 2008-03-10T21:00:19
2008-03-11T06:19:45 2008-03-15T05:30:18
2008-09-02T10:15:27 2008-09-03T22:52:24
2008-09-03T22:53:29 2008-09-05T06:56:19
2008-09-05T06:57:08 2008-09-06T15:55:24
2008-09-06T15:56:13 2008-09-08T01:39:24
2008-09-08T09:08:09 2008-09-09T21:33:19
2008-09-14T19:35:07 2008-09-16T00:50:14
2008-09-16T00:51:19 2008-09-18T04:44:24
2008-09-18T04:46:49 2008-09-19T07:32:20
2008-09-19T07:33:05 2008-09-20T09:56:13
2008-09-22T06:47:49 2008-09-23T08:07:17
2008-09-23T08:08:10 2008-09-24T09:21:13
2008-09-24T09:27:54 2008-09-24T22:30:24
2009-03-06T02:39:11 2009-03-09T02:55:24
2009-03-21T02:03:27 2009-03-23T18:26:09
b = −2◦
2007-09-29T01:40:51 2007-10-02T14:10:16

Exp.
(ks)

XIS2

SCI

Object
Name

44.7
37.4
43.5
3.0
42.9
76.4
106.9
88.4
63.0
61.4
44.8
129.6
51.2
50.7
19.6
22.0
21.9
19.3
21.9
20.5
96.4
25.2
25.6
27.5
26.1
27.2
26.5
63.3
62.2
64.8
134.8
51.4
139.1
23.2
21.5
19.9
20.6
22.6
22.0
23.8
20.9
21.6
23.0
70.5
72.6
79.0
138.8
52.2
53.7
52.4
53.1
57.6
57.7
104.8
55.4
56.8
52.2
51.3
29.4
140.6
110.8

ok
ok
ok
ok
ok
ok
ok
ok
ok
ok
ok
ok
ok
ok
ok
ok
ok
ok
ok
ok

oﬀ
oﬀ
oﬀ
oﬀ
oﬀ
oﬀ
oﬀ
oﬀ
oﬀ
on
on
on
on
oﬀ
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on

Sgr A East
Sgr A west
Sgr A west
KS1741-293(GCBGD2)
Sgr A East
Sgr B2
SGR C
GC CENTER #2
GALACTIC CENTER
GC SGR B NORTH
GC SGR B NORTHGC SOUTHGC SOUTH BGDHESS J1745-303
GALACTIC CENTERGALACTIC CENTERGALACTIC CENTERGALACTIC CENTERGALACTIC CENTERGALACTIC CENTER
GC SGR B EASTGALACTIC CENTERGALACTIC CENTER GC2GALACTIC CENTER GC3
GALACTIC CENTER GC9GALACTIC CENTER GC10GALACTIC CENTER GC11
GC SGR D NORTH
GC SGR D
GC SGR D EAST
(L,B)=(0.25,-0.27)
GALACTIC CENTER
SGR D SNR
GC14
GC15
GC16
GC17
GC18
GC19
GC20
GC21
GC22
GC23
HESS J1745-303 1
HESS J1745-303 2
HESS J1745-303 3
GC G0.9+0.1
GC LARGEPROJECT1GC LARGEPROJECT2GC LARGEPROJECT3GC LARGEPROJECT4GC LARGEPROJECT5GC LARGEPROJECT6GC LARGEPROJECT7GC LARGEPROJECT8GC LARGEPROJECT9GC LARGEPROJECT10GC LARGEPROJECT11GC LARGEPROJECT12EXTENDED CHIMNEY
ARCHES CLUSTERonGALACTIC BULGE

136.8

4.1. OBSERVATION & DATA REDUCTION

37

reﬂected X-rays from the Earth) and the South Atlantic Anomaly (to minimum the NXB). The
eﬀective exposures after this screening of each data were shown in table 4.1.
We conﬁrmed that the center energy of the Mn I Kα line of the onboard calibration sources(55 Fe) is 5.895 eV within the calibration uncertainty of 20 eV during the observation period.

4.1.3

Response Function of the XIS

Using xissimarfgen [34] and xisrmfgen, the eﬀective area of the XRT and the response of the
XIS for spectra were calculated, and ancillary response ﬁles (ARFs) and redistribution matrix
ﬁles (RMFs) were made. ARFs and RMFs are telescope and detector response ﬁles used ina spectral analysis tool, xspec11. The ARFs and RMFs depend on the observation date, thetarget positions on the CCD and in the sky, and the attitudes of the Suzaku satellite during theobservation.
Since the responses of the FIs are almost same, we merged the FI spectra into one for thefollowing spectral analysis. The RMFs and ARFs were merged for the FIs with addrmf andaddarf.

4.1.4

BackgroundThe background of the GDXE observation composes of mainly the NXB and the cosmic X-ray
background (CXB). The detail of NXB was summarized in section 3.3.4. The NXB spectra sortedby the COR were made from the same source region in the detector coordinate using xisnxbgen[125]. The NXB spectra were averaged for the FIs. These NXB spectra were subtracted from theobserved GDXE spectra in the following spectral ﬁttings. The CXB is celestial X-rays comingfrom all over the sky almost uniformly. The CXB is due to unresolved extra-galactic X-ray
sources, the bulk of which are active galactic nuclei (AGN) [24, 70]. According to Kushino etal. [58], the averaged CXB is described with a power-low model with a photon index Γ = 1.4
and the intensity is 5.4 × 10−15 erg s−1 cm−2 arcmin−2 in the 2–10 keV band. This CXB model
was used in the following spectral ﬁttings.
In ﬁgure 4.1, we compare the observed spectra of the GDXE without the background subtraction, with the spectra of the reconstructed NXB and the GDXE spectra subtracted the NXB
and the modeled CXB spectra. Here, the interstellar absorption of the CXB is assumed to be
NH = 3 × 1022 cm−2 . The spectra of region 5 and b = −2◦ (section 4.3) are chosen as examplesof dark GDXE regions. We did not use the BI data above ∼ 8 keV in the following analyses
because the signal to noise ratio is low as shown in ﬁgure 4.1. Note that even in the dark GDXE
regions, the intensity ratios of the CXB to the GDXE are 10 % and 20 % at most for the GC
and b = −2◦ regions, respectively.

4.1.5

Exclusion of Bright Sources

Bright sources listed in table 4.2 were excluded to make the GDXE spectra. In the case of thedata at b = −2◦ , seven point sources were found in the FOV and summarized in table 4.3. The
circles with a radius of 1′ .5 centered at the point sources were excluded.

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

5

10

0.01

(b)

10−3
10−4

0.01

Normalized counts/s/keV

0.1

(a)

10−3
10−4

Normalized counts/s/keV

0.1

38

5

10
Energy (keV)

0.01

(d)

10−3
10−4

0.01

Normalized counts/s/keV

0.1

(c)

5

10
Energy (keV)

10−3
10−4

Normalized counts/s/keV

0.1

Energy (keV)

5

10Energy (keV)

Figure 4.1: Comparison of the spectra of the GDXE, NXB and CXB. Black: observed spectra of
the GDXE without the background subtraction. Red: the reconstructed NXB. Green: GDXE
spectra subtracted the NXB. Blue: modeled CXB spectra. (a) (b) Spectra at region 5 of theFIs and the BI. (c) (d) Spectra at b = −2◦ of the FIs and the BI, respetively.

4.1. OBSERVATION & DATA REDUCTION

39

Table 4.2: List of the bright sources excluded to make the GDXE spectra.
Position∗

Radius∗lb(arcmin)Great Annihilator−0◦ .887 −0◦ .097
4◦◦
CXOGCS J174445.5–295042 −0 .774 −0 .062
2◦◦
2E 1742.9–2929−0 .439 −0 .383
5◦◦
KS 1741–293−0 .432 −0 .074
3
Sgr A East−0◦ .047 −0◦ .053
3
Arches cluster0◦ .1230◦ .024
2◦◦
1E 1743.1–2843
0 .256 −0 .029
3◦◦
G0.61+0.01
0 .6200 .0164CXOGC J174645.3–281546
0◦ .7120◦ .141
2
HD 161507
0◦ .8370◦ .0852◦◦
G0.9+0.1
0 .866
0 .077
2
SAX J1748.2-2808
0◦ .994 −0◦ .086
2
AX J1749.1–2733
1◦ .592 −0◦ .062
2
∗ The position and the radius show those of the excluded circles.
Name

Table 4.3: List of the seven point sources excluded from the b = −2◦ data.
Identiﬁcation∗

Positionlb−0◦ .079−1◦ .945

None
2XMM J175339.3–300334−0◦ .020 −2◦ .125 2XMM J175358.1–300234
0◦ .002 −2◦ .053 2XMM J175344.7–295905
0◦ .004 −1◦ .985 2XMM J175327.8–295719
0◦ .046 −2◦ .007 2XMM J175341.2–295521
0◦ .089 −2◦ .077 2XMM J175402.7–295524∗ We searched corresponding objects in the 2XMM catalog [139].
−0◦ .067−2◦ .073

40

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

4.2

Spatial Distribution of the GDXE

4.2.1

Intensity Proﬁles along the Galactic Longitude

To study the intensity proﬁle of the Fe Kα line along the galactic longitude, X-ray spectra
were made from the rectangles show in ﬁgure 4.2. The Galactic latitude coordinate of the eachrectangle is same as that of Sgr A∗ (b = −0◦ .046) [100]. The size of the each rectangle is 6′ × 12′ .

Figure 4.2: Regions from which spectra were extracted along the galactic longitude. They areoverlaid on the X-ray image in the 6.55-6.8 keV band.
The spectra of the NXB, the ARFs and the RMFs were generated in the way described in
section 4.1.3. We ﬁtted a phenomenological model below to the spectra subtracted the NXB
spectra in the 5–10 keV band:
4 Gaussians + P L × Abs1 + CXB × Abs2.

(4.1)

P L is a power-law model. Abs is the interstellar absorption model. The line centers of thegaussians were ﬁxed to 6.40 (Fe I Kα), 6.68 (Fe XXV Kα), 6.97 (Fe XXVI Kα) and 7.06 (Fe I
Kβ) keV according to APEC model [119] for the ionized lines and Kaastra & Mewe [40] for theneutral lines. Width of the each line was ﬁxed to 0 eV. The intensities of the Fe I Kα, Fe XXV
Kα, and Fe XXVI Kα lines in the models were free in the ﬁttings. The intensity of the Fe I
Kβ line was ﬁxed to 0.125 times that of the Fe I Kα line according to [40]. The normalizationand photon index Γ of the power-law model were free. CXB is a power-law model as the CXB
and the normalization and Γ were ﬁxed as described in section 4.1.4. Abs = exp(−σ(E) × NH )where E is the energy of the incident photon and σ(E) is the cross section of the interstellar
absorption. We calculated σ(E) referring to [1] and [4]. NH is the column density of hydrogenin the interstellar matter. The column density of Abs1 was free. The column density of Abs2
was the interstellar absorption of the CXB and ﬁxed to be NH = 1.2 × 1023 cm−2 , which is
twice the typical interstellar absorption toward the GC region, NH = 6 × 1022 cm−2 [96, 110].
The FIs and BI spectra of each region were ﬁtted simultaneously. The spectra of the same
region but of diﬀerent observations were ﬁtted simultaneously using diﬀerent ARFs and RMFs.
Some examples of the spectra and the best-ﬁt models are shown in ﬁgure 4.3.

4.2. SPATIAL DISTRIBUTION OF THE GDXE
lm1.156 FI

χ

0.1
0.01
10−3
2 10−4
0

10−3

0.01

0.1

Normalized counts s−1 keV−1

501051010−2

χ−4 −2 0 2 10−4

Normalized counts s−1 keV−1

l0.244 FI
B:100027010 R:100048010 G:102013010 B:102013010

41

Energy (keV)

Energy (keV)

Figure 4.3: Examples of the spectra and the best-ﬁt models. The data of the diﬀerent observations are shown in the diﬀerent colors. Only the FIs spectra are shown for the simplicity
although we simultaneously ﬁtted the FIs and BI spectra. The vertical error bars show the 1σ
errors. Left : l = 0◦ .244. Right: l = −1◦ .156.

Fe XXV Kα and Fe XXVI Kα
The intensity proﬁles of the Fe XXV Kα and Fe XXVI Kα lines were shown in ﬁgure 4.4 in which
the results of |l| > 8◦ by Yamauchi et al. [143] are also added. Note that the galactic latitudesof the data of [143] were −0◦ .05 < b < 0◦ .7 and diﬀerent from those of ours (b = −0◦ .046).
The data of [143] were corrected assuming that the scale height of the intensity proﬁle along thegalactic latitude is 0◦ .5 according to Kaneda et al. [41].
We ﬁtted phenomenological models to the proﬁles to obtain a typical e-folding scale along
the galactic longitude hl (degree). We ﬁrstly made a ﬁtting with a one-component model; Al ×
exp(−l∗ /hl ). The symbol l∗ denotes the distance from Sgr A∗ (l = −0◦ .056) along the Galactic
plane measured in the unit of degree, i.e. l∗ = |l+0◦ .056|. Al (photons s−1 cm−1 arcmin−2 ) is theline intensity at l∗ = 0◦ . The data points of l∗ > 2◦ , however, excess the best-ﬁt one-component
models. Thus we tried a two-component model; Al1 × exp(−l∗ /hl1 ) + Al2 × exp(−l∗ /hl2 ) . The
ﬁttings are signiﬁcantly improved. χ2 /d.o.f. are 2805.6/59 (Fe XXV Kα) and 255.1/59 (FeXXVI Kα) with the one-component models, and 1946.9/57 (Fe XXV Kα) and 217.0/57 (FeXXVI Kα) with the two-component models. The best-ﬁt model and parameters are shown inﬁgure 4.4 and table 4.4.

Fe I Kα
The intensity proﬁle of Fe I Kα line along the galactic longitude is shown in ﬁgure 4.5. The
results of |l| > 8◦ by Yamauchi et al. [143] are also added after the correction assuming thatthe scale height of the intensity proﬁle along the galactic latitude is 0◦ .5. The intensity proﬁleof the Fe I Kα line is clumpy compared with the smooth proﬁle of the Fe XXV Kα line. The
bright clumps corresponds to radio-observed giant molecular clouds in the CMZ. It is diﬃcult
to model the eﬀect of the molecular clouds so the intensity proﬁle of the Fe I Kα line was not
ﬁtted with phenomenological models.

42

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

Observed photons cm-2 s-1 arcmin-2

Fe XXV Kα
10

-5

West
East
Yamauchi+09 corrected

10-6

10-7

10-8
0.1
1
10
*
Distance from Sgr A along the longitude l* (degree)

Observed photons cm-2 s-1 arcmin-2

Fe XXVI Kα
10

-5

West
East
Yamauchi+09 corrected

10-6

10-7

10-8
0.1
1
10
*
Distance from Sgr A along the longitude l* (degree)

Figure 4.4: Intensity proﬁles of the Fe XXV Kα (top) and Fe XXVI Kα line (bottom) along
the galactic latitude. The error bars show the 1σ errors. The black lines show the best-ﬁt
two-component models. The cyan and orange lines show the components with e-folding scales
hl1 and hl2 , respectively.
Table 4.4: Best-ﬁt parameters of the two-component model to the intensity proﬁles of the highly
ionized Fe Kα lines along the galactic longitude.∗

Fe XXV Kα
Fe XXVI Kα

Al1 (10−7 )†27.3 ± 1.4
10.4 ± 0.4

Al2 (10−7 )†1.6 ± 0.3
0.26 ± 0.08

hl1 (degree)
0.43 ± 0.03
0.48 ± 0.02

hl2 (degree)
33 ± 2033 ± 37† The uncertainties show the 1σ errors. ∗ The unit is photons s−1 cm−2 arcmin−2 .

4.2. SPATIAL DISTRIBUTION OF THE GDXE

43

Observed photons cm-2 s-1 arcmin-2

Fe I Kα
-5

10

West
East
Yamauchi+09 corrected

10-6

10-7

10-8
0.1
1
10
*
Distance from Sgr A along the longitude l* (degree)

Figure 4.5: Intensity proﬁle of the Fe I Kα line. The error bars show the 1σ errors.

Intensity in the 5–10 keV Band
The intensity proﬁle of the GDXE in the 5–10 keV band is shown in ﬁgure 4.6. The results of|l| > 8◦ by Yamauchi et al. [143] are also added after the correction assuming that the scaleheight of the intensity proﬁle along the galactic latitude is 0◦ .5. This intensity includes bothof the line and continuum emissions. The CXB and the NXB were already subtracted. Theinterstellar absorption is not corrected. We tried two-component model ﬁttings. We did not use
the data points of l = −0◦ .96 ∼ −0◦ .75 for the ﬁttings because they were clearly contaminatedby the Great Annihilator at (l, b) = (−0◦ .89, −0◦ .10). The best-ﬁt model and parameters areshown in ﬁgure 4.6 and table 4.5.

4.2.2

Intensity Proﬁles along the Galactic Latitude

To study the intensity proﬁle of the Fe K α lines along the Galactic latitude, X-ray spectra
were made from the rectangle regions show in ﬁgure 4.7. The regions were selected along the
Galactic latitude at l = −1◦ .16, −0◦ .84, −0◦ .56, −0◦ .17, 0◦ .17, 0◦ .50, 0◦ .82, 1◦ .17. The size ofeach rectangle region is 12′ × 6′ .
In the case of the data of b = −2◦ , the spectrum was obtained from a circle with a radiusof 7′ .4 at the center of the FOV. This radius was chosen to avoid the area illuminated by thecalibration source on the CCD chip.
The NXBs, ARFs and RMFs were generated in the way described in section 4.1.3. We ﬁtted
a phenomenological model shown in equation 4.1 to the spectra subtracted the NXB in the 5–10
keV band. Some examples of the spectra and the ﬁtting results are shown in ﬁgure 4.8. Thus
the intensity proﬁles of the three Fe Kα lines and the 5–10 keV band for each longitude wereobtained. The result of b = −2◦ was added to the proﬁle of l = −0◦ .17.

44

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

Observed photons cm-2 s-1 arcmin-2

5-10 keV intensity
-12

10

West
East
Yamauchi+09 corrected

10-13

10-14

10-15

0.1

1

10

*

Distance from Sgr A along the longitude l* (degree)

Figure 4.6: Intensity proﬁle in the 5–10 keV band along the galactic longitude. The error bars
show the 1σ errors. The black lines show the best-ﬁt two-component models. The cyan andorange lines show the components with e-folding scales hl1 and hl2 , respectively.
Table 4.5: Best-ﬁt parameters of the two-component model to the intensity proﬁle in the 5–10
keV band along the galactic longitude.∗Al1 (10−13 )†3.6 ± 0.2

Al2 (10−13 )†0.19 ± 0.07

hl1 (degree)
0.43 ± 0.04

hl2 (degree)
30 ± 62† The uncertainties show the 1σ errors. ∗ The unit is erg s−1 cm−2 arcmin−2 .

Figure 4.7: Regions from which spectra were made along the galactic latitude. They are overlaid
on the XIS image in the 6.55-6.8 keV band.

4.2. SPATIAL DISTRIBUTION OF THE GDXE

45

l0.17 bm0.246 FI

l1.17 bm0.246 FI

χ

0.1
0.01
10−3

Normalized counts s−1 keV−1

0.1
0.01
10−3

Energy (keV)−2 0 2 10−4χ502020010−2 0 2 10−4

Normalized counts s−1 keV−1

502022010

Energy (keV)

Figure 4.8: Examples of the spectra and the best-ﬁt models. Only the FIs spectra are shownfor the simplicity although we simultaneously ﬁtted the FIs and BI spectra. The vertical error
bars show the 1σ errors. Left : (l, b) = (0◦ .17, −0◦ .246). Right: (l, b) = (1◦ .17, −0◦ .246).

Fe XXV Kα

The intensity proﬁles of the Fe XXV Kα line for the each longitude are shown in ﬁgure 4.9.
We ﬁtted phenomenological models of Ab × exp(−b∗ /hb ) to the proﬁles to obtain typical scaleheight along the galactic latitude hb . Ab is the line intensity at b∗ = 0◦ . The symbol b∗denotes the distance from the Galactic plane (b = −0◦ .046) measured in the unit of degree, i.e.
b∗ = |b + 0◦ .046|. Note that the data point of b = −2◦ is added to the proﬁle in l = −0◦ .17. The
best-ﬁt results are shown in ﬁgure 4.9.
For the intensity proﬁles in l = ±0◦ .17, the data points in b∗ > 0.7◦ excess the best-ﬁt
models systematically. Thus a two-component models was tried for the proﬁles in l = ±0◦ .17;
Ab1 ×exp(−b∗ /hb1 )+Ab2 ×exp(−b∗ /hb2 ). The ﬁttings are signiﬁcantly improved with this model;the χ2 /d.o.f. are 273.1/19 (l = −0◦ .17) and 308.8/18 (l = 0◦ .17) with the one-component modeland 174.3/17 (l = −0◦ .17; hb1 = 0◦ .16 ± 0◦ .03, hb2 = 1◦ .1 ± 0◦ .2) and 255.8/16 (l = 0◦ .17; hb1 =0◦ .18 ± 0◦ .02, hb2 = 1◦ .0 ± 0◦ .5) with the two-component model (1σ errors).
Figure 4.10 shows the relation between the scale height hb of the one-component model andl∗ . The scale height hb increase with l∗ . It implied that not only the proﬁles in l = ±0◦ .17
but also the other proﬁles would have two scale-height components and the larger scale-heightcomponents dominate at large l∗ . Although the second component is not statically required, we
ﬁtted the two-component model to all the proﬁles in ﬁgure 4.9. The proﬁles in l∗ > 0◦ .2, however,
did not require two scale-height components. Thus we ﬁtted the every proﬁles simultaneously,
and hb1 and hb2 were common for the every longitudes but Ab1 and Ab2 were free for the eachlongitude. Figure 4.11 shows the ﬁtting results. The best-ﬁt parameters of hb1 and hb2 are0◦ .21 ± 0◦ .02 and 1◦ .1 ± 0◦ .7 (1σ errors). The χ2 /d.o.f. is 681.4/76. This is better than the total
χ2 /d.o.f. of 812.9/78 of the one-component model ﬁtting. The best-ﬁt parameters are shown intable 4.6. The relations of Ab1 , Ab2 and the galactic latitude are shown in ﬁgure 4.12. Figure
4.12 suggests that the Ab1 decreases along with the galactic longitude but the Ab2 is almost
constant.

-2

-2 -1

-0.5

0

0.5

1

10

-6

10

-7

10

-8

h=0.367481+/- 0.0542727

-2

-1.5

-1

-2 -1

10

-0.5

0

0.5

1

!

Fe XXV K! l=1.17-5-6

h=0.579654+/- 0.127115

-2 -1

10

10-7
-8

10

-5

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1!

Fe XXV K! l=0.5

10

-6

10

-7

10

-8

h=0.3389+/- 0.0472087

-2

-1.5

-1

-0.5

0

h=0.340619+/- 0.0471012

10

-7

10

10-8

-2

0.5

Galactic latitude (degree)

1

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)!

Fe XXV K! l=-0.17
-5

10

-6

h=0.260226+/- 0.0197034

10

-7

10

-8

10

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)!

Fe XXV K! l=0.82
-5

10

-6

h=0.410528+/- 0.0480402

10

10-7
-8

10

-2

-2

10

-6

-2

Galactic latitude (degree)

10!

Fe XXV K! l=-0.84
10-5

-2

!

Fe XXV K! l=-0.56
-5

-2 -1

-2 -1

Observed photons cm s arcmin-1

Galactic latitude (degree)

-2
-2 -1

Observed photons cm s arcmin-1.5

Observed photons cm s arcmin

10-8

Observed photons cm s arcmin-7

Observed photons cm s arcmin

10

h=0.51463+/- 0.0307245

Observed photons cm s arcmin

10

-6

-2

-2 -1

Observed photons cm s arcmin!

Fe XXV K! l=-1.16
10-5

-2
-2 -1

Observed photons cm s arcmin-2

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION-2

46-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1!

Fe XXV K! l=0.17
-5

10

-6

h=0.236677+/- 0.0232183

10

-7

10

-8

10

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

Figure 4.9: Intensity proﬁles of the Fe XXV Kα line along the galactic latitude. The error bars
show the 1σ errors. The solid lines show the best-ﬁt one-component models. The best-ﬁt scaleheights hb are shown with the 1σ errors.

4.2. SPATIAL DISTRIBUTION OF THE GDXE

47

Fe XXV Kα with one-component model
0.8
0.7

East
West

hb (degree)

0.6
0.5
0.4
0.3
0.2
0.1
0
0
0.2
0.4
0.6
0.8
1
1.2
1.4
*
Distsance from Sgr A along the longitude l* (degree)

Figure 4.10: Relation between l∗ and hb of the Fe XXV Kα line intensity proﬁles along thegalactic latitude with the one-component model. The error bars show the 1σ errors.
Fe XXVI Kα
The intensity proﬁles of the Fe XXVI Kα line along the galactic latitude are shown in ﬁgure 4.13.
We ﬁtted phenomenological models to the proﬁles to obtain typical scale heights. The intensity
proﬁles were ﬁtted with a one-component model ﬁrst but those in l = ±0.17◦ required a twocomponent model. Thus the every proﬁles were ﬁtted simultaneously with the two-component
models in the same way as the Fe XXV Kα line. The best-ﬁt model is shown in ﬁgure 4.13.
The χ2 /d.o.f of 219.4/76 with the two-component models is better than 296.5/78 with the onecomponent models. The best ﬁt parameters are shown in table 4.7 and summarized in ﬁgure
4.14.
Fe I Kα
The intensity proﬁles of the Fe I Kα line along the galactic latitude are shown in ﬁgure 4.15.
Note that the Fe I Kα line is not detected from the b = −2◦ data signiﬁcantly. The upper limitsof the line intensity and the equivalent width in the 90 % conﬁdence level are 1.2 × 10−8 photons
s−1 arcmin−2 and 97 eV. The intensity proﬁles of the Fe I Kα line seem asymmetric between
the north and the south. We applied a one-component model with diﬀerent scale heights fornorth and south. (hN for b > −0◦ .046 and hS for b ≤ −0◦ .046 ). The best-ﬁt results are shownin ﬁgure 4.15 and table 4.8. The χ2 /d.o.f. is smaller than that of the previous symmetry onecomponent model. The intensity proﬁles of the Fe XXV Kα line were also ﬁtted with this model
for comparison. The comparison of χ2 /d.o.f. is shown in table 4.9. In the case of the Fe XXV
Kα line emission, the χ2 /d.o.f. does not change so much between the symmetry and asymmetry
models.
Intensity in the 5–10 keV Band
The intensity proﬁle of the GDXE in the 5–10 keV band is shown in ﬁgure 4.16. This intensity
includes both of the line and continuum emissions. The CXB and the NXB were subtracted. The

10-8

-2 -1

-7

-2

-0.5

0

0.5

1

-6

10

-7

10

-8

-2 -1

10

-2

-1.5

-1

-0.5

0

0.5

1

!

10

-6

-2 -1

Fe XXV K! l=1.17
10

10-7
10

-8

10

-5

10

-6

10

-7

10

-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

!

-2 -1

Fe XXV K! l=0.5

-2

-1.5

-1

-0.5

0

-6

10

-7

10

10-8

-2

0.5

Galactic latitude (degree)

1

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)!

Fe XXV K! l=-0.17
-5

10

-6

10

-7

10

-8

10

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)-2

Galactic latitude (degree)

-5

10

-2

!

10

!

Fe XXV K! l=-0.84-5!

Fe XXV K! l=0.82
-5

10

-6

10

10-7
-8

10

-2

-2

-2 -1

Observed photons cm s arcmin-1

Fe XXV K! l=-0.56

-2
-2 -1

Observed photons cm s arcmin-1.5

Galactic latitude (degree)-5

Observed photons cm s arcmin

10

Observed photons cm s arcmin-6

Observed photons cm s arcmin

10

Observed photons cm s arcmin

10

-2

-2 -1

Observed photons cm s arcmin!

Fe XXV K! l=-1.16
-5

-2
-2 -1

Observed photons cm s arcmin-2

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION-2

48-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1!

Fe XXV K! l=0.17
-5

10

-6

10

-7

10

-8

10

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

Figure 4.11: Same as ﬁgure 4.9 but with the best-ﬁt two-component models. The error bars
show the 1σ errors. The solid lines show the best-ﬁt two-component models. The dashedand alternate-long-and-short-dash lines show the components with scale heights hb1 and hb2 ,respectively.

4.2. SPATIAL DISTRIBUTION OF THE GDXE

49

Table 4.6: Best-ﬁt parameters of the two-component model to the intensity proﬁles of the Fe
XXV Kα line along the galactic latitude.∗l (degree)
-0.17
-0.56
-0.84
-1.16

Ab1 (10−6 )†1.99 ± 0.18
0.53 ± 0.18
0.90 ± 0.13
0.19 ± 0.13

Ab2 (10−6 )†0.21 ± 0.16
0.43 ± 0.16
0.07(< 0.14)
0.25 ± 0.09

l (degree)
0.17
0.50
0.82
1.17

hb1 (degree)
0.21 ± 0.03

Ab1 (10−6 )†1.86 ± 0.17
0.65 ± 0.18
0.36 ± 0.17
0.05(< 0.19)

Ab2 (10−6 )†0.14 ± 0.12
0.43 ± 0.16
0.29 ± 0.13
0.31 ± 0.11

hb2 (degree)
1.1 ± 0.7

† The uncertainties show the 1σ errors. ∗ The unit is photons s−1 cm−2 arcmin−2 .

Observed photons cm-2 s-1 arcmin-2

Fe XXV Kα
-5

10

Ab1
Ab2

10-6

10-7

10-8
0
0.2
0.4
0.6
0.8
1
1.2
1.4
*
Distance from Sgr A along the longitude l*(degree)

Figure 4.12: Relations among l∗ , Ab1 and Ab2 of the Fe XXV Kα line intensity. The error bars
show the 1σ errors.

-2 -1

-7

10

-2 -1

10

-6

10

10-7

-1.5

-1

-0.5

0

0.5

1

Fe XXVI K! l=1.17!
-5

-2 -1

10

10-6
-7

10

10-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)!

Fe XXVI K! l=0.5
-5

-2 -1

10

10-6
10-7
-8

10

-2

10-6
10

-7

10-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)1-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

Fe XXVI K! l=-0.17!
10

-5

10

-6

10-7
10-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)-2

Galactic latitude (degree)

Observed photons cm s arcmin

-2

-5

-2

Fe XXVI K! l=-0.56!

10-8

!

Fe XXVI K! l=-0.84
10

Fe XXVI K! l=0.82!
10

-5

10-6
10

-7

10-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

-2

-2 -1

Observed photons cm s arcmin

1

-5

-2
-2 -1

Observed photons cm s arcmin-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

Observed photons cm s arcmin-2

Observed photons cm s arcmin

10-8

Observed photons cm s arcmin

10-6

-2

-2 -1

Observed photons cm s arcmin!

Fe XXVI K! l=-1.16
-5

10

-2
-2 -1

Observed photons cm s arcmin-2

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

-2

50

!

Fe XXVI K! l=0.17
10

-5

10-6
10-7
10

-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

Figure 4.13: Intensity proﬁles of the Fe XXVI Kα line along the galactic latitude. The error
bars show the 1σ errors. The solid lines show the best-ﬁt two-component models. The dashedand alternate-long-and-short-dash lines show the components with scale heights hb1 and hb2 ,respectively.

4.2. SPATIAL DISTRIBUTION OF THE GDXE

51

Table 4.7: Best-ﬁt parameters of the two-component model to the intensity proﬁles of the Fe
XXVI Kα line along the galactic latitude.∗l (degree)
-0.17
-0.56
-0.84
-1.16

Ab1 (10−6 )†0.76 ± 0.07
0.26 ± 0.09
0.35 ± 0.06
0.03(< 0.09)

Ab2 (10−6 )†0.05(< 0.11)
0.13 ± 0.08
0.03(< 0.06)
0.06 ± 0.03

l (degree)
0.17
0.50
0.82
1.17

hb1 (degree)
0.21 ± 0.03

Ab1 (10−6 )†0.64 ± 0.08
0.18 ± 0.09
0.19 ± 0.08
< 0.05Ab2 (10−6 )†0.05(< 0.10)
0.16 ± 0.08
0.05(< 0.10)
0.09(< 0.25)

hb2 (degree)
1.2(< 2.4)

† The uncertainties show the 1σ errors. ∗ The unit is photons s−1 cm−2 arcmin−2 .

Observed photons cm-2 s-1 arcmin-2

Fe XXVI Kα
-5

10

Ab1
Ab2

10-6

10-7

10-8
0
0.2
0.4
0.6
0.8
1
1.2
1.4
*
Distance from Sgr A along the longitude l*(degree)

Figure 4.14: Relations among l∗ , Ab1 and Ab2 of the Fe XXVI Kα line intensity. The error bars
show the 1σ errors.

-2 -1

10-6
-7

10

-2 -1

10

-6

10

10-7

-2

-1.5

-1

-0.5

0

0.5

1

Fe I K! l=1.17!
-5

-2 -1

10

10-6
-7

10

10-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)!

Fe I K! l=0.5
-5

-2 -1

10

10-6
10-7
-8

10

-2

10-6
10

-7

10-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

1

Fe I K! l=-0.17!
10

-5

10

-6

10-7
10-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)-2

Galactic latitude (degree)

Observed photons cm s arcmin

10-8

10

Fe I K! l=0.82!
10

-5

10-6
10

-7

10-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

-2

-2 -1

Observed photons cm s arcmin

Fe I K! l=-0.56!
-5

-2
-2 -1

Observed photons cm s arcmin

1!

Fe I K! l=-0.84-5-2-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

Observed photons cm s arcmin-2

Observed photons cm s arcmin

10-8

Observed photons cm s arcmin

10

-2

-2 -1

Observed photons cm s arcmin!

Fe I K! l=-1.16
-5

-2
-2 -1

Observed photons cm s arcmin-2

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

-2

52

!

Fe I K! l=0.17
10

-5

10-6
10-7
10

-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

Figure 4.15: Intensity proﬁles of the Fe I Kα lines along the galactic latitude. The error bars
show the 1σ errors. The solid lines show the best-ﬁt asymmetry one-component models.

4.2. SPATIAL DISTRIBUTION OF THE GDXE

53

Table 4.8: Best-ﬁt parameters of the asymmetric one-component model to the intensity proﬁlesof the Fe I Kα line along the galactic latitude.∗l (degree)
-1.16
-0.84
-0.56
-0.17
0.17
0.5
0.82
1.17

hN (degree)
0.27 ± 0.07
0.30 ± 0.09
0.14 ± 0.05
0.17 ± 0.02
0.11 ± 0.03
0.14 ± 0.02
0.11 ± 0.02
0.34 ± 0.12

A† (10−6 )0.20 ± 0.03
0.34 ± 0.07
1.03 ± 0.12
0.75 ± 0.05
3.26 ± 0.44
1.00 ± 0.08
0.96 ± 0.09
0.29 ± 0.04

hS (degree)
0.52 ± 0.11
0.28 ± 0.05
0.49 ± 0.27
0.16 ± 0.01
0.11 ± 0.01
0.34 ± 0.12
> 0.780.66 ± 0.49

∗ The uncertainties show the 1σ errors. † The unit is photons s−1 cm−2 arcmin−2 .

Table 4.9: Comparison between the symmetric and asymmetric one-component models.∗
Fe I KαSymmetric Asymmetric2409/771315/69

χ2 /do.f.

Fe XXV Kα
Symmetric Asymmetric657/77
611/69∗ The data of b = −2◦ is omitted for the ﬁttings.

interstellar absorption is not corrected. The proﬁle resembles to those of the highly ionized Fe Kαlines. Thus we tried two-component model ﬁttings in the same way as the highly ionized Fe Kαlines. The proﬁles of l = −0◦ .56 and −0◦ .84 are, however, clearly contaminated by 2E 1742.9–
2929 at (l, b) = (−0◦ .44, −0◦ .38) and the Great Annihilator at (l, b) = (−0◦ .89, −0◦ .10) eventhough the bright regions around them were excluded from the spectra. Thus we did not use the
proﬁles of l = −0◦ .56 and −0◦ .84 for the ﬁttings. We point out that the spectra of these brightsources show no Fe Kα lines [62, 109], so the proﬁles of the Fe Kα lines are not aﬀected by them.
For the ﬁttings, three data points of b∗ < 0◦ .2 in the proﬁle at l = 0◦ .17 and those of b = 0◦ .054
– 0◦ .154 in the proﬁle of l = 0◦ .82 were also excluded because they were contaminated by the
local structures around the Sgr A region, The best-ﬁt model is shown in ﬁgure 4.16. The bestﬁt parameters are shown in table 4.10 and summarized in ﬁgure 4.17.

4.2.3

Two-Dimensional Model for the Intensity Proﬁle

The relations between l∗ and Ab1 or Ab2 (the red and green data points in ﬁgure 4.12, 4.14 and4.17) match the proﬁles of the Al1 - and Al2 -components (the cyan and orange lines in ﬁgure
4.4 and 4.6) though we ﬁtted the two-component models to the intensity proﬁles along thegalactic longitude and latitude, independently. It indicates that the two-dimensional intensity
distribution can be modeled well with the next function;(l∗I(l, b) = A1 × exp −
hl1

)

(

b∗
× exp −
hb1

)

(

l∗
+ A2 × exp −
hl2

)

(

b∗
× exp −
hb2

)

(4.2)

We ﬁtted the two-dimesional model to the intensity proﬁles of the highly ionized Fe Kα line
and the 5–10 keV band, respectively. The proﬁles along the galactic longitude and latitude

54

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION°°

5-10 keV intensity l=-1.16

-13

10

-14

10-15

-2 -1

erg cm s arcmin-2 -1

erg cm s arcmin

10

-12

10

-13

10

-14

-2

10

5-10 keV intensity l=-0.84

-2

10

-12

-2

-1.5

-1

-0.5

0

0.5

10-15

1

-2

Galactic latitude (degree)

-13

10

-14

10

-15

10

-12

10

-13

-2 -1

erg cm s arcmin-2 -1

erg cm s arcmin

10

-12

10

-13

10

-14

-2

-1.5 -1 -0.5 0 0.5
Galactic latitude (degree)

1

10

-15

10

-12

10

-13

-2

°

-1.5 -1 -0.5 0 0.5
Galactic latitude (degree)

1

-2
-2 -1

erg cm s arcmin

-2
-2 -1

erg cm s arcmin

15-10 keV intensity l=0.82

10-14
10

-15

-2

-1.5

-1

-0.5

0

0.5

10-14
10

1

-15

-2

Galactic latitude (degree)

-1

-0.5

0

0.5

1

5-10 keV intensity l=0.17°

-12

-12

erg cm s arcmin

-2

10

-2

10

-1.5

Galactic latitude (degree)

5-10 keV intensity l=0.5°

10-13
10

-14

10

-15

-2 -1

-2 -1

0.5

°

5-10 keV intensity l=1.17

erg cm s arcmin

0

-2

10

-0.5

5-10 keV intensity l=-0.17°

-2

10

-1

Galactic latitude (degree)

5-10 keV intensity l=-0.56°
-12

-1.5

-2

-1.5 -1 -0.5 0 0.5
Galactic latitude (degree)

1

10-13
10

-14

10

-15

-2

-1.5 -1 -0.5 0 0.5
Galactic latitude (degree)

1

Figure 4.16: Intensity proﬁles in the 5–10 keV band along the galactic latitude. The error
bars show the 1σ errors. The solid lines show the best-ﬁt two-component models. The dashedand alternate-long-and-short-dash lines show the components with scale heights, hb1 and hb2
respectively.

4.2. SPATIAL DISTRIBUTION OF THE GDXE

55

Table 4.10: Best-ﬁt parameters of the two-component model to the intensity proﬁles in the 5–10
keV band along the galactic latitude.∗l (degree)
-0.17
-0.56
-0.84
-1.16

Ab1 (10−13 )†1.99 ± 0.15––
0.25 ± 0.12

Ab2 (10−13 )†0.29 ± 0.16––
0.30 ± 0.08

l (degree)
0.17
0.50
0.82
1.17

Ab1 (10−13 )†1.85 ± 0.21
1.03 ± 0.17
0.48 ± 0.16
0.06(< 0.19)

hb1 (degree)
0.21 ± 0.03

Ab2 (10−13 )†0.22 ± 0.11
0.36 ± 0.15
0.44 ± 0.13
0.37 ± 0.10

hb2 (degree)
1.2 ± 0.5

† The uncertainties show the 1σ errors. ∗ The unit is ergs s−1 cm−2 arcmin−2 .

5-10 keV intensity-12

erg cm-2 s-1 arcmin-2

10

Ab1
Ab2

10-13

10-14

10-15
0
0.2
0.4
0.6
0.8
1
1.2
1.4
*
Distance from Sgr A along the longitude l*(degree)

Figure 4.17: Relations among l∗ , Ab1 and Ab2 of the 5–10 keV intensity proﬁles. The error bars
show the 1σ errors.

56

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

Table 4.11: Best-ﬁt parameters of the two-dimensional model shown in equation 4.2.∗

Fe XXV Kα
Fe XXVI Kα
5–10 keV intensity

A1 †26.3 ± 0.9
10.0 ± 0.3
3.4 ± 0.2A2 †1.6 ± 0.2
0.25 ± 0.07
0.20 ± 0.04

hl1 (degree)
0.47 ± 0.02
0.51 ± 0.02
0.48 ± 0.04

hb1 (degree)
0.22 ± 0.01
0.22 ± 0.01
0.19 ± 0.02

hl2 (degree)
30 ± 1047(< 102)55(< 133)hb2 (degree)
1.7 ± 0.7
17(< 95)
2.5 ± 1.7

∗ The uncertainties show the 1σ errors. † The units are 10−7 photons s−1 cm−2 arcmin−2 for

the Fe XXV Kα and Fe XXVI Kα lines, and 10−13 erg s−1 cm−2 arcmin−2 for the 5–10 keV
intensity.
were ﬁtted simultaneously. We did not use the data points in the 5–10 keV intensity proﬁleat l = −0◦ .56 and −0◦ .84, and l = −0◦ .96 ∼ −0◦ .75 because they are contaminated by brightsources. The ﬁtting results show in ﬁgure 4.18 – 4.21. The best-ﬁt parameters are summarizedin table 4.11.

4.3

Detailed Spectral Analysis

We merged the spectra of the regions shown in ﬁgure 4.22 to increase statistics. New spectra 1–5
and A–G were shown in ﬁgure 4.23 and 4.24. We called the regions 1–5 and A–G as ”on-plane”
and ”oﬀ-plane” data, respectively. The weighted averages of the RMFs were made using thecounts of the spectra as weights. The ARFs were summed. The central region of l∗ < 0◦ .3 hasbeen studied well with the Suzaku data by Koyama et al. [51, 55] and thus we did not analyze
this region. In the discussion (chapter 7), their results will be referred. We ﬁtted the following
phenomenological model to the merged spectra according to Koyama et al. [51];Abs1 × (P L1 + 10 Gaussians) + Abs2 × CXB.

(4.3)

Compared with the previous model (equation 4.1), weak lines were added, and the center energies
and widths of the Fe Kα lines were free. The center energies and intensities of the weak lines
were constrained according to APEC model [119]. CXB is a power-law model as the CXB
and the normalization and Γ were ﬁxed as described in section 4.1.4. Abs2 was ﬁxed to be
NH = 1.2 × 1023 cm−2 . The emission lines in this model are summarized in table 4.12.
The best-ﬁt models are shown in ﬁgure 4.23 and 4.24. The best-ﬁt parameters are describedin table 4.13 and 4.14. The proﬁles of the line intensity ratio of Fe XXVI Kα to Fe XXV Kαalong the galactic longitude and latitude are shown in ﬁgure 4.25. The result of b = −2◦ is also
added. As described in section A, the line intensity proﬁle ratio corresponds to the temperatureof the plasma. Corresponding temperatures were calculated according to APEC code. The
results are shown in ﬁgure 4.25. The ratio decrease along with l∗ signiﬁcantly. On the other
hand, the ratio does not change along the galactic latitude over b∗ = 0◦ ∼ −2◦ signiﬁcantly.
The proﬁles of the center energies of the three Fe Kα lines are shown in ﬁgure 4.26. The center
energy of Fe XXV Kα seems to get smaller along with l∗ . The proﬁles of the equivalent widths(EWs) of the three Fe Kα are lines shown in ﬁgure 4.27. The EW of the Fe XXV Kα line seemto increase along with l∗ . On the other hand, the EW of the Fe I Kα line decrease along withl∗ . Further discussion using these parameters are made in chapter 7.

4.3. DETAILED SPECTRAL ANALYSIS

57

Observed photons cm-s s-1 arcmin-2

Fe XXV Kα
-5

10

West
East
Yamauchi+09 corrected

10-6

10-7

10-8
0.1
1
10
*
Distance from Sgr A along the longitude l* (degree)

Observed photons cm-s s-1 arcmin-2

Fe XXVI Kα
-5

10

West
East
Yamauchi+09 corrected

10-6

10-7

-8

10

0.1
1
10
*
Distance from Sgr A along the longitude l* (degree)
5-10 keV intensity-12

erg cm-2 s-1 arcmin-2

10

West
East
Yamauchi+09 corrected

10-13

10-14

-15

10

0.1
1
10
*
Distance from Sgr A along the longitude l* (degree)

Figure 4.18: Same as ﬁgure 4.4 and 4.6 but with the best-ﬁt two-dimensional model shown inequation 4.2. Top: Fe XXV Kα line. Middle: Fe XXVI Kα line. Bottom:5–10 keV intensity.
The error bars show the 1σ errors. The black lines show the best-ﬁt two dimensional models.
The cyan and orange lines show the components with A1 and A2 respectively.

10-7
10

-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

Fe XXV K! l=-0.56!
10

-5

10-6
10

-7

10-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

Fe XXV K! l=1.17!
10

-5

10-6
10-7
10-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)!

Fe XXV K! l=0.5
10

-5

10-6
10-7
10

-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

Observed photons cm-s s-1 arcmin-2-6

Observed photons cm-s s-1 arcmin-2

10

Observed photons cm-s s-1 arcmin-2

Observed photons cm-s s-1 arcmin-2
Observed photons cm-s s-1 arcmin-2
Observed photons cm-s s-1 arcmin-2

Fe XXV K! l=-1.16!
10-5

Observed photons cm-s s-1 arcmin-2

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

Observed photons cm-s s-1 arcmin-2

58

Fe XXV K! l=-0.84!
10-5
10

-6

10-7
10

-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

Fe XXV K! l=-0.17!
10

-5

10-6
10

-7

10-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

Fe XXV K! l=0.82!
10

-5

10-6
10-7
10-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)!

Fe XXV K! l=0.17
10

-5

10-6
10-7
10

-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

Figure 4.19: Same as ﬁgure 4.9 but with the best-ﬁt two-dimensional model shown in equation
4.2. The error bars show the 1σ errors. The solid lines show the best-ﬁt two-component models. The dashed and alternate-long-and-short-dash lines show the components with A1 and A2
respectively.

10-7
10

-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

Fe XXVI K! l=-0.56!
10

-5

10-6
10

-7

10-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

Fe XXVI K! l=1.17!
10

-5

10-6
10-7
10-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)!

Fe XXVI K! l=0.5
10

-5

10-6
10-7
10

-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

Observed photons cm-s s-1 arcmin-2-6

Observed photons cm-s s-1 arcmin-2

10

Observed photons cm-s s-1 arcmin-2

Fe XXVI K! l=-1.16!
10-5

59

Observed photons cm-s s-1 arcmin-2

Observed photons cm-s s-1 arcmin-2

Observed photons cm-s s-1 arcmin-2

Observed photons cm-s s-1 arcmin-2

Observed photons cm-s s-1 arcmin-2

4.3. DETAILED SPECTRAL ANALYSISFe XXVI K! l=-0.84!
10-5
10

-6

10-7
10

-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

Fe XXVI K! l=-0.17!
10

-5

10-6
10

-7

10-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

Fe XXVI K! l=0.82!
10

-5

10-6
10-7
10-8

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)!

Fe XXVI K! l=0.17
10

-5

10-6
10-7
10

-8

-2

-1.5 -1 -0.5 0
0.5Galactic latitude (degree)

1

Figure 4.20: Same as ﬁgure 4.13 but with the best-ﬁt two-dimensional model shown in equation
4.2. The error bars show the 1σ errors. The solid lines show the best-ﬁt two-component models. The dashed and alternate-long-and-short-dash lines show the components with A1 and A2
respectively.

60

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

10
10

-13

5-10 keV intensity l=-0.84°
erg cm-2 s-1 arcmin-2

erg cm-2 s-1 arcmin-2

5-10 keV intensity l=-1.16°
-12

10-14
10

-15

-2

-1.5

-1

-0.5

0

0.5

10

-12

10

-13

10-14
10

1

-15

-2

Galactic latitude (degree)

erg cm-2 s-1 arcmin-2

erg cm-2 s-1 arcmin-2

10

-14

10-15

-2

-1.5 -1 -0.5 0 0.5
Galactic latitude (degree)

10

-14

-2

erg cm-2 s-1 arcmin-2

erg cm-2 s-1 arcmin-2-1.5 -1 -0.5 0 0.5
Galactic latitude (degree)1

5-10 keV intensity l=0.82°

10-14

-2

-1.5

-1

-0.5

0

0.5

10-13
10-14
10-15

1

-2

Galactic latitude (degree)

°

-0.5

0

0.5

1

°

erg cm-2 s-1 arcmin-2

10

10-13
10-14
-15

-1.5 -1 -0.5 0 0.5
Galactic latitude (degree)-1

5-10 keV intensity l=0.17

-12

-2

-1.5

Galactic latitude (degree)

5-10 keV intensity l=0.5
erg cm-2 s-1 arcmin-2

1

10-12

10-13

10

0.5

10-13

5-10 keV intensity l=1.17°

10

0

-12

10-15

1

10-12

10-15

-0.5

5-10 keV intensity l=-0.17°

-12

10-13
10

-1

Galactic latitude (degree)

5-10 keV intensity l=-0.56°
10

-1.5

1

-12

10-13
10-14
10

-15

-2

-1.5 -1 -0.5 0 0.5
Galactic latitude (degree)

1

Figure 4.21: Same as ﬁgure 4.16 but with the best-ﬁt two-dimensional model shown in equation
4.2. The error bars show the 1σ errors. The solid lines show the best-ﬁt two-component models. The dashed and alternate-long-and-short-dash lines show the components with A1 and A2
respectively.

4.3. DETAILED SPECTRAL ANALYSIS

61

Figure 4.22: Regions where spectra were merged. They are overlaid on the X-ray image in the
6.55-6.8 keV band.

Table 4.12: Emission lines included in the phenomenological model (equation 4.3).
Line

Fe I KαFe XXV Kα
Fe XXVI Kα
Fe I Kβ
Ni I KαNi XXVII Kα
Fe XXV Kβ
Fe XXVI Kβ
Fe XXV Kγ
Fe XXVI Kγ

Line center
Theoretical
Constraint
value∗ (eV)
6400
free
6680
free
6966
free
7059
[Fe I Kα] × 1.103
7475[Fe I Kα] × 1.168
7821
free
7881
[Ni XXVII Kα] + 60
8251
free
8295
[Fe XXVII Kβ] + 44
8700
[Fe XXVII Kβ] + 449

Width
σ (eV)

Intensity

free
free
free
0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)

free
free
free
[Fe I Kα] × 0.125†
free
free
free
free
free
free

∗ We referred to the APEC model [119] for ionized lines and for neural lines [40].

62

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

Region 2
0.1
0.01
10−3

Counts s−1 keV−1χ5Black:FI Red:BI−4 −2 0 2 10−4

0.01
10−3
−4 −2 0 2 10

χ−4Counts s−1 keV−1

0.1

Region 1
Black:FI Red:BI

105Energy (keV)

Region 4χ−4 −2 0 2 10

0.01
10−3

Counts s−1 keV−1

10

−3

5

Black:FI Red:BI−4 −2 0 2 10−4

0.01

0.1

Black:FI Red:BI−4

Counts s−1 keV−1

0.1

Region 3χ10Energy (keV)

105Energy (keV)

10Energy (keV)

χ

0.01
10−3
−4 −2 0 2 10−4

Counts s−1 keV−1

0.1

Region 5
Black:FI Red:BI

5

10Energy (keV)

Figure 4.23: On-plane merged spectra. The vertical error bars show the 1σ errors. The best-ﬁt
results of the phenomenological model are also shown. Black: the FIs spercta. Red: the BI
spectra.

4.3. DETAILED SPECTRAL ANALYSIS

63
Region B
0.1
0.01
10−3

Counts s−1 keV−1χ5Black:FI Red:BI−4 −2 0 2 10
4 −4

0.01
10−3
−4 −2 0 2 10
4

χ

−4

Counts s−1 keV−1

0.1

Region A
Black:FI Red:BI

10

5

10

Energy (keV)

Energy (keV)
Region D
0.1
0.01
10−3

Counts s−1 keV−1χ5Black:FI Red:BI−4 −2 0 2 10
4 −4

0.01
10

−3

−4 −2 0 2 10
4

χ

−4

Counts s−1 keV−1

0.1

Region C
Black:FI Red:BI

10

5

10

Energy (keV)

Energy (keV)
Region F

0.01

0.1
10

10−3

Counts s−1 keV−1χ5Black:FI Red:BI−4 −2 0 2 10
4 −4

0.01
10

−3

−4 −2 0 2 10
4

χ

−4

Counts s−1 keV−1

0.1

Region E
Black:FI Red:BI

5

10Energy (keV)

Energy (keV)
Region H

10
Energy (keV)

10−3

0.01

0.1
χ

5

Black:FI Red:BI−4 −2 0 2 10
4 −4Counts s−1 keV−1

0.01
10−3
−4 −2 0 2 10
4 −4

χCounts s−1 keV−1

0.1

Region G
Black:FI Red:BI

5

10Energy (keV)

Figure 4.24: Oﬀ-plane merged spectra. The vertical error bars show the 1σ errors. The best-ﬁt
results of the phenomenological model are also shown. Black: the FIs spercta. Red: the BI
spectra.

64

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

Table 4.13: Best-ﬁt parameters of the phenomenological model to the spectra in 5–10 keV bandof the on-plane data.∗
RegionFe I Kα
Energy (eV)Width (eV)Intensity† (10−7 )EW (eV)Fe XXV Kα
Energy (eV)Width (eV)Intensity† (10−7 )EW (eV)Fe XXVI Kα
Energy (eV)Width (eV)Intensity† (10−7 )EW (eV)Ni I KαIntensity† (10−7 )Ni XXVII Kα
Energy (eV)Intensity† (10−7 )Fe XXV Kβ
Intensity† (10−7 )Fe XXVI Kβ
Energy (eV)Intensity† (10−7 )Fe XXV Kγ
Intensity† (10−7 )Fe XXVI Kγ
Intensity† (10−7 )ΓNormalization§ (10−4 )
NH (1022 cm−2 )Intensity# (10−14 )χ2 /d.o.f.

1

2

3

4

5

6404+5−3< 22
2.78+0.18−0.12
235±136396+2−1<7
17.9+0.4
−0.3
494±9

6404+7−2< 21
14.1+1.0−0.6
341±196400+10−10< 34
1.58+0.19−0.20
168±206404+21−20< 65
0.39+0.12−0.05
123±26

6668+4−2
26+7−12
4.39+0.15−0.20
415±16

6671+2−2
20+5−8
13.8+0.3
−0.4
419±9

6684+6−551+7−10
17.8+0.7−1.3
474±23

6670+5−4< 32
3.93+0.22−0.27475±296664+6−4< 24
1.44+0.07−0.10
504±286959+2−11< 21
1.45+0.15−0.12
154±146950+8−3< 12
4.67+0.27
−0.25
157±8

6966+17−7< 55
6.72+1.74−0.48
197±286970+18−10< 33
1.06+0.18−0.18
147±256952+17−22< 44
0.35+0.07−0.07
134±26< 0.130.95+0.51−0.28< 1.27< 0.54< 0.107756+45−156< 0.177790+35−190
0.77+0.26−0.25

7776+33−176
1.50+0.69−0.627803+39−59
0.36+0.21−0.25

7811+25−211
0.13+0.10−0.06

0.43+0.15−0.19

0.79+0.18−0.49

1.50+0.53−1.18

0.36+0.24−0.30

0.13+0.08−0.118203+199−60
0.34+0.21−0.198251+30−31
0.95+0.43−0.328275+181−79< 1.418251+25−68
0.77+0.28−0.418255+186−55
0.28+0.11−0.14< 0.240.34+0.51−0.26

1.42+0.69−1.17< 0.23< 0.090.41+0.15−0.28
2.66+0.02−0.03
1.64+0.02
−0.02
18+1
−1
4.3
219/179

0.63+0.37−0.37
2.33+0.01−0.01
2.72+0.02
−0.02
19+1
−1
13.9
395/248< 1.512.31+0.03−0.02
3.03+0.05
−0.05
18+1
−2
16.1
365/322< 0.27
3.10+0.03−0.05
2.97+0.08
−0.06
22+2
−2
3.1
139/106< 0.22
2.46+0.04−0.04
0.31+0.01
−0.01
4+2
−2
1.6
150/132∗ The uncertainties show the 90 % conﬁdence level. † The unit is photons s−1 cm−2 arcmin−2 .

The interstellar absorption is corrected. § The unit is photons s−1 cm−2 arcmin−2 keV−1 at 1
keV. # The unit is ergs s−1 cm−2 arcmin−2 in the 5–10 keV band. The interstellar absorption
is not corrected.

4.3. DETAILED SPECTRAL ANALYSIS

65

Table 4.14: Best-ﬁt parameters of the phenomenological model to the spectra in 5–10 keV band∗.of the oﬀ-plane data.
RegionFe I Kα
Energy (eV)Width (eV)Intensity† (10−7 )EW (eV)Fe XXV Kα
Energy(eV)Width (eV)Intensity† (10−7 )EW (eV)Fe XXVI Kα
Energy (eV)Width (eV)Intensity† (10−7 )EW (eV)Ni I Kα#Intensity† (10−7 )Ni XXVII Kα
Energy‡ (eV)Intensity† (10−7 )Fe XXV Kβ
Intensity† (10−7 )Fe XXVI Kβ
Energy‡ (eV)Intensity† (10−7 )Fe XXV Kγ
Intensity† (10−7 )Fe XXVI Kγ
Intensity† (10−7 )ΓNormalization§ (10−4 )
NH (1022 cm−2 )Intensity# (10−14 )χ2 /d.o.f.

A

B

C

D

E

F

G

H

6407+11−7< 26
3.17+0.44−0.24
138 ±146398+7−5< 31
2.45+0.20−0.17
135±106425+22−21< 66
0.77+0.20−0.26
112±336426+32−24< 68
0.75+0.26−0.21
101±316402+44−102< 96
0.40+0.26
−0.15
98±37

6378+9−10< 28
2.27+0.30−0.30
130±176404+19−104< 49
1.02+0.26−0.23
119±286408+192−8< 33
1.57+0.21−0.20
150±21

6673+5−2
33+7−8
9.90+0.39−0.44477±206665+3−125+6−7
8.03+0.25−0.23
492±156679+9−7< 33
3.04+0.26−0.31
494±496670+9−9< 35
2.49+0.31−0.23
370±43

6675+13−13< 44
1.42+0.19−0.21392±77

6668+7−3
26+12−11
7.51+0.41−0.39482±286667+9−537+10−20
4.21+0.38−0.34
560±516680+4−5< 23
3.48+0.26−0.22
374±31

6966+5−9< 23
3.80+0.35−0.30
204±17

6957+4−7< 15
2.73+0.20−0.18
187±13

6966+17−22< 63
1.16+0.26−0.21
215±436985+23−35< 62
0.69+0.22−0.19
116±356967+37−36< 80
0.44+0.18−0.16
139±56

6967+12−10< 34
2.51+0.31−0.32
180±226964+26
−18
58+29
−33
1.52+0.39−0.32
234±556962+16−15< 37
1.14+0.22−0.18
137±24< 0.51< 0.26< 0.120.49+0.43−0.23< 0.48< 0.16< 0.26< 0.137768+35−168
1.03+0.21−0.45

7810+40−210
0.70+0.14−0.24

7821+23−221
0.30+0.24−0.227821(ﬁxed)
0.29+0.20−0.277821(ﬁxed)
0.29+0.34−0.207759+51−159
0.75+0.39−0.347821(ﬁxed)
0.57+0.38−0.407821(ﬁxed)
0.13+0.21−0.13

1.03+0.25−0.59

0.70+0.22−0.27< 0.52< 0.34< 0.21< 0.53< 0.78< 0.428248+25−37< 0.858242+28−19
1.04+0.24−0.308254+746−254< 0.458251(ﬁxed)< 0.358251(ﬁxed)< 0.318209+197−31
1.11+0.39−0.608251(ﬁxed)< 1.168251(ﬁxed)< 0.300.67+0.33−0.59< 0.28< 0.27< 0.39< 0.31< 0.31< 1.05< 0.53< 0.44
0.49+0.37−0.38
3.31+0.27
2.67+0.05
−0.13
−0.04
1.49+0.22
4.03+1.27
−0.42
−1.24
19+3
15+2
−2
−3
3.1
3.8
72/69
86/92
∗ The uncertainties show the 90 % conﬁdence level. † The unit is photons s−1 cm−2 arcmin−2 .
0.49+0.50−0.42
2.51+0.03−0.02
2.45+0.03
−0.04
17+1
−1
8.7
220/166

0.54+0.22−0.31
2.59+0.02−0.02
2.20+0.03
−0.02
14+1
−1
7.2
245/226< 0.403.08+0.06−0.32
2.14+0.07
−0.07
16+3
−4
2.6
138/111< 0.402.63+0.10−0.05
0.99+0.04
−0.05
13+3
−3
2.9
132/104< 0.362.84+0.38−0.33
0.79+0.17
−0.03
5+4
−5
1.8
66/76

1.22+0.44−0.54
2.54+0.04−0.03
1.95+0.04
−0.04
14+1
−2
6.9
212/188

The interstellar absorption is corrected. # The center energy of the Ni I Kα line was free for‡ The center
the region D because it was relatively strong. The center energy is 7598+111−58 eV.
energies of the Ni XXVII Kα and Fe XXVI Kβ lines were ﬁxed for the region D, E, G and H
because the statistics is poor. § The unit is photons s−1 cm−2 arcmin−2 keV−1 at 1 keV. #The unit is ergs s−1 cm−2 arcmin−2 in the 5–10 keV band. The interstellar absorption is notcorrected.

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

Ratio of [Fe XXVI Kα] to [Fe XXV Kα]

66

West
East

0.5
0.4

kT=7 keV

0.3
kT=6 keV
0.2
kT=5 keV
0.1

Ratio of [Fe XXVI Kα] to [Fe XXV Kα]

0
0.5
1
1.5
2
2.5
3
3.5
*
Distance from Sgr A along the longitude l* (degree)

0.5
0.4

kT=7 keV

0.3

A
B
C
D
E
F
G
H
b=2°

kT=6 keV
0.2
kT=5 keV
0.1
-2

-1.5

-1 -0.50
0.5
1Galactic latitude b(degree)

1.5

Figure 4.25: Proﬁle of the line intensity ratio of Fe XXVI Kα to Fe XXV Kα. The error bars
show the 90 % conﬁdence level. Top: On-plane. Bottom: Oﬀ-plane.

4.3. DETAILED SPECTRAL ANALYSIS

67

Fe I KαFe I Kα

6.44
Line center (keV)Line center (keV)

6.44
6.43
6.42
6.41
6.4

A
B
C
D
E
F
G
H

6.45

West
East

6.45

6.43
6.42
6.41
6.4

6.39

6.39

6.38

6.38
-1.5

0
0.5
1
1.5
2
2.5
3
3.5
Distance from Sgr A* along the longitude l* (degree)-1

6.7
Line center (keV)Line center (keV)

6.69
6.68
6.67
6.66

A
B
C
D
E
F
G
H

6.71

West
East

6.7

6.69
6.68
6.67
6.66
6.65

6.65

6.64

6.64

-1.5

0
0.5
1
1.5
2
2.5
3
3.5
Distance from Sgr A* along the longitude l* (degree)

-1

7

7.01
West
East

1.5

6.99
6.98
6.97
6.96
6.95

A
B
C
D
E
F
G
H

7
Line center (keV)

7.01

-0.5
0
0.5
1
Galactic latitude b(degree)
Fe XVVI Kα

Fe XXVI Kα

Line center (keV)

1.5

Fe XXV Kα

Fe XXV Kα
6.71

-0.5
0
0.5
1
Galactic latitude b(degree)

6.99
6.98
6.97
6.96
6.95
6.94

6.94
6.93
0
0.5
1
1.5
2
2.5
3
3.5
Distance from Sgr A* along the longitude l* (degree)

6.93
-1.5

-1

-0.5
0
0.5
1
Galactic latitude b(degree)

1.5

Figure 4.26: Proﬁle of the center energies of the Fe Kα lines. The error bars show the 90 %conﬁdence level. Top: Fe I Kα. Middle: Fe XXV Kα. Bottom: Fe XXVI Kα. Left: On-plane.
Right: Oﬀ-plane.

68

CHAPTER 4. GALACTIC DIFFUSE X-RAY EMISSION

Fe I KαFe I Kα
West
East

500

400
EW (eV)

400
EW (eV)

A
B
C
D
E
F
G
H

500

300

300

200

200

100

100
-2

0
0.5
1
1.5
2
2.5
3
3.5
Distance from Sgr A* along the longitude l* (degree)-1.5-1 -0.50
0.5
1Galactic latitude b(degree)
Fe XXV Kα

Fe XXV Kα
West
East

500
400

A
B
C
D
E
F
G
H
°
b=2

600

EW (eV)EW (eV)

600

500
400
300

300

-2

0
0.5
1
1.5
2
2.5
3
3.5
Distance from Sgr A* along the longitude l* (degree)-1.5-1 -0.50
0.5
1Galactic latitude b(degree)

West
East

200

A
B
C
D
E
F
G
H
b=2°

300
EW (eV)300

1.5

Fe XXVI Kα

Fe XXVI Kα

EW (eV)

1.5

200

100

100
0
0.5
1
1.5
2
2.5
3
3.5
Distance from Sgr A* along the longitude l* (degree)

-2

-1.5

-1 -0.50
0.5
1Galactic latitude b(degree)

1.5

Figure 4.27: Proﬁle of the equivalent widths of the Fe Kα lines. The error bars show the 90 %conﬁdence level. Top: Fe I Kα. Middle: Fe XXV Kα. Bottom: Fe XXVI Kα. Left: On-plane.
Right: Oﬀ-plane.

Chapter 5

Sgr A East
In this chapter, we show the result of the Suzaku observation of Sgr A East. As reviewed in
section 2.6, previous observations [92, 112] suggest that this SNR have a peculiar hot component
comparable to the GDXE. We took advantage of the high energy resolution of the Suzaku to
decide the accurate temperature of Sgr A East without the uncertainty of the possible contamination of the non-thermal component. In section 5.1, we describe about the observations andthe data reduction. In section 5.2, we show the result of the spectral analysis. In section 5.3,
we make the short discussion on Sgr A East.

5.1

Observation & Data Reduction

Sgr A East was in the XIS ﬁeld of view in four pointing observations on 2005 September 23–
24 and 29–30 at the aim-points of (α, δ)2000 = (266◦ .51, −28◦ .93) and (266◦ .30, −29◦ .17). The
eﬀective exposure time after removing the epoch of low Earth elevation angles less than 5
degrees, day Earth elevation angles less than 20 degrees, and the South Atlantic Anomaly was
about 90–95 ksec.
The data were taken with the normal clocking full window mode. The SCI technique was
not used because Sgr A East was observed almost just after the launch and the CTI was verysmall. A composite image that adds all of the XIS data is shown in ﬁgure 5.1. The bright spotnear the corner of the XIS ﬁeld of view is Sgr A East. We extracted data from a circle of 1′ .6
radius, added all of the data from three FI CCDs and made a composite X-ray spectrum. The
spectrum of the BI CCD was separately made. Baganoﬀ et al. [3] reports a signiﬁcant diﬀuse
emission not only from the central region of the Sgr A East shell but also from the Sgr A Westregion located at the west in the shell. The diﬀuse emission from Sgr A West is not clear if it
is a part of the Sgr A East SNR or the central black hole Sgr A∗ or other origins. Due to thelimited angular resolution of the Suzaku XRT (half power diameter ∼2′ ), we can not spatiallyresolve Sgr A West from Sgr A East. Therefore, the Sgr A East spectrum we refer to in thispaper includes all of the diﬀuse emission inside the shell of the Sgr A East.
The non-X-ray background (NXB) is made from the night Earth data in the same detector
area with the same data-selection criteria as those of Sgr A East. The largest background of Sgr
A East is the GDXE, which emit the strong Kα lines from FeI, FeXXV and FeXXVI. Koyama etal. [50] reports that the ﬂux ratio of the latter two lines along the Galactic plane (the constantb = −0◦ .046 line) across Sgr A East is almost constant in the l = 0◦ .1 to l = −0◦ .4 area.
This means that the spectral shape of the GDXE is essentially the same near Sgr A East. The
69

70

CHAPTER 5. SGR A EAST

0.100

0.000

-0.100Longitude (l)-0.200-0.100Latitude (b)Figure 5.1: The XIS image near the Sgr A East region in the 2–8 keV band. The white circle of
1′ .6 radius is the source region. The cross inside is the position of Sgr A∗ (the Galactic center).
The rectangles are the unit area extracting the ﬂuxes of the GDXE and of point sources. The
dotted squares indicate the GDXE background regions used for Sgr A East.
2–8 keV ﬂuxes of the GDXE (small triangles in ﬁgure 5.2) are smooth near Sgr A East. We
also note that Park et al.[91] shows the spatial distributions of Fe XXV Kα and Fe I Kα line
emissions with arc-second resolution. The high resolution Chandra image supports that the
GDXE emission is smoothly distributed across Sgr A East. We made the GDXE backgroundtaken from the regions shown by the dotted squares in ﬁgure 5.1, or by the solid bars in ﬁgure
5.2. This GCDX background was subtracted from the Sgr A East spectrum (in the 1′ .6 radiusregion), where the vignetting was corrected.
The second-largest background is an ensemble of point sources. We simulated the pointsource ﬂux in the 2.0–8.0 keV band observed with XIS, putting the positions and ﬂuxes (2.0–
8.0 keV band) of the Chandra cataloged sources [78] in and around the 13 rectangles of ﬁgure 5.1.
After convolving with the response of the XRT+XIS, we extracted the integrated point-source
ﬂuxes from each rectangle. The results are plotted along the Galactic plane (the constantb = −0◦ .046 line), as is shown in ﬁgure 5.2. The mean ﬂux of the integrated point sourcesin this background regions (the regions shown by the solid bars in ﬁgure 5.2) is 0.5×10−5
photons cm−2 arcmin−2 s−1 (2–8 keV). This ﬂux was already subtracted in the process of the
GDXE subtraction. However the integrated point source ﬂux in the source (Sgr A East) region(the 1′ .6 radius region in ﬁgure 5.1) is 3×10−5 photons cm−2 arcmin−2 s−1 (2–8 keV) (see ﬁgure2). Therefore, we further subtracted the excess ﬂux of 2.5×10−5 photons cm−2 arcmin−2 s−1(2–8 keV) from the source spectrum. In this subtraction, we assumed that the integrated point
source spectrum is that obtained by Muno et al. [80].

5.2

Spectral Analysis

5.2.1

Emission Lines

The background (GDXE and point sources) subtracted spectra are shown in ﬁgures 5.3–5.5. We
can see many emission lines. To study these lines, we ﬁt the spectra with a phenomenologicalmodel: a thermal bremsstrahlung plus Gaussian lines with absorption edges of the main atoms.
Fits for the low energy part (2–5 keV) and the high energy part (5–9 keV) were made separately.
The best-ﬁt models and parameters are given in ﬁgures 5.3 and 5.4, and in tables 5.1 and 5.2.

71

100

10

10

−6

Photons s

−1

cm

−2

arcmin−2

5.2. SPECTRAL ANALYSIS

1

0.1

o

0
-0.1
Galactic longitude (l)
o

o

-0.2

o

Figure 5.2: Surface brightness distributions of point sources in the 2–8 keV band (circles). The
surface brightness in the 2–8 keV band of the GDX along the galactic plane (b = −0◦ .046) are
given by small triangles, while that in the Sgr A East region (1′ .6 radius) is shown with thelarge triangle. The solid and dotted arrows show the positions of Sgr A East and Sgr A∗ (theGalactic center), respectively. The solid bars indicate the regions of the GDXE backgrounds for
Sgr A East.

The Kα lines from S XV, Ar XVII, Ca XIX and Fe XXV (helium-like atoms) have been reported
by Sakano et al.[112] and Maeda et al.[63], and the presence of Kα lines from S XVI, Ar XVIII
and Fe XXVI (hydrogenic atoms) have been suggested by Sakano et al.[112]. We clearly foundthese emission lines and newly detected a Kα line from Ni XXVII (helium-like nickel), as wellas Kβ lines from SXV, Ar XVII and Fe XXV, and a Kγ line from Fe XXV.

5.2.2

Sulfur Line Diagnostics

The Kα ﬂux ratio of SXVI/SXV is 0.20+0.03−0.06 (in this chapter, all errors are 90% conﬁdence
intervals for one interesting parameter, unless otherwise described). This ratio indicates that+0.02
the ionization temperature is 1.1 +0.1
−0.1 keV. The ﬂux ratio of Kβ/Kα from SXV is 0.08 −0.03 ,+0.7which constrains the electron temperature to be 1.4−0.8 keV. Since the center energy of theKα line from SXV decreases as decreasing plasma temperature by increasing contribution of
the satellite lines, we can check the plasma temperature using the center energy. The plasma
temperature determined from the line centroid is in the range of 0.2–0.6 keV. It seems a little
lower than the temperature determined from the line ratio. However if we take into account thesystematic error of the XIS gain (∼ ±5 eV in this observation data), the allowable range becomes
0.2–2.0 keV , which is consistent with that determined from the line ratio. Therefore, from theemission line alone, we deduce that Sgr A East has at least a ∼1-keV plasma in collisional
ionization equilibrium (CIE).

CHAPTER 5. SGR A EAST

0.01

Ar XVII Kα
S XV Kα

Ca XIX Kα
Ar XVII Kβ

S XVI KαS XV Kβ Ar XVIII Kα−2 0 2χCounts s−1 keV−1

0.1

72

2

3

4

5

0.01

Ar XVII Kα

Ca XIX KαS XV KαS XVI KαS XV Kβ Ar XVIII Kα−2 0 2χCounts s−1 keV−1

0.1

Energy (keV)

2

3

4

5

Energy (keV)

Figure 5.3: Top: X-ray spectrum of the sum of 3 FI CCDs with the best-ﬁt phenomenologicalmodel in the 2.0–5.0 keV band. Bottom: same as the top, but of the BI CCD.

5.2.3

Iron and Nickel Line Diagnostics

The Kα ﬂux ratio of Fe XXVI /Fe XXV is 0.05+0.01−0.02 . This indicates that the ionization tempera+0.3
ture is 3.1−0.4 keV. The ratio Kβ / Kα of Fe XXV can not be determined accurately, because theKβ line of Fe XXV is contaminated by the nearby Ni XXVII Kα line. We therefore examined
the ﬂux ratio of Fe XXV Kα vs Ni XXVII Kα plus Fe XXV Kβ. The observed ratio is 0.10 +0.02−0.03 .
If the abundance ratios of Fe and Ni are solar, we can then constrain the electron temperature
to be 2.5 +1.0
−0.8 keV.
The center energy of Fe XXV Kα is 6650 (6648–6653) eV, which is signiﬁcantly lower thanthat from an ∼3 keV temperature plasma (e.g. MEKAL [39, 61] and APEC [119] models give6673eV and 6681 eV, respectively). In fact this line-center energy gives a plasma temperatureof 1.4+0.1−0.1 keV, consistent with those determined from the line -ﬂux ratio of sulfur. Furthermore,
the line width of 44 eV is signiﬁcantly larger than that observed from the GDXE [51]. These
facts indicate that the Fe XXV Kα may include lines from lower ionization iron associated witha lower temperature plasma.

73

0.01
10−3

Fe XXV Kα
Fe XXVI Kα

Fe XXV Kβ
Fe XXV Kγ−4−2 0 2

χCounts s−1 keV−1

0.1

5.2. SPECTRAL ANALYSIS

6

7Energy (keV)

8

0.01
10−3

Fe XXV Kα

Fe XXV Kβ

Fe XXVI Kα−4−2 0 2

χCounts s−1 keV−1

0.1

5

5

6

7
Energy (keV)

8

Figure 5.4: Same as ﬁgure 3, but in the 5–9 keV band.

The errors in the parentheses are at 90% conﬁdence level.
a Fixed to 30 × (E/5895)0.5 eV, where the width of the calibration sources (Mn I Kα at 5895 eV) is 30 eV.
b Fixed to E (Ar XVII Kα) + 200 eVc Fixed to E (Ca XIX Kα) - 190 eV∫
d The units of 3.02 × 10−15 /(4πD 2 ) n n dV , where D, n and n are distance to the source (cm), the electron density (cm−3 ), and the ion
e IeIdensity (cm−3 ), respectively.
e Abundances relative to solar.
f Flux is uncorrected for absorption, while intensity and luminosity are corrected.

242/266 = 0.91

Luminosity (1.5–5.0 keV) f

photons s−1 cm−2
1.1 × 10−3

Flux (1.5–5.0 keV) f

χ2 /d.o.f

ZS e
2.9 (2.7–3.3)ZSi e
3.2 (3.1–3.3)

NH (× 1022 cm−2 )
9.0 (8.9–9.1)

Absorption

ergs s−1
3.9 × 1035

Norm. d4.34 (4.30–4.39) ×10−2

Temperature(keV)
1.89 (1.86–1.93)

Thermal bremss

Intensity f(photons s−1 cm−2 )9.53 (8.66–10.1) ×10−4
1.86 (1.38–2.23) ×10−4
7.47 (4.69–9.64) ×10−5
2.27 (2.07–2.51) ×10−4
1.46 (0.08–2.61) ×10−5
1.86 (0.93–2.68) ×10−55.74 (5.03–6.68) ×10−5

Energy(eV)2446 (2440–2449)
2625 (2603–2631)
2869 (2840–2889)
3124 (3116–3130)
3324 b3692 c3882 (3873–3890)

Line
Identiﬁcation
S XV KαS XVI KαS XV KβAr XVII Kα
Ar XVIII Kα
Ar XVII Kβ
Ca XIX Kα

Width
(eV)
22 (16–27)
20 (ﬁxed) a
21 (ﬁxed) a
47 (39–56)
23 (ﬁxed) a
24 (ﬁxed) a
24 (10–39)

Table 5.1: Best ﬁt parameters of the phenomenological model in the 2.0–5.0 keV band

74
CHAPTER 5. SGR A EASTAbundance relative to solar.

Flux is uncorrected for absorption, while intensity and luminosity are corrected.

c

d

e

ne nI dV , where D, ne and nI are distance to the source (cm), the electron density (cm−3 ), and the ion density (cm−3 ), respectively.

The units of 3.02 × 10−15 /(4πD2 )b∫This line is a mixture of Ni XXVII Kα (∼7.82 keV) and Fe XXV Kβ (∼7.88 keV).

Fixed to 30 × (E/5895)0.5 eV, where the width of the calibration sources (Mn I Kα at 5895 eV) is 30 eV.

a

The errors in the parentheses are at 90% confidence level.

194/155 = 1.25

ergs s−1
7.4 × 1034

photons s−1 cm−2
7.1 × 10−4

Flux (5.0–9.0 keV) e

χ2 /d.o.f

ZFe d
2.8 (2.3–3.2)

NH (×1022 cm−2 )
9.0 (ﬁxed)

Absorption
Luminosity (5.0–9.0 keV) e

Norm. c9.9 (9.7–10) ×10−3

Temperature(keV)
4.5 (4.3–4.6)

Thermal bremss

Intensity e(photons s−1 cm−2 )2.23 (2.16–2.29)×10−4
1.09 (0.75–1.40)×10−5
2.12 (1.49–2.53)×10−5
1.32 (0.66–1.64)×10−5

Energy(eV)6650 (6648–6653)
6956 (6938–6981)
7820 (7790–7847)
8238 (8199–8286)

Line
Identiﬁcation
Fe XXV Kα
Fe XXVI Kα
Fe XXV Kβ a
Fe XXV Kγ

Width
(eV)
44 (41–47)
0 (<34)83 (53–113)
35 (ﬁxed) b

Table 5.2: The best ﬁt parameters of the phenomenological model in the 5.0–9.0 keV band

5.2. SPECTRAL ANALYSIS
75

76

5.2.4

CHAPTER 5. SGR A EAST

Thin Thermal Plasmas and Power-Law Component

From the emission line information (sections 5.2.2 and 5.2.3), Sgr A East requires at least a
two-temperature plasma. The lower temperature plasma is CIE, while no constraint on the hightemperature plasma is obtained from the line information. We therefore ﬁt the full energy rangeof the spectrum with a model of a ∼1-keV CIE plasma plus a higher temperature plasma. Since
the spectrum has a systematic line broadening due to the incomplete response function[50], we
artiﬁcially add a 1400 km s−1 random motion in the plasma code (APEC) to compensate forthe systematic error. The reduced χ2 has 2 local minima: ∼43–58 keV with nt >1013 s cm−3
(CIE plasma) , and ∼37–54 keV with nt=1011 s cm−3 (non-ionization equilibrium: NEI plasma),where nt is the ionization parameter. Because these temperatures are outside the XIS bandpass, the parameter ranges for the ﬁt are unreliable. The reduced χ2 values are 1.1 with 460–470
degree of freedoms (d.o.f.) at both the local minima. The two-CIE model can not reproduce
the Fe XXV Kβ line ﬂux at ∼ 7.88 keV, while the CIE+NEI model fails to ﬁt the Ni XXVII
Kα line ﬂux at ∼ 7.82 keV. In both cases, signiﬁcant residuals are found above the Fe XXV Kαenergy (7–9 keV). For the CIE case, absorption edge of Fe is 0.5 solar, while for the NEI case,
it is as large as 3.4 of solar. These unrealistic parameters and/or signiﬁcant residuals above the
Fe XXV Kα energy surely come from the artiﬁcial plasma ﬁtting for the hard X-ray band. We
therefore added a power-law component to a two-CIE model in order to naturally explain thehigh energy band. Although no essential change is found below 7 keV, a signiﬁcant improvement
is noted in the 7–9 keV band, with the reduced χ2 of 1.07 (for 453 d.o.f.). The best ﬁt results
are given in ﬁgure 5.5 and table 5.2.4.

5.3

Short Discussion for Sgr A East

5.3.1

Thermal Plasma

The X-ray luminosity of the thermal plasma is 4.2×1035 erg s−1 at a distance of 8 kpc. Assuming
a uniform density plasma in a sphere of radius (1′ .6), we estimate the plasma density to be 4.7
H cm−3 for the 1-keV component and 0.6 H cm−3 for the 6-keV component. The dynamical age
estimated from the SNR size divided by the sound velocity in the high energy plasma is about
4000 years. These arguments are consistent with the previous studies [63, 92, 112].

5.3.2

Hard Tail

The Sgr A East spectrum requires a hard tail component with a power-law index of 0.87 (0.84–0.91), although Park et al. [92] and Sakano et al. [112] detected no power-law component in the
Chandra or XMM spectra. Since the non-X-ray background of XIS is far lower than those of
Chandra and XMM, particularly, in the high energy band (∼≥ 6 keV, see [50]), where the hard
tail is manifested, we believe that our result is more reliable.
Since the response function of the XIS at a high energy of around 10 keV is not well established yet, the absolute value of the photon index may have a larger error than the statisticalone. Nevertheless the presence of a hard tail is not in doubt. The hard tail ﬂux in the 7–9
keV band is 1.2 times that of the thin thermal components. Thus, an important question is
the origin of the large hard tail. One possibility of the hard tail origin is integrated emissionof point sources fainter than the limiting ﬂux (3×10−15 erg cm−2 s−1 ) of the Chandra deep

77

0.01
10−3

VAPEC[1]VAPEC[2]Power−law−4−2 0 2

χCounts s−1 keV−1

0.1

5.3. SHORT DISCUSSION FOR SGR A EAST

2

5

0.01
10−3

VAPEC[1]VAPEC[2]Power−law−4−2 0 2

χCounts s−1 keV−1

0.1

Energy (keV)

2

5Energy (keV)

Figure 5.5: Top: the X-ray spectrum of the sum of 3 FI CCDs with the best-ﬁt two-temperatureplasma plus power-law model. Bottom: same as the top, but of the BI CCD.

483/453 = 1.07

Abundances relative to solar.

Fixed to the abundance of Fe.

Flux is not corrected for absorption, while luminosity is corrected.

Filling factor (0 < f < 1). Distance of 8 kpc is assumed (Reid et al. 1993).

d

e

f

g

ne nI dV , where D, ne and nH are distance to the source (cm), the electron density (cm−3 ), and the hydrogen density (cm−3 ), respectively.

The units of 10−14 /(4πD2 )c∫Random motion of 1400 km s−1 velocity is added to the APEC models to compensate the systematic line broadening obtained with the calibration line (MnI Kα).

b

ZFe d
1.33 (1.26–1.40)

The unit is photons keV−1 cm−2 s−1 at 1 keV.

iron (M⊙ )0.15f 0.5

ergs s−1
4.5 × 1035

ZS d
2.9 (2.7–3.0)aThe errors in the parentheses are at 90% confidence level.

χ2 /d.o.f

total (M⊙ )27f 0.5

Luminosity (2.0–9.0 keV) f

photons s−1 cm−2
1.8 × 10−3

Flux (2.0–9.0 keV) f

Plasma mass g

ZSi d
3.2 (3.1–3.3)

NH (× 1022 cm−2 )9.66 (9.58–9.74)

Absorption

Norm. a3.5 (3.3–3.7)×10−4

APEC[1]APEC[2]
1.21 (1.18–1.30)
6.0 (5.5–6.4)
0.165 (0.161–0.169)
2.7 (2.5–2.9)×10−3
1.17 (1.07–1.24)
1.20 (1.07–1.33)
1.18 (1.03–1.33)
2.63 (2.54–2.70)
2.63 eindex
0.87 (0.84–0.91)

Thin thermal plasma b
Temperature (keV)
Norm. c
ZS d
ZAr d
ZCa d
ZFe d
ZNi d

Power-law

Table 5.3: Best ﬁt parameters of the two-temperature plasma plus power-law model in the 2.0–9.0 keV band

78
CHAPTER 5. SGR A EAST

5.3. SHORT DISCUSSION FOR SGR A EAST

79

exposure (Muno et al. [78]). The fact that the power-law index of the hard tail spectrum ofΓ ∼0.8–0.9 is similar to that from the point sources of Γ ∼0.8–0.9[80] may favor this possibility.
The integrated ﬂux of point sources above the limiting luminosity in the Sgr A East region (a1′ .6 radius) is 2.2×10−12 erg cm−2 s−1 . The ﬂux of the hard tail is calculated to be 2.1×10−12
erg cm−2 s−1 (2–8keV) , nearly the same as that of the integrated point sources. Muno et al.[78]reports that power-index of the logN -logS slope below the ﬂux (S) of 8×10−7 photons cm−2 s−1(6.4×10−15 erg cm−2 s−1 ) is 1.7 for point sources near the GC. If this luminosity distribution is
the same for point sources in the Sgr A East region, and extends by a factor of 5 lower ﬂux thanthe current Chandra detection limit, then point sources would account for most of the hard tailﬂux.
The other possibility is clumps and ﬁlaments found by Chandra in Sgr A East, becausesigniﬁcant fractions of these ﬁlaments have non-thermal X-ray spectra [49, 71]. In this case, thehard tail must be closely related to a high energy cosmic ray origin. Sgr A East locates very
near the GC, which would be a cite of extreme physical conditions. Therefore, more extensive
study of the hard X-ray tail using HXD [121] on board Suzaku, for example, is highly required.

5.3.3

Chemical Compositions

The chemical compositions averaged over the SNR are nearly solar for S, Ar and Ca, while
Fe is overabundant, which is consistent with the Chandra and XMM results [63, 112, 92] The
abundance of Fe averaged over the whole SNR is 2–3 of solar. We estimated the total mass ofiron and hot gas to be 0.15f 0.5 M⊙ and 27f 0.5 M⊙ , where f is ﬁlling factor. This is consistent
with the result of Park et al. [92], who found that the mass of Fe was <0.27M⊙ within 1′ radius.
As for the interstellar absorption, on the other hand, the Si and S abundances in the ISM
are 3 of solar, while iron is moderately overabundant. Similar over abundance are found fromother X-ray objects near the GC [51, 130], and hence would be common phenomena.

Chapter 6

Suzaku J1740.5–3014
We found a new X-ray source, Suzaku J1740.5–3014, in the data set of the GC observation.
In this chapter, we report the results of the analysis of Suzaku J1740.5–3014. If it is locatedin the intermediate region between the GC and the ridge, it is a precious sample because theproperty of the point sources around this region has not been studied well. In section 6.1, we
describe about the observations and the data reduction. In section 6.2, we show the result of
the analysis. In section 6.3, we discuss the properties and distance of Suzaku J1740.5–3014.

6.1

Observation & Data Reduction

With the Suzaku XIS, regions near a radio pulsar PSR B1737−30 [21] were observed twice. The
observation logs are listed in table 6.1.
Table 6.1: Observation data list.

TargetObs. ID
Start Time
Stop Time
GC LARGEPROJECT15 503021010 2008-10-04 03:44:03 2008-10-05 10:57:24
HESSJ1741-B
503077010 2009-02-26 01:00:60 2009-02-27 11:35:19* Eﬀective exposure after the data screening described in the text.

Exp. (ks)*
50.1
46.2

The XIS observations were made in the normal clocking full window mode. We used cleaned
event data distributed from DARTS 1 . The data processing versions are 2.2.11.22 and 2.2.11.24
for the ﬁrst and second observations, respectively. The data during the epoch of low Earthelevation angles less than 5 degrees, day Earth elevation angles less than 20 degrees, and theSouth Atlantic Anomaly was removed. The eﬀective exposure times are listed in table 6.1. The
SCI technique was used in these observations. We conﬁrmed that the spectral resolutions at 5.9
keV were ∼150 and ∼200 eV (FWHM) for the FI and BI CCDs in these observations using theonboard calibration sources.
Due to a star tracker (STT) problem, the Suzaku attitude during the second observation
was not locked at the programmed position, and hence drifted by ∼ 2′ in the direction of theright ascension between the start and stop time of the second observation.
1<http://www.darts.isas.jaxa.jp/astro/suzaku.>

81

82

CHAPTER 6. SUZAKU J1740.5–3014

6.2

Analysis

We analyzed the data using the software package HEASoft 6.5.1. In this chapter, uncertainties
are quoted at the 90 % conﬁdence range unless otherwise stated.

6.2.1

Imaging

To make X-ray images, we used only the data of October 2008 when the STT operated correctly.
The data of XIS 0, 1 and 3 were merged. We obtained COR-sorted NXB images using xisnxbgen[125]. After subtracting the NXB image, we divided the X-ray image by an image of the ﬂatﬁeld to correct vignetting. The image of the ﬂat ﬁeld was made with xissim [34].
Since Suzaku has some position error [135], we ﬁne-tuned the nominal Suzaku coordinate
using the positions of catalogued visible and near-infrared stars. Since the X-ray from normalstars is generally soft, we made 0.5−2.0 keV band image. In the X-ray image, there are threepoint sources which corresponds to catalogued visible and near-infrared stars. Using the threestars, we shifted the Suzaku coordinate as is shown in ﬁgure 6.1. Before this correction, averaged
diﬀerence between the coordinates of cataloged stars and X-ray point sources was 40′′ ± 11′′ ,where the error shows the statical uncertainty of the peak positions of the X-ray point sources.
After this correction, the averaged diﬀerence became 9′′ .
No signiﬁcant X-ray is found from the radio pulsar PSR B1737−30, but a bright new sourceis discovered at about 90′′ east of PSR B1737−30. From the ﬁne tuned position, we designate
this source as Suzaku J1740.5−3014. The position of Suzaku J1740.5−3014 is (α, δ)2000 = (17 :40 : 35.6, −30 : 14 : 16) with the uncertainty of 14′′ . In the Galactic coordinate, the position is(l, b) = (−1◦ .682, 0◦ .245).
Using the ﬁne-tuned position, we constructed the 1−9 keV band image in ﬁgure 6.2. The
peak positions of Suzaku J1740.5−3014 are same in the 0.5−2 and 1−9 keV band images.

(b)

(a)

0.4

0.3

Galactic latitude (degree)

Galactic latitude (degree)

0.4

1
2

0.2
3
1

0.3

2

0.2
3
1

0.1

0.1

0.0

0.0

358.5

358.4

358.3

358.2

Galactic longitude (degree)

1

358.5

358.4

358.3

358.2

Galactic longitude (degree)

Figure 6.1: X-ray image in the 0.5−2 keV band before (a) and after the coordinate correction (b).
The numbered circles represent catalogued stars used as references of the position. (1) DENIS-P
J174047.6−300613 (2) HD 316162 (3) DENIS-P J174054.9−301915 [16, 28]. The rectangle markshows the position of PSR B1737−30 [21].

6.2. ANALYSIS

83

Galactic latitude (degree)

0.4

0.3

0.2

0.1

0.0
358.5

358.4

358.3

358.2

Galactic longitude (degree)
Figure 6.2: X-ray image in the 1−9 keV band after the coordinate correction. The source andbackground regions are shown by the solid circle and dash annuals, respectively.

6.2.2

Spectrum

We made spectra of the two observation separately. For the ﬁrst observation, the spectrum of
Suzaku J1740.5−3014 for each sensor was extracted from the solid circle region in ﬁgure 6.2
with a radius of 1′ .5. Background spectra were evaluated from the dashed annulus region inﬁgure 6.2 having inner and outer radii of 2′ and 4′ . The background spectra were subtracted
from the spectra of Suzaku J1740.5−3014.
In the second observation, however, the STT was troubled, and the attitude had a large
ambiguity. To correct it, we divided the second observation into six epochs. We made an X-ray
image of each epoch, and conﬁrmed that the attitude drift caused by the STT problem wasless than 20′′ in each epoch. This value is as small as the nominal positional uncertainty of the
XIS [135]. Then we corrected the coordinates of each epoch according to the image of the ﬁrst
observation (ﬁgure 6.1). After the correction, we obtained the spectra for the second observationfrom the regions shown in ﬁgure 6.2.
The spectra of the two FIs are co-added because the response functions are almost same
between the FIs. We made RMF ﬁles using xisrmfgen, and ARF ﬁles using xissimarfgen [34]for each observation.
To check the diﬀerence of the spectra between the two observations, we divided the spectraby the corresponding ARF ﬁles to correct the diﬀerence of the eﬀective area and took the ratioof the spectra of the two observations. The ratio is consistent with unity by χ2 −test in the 90%
conﬁdence level (χ2 /d.of.=89/98) and thus we judged that the spectra of the two observations
are not diﬀerent. We combined the spectra of XIS of the both observations to increase thestatistics. The RMF ﬁles and ARF ﬁles were also combined. The X-ray spectrum is shown inCHAPTER 6. SUZAKU J1740.5–3014

10−3−2 0 2 10−4χCounts s−1 keV−1

0.01

84

1

2

5

10

Energy (keV)

Figure 6.3: Background-subtracted spectrum of Suzaku J1740.5−3014 for the FI CCD. The
source spectrum was made from the solid circle in ﬁgures 6.2, while the local background spectrum was taken from the dash annals. The vertical error bars of each data point are the 1σ
error. The solid line is the best-ﬁt result of two-component CIE model (table 6.2(b)). The dashand dot lines represent high and low temperature CIE components. The dash-dot lines are Fe I
K α and β.
ﬁgure 6.3. We show only the spectrum of the FI CCD to be easy-to-see.
The X-ray spectra exhibit three lines at 6.4, 6.7, and 7.0 keV. In order to identify these lines,
we ﬁtted the spectra in the 5–10 keV band with a phenomenological model of an absorbed powerlaw model plus three narrow Gaussian lines. The cross section of the photoelectric absorption
was obtained from [76]. The best-ﬁt line energies with the equivalent widths (EW ) of the lines
+30
+60
+0.04
are 6.40+0.04−0.02 keV with 210−50 eV, 6.67 ± 0.02 keV with 190−40 eV, and 6.95−0.05 keV with+40
130−50 eV, respectively. From the line energies, these lines are identiﬁed as the Kα lines from
Fe I, XXV and XXVI.
We therefore ﬁtted the X-ray spectra of the FI and BI sensors simultaneously with a collisional ionization equilibrium (CIE) plasma model (APEC: [119]) for highly ionized iron lines
plus two Gaussian lines for Fe I Kα and Kβ lines. The center energy ratio of the two Gaussianlines was ﬁxed to be 1:1.103 and the ﬂux ratio to be 1:0.125 [40]. The abundance of the CIE
plasma model was free, while the relative ratios among various elements were ﬁxed to be thesolar value. The solar abundance of each element was obtained from [1]. The model did not ﬁtwell with χ2 /d.o.f. = 363.3/166. We also tried two two-component CIE plasma model wherean absorption column density and an abundance are common to the two components. However,this model did not describe the spectra with χ2 /d.o.f. = 275.7/164.
Since large residuals were found in the low energy band, we applied a one-component CIE
plasma model with partial absorption. This model gave a marginal χ2 /d.o.f. of 204.9/164. The
best-ﬁt parameters are listed in table (a).
We ﬁnally tried a two-component CIE plasma model with the independent absorptions anda common abundance, and obtained a reasonable ﬁt with χ2 /d.o.f. = 190.6/163. The best-ﬁt
parameters are given in table (b), and the best-ﬁt model together with the spectrum is shownin ﬁgure 6.3.

6.2. ANALYSIS

85

Table 6.2: Best-ﬁt parameters of the model ﬁttings for the background-subtracted spectra∗ .
(a) one-component CIE plasma model with partial absorption.
Model: Abs1×{(1-α)+α×Abs2)×(APEC+Neutral Iron Lines)}
Iron Lines
Fe I KαFe I Kβ+0.03
Energy (keV)
6.40−0.02
7.06†
Flux (10−6 ph s−1 cm−2 )8.8+1.7
1.1†−1.3
Parameter
APEC
kT (keV)
9.5+1.2−0.9
Abundance (solar)
0.49+0.11−0.10
Normalization‡
2.8+0.1−0.2
Parameter
Abs1
Abs2
+0.1
22
−2
NH (10 cm )1.7−0.2
16 ± 3Covering factor α−
0.71+0.02−0.03
Flux (10−12 erg s−1 cm−2 )§2.0 (2.2)
χ2 /d.o.f. = 204.9/164(b) two-component CIE plasma model
Model: Abs1×(APEC1+Neutral iron lines)+Abs2×APEC2Iron Lines
Fe I KαFe I Kβ+0.03
Energy (keV)
6.40−0.02
7.06†
+1.7
−6
−1
−2
Flux (10 ph s cm )8.4−1.4
1.0†
Parameter
APEC1
APEC2
kT (keV)64(≥ 44)6.0 ± 1.1
+0.2∥
Abundance (solar)
1.3−0.3
Normalization‡
1.6+0.2
0.62+0.04
−0.1
−0.03
Parameter
Abs1
Abs2
NH (1022 cm−2 )11+2
1.6 ± 0.1
−1
−12
−1
−2
§
Flux (10erg s cm )2.1 (2.3)
2χ /d.o.f. = 190.6/163∗ Errors show the 90% conﬁdence level.
† The energy and ﬂux of Fe I Kβ line are ﬁxed to be 1.103 and 0.125 times those of Fe I Kαline, respectively.
∫‡ The units of 10−17 /(4πD 2 ) n n dV , where D, n and ne He
H are the distance to the source−3(cm), the electron density (cm ), and the hydrogen density (cm−3 ), respectively.
§ Observed ﬂux in the 2–10 keV band. Values in the parentheses are absorption-corrected ﬂuxes,
where we use NH = 1.7 × 1022 cm−2 for the case (a) and NH = 1.6 × 1022 cm−2 for the case (b)as an interstellar absorption.
∥ The abundance is common to the two APEC components.

CHAPTER 6. SUZAKU J1740.5–3014

0.2
0.1
0

Power [(counts s−1)2 Hz−1]

0.3

86

10−4

10−3
Frequency (Hz)0.01

Figure 6.4: Power spectrum in the 1–9 keV band obtained by the FFT analysis. This is theresult of the combined data of the two observations.

6.2.3

TimingWe extracted light curves of Suzaku J1740.5−3014 from the circular region in ﬁgure 6.2. We
made the light curves of the two observations separately. The light curves of the two FI sensors
were also combined. However, we did not use the BI data, because the NXB dominates the data
above 7 keV, and the NXB is time valuable. We searched for pulsations in the 1–9 keV band fromthe FI data. The fast Fourier transform (FFT) analysis revealed clear peaks at ∼ 2.31 × 10−3
Hz from the light curves of the both observations We then searched for an accurate pulse periodwith the folding technique, and found pulse periods of 432.1 ± 0.2 s and 431.9 ± 0.3 s from thedata of October 2008 and February 2009. The pulse periods are consistent with the errors andthus we combined the light curves of the two observations to increase the statistics. We made a
coherent FFT analysis and a folding search for the combined light curves. The obtained power
spectrum is shown in ﬁgure 6.4. The pulse period is 432.1 ± 0.1 s. The error of the pulse period
was estimated in reference to [59], considering the highest harmonic.
Since we found the two-component CIE plasma in the spectral analysis, we made folded
pulse proﬁles in the 1–3 keV and 4–8 keV bands; they represent the low- and high-temperaturecomponents, respectively. These folded pulse proﬁles are shown in ﬁgure 6.5. We calculated theratio of the low-energy proﬁle to the high-energy proﬁle, and the ratio is ﬁtted with a constantvalue with χ2 /d.o.f.= 38.7/25. Thus we reject a hypothesis that the low-energy proﬁle is aconstant fraction of the high-energy proﬁle with the 96% conﬁdence level. The pulse proﬁle ofthe high-temperature plasma is quasi-sinusoidal with small humps on both sides of the main
peak, while that of the low-temperature plasma is more complicated with nearly equal threepeaks.

6.3. SHORT DISCUSSION FOR SUZAKU J1740.5-3014

87(a) 1-3 keV(b) 4-8 keV

0.03

0.07
0.06
0.05
-1

Count sCount s

-1

0.02

0.01

0.04
0.03
0.02
0.01

0

0
0

0.5

1
Phase

1.5

2

0

0.5

1
Phase

1.5

2

Figure 6.5: Pulse proﬁles in the 1–3 keV band (a) and in the 4–8 keV band (b). Folded period
is 432.1 s. The vertical error bars of each data point are the 1σ error. The dotted lines show
the background levels.

Figure 6.6: Positions of Suzaku J1740.5–3014, AX J1740.5–3014 and SAX J1740.5–3013[110, 30].
The circles show the 90% error uncertainties of the positions. They are overlaid on the Suzaku
image of 1–9 keV.

6.3

Short Discussion for Suzaku J1740.5-3014

6.3.1

Identify of Suzaku J1740.5-3014 by Comparison with Previous X-ray
Observations

We searched catalogs of X-ray sources to identify Suzaku J1740.5-3014. Two sources, AX
J1740.5–3014[110] and SAX J1740.5–3013[30, 31] were found near Suzaku J1740.5–3014. The
positions of them are shown in ﬁgure 6.6 where the circles show the 90% uncertainties of thepositions. The spectral properties of the each observation are summarized in table 6.3. The
positions and spectral properties almost match in the uncertainties so Suzaku J1740.5-3014, AX
J1740.5–3014 and SAX J1740.5–3013 are probably a same object. The detection of Fe Kα lines
and the pulsation had not been reported in the previous observations. Due to the high energy
resolution and high statistics of the Suzaku data, we can ﬁnd them for the ﬁrst time.

88

CHAPTER 6. SUZAKU J1740.5–3014

Table 6.3: Comparison of Suzaku J1740.5-3014, AX J1740.5–3014 and SAX J1740.5–3013.
ObjectAX J1740.5–3014
SAX J1740.5–3013
Suzaku J1740.5−3014

Flux∗
2.5
2.4
2.1

Γ
3.4+7
−3.7
2.4+1.5−2.1
0.8+0.1−0.1

NH (1022 cm−2 )10 (< 40)
1.5+0.5−0.7
1.5+0.3−0.2

Obs. date
Sep. 95
Apr. 98
Oct. 08 + Feb. 09

Ref.
[110]
[31]
This work∗ The unit is ergs s−1 cm−2 . The interstellar absorption is not corrected.

6.3.2

Suzaku J1740.5-3014 as a New Intermediate Polar Candidate

The X-ray spectra of Suzaku J1740.5−3014 exhibit three lines at 6.4, 6.7, and 7.0 keV. The
best-ﬁt EW values of the lines are ∼210, ∼190, and ∼130 eV, respectively. Ezuka and Ishida[20] reported that mCVs have three iron Kα lines at 6.4, 6.7, and 7.0 keV with mean EW valuesof ∼100, ∼ 200, and ∼100 eV, which are nearly the same as those of Suzaku J1740.5−3014.
We found the coherent and clear pulsation of 432.1 s from Suzaku J1740.5−3014. Thispulse period is typical for mCVs[106]. Although we found no other periodicity due to orbital
modulation, this 432.1 s period would be due to spin, because possible orbital modulation islarger than this value. Furthermore, the pulse proﬁle is energy dependent, which may not be
appeared if the pulse is due to orbital modulation.
The overall spectrum was well ﬁtted with the two-component CIE plasma model with temperatures of kT = 6.0 and 64 keV. The spectrum can be ﬁtted with the one-component CIE
plasma model with partial absorption marginally. Both models require the Kα and Kβ lines fromneutral iron. Ezuka and Ishida [20] also reported that the spectra of mCVs can be described bya thin thermal plasma model with a mean temperature of ∼20 keV plus the 6.4 keV line. The
one-component CIE plasma model with partial absorption gave a temperature of kT = 10 keV,which resembles to the general feature of mCVs. We, however, favor the two-component CIE
plasma model, because the χ2 /d.o.f. value is better and the energy dependent pulse proﬁle isdiﬃcult to be explained with the one-component CIE plasma model.
In the two-component CIE model, the absorption for the low-temperature plasma is NH =(1.6 ± 0.1) × 1022 cm−2 , which may be due to the interstellar absorption. Assuming the uniform
density toward the GC, we can estimate the source distance to be ∼2 kpc. Here the absorption
and distance to the GC are thought to be NH = 6 × 1022 cm−2 [96, 110] and 8 kpc (e.g. [66]). Then the absorption-corrected ﬂux of 2.1 × 10−12 erg cm−2 s−1 in the 2−10 keV band isconverted to the source luminosity of 1 × 1033 erg s−1 in the 2–10 keV band. This luminosity is
typical for an intermediate polar (IP) of ∼ 1031 –1033 erg s−1 [20].
All these facts favor an idea that Suzaku J1740.5−3014 is a mCV, especially an IP, not apolar, because the 432.1 s period is too short as a synchronized orbital-spin period (≥ 4000 s for
all catalogued polars by Ritter & Kolb [106]). In ﬁgure 6.4, a sub peak appear to be present at
the higher-frequency side of the main peak. The frequency of the sub peak is ∼ 2.48 × 10−3 Hz.
It might be a beat of the spin and orbital periods observed in some IPs (e.g. [88]). The frequencyof the beat is Ω ± ω; Ω and ω are the frequency of the orbital and spin motions. If the sub peakis a beat Ω + ω, ω is 2.31 × 10−3 Hz and thus Ω is ∼ 1.7 × 10−4 Hz or the period is 5.9 × 103 s(∼ 1.6 hours). The orbital periods of the IPs catalogued by [106] are 5.1 × 103 ∼ 5.6 × 104 s andthus the period of 5.9 × 103 s can be the orbital period. The period is , however, also very close

6.3. SHORT DISCUSSION FOR SUZAKU J1740.5-3014

89

to the orbital period of the Suzaku satellite so there is a possibility that it is an artiﬁcial one.
From the absorption of the low-temperature component, Suzaku J1740.5−3014 would be aforeground source, not a member of the GC sources.
We suggest that a large absorption for the 64 keV plasma is due to the circum-stellar gas,
which could be up to several times 1023 cm−2 [20]. This large amount of circum-stellar gas cannaturally explain the origin of the strong 6.4 keV line from neutral iron (see e.g. [20]).
The spectrum of the standard model with a cylindrical emission region is given by multitemperature plasma components with a power-law emission measure distribution of index ∼-0.5[32]. The multi-component plasma can be approximated by two representative temperatures.
The ratio of the volume emission measures of the 6.0 keV plasma to the 64 keV plasma is 1:2.6.
This ratio does not follow the power-law relation of ∼-0.5. It is also strange that the 6.0 keV
plasma emission does not suﬀer the large absorption, if that is attributable to the circum-stellar
gas. We thus suspect a standard emission model and geometry should be modiﬁed for this new
IP. Detailed modeling, however, is beyond the scope of this thesis.

Chapter 7

Discussion
In this chapter, we make discussions based on the results in the previous chapters. In section
7.1 and 7.2, the spatial distribution and the spectra of the GDXE are discussed. In section 7.3,
the intensity of the GDXE is compared with that of point sources detected with Chandra. In
section 7.4 and section 7.5, Sgr A East and Suzaku J1740.5-3014 are compared with other SNRs
and CVs, and their contribution to the GDXE is discussed.

7.1

Spatial Distributions of the Fe Kα Line Emissions of the
GDXE

7.1.1

Uncertainty of the Intensity Measurement

Here we evaluate the uncertainties of the intensity measurement. The eﬀect of the interstellar
absorption is small in the band of the Fe Kα lines because the photoelectric absorption decreases
with the incident energy to -2.7 power, and the energies of the Fe Kα lines are just below the Fe
K edge. Only 9% of the incident photons with the energy of 6.5 keV are absorbed by interstellar
medium of NH = 6 × 1022 cm−2 . The uncertainty of the interstellar absorption on the Fe Kαline proﬁles is this level and does not aﬀect the following discussion. In addition, the Fe Kα line
emissions are little contaminated by bright point sources and thus a good tracer for the GDXE.
Sgr A East is, however, bright in the Fe XXV Kα line. The region of Sgr A East was excluded
but the spectra of the GDXE near it are contaminated. We estimated the contamination using
the XIS Monte Carlo simulator xissim [34]. We inputted the ﬂux of the Fe XXV Kα line
emission obtained in chapter 5 and calculated the leak from Sgr A East to the neighbor region.
As a result, less than 15% of the Fe XXV Kα line emission in l∗ < 0◦ .15 comes from Sgr A
East. In the outer region (l∗ > 0◦ .15), the contamination from Sgr A East is less than 1%. The
contamination from Sgr A East is ignorable in the following discussion.

7.1.2

Comparisons with Stellar Mass Distribution Observed in the Infrared

We compared the intensity proﬁle of the GDXE with the stellar mass distribution model of the
Milky Way galaxy, originally made by Muno et al. [81], based on infrared observations of IRAS
,COBE [60] and IRT of Spacelab 2 [42]. The major emission sources at near infrared (NIR)
wavelength (1 – 7µ m) are stars and thus NIR is a good tracer of the stellar mass. The stellar
mass model consists of the components of the nuclear bulge (NB), Galactic bulge (GB) and91

92

CHAPTER 7. DISCUSSION

Galactic disk (GD). The NB is composed of the nuclear stellar cluster and the nuclear stellar
disk. The further detail of the stellar mass model is shown in Muno et al. [81]. This model has
the uncertainty of about 25% in the 1σ conﬁdence level. The proﬁles of the stellar mass model
are shown in ﬁgure 7.1. We show only the projections along l = −0◦ .056 and b = −0◦ .046 but
this model is actually a two-dimensional model.

-2

Mass Surface Density (MSolar arcmin )Profile along the longitude (b=-0°.046)
10

7

10

6

Stellar Mass Model
Nuclear Stellar Cluster
Nuclear Stellar Disk
Galactic Disk
Galactic Bulge

105

10

4

0.01
0.1
1
10
100
*
Distance from Sgr A along the longitude l* (degree)-2

Mass Surface Density (MSolar arcmin )Profile along the latitude (l=-0°.056)
10

7

106

10

5

10

4

Stellar Mass Model
Nuclear Stellar Cluster
Nuclear Stellar Disk
Galactic Disk
Galactic Bulge

0.01
0.1
1
*
Distance from Sgr A along the latitude b* (degree)

Figure 7.1: Stellar mass model along the galactic longitude (top, b = −0◦ .046) and latitude(bottom, l = −0◦ .056). This model is based on the infrared observations [60, 42] The detail of
this model is described in Muno et al. [81].
.
We ﬁtted the stellar mass model to the intensity proﬁles of the Fe Kα lines. The normalization of the model was a free parameter. The intensity proﬁles along the galactic longitude andlatitude were used for the ﬁttings simultaneously. The best-ﬁt result of the Fe XXV Kα line isshown in ﬁgure 7.2 and 7.3. Both of the proﬁles of the GDXE and the stellar mass similarly
decrease at l = 1◦ –2◦ which corresponds to the border of the NB and GD. The proﬁles of the
GDXE and the stellar mass are, however, clearly diﬀerent. The data points excess the best-ﬁt
model in the inner region (l∗ < 1◦ ) but fall below in the outer region (l∗ > 2◦ ) and the b = −2◦

7.1. SPATIAL DISTRIBUTIONS OF THE FE Kα LINE EMISSIONS OF THE GDXE

93

region.
°-s -1

Observed photons cm s arcmin-2

Fe XXV Kα b=-0.046 Fitted with full data
10-5

10

-6

10

-7

10

-8

West
East
Stellar Mass Model

0.1
1
10
Distance from Sgr A* along the longitude l*(degree)

Figure 7.2: The result of the ﬁtting the stellar mass model to the proﬁle of the Fe XXV Kα line
emission along the galactic longitude. The magenta lines show the 1σ uncertainty.

Revnivtsev et al. [102] reported that the Fe Kα line intensity proﬁle agree with the NIR
proﬁle well in l∗ = 1◦ –100◦ with the RXTE data. We also excluded the data points in the inner2◦ × 1◦ region and ﬁtted the stellar mass model to the Fe Kα line intensity proﬁles. The results
are shown in ﬁgure 7.4–7.7. The outer region (l∗ > 1◦ or b∗ > 0◦ .5) is ﬁtted well. On the other
hand, the Fe Kα line intensities excess the stellar mass model in the inner region.
Revnivtsev et al. [102] reported that, in l∗ = 1◦ –100◦ , the ratio of the Fe Kα line intensity to
the 4.9 µm NIR intensity is IFeKα [photon s−1 cm−2 deg−2 ] =4.7 ×10−5 I4.9µm [MJy str−1 ]. Notethat Revnivtsev et al. [102] did not resolve the three Fe Kα lines because they used a proportionalcounter (PCA) on the RXTE, whose energy resolution is ∼ 18 % at 6 keV. To compare theresult of Revnivtsev et al. with ours, we summed the intensities of the three Fe Kα lines andconvert our ﬁtting results in the outer region to IFeKα [photon s−1 cm−2 deg−2 ]=(3.6 ± 0.9)×10−5 I4.9µm [MJy str−1 ]. The error shows the uncertainty of the stellar mass model. The ratioof the Fe Kα line intensity to the NIR intensity we obtained is almost consistent with the resultof [102] in the uncertainty.
For the inner region (l∗ < 1◦ ), Revnivtsev et al.[102] did not obtain reliable results becauseof the low angular resolution (FWHM ∼ 1◦ ) of the RXTE satellite. We revealed that the inner
region, the intensity proﬁles of the Kα lines do not match the NIR proﬁle and excess it about
40 % in l∗ = 0◦ .1–1◦ .
We should note that this discussion depends on the uncertainty of the stellar mass model.
The stellar mass model is based on the result of COBE whose angular resolution is low (FWHM∼ 0.7◦ ) and thus has large uncertainty 1 . If we consider the 90 % uncertainty of the stellar massmodel, the data points and the model in ﬁgure 7.2 and 7.3 roughly match and thus the excess of
the GDXE in l∗ < 1◦ is not clear. A more accelerate stellar mass model is required to compare
The ground-observation data with the high resolution of less than 1′′ is used to build the stellar mass modelof the inner 30 pc (R < 0◦ .2). See Launhardt et al. [60] in detail.
1

94

CHAPTER 7. DISCUSSION

!

!

-5

10

-6

10

-7

10

-8

Fe XXV K! l=0.5 Fitted with full data
Observed photons cm-s s-1 arcmin-2-s -1

Observed photons cm s arcmin-2

Fe XXV K! l=0.17 Fitted with full data
10

-2

-1.5

-1

-0.5

0

0.5

1

10

-5

10

-6

10

-7

10

-8

-2

-1.5

Galactic latitude (degree)!

Observed photons cm-s s-1 arcmin-2

-2
-s -1

Observed photons cm s arcmin

10-7

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

1

10-7

10-8

-2

-6

10

-7

10

-8

10

-5

10

-6

10

-7

10

-8

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

1

-1

-0.5

0

-2
-s -1

Observed photons cm s arcmin

Galactic latitude (degree)

0.5

0.5

1

-5

10

-6

10

-7

10

-8

10

-5

10

-6

10

-7

10

-8

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

1

Fe XXV K! l=-1.16! Fitted with full data
Observed photons cm-s s-1 arcmin-2

-1.5

-1
-0.5
0
Galactic latitude (degree)!

10

Fe XXV K! l=-0.84! Fitted with full data

-2

-1.5

Fe XXV K! l=-0.56 Fitted with full data
Observed photons cm-s s-1 arcmin-2

-2
-s -1

Observed photons cm s arcmin

10

1

10-6

!

-5

0.5

10-5

Fe XXV K! l=-0.17 Fitted with full data
100

Fe XXV K! l=1.17 Fitted with full data

10-6

-2

-0.5

!

Fe XXV K! l=0.82 Fitted with full data
10-5

10-8

-1

Galactic latitude (degree)

1

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

Figure 7.3: The result of the ﬁtting the stellar mass model to the proﬁle of the Fe XXV Kα line
emission along the latitudes. The magenta lines show the 1σ uncertainty.

7.2. SPECTRA OF THE GDXE

95

with our results.

7.1.3

Comparisons with Distribution of X-ray Point Sources Observed with
ChandraWe compare the intensity proﬁle of the Fe XXV Kα line with that of X-ray point sources
obtained with Chandra. Muno et al. made point source catalog with the Chandra 2 Ms data of
the 2◦ × 0◦ .8 GC region. It contains 9017 point sources with the photometry data (e.g. hardness
ratio, photon ﬂux in the 0.5–2 and 2–8 keV bands etc.) [82]. Using the Chandra catalog, weconstructed the number distribution of the point sources in the regions shown in ﬁgure 7.8.
To preserve completeness limit, we selected the point sources brighter than 2 × 10−6 photons
cm−2 s−1 (0.5–8 keV). The number distributions of the point sources are shown in ﬁgure 7.9.
The intensity proﬁle of the Fe XXV Kα line in the arbitrary unit are plotted in ﬁgure 7.9 forcomparison. Koyama et al. reported that the intensity proﬁle of the Fe XXV Kα line at the Sgr
A region (l∗ < 0◦ .3, b∗ < 0◦ .05) has a gradual increase toward Sgr A∗ compared with the sharp
proﬁle of point sources [50]. In the region of l∗ < 0◦ .8, b∗ < 0◦ .4, however, we do not found clear
diﬀerence between the proﬁles of the Fe XXV Kα line intensity and point sources.

7.2

Spectra of the GDXE

Here we discuss the variation of the GDXE spectra in −3◦ < l < 2◦ . We concentrate on the line
intensity ratio of Fe XXVI Kα to Fe XXV Kα and equivalent widths. They are little aﬀected by
the uncertainty of the interstellar absorption and the systematic errors of the detector because
they are relative relations of the photon intensities in the close energy bands.

7.2.1

Line Intensity Ratios of Fe XXVI Kα to Fe XXV KαKoyama et al. [51] and Yamauchi et al. [143] reported that the typical line intensity ratios of Fe
XXVI Kα to Fe XXV Kα in the Sgr A region (l∗ < 0.3◦ ) and the ridge (l∗ > 8◦ ) are 0.33 ± 0.02
and 0.22 ± 0.06, respectively. We revealed the line intensity ratio in the transition region. The
proﬁle of the line intensity ratio along the galactic longitude is shown in the top panel of ﬁgure
4.25. The line intensity ratio decrease along with l∗ signiﬁcantly. The hypothesis of the constantratio is rejected with the 99.5 % conﬁdence χ2 -test.
The ratio at l∗ = 3◦ is almost same as that in the ridge. The transition of the ratio happensin l∗ = 1◦ –2◦ which corresponds to the border between the NB and the GD as shown in the left
panel of ﬁgure 7.1. Figure 4.18 shows that the e-falling scale of the intensity proﬁle of the Fe
XXV Kα line also changes around l∗ = 1◦ –2◦ . These facts suggest that the GDXE contains twocomponents which accompany the NB and the GD, respectively. The simplest interpretation of
the variation of the line intensity ratio is that the temperature of the plasma emitting the GDXE
decreases from kT =7 keV to 6 keV along l∗ . If there are the two components of the GDXE,
the NB component is hotter than the GD component. If the GDXE is truly diﬀuse plasma, theresult show the spatial diﬀerence of the plasma temperature. If the origin of the GDXE is theintegration of the point sources, the natures of the faint X-ray point sources might be diﬀerent
be the GC (or the NB) and the ridge (or the GD).
The scale heights of the line intensity proﬁles at l = ±0◦ .17 (ﬁgure 4.19) of Fe XXV Kαchanges around b∗ = 0◦ .5–1◦ which corresponds to the border between the NB, and the GD and

96

CHAPTER 7. DISCUSSION-s -1

Observed photons cm s arcmin

-2

°
°
°
Fe I Kα b=-0.046 Fitted with l*>1 or b*>0 .5

10-5

10

West
East
Stellar Mass Model

-6

10-7

10

-8

10

-5

10

-6

0.1
1
10
Distance from Sgr A* along the longitude l*(degree)-s -1

Observed photons cm s arcmin

-2

°
°
°
Fe XXV Kα b=-0.046 Fitted with l*>1 or b*>0 .5

West
East
Stellar Mass Model

10-7

10

-8

10

-5

10

-6

0.1
1
10
Distance from Sgr A* along the longitude l*(degree)-s -1

Observed photons cm s arcmin

-2

°
°
°
Fe XXVI Kα b=-0.046 Fitted with l*>1 or b*>0 .5

West
East
Stellar Mass Model

10-7

10

-8

0.1
1
10
Distance from Sgr A* along the longitude l*(degree)

Figure 7.4: The results of the ﬁtting the stellar mass model to the intensity proﬁles of the FeKα lines along the galactic longitude. The data points in the inner 2◦ × 1◦ region were not used
for the ﬁttings. The magenta lines show the 1σ uncertainty. Top: the Fe I Kα line. Middle: theFe XXV Kα line. Bottom: the Fe XXVI Kα line.

7.2. SPECTRA OF THE GDXE

!

97

!

!

!

-5

10

-6

10

-7

10

-8

-2

-1.5

-1

-0.5

0

0.5

1

10

-5

10

-6

10

-7

10

-8

-2

-1.5

Galactic latitude (degree)

!

!

!

Observed photons cm-s s-1 arcmin-2

-2
-s -1

Observed photons cm s arcmin

10-7

-1
-0.5
0
Galactic latitude (degree)

!

0.5

!

1

10

-6

10

-7

10

-8

10

-5

10

-6

10

-7

10

-8

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

-1

-0.5

0

-2
-s -1

Observed photons cm s arcmin

Galactic latitude (degree)

0.5

!

!

10-8

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

!

0.5

!

1

!

Fe I K! l=-0.56 Fitted with l*>1 or b*>0 .5

1

10

-5

10

-6

10

-7

10

-8

10

-5

10

-6

10

-7

10

-8

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

1

Fe I K! l=-1.16! Fitted with l*>1! or b*>0!.5
Observed photons cm-s s-1 arcmin-2

-1.5

1

10-7

Fe I K! l=-0.84! Fitted with l*>1! or b*>0!.5

-2

0.5

10-6

!

Observed photons cm-s s-1 arcmin-2

-2
-s -1

Observed photons cm s arcmin

-5

0

10-5

Fe I K! l=-0.17 Fitted with l*>1 or b*>0 .5
10

-0.5

Fe I K! l=1.17 Fitted with l*>1 or b*>0 .5

10-6

-1.5

-1

!

Fe I K! l=0.82 Fitted with l*>1 or b*>0 .5-2!

Galactic latitude (degree)

10-5

10-8

!

Fe I K! l=0.5 Fitted with l*>1 or b*>0 .5
Observed photons cm-s s-1 arcmin-2-s -1

Observed photons cm s arcmin-2

Fe I K! l=0.17 Fitted with l*>1 or b*>0 .5
10

1

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

Figure 7.5: The result of the ﬁtting the stellar mass model to the proﬁle of the Fe I Kα line
emission along the galactic latitude. The data points in the inner 2◦ × 1◦ region were not used
for the ﬁtting. The magenta lines show the 1σ uncertainty.

98

CHAPTER 7. DISCUSSION

!

!

!

!

-5

10

-6

10

-7

10

-8

-2

-1.5

-1

-0.5

0

0.5

1

10

-5

10

-6

10

-7

10

-8

-2

-1.5

Galactic latitude (degree)

!

!

!

Observed photons cm-s s-1 arcmin-2

-2
-s -1

Observed photons cm s arcmin

10-7

-1
-0.5
0
Galactic latitude (degree)

!

0.5

!

1

10

-6

10

-7

10

-8

10

-5

10

-6

10

-7

10

-8

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

-1

-0.5

0

-2
-s -1

Observed photons cm s arcmin

Galactic latitude (degree)

0.5

!

!

10-8

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

!

0.5

!

1

!

Fe XXV K! l=-0.56 Fitted with l*>1 or b*>0 .5

1

10

-5

10

-6

10

-7

10

-8

10

-5

10

-6

10

-7

10

-8

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

1

Fe XXV K! l=-1.16! Fitted with l*>1! or b*>0!.5
Observed photons cm-s s-1 arcmin-2

-1.5

1

10-7

Fe XXV K! l=-0.84! Fitted with l*>1! or b*>0!.5

-2

0.5

10-6

!

Observed photons cm-s s-1 arcmin-2

-2
-s -1

Observed photons cm s arcmin

-5

0

10-5

Fe XXV K! l=-0.17 Fitted with l*>1 or b*>0 .5
10

-0.5

Fe XXV K! l=1.17 Fitted with l*>1 or b*>0 .5

10-6

-1.5

-1

!

Fe XXV K! l=0.82 Fitted with l*>1 or b*>0 .5-2!

Galactic latitude (degree)

10-5

10-8

!

Fe XXV K! l=0.5 Fitted with l*>1 or b*>0 .5
Observed photons cm-s s-1 arcmin-2-s -1

Observed photons cm s arcmin-2

Fe XXV K! l=0.17 Fitted with l*>1 or b*>0 .5
10

1

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

Figure 7.6: The result of the ﬁtting the stellar mass model to the proﬁle of the Fe XXV Kα line
emission along the galactic latitude. The data points in the inner 2◦ × 1◦ region were not used
for the ﬁtting. The magenta lines show the 1σ uncertainty.

7.2. SPECTRA OF THE GDXE

!

99

!

!

!

-5

10

-6

10

-7

10

-8

-2

-1.5

-1

-0.5

0

0.5

1

10

-5

10

-6

10

-7

10

-8

-2

-1.5

Galactic latitude (degree)

!

!

!

Observed photons cm-s s-1 arcmin-2

-2
-s -1

Observed photons cm s arcmin

10-7

-1
-0.5
0
Galactic latitude (degree)

!

0.5

!

1

10

-6

10

-7

10

-8

10

-5

10

-6

10

-7

10

-8

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

-1

-0.5

0

-2
-s -1

Observed photons cm s arcmin

Galactic latitude (degree)

0.5

!

!

10-8

-1.5

-1
-0.5
0
Galactic latitude (degree)

!

1

1

-2

0.5

!

1

!

Fe XXVI K! l=-0.56 Fitted with l*>1 or b*>0 .5
10

-5

10

-6

10

-7

10

-8

10

-5

10

-6

10

-7

10

-8

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

1

Fe XXVI K! l=-1.16! Fitted with l*>1! or b*>0!.5
Observed photons cm-s s-1 arcmin-2

-1.5

1

10-7

Fe XXVI K! l=-0.84! Fitted with l*>1! or b*>0!.5

-2

0.5

10-6

!

Observed photons cm-s s-1 arcmin-2

-2
-s -1

Observed photons cm s arcmin

-5

0

10-5

Fe XXVI K! l=-0.17 Fitted with l*>1 or b*>0 .5
10

-0.5

Fe XXVI K! l=1.17 Fitted with l*>1 or b*>0 .5

10-6

-1.5

-1

!

Fe XXVI K! l=0.82 Fitted with l*>1 or b*>0 .5-2!

Galactic latitude (degree)

10-5

10-8

!

Fe XXVI K! l=0.5 Fitted with l*>1 or b*>0 .5
Observed photons cm-s s-1 arcmin-2-s -1

Observed photons cm s arcmin-2

Fe XXVI K! l=0.17 Fitted with l*>1 or b*>0 .5
10

-2

-1.5

-1

-0.5

0

0.5

1

Galactic latitude (degree)

Figure 7.7: The result of the ﬁtting the stellar mass model to the proﬁle of the Fe XXVI Kαline emission along the galactic latitude. The data points in the inner 2◦ × 1◦ region were notused for the ﬁtting. The magenta lines show the 1σ uncertainty.

100

CHAPTER 7. DISCUSSION

Figure 7.8: The regions where we constructed the number distribution of the X-ray point sources
using the Chandra data. They are overlaid on the X-ray image in the 6.55–6.8 keV band obtainedwith the Suzaku XIS.

GB (the bottom panel of ﬁgure 7.1). It indicates that the GDXE might have a component whichaccompany the GB. On the other hand, the proﬁle of the line intensity ratio of Fe XXVI Kα to
Fe XXV Kα along the latitude (the top panel of ﬁgure 4.25) does not change over b∗ = 0◦ – 2◦
signiﬁcantly. The line intensity ratios of the GB and GC components might be similar but thecurrent oﬀ-plane data are insuﬃcient to clarify the diﬀerence between the line intensity ratios
between the GC and GB. More observations for the large b∗ regions with the long exposures will
reveal it. It is a future work.

7.2.2

Relation between the Equivalent Widths of the Fe I Kα and Fe XXV
Kα Line Emissions

Koyama et al. [55] analyzed the GDXE spectra in the Sgr A region (l∗ < 0◦ .3) and found theanti-correlation between the equivalent widths of the Fe I Kα and Fe XXV Kα lines. Using
this relation, the GDXE in the Sgr A region was decomposed phenomenologically into two
components; the Fe I Kα line emission with its associated continuum (6.4-component) and Fe
XXV Kα line emission with its associated continuum (6.7-component). The suzaku data of the
Sgr C region were analyzed with the same method by Nakajima et al. [84].
We used the same method to decompose the GDXE into the 6.4- and 6.7-components. The
relation between the equivalent widths of the Fe I Kα (EW6.4 ) and Fe XXV Kα (EW6.7 ) lines
are shown in ﬁgure 7.10. We found two groups in the on-plane data (the left panel of ﬁgure7.10). One group consists of region 2 and 3, which correspond to the giant molecular regions(Sgr B and Sgr C, respectively). This group has large equivalent widths of the Fe I Kα line.
The other group consists of region 1, 4 and 5, which are the regions outside the NB (l∗ >1◦ ). It has an clear anti-correlation between EW6.4 and EW6.7 . The best-ﬁt linear function
is EW6.7 + 0.83(±0.07) × EW6.4 = 0.61(±0.06) (keV). We also plot the average of the data in|l| > 8◦ obtained by Yamauchi et al. [143] in ﬁgure 7.10. The data point in |l| > 8◦ is locatedon the relation of region 1, 4 and 5. The equivalent width of the Fe I Kα line to its associated′ ), which is obtained as the EW
continuum (EW6.4
6.4 in the limit of EW6.7 → 0, is 0.74 ± 0.04′
keV. The EW6.4 is rid of the eﬀect of the continuum associated with the Fe XXV Kα line and7.2. SPECTRA OF THE GDXE

101
b=-0°.046
Point sources
Fe XXV Kα (arbitrary)

Sources arcmin

-2

1

0.1

0.8

0.6

0.4

0.2
0 -0.2 -0.4 -0.6 -0.8
Galactic Longitude (degree)-1l=0°.17
Point sources
Fe XXV Kα (arbitrary)
Sources arcmin

-2

1

0.1

-0.4

-0.2
0
0.2
Galactic Latitude (degree)

0.4

l=-0°.17
Point sources
Fe XXV Kα (arbitrary)
Sources arcmin

-2

1

0.1

-0.4

-0.2
0
0.2
Galactic Latitude (degree)

0.4

Figure 7.9: The number distributions of the X-ray point sources detected by Chandara. The
proﬁles were measured in the regons shown in ﬁgure 7.8. The proﬁles of the Fe XXV Kα line
emission in the arbitrary unit are also plotted for comparison.

102

CHAPTER 7. DISCUSSION′ , which is deﬁned asthus directly reﬂects the emission process of the Fe I Kα line. The EW6.7
the EW6.7 in the limit of EW6.7 → 0, is 0.60 ± 0.06 keV. The results of the Sgr A, C regions′ is similarly ∼ 0.6
and l∗ = 1◦ –3◦ (region 1, 4 and 5) are summarized in table 7.1. The EW6.7′ is 0.74 ± 0.04 keV
keV in the Sgr A, C regions and l∗ = 1◦ –3◦ . On the other hand, the EW6.4◦◦
in l∗ = 1 –3 , signiﬁcantly smaller than those in the Sgr A and C regions (1.2 and 2 keV).
This result indicates that the origins of the Fe I Kα are diﬀerent between the inner and outside
regions of the NB or the CMZ.
The relation between EW6.4 and EW6.7 of the oﬀ-plane data is shown in the right panel ofﬁgure 7.10. There is no clear relation. The near center and north east regions (region A, B, C
and F) seem to match the relation of the on-plane data (dotted line). The south west region(region D, E, H and b = −2◦ ), however, does not follow the relation. The equivalent widths ofregion D, E and H are close to those of the point sources in the Sgr A region (EW6.4 ∼ 140 eV,
EW6.7 ∼ 400 eV) [80]. The GDXE of these regions might be the integration of point sources.
The spectra of the point sources in these regions have not been studied well and might bediﬀerent from those in the Sgr A region.

600

600
2

500

400

EW6.4 (eV)

EW6.4 (eV)

500

3
300
1
200

A
B
C
D
E
F
G
H
°
b=-2

400
300
200

4
5

100
0
200

100

300

400 500
EW6.7 (eV)

600

700

0
200

300

400 500
EW6.7 (eV)

600

700

Figure 7.10: The relation between the equivalent widths of the Fe XXV Kα (EW6.7 ) and Fe I
Kα (EW6.4 ) lines. The error bars show the 90 % conﬁdence level. The dotted line shows thelinear relation determined by region 1, 4 and 5. The magenta and cyan lines show the relationsin the Sgr A and C regions [55, 84], respectively. Left: the on-plane data. The green data pointshows the averaged result in |l| > 8◦ [143]. Right: same but for the oﬀ-plane data.

7.2.3

Center Energy of the Fe XXV Kα Line Emission

We comment on the proﬁles of the center energies shown in ﬁgure 4.26. The center energy of Fe
XXV Kα becomes lower at large l∗ . The center energy of Fe XXV Kα in region 5 is 6664+6−4 eV
which is closer to that of the charge exchange origin than that of the CIE plasma origin [51].
The systematic errors are, however, large and thus the shift of the line center is not reliable. As
described in section 3.3.4, the energy scale calibration has large uncertainty due to the CTI. The7.3. COMPARISON BETWEEN THE INTENSITIES OF THE GDXE AND X-RAY POINT SOURCES D

Table 7.1: The equivalent widths of the Fe XXV Kα and I Kα lines to their associated continua.∗
RegionSgr A
Sgr Cl∗ = 1◦ –3◦

′ [EW
EW6.7
6.4 → 0] (eV)620 ± 70530 ± 60600 ± 60

′ [EW
EW6.4
6.7 → 0] (eV)1200 ± 200
2400+2300−700
740 ± 40

Ref.
[55]
[84]
This work∗ The uncertainties show the 90 % conﬁdence level.

distinction of the Fe XXV Kα line centers between the charge exchange and the CIE plasma
comes from the diﬀerence of the line ﬂux ratio of the resonance, intercombination and forbiddenlines. The Soft X-ray Spectrometer (SXS) onboard the ASTRO-H satellite will resolved theselines with the energy resolution of ≤ 7 eV and reveal whether this shift of the line center energy
we observed is true or not.

7.3

Comparison between the Intensities of the GDXE and X-ray
Point Sources Detected with Chandra

We compare the intensities of the GDXE with that of integrated point sources detected with
Chandra [82]. Muno et al. [78] and Revnivtsev et al. [105] made similar comparisons for the
Sgr A region. We further analyzed region 2, 3, and A, which the catalog [82] covers. In the
Chandra catalog [82], photon ﬂux in the 2–8 keV band and hardness ratio of each point source
are listed. Assuming that power-low spectrum for the point sources, we converted the photon
ﬂux and the hardness ratio given in [82] into energy ﬂux in the 5–10 keV band. The results are
summarized in ﬁgure 7.11 and table 7.2. Comparisons with the integrated intensity of the point
sources Ipointsources and the observed GRXE intensity IGDXE show that 4–15% of the GDXE isresolved into point sources.
The log N-log S distributions in the Sgr A region was studied by Muno et al [82]. The number
density per solid angle of point sources brighter than the ﬂux F is proportional to F −α whichis ∼ 1.6 for the Sgr A region. Assuming this relation, we extrapolated from the sensitivity limitof the current observation Flimit and estimated the necessary limit ﬂux (F100%GDXE ) to resolve
all of the GDXE into point sources (ﬁgure 7.11). The results are showed in table 7.2.
The necessary limit ﬂuxes are 5 × 10−18 – 2 × 10−16 erg s−1 cm−2 , 2–3 order smaller thanthe current limits. The region 2 and 3 corresponds to the Sgr B and C regions which containthe X-ray reﬂection nebula and the strong 6.4-component emission. This might be the reason
why the F100%GDXE of region 2 and 3 is much smaller than that of region A.

7.4

Properties of Sgr A East and Other GC SNR and Their
Contribution to the GDXE

In this thesis, we successfully determined the temperature of Sgr A East is free from the nonthermal component. Sgr A East has a peculiarly hot component (kT ∼ 6 keV) as a SNR, but
its line intensity ratio of Fe XXV Kα to Fe XXVI Kα is 0.05, and smaller than that (0.2–0.4)
of the GDXE.

CHAPTER 7. DISCUSSION

Region A
10

-12

GDXE

-2

Iponitsources(>F) (erg s cm arcmin , 5-10 keV) Iponitsources(>F) (erg s cm arcmin , 5-10 keV) Iponitsources(>F) (erg s cm arcmin , 5-10 keV)

104

10-14
10-15

10

-16

10

-18

-17

10

10

-16

-1

-15

10
-2

10

-14

-13

10

-12

-13

10

-12

-13

10

-12

10

-2

F (erg s cm arcmin , 5-10 keV)
Region 2
10

-12

GDXE

-2

-1

-2

10-13

10-14
10-15

10

-16

10

-18

-17

10

10

-16

-1

-15

10
-2

10

-14

10

-2

F (erg s cm arcmin , 5-10 keV)
Region 3
10

-12

GDXE

-2

-1

-2

10-13

-1

-2

10-13

10

-14

10-15

10

-16

10

-18

-17

10

10

-16

-1

-15

10
-2

10

-14

10

-2

F (erg s cm arcmin , 5-10 keV)

Figure 7.11: Comparison between the intensities of the GDXE and integrated point sources
detected with Chandra. The horizontal and vertical axes show the ﬂux of point sources F
and the integrated intensity of the point sources brighter than F . The green line shows theintensity of the GDXE. The dotted line shows the α = 1.6 relation extrapolated from (Flimit ),Ipointsources (> Flimit ).

7.5. PROPERTIES OF SUZAKU J1740.5-3014 AND OTHER POINT SOURCES AND THEIR CONTRIB

Table 7.2: Comparison between the Intensities of the GDXE and integrated point sources detected with Chandra.
RegionA
2
3

IGDXE ∗ (10−14 )
8.7
13.9
16.1

Ipointsources ∗ (10−14 )
1.3
0.56
0.75

Flimit †1.6 × 10−14
6.5 × 10−15
1.2 × 10−14

F100%GDXE †2.4 × 10−16
4.5 × 10−18
1.3 × 10−17∗ The unit is observed erg s−1 cm−2 arcmin−2 (5–10 keV). † The unit is observed erg s−1 cm−2

(5–10 keV).
Table 7.3: Properties of supernova remnants in the GC region.

Sgr A East
G359.79-0.26
G359.77-0.09
G359.41-0.21
G0.61+0.01
G1.2-0.0
G1.02-0.18

kT (keV)1.21±0.01, 6.0+0.5−0.6
1.02+0.03−0.04
0.66+0.08−0.09
0.90+0.2−0.3
3.2+2.3−0.9
0.9 ± 0.2
1.9+0.5−0.3

LX (erg s−1 )4.2×1035
3.6×1035
1.1×1035
2.2×1034
1.5×1034
1.4×1035
4.2×1033

Thermal energy (erg)
2.8×1050
1.7×1050
1.9×1050
3.9×1049
2.4×1049< 1051
3.8×1049

Ref.
This work
[74]
[75]
[133]
[53]
[113]
[86]

The properties of other SNRs in the GC region are summarized in table 7.3. Their temperatures (kT ∼ 1–3 keV) are also too low to explain the line intensity ratio of Fe XXV Kα to Fe
XXVI Kα of the GDXE. It is impossible to explain the spectra of the GDXE only by the GC
SNRs.

7.5

Properties of Suzaku J1740.5-3014 and Other Point Sources
and Their Contribution to the GDXE

We studied Suzaku J1740.5-3014 as a candidate source in the intermediate region between the
GC and the ridge but it is likely to be in the foreground region (distance ∼ 2 kpc). The propertiesof Suzaku J1740.5-3014, known CVs in the vicinity of the Sun, and the point sources in the Sgr
A region are summarized in table 7.4.
The integrated spectrum of the point sources in the Sgr A region (hereafter “the Sgr A point
sources”) has about 2–3 times larger equivalent widths of the Fe Kα lines than those of known
CVs. On the other hand, It indicates that the metal abundance of the Sgr A point sources might
be systematically large. The mean metal abundance of known CVs is 0.3 solar [20] and thus, if
the most of the Sgr A point sources are CVs as shown by Muno et al. [80], the metal abundanceshould be ∼ 1 solar. Such a high-abundance CVs might belong to the NB, and ordinary lowabundance CVs might belong to the Galactic bulge and ridge. Systematic observation of X-ray
point sources in the intermediate region between the GC and the ridge will reveal that because
this region is the transition point. It is one of the future works.

106

CHAPTER 7. DISCUSSION

Table 7.4: Properties of Suzaku J1740.5-3014, known CVs in the vicinity of the Sun, and the
GC point sources∗ .

Suzaku J1740.5-3014
Average of non-mCVs
Average of polar CVs
Average of IP CVs
Sgr A point sources†

EW6.4 (eV)210+30
−50
∼ 50
∼ 120
∼ 130
137 ± 21

EW6.7 (eV)190+60
−40
∼ 130
∼ 160
∼ 120
404 ± 59

I6.97 /I6.7
0.68 ± 0.30
∼ 0.48∼ 0.67∼ 0.580.56 ± 0.13

LX (erg s−1 )1 × 1033
1029.5 –1032.5
1030 –1033
1031 –1033–

Ref.
This work[99, 80][20][20, 80][80]∗ The error shows 90% conﬁdence level. † The values of faint sources (C < 500) in table 4 of [80].

The line intensity ratio of Fe XXVI Kα to Fe XXV Kα is not diﬀerent among the Sgr A point
sources and known CVs (non-magnetic CVs, IPs, polars) in the vicinity of the Sun. The line
intensity ratio of Sgr A point sources and known CVs is too high (∼ 0.6) for that of the GDXE(∼ 0.2–0.4). The characteristics of the Sgr A point sources and known CVs are the strong Fe
XXVI Kα line emission. On the other hand, the known Galactic diﬀuse objects do not emit thestrong Fe XXVI Kα line emission. Most of the Fe XXVI Kα line emission of the GDXE mightcome from point point sources.

Chapter 8

Conclusion
With the Suzaku satellite, we performed a study of the GDXE in the region from the GC to theridge and the bulge (−3◦ < l < 2◦ , −2◦ < b < 1◦ ). We also studied Sgr A East, the SNR locatedin the GC, and a newly found X-ray point source, Suzaku J1740.5−3014, in the data set.
The intensity proﬁles of the highly ionized Fe Kα lines and the continuum in the GDXE
including the data points in the ridge (l∗ > 8◦ ) are described with phenomenological twoexponential-component models.
The e-folding scale of the intensity proﬁle along the longitude changes from ∼ 0◦ .5 to ∼ 30◦
around l∗ = 1◦ –2◦ . The line intensity ratio of Fe XXVI Kα to Fe XXV Kα decrease from ∼ 0.4
to ∼ 0.2 with the distance from the GC along the longitude. The line intensity ratio at l∗ = 3◦
is almost same as that in the ridge. The transition of the ratio also happens around l∗ = 1◦ –2◦ .
The region around l∗ = 1◦ –2◦ corresponds to the border between the NB and the GD. These
facts suggest that the GDXE contains two components which accompany the NB and the GD
respectively. The scale height of the intensity proﬁle along the latitude changes from ∼ 0◦ .2 to∼ 3◦ around b∗ = 0.◦ 5–1◦ . On the other hand, the line intensity ratio of Fe XXVI Kα to Fe
XXV Kα is not diﬀerent in b∗ = 0◦ –2◦ , signiﬁcantly.
The equivalent widths of the Fe I Kα and Fe XXV Kα lines are anti-correlated in the GD
regions (l∗ = 1◦ –3◦ , b∗ < 0◦ .1), From this relation, the equivalent widths of the Fe I Kα and′ and EW ′ are obtained. The EW ′ isFe XXV Kα lines to their associated continua, EW6.4
6.7
6.4
signiﬁcantly smaller than those of the previously reported results in the inner NB regions (Sgr′ is almost same among all the regions. It suggests that the origin
A and C) though the EW6.7
of the Fe I Kα line emission might be diﬀerent between the NB and GD.
The non-thermal component of Sgr A East was detected for the ﬁrst time. Sgr A East has
two-temperature plasma of kT ∼ 6 keV and ∼ 1.2 keV. The temperature of kT ∼ 6 keV ispeculiarly hot as a SNR and close to that of the GDXE. However, the total line intensity ratioof Fe XXVI Kα to Fe XXV Kα is 0.05, and too low to explain that (0.2–0.4) of the GDXE.
From Suzaku J1740.5−3014, the three Fe Kα lines and the coherent and clear pulsation of
432.1 s are detected. These features are typical for an IP. The interstellar absorption indicatesthat Suzaku J1740.5−3014 is located not in the GC region but in the distance of ∼ 2 kpc fromthe Sun. The luminosity of ∼ 1033 erg s−1 is also typical for an IP.
We compared the GDXE spectra with those of Sgr A East, Suzaku J1740.5−3014, otherCVs, other SNRs and point sources in the GC region. The spectra of these sources are diﬀerent
from that of the GDXE.
The results of this thesis are summarized in table 8.1 with bold type.
107

108

CHAPTER 8. CONCLUSION

Table 8.1: Same as table 2.3 but the results of this thesis are shown in bold type. The results
written in red (blue) support whether diﬀuse (point sources) origin.

Sgr A region|l|<0°.3

Between GC &Ridge |l|=0°.3–3°

Ridge|l|>8°

Comparison
between the
profiles of the
GDXE andpoint sources• The profile of
Fe XXV K! isdifferent fromthat of X-ray
point sources
detected with
Chandra.

•The profile of Fe
XXV K! in |l| <1° isdifferent from thatof the NIR but notdistinguished withthat of X-ray point
sources.

•The profile of the
GDXE in |l|>1°
matches that of the
NIR.

Equivalent
widths (EWs)of Fe K!

•The EWs of
the GDXE are
too large asknown CVs.
•The EWs of
the pointsources in the
GC are as large
as those of the
GDXE.

•The EWs of the
GDXE are too large
as known CVs.
!Suzaku J1740.5–3014 was studied
but it is aforeground source.
Its EWs are similar
to that of known
CVs.

!The EWs of the
GDXE are toolarge as known
CVs.
•The EWs ofknown CVs near
the Sun are too
small as the
GDXE.

EW of FeXXV K! v.s.
EW of Fe I K!

!They are anticorrelated.

!They are anti!Not studied.
correlated in |l|=1°–3°, |b| < 0°.1 but therelation is different
from those in |l| <1°.

Intensity ratioof Fe XXVI K! to Fe XXV
K!

!0.33
•The ratio of
Sgr A East is
0.05.

•It changes from
~0.3 to ~0.2 in theborder between the
NB and the GD.

!0.22

!0.42

Fraction of
the GDXE
resolved into
point sources.

•40%!4–15%.

•19%

•85%

Bulge
|b|>1°

Appendix A

X-ray Emission from Thin-Thermal
PlasmaThe spectrum of X-ray emission from thin thermal plasma composed of continuum and line
emission components. The continuum component is bremsstrahlung emission by thermal electrons. The power of the bremsstrahlung per unit time, volume and frequency for typical thermal
plasma emitting X-ray is written as
32πe6
dW
=
3/2
dtdV dν3hme c3(∑)

z az ne nH (hν)−1/2 exp(−hν/kTe ).
2

(A.1)

z

Here, kTe , me and ne are the temperature, mass and density of electron. nH is the hydrogendensity. z and az are the atomic number and the abundance relative to hydrogen of the element
Z [107]. The spectrum has a cutoﬀ at hν ∼ kTe so we can determine the plasma temperaturefrom the shape of the observed continuum. Combined with the temperature, we can calculate
ne nH from the luminosity. If a non-thermal component exists, the temperature determinationfrom the continuum has the ambiguity because the non-thermal component is often diﬃcult tobe distinguished with high-temperature bremsstrahlung.
The line emissions also indicate the plasma temperature. We review the line emission process
here. For low-density (ne < 1014 cm−3 ) plasma, the coronal approximation can be applied. In
this approximation, ions are essentially in their ground states. After collisional excitation, theion emits photons and becomes in the ground state before the next collision. Considering only
the dominant process of the collisional excitation, the volume line emissivity of a particular linetransition of a → b is written as
PZab+i = ne nH az ηZi SZga+i BZab+i .
∑

(A.2)

Here ηZi is the ionic fraction of Z+i ( zi=0 ηZi = 1) and a function of the plasma temperature.
i is shown in ﬁgure A.1 for example. We can inversely deﬁne the ionization temperature
ηFefrom the ionic fractions of given plasma. In CIE plasma, the ionization temperature equals
to the electron temperature. Plasma in which the ionization temperature is lower than theelectron temperature is often found in young supernova remnants (e.g. SN 1006 [141]), andcalled non-equilibrium ionization (NEI) plasma. SZga+i is the rate coeﬃcient for electron impact
excitation of ions Z+i from its ground state to its excited state a. SZga+i is calculated by averaging
the corresponding cross sections over the Maxwellian electron velocity distribution and thus a
109

110

APPENDIX A. X-RAY EMISSION FROM THIN-THERMAL PLASMA

function of the electron temperature. BZab+i is the radiative branching ratio of the transitiona → b among all possible transitions from level a.

Figure A.1: Iron ionic fractions at ionization equilibrium as a function of plasma temperature[2, 124].

0.9

0.09

0.8

0.085

0.7

[S XV Kβ]/[S XV Kα][Fe XXVI Kα]/[Fe XXV Kα]

The intensity ratio between the emission lines from the ions of the same element (e.g. Fe
XXVI Kα/Fe XXV Kα, Fe XXV Kβ/Fe XXV Kα, etc.) is written as a combination of ηZi , SZga+i
and BZab+i because ne , nH and az are canceled. ηZi and SZga+i are functions of the temperatureand BZab+i is constant. Thus, we can obtain the plasma temperature from the observed line
intensity ratio. We calculated the some examples of the relations between the line ratio and thetemperature shown in ﬁgure A.2 with CIE plasma model code APEC [119].

0.6
0.5
0.4
0.3
0.2

0.08
0.075
0.07
0.065
0.06
0.055

0.1

0.05

0

0.045

2

3

4

5
6
7
Temperature (keV)

8

9

10

0.6

0.8

1

1.2
1.4
Temperature(keV)

1.6

1.8

Figure A.2: Relations between the temperature and the line intensity ratio.
[Fe XXVI Kα]/[Fe XXV Kα]. Right: [S XV Kβ]/[S XV Kα].

Left:

The equivalent width is the ratio of the line ﬂux to the ﬂux per unit energy of the underlying
continuum at the line energy. Equations A.1 and A.2 show that the equivalent width is a functionof az and the temperature. If we obtained the plasma temperature from the line intensity ratioor the continuum, we can determine the metal abundance az from the equivalent width. We
show an example of the relation between the line equivalent width and the temperature in ﬁgure
A.3 assuming the solar abundance with the APEC code.
A He-like Kα emission line (e.g. Fe XXV Kα) has ﬁne structure of the resonance (w:1s2

111

Equivalent width of Fe XXV Kα (eV)

1800
1600
1400
1200
1000
800
600
400
200
0

2

4
6
Temperature (keV)

8

10

Figure A.3: Relations between the temperature and the equivalent width of the Fe XXV Kαline.
1 S –1s2p 1 P ), the intercombination (x+y:1s2 1 S –1s2p 3 P ), the forbidden (z:1s2 1 S –1s2s
0
1
0
2,1
0
3 S ) lines (ﬁgure A.4) and the satellite lines (s:Li-like 1s2 2p–1s2p2 etc.). The center energies are1higher in the order of the resonance, intercombination and forbidden lines. The center energiesof the satellite lines are lower than that of the resonance line. With the higher temperature, therelative intensity ratio of the resonance line to the sum of the intercombination and forbiddenlines become higher and the intensity of the satellite lines get lower. The blend of the resonance,intercombination, forbidden and satellite lines is observed as a single He-like Kα line with a
CCD camera, because its energy resolution (FWHM ∼ 150 eV) is not enough to resolve theﬁne structure. Thus, the center energy of an observed He-like Kα line indicates the plasma
temperature. The relation between the temperature and the center energy of the Fe XXV Kαline calculated with the APEC code is shown in ﬁgure A.5.
6690
Centroid of Fe XXV Kα (eV)

6680
6670
6660
6650
6640
6630
6620
6610
6600
1

1.5

2

2.5
3
3.5
Temperature(keV)

4

4.5

5

Figure A.4: Scheme of the ﬁrst energy levels Figure A.5: Relation between the temperatureand the center energy of the Fe XXV Kα line.
of He-like ions [124].

Appendix B

CTI Correction and In-Orbit
Calibration for SCI
The CTI correction method and in-orbit calibration of the SCI technique are described in thischapter. The SCI technique was in practical use for the ﬁrst time in the Suzaku XIS. It is the
key technique to decrease the systematic error of the long-period observation data used in thisthesis. These results are also reported in Uchiyama et al.[134].

B.1

Property of the CTI in the SCI Mode and the “Sawtooth”
CTI model

In the SCI mode of the Suzaku XIS, a charge is injected into every 54th row. The amounts of
the injected charge into each pixel are equivalent to the X-ray energy of ∼6 keV and of ∼2 keV,for the front-illuminated (FI) and back-illuminated CCD (BI), respectively.
We, especially Nakajima et al. [83] and Ozawa et al. [89], have revealed the properties of
the CTI in the SCI-oﬀ mode as below. A charge packet in the XIS loses its charge during (a)the fast transfer (24 µs pixel−1 ) in the imaging area, (b) the fast transfer in the frame-store
region, (c) the slow transfer (6.7 ms pixel−1 ) in the frame-store region, (d) the slow transfer
to the read-out node. The CTI depends on many parameters such as the transfer speed, thenumber density of the charge traps [26], pixel size, and the structure of transfer channel. The
imaging area and the frame-store region may have diﬀerent trap densities since the pixel size
are diﬀerent among these areas. Also the frame-store region is covered by the shield and is not
exposed to the radiation directly. The transfer (d) is examined by using the calibration source
data taken in 2007 April, and no signiﬁcant decrease of the pulse height along the ActX axis
is found. It shows that the charge loss in the transfer (d) is ignorable. We deﬁne that i is thetransfer number in the imaging area (i = ActY + 1; here, ActY is a coordinate value where
an incident X-ray generates a charge packet). Then the relation between P H ′ (i) and P Ho isexpressed as
P H ′ (i) = P Ho (1 − ca )i (1 − cb )1024−i (1 − cc )i ,∼ P Ho {1 − i(ca − cb + cc ) − 1024cb } ,(B.1)

where ca , cb and cc are the CTI values in the transfers (a), (b), and (c), respectively. Here the fact
113

114

APPENDIX B. CTI CORRECTION AND IN-ORBIT CALIBRATION FOR SCI

Notationi
P Ho
P H ′ (0)P H ′ (i)jc1 (i)
c1t , c1b
c2
−s(j)
−st
−sb
β

Table B.1: The list of notation.
Meaning
Transfer number.
Original pulse height.
Readout pulse height of a pixel at i = 0.
Readout pulse height of a pixel at i.
Row number between a pixel and its preceding CI row. j = i mod 54.
CTI 1 of equation B.2 for a pixel at i.
The c1 values at the peak and valley of the sawtooth (see ﬁgure 1a).
CTI 2 of equation B.2 for a pixel at i.
Slope for a pixel with j. It is equal to c1 (i) · P Ho .
Slope for the tops of the sawtooth. It is equal to c1t · P Ho .
Slope for the bottoms of the sawtooth. It is equal to c1b · P Ho .
CTI depends on the pulse height as c1,2 ∝ P Ho−β . See equation B.3.

that the CTI values are much smaller than 1 is used. Thus we can separate the charge loss into
i-dependent component (the second term in the right-hand side of equation B.1) and constantcomponent (the third term). We therefore substitute CTI to CTI 1 (the former component)
and CTI 2 (the latter component), which have the CTI values of c1 = ca − cb + cc and c2 = cb ,respectively . Then equation B.1 can be written as
P H ′ (i) ∼ P Ho (1 − ic1 − 1024c2 ).

(B.2)

Since the CTI values (c1 , c2 ) depend on the amount of transfer charge, we assume the CTI
dependence on the PH (proportional to the charge) as
c1 = k1 (P H)−βand,c2 = k2 (P H)−β ,(B.3)

where k1 and k2 are scale factors for the CTI 1 and CTI 2, and the index β is common to CTI 1
and CTI 2.
To make clear for further descriptions of the SCI, relevant terminologies and notations are
summarized in table B.1.
The CTI in the SCI mode is assumed to consist of two components as CTI 1 and CTI 2
following the result in SCI-oﬀ mode. The relation among readout pulse height (P H ′ ), original
pulse height generated by an X-ray of energy E (P Ho at E), CTI 1 (c1 ), CTI 2 (c2 ), and transfer
number i are formalized in equation B.2, where the detector coordinates are deﬁned as ActX
and ActY . The CCD consists of four segments (A, B, C, and D); each segment consists of 256
columns along ActX and has a dedicated read-out node.
In the SCI mode, a radiation-induced trap is ﬁlled with a probability p when a sacriﬁcial
charge passes through the trap. The ﬁlled trap does not capture a signal charge produced by an
X-ray. However, the electron once ﬁlling the trap will be re-emitted with the time scale τ . Thus
the probability of the trap holding the electron is proportional to p · exp(−t/τ ), where t is thetime elapsed since the sacriﬁcial charge passes the trap. Thus it is reasonable to assume that apixel which is j rows away from its preceding charge injected row has the c1 value proportional
to 1 − p · exp(−δt · j/τ ), where δt is time for one vertical transfer; δt · j represents a time lag

B.2. CTI MEASUREMENT IN ORBIT

(a)
c1b

c1

(b)
Pulse height (ch)Charge-injected rows

115

c1( i )
c1t

Charge-injected rows

PHo
due to CTI 2
PH!(0)
PH!(i)-stj-s( j ) slope(due to CTI 1)-sbj

54 rows

54 rowsii(transfer number)(transfer number)

Figure B.1: Our model of the “sawtooth” variation. We assume the c1 -i relation show in (a),and it generates the P H ′ -i relation shown in (b).
between the pixel and the charge injected row. We assume δt · j/τ is a small value, and hence
1 − p · exp(−δt · j/τ ) is a linear function of j approximately. The charge is injected into every
54th row, and hence j = i mod 54 1 . We, consequently, can model the c1 -i relation of the SCI
mode as,c1 (i) = c1t +c1b − c1t
· (i mod 54).
54

(B.4)

Equation B.4 is a periodic sawtooth function as demonstrated in ﬁgure B.1a. This c1 -i
relation generates the P H ′ -i relation in ﬁgure B.1b, and the shape is nicely reproduced by theground experiments using the heavily damaged CCD as demonstrated in ﬁgure 15 of [128].
The sawtooth distribution is also found in the XIS in-orbit. We measured the P H ′ -i relations
using the onboard calibration source data in October 2006 with an eﬀective exposure time of
1 Ms, and show the results in ﬁgures B.2a and b. Figure B.2 shows that the sawtooth of the
BI sensor is shallower than that of the FI sensors. The results based on the data in February
2008 are also shown in ﬁgures B.2c and d. Compared with October 2006, the sawtooth became
deeper. The pulse height just after the charge injected rows changed more in the case of the BI
sensor than the FI sensors. This diﬀerence might be resulted in part from the smaller amountof charge injected in the BI sensor.
Our sawtooth CTI model can represent the complicated relation between P H ′ and i, andmake it possible to convert P H ′ (i) to P Ho with only three parameters, c1t , c1b , and c2 . Our
goal is to convert P H ′ (i) to P Ho by deciding the three parameters.

B.2

CTI Measurement in Orbit

B.2.1

Calibration Data

To study the CTI in the SCI mode, we analyzed the data of the onboard calibration sources
55 Fe, the Perseus cluster of galaxies, and 1E 0102.2−7219, whose properties are summarized
below.
1

The modular arithmetic i mod 54 = j means that when i is divided by 54, it leaves j as the remainder,hence i − j is divisible by 54. For example, 111 mod 54 = 3. Strictly speaking, charges are injected in rows of
i = 54 · n − 1 (n = 0, · · · , 18) and 1023 for full window mode. Thus j should be (i + 1) mod 54. We nevertheless use
this simplified description to avoid unnecessary confusion. We treated j properly for actual CTI measurement.

116

APPENDIX B. CTI CORRECTION AND IN-ORBIT CALIBRATION FOR SCI(b) XIS 1(BI) Segment A Oct. 2006
1520

1630

1510
Pulse height (ch)Pulse height (ch)(a) XIS 0(FI) Segment A Oct. 2006
1640

1620
1610
1600

830

1490
1480

Readout
Corrected

1590
800

1500

860
890
920
i (transfer number)

950

Readout
Corrected
1470
800
830

980

1520

1630

1510

1620
1610
1600

830

980

1500
1490
1480

Readout
Corrected

1590
800

950

(d) XIS 1(BI) Segment A Feb. 2008

1640

Pulse height (ch)Pulse height (ch)(c) XIS 0(FI) Segment A Feb. 2008

860
890
920
i (transfer number)

860
890
920
i (transfer number)

950

980

Readout
Corrected
1470
800
830

860
890
920
i (transfer number)

950

980

Figure B.2: Pulse height of the MnI Kα line from the onboard calibration source as a functionof i: October 2006 (a and b) and February 2008 (c and d). We show the results of the segment
A in XIS 0, 1 as typical examples. X-ray events of grade 02346 were analyzed. Black and red
marks are data before and after our new CTI correction, respectively.
Onboard Calibration Source 55 Fe
The onboard calibration source 55 Fe produces the Mn I Kα line. The center energy is 5.895 keVand the intrinsic line width is negligible for our purposes. Thus the observed line width directly
reﬂects the energy resolution of the XIS. We used the archival trend data of the calibration
source distributed by the Data Archives and Transmission System (DARTS)2 of ISAS/JAXA.
The data were obtained between August 2006 and March 2008, and the total eﬀective exposure
is about 23.7 Ms.
The Perseus Cluster of Galaxies
The Perseus cluster of galaxies is one of the X-ray brightest clusters in the sky at redshift
z = 0.0176. It has the X-ray spectrum of thermal plasma with strong Kα lines of FeXXV. Its
radius is about 15′ [13]. Since the XIS combined with the X-Ray Telescope (XRT) onboard
Suzaku has a ﬁeld-of-view (FOV) of 18′ × 18′ on the sky, the X-ray emission from the cluster
ﬁlls the entire ﬁeld of view. The temperature of the plasma changes smoothly from kT ∼ 4 keV
to ∼ 7 keV toward the outer region. The line center energy of the FeXXV Kα triplet is almost
constant (∼6.56 keV at z = 0.0176) within the temperature range, and the line is suitable
to measure the CTI. However, we do not use this line to determine the absolute energy scalebecause of the uncertainty of the absolute energy of the line. We have observed the Perseus
cluster with the SCI mode every six months for calibration.
2

http://darts.isas.jaxa.jp/index.html.en

B.2. CTI MEASUREMENT IN ORBIT

Obs. ID

101012010
101012020
102011010
102012010
101005090
101005110
102001010
102002010
102003010
102004010
102005010
102022010
102006010

117

Table B.2: Log of Observations for Calibration.
Observation time (UT)
Exposure time
Start
End
[ks]
The Perseus cluster
2006/08/29 18:55:07 2006/09/02 01:54:19
50.0
2007/02/05 15:57:48 2007/02/06 14:30:14
43.9
2007/08/15 12:40:49 2007/08/16 11:27:22
42.3
2008/02/07 02:09:42 2008/02/08 10:30:19
41.6
1E 0102.2−7219
2006/12/13 18:53:16 2006/12/14 03:04:19
28.2
2007/02/10 22:13:47 2007/02/11 19:30:14
36.2
2007/04/10 10:35:08 2007/04/10 19:30:19
18.1
2007/06/13 10:10:12 2007/06/14 03:31:19
27.9
2007/08/12 05:21:09 2007/08/13 03:45:24
39.5
2007/10/25 12:24:45 2007/10/26 09:00:14
26.2
2007/12/01 19:25:40 2007/12/02 09:50:19
24.8
2008/02/14 16:57:28 2008/02/16 03:10:24
26.5
2008/03/15 05:43:27 2008/03/15 20:45:24
28.2

1E 0102.2−7219
This is one of the brightest supernova remnants (SNR) in the Small Magellanic Cloud. Chandra
resolved its ring-like structure of the 30′′ -diameter [22] though it is regarded as a point sourcewith the spatial resolution of Suzaku. The spectrum of 1E 0102.2−7219 has strong lines fromhighly-ionized O, Ne and Mg. These lines are resolved to resonance, forbidden and intercombination lines by the XMM-Newton RGS, and thus the accurate energies of the line centroids
are known [98]. 1E 0102.2−7219 has been used for calibrations in a soft-energy band by several
satellites, and an empirical model for calibrations has been constructed [94]. We have observed
1E 0102.2−7219 with the SCI mode regularly for the calibration of the low energy gain and thecontamination.
XIS 2 suddenly showed an anomaly on November 9, 2006, and it has not been operated sincethen. Thus only the data of XIS 0, 1 and 3 are studied.
All data were acquired with the normal clocking mode and the full window option using the
SCI. The editing mode was 3 × 3 or 5 × 5. [50] provide details of these modes.
As mentioned in [50], a small fraction of the charge in a pixel is left behind (trailed) to thenext pixel during the transfer. This charge-trail phenomenon changes the spatial extent of an
X-ray event. All data were corrected for the phenomenon based on the in-orbit data 3 , otherwisesome X-ray events would be judged as grade 7, and the detection eﬃciency would decrease.
The observations of the celestial objects are summarized in table B.2.
3

The details about the charge-trail correction is shown in
http://xmm2.esac.esa.int/external/xmm sw cal/icwg/
presentations/Suzaku XIS.pdf

118

B.2.2

APPENDIX B. CTI CORRECTION AND IN-ORBIT CALIBRATION FOR SCI

Determination of the CTI Parameters

The procedures for the CTI determination are as follows:• Step 1: deciding P Ho at 5.895 keV and 6.56 keV.
• Step 2: measuring c2 for P Ho at 6.56 keV.
• Step 3: measuring c1t and c1b for P Ho at 5.895 keV or 6.56 keV.
• Step 4: deciding the CTIs for any P H ′ values.
In the case of the SCI-oﬀ mode, we can measure the CTI of each column by the checker ﬂag
CI [83, 89]. Since the checker ﬂag CI is a complicated operation, we have not used this techniquein the SCI mode. We, therefore, measured the averaged CTI of each segment.
In the normal analysis of the XIS data, both of single-pixel and multi-pixel events (grade 0
and 2346 events; see [50]) are used. To determine the CTI, however, we analyzed only thegrade 0 events; if we use the multi-pixel events, it is diﬃcult to measure the CTI correctly,
because the CTI depends on the amount of transferred charge, and the amounts of charge ineach pixel comprising the multi-pixel event is diﬀerent with each other. We will mention how
to correct the data of grade 02346 events in step 4.

Step 1: P Ho at 5.895 keV and 6.56 keV
We determined the P Ho s of the FeXXV Kα line from the Perseus cluster (6.56 keV) by using
the data of the ﬁrst SCI observation on 29 August, 2006. Because it is impossible to measure
the c2 value in August 2006, we assumed that value is equal to zero. Even if this assumption
would not be reasonable, we can cover it by adjusting the P Ho -E relation. Then the P H ′ ati = 0 (P H ′ (0)) is equal to P Ho .
Since the statistics are limited, we divided each segment to four regions along the ActY axis,
and obtained the center pulse height of the iron line from each region. An example of P H ′ asa function of i is shown in ﬁgure B.3. By ﬁtting the data with a linear function of i, the P Hovalue was obtained as P H ′ (0)For the segments A and D, we obtained the P Ho value of the Mn I Kα line (5.895 keV)using the data of the calibration sources in August, 2006. We obtained the P H ′ -i relation, but
in this case, the i values are limited at around 900 because the calibration sources irradiate
only the two far-end corners from the read-out node of the imaging area [50]. We extrapolated
the relation to i = 0 by using a linear function, the slope of which was ﬁxed to a value sMncalculated as follows. Because the slope represents an averaged −c1 · P Ho , the slope should be
proportional to P Ho1−β . Then we calculated sMn from the slope for the 6.56 keV data (sFe ) as
sMn = sFe · (5.895/6.56)1−β = sFe · (0.90)1−β .

Step 2: c2 for P Ho at 6.56 keV
We obtained the P H ′ (0) values for all data of the Perseus cluster (table B.2) with the same
method of step 1. The P H ′ -i relations are shown in ﬁgure B.3. As is found in ﬁgure B.3,
P H ′ (0) is decreased with time; P H ′ (0) becomes diﬀerent from P Ho . We assumed the diﬀerence
P Ho − P H ′ (0) is attributable to c2 . Assuming that c2 is a linear function of time, the increasingrate of c2 for P Ho (at 6.56 keV) is typically ∼ 2×10−6 yr−1 for the FI sensors and ∼ 6×10−6 yr−1
for the BI sensor.

B.2. CTI MEASUREMENT IN ORBIT

119

PH′(i) (ch)

1680

1670

1660
Aug. 2006
Feb. 2007
Aug. 2007
Feb. 2008

1650
0

200

400
600
i (transfer number)

800

1000

Figure B.3: P H ′ as a function of i obtained using the FeXXV Kα line from the Perseus cluster.
The result of the XIS 1 segment C is shown as a typical example.

Step 3: c1t and c1b for P Ho at 5.895 keV and 6.56 keV
For the segments A and D, we determined st and sb for P Ho at 5.895 keV by using the dataof the calibration sources. For the segments B and C, we measured those for P Ho at 6.56 keV
by using the Perseus data. We ﬁtted the P H ′ -i relation with the sawtooth function shown inﬁgure B.1b, and determined st and sb . In the ﬁtting, we ﬁxed the P Ho and c2 to the values
obtained in step 1 and step 2. Typical examples of the ﬁtting are shown in ﬁgures B.4 and B.5.
Using st and sb , we calculated c1t and c1b .
In the case of the FI sensors, the c1t and c1b values for P Ho at 5.895 keV at August 2006
are typically < 10−7 and ∼ 3 × 10−6 , respectively. The increasing rates are ∼ 1 × 10−6 yr−1
and ∼ 5 × 10−6 yr−1 , respectively. In the case of the BI sensor, the parameters c1t and c1b at
August 2006 are ∼ 6 × 10−6 and ∼ 1 × 10−5 , and the increasing rates are ∼ 5 × 10−6 yr−1 and∼ 8 × 10−6 yr−1 , respectively. Typical examples of the time evolution of the c1 values for P Ho
at 5.895 keV are shown in ﬁgure B.6.
The CTI of the BI sensor is larger, and increases more rapidly than that of the FI sensors.
It might be the reason that the amount of the injected charge to the BI sensor is less than thatof FIs.

Step 4: CTI for any P H ′ values
For an event of P H ′ , we calculated the CTI with an equation of c1,2 × {P H ′ /(P Ho at 5.895 keV
or 6.56 keV)}−β . Here, we assumed the same P Ho -c1,2 relation as that of the SCI-oﬀ mode andused the same β (typically ∼ 0.25; [89]).
Since we obtained the CTI for all pulse-height values, we can now correct the CTI of multipixel events (grade 2346 events). By using the Mn I Kα line of the calibration sources or the
Fe XXV Kα line from the Perseus cluster, we examined the CTI correction for the grade 02346
events. Figure B.7 shows the line centroid of the grade 0 and grade 02346 events as a function oftime. While the line centroid of the grade 0 events is temporally constant, that of the grade 02346

120

APPENDIX B. CTI CORRECTION AND IN-ORBIT CALIBRATION FOR SCI

PH′(i) (ch)

1620

1610

1600
800

830

860
890
920
i (transfer number)

950

980

Figure B.4: Best-ﬁtting result of the sawtooth function for the MnI Kα line of the calibration
source. The result of the XIS 0 segment A in February 2008 is shown as a typical example.
Only grade 0 events were used.

PH′(i) (ch)

1820

1810

1800
200

300

400

500
600
700
i (transfer number)

800

900

Figure B.5: Best-ﬁtting result of the sawtooth function for the FeXXV Kα line from the Perseus
cluster. The result of the XIS 0 segment C in February 2008 is shown as a typical example.
The white and black circles show the pulse height of regions with smaller and larger j values,
respectively. Only grade 0 events were used.

B.3. ENERGY SCALE UNIFORMITY AND RESOLUTION IN THE SCI MODE

-5

2×10

121

c1b
c1t

c1

XIS 1(BI)10-5
XIS 0(FI)
0
400

500

600
700
800
900
Time since 2005-7-10 (day)1000

Figure B.6: Time history of c1t and c1b for P Ho at 5.895 keV. The segment A of XIS 0 and 1
are shown as typical examples.
events increases with time. This means that the CTI parameters for the grade 2346 events are
smaller than those for the grade 0 events. The CTI of the grade 2346 events can be smaller,
because the split charge acts as the sacriﬁcial charge.
We then ﬁne-tuned the parameters c1t , c1b , and c2 by multiplying a common time-independent
factor. We determined the factor of each segment so that the line center of the grade 02346
events becomes temporally constant. The factors are typically 0.9. Figure B.7 also shows thatthe result of the grade 02346 data corrected with the ﬁne-tuned CTI parameters. The adjusted
pulse height becomes constant. However, as a result, the value becomes diﬀerent from P Ho measured in step 1. We regarded this adjusted pulse height as P Ho at 5.895 keV, and determined
the P Ho -E relation.
To check the CTI correction in the low-energy band, we applied the ﬁne-tuned CTI parameters to the grade 02346 data of 1E 0102.2−7219. We ﬁtted the spectrum with the empiricalcalibration model4 . Figure B.8 shows the pulse height of the OVIII Kα line after the CTI
correction. We can see that the line centroid is constant, which supports the validity of our
method.

B.3

Energy Scale Uniformity and Resolution in the SCI Mode

The pulse height of the MnI Kα line after our new CTI correction is shown in ﬁgure B.2. The
sawtooth structure disappeared, and hence our new method greatly reduces the variation of thepulse height. Comparing between October 2006 and February 2008 shows that the CTI variation
with time is also corrected properly.
The P Ho -E relation for each segment is assumed to be the same function form with a similar
ﬁne-tuning process of the absolute energy as [89]: a model of two slopes crossing at the energy of
the Si-K edge (1.839 keV) with the same ratio of the two slopes as that obtained in the groundexperiments [50].
4See http://cxc.harvard.edu/acis/E0102/

APPENDIX B. CTI CORRECTION AND IN-ORBIT CALIBRATION FOR SCICorrected pulse height (ch)

122

1520

Grade 0
Grade 02346
Adjusted grade 02346

1515

1510
400

500

600
700
800
Time since 2005-7-10 (day)

900

1000

Figure B.7: Time history of the line centroid of MnI Kα line. We show the result of the
XIS 1 segment A as a typical example. The triangle and circle marks show the data of grade 0
and 02346 after the CTI correction with the parameters based on only the grade 0 events. The
rectangle marks show the data of grade 02346 events after the CTI correction with the ﬁne-tuned
parameters as described in the text.

Corrected pulse height (ch)

170

165

~ 4 eV

160
400

500

600
700
800
Time since 2005-7-10 (day)

900

1000

Figure B.8: Time history of the line centroid of OVIII Kα (0.653 keV). We show the result of
the XIS 1 segment C as a typical example. The results are based on the grade 02346 events
corrected with the ﬁne-tuned CTI parameters as described in the text.

B.3. ENERGY SCALE UNIFORMITY AND RESOLUTION IN THE SCI MODE

123

(a)

Mn I Kα FWHM (eV)

180

170

XIS 1(BI)
Oct.
2006

160

Mar.
2008

150
XIS 0(FI)
140
400

500

600
700
800
Time since 2005-7-10 (day)

900

1000

600
700
800
Time since 2005-7-10 (day)

900

1000

(b)

O VIII Kα FWHM (eV)70

XIS 1(BI)60

XIS 0(FI)
50

400

500

Figure B.9: Time history of the energy resolution (FWHM) at MnI Kα (5.895 keV) (a) andat OVIII Kα (0.653 keV) (b). The averaged value of the segments A and D are shown foreach sensor. The results are based on the grade 02346 events corrected with the ﬁne-tuned CTI
parameters as described in the text.

124

APPENDIX B. CTI CORRECTION AND IN-ORBIT CALIBRATION FOR SCI

The time histories of the energy resolution of the corrected data are shown in ﬁgure B.9. At
the high energy (5.895 keV), the energy resolution of XIS 0 at March 2008 is improved from∼160 eV to ∼155 eV by the sawtooth correction. On the other hand, the energy resolution of
XIS 1 at March 2008 is not improved in spite of the sawtooth correction, and stays at ∼175 eV.
At the low energy (0.653 keV), no clear eﬀect of the sawtooth correction is seen in both FI and
BI. The energy resolutions at March 2008 are ∼53 eV (XIS 0) and ∼62 eV (XIS 1) independentlyof the correction.
As we mentioned in section B.1, the sawtooth of the BI sensor is shallower than that of the FI
sensors. We think it causes the diﬀerence of the eﬀect of the sawtooth correction between BI and
FI. We speculate, at the low energy, the eﬀect of the sawtooth correction was not seen because
the original degradation of the energy resolution was small. In fact, the energy resolutions at
OVIII Kα are almost same between the SCI and SCI-oﬀ mode.

The results of our new method to correct the sawtooth structure has been implemented to thesoftware package released by HEASARC5 since HEAsoft version 6.3. Now, all of the XIS dataof the SCI mode after the processing version 2.0 are corrected by the sawtooth method.

5

http://heasarc.gsfc.nasa.gov/

Bibliography[1] Anders, E., & Grevesse, N. 1989, Geochim. Cosmochim. Acta, 53, 197[2] Arnaud, M., & Raymond, J. 1992, ApJ, 398, 394[3] Baganoﬀ, F. K., et al. 2003, ApJ, 591, 891[4] Balucinska-Church, M., & McCammon, D. 1992, ApJ, 400, 699[5] Bautz, M. W., Kissel, S. E., Prigozhin, G. Y., LaMarr, B., Burke, B. E., & Gregory, J. A.
2004, Proc. SPIE, 5501, 111[6] Bautz, M. W., et al. 2007, Proc. SPIE, 6686, 66860Q[7] Balucinska-Church, M., & McCammon, D. 1994, VizieR Online Data Catalog, 6062, 0[8] Bearden, J. A. 1967, Reviews of Modern Physics, 39, 78[9] Becker, R. H., White, N. E., Boldt, E. A., Holt, S. S., & Serlemitsos, P. J. 1980, ApJL, 237,
L77
[10] Bradt, H. V., Rothschild, R. E., & Swank, J. H. 1993, A&AS, 97, 355[11] Breitschwerdt, D., McKenzie, J. F., & Voelk, H. J. 1991, A&A, 245, 79[12] Burke, B. E., Gregory, J. A., Bautz, M. W., Prigozhin, G. Y., Kissel, S. E., Kosicki, B. B.,
Loomis, A. H., & Young, D. J. 1997, IEEE Transactions on Electron Devices, 44, 1633
[13] Churazov, E., et al. 2003, ApJ, 590, 225[14] Cooke, B. A., Griﬃths, R. E., & Pounds, K. A. 1970, Non-Solar X- and Gamma-Ray
Astronomy, 37, 280[15] Cropper, M. 1990, Space Science Reviews, 54, 195[16] Cutri, R. M., et al. 2003, The IRSA 2MASS All-Sky Point Source Catalog, NASA/IPAC
Infrared Science Archive. http://irsa.ipac.caltech.edu/applications/Gator/[17] Diehl, R., et al. 2006, Nature, 439, 45[18] Ebisawa, K., et al. 2005, ApJ, 635, 214[19] Ebisawa, K., et al. 2008, PASJ, 60, 223[20] Ezuka, H. & Ishida, M. 1999, ApJS, 120, 277
125

126

BIBLIOGRAPHY[21] Fomalont, E. B., Goss, W. M., Manchester, R. N., & Lyne, A. G. 1997, MNRAS, 286, 81[22] Gaetz, T. J., Butt, Y. M., Edgar, R. J., Eriksen, K. A., Plucinsky, P. P., Schlegel, E. M.,& Smith, R. K. 2000, ApJL, 534, L47[23] Giacconi, R., Gursky, H., Paolini, F. R., & Rossi, B. B. 1962, Physical Review Letters, 9,439[24] Giacconi, R., et al. 2002, ApJS, 139, 369[25] Gillessen, S., Eisenhauer, F., Fritz, T. K., Bartko, H., Dodds-Eden, K., Pfuhl, O., Ott, T.,& Genzel, R. 2009, ApJL, 707, L114
[26] Hardy, T., Murowinski, R., & Deen, M. J. 1998, IEEE Transactions on Nuclear Science,
45, 154[27] Hatsukade, I., Tsunemi, H., Yamashita, K., Koyama, K., Asaoka, Y., & Asaoka, I. 1990,
PASJ, 42, 279[28] Hog, E., Kuzmin, A., Bastian, U., Fabricius, C., Kuimov, K., Lindegren, L., Makarov,
V. V., & Roeser, S. 1998, A&A, 335, L65[29] Hyodo, Y., Tsujimoto, M., Koyama, K., Nishiyama, S., Nagata, T., Sakon, I., Murakami,
H., & Matsumoto, H. 2008, PASJ, 60, 173[30] Huovelin, J., Schultz, J., Vilhu, O., Hannikainen, D., Muhli, P., & Durouchoux, P. 1999,
A&A, 349, L21[31] Huovelin, J., Schultz, J., Vilhu, O., Hannikainen, D., Muhli, P., & Durouchoux, P. 2002,
arXiv:astro-ph/0208179[32] Ishida, M., Makishima, K., Mukai, K., & Masai, K. 1994, MNRAS, 266, 367[33] Ishida, M., Okada, S., Hayashi, T., Nakamura, R., Terada, Y., Mukai, K., & Hamaguchi,
K. 2009, PASJ, 61, 77[34] Ishisaki, Y., et al. 2007, PASJ, 59, 113[35] Inui, T., Koyama, K., Matsumoto, H., & Tsuru, T. G. 2009, PASJ, 61, 241[36] Jansen, F., et al. 2001, 365, L1[37] Jansky, K. G. 1933, Proceedings of the IRE Volume 21, Issue 10, Oct. 1933 Page(s):1387 1398
[38] Jones, L. R., & Pye, J. P. 1989, MNRAS, 238, 567[39] Kaastra, J.S., 1992, An X-Ray Spectral Code for Optically Thin Plasmas (Internal SRONLeiden Report, updated version 2.0)[40] Kaastra, J. S., & Mewe, R. 1993, A&AS, 97, 443[41] Kaneda, H., Makishima, K., Yamauchi, S., Koyama, K., Matsuzaki, K., & Yamasaki, N. Y.
1997, ApJ, 491, 638

BIBLIOGRAPHY

127[42] Kent, S. M., Dame, T. M., & Fazio, G. 1991, ApJ, 378, 131[43] Kelley, R. L., et al. 2007, PASJ, 59, 77[44] Koyama, K., Makishima, K., Tanaka, Y., & Tsunemi, H. 1986, PASJ, 38, 121[45] Koyama, K., Ikeuchi, S., & Tomisaka, K. 1986, PASJ, 38, 503[46] Koyama, K., Tsunemi, H., Becker, R. H., & Hughes, J. P. 1987, PASJ, 39, 437[47] Koyama, K., Awaki, H., Kunieda, H., Takano, S., & Tawara, Y. 1989, Nature, 339, 603[48] Koyama, K., Maeda, Y., Sonobe, T., Takeshima, T., Tanaka, Y., & Yamauchi, S. 1996,
PASJ, 48, 249[49] Koyama, K., Senda, A., Murakami, H., & Maeda, Y. 2003, Chinese Journal of Astronomy
and Astrophysics Supplement, 3, 297[50] Koyama, K., et al. 2007, PASJ, 59, 23[51] Koyama, K., et al. 2007, PASJ, 59, 245[52] Koyama, K., Uchiyama, H., Hyodo, Y., Matsumoto, H., Tsuru, T. G., Ozaki, M., Maeda,
Y., & Murakami, H. 2007, PASJ, 59, 237[53] Koyama, K., et al. 2007, PASJ, 59, 221[54] Koyama, K., Inui, T., Matsumoto, H., & Tsuru, T. G. 2008, PASJ, 60, 201[55] Koyama, K., Takikawa, Y., Hyodo, Y., Inui, T., Nobukawa, M., Matsumoto, H., & Tsuru,
T. G. 2009, PASJ, 61, 255[56] Krause, M. O., & Oliver, J. H. 1979, Journal of Physical and Chemical Reference Data, 8,
329[57] Krasnopolsky, V. A., Greenwood, J. B., & Stancil, P. C. 2004, Space Science Reviews, 113,
271
[58] Kushino, A., Ishisaki, Y., Morita, U., Yamasaki, N. Y., Ishida, M., Ohashi, T., & Ueda, Y.
2002, PASJ, 54, 327[59] Larsson, S. 1996, A&AS, 117, 197[60] Launhardt, R., Zylka, R., & Mezger, P. G. 2002, A&A, 384, 112[61] Liedahl, D. A., Osterheld, A. L., & Goldstein, W. H. 1995, ApJL, 438, L115[62] Lutovinov, A. A., Grebenev, S. A., Pavlinsky, M. N., & Sunyaev, R. A. 2001, Astronomy
Letters, 27, 501[63] Maeda, Y., et al. 2002, ApJ, 570, 671[64] Matsumoto, H., et al. 2006, Proc. SPIE, 6266,

128

BIBLIOGRAPHY[65] Matteucci, F., Panagia, N., Pipino, A., Mannucci, F., Recchi, S., & Della Valle, M. 2006,
MNRAS, 372, 265[66] Matsunaga, N., Kawadu, T., Nishiyama, S., Nagayama, T., Hatano, H., Tamura, M., Glass,
I. S., & Nagata, T. 2009, MNRAS, 399, 1709
[67] Mewe, R. 1996, Sol. Phys., 169, 335[68] Mezger, P. G., Duschl, W. J., & Zylka, R. 1996, A&A Rev., 7, 289[69] Mitsuda, K., et al. 2007, PASJ, 59, 1[70] Moretti, A., Campana, S., Lazzati, D., & Tagliaferri, G. 2003, ApJ, 588, 696[71] Morris, M., et al. 2003, Astronomische Nachrichten Supplement, 324, 167[72] Mori, H., et al. 2005, PASJ, 57, 245[73] Mori, H., et al. 2007, Proc. SPIE, 6686,[74] Mori, H., Tsuru, T. G., Hyodo, Y., Koyama, K., & Senda, A. 2008, PASJ, 60, 183[75] Mori, H., Hyodo, Y., Tsuru, T. G., Nobukawa, M., & Koyama, K. 2009, PASJ, 61, 687[76] Morrison, R. & McCammon, D. 1983, ApJ, 270, 119[77] Murakami, H., Koyama, K., Tsujimoto, M., Maeda, Y., & Sakano, M. 2001, ApJ, 550, 297[78] Muno, M. P., et al. 2003, ApJ, 589, 225[79] Muno, M. P., et al. 2004, ApJ, 613, 326[80] Muno, M. P., et al. 2004, ApJ, 613, 1179
[81] Muno, M. P., Bauer, F. E., Bandyopadhyay, R. M., & Wang, Q. D. 2006, ApJS, 165, 173[82] Muno, M. P., et al. 2009, ApJS, 181, 110[83] Nakajima, H., et al. 2008, PASJ, 60, 1[84] Nakajima, H., Tsuru, T. G., Nobukawa, M., Matsumoto, H., Koyama, K., Murakami, H.,
Senda, A., & Yamauchi, S. 2009, PASJ, 61, 233[85] Nobukawa, M., et al. 2008, PASJ, 60, 191[86] Nobukawa, M., Sawada, M., Matsumoto, H., Tsuru, T. G., & Koyama, K. 2009, Advancesin Space Research, 43, 1045 2008, PASJ, 60, 1[87] Nobukawa, M. et al. 2009 in prep.
[88] Norton, A. J., Watson, M. G., King, A. R., Lehto, H. J., & McHardy, I. M. 1992, MNRAS,
254, 705[89] Ozawa, M., et al. 2009, PASJ, 61, 1[90] Paul, J. M. 1971, Nuclear Instruments and Methods, 94, 275

BIBLIOGRAPHY

129[91] Park, S., Muno, M. P., Baganoﬀ, F. K., Maeda, Y., Morris, M., Howard, C., Bautz, M. W.,& Garmire, G. P. 2004, ApJ, 603, 548[92] Park, S., et al. 2005, ApJ, 631, 964[93] Patterson, J. 1994, PASP, 106, 209[94] Plucinsky, P. P., et al. 2008, Proc. SPIE, 7011
[95] Porquet, D., Decourchelle, A., & Warwick, R. S. 2003, A&A, 401, 197[96] Predehl, P. 1995, Advances in Space Research, 16, 159[97] Prigozhin, G., Burke, B., Bautz, M., Kissel, S., & Lamarr, B. 2008, IEEE Transactions on
Electron Devices, 55, 2111
[98] Rasmussen, A. P., Behar, E., Kahn, S. M., den Herder, J. W., & van der Heyden, K. 2001,
A&A, 365, L231
[99] Rana, V. R., Singh, K. P., Schlegel, E. M., & Barrett, P. E. 2006, ApJ, 642, 1042
[100] Reid, M. J., & Brunthaler, A. 2004, ApJ, 616, 872[101] Revnivtsev, M., Sazonov, S., Gilfanov, M., Churazov, E., & Sunyaev, R. 2006, A&A, 452,
169[102] Revnivtsev, M., Molkov, S., & Sazonov, S. 2006, MNRAS, 373, L11[103] Revnivtsev, M., & Sazonov, S. 2007, A&A, 471, 159[104] Revnivtsev, M., Vikhlinin, A., & Sazonov, S. 2007, A&A, 473, 857[105] Revnivtsev, M., Sazonov, S., Churazov, E., Forman, W., Vikhlinin, A., & Sunyaev, R.
2009, Nature, 458, 1142
[106] Ritter, H., & Kolb, U. 2003, A&A, 404, 301[107] Rybicki, G. B., & Lightman, A. P. 1979, New York, Wiley-Interscience, 1979. 393 p.,[108] Ryu, S. G., Koyama, K., Nobukawa, M., Fukuoka, R., & Tsuru, T. G. 2009, PASJ, 61,751[109] Sakano, M., Imanishi, K., Tsujimoto, M., Koyama, K., & Maeda, Y. 1999, ApJ, 520, 316[110] Sakano, M., Koyama, K., Murakami, H., Maeda, Y., & Yamauchi, S. 2002, ApJS, 138, 19[111] Sakano, M., Warwick, R. S., Decourchelle, A., & Predehl, P. 2003, MNRAS, 340, 747[112] Sakano, M., Warwick, R. S., Decourchelle, A., & Predehl, P. 2004, MNRAS, 350, 129[113] Sawada, M., Tsujimoto, M., Koyama, K., Law, C. J., Tsuru, T. G., & Hyodo, Y. 2009,
PASJ, 61, 209[114] Schwope, A. D., Brunner, H., Buckley, D., Greiner, J., Heyden, K. v. d., Neizvestny, S.,
Potter, S., & Schwarz, R. 2002, A&A, 396, 895

130

BIBLIOGRAPHY[115] Sazonov, S., Revnivtsev, M., Gilfanov, M., Churazov, E., & Sunyaev, R. 2006, A&A, 450,
117[116] Sciortino, S. 2008, Memorie della Societa Astronomica Italiana, 79, 192[117] Serlemitsos, P. J., et al. 2007, PASJ, 59, 9[118] Sidoli, L., Belloni, T., & Mereghetti, S. 2001, A&A, 368, 835[119] Smith, R. K., Brickhouse, N. S., Liedahl, D. A., & Raymond, J. C. 2001, ApJL, 556, L91[120] Strüder, L., et al. 2001, A&A, 375, L5[121] Takagi, S., Ph.D thesis, Kyoto University, 2006[122] Tanaka, Y., Koyama, K., Maeda, Y., & Sonobe, T. 2000, PASJ, 52, L25[123] Tanaka, Y. 2002, A&A, 382, 1052
[124] Tatischeﬀ, V. 2003, EAS Publications Series, 7, 79[125] Tawa, N., et al. 2008, PASJ, 60, 11[126] Takahashi, T., et al. 2007, PASJ, 59, 35[127] Tibolla, O., Chaves, R. C. G., de Jager, O., Domainko, W., Fiasson, A., Komin, N., &Kosack, K. 2009, arXiv:0907.0574[128] Tomida, H., et al. 1997, PASJ, 49, 405[129] Tsuboi, M., Handa, T., & Ukita, N. 1999, ApJS, 120, 1[130] Tsujimoto, M., Hyodo, Y., & Koyama, K. 2007, PASJ, 59, 229[131] Tucker, W. H. 1970, Ap&SS, 9, 315[132] Tsunemi, H., Yamashita, K., Masai, K., Hayakawa, S., & Koyama, K. 1986, ApJ, 306, 248[133] Tsuru, T. G., Nobukawa, M., Nakajima, H., Matsumoto, H., Koyama, K., & Yamauchi,
S. 2009, PASJ, 61, 219[134] Uchiyama, H., et al. 2009, PASJ, 61, 9[135] Uchiyama, Y., et al. 2008, PASJ, 60, S35[136] van Loon, J. T., et al. 2003, MNRAS, 338, 857[137] Warwick, R. S., Turner, M. J. L., Watson, M. G., & Willingale, R. 1985, Nature, 317, 218[138] Weisskopf, M. C., Brinkman, B., Canizares, C., Garmire, G., Murray, S., & Van Speybroeck, L. P. 2002, PASP, 114, 1[139] Watson, M. G., et al. 2009, A&A, 493, 339[140] Yamaguchi, H., et al. 2006, Proc. SPIE, 6266,

BIBLIOGRAPHY

131[141] Yamaguchi, H., et al. 2008, PASJ, 60, 141[142] Yamauchi, S., & Koyama, K. 1993, ApJ, 404, 620[143] Yamauchi, S., Ebisawa, K., Tanaka, Y., Koyama, K., Matsumoto, H., Yamasaki, N. Y.,
Takahashi, H., & Ezoe, Y. 2009, PASJ, 61, 225[144] Yusef-Zadeh, F., Melia, F., & Wardle, M. 2000, Science, 287, 85[145] The
SuzakuTechnical
Description
Announcement
http://www.astro.isas.jaxa.jp/suzaku/proposal/ao5/suzaku td/ofOpportunity#5,

Acknowledgement
First of all, I am deeply grateful to my supervisor, Prof. K. Koyama, Assoc. Prof. T. Tsuru,
and Res. Assoc. H. Matsumoto for their continual guidance, advice and support throughout mygraduate career. Their patience and sometimes strict guidance led me to complete this work.
In particular, Tsuru-san checked and revised this thesis until really late night even in the
New Year day. I sincerely appreciate and want to say sorry to him and his family.
I thank all the members of the current and past Suzaku team, especially the XIS team. The
quality of the GC data used in this thesis is excellent and it is due to a lot of their eﬀorts.
I thank all the members of the current and past CR laboratory during my graduate career.
Special thanks to M. Nobukawa and S. Kabuki. Thanks to my ”comrade-in-arms”, K. Ueno, T.
Mizukami, and M. Ozawa.
I really thank my family.
I thank the Suzaku satellite who provides the data used in this thesis and is working alone
in orbit at this moment.

133