X-RAY STUDY ON PLASMA OUTFLOWSFROM THE GALACTIC CENTER
by

Shinya Nakashima
Department of Physics, Graduate School of Science, Kyoto University
Kitashirakawa Oiwake-cho, Sakyo-ku, Kyoto, 606-8502, Japanshinya@cr.scphys.kyoto-u.ac.jp

A dissertation submitted to the Department of Physics,
Graduate School of Science, Kyoto university
in partial fulfillment of the requirements
for the degree of Doctor of Philosophy in physics
January, 2014

Abstract
The center of our Galaxy (Galactic center; GC) has a lower activity than that of active galactic nuclei or starburst galaxies at present. The current luminosity of the central
supermassive black hole, Sagittarius (Sgr) A*, is eight orders of magnitude lower than its
Eddington luminosity of 1044 erg s−1 . Moreover, on-going explosive star formation is notfound. However, signatures of past GC activities have been reported recently. For instance,
X-ray observations have discovered fluorescent lines of neutral iron from dense molecular
clouds in the GC. They are likely to be caused by reflections of past Sgr A* flares occurring
100–1000 yr ago. The luminosity of the flares are estimated as ∼1039 erg s−1 . Gamma-ray
observations also revealed the giant bubbles, which suggest another Sgr A* flare or a nuclear
starburst 106 –107 yr ago. Such observational evidence of past GC activities is still limited,
so that there is a large gap in the activity history between 103 and 106 yr ago. Our aim is
to find other relics of GC activities and to clarify the history of them.
Previous X-ray observations have concentrated on the Galactic plane region. Contrastingly, we focused on the Galactic bulge region because it is less contaminated by the strong
GC background emission and is more suitable to detect phenomena associated with the GC
activities. We performed the survey of the off-plane region with Suzaku, and found the large
diffuse plasma to the south and north of the GC, which are referred to as GC South and GC
North, respectively.
From the spectral analysis of GC South, we discovered strong radiative recombination
continua of highly ionized Si and S. We suggest that the plasma is in a recombining phase,
which is not expected in the standard shock-heating scenario. Indeed, a recombining plasma(RP) model well represents the spectrum, and gave us the electron temperature of 0.5 keV,
the initial ionization temperature of 1.6 keV. The recombination timescale of 105 yr indicatesthat GC South was generated ∼105 yr ago. The spectrum suffers from strong interstellar
absorption with NH ∼ 1 × 1022 H cm−2 , which is consistent with the GC distance. Then themass and thermal energy of the plasma reaches 3000 M and 8×1051 erg, respectively, whichis two orders of magnitude higher than that of a typical Galactic supernova remnant. We
also investigated the spatial variation of the plasma parameters, and found that the electron
temperature and density decrease as the distance from the GC becomes large, changing from
0.4 keV to 0.8 keV, and from 0.06 to 0.12 cm3, respectively. The plasma become a collisional

ionization equilibrium (CIE) state near the GC.
The spectra of GC North is represented by a CIE plasma with the electron temperatureof 0.7 keV. The absorption column densities of NH ∼ 3 × 1022 H cm−2 indicates that GC
North is also located at the GC distance. The mass and thermal energy are found to be
900 M and 3 × 1051 erg, respectively. The dynamical time scale of the plasma is estimatedas ∼105 yr. These parameters are comparable to those of GC South.
We consider that a past GC activity ∼105 yr ago generated GC South/North on the basisof large thermal energies in both the plasma and the RP in GC South.

Contents
1 Introduction

12 Review
2.1 Overview of the Milky Way Galaxy . . . . . . . . . . . . . . . . . . . . . . .
2.2 Galactic Center Environment . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2.1 Characteristics of the interstellar medium . . . . . . . . . . . . . . .
2.2.2 X-ray emitting diffuse plasma . . . . . . . . . . . . . . . . . . . . . .
2.2.3 Magnetic field . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2.4 Star formation in the GC . . . . . . . . . . . . . . . . . . . . . . . .
2.3 Supermassive Black Hole Sagittarius A* . . . . . . . . . . . . . . . . . . . .
2.3.1 Evidence for the central supermassive black hole . . . . . . . . . . . .
2.3.2 Emissions from Sgr A* . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.3 Evidence of past activity . . . . . . . . . . . . . . . . . . . . . . . . .
2.4 Subject of This Thesis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

3
4
5
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8
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3 Instruments
3.1 Overview of Suzaku . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2 X-ray Telescope (XRT) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.1 X-ray optics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.2 Telescope design . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.3 In-flight performance . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3 X-ray Imaging Spectrometer (XIS) . . . . . . . . . . . . . . . . . . . . . . .
3.3.1 Overview of the XIS system . . . . . . . . . . . . . . . . . . . . . . .
3.3.2 Device description of the CCD camera . . . . . . . . . . . . . . . . .
3.3.3 In-flight performance and calibration . . . . . . . . . . . . . . . . . .
3.4 Advantage of Suzaku in the GC survey . . . . . . . . . . . . . . . . . . . . .

25
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31
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36
41

4 X-ray Images of the South and North of the GC
4.1 Suzaku X-ray Images of the Galactic Center . . . . . . . . . . . . . . . . . .
4.1.1 Observations and data reduction . . . . . . . . . . . . . . . . . . . .

43
44
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i

CONTENTS
4.1.2 Construction of X-ray images . . . . . . . . . . . . . . . . . . . . . .
ASCA X-ray Image of the Galactic Center . . . . . . . . . . . . . . . . . . .

44
45

5 Analysis on GC South
5.1 Analysis Method . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.2 Background Estimation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3 Spectrum Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3.1 Brightest region spectrum . . . . . . . . . . . . . . . . . . . . . . . .
5.3.2 Spatially resolved spectra . . . . . . . . . . . . . . . . . . . . . . . .
5.4 Physical Properties of GC South . . . . . . . . . . . . . . . . . . . . . . . . .

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6 Analysis on GC North
6.1 Analysis Method . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.2 Background Estimation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.3 Spectrum Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.3.1 Merged spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.3.2 Spatially resolved spectra . . . . . . . . . . . . . . . . . . . . . . . .
6.4 Physical Properties of GC North . . . . . . . . . . . . . . . . . . . . . . . .

69
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7 Discussion
7.1 Summary of the Obtained Parameters . . . . . . . . . . . . . . . . . . . . .
7.2 Origin of GC South . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.2.1 Energy source . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.2.2 Mechanism to form the recombining plasma . . . . . . . . . . . . . .
7.3 Origin of GC North . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.4 Connection between GC South and GC North . . . . . . . . . . . . . . . . .

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8 Conclusion

89

A Emission from a Thermal Plasma
A.1 Ionization State in a Plasma . . . . . . . . . . . . . . . . . . . . . . . . . . .
A.2 Radiation Process in a Thermal pPasma . . . . . . . . . . . . . . . . . . . .
A.3 Spectrum of a Recombining Plasma . . . . . . . . . . . . . . . . . . . . . . .

91
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93

4.2

–ii–

List of Figures
2.1 Schematic illustration of the Galaxy. . . . . . . . . . . . . . . . . . . . . . .
2.2 Longitude-velocity map of the 12 CO emission. . . . . . . . . . . . . . . . . .
2.3 Density profile of the ISM along the Galactic latitude at just the GC. . . . .
2.4 X-ray spectra of the GC. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.5 Suzaku X-ray images of the GC. . . . . . . . . . . . . . . . . . . . . . . . . .
2.6 Radio image of the GC. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.7 Orbits of S-Stars. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.8 Sgr A* spectra in the radio band. . . . . . . . . . . . . . . . . . . . . . . . .
2.9 Intrinsic size of Sgr A*. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.10 X-ray image and spectrum of Sgr A*. . . . . . . . . . . . . . . . . . . . . . .
2.11 SED of Sgr A*. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.12 Suzaku X-ray image in the Fe Iband. . . . . . . . . . . . . . . . . . . . . . .
2.13 X-ray spectrum of the XRN in Sgr A. . . . . . . . . . . . . . . . . . . . . . .
2.14 Time variation of the 6.4 keV line and continuum in the Sgr B2 region. . . .
2.15 Light curve of Sgr A*. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.16 Images of the Fermi bubbles. . . . . . . . . . . . . . . . . . . . . . . . . . .
2.17 Spectrum of the Fermi bubbles. . . . . . . . . . . . . . . . . . . . . . . . . .
2.18 Schematic diagram of past GC activities. . . . . . . . . . . . . . . . . . . . .

4
5
7
9
10
11
13
14
14
16
18
19
20
20
21
22
23
24

3.1 Schematic illustration of Suzaku. . . . . . . . . . . . . . . . . . . . . . . . . .
3.2 Schematic view of the Wolter type I telescope. . . . . . . . . . . . . . . . . .
3.3 Photos of the XRT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.4 Effective are of the XRT. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.5 Calibration image of SS Cygni. . . . . . . . . . . . . . . . . . . . . . . . . .
3.6 Schematic view of the stray light. . . . . . . . . . . . . . . . . . . . . . . . .
3.7 Images of the XRT with and without the pre-collimator. . . . . . . . . . . .
3.8 Photo of the XIS camera body. . . . . . . . . . . . . . . . . . . . . . . . . .
3.9 Schematic overview of the XIS system. . . . . . . . . . . . . . . . . . . . . .
3.10 QE of the XIS. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

26
28
29
30
31
32
32
33
35
36

iii

LIST OF FIGURES
3.11 Long term trend of 55 Fe centroid energy and width before charge-trail andCTI correction. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.12 Schematic illustration of the CTI. . . . . . . . . . . . . . . . . . . . . . . . .
3.13 Long term trend of 55 Fe centroid energy after charge-trail and CTI correction.
3.14 Long term trend of the line with after charge-trail and CTI correction. . . .
3.15 XIS spectra of RX J1856.5−3754 taken from 2005 Oct. to 2012 Apr. . . . .
3.16 Long term trend of the optical depth of the contamination material. . . . . .
3.17 Spectra of the NXB. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

37
38
39
39
40
41
42

4.1
4.2
4.3

X-ray images of the GC region taken by Suzaku. . . . . . . . . . . . . . . . .
Projection of the Suzaku image. . . . . . . . . . . . . . . . . . . . . . . . . .
X-ray image of the GC region taken by ASCA. . . . . . . . . . . . . . . . . .

45
46
47

5.1
5.2
5.3
5.4
5.5
5.6
5.7
5.8
5.9

Close-up view of the GC South region. . . . . . . . . . . . . . . . . . . . . .
Spectra and the best-fit models of reference regions. . . . . . . . . . . . . . .
Profile of the absorption column density. . . . . . . . . . . . . . . . . . . . .
Spectrum extracted from the entire source region. . . . . . . . . . . . . . .
Spectra of the S-Src 1–10 regions with the tentative model. . . . . . . . . . .
Spectra of the S-Src 1–10 regions with the modified model. . . . . . . . . . .
Absorption column densities along the Galactic latitude. . . . . . . . . . . .
Best-fit parameters along the Galactic latitude. . . . . . . . . . . . . . . . .
Spatial profiles of derived physical parameters . . . . . . . . . . . . . . . . .

51
53
54
57
60
62
64
65
67

6.1
6.2
6.3
6.4

Close-up view of the GC North region. . . . . . . . . . . . . . . . . . . . . .
Spectra of the reference region. . . . . . . . . . . . . . . . . . . . . . . . . .
Merged Spectra of GC North. . . . . . . . . . . . . . . . . . . . . . . . . . .
Spectra extracted from N-Src 1 and N-Src 2. . . . . . . . . . . . . . . . . . .

70
72
74
75

7.1

Absorption column densities along the Galactic latitude obtained from the
GC South and GC North. . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Spatial profiles of the surface brightness, electron temperature, electron density, and thermal pressure along the Galactic latitude. . . . . . . . . . . . . .
Correlation between the temperature and density of GC South. . . . . . . .
Emission of CO (J=2–1) overlaid with the X-ray contours. . . . . . . . . . .
Revealed activity history of the GC. . . . . . . . . . . . . . . . . . . . . . . .

7.2
7.3
7.4
7.5

A.1 Spectrum model of a thermal plasmas with kTe = 0.5 keV. . . . . . . . . . .

–iv–

79
80
83
85
87
94

List of Tables
2.1

Typical ISM characteristics in the Galactic center and the Disk

. . . . . . .

6

5.1 Observation logs of the GC South region. . . . . . . . . . . . . . . . . . . . .
5.2 Fitting results of the spectra in the reference regions. . . . . . . . . . . . . .
5.3 Fitting results of the brightest-region spectrum with the NEIJ model. . . . .
5.4 Average charge state of each element. . . . . . . . . . . . . . . . . . . . . . .
5.5 Fitting results of the spatially resolved spectra with the NEIJ model. . . . .
5.6 Physical parameters of each region derived from the spatially resolved spectrum fitting. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

52
55
58
59
63

6.1
6.2
6.3
6.4

Observation logs of the GC North region. . . . . . . . . . . . . . . . . . . . .
Fitting results of the spectra in the reference regions for GC North. . . . . .
Fitting results of the merged spectrum. . . . . . . . . . . . . . . . . . . . . .
Physical parameters of the N-Src 1 and N-Src 2. . . . . . . . . . . . . . . . .

71
72
74
76

7.1

Typical physical parameters of GC South and GC North . . . . . . . . . . .

78

v

66

LIST OF TABLES–vi–

Chapter 1
Introduction
Feeding and feedback of activities in the center of a galaxy, including a supermassive blackhole (SMBH), are considered to play important roles in the context of evolution of galaxies.
For its proximity, the center of our Galaxy (Galactic center; GC), is the best laboratory
to investigate the details of these mechanisms. The GC has a lower activity than that ofactive galactic nuclei (AGNs) or starburst galaxies at present. The current luminosity of thecentral SMBH, Sagittarius (Sgr) A*, is eight orders of magnitude lower than its Eddington
luminosity of 1044 erg s−1 (e.g., Baganoff et al. 2003). Moreover, on-going explosive starformation is not found (e.g., Longmore et al. 2013). However, evidence of past GC activities
have been reported recently. For instance, X-ray observations discovered fluorescent lines ofneutral iron from dense molecular clouds in the GC (Koyama et al. 1996). They are likely
to be caused by reflections of past Sgr A* flares occurring 100–1000 yr ago, which shouldhave the luminosity of 1039 erg s−1 (Ryu et al. 2013). Gamma-ray observations also revealed
the giant bubbles, which suggest another Sgr A* flare or a nuclear starburst 106 –107 yr ago(Su et al. 2010). Such observational evidence of past GC activities is still limited, so that
there is a large gap in the activity history between 103 and 106 yr ago. Is there any otherevidence in the GC region?
In order to approach this subject, we focused on the Galactic bulge region, which isless contaminated by the strong GC background emission and is more suitable to detect
phenomena associated with the GC activities. We performed the survey of this region with
Suzaku, and found the large diffuse plasma to the south and north of the GC, possible relicsof a past GC activity. In this thesis, we analyzed these emissions.
In Chapter 2, we review the specific environment of the GC and evidence of past GC
activities. Performances of the instruments onboard Suzaku are summarized in Chapter 3.
We show the GC image taken by the Suzaku survey and clearfy the diffuse X-ray emissions
extending to the north and south of the GC in Chapter 4. Then, the results of spectral
analyses on these emissions are descrived in Chapter 5 and Chapter 6. Using these results, we
1CHAPTER 1. INTRODUCTION
discuss the possible origins in Chapter 7. Finally, we described the summary and conclusionof this thesis in Chapter 8.
Throughout this thesis, we mainly use the Galactic coordinates, where the horizontal
and vertical axes are the Galactic longitude (l) and latitude (b), respectively. We refer tothe positive directions of l and b as east and north, respectively. We adopt the distance to
the GC to be 8 kpc. Errors are quoted at the 90% confidence level otherwise noted.

–2–Chapter 2
Review
In this chapter, we review the specific environment of the GC including the central SMBH,
and evidence of past GC activities to clarify the subject of this thesis.

Contents
2.1

Overview of the Milky Way Galaxy . . . . . . . . . . . . . . . .

4

2.2

Galactic Center Environment . . . . . . . . . . . . . . . . . . . .

5

2.2.1

Characteristics of the interstellar medium . . . . . . . . . . . . . .

5

2.2.2

X-ray emitting diffuse plasma . . . . . . . . . . . . . . . . . . . . .

6

2.2.3

Magnetic field . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

8

2.2.4

Star formation in the GC . . . . . . . . . . . . . . . . . . . . . . .

11

Supermassive Black Hole Sagittarius A* . . . . . . . . . . . . .

12

2.3.1

Evidence for the central supermassive black hole . . . . . . . . . .

12

2.3.2

Emissions from Sgr A* . . . . . . . . . . . . . . . . . . . . . . . . .

12

2.3.3

Evidence of past activity . . . . . . . . . . . . . . . . . . . . . . . .

17

Subject of This Thesis . . . . . . . . . . . . . . . . . . . . . . . .

23

2.3

2.4

3

CHAPTER 2. REVIEW(a) Edge-on view(b) Face-on view

Halo

Outer Arm

s Arm
rseu
Pe

Disk

Far 3
kp
cA

Ne

rm

Sun

a

r3

8 kpc

kpc A

N or

30 kpc
Ca

rin a

rm

ma Ar mS cutu m-C

us

4 kpc

Sun

ta

ur

Bulge

en

- S a gitta ri u s

Figure 2.1. Schematic illustration showing the structure of the Galaxy in (a) the edge-on viewand (b) the face-on view.

2.1

Overview of the Milky Way Galaxy

The Milky Way Galaxy (the Galaxy), harboring our solar system, is a typical barred spiralgalaxy belonging to the Local Group, where three large (the Galaxy, M31 and M33) and over
30 small galaxies are contained. Since the solar system is situated within the Galaxy, we
are not able to directly see its overall morphology. Multi-wavelength observations, especiallyusing the 21 cm and 12 CO line, enable us to know the dynamics of the interstellar medium(ISM) and the structure of the Galaxy.
As shown in Figure 2.1(a), the Galaxy consists of the disk, the bulge, and the halo. The
solar system is located at the 8 kpc distance from the GC and 40 pc above the plane of the
disk. The disk is the main component of the Galaxy containing the most of Population I
stars. It has the 30 kpc diameter and the 1 kpc thickness with the mass of ∼1011 M . The
bulge is the spheroid located at the center of the disk with the major and minor axes of
4 kpc and 3 kpc, respectively. In the face-on view (Figure 2.1(b)), the bulge is flattened likea bar with the 25◦ orientation angle to the line of sight. The mass of the bulge is estimatedto be ∼1010 M . The disk and the bulge are surrounded by the ellipsoidal halo with theradius of 20 kpc. Its mass is found to be ∼4 × 1011 M from the rotation curve (therefore
including dark matter). Population II stars dominate in the bulge and halo.
Figure 2.1(b) shows the spiral arms in the disk, which have been traced by doppler shiftsof the emission lines from H I and 12 CO. Figure 2.2 shows the l-v diagram of the 12 CO line–4–

2.2. GALACTIC CENTER ENVIRONMENT

Sagittairus
Tangent

Carina Arm

Local Arm

pc

m
r Ar

m

Ar

Oute

3k

Arm

ar

PerseusNe

LSR radial velocity (km/s)

ScutumTangent

Perse

us Ar
m

Centaurus
Norma Tangent
Tangent

Galactic longitude (degree)

Figure 2.2. Longitude-velocity map of the 12 CO emission. The horizontal and vertical axes are
the Galactic longitude and the local standard of rest (LSR) radial velocity, respectively. The centerof the figure is (l, v) = (0.◦ 0, 0.0 km s−1 ) The upper and lower half of the panel indicates positive
and negative LSR velocities, respectively. (From Dame et al. 2001.)

survey covering the entire Galactic plane. Its symmetrical shape to the center of the image(l = 0, v = 0) reflects the rotation of the Galaxy. The arc structures correspond to spiralarms, and are named after the constellations. The arms in the first and third quadrants andthe second and fourth quadrants are within and beyond the solar circle, respectively. The
GC direction shows very large velocity dispersion, which suggests distinctive condition in
the GC.

2.2

Galactic Center Environment

2.2.1

Characteristics of the interstellar medium

The ISM in the central hundreds parsecs of the Galaxy has different physical characteristics from that in the disk. Table 2.1 shows typical parameters of the ISM in the GC andthe disk. In general, the GC ISM is dense and high temperature. The H I-to-H2 ratio in the
GC is also extremely smaller (∼0.05) than that in the disk (∼1) (Ferrière et al. 2007).
The predominant medium in the GC is dense molecular clouds (MCs), the so-called
central molecular zone (CMZ; see Morris & Serabyn 1996 for a review), which contains ∼5%
mass of the total MCs in the Galaxy. Its temperature and density reach 70 K and 104 cm−3 ,respectively. Such a extremely warm and dense molecular is hardly seen in the disk except
for molecular cores. Moreover, the velocity dispersion of the CMZ is 15–50 km s−1 . The
large turbulent pressure might be in equilibrium with those of the magnetic fields and the
hot plasmas.
–5–CHAPTER 2. REVIEW
Table 2.1. Typical ISM characteristics in the Galactic center and the Disk

Molecular (H2 )Neutral (H I)

Warm (H II)

Hot (H II)

Galactic center(Dense)T (K)70n (cm−3 )104
M (M ) 5 × 107(Diffuse)
150
102
107

100
10
106

104
10
105

107 –108
0.1
103

104
0.1
108

107
10−3
107

Disk region
T (K)n (cm−3 )M (M )

10
102
109

100
0.3
109

Notes:
Molecular: Oka et al. (1998); Tsuboi et al. (1999) for the GC. Dame (1993) for the disk.
Neutral: Ferrière et al. (2007) for the GC. Rohlfs & Kreitschmann (1987) for the disk.
Warm: Roy (2013) for the GC. Haffner et al. (2009) for the disk. Not including HII regions.
Hot: Uchiyama et al. (2013) for the GC. Snowden et al. (1997) for the disk. Not including SNRs

Another distinctive component in the GC is the very hot plasma. Its mass is very small(∼103 M ) but its thermal pressure is comparable to that of the CMZ, and thus it has aconsiderable effect on the GC environment. Details of its nature are described in Section
2.2.2.
The spatial distribution of the ISM in the GC region was modeled by Ferrière et al.
(2007). Figure 2.3 shows the density profile along the Galactic latitude at just the GC. The
cold matters (H2 and H I) are highly concentrated within 100 pc from the Galactic plane
while the ionized medium is sparse and extended.

2.2.2

X-ray emitting diffuse plasma

Early X-ray observations detected not only a number of luminous X-ray binaries but alsounresolved emissions on the Galactic plane. From the EXOSAT observations, Warwick et al.
(1985) reported that this unresolved emission extends within |l| < 40◦ , and thus referred
to it as the Galactic ridge X-ray emission (GRXE). Koyama et al. (1986) discovered theintense 6.7 keV (Fe XXV) line from the GRXE spectrum, and suggested the origin of an
optically-thin plasma with the temperature of 5–10 keV. Subsequently, Koyama et al. (1989)
found that the 6.7 keV line is smoothly distributed within |l| < 60◦ and has a strong peakat |l| < 1◦ . The peak intensity is 10 times larger than that of the outer region. This intense–6–2.2. GALACTIC CENTER ENVIRONMENT

Figure 2.3. Density profile of the ISM along the Galactic latitude at just the GC. The red, blue,and green lines show the molecular, neutral, and ionized phase. (From Ferrière et al. 2007.)

GC emission is referred to as the Galactic center X-ray emission (GCXE) to distinguish itfrom the GRXE.
An extensive study on the GCXE was carried out by ASCA, which is the first satellite
using X-ray CCDs with the imaging capability up to 10 keV. As shown in Figure 2.4(a),ASCA resolved the ∼6.7 keV line seen in the Tenma and Ginga spectrum into the compositionof three lines: 6.4 keV (Fe I), 6.7 keV (Fe XXV), and 6.9 keV (Fe XXVI) lines (Koyama et al.
1996). The 6.7 keV and 6.9 keV lines likely come from a thermal plasma with kT = 7 keV assuggested in the previous studies, whereas the 6.4 keV line should be emitted from neutral
cold matters by another mechanism. The nature and origin of the 6.4 keV line is discussed
in Section 2.3.3.
As shown in Figure 2.5(a), the 6.7 keV emission is smoothly distributed over the hundreds
parsecs of the GC. If it is of a diffuse plasma origin, its density and total thermal energy
is found to be 0.05 cm−3 and 1053 erg s−1 , respectively (Uchiyama et al. 2013). The soundvelocity of the 7 keV plasma (∼ 1500 km s−1 ) exceeds the escape velocity of the GC region,
so that the plasma leaves from the GC in ∼2 × 104 yr. Therefore, one supernova (SN) per
200 yr (equivalent to 1041 erg s−1 ) is required to sustain the plasma. However, this SN rate isunreasonably high because one SN per 1000 yr is expected in this region scaling the Galactic
SN rate (one SN per 50 yr) by the mass ratio (5%). This is the most serious problem for the
GCXE.
In order to solve this energy problem, other scenarios have been proposed. One is the–7–

CHAPTER 2. REVIEW
superposition of unresolved point sources consisting of intermediate polars or active stars,
which emit the 6.4 keV, 6.7 keV, and 6.9 keV lines. Indeed, Revnivtsev et al. (2009) argued
that more than 80% of the 6–7 keV emission in the bulge region (l = 0.◦ 0, b = −1.◦ 4) is
resolved into discrete sources by Chandra. However, Chandra resolved only 30% of the total
flux into discrete sources near Sgr A* (Muno et al. 2004). Moreover, Nishiyama et al. (2013)
compared the spatial distribution between the X-ray surface brightness and stellar massdensity, and found ∼50% discrepancy between them at the GC.
Another scenario, the origin of charge exchange recombination, was claimed by Tanaka
et al. (2000). In this scenario, interaction between neutral hydrogen and Fe XXV or Fe XXVI
in cosmic-rays causes electron capture into a excited state of Fe ions, and then de-excitationof this electron emits the 6.7 keV or 6.9 keV lines. The 6.7 keV line mainly consists of theforbidden line (6636 eV) and the resonance line (6700 eV), which are predominantly due torecombination and collisional excitation, respectively. In a collisional ionization equilibrium
plasma, the resonance line dominates and hence center energy of the Fe XXV line is ∼6700 eV.
If the charge exchange recombination is the case, the forbidden line is enhanced and center
energy of the Fe XXV line is close to the 6636 eV. Using the Suzaku high-quality spectrum(Figure 2.4(b)), Koyama et al. (2007b) found a centroid of the 6.7 keV line is 6680 ± 1 eV,
which suggests that the 6.7 keV line is not due to charge exchange but collisional excitation.
Therefore, the charge exchange scenario is no favored. The origin of the 7 keV thermal
plasma is still an open question.
In addition, K-shell emission lines from highly-ionized Si, S, Ar, and Ca are also detectedin the GCXE spectrum (see Figure 2.4(a)), which are typical features of the thin thermal
plasma with kT ∼ 1 keV. As shown in Figure 2.5(b), the spatial distribution of the S line ispatchy in contrast to the 7 keV plasma. Some of them are X-ray SNRs reported by many
authors (Koyama et al. 2007a; Nobukawa et al. 2008; Tsuru et al. 2009; Mori et al. 2009;
Sawada et al. 2009; Nakashima et al. 2010; Ohnishi et al. 2011). The mean density andtotal thermal energy of the 1 keV plasma are 0.07 cm−3 and 1052 erg s−1 , respectively. Since
the dynamical timescale of the1 keV plasma is ∼ 105 yr, one SN per 104 yr is required tosustain the plasma. It is consistent with the expected SN rate in this region. The total flux
and luminosity in the 2–10 keV band from the 1◦ square field are ∼ 10−9 erg s−1 cm−2 and∼ 1037 erg s−1 , respectively.

2.2.3

Magnetic field

As shown in Figure 2.6, radio continuum observations have revealed numerous filamentary structures in the GC region (see Morris 2007 for a review). Their radio polarization
indicates that the emissions are synchrotron radiation from the magnetic flux tube filledwith non-thermal electrons. In addition, the position angles of the polarization shows that–8–

2.2. GALACTIC CENTER ENVIRONMENT
(a)

(b)

Figure 2.4. (a) ASCA spectrum of GCXE in the 1–10 keV band. (From Koyama et al. 1996.) (b)Suzaku spectrum of GCXE in the 5.5–12 keV band with the phenomenological model (Gaussians
plus power-law functions). (From Koyama et al. 2007b.)

the magnetic filed along the filaments. Most of the filaments elongated from north to southwith < 0.5 pc in width and tens of parsecs in length. Therefore, globally poloidal magnetic
fields in the GC is suggested.
The most outstanding feature located at l = 0.◦ 15 is referred to as the “Arc”. The
Arc seemed to interact with nearby molecular cloud, whose density is ∼103 cm−3 with thevelocity dispersion of ∼15 km s−1 (Yusef-Zadeh & Morris 1987). The Arc shows almost
invariant curvature and absence of distortion against the large turbulent pressure of thecloud (10−8 erg cm−3 ). Therefore, Yusef-Zadeh & Morris (1987) argued that the magnetic
filed of the Arc should be &1 mG to overcome the turbulent pressure. The cooling time ofthe synchrotron radiation is described as

−1 
−2
Ee
B
4
yr,
(2.1)
tsync = 1.3 × 10
1 GeV
1 mG
where Ee and B is the electron energy and the magnetic flux density. Assuming the electrons
diffuse along the magnetic field with Alfvèn speed ofBne −0.5
vA = 2200
km s−1 ,
(2.2)
−3
1 mG
1 cm
where ne is the electron density, 1 GeV electrons can travel vA tsync ∼ 30 pc before losing
their energy. It is the same order of magnitude as the observed Arc length (60 pc).
LaRosa et al. (2005) found a diffuse synchrotron emission without filamentary structuresin the low frequency band (74 MHz). Assuming equipartition of the energy between the–9–CHAPTER 2. REVIEW

0.5(a) Fe XXV

Galactic latitude (degree)

0.0

-0.5
0.5 (b) 2.3-2.6 keV band

0.0

-0.5

1.0

0.5

0.0
359.5
Galactic longitude (degree)

359.0

Figure 2.5. X-ray images in the 2◦ × 1◦ field taken by Suzaku. (a) Continuum subtracted 6.7 keVline (Fe XXV) image. (b) 2.3-2.6 keV, including S XIII line, band image.

–10–

2.2. GALACTIC CENTER ENVIRONMENT

(a)
00 20

(b)

HII 0.81+0.18
SNR 0.3+0.0

15

NRFs

Radio Arc

Sgr C

Sgr E

GALACTIC LAT.

10
05
00
-00 051015
20
25

Sgr B2

Sgr B1The Snake

SNR 0.9+0.1

Sgr A

30
00 30

00
GALACTIC LONG.

HII 359.29-0.26-00 30-01 00Figure 2.6. (a) Radio image of the 2◦ × 1◦ GC region in the 20 cm band. Significant features arelabeled. (b) Schematic diagram of the distribution of non-thermal filaments in the GC. (Both fromYusef-Zadeh et al. 2004.)

magnetic filed and the cosmic-rays, the authors obtained the magnetic filed of 6 µG, whichis extremely lower than that obtained from the non-thermal filament. On the other hand,
Crocker et al. (2010) combined the gamma-ray upper limit taken by EGRET with the low
frequency radio data, and derived the lower limit of 50 µG for the inner 400 pc region. They
argued that the GC magnetic filed is not in equipartition.
Not only the poloidal structures seen in the non-thermal filaments but also toroidal
configuration is shown in near-infrared polarization (Nishiyama et al. 2010). Such toroidal
magnetic fields are also found in the dense molecular clouds (Novak et al. 2003). Therefore,
global configuration of the magnetic fields are under debate.

2.2.4

Star formation in the GC

From the observations of massive stars, Immer et al. (2012) estimated the star formation
rate (SFR) in the CMZ (450 pc region) to be 0.08 M yr−1 over the last 106 yr, which is ∼8%of the SFR in the entire Galaxy (∼1 M yr−1 , Robitaille & Whitney 2010) and is consistent
with the mass ratio of the CMZ to the entire Galaxy. This value is consistent with the mass
ratio of the CMZ to the molecular gas in the entire Galaxy. However, Longmore et al. (2013)claimed that this SFR is too small considering the correlation between density of MC andan SFR; much higher SFR is expected in the GC because the CMZ is denser than the disk(see Section 2.2.1). One possible explanation is that the large turbulent pressure of the CMZ
suppresses the star formation (Longmore et al. 2013).
In the inner 30 pc of the GC, there are three young star clusters, the so-called central,
Quintuplet, and Arches clusters, whose total mass is 5 × 104 M . Since the age of theseclusters are in the range of 2–7 Myr, star formation late at that time is estimated to be–11–CHAPTER 2. REVIEW
5 × 104 M /5 Myr = 0.01 M yr−1 (Figer et al. 2004). The authors also investigated thestellar luminosity function including intermediate and old stars in the central 30 pc of the
Galaxy, and found that continuos star formation rate of 0.02 M yr−1 between 107 and 1010
year ago represents the obtained luminosity function.
In contrast to the above continuous star formation rate, Yusef-Zadeh et al. (2009) investigated various SFR indicators in the GC, and reported that the SFR 105 –107 yr ago is∼0.1 M yr−1 while it decreases to ∼0.01 M yr−1 in the last 105 yr. Yasuda et al. (2008)
also claimed that the current star-formation activity is lower than in past based on theinfrared spectroscopy. Therefore, the star formation history in the GC is still controversial.

2.3

Supermassive Black Hole Sagittarius A*

2.3.1

Evidence for the central supermassive black hole

Based on the discovery of quasars in the 1960s, it is predicted that nearby galaxies, including the Galaxy, have SMBHs in their center (Lynden-Bell 1969; Lynden-Bell & Rees
1971). Radio observations detected the compact (sub-arcsecond) and high brightness temperature (107 K) source at the dynamical center of the Galaxy (Balick & Brown 1974). Sucha source is not seen in other part of the Galaxy. This peculiar object was suspected to be a
SMBH in our Galaxy, and has been referred to as Sgr A* (Brown 1982). However, there islittle signature of the SMBH in other-wavelength emissions. (see Section 2.3.2).
Gas and stellar dynamics near Sgr A* is a powerful tool to identify the Galactic SMBH.
The most plausible evidence of the central SMBH is provided by the measurements of theproper motions for the “S-stars” in the central cluster, which are located within 0.500 of Sgr
A* (Figure 2.7(a); Ghez et al. 2005). High-resolution near-infrared observations with laser
guide-star adaptive optics revealed the precise Kepler orbits of S-stars as shown in Figure
2.7(b) (Schödel et al. 2002; Ghez et al. 2005). In order to explain the observed orbits, there
should be a mass of ∼4 × 106 M within 0.01 pc at a distance of 8 kpc from the Sun (Ghezet al. 2008).

2.3.2

Emissions from Sgr A*

Radio
After the discovery of Sgr A*, many radio observations have been performed. The GHz
band flux of Sgr A* is variable by a few tens of percent on the time scales of years to days,
and the slope of its spectrum (Sν ∝ ν α ) also varies between α = 0.1 and 0.6 (Brown &–12–

2.3. SUPERMASSIVE BLACK HOLE SAGITTARIUS A*

(a)

(b)

Figure 2.7. (a) Central 100 × 100 K-band (2.2 µm) image taken by the W. M Keck I 10 m telescope.
Detected S-stars and the position of Sgr A* are lebeld. (b) Astrometric positions and orbital fitsof S-stars as a result of eight-year observations between 1995 and 2003. (From Ghez et al. 2005.)

Lo 1982). In the <1 GHz regime, turnover of the spectrum slope is suggested, but theentire Sgr A* region suffers from free-free absorption and no signal is detected at 330 MHz(Pedlar et al. 1989). Duschl & Lesch (1994) claimed that a time-averaged spectrum inthe 1–1000 GHz band shows α ∼ 1/3, which indicates synchrotron emission from monoenergetic electron distribution (Figure 2.8(a)). However, a simultaneous observation using
VLA, BIMA, Nobeyama, and IRAM revealed the excess in the millimeter to sub-millimeterband (Falcke et al. 1998; Melia & Falcke 2001). Figure 2.8(b) shows a spectrum including
infrared upper limits. We can see a hump structure, the so-called “submillimeter bump”.
Superior spatial resolution (<0.1 mas = 1 × 1013 cm) of radio interferometers enable us
to investigate the intrinsic size of Sgr A*, altough an observed size suffer from scattering
by ionized electrons in the ISM. The red points in Figure 2.9 shows the observed sizes of
Sgr A* as a function of wavelength. The black solid line is the λ2 scattering law derived byBower et al. (2006). The scattering effect dominates in the λ > 1 cm band while a source
size appears in the λ < 1 cm band. After correcting the scattering effect (green points), theintrinsic size of Sgr A* at 1.3 mm is 37 µas, which corresponds to a few of Schwarzschildradius ( 2GM/c2 ∼ 2 × 1012 cm ∼ 10 µas) for Sgr A* (Doeleman et al. 2008). The authors
also confirmed the λ1.4 dependence of the intrinsic size and suggested the stratified emission
region.
–13–

CHAPTER 2. REVIEW

(a)

(b)
α ~ 0.6

α ~ 0.28
submillimeter bump

Figure 2.8. Radio spectra of Sgr A*. (a) Adopted from Duschl & Lesch (1994). The verticalbars show the range of time variability. The dashed-dotted line represents the best-fit power-law
function with the index of ∼1/3. (b) Adopted from Melia & Falcke (2001). The data between
1–300 GHz were taken simultaneously (Falcke et al. 1998). There are the remarkable breaks of thespectrum at ∼60 GHz and ∼1000 GHz.

att
er
in

g

law

100

FWHM size (mas)

λ2
sc

10

Source-dominated1

Scattering-dominated
0.1

0.01

0.1

1
Wavelength (cm)10

Figure 2.9. The angular sizes of the major axis of Sgr A*. The solid line shows the λ2 scatteringlaw. The scattering effect dominates below this line. The red and green points are observed andscattering-corrected sizes, respectively. (From Doeleman et al. 2008.)–14–

2.3. SUPERMASSIVE BLACK HOLE SAGITTARIUS A*
X-ray
X-ray observatories had been expected to detect a loud activity, like AGNs, from Sgr A*.
However, Einstein, ROSAT, ASCA, and BeppoSAX reported the upper limit of <1036 erg s−1 ,contrary to the expectation (e.g., Watson et al. 1981; Predehl & Truemper 1994; Koyama
et al. 1996). Limited spatial resolution at most ∼ 1000 of ROSAT/PSPC makes it difficult
to eliminate contamination from nearby X-ray sources, and hence X-ray characteristics of
Sgr A* bad been veiled for a long time.
Chandra firstly resolved the X-rays from Sgr A* within 0.300 and found the surprisingly
low luminosity of 1033 erg s−1 , which is comparable to or even lower than other X-ray binaries(Baganoff et al. 2003, Figure 2.10(a)). The detected angular size is significantly larger than
the instrumental response with the intrinsic size of 1.00 4 (FWHM), which is consistent with theestimated Bondi-Hoyle accretion radius of 100 –200 (see Eq.2.3). The X-ray spectrum reported
by Baganoff et al. (2003) is represented by either a power-law function with the photon index(Γ) of 2.7 or an optically thin thermal plasma model with kT = 1.9 keV (Figure 2.10(b)).
The absorption column density is found to be 1 × 1023 H cm−2 .
Subsequent X-ray observations by Chandra and XMM-Newton have detected flares of
Sgr A*, but it reached at most 1035 erg s−1 (Porquet et al. 2003; Nowak et al. 2012). The
spectrum became harder (Γ ∼ 2) in the flare state, although the photon index is strongly
coupled with the absorption column density. The 3 Ms Chandra observation performed
during 2012 revealed that flares occurred every one day with the luminosity range of 1034 –1035 erg s−1 and lasted from a few hundred to eight kilo secounds (Neilsen et al. 2013). Thereis no difference in spectral hardness ratios between the faint and bright flares. All of the
X-ray flares associate with infrared flares (Markoff 2010).

Infrared
The identification of Sgr A* in the infrared band had been difficult because of confusion
with stellar sources or thermal dust emissions. We can only obtain upper limit in the midinfrared band (e.g., Serabyn et al. 1997; Cotera et al. 1999). In the near-infrared band,
recent development of adaptive optics enables us to resolve the Sgr A* emission. Ghez et al.
(2004) detected a variable point source within 18 mas of Sgr A* in the L0 (3.8 µm) band.
The observed luminosity varied in the range of 4–17 mJy on four separate nights. No othersources in this region show such large (factor of 4) and fast variation. The short timescaleof the flux variation indicates that the emission comes from close to Sgr A* within 5 AU(∼8 × 1013 cm).
–15–

CHAPTER 2. REVIEW

(a)

(b)

Figure 2.10. (a) Chandra 0.5–7 keV X-ray image of the central 10 × 10 field of the GC. The
position of Sgr A* is marked with the cross points. (b) X-ray spectrum of Sgr A* with the best-fit
thin thermal plasma model, where temperature and absorption column density are 1.9 keV and1.2 × 1023 H cm−2 , respectively. (Both from Baganoff et al. 2003)

Emission mechanism
Figure 2.11 shows the spectrum energy distribution (SED) of Sgr A*. The observed
bolometric luminosity of Sgr A* is ∼1036 erg s−1 , which is eight orders of magnitude lower
than its Eddington luminosity of 1044 erg s−1 (Figure 2.11; Genzel et al. 2010; Markoff 2010.
This Eddington ratio of 10−8 is even four orders of magnitude lower than that of a typical
low-luminosity AGN (LLAGN). What causes the extremely low activity of Sgr A*?
One possibility is a low accretion rate to Sgr A*. The IR observations discovered many
early-type stars in the inner parsec of the GC (Krabbe et al. 1991; Najarro et al. 1994;
Krabbe et al. 1995). The total mass-loss rate of the stars is ∼10−3 M yr−1 with the windvelocity (vw ) and density (nw ) of 500–2000 km s−1 and 103 –104 cm−3 , respectively (e.g.,
Martins et al. 2007). The simple Bondi-Hoyle accretion flow model predicts the accretion2
rate of ṀBH = πRAmH nw vw , where RA ≡ 2GM/vw2 is the accretion radius and mH is thehydrogen mass (Ruffert & Melia 1994; Coker et al. 1999). Therefor, the accretion rate fromthe stellar winds is


−2
M
vw
RA = 0.04
pc (∼ 100 )(2.3)
4.5 × 106 M
1000 km s−1

2 


RA
nw
vw
−4
M yr−1 .
(2.4)
Ṁ = 1.3 × 10
0.04 pc
103 cm−31000 km s−1
Then, assuming the radiative efficiency η = 0.1 as the standard accretion disk, we find the–16–

2.3. SUPERMASSIVE BLACK HOLE SAGITTARIUS A*
luminosity of Sgr A* to be
L = η Ṁ c2 = 5.6 × 1042

Ṁ
−4
10 M yr−1!
erg s−1 ,(2.5)

which is 6–7 orders of magnitude higher than the observed bolometric luminosity. Therefore,
the low luminosity of Sgr A* is not explained by the estimated accretion rate.
Another possible cause is extremely low radiative efficiency (η = 10−7 ). Such a radiatively
inefficient model is proposed by many authors, the so-called “advection-dominated accretion
flow (ADAF)” (see Narayan & McClintock 2008 for a review). Figure 2.11(a) shows themodel spectrum of the ADAF. Narayan et al. (1995) firstly proposed the ADAF model toexplain both the submillimeter bump and the low X-ray luminosity. The submillimeterbump is explained by synchrotron emission from thermal plasma in the inner accretion
flow, while X-rays come from bremsstrahlung of thermal plasma in the outer region. Thismodel marginally represents the observed SED, but underestimates the emissions in the
GHz band and the infrared band. Yuan et al. (2003) added the small fraction (∼1.5%)of non-thermal electrons as a tail of the thermal electron distribution. They argued thatsynchrotron emission from the non-thermal electrons successfully represents low-frequency
radio and infrared emissions.
Motivated by the discoveries of weak jets in the nearby LLAGNs, another model, where
jet is the predominant radiative source, is also proposed (see Markoff 2010 for a review).
This model well represents the radio to infrared SED including the GHz inverted spectrumby synchrotron emission from non-thermal electrons in a jet. X-ray emission in this model
is considered to be synchrotron self-Compton.
Thses models are not excluded by the observational data at this time, and hence theemission mechanism for Sgr A* is still controversial.

2.3.3

Evidence of past activity

As shown in the previous section, Sgr A* is very quiescent now. However, recent observations have discovered relics of past Sgr A* activities. We introduce these observational
evidence in this section.
X-ray reflection nebula
The X-ray observation by ASCA discovered strong 6.4 keV (Fe I) line in the GC region(Figure 2.4; Koyama et al. 1996). As shown in Figure 2.12, spatial distribution of the 6.4 keVlines is clumpy and coincides with the dense MCs found by the CS (J=1–0) line observation–17–

CHAPTER 2. REVIEW

2
Quiescent spectrum: “Leaky” RIAF
Quiescent spectrum: Jet
Flaring SSC component

0
-2
-4
-6
-8
-10

0

2

4

6

8

10

Figure 2.11. (a) Spectral energy distribution of Sgr A* with the ADAF model constructed by Yuan
et al. (2003). The dashed, dotted, dashed-doted lines represent individual emission mechanisms,
and the solid line shows the total emissions. (from Genzel et al. 2010.) (b) The same as (a) but
with the jet model. You should note that the vertical axis is not νFν but Fν . (from Markoff 2010.)(Tsuboi et al. 1999). Since MCs are too cold to emit X-rays, they should be irradiated withhigh energy photons or cosmic-rays.
The particles exciting the cold MCs had been debated for a long time. Figure 2.13 shows
the typical spectrum of the 6.4 keV clump derived from the Sgr A region. The equivalent
width (EW) of the 6.4 keV line with respect to the associated hard continuum is ∼1 keV.
It is consistent with the predicted EW of the photon or proton origin (∼1 keV) rather thanthat of electron origin (∼400 eV). The strong Fe absorption edge in the hard-continuum
indicates the photon origin rather than the proton origin. Nobukawa et al. (2010) found
not only the 6.4 keV line but also the neutral Ar, Ca, Cr, Mn, and Ni lines (Figure 2.13).
The EWs of all the lines are consistent with the photon origin. In addition to the spectrum
shape, Inui et al. (2009) discovered that the flux and morphology of 6.4 keV lines in the Sgr
B2 molecular cloud, whose spatial scale is a few parsec (∼10 light yr), changed in 10 yr.
This rapid variation is only achieved by photon irradiation, because charged particles whichmainly contribute to excitation of neutral Fe (∼10 keV for electrons and ∼10 MeV for
protons) have significantly lower velocity than the light speed. Nobukawa et al. (2011) also
confirmed that the 6.4 keV flux and associated hard-continuum in Sgr B2 deceased by halfin between 2005 and 2009, while no time variation on the 6.7 keV line flux was found (seeFigure 2.14). Therefore, the photon irradiation is favor for the 6.4 keV clumps. We refer tothe clumps as the X-ray reflection nebula (XRN).
From the intensity of 6.4 keV lines, the required luminosity for X-ray sources illuminating
the MCs is estimated to be ∼1037 × (D/10 pc)−2 erg s−1 , where D is the distance between–18–

2.3. SUPERMASSIVE BLACK HOLE SAGITTARIUS A*

Galactic latitude (degree)

0.2(a) 6.4 keV
Sgr B20.0

Sgr B1-0.2

0.2

Sgr A*Sgr C

M359.23-0.04

Sgr A(b) CS (J=1-0)
Sgr A*

0.0
-0.2
1.0

0.5
0.0
359.5
Galactic longitude (degree)

359.0

Figure 2.12. (a) Suzaku 6.4 keV image of the 2◦ × 0.8◦ field of the GC. (b) Image of CS (J=1–0),
which traces dense MCs. (From Tsuboi et al. 1999.)

the XRN and the illuminating source. One of the XRNe, M359.23−0.04, can be illuminatedby the nearby low-mass X-ray binary 1E 1740.7−2942 (Nakashima et al. 2010). However, theother XRNe have no bright nearby sources providing this luminosity. A remaining possibility
is past Sgr A* flare of >1039 erg s−1 .

When we have a line-of-sight distance toward an XRN and a distance between the XRN
and Sgr A*, we can derive a period of past Sgr A* flare. In order to estimate the lineof-sight position of the XRNe, Ryu et al. (2009) developed a new method using the partial
absorption of the 1 keV plasma in the GCXE. Applying this method to some 6.4 keV clumps,
the authors determined the line-of-sight position the XRNe (Figure 2.15(a)). Based on thisresults, past light curve of Sgr A* is derived as shown in Figure 2.15(b); Sgr A* had been
continuously active with 1039 erg s−1 100-1000 yr ago.
–19–

CHAPTER 2. REVIEW

Figure 2.13. Spectrum of the 6.4 keV clump in the Sgr A region. The red lines represent thecontamination of the 1 keV and 7 keV plasmas in the GCXE (see Section 2.2.2). The blue lines are
the emissions from neutral atoms and the associated continuum emission. (From Nobukawa et al.
2010.)

0.1

b

0.0

(a)

Sgr B2

M0.74-0.09

b

-0.1
-0.2

0.1

0.7

l

0.6

0.0

-0.1

0.8
0.1

b

-0.1
-0.2

0.7

l

0.6

0.5

Hard continuum
2009

(d)

0.0
-0.1
-0.2

0.8
0

0.0

0.5

Hard continuum
2005

(c)

FeI Kα
2009

(b)

-0.2

Background
0.8

b

0.1

FeI Kα
2005

0.7

l

0.6

0.5

0.8

0.5

0.7

l

0.6

0.5
1.0

Figure 2.14. Suzaku X-ray images of the Sgr B2 region. The upper panels shows the 6.4 keVband image in (a) 2005 and (b) 2009. The lower panels shows the hard-continuum band image in(c) 2005 and (d) 2009. (From Nobukawa et al. 2011.)–20–

2.3. SUPERMASSIVE BLACK HOLE SAGITTARIUS A*
5

600
T*
(yr)

Y (ly)

200
0

C2
D
B2

-200

B1
M
0.74

-400
To
Sun

-600
600 400 200

800

3.5

600

3
2.5

Sgr A*
C3

4

400
C1

50

200
E

0 -200 -400 -600
X (ly)

2
1.5
1

LX* (erg s-1)

Y=–X2/2cT*+cT* /2

I6.4 keV (e-6 photon/cm2/s/arcmin2)

400

1040

4.5

E

M0.74

C1

B2

D
C2

1039

B1

C3

0.5

-700

-600

-500

-400
-300
-200
T * (years from 2010)-1001033–350Figure 2.15. (a) Map of the XRN positions in the face-on-view constructed by Ryu et al. (2013).
(b)Light curve of the past Sgr A* flares. (From Ryu et al. 2013.)

Fermi bubbles
All-sky survey performed by the Fermi gamma-ray observatory discovered large bipolar
emissions, which has a bubble-like morphology extending 50◦ north and south of the GC withthe width of 40◦ (Figure 2.16(a)). Therefore, they are referred to as the “Fermi bubbles”.
Surface brightness in the bubbles is nearly uniform, so that bubble shape is neither simple
shell nor sphere. As shown in Figure 2.17, Gamma-ray spectrum of the bubbles has spectral
photon index of ∼2, which is harder than that of the disk region. No significant spatial
variation of the spectrum in both of the north and south bubbles is found. Although it isdifficult to determine the accurate distance toward the Fermi bubbles, their considerably
symmetric morphology to the GC strongly suggests that the a past GC activity formed the
Fermi bubbles.
There are some features possibly associated with the Fermi bubble in other wavelength.
The North polar spur, which is seen in the low-frequency radio image (Figure 2.16(b)) andthe X-ray image (Figure 2.16(c)), is located eastern edge of the northern Fermi bubble.
Kataoka et al. (2013) proposed that the North polar spur was formed by a weak shock
driven by the bubble. On the other hand, the Loop I , which is also seen in the radio and
X-ray images, is less associated with the gamma-ray emissions. The GHz band radio image(Figure 2.16(d)) shows the biconical features above and below the GC, the so-called “WMAP
haze”. It coincides with the foot points of the Fermi bubbles. The WMAP haze exhibits hard
spectrum probably due to synchrotron emission from harder electron distribution (Figure
2.17(b)). Assumed the Combined spectrum of the Fermi bubbles and the WMAP haze ismarginally represetned by single power law electron energy distribution with the index of
2.0–2.5 (Figure 2.17): the gamma-rays are emitted by inverse Compton with the Galactic
radiation fields and radio continuum is emitted by synchrotron with the magnetic filed of–21–

CHAPTER 2. REVIEW

Fermi 1-5 GeV

3.0

40

50

50

30

0

0

20

1.0
-50

-50

2.0
1.5

10

0.5
50

0

-50

50

Rosat Band 6 and 7

400

0

-50

WMAP 23 GHz haze

250
200

50

0.06
0.04

50

0.02
0

0

0.00

150
-50

-50

-0.02

100
50

0

mK

10-6 counts s-1 arcmin-2

350
300

K

keV cm-2 s-1 sr-1

2.5

Haslam 408 MHz

50

0

-50

50

0

-50

-0.04

Figure 2.16. Images of the Fermi bubbles and possibly associated features in other wavelength.
(a) Fermi image in the 1–5 GeV band. The dotted green, blue, and red curves show the bubble
features, the North polar spur, and Loop I, respectively. (b) Haslam 408 MHz band image. (c)ROSAT 1–2 keV band image. (d) WMAP 23 GHz band image. (From Su et al. 2010.)–22–

2.4. SUBJECT OF THIS THESIS

Inverse Compton Scattering
0.1 GeV to 1000.0 GeV-5

101000.00
γ=1.5
γ=1.8
γ=2.4
γ=3.0

10-6

10-7
10-8
0.1

Synchrotron

WMAP 23 GHz

10.00
1.00
γ=1.5
γ=1.8
γ=2.4
γ=3.0

0.10

1.0

10.0
Eγ [GeV]

100.0

B= 10.0μG

100.00Iν [kJy/sr]

E2 dN/dE [GeV cm-2 s-1 sr-1]

10-4

1000.0

0.01
0.1

1.0

10.0
100.0
Frequency [GHz]1000.0Figure 2.17. (a) Spectrum of Fermi bubbles. (b) Spectrum of WMAP haze. (From Su et al.
2010).

∼10 µG. This model implies the total energy of the bubbles reaches 1055 –1056 erg (Carretti
et al. 2013; Kataoka et al. 2013).
The origin of the Fermi bubble is still under debate. Two possible scenario is suggested
at present: a Sgr A* jet or a star-formation-driven outflow. In the context of a Sgr A* jet, aflare of sub-Eddington luminosity (1042 erg s−1 ) a few Myr ago is proposed (Su et al. 2010;
Zubovas et al. 2011). In the scenario of a star-formation-driven outflow, a wind velocity is
estimated to be 1000 km s−1 , and then the dynamical timescale is 10 Myr (Su et al. 2010;
Carretti et al. 2013). If this is the case, energy injection of 1041 erg s−1 is required to form
the Fermi bubbles.

2.4

Subject of This Thesis

As shown Section 2.3.3, recent high-energy observations have been unveiled the past GC
activities. However, the samples are still limited and there is a large gap in the activity
history between 103 and 106 years ago (see Figure 2.18). Our aim is to find other relics ofGC activities and to reveal the history of them.
Past X-ray studies of the GC are mostly concentrated on the Galactic plane region asreviewed in this chapter. However, a signature of past GC activity may be conceal by thestrong GCXE in the on-plane region. Therefore, We focused on the Galactic bulge region,
which is less contaminated by the GCXE and is more suitable to extract pure phenomena
associated with the GC activities. We performed the survey of this region with the X-ray
observatory of Suzaku, whose details are represented in Chapter 3.

–23–

CHAPTER 2. REVIEWFBRelased energy rate (erg s-1)

1042

(Sgr A* jet origin)

FB
1040 (Starburst origin)

?

1038

XRN

10% SN rate(r < 50 pc)
Current
Sgr A*

1036

Flare

1034
Quiescent

1032

107

106

105

104

103

102

101

Look back time (years)

Figure 2.18. Schematic diagram of past GC activities.

–24–Chapter 3
Instruments
In this chapter, we review Suzaku and its onboard instruments, which were used for theobservations of the Galactic center.

Contents
3.1

Overview of Suzaku . . . . . . . . . . . . . . . . . . . . . . . . . . .

26

3.2

X-ray Telescope (XRT) . . . . . . . . . . . . . . . . . . . . . . . .

27

3.2.1

X-ray optics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

27

3.2.2

Telescope design . . . . . . . . . . . . . . . . . . . . . . . . . . . .

27

3.2.3

In-flight performance . . . . . . . . . . . . . . . . . . . . . . . . . .

28

X-ray Imaging Spectrometer (XIS) . . . . . . . . . . . . . . . . .

31

3.3.1

Overview of the XIS system . . . . . . . . . . . . . . . . . . . . . .

31

3.3.2

Device description of the CCD camera . . . . . . . . . . . . . . . .

33

3.3.3

In-flight performance and calibration . . . . . . . . . . . . . . . . .

36

Advantage of Suzaku in the GC survey . . . . . . . . . . . . . . .

41

3.3

3.4

25

CHAPTER 3. INSTRUMENTS
(a)
5.4 m

6.5 m

(b)

(c)

570 kmFigure 3.1. (a) Appearance of Suzaku in the orbit. (b) Schematic view of the configuration of
instruments and satellite systems. (c) The Suzaku orbit.

3.1

Overview of Suzaku

Suzaku is the fifth japanese X-ray astronomy satellite lunched in 2005 July 10 fromthe Uchinoura Space Center (USC) by Japan Aerospace Exploration Agency (Mitsuda et al.
2007). Figure 3.1 shows a schematic view of Suzaku and its orbit. The length of the spacecraft
is 6.5 m after extending the extensible optical bench (EOB). The total weight is ∼1700 kg. It
was put into a near-circular orbit at 570 km altitude with an inclination angle of 31 (Figure
3.1(c)). An orbital period is 96 minuets. The low earth orbit has a merit in less exposureof cosmic-rays, while celestial targets are occulted by the Earth for one third of each orbit
except for targets near the orbital poles. Tracking and data down-link are managed in the
USC. The satellite passes over the USC 4–5 times a day for ∼10 minutes per pass. Attitudeof the satellite is controlled by two star tracker systems, and ∼10 minuets are required todirect the satellite toward a target.
–26–

3.2. X-RAY TELESCOPE (XRT)
Suzaku has a capability of high-sensitivity and wide-band spectroscopy. It carries two distinct scientific instruments: the X-ray Imaging Spectrometer (XIS; Koyama et al. 2007c) andthe Hard X-ray Detector (HXD; Takahashi et al. 2007). The XIS is imaging-spectrometerusing a X-ray CCD camera, which is mounted on the focal plane of a dedicated X-ray Telescope (XRT; Serlemitsos et al. 2007). The XIS has a large effective area and moderate energy
resolution in the 0.2–12 keV band. The HXD is non-imaging collimated detector for hard
X-rays covering the bandpass of 10–600 keV. The XIS and the HXD work in complementaryand covers the wide bandpass. Suzaku has also carried the micro-calorimeter system, named
the X-ray Spectrometer (XRS), but the XRS unfortunately failed to work on 2005 August
8 because of liquid helium loss. In the following sections, we described details of the XRT
and the XIS, which were mainly used in this thesis.

3.2

X-ray Telescope (XRT)

3.2.1

X-ray optics

To focus visible light, we usually use refraction by lenses whose refraction index is muchlarger than 1. However, a refraction index depends on wavelength, and it becomes nearly 1(and smaller than 1) in the X-ray band. It is because that we are not able to use lenses forfocusing X-rays. Therefore, X-ray optics used either reflection, diffraction or interference.
X-ray telescopes commonly use the reflection by mirrors at very shallow angles of <1◦(grazing incidence reflection). Based on this idea, Wolter (1952) proposed the some designs
composed of two-stage reflection mirrors realizing a wide filed-of-view (FOV). The mostpopular design for X-ray astronomy is the Wolter type I telescope, where primary andsecondary mirrors are aligned with paraboloid and confocal hyperboloid, respectively as
shown in Figure 3.2. Since one grazing incidence mirror has very small collecting area,
a nesting technique of thin mirrors is necessary to increase a photon-collecting capability.
However, nesting causes degradation of angular resolution because accurate alignment of
all nested mirrors is difficult. There is a tradeoff between an effective area and angularresolution.

3.2.2

Telescope design

The Suzaku XRT is an Wolter type I X-ray telescope consisting of closely nested thin-foilreflectors, which realize a light telescope with a large effective area. In the XRT, the shapesof the primary and secondary reflectors, which are originally a paraboloid and hyperboloid,
respectively, are both approximated by cones. This approximation limits the spatial resolution as l tan θ, where l and θ are the axial length and the grazing angle, respectively. For the–27–

CHAPTER 3. INSTRUMENTS

Figure 3.2. (a) Schematic view of an Wolter type I telescope consisting of nested thin-foil reflectors: (a) front view and (b) the cross section of the telescope (Mori et al. 2005).

Suzaku configuration, it is calculated as 1800 in half power diameter (HPD). This degradationis negligible compared to other degradation factors such as the alignment of mirrors.
Figure 3.3(a) is a photo of one XRT module, which contains 1400 reflectors with 175
nested shells. The telescope has a weight of 19.7 kg, a height of 279 mm, a geometrical areaof 873 cm2 , and a focal length of 4.75 m. An innermost and an outermost diameters of thenested shells are 118 mm and 399 mm, respectively. An incident angle of each reflector is
between 0.◦ 18–0.◦ 60. Each reflector was produced through a replication method in which athin (∼0.1 µm) gold layer sputtered on a highly smooth glass tube was transferred to analuminum substrate (152 µm thickness and 101.6 mm axial length) with an epoxy adhesive(25 µm thickness). Figure 3.3(b) shows a photo of reflectors.
In a incident side of the telescope, a pre-collimator to mitigate stray light (see following section) is mounted on, and covered by a thermal shield to isolate the XRT thermallyfrom space and to reflect infrared radiation from the interior of the spacecraft. The XRT
temperature is kept at ∼20 ◦ C.

3.2.3

In-flight performance

Effective Area
Figure 3.4(a) shows the simulated effective area of the one XRT module as a functionof incident X-ray energy. Strong absorption edges in the 0.2–2 keV band comes from transmittance of the thermal shield, which consists of a 0.24 µm polyethylene terephthalate film
(C10H8O4) coated with an 30 nm aluminum layer. Other absorption edges in the 2–3 keV–28–

3.2. X-RAY TELESCOPE (XRT)(a) XRT module(b) Reflectors

Figure 3.3. Photos of (a) one XRT module and (b) reflectors nested in the XRT.

band and at 12 keV are due to Au coated on the reflectors. Using a observation of the
Crab nebula, the absolute photon flux, photon index, and absorption column densities werechecked. The results are consistent with the standard value compiled by Toor & Seward(1974).
The effective are of the XRT decreases with increasing the off-axis angle because incident
X-ray are partially obscured by the backside of the neighbor reflectors. This effect is more
sever at higher energies. It is called the “vignetting” effect. The amount of vignetting effectis also calibrated by the observations of the Crab nebula with various off-axis angles. The
resultant curve is shown in Figure 3.4(b). We can see steeper degradation for higher energy
X-rays. At the 80 off-axis angle, which corresponds to the edge of the XIS detector, theeffective area is half of that at the on-axis. The effective are is calibrated within ∼ 5% overthe FOV.

Angular resolution
The imaging capability of the XRT was investigated by the observation of SS Cygni,
which is a moderately bright point source suitable for this calibration. Figure 3.5(a) shows
the observation image. The shape of the image is asymmetric compared with the ray-trace
simulation shown in Figure 3.5(b), even if the ray-trace simulation is fine-tuned for eachquadrant. The spiky structures in the observation image is due to shadows of the supportingbars and misalignments of the reflectors.
Figure 3.5(c) and Figure 3.5(d) show the point spread function (PSF) and encircled energyfunction (EEF) derived from the observational result, respectively. The large difference
–29–

500

CHAPTER 3. INSTRUMENTS(a) Energy dependence of the effective area(b) Curve of the vignetting effect

C K-edge200

Au M-edges
O K-edge

100

Effective Area (cm−2)

Al K-edge

50

Au L-edge
0.2

0.5

1
2
Energy (keV)

5

10

Figure 3.4. (a) Simulated on-axis effective are of one XRT module as a function of X-ray energy.
Obvious absorption edges are labeled. (b) Curves of the vignetting effect. The count rate derived
from the Crab nebular observations are shown in symbols, and the results of ray-trace simulations
are shown in the solid lines.

between the observation and the simulation nearby the image center is due to the attitude
error of the satellite control Uchiyama et al. (2008). The EEF in the range of 10 –60 iscalibrated within 4%. The half power diameter of the XRT is ∼2.◦ 0

Stray light

The closely nested structure causes unexpected reflections of X-rays coming from the offaxis angle, the so-called “stray light”. As shown in Figure 3.6, the two passes are possible.
One is the secondary reflection where incident photons are not reflected by a primary mirror
but only by a secondary mirror, and another is the backside reflection where photons arereflected by the backside of a primary mirror. The stray light become serious problem when
we observe a weak emission with a bright point source out of FOV.
For the Suzaku XRT, the pre-collimator is installed on the incident side, and then thestay light is remarkably suppressed compared to the ASCA X-ray telescope, which has thebasically same structure except for the pre-collimator Serlemitsos et al. (1995). As shown in
Figure 3.7, the stray light due to the secondary reflection is fairly reduced with pre-collimator,
but still remains at 200 and 600 off-axis angles (Mori et al. 2005).
–30–

3.3. X-RAY IMAGING SPECTROMETER (XIS)(a) Observation image(b) Ray-trace simulation image(c) Point spread function(d) Encircled energy function

Figure 3.5. (a) XIS image taken from the observation of SS Cygni. (b) Ideal image obtained froma ray-trace simulation. (b) Point spread function and (d) encircled energy function. For both thefigures, black and green lines are the observational results and simulated results, respectively. The
lower panels show the residuals between the observation data and the simulation data.

3.3

X-ray Imaging Spectrometer (XIS)

3.3.1

Overview of the XIS system

The XIS is X-ray imaging-spectrometers consisting of four set of CCD camera systems.
The system has been developed and maintained by a collaboration of Japan (ISAS/JAXA,
the University of Kyoto, Osaka, Miyazaki, and Kogakuin) and the United States (the Massachusetts Institute of Technology (MIT)).
Figure 3.8 is a photo and a schematic view of one camera system. The system has aweight of 48.7 kg and consumes 67 W at a bus voltage of 50 V during normal operations.
The CCD cameras are cooled at −90 ◦ C by thermoelectric coolers to reduce dark current.
–31–

CHAPTER 3. INSTRUMENTS

Figure 3.6. Schematic view of the incident X-ray passes. (a) normal double reflection, (b)secondary reflection, and (c) backside reflection are illustrated. The stray light is formed by thepasses of (b) and (c).

Figure 3.7. Ground measurement of the stray light with and without the pre-collimator (Moriet al. 2005).

–32–

3.3. X-RAY IMAGING SPECTROMETER (XIS)

Figure 3.8. Photo (left) and a schematic view (right) of one camera system of the XIS.

The incidence of optical and UV photons are prevented by an optical blocking filter (OBF),
which is installed on the surface of the CCDs. Calibration sources of 55 Fe, whose half life is
2.73 years, are located above the CCDs and illuminate corners of them to monitor in-flight
performance.
The CCD is clocked by the analog electronics (AE) system, which supplies the various
voltages in the CCD. Obtained analog X-ray signals are also transferred to the AE and converted to digital signals. The digitized signals are then transferred to the digital electronics(DE) system, where raw pixel data are temporally stored, event detection is carried out,telemetry packets are made, and finally sends to the main digital processor of the satellite.

3.3.2

Device description of the CCD camera

X-ray CCDs have two-dimensional array of photo sensitive sells (pixels) which produce
a number of electron-hole pairs proportional to the energy of incident X-ray photons. After
exposure, signal charges are transferred to the readout gates clocking the gate voltage. Since
CCDs achieve the fine pixel pitch (∼20 µm) and the fano-limited energy resolution (∼130 eV
at FWHM for 6 keV X-rays), they are widely used in the modern X-ray observatories: Suzaku,
Chandra, XMM-Newton, and Swift.
The XIS used a MOS-type frame transfer CCD fabricated by the MIT Lincoln Laboratory.
Transfer clocking is three phase operation. It a format of 1024 × 1024 pixels with a pixels–33–

CHAPTER 3. INSTRUMENTS
pitch of 24 µm (Figure 3.9). Combined with the XRT, the XIS has a 17.0 8 × 17.0 8 field-of-view.
The frame-store area is divided into four segments, which have own readout gates. The
charge transfer rate is fixed at 24.4 µs pixel−1 . In the normal mode, 24.4×1024×(256+20) ∼6.9 s (including under- and over-clocked pixels) is required to readout the stored signals froma gate in each segment. The timing resolution is limited by this charge transfer time. In
order to increase the timing resolution, XIS CCD has window options (1/4 and 1/8) and a
P-sum mode. In the window options, effective pixels are limited to 1/4 or 1/8 of the entire
area in order to reduce the charge transfer time. In the P-sum mode, signals in a column
is summed in one pixel, so that the charge transfer time is fairly reduced while imaging
capability is lost.
Three CCDs are front-illuminated (FI) type (XIS0, 2, and 3) and the other one is backilluminated (BI) type (XIS1). The quantum efficiency (QE) of both type of CCDs is shownin Figure 3.10(a) The FI CCD is superior in the high energy band for a thick depletion layerof ∼65 µm, but its entrance window has a MOS gate structures made of 0.28 µm Si and0.44 µm SiO2 , which severely degrades the QE below 1 keV. On the other hand, the BI CCD
has very thin entrance window (5 nm HfO2 , 1 nm Ag, and 3 nm SiO2 ), which is made by
the chemisorption charging process, and realizes high QE for low energy X-rays. However,back-thinning process to achieve the full-depletion reduce the thickness of effective layer to∼42 µm. Therefore, the FI and BI CCDs work complementary. Effective areas combinedwith the XRT are also shown in Figure 3.10(b). The chemisorption charging process also
improved the charge correction efficiency for low energy X-rays, which are absorbed at thesurface of the CCD. Therefore, spectral shape of the XIS BI CCD is fairly better than thatof Chandra/ACIS and XMM-Newton/EPIC (Burke et al. 2004).
In the space environment, charge transfer inefficiency (CTI) in a CCD gradually increases,
because charge traps are made due to bombardment of high energy cosmic-rays (e.g., Yamashita et al. 1997). In order to mitigate the radiation damage, the XIS CCD is equippedwith function to inject artificial charges. A charge injection gate is located at the opposite
side of the readout gate. Before an exposure, an arbitrary amount of charges (∼1000 e− )are injected into every 54th rows. These artificial charges fill traps in charge transfer, andprevent signal charges to be captured by the traps. This technic called the “spaced-raw
charge injection (SCI)”.
XIS2 suddenly showed anomaly at 1:03 UT on Nov 9, 2006. About 2/3 of the image wasflooded with a large amount of charge, which was leaked somewhere in the imaging region. It
is probably caused by a micro-meteorite impact. Thus, entire XIS2 have not been functional
for scientific observations. A possible micro-meteorite impact also occurred at 2:00 UT on
June 23, 2009, and a part of segment A of XIS0 have not been functional. Since the othersegment of XIS0 was not affected by this anomaly, we used the 3/4 of XIS0 for analysis in
this thesis.
–34–

3.3. X-RAY IMAGING SPECTROMETER (XIS)1024 pixels

Charge injection
gate

A

B

C

D

256, 128 or 64

1024 pixels

Exposure Area(no physicalboundaries
betweensegments)

XIS-Sensor

pixels

Optional Window

XIS-AE/TCE

Act Y

control temperaturedrive CCD
read-out data
send data to DE

Act X

Read-outgates

256
pixels256
pixels

3
1

4
2

4
2

3
1

256
256
pixels pixels

3
1

4
2

4
2

1024 pixels

Frame Store
Area

Read-out
gates

3
1

Read-out (through AE)

Raw Y(0-1023)

XIS-DE

1

2

1

2

1

2

1

2

3

4

3

4

3

4

3

4

ActivePixels(256
columns)

ActivePixels(256
columns)

Satellite
Data
Processor(DP)

ActivePixels(256
columns)

ActivePixels(256
columns)

H-OverClocked
Pixels(16 columns)
Memory(Pixel RAM)

Event
Detection
Raw X

Dummy pixels(two columns)

Raw X

Raw X

Raw X

Dummy pixels(two columns)
Copied Pixels(two columns)

Figure 3.9. Schematic overview of the XIS system.

–35–

Data
Edit
Main
Processing
Unit
(MPU)
Event
RAM
Pixel
Processing
Unit
(PPU)

CHAPTER 3. INSTRUMENTS
(a)

(b)

Figure 3.10. (a) The QE as a function of incident energy, calculated using best estimate valuesof the thickness of dead layers and depletion layer. The solid line represents the FI CCD (XIS 0)
and the broken line is for the BI CCD (XIS 1). (b) The effective are of the XIS combined with the
XRT. The black and red lines represents the FI CCD and BI CCD, respectively.

3.3.3

In-flight performance and calibration

SCI operation

Increasing CTI degrade the signal pulse hight. Since the amount of the degradation
varies from pixel to pixel, the width of the pulse hight also increase. Figure 3.11 shows thetime history of the pulse hight and its width monitored by the 55 Fe calibration source. Just
after the launch, the signal pulse hight of the calibration source is calibrated to be 5894 eV(Mn Kα line), and the energy resolution is ∼130 eV. However, both degrade rapidly between
the launch and the mid 2006.
The SCI technique has been applied since August 2006, and the gain and the energy
resolution are successfully recovered to the performance of nearly just after the launch.
Although performance gradually degrades even in the SCI mode, its slope is considerably
small compared to the operation without the SCI. The slopes are different among FIs (XIS0and XIS3) and BI (XIS1) from the mid 2006 to the late 2010. It is due to the differenceof the amount of injected charges: ∼1640 e− (equivalent to 6 keV signal) for the FIs and∼550 e− (equivalent to 2 keV signal) for the BI. Increasing charge traps increase the CTI.
For the BI sensor since 2011, the amount of injected charges increase to the same amountfor the FI sensors, and then the performance became comparable to the FIs.
–36–

3.3. X-RAY IMAGING SPECTROMETER (XIS)

Figure 3.11. Long term trend of 55 Fe centroid energy and width before trail and CTI correction.
The data after calibration are shown in Figure 3.13.

CTI correction
As shown in Figure 3.11 pulse hight degradation occurs even applying the SCI technic,
and hence the energy scale calibration from pixel to pixel is necessary to bring out the best
performance of the XIS. Figure 3.12 shows a pulse hight profile along the charge transfer
number (row pixel number) shows the saw-tooth like shape at 54 pixel intervals. It isbecause that trapped artificial charges are detrappted in a certain time, and the efficiencyof the charge injection decreases with increasing the distance from the injected row.
In order to correct the CTI including the saw-tooth like profile, Uchiyama et al. (2009)
composed a model of the observed pulse hight (P H) as a function of the elapsed time (t),row pixel number (k), and intrinsic pulse hight (P H0 ):P H(t, i, P H0 ) = [1 − c1 (t, i, P H0 )]i [1 − c2 (t, P H0 )]1024 P H0
' [1 − i × c1 (t, i, P H0 )] [1 − 1024 × c2 (t, P H0 )] P H0

(3.1)
(3.2)

where c1 and c2 are the CTI per one transfer due to the slow and the fast transfer, respectively.
Assuming c1 and c2 is  1, Eq.3.1 is approximated by Eq.3.2. Then, c1 and c2 are represented
as below:
c1b (t) − c1t (t)c1 (t, i, P H0 ) = c1t (t) +(i mod 54) P H0−β(3.3)
54c2 (t, P H0 ) = c02 (t)P H0−β
–37–

(3.4)

CHAPTER 3. INSTRUMENTS

Figure 3.12. Schematic view of (a) CTI1 and (b) observed pulse hight as a function of the chargetransfer number.

where c1t , c1b are c1 values at the peak and valley of the sawtooth normalized by a certain
P H0 , c02 (t) is c2 values for a certain P H0 , and β is a parameter showing dependence of the
CTI on P H0 . Sawada et al. (2012) improved this model introducing the time-variable β
and the discontinuous slope of c1t and c1b . Figure 3.13 shows the the center of emission
lines after the correction of the CTI using the above model. The data well agree with thereference values.
As shown in Figure 3.14, the energy resolution is considerably improved after the energy
scale correction compared with Figure 3.11. The observed width of F W HM is given byq
2
2
F W HM =+ F W HMadd(3.5)
F W HMinit
where F W HMinit is the width at the launch and F W HMadd is the additional width due tothe CTI degradation. F W HMadd is then modeled as a function of the X-ray energy (E) andthe time after launch (t) as below.
F W HMadd =q

c20 + {c1s (t − t0 ) + c1o }2 E + {c2s (t − t0 ) + c2o }E 2

(3.6)

where t0 is the time at the launch and c0 , c1s , c1o , c2s , c2o are parameters larger than 0. The
obtained model is also shown in Figure 3.14 by the solid line.

Contamination
Figure 3.15 shows spectra of RX J1856.5−3754 observed in every six months. We can see
significant degradation of the detection efficiency below 1 keV. It is caused by the absorption
due to contaminating materials accumulated on the OBF, which is cooler than the other part–38–

3.3. X-RAY IMAGING SPECTROMETER (XIS)5920

Blue: Segment A Red: Segment D

XIS 0

Mn K Line center (keV)

5900
5880
5920

XIS 1

5900
5880
5920

XIS 3

5900
5880

2008

2010
Time (year)

2012

Figure 3.13. Time variability for the line center of the Mn Kα line (55 Fe calibration source)
after the CTI correction. The top, middle, and bottom panels indicates XIS0, XIS1, and XIS3,
respectively.

150

5.9 keV line

250

Line width at FWHM (eV)Line width at FWHM (eV)300

TrendArchive MnK
XIS 3, Seg. A (red) D (blue)

200
150
100

2007

2

2.5

2008

2009

3
Time (108 s)

2010

2011

3.5

4

E0102 NeHe
XIS 3, Seg. B (orange) C (green)

0.92 keV line
100

50

0

2007

2

2.5

2008

2009

3
Time (108 s)

2010

2011

3.5

4

Figure 3.14. Time variability for the line width of (a) Mn Kα line and (b) He-like Ne Kα line
after the CTI correction. The solid lines show models of the line width.

–39–

CHAPTER 3. INSTRUMENTS

normalized counts s−1 keV−1

RXJ1856 xissimarfgen−with−HCNO−model(v2012−07−19)
2005Oct
2006Mar
2006Oct
2007Oct
2008Mar
2008Oct1

0.1
2009Oct
2010Mar2010Oct
2011Oct
2012Apr

0.01
2

ratio

1.5
1
0.5
0

0.5
Energy (keV)

1

Figure 3.15. XIS spectra of RX J1856.5−3754 observed from 2005 Oct. to 2012 Apr.

of the satellite. Therefore, absorption column density and chemical composition of the contamination are measured using the observation data of RX J1856.5−3754, PKS 2155−304,
and 1E 0102.2−7219. The resultant optical depth as a function of time is shown in Figure
3.16. For 0.65 keV X-rays, the detection efficiency is 0.25 times of that at just after launch.

Non X-ray backgroundIn the orbit, comic-rays produce background signals, the so-called the non X-ray background (NXB). Most of the NXB is due to direct ionization loss of cosmic-rays. Such events
make long tailed images, and thus they are eliminated by the imaging analysis referred toas the grade selection. However, fluorescence X-rays generated by the interaction between
cosmic-rays and the camera body of the XIS are not distinguished from the celestial X-rays.
Figure 3.17 shows the NXB spectra after the grade selection. We can see fluorescence
X-rays of Al, Si, Ni, and Au, which are used for the camera components. There are also Mn
K shell emission lines coming from the 55 Fe calibration source. The constant components
are due to Compton events. The large bump-like feature is also seen in the XIS1 spectrum.
It is due to low-energy comic-rays.
–40–

3.4. ADVANTAGE OF SUZAKU IN THE GC SURVEY

Figure 3.16. Optical depth of the contamination material measured by the calibration observations.

3.4

Advantage of Suzaku in the GC survey

The main advantage of the XIS in Suzaku is the several factors smaller and more stable
NXB than the ACIS in Chandra (Garmire et al. 2003) or the EPIC in XMM-Newton (Turner
et al. 2001). Combined with its large effective area and good energy resolution, the XIS has
a unprecedented capability to detect diffuse and faint X-ray emissions. Therefore, Suzaku is
the most suitable to explore a relics of past GC activity which is concealed by the GCXE or
GRXE.

–41–

CHAPTER 3. INSTRUMENTS

Figure 3.17. Spectra of the non X-ray background derived from XIS0 (the thick line) and XIS1(the thin line).

–42–Chapter 4
X-ray Images of the South and Northof the GC
As described in Chapter 2, previous X-ray studies on the GC have predominantly focusedon the Galactic plane region (|b| < 0.◦ 5), and paid less attention to the off-plane region(|b| > 0.◦ 5). In order to open up this frontier, we performed Suzaku observations of the offplane region. Using the obtained data, we constructed X-ray images covering a few hundreds
parsecs of the GC region (Section 4.1). Since some regions have not been observed by Suzaku
yet, we also analyzed archival survey data taken by ASCA for a supplement (Section 4.2).

Contents
4.1

4.2

Suzaku X-ray Images of the Galactic Center . . . . . . . . . . . .

44

4.1.1

Observations and data reduction . . . . . . . . . . . . . . . . . . .

44

4.1.2

Construction of X-ray images . . . . . . . . . . . . . . . . . . . . .

44

ASCA X-ray Image of the Galactic Center

43. . . . . . . . . . . .

45

CHAPTER 4. X-RAY IMAGES OF THE SOUTH AND NORTH OF THE GC

4.1

Suzaku X-ray Images of the Galactic Center

4.1.1

Observations and data reduction

We observed the GC region with Suzaku from 2005 to 2012. The XIS was operated in
the normal clocking mode without either the window nor burst options. The SCI technic
was applied for observations taken after 2006 October.
We used the software package HEASoft version 6.12 for data reduction. After reprocessing
the data with the calibration database released in 2011 October 10, we removed events taken
during the South Atlantic Anomaly passages, night-Earth elevation angles below 5◦ , and dayEarth elevation angles below 20◦ . The entire region of XIS2 and the Segment A of XIS0
have not been functional because of anomalies in 2006 November and 2009 June, respectively.
These disabled regions were eliminated in the following analysis.

4.1.2

Construction of X-ray images

We made mosaic images of the GC region using the XIS data. A raw-data image inthe selected energy band was derived from each observational data. The NXB images werecompiled by xisnxbgen. In order to correct difference of a exposure time and the vignetting
effect, we constructed a exposure map of each observation using the ray-trace simulator,
xissim. We then reprojected the assembled raw images, NXB images, and exposure mapsof all the observations on a reference frame centered at (l, b) = (0.◦ 0, −1.◦ 0), individually.
After subtracting the tiled NXB image from the tiled raw image, we divided it by the tiled
exposure map. The data of all CCDs were co-added to increase photon statistics. Image
pixels were binned with 12 × 12 and smoothed with a Gaussian kernel of σ = 4800 .
Figure 4.1 shows the obtained images in the 0.5–1.0 keV (soft), the 1.0–3.0 keV (middle),
and the 5.0–8.0 keV (hard) band. Sgr A* is located at the center of the images, (l, b)= (−0.◦ 056, −0.◦ 046). In the Galactic plane of the middle and hard bands, we see brightemissions consisting of the GCDX, SNR candidates, and point sources artificially extending
due to the PSF of the XRT. Contrastingly in the soft band, on-plane region is dim becauseof strong interstellar absorption toward the GC direction (NH ∼ 6 × 1022 cm−2 ).
We then focused on the off-plane region in the soft and middle band images. We foundlarge extended emission to the ∼1.◦ 5 south of Sgr A*, which is not seen in the hard band. Inorder to confirm this excess quantitatively, we made projections of the middle band image
along l = 0.◦ 0 (source) and along l = −1.◦ 5 (reference). The obtained profiles are shown in
Figure 4.2. The surface brightness of the source region is about 5 times larger than thatof the reference region at b ∼ −1.5, whereas consistent with the reference region beyondb = −2.◦ 0. Therefore, this excess is not apparent due to low extinction in the off-plane–44–

4.2. ASCA X-RAY IMAGE OF THE GALACTIC CENTER

2.0(a) Soft (0.5-1.0 keV) band(b) Middle (1.5-3.0 keV) band(c) Hard (5.0-8.0 keV) band

SLX 1737-282

Galactic latitude (degree)

1.0

GC North
Sgr A*

0.0

Sgr A*

G359.1-0.05

1E 1743.1-2843

Sgr A*

1E 1740.7-2942

G359.1-0.5
Galactic center diffuse X-ray−1.0

GC SouthA1742-294

G359.0-0.9

−2.0
−3.0
−4.0
1.0

0.0

359.0

1.0

0.0

359.0

1.0

0.0

359.0

Galactic longitude (degree)

Figure 4.1. Suzaku X-ray images of the GC region in the (a) soft (0.5–1.0 keV), (b) middle (1.0–
3.0 keV), and (c) hard (5.0–8.0 keV) bands. The projections shown in Figure 4.2 were extracted
from the the green solid (source) and dashed (reference) rectangles. Some cataloged sources arelabeled in white. The emissions analyzed in the following chapters are labeled in magenta.

region. It is identical with the plume like excess in the ROSAT all sky survey reported by
Wang (2002). Hereafter, we refer to this emission as “GC South”. Detailed spectral analysis
of GC South is performed in Chapter 5.
We also found a strong extended emission near (l, b) = (0.◦ 3, 0.◦ 5) in the middle band.
It seems to be a counterpart to GC South, but limited Suzaku FOVs make it difficult to
identify the overall morphology of the emission. This excess is not seen in either the soft nor
hard band.

4.2

ASCA X-ray Image of the Galactic Center

The 5◦ × 5◦ field in the GC region was surveyed by the Gas Imaging Spectrometers (GIS)
onboard ASCA (Tanaka et al. 1994) as reported in Sakano et al. (2002). Details of theobservations are summarized in that paper. In order to clarify the overall morphology of theoff-plane emissions seen in the Suzaku images, we reanalyzed the GIS image composed by–45–

Surface brightness (arbitrary unit)

CHAPTER 4. X-RAY IMAGES OF THE SOUTH AND NORTH OF THE GC

10

-4

10

-5

Sgr A

projection along l = 0°.0along l = −1°.5
GC South

-4

-3.5

-3

-2.5
-2
-1.5
-1
Galactic latitude (degree)

-0.5

0

0.5

Figure 4.2. Projections of the middle band image extracted from the source (red) and reference(black) regions. The source and background region are along l = 0.◦ 0 along l = −1.◦ 5, respectively(see Figure 4.1).

Sakano et al. (2002).
Figure 4.3(a) shows the 1.5–3.0 keV band image taken by the GIS/ASCA. Althoughphoton statistics and spatial resolution are poorer than the Suzaku data, overall features are
consistent with that in the Suzaku middle band image. We discovered not only GC Southbut also diffuse emission extending toward the north of Sgr A*, which is partially seen in the
Suzaku image. These emissions show a S-shaped morphology and seemed to connect near
Sgr A*.
In order to quantitatively evaluate this northern excess, we derived projections of theimage along l = 0.◦ 25 (source) and l = −0.◦ 75 (reference) as shown in Figure 4.3(a). The
obtained profiles are shown in Figure 4.3(b). We excluded the range of b = 1.◦ 2–1.◦ 3 in thesource projection, because there is a bright point source, SLX 1737−282. We found that
surface brightness of source region at b ∼ 0.◦ 5 is larger than that of the reference region by afactor of 2. Therefore, this excess is also not apparent due to low extinction. Hereafter, we
refer to this excess as “GC North”. Detailed spectrum analysis of GC North is described in
Chapter 6.

–46–

4.2. ASCA X-RAY IMAGE OF THE GALACTIC CENTER(a) ASCA 1.5-3.0 keV band
2.0
SLX 1737-282

1.5

10

Surface brightness (arbitrary unit)

Galactic latitude (degree)

1.0
GC North
0.5
0.0

Sgr A*

-0.5
-1.0
-1.5

0.5

0.0

359.5

projection along l = 0.25
along l = −0.75
SLX 1737-282

GC South
1.0(b) Spatial profile

GC North

10

-2.0

-4

-5

0.5

1

1.5

Galactic latitude (degree)

2

359.0

Galactic longitude (degree)

Figure 4.3. (a) ASCA X-ray image of the GC region in the 1.5–3.0 keV band. Sgr A* is located at
the center of the image and indicated by the cross point. The dotted S-shaped curve is a guide-eyeof GC North and GC South. The solid and dotted green rectangles shows the projection extracted
regions. (b) Spatial profile of the image. The red data is along l = 0.◦ 25 (the solid rectangle in (a)),
which includes GC North, and the black data is along l = −0.◦ 75 (the dotted rectangle in (a)). The
bright point source, SLX 1737−282, is excluded from the red data.

–47–

CHAPTER 4. X-RAY IMAGES OF THE SOUTH AND NORTH OF THE GC–48–

Chapter 5
Analysis on GC SouthIn the previous chapter, we found the large diffuse emission extending toward the south of
Sgr A*. The nature and origin of this emission have not been uncertain. In this chapter, we
present the spectral analysis using Suzaku data, which provide us sufficient photon statistics
to address this question.

Contents
5.1

Analysis Method . . . . . . . . . . . . . . . . . . . . . . . . . . . .

50

5.2

Background Estimation . . . . . . . . . . . . . . . . . . . . . . . .

50

5.3

Spectrum Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . .

56

5.3.1

Brightest region spectrum. . . . . . . . . . . . . . . . . . . . . .

56

5.3.2

Spatially resolved spectra . . . . . . . . . . . . . . . . . . . . . . .

59

Physical Properties of GC South . . . . . . . . . . . . . . . . . .

64

5.4

49

CHAPTER 5. ANALYSIS ON GC SOUTH

5.1

Analysis Method

The close-up view of the GC South area are shown in Figure 5.1 to declare the analysis
region. Spectrum analysis in this section was performed by three steps. We first estimatedthe Galactic X-ray background (BG) using the reference regions (white squares in Figure5.1). We then analyzed the brightest region (green rectangle in Figure 5.1(b)) to reduce
contamination of the BG and to derive precise physical properties of GC South. Finally, we
divided the broad area into ten annuli as shown in Figure 5.1(a), and investigated spatial
variations of plasma parameters based on the analysis of the brightest region. Logs of theobservations used in this chapter are summarized in Table 5.1.
We used the SPEX software version 2.04.00 (Kaastra et al. 1996) in our spectral analysis.
Redistribution matrix files (RMFs) and ancillary response files (ARFs) were generated byxisrmfgen and xissimarfgen, respectively. Spectral data of XIS0 and XIS3 (FIs) wereco-added since they have nearly the same responses. The energy band affected by thecontamination (<0.7 keV) and that around the neutral Si K-shell edge (1.7–1.8 keV) wereignored because of calibration uncertainties.

5.2

Background Estimation

X-ray sources in the GC region are affected by the strong Galactic BG, whose X-ray flux
depends on position with various components. Thus, careful evaluation of the BG emission
is essential. Uchiyama et al. (2011, 2013) extensively studied this emission with Suzaku, anddetermined the spatial distribution along the Galactic latitude and longitude. On the basisof the spatial distribution, the authors divided the BG emission into two components: the
GCXE and the GRXE (see Section 2.2.2). Uchiyama et al. (2013) reported that the GCXE
is elliptical with a size of ∼ 0.◦ 6 × 0.◦ 3 while the GRXE is at least one order of magnitudelarger scale. Therefore, GC South is located in the region where contributions from the
GRXE are dominant.
We extracted X-ray spectra from the four reference regions shown in Figure 5.1 (Ref 1–4).
These spectra exhibit many emission lines of various elements that are typical of the GRXE.
We therefore constructed a BG model spectrum based on the GRXE model reported by
Uchiyama et al. (2013). The model includes a high temperature plasma component (HP),a low temperature plasma component (LP), and a foreground emission component (FE), allof which are modeled by collisional ionization equilibrium (CIE) plasmas. Although the FE
consists of 0.59 keV and 0.09 keV plasmas in the study of Uchiyama et al. (2013), we applieda single ∼0.59 keV plasma because the contribution of the 0.09 keV plasma is negligible in
the energy range used in our analysis (0.7–10 keV). Free parameters are the temperatures–50–

5.2. BACKGROUND ESTIMATION(a) Soft (0.5-1.0 keV) band
Sgr A*−0.5

S-Src 8

−1.5

−2.5

S-Src 9

0.5

Ref 1

S-Src 1
S-Src 2
S-Src 3S-Src 4
S-Src 5
S-Src 6
S-Src 7

−1.0

Ref 2
Ref 3
Ref 4

−0.5

Ref 1
Ref 2

−1.0
brightest region

−1.5

−2.5

359.5 359.0 358.5

Ref 3

42’

Ref 4
16’

−2.0

S-Src 10

0.0

Sgr A*

0.0
Galactic latitude (degree)

Galactic latitude (degree)

0.0

−2.0

(b) Middle (1.5-3.0 keV) band

0.5

0.0

359.5 359.0 358.5

Galactic longitude (degree)

Galactic longitude (degree)(c) Hard (5.0-8.0 keV) band
Sgr *A

Galactic latitude (degree)

0.0
−0.5

Ref 1
Ref 2

−1.0

Ref 3

−1.5

Ref 4

−2.0
−2.5

0.5

0.0 359.5 359.0 358.5
Galactic longitude (degree)

Figure 5.1. X-ray mosaic images of the close-up view of the GC South region. (a) Soft band
image. The spectrum extracted regions for the spatially resolved analysis are shown in the green
annuli (see Section 5.3.2). (b) Middle band image. The green square shows the brightest part of
GC South, where we extracted a spectrum in the analysis of Section 5.3.1. (c) Hard band image. In
all the images, the cyan cross indicates the position of Sgr A*, and the white squares are reference
regions used to estimate the BG spectrum (see Section 5.2).

–51–

CHAPTER 5. ANALYSIS ON GC SOUTH
Table 5.1. Observation logs of the GC South region.
Sequence #α2000.0

Aim Pointδ2000.0lObs. Date(yyyy-mm-dd)b

Exposure(ks)

Obs. Name

23
21
20
21
118
59
100
47
55
51
103
54
51
52
52
52
53

GC14
GC15
GC16
GC17
GALACTIC BULGE
GC(0,-1.6)
L,B=0.107,-1.425
GALACTIC BULGE1GALACTIC BULGE2GALACTIC BULGE3RADIO ARC SOUTH
RADIO ARC S E
RADIO ARC S W
GC SOUTH 1
GC SOUTH 2
GC SOUTH 3
GC SOUTH 4

53
51
41
51

HESSJ1745-303B
HESSJ1745-303C
GALACTIC BULGE7GALACTIC BULGE8

Source regions
502002010
502003010
502004010
502005010
502059010
503081010
504050010
504088010
504089010
505078010
507016010
507017010
507018010
507028010
507029010
507030010
507031010

267.◦ 16266.◦ 96267.◦ 48267.◦ 29268.◦ 38268.◦ 05267.◦ 86267.◦ 22267.◦ 56267.◦ 95266.◦ 98267.◦ 19266.◦ 77268.◦ 10268.◦ 37268.◦ 18268.◦ 41−29.◦ 14−29.◦ 42−29.◦ 31−29.◦ 60−29.◦ 96−29.◦ 76−29.◦ 58−29.◦ 37−29.◦ 60−29.◦ 80−29.◦ 07−28.◦ 77−29.◦ 37−29.◦ 48
−29.◦ 71−29.◦ 20−29.◦ 450.◦ 17359.◦ 840.◦ 17359.◦ 830.◦ 050.◦ 03359.◦ 210.◦ 00359.◦ 95359.◦ 95
0.◦ 150.◦ 50359.◦ 800.◦ 300.◦ 220.◦ 570.◦ 46504002010
504003010
504090010
504091010

266◦ .06−30◦ .55

358.◦ 46266◦ .39−30◦ .62358.◦ 56266◦ .68
267◦ .09−30◦ .83−31◦ .05358.◦ 51358.◦ 50−0.◦ 67−0.◦ 67−1.◦ 00−1.◦ 00−2.◦ 00−1.◦ 65−1.◦ 95−0.◦ 83−1.◦ 21−1.◦ 60−0.◦ 50−0.◦ 50−0.◦ 50−1.◦ 55−1.◦ 87−1.◦ 47−1.◦ 77

2007-10-09
2007-10-10
2007-10-10
2007-10-11
2007-09-29
2009-03-09
2010-03-06
2009-10-14
2009-10-09
2011-03-04
2012-09-16
2012-09-29
2012-09-30
2012-09-20
2012-09-22
2012-10-05
2012-10-06

Reference regions−0.◦ 59−0.◦ 87−1.◦ 19−1.◦ 60

2010-02-27
2010-02-25
2009-10-13
2009-09-14(kTHP , kTLP , kTFE ) and the emission measures (EMHP , EMLP , EMFE ) of the HP, LP, andFE. The metal abundances relative to solar values (Anders & Grevesse 1989) were common
between the HP and LP. Those of Mg, Si, S, Ar, Ca, and Fe were free parameters while the Ne
and Ni abundances were linked to Mg and Fe, respectively. The other elements were fixed tothe solar values. In the FE, the metal abundances of Ne and Mg were free parameters while
those of the other elemetns were linked and allowed to vary. In order to reproduce the Fe I
Kα emission in the spectra, we also added a Gaussian component with a fixed line centroidof 6.4 keV. The continuum emission associated with the Fe I Kα emission was included in
the HP. In addition, we need to include the cosmic X-ray background (CXB) even though
its contribution is small. The CXB was represented by a single power-law function with aphoton index and surface brightness (2–10 keV) of 1.41 and 6.38 × 10−8 erg cm−2 s−1 sr−1 ,respectively (Kushino et al. 2002). The explicit form of our BG model is therefore
AGC × (HP + LP + Gaussian) + AFE × FE + (AGC )2 × CXB,

(5.1)

where AGC (column density of NH GC ) and AFE (column density of NH FE ) represent the interstellar absorptions toward the GC and for the FE component, respectively. The absorption–52–

5.2. BACKGROUND ESTIMATIONCounts s−1 keV−1

Ref 2

10−2

5
0
−5

χ

χ

Counts s−1 keV−1

Ref 1
10−1

1

2
Energy (keV)

5

10

10−1

10−2

5
0
−5

Counts s−1 keV−1

10−2

5
0
−5

2
Energy (keV)

5

10

1

2
Energy (keV)

5

10

Ref 4

10−1

χ

χ

Counts s−1 keV−1

Ref 3

1

1

2
Energy (keV)

5

10

10−1

10−2

5
0
−5

Figure 5.2. Spectra and the best-fit models of reference regions. Only the co-added FI spectra andmodels are displayed for visibility, although the FI and BI spectra were simultaneously fitted. The
red, orange, purple, blue, and gray lines represent the HP, LP, Gaussian, FE, and CXB, respectively(see text).

for the CXB is twice that of NH GC , because the CXB is of extragalactic origin behind the
GC region.
We simultaneously fitted the reference-region spectra with this model. Since the fluxin the term (HP + LP + Gaussian) depends on the position as shown in Uchiyama et al.
(2013), we multiplied by a constant factor as
AGC × (HP + LP + Gaussian) × f actor,(5.2)

where f actor for Ref 1 was fixed to 1, and those for the other regions were free parameters.
The NH GC values were also independently allowed to vary while the other parameters werecommon among the regions. As a result of our fitting, the spectra were well reconstructedwith χ2 /d.o.f = 1415/1107. The best-fit models and parameters are shown in Figure 5.2 andTable 5.2, respectively. We note that the HP, LP and FE parameters are slightly different
from those of Uchiyama et al. (2013). This is due to the difference of the models andselected regions; the GRXE spectrum of Uchiyama et al. (2013) is a composite from theon-plane regions of |l| > 2.◦ 0 and |b| < 0.◦ 5, while our BG regions are far from the plane of–53–Hydrogen column density (1022 cm-2)

CHAPTER 5. ANALYSIS ON GC SOUTH
5
infrared extinction(l = −1°.5)

2

X-ray absorption(reference region)
infrared extinction( l = 0°)

1
X-ray absorption(brightest region)

0.5
−1.6

−1.4
−1.2
−1.0
−0.8
Galactic latitude (degree)−0.6

Figure 5.3. Profile of the absorption column density. The blue circles indicate NH GC obtained
from the best-fit values of the reference regions. The blue broken line is the exponential-folding
model of NH GC (see text). The blue and red shaded areas show the near-infrared extinctions (AKs )including the errors along l = −1◦ .5 and l = 0◦ , respectively (Gonzalez et al. 2012). The nearinfrared extinctions were converted to NH with a scaling factor of NH /AKs = 1.8×1022 cm−2 mag−1 .
The red circles indicate the fitting results of the spatially resolved source spectra.

l = −1.◦ 4 and −1.◦ 6 < b < −0.◦ 6. One might claim that the abundances of some elements
are significantly smaller than the solar values. We approximated the BG spectra using
simplified plasma models based on the GRXE model near the GC (Uchiyama et al. 2013).
In the Galactic bulge regions (our observation fields) however, the possible contribution ofunresolved faint point sources (e.g. Revnivtsev et al. 2009), CXB fluctuation (Kushino et al.
2002), or unknown non thermal emissions may not be ignored. These sources increase thecontinuum flux in the BG spectrum, and hence reduce the apparent abundances. In orderto verify this possibility, we fitted the spectra by fixing the abundances of all the elementsin the relevant plasma models (LP, HP, and FE) to the solar values, and adding two powerlaw components; one is convolved with AGC , and the other is convolved with AFE . Then
we obtained a reasonable fit (χ2 /d.o.f = 1654/1116) with no significant change for all theparameters in Table 5.2.

The best-fit f actor values decrease with increasing the Galactic latitude (b) with an
exponential-folding scale of 3.◦ 3 ± 0.◦ 8, which is roughly consistent with those of Uchiyama–54–

5.2. BACKGROUND ESTIMATION
Table 5.2. Fitting results of the spectra in the reference regions.
Parameter (Unit)

Ref 1

Ref 2

Ref 3

Ref 4

— HP + LP + Gaussian —NH GC (1022 cm−2 )f actorkTHP (keV)
kTLP (keV)∗
EMHP(1058 cm−3 )∗EMLP(1058 cm−3 )Mg = Ne (solar)
Si (solar)
S (solar)
Ar (solar)
Ca (solar)
Fe = Ni (solar)
Other Elements (solar)
Fe I Kα flux (10−6 ph cm−2 s−1 )3.40 ± 0.11
1 (fixed)

2.02 ± 0.09 1.37 ± 0.08
0.85 ± 0.03 0.71 ± 0.02
7.6+0.4−0.3
0.69 ± 0.03
0.38 ± 0.01
0.45+0.09−0.08
0.43+0.14−0.12
0.94+0.17−0.14
1.6 ± 0.3
3.5 ± 0.8
2.8 ± 1.0
0.53 ± 0.04
1 (fixed)
6.4 ± 1.2

1.33 ± 0.08
0.77 ± 0.02

— FE —NH FE (1022 cm−2 )kTFE (keV)∗(1058 cm−3 )EMFE
Ne (solar)
Mg (solar)
Other Elements (solar)

0.42 ± 0.01
0.33 ± 0.04
0.86+0.66−0.33
0.06 ± 0.02
0.19+0.05−0.04
0.022+0.014−0.009Notes:R
Emission measure assuming the distance of 8 kpc. EM = ne np dV , where ne , np , and V , are theelectron density, the proton density, and the emitting volume, respectively.

∗

et al. (2013). The NH GC values also decrease with b, as phenomenologically given by


|b|
22
NH GC = 6.1 ± 0.7 × 10 exp − ◦cm−2 .
(5.3)
0. 87 ± 0.◦ 10
This model and the best-fit NH GC are plotted in Figure 5.3 by the blue broken line andcircles, respectively. We compared these data with the near-infrared extinction (AKs ) of
the stars located in the GC (Gonzalez et al. 2012). The mean AKs and best-fit NH GC in
the region of −1.◦ .6 < b < −1.◦ .2 are 0.76 mag and 1.35 × 1022 cm−2 , respectively. Thus, aconversion factor NH /AKs was estimated to be 1.8 × 1022 cm−2 mag−1 , which is within therange reported in the literature (e.g., van den Berg et al. 2009; Güver & Özel 2009). With
this conversion factor, the AKs profile is superimposed on Figure 5.3 as the blue shaded area.
We confirmed that the NH GC model well follows the over-all trend of the AKs profile.
–55–

CHAPTER 5. ANALYSIS ON GC SOUTH
Figure 5.3 also shows AKs in the source region (l = 0.◦ 0) with the red shaded area (Gonzalez et al. 2012). We found that the AKs in the source region is smaller than that in thereference region by a factor of ∼0.55 in the −1.◦ 6 < b < −1.◦ 2 range. Assuming the sameb-dependence as in the reference regions, we estimated NH GC in the source region by normalizing Equation (5.3) by a factor of 0.55. This estimated model is shown by the broken
red line in Figure 5.3.

5.3

Spectrum Analysis

5.3.1

Brightest region spectrum

We first investigated a spectrum extracted from the brightest region (420 ×160 rectangularregion) in order to increase the signal-to-background ratio and to reveal the precise physicalproperties. Figure 5.4(a) shows the extracted spectrum. The hard X-ray band above 5 keV
is dominated by the BG emission as shown in Figure 5.4(c). In order to subtract these BG
components, we applied the BG model with the best-fit values obtained from the reference
regions (Section 5.2). Only the surface brightness and NH GC for the (HP + LP + Gaussian)component should be different from those in the reference regions. Therefore, f actor wastreated as a free parameter, and the NH GC value was fixed to 0.69 × 1022 cm−2 , which is
the mean value of the NH GC profile in the source regions (the broken red line in Figure 5.3).
The parameters for the FE were the same as those in the reference regions.
The BG model with these parameters well reproduced the 5–10 keV band spectrum, but
it left a large data excess in the 0.7–5.0 keV band. This excess shows emission lines fromhighly ionized atoms, and hence is of thermal plasma origin. We then attempted a fittingwith a model of combining CIE plasma and the BG model. The free parameters in the CIE
plasma were NH , kT , and the abundances of the major elements (Ne, S, Si, Ar, and Fe). The
abundances of O, Ca, and Ni were linked to those of Ne, Ar, and Fe, respectively. Those of
the other elements were fixed to the solar values.
Spectral fitting with the composite model gave a best-fit CIE temperature of kT =0.62 keV, but failed to reconstruct the data with χ2 /d.o.f = 2617/352 (Figure 5.4(b)). The
most prominent residuals in Figure 5.4(b) are the line and bump at 0.8 keV and 1.2–1.3 keV,
respectively. These are known to originate from incomplete atomic data for the Fe L-shellcomplex in the current plasma model. Gu et al. (2007) claimed that the line intensity ratioof 3s→2p (∼0.7 keV) over 3d→2p (∼0.8 keV) for Fe XVIII has an uncertainty. Brickhouse
et al. (2000) and Audard et al. (2001) reported that the SPEX plasma code underestimates
the flux around 9.6–10.6 Å (1.17–1.29 keV), since some Fe L-shell transitions (n = 6, 7, 8 → 2
for Fe XVII, n = 6, 7 → 2 for Fe XVIII, and n = 6 → 2 for Fe XIX) are missing. We thus–56–

5.3. SPECTRUM ANALYSIS(a)

Recombining plasma model
Missing Fe L-shell complex
Background(b)χ2/d.o.f = 2617/352

Counts s−1 keV−1

100

10−1

10−2

χ

5
0
−5

χ

5

Fe-L

Fe-L
χ2/d.o.f = 1378/350

0
−5
5

χ

(c)

Si XIV Si-RRC

(d)

S-RRC
χ2/d.o.f = 599/348

0
−5
1

2
Energy (keV)

5

10

Figure 5.4. (a) Spectrum extracted from the entire source region. Only the co-added FI spectra
and the models are displayed for visibility, although the FI and BI spectra were simultaneously
fitted. The red, purple, and gray solid lines indicate the best-fit RP model, the missing Fe L-shelllines, and the background model, respectively. The χ2 values fitted with each model are also shown:(b) CIE plasma model, (c) CIE plasma and additional Fe-L shell transients model, and (d) RP and
additional Fe-L shell transients model (best-fit).

phenomenologically added Gaussians at 0.8 keV and 1.2 keV. The FWHM of the 1.2 keVline was fixed to 120 eV to approximate multiple lines. Then, the fit was improved with
χ2 /d.o.f = 1378/350, but large residuals still remained (Figure 5.4(c)). We conclude that
the CIE model for GC South plasma is invalid.
In Figure 5.4(c), the data excess at 2.0 keV corresponds to Si XIV Kα. This indicates that
the CIE model (kT = 0.62 keV) underestimates the ionization state of Si. In addition, we
can clearly observe hump-like features around 2.5–3.0 keV and 3.2–3.8 keV. These features
are most likely due to the radiative recombination continua (RRCs) of Si and S, which areobserved in some Galactic supernova remnants (SNRs; see Sawada & Koyama 2012; Uchida
et al. 2012 and references therein). Although the residual of Si XIV Kα can be solved by amulti-temperature CIE model, the RRCs cannot be reproduced.
–57–

CHAPTER 5. ANALYSIS ON GC SOUTH
Table 5.3.
Fitting results of thebrightest-region spectrum with the NEIJ
model.
Parameter

Unit

Value

NH
kTe
kTinit
nt
EM ∗
Ne =O
Mg
Si
S
Ar = CaFe = Ni(1022 cm−2 )
(keV)
(keV)
(1011 s cm−3 )(1058 cm−3 )(solar)(solar)(solar)(solar)(solar)(solar)0.70 ± 0.03
0.46 ± 0.02
1.63+0.27−0.18
5.3 ± 0.5
9.5 ± 1.0
0.39+0.05−0.04
0.81 ± 0.07
0.71 ± 0.05
0.79+0.09−0.08
1.9+0.6−0.5
0.10 ± 0.02

Notes:* Emission measure assuming the distance of 8 kpc.

The RRC residuals strongly indicate that the plasma is in the recombining phase. We thus
applied a non-equilibrium ionization model (NEIJ in SPEX), that successfully reproduces the
RRC features observed in the SNRs W44 and G346.6−0.2 (Uchida et al. 2012; Yamauchi
et al. 2013). This model traces the evolution of the ionization state when the initial ionization
temperature (kTinit ) and the electron temperature (kTe , assumed to be constant) are given;
the following evolution of the ionization state is traced by the relaxation timescale of ne t,
where ne and t are the number density of electrons and the elapsed time, respectively. If thebest-fit kTinit is ∼0, then this model describes an ionizing plasma (NEI in XSPEC), while if
kTinit > kTe , it describes a recombining plasma (RP).

The free parameters in the NEIJ model were kTinit , kTe , ne t, and the abundances ofthe major elements (the same as in the CIE fitting). The NEIJ model gave a good fit
with χ2 /d.o.f = 599/348. The best-fit kTinit , kTe , and ne t are 1.6 keV, 0.46 keV, and5.3 × 1011 s cm−3 , respectively (see Table 5.3). This best-fit model well reproduced the
RRC features around 2.5–3.0 keV and 3.2–3.8 keV, as shown in Figure 5.4(d).

For comparison, we list the average charge state of each element in the best-fit NEIJ
model and that of the CIE at the same electron temperature (Table 5.4). We confirmed that
GC South is really an RP.
–58–

5.3. SPECTRUM ANALYSIS
Table 5.4. Average charge state of each element.
model

Ne

Mg

Si

S

Ar

Fe

best-fit NEIJ
CIE (kT = 0.46 keV)

9.56
9.30

11.1
10.3

12.7
12.0

14.2
13.9

15.7
15.6

17.0
16.4

Notes: Average charge state of 0 means a neutral atom.

5.3.2

Spatially resolved spectra

We found that the recombining process dominates in the brightest part of the plasma. In
order to study the spatial variation of the plasma parameters, we then divided the entire GC
South field into ten annual regions as shown in Figure 5.1(a), and extracted their spectra(S-Src 1 to S-Src 10). Bright point sources were excluded to avoid the contamination of them.
The factor in the BG model was left free as done in the fitting of the brightest part. The
absorption column density (NH GC ) of each region was fixed to the value estimated from AKs ,which is converted by the factor of 1.8 × 1022 cm−2 mag−1 . We fixed kTinit in the NEIJ (RP)model to 1.6 keV, while allowed to vary kTe and ne t to study their spatial variation. The
elemental abundances were fixed to the best-fit values of the brightest region (Table 5.3).
The Fe L-shell lines at ∼0.8 keV and ∼1.2 keV were not added at first. We individually
fitted each spectrum with this model (hereafter the tentative model).
Figure 5.5 shows the fitting results of the S-Src 1–10 spectra with the tentative model.
In the S-Src 1-3 spectra, we found large residuals especially in the 0.5-1.0 keV band. One
possible cause is the statistical error (∼15%) of AKs , which is used to estimate NH GC . We
attempted to treat NH GC as a free parameter in this range but the fits are not improved.
Another cause is the spatial fluctuation of the FE flux. Ryu (2013) investigated GCXE
spectra in various on-plain regions (Sgr B2, Sgr C, Sgr D, and Sgr E), and reported that the
FE fluxes vary within a factor of 2.4. Moreover, we found local excess of Hα emission in thisregion, which mainly traces the foreground warm matter, from the archival data of SHASSA(Gaustad et al. 2001). Therefore, we treated the normalization (emission measure) of the
FE component as a free parameter in the fitting of the S-Src 1–3 spectra. Then, the FE
fluxes were renormalized by a factor of 2.1–2.6, and spectra in the 0.5–1.0 keV band wererepresented. These factors are consistent with that reported by Ryu (2013).
In the S-Src 4-10 spectra, we also found Gaussian-like residuals at ∼0.8 keV and ∼1.2 keV.
These are most likely due to the missing Fe L-shell lines as shown in the previous section.
Thus, we also added the two Gaussians to the model in the same manner as the fitting of
the brightest part (hereafter the modified model).
The modified model fairly improved the fits as shown in Figure 5.6. The obtained parameters and reduced χ statistics are summarized in Table 5.5. The surface brightness in–59–

CHAPTER 5. ANALYSIS ON GC SOUTH

0.1

Counts sï1 keVï1

ï5

0

r

5

0.01

0.1
0.01
0
ï5

r

5

Counts sï1 keVï11SïSrc21SïSrc1

1

2
Energy (keV)

5

10

10

1

2
Energy (keV)

5

10

1

2
Energy (keV)

5

10

1
0.1

Counts sï1 keVï1

5
0

r

ï5
2
Energy (keV)

5

10

0.1
0.01
5
0
ï5

ï5

0

r

5

0.01

0.1

Counts sï1 keVï11SïSrc61

SïSrc5Counts sï1 keVï1

5

0.01

0.1

Counts sï1 keVï1

0.01
5
0

r

ï5

1

r

2
Energy (keV)SïSrc41SïSrc3

1

1

2
Energy (keV)

5

10

Figure 5.5. Spectra of the S-Src 1–10 regions fitted with the tentative model (see text). Only the
co-added FI spectra and models are displayed for visibility, although the FI and BI spectra were
simultaneously fitted.

–60–

5.3. SPECTRUM ANALYSIS

0.1

Counts sï1 keVï1

ï5

0

r

5

0.01

0.1
0.01
0
ï5

r

5

Counts sï1 keVï11SïSrc81SïSrc7

1

2
Energy (keV)

5

10

5

10

1

2
Energy (keV)

5

10

1
0.1

Counts sï1 keVï1

ï5

0

r

5

0.01

0.1

Counts sï1 keVï1

0.01
5
0
ï5

r

2
Energy (keV)SïSrc101SïSrc91

12
Energy (keV)

5

10

Figure 5.5. Continued.

the 5–8 keV band derived from the best-fit of the source and reference regions agrees withthe spatial profile composed by Uchiyama et al. (2013). This supports the reliability of our
method of the BG estimation and subtraction. In order to investigate the impact of the
FE-flux variation on the fitting results of S-Src 4–10, we tried to increase the FE flux by afactor of 2, and confirmed that the best-fit parameters are consistent with that in Table 5.5.
Therefore systematic errors due to the spatial variation of the FE flux are negligible in the
S-Src 4–10 regions.
We plotted the best-fit NH and estimated NH GC on Figure 5.7 as a function of the Galactic
latitude, and found that they are reasonably consistent. Figure 5.8(a)–(c) shows spatial
distributions of kT , ne t, and the surface brightness of the plasma. All these parameters
decrease as higher the Galactic latitude. We found that plasmas in S-Src 1–3 are in a CIE
state whereas those in S-Src 4–10 are RPs. To evaluate a degree of over-ionization, we
calculated the difference of the average ionization state of Si between the best-fit model andthe CIE model assuming the obtained kTe . If the difference becomes a positive value, theplasma is in an over-ionized state (RP). The result is shown in Figure 5.8(d). The degree ofover-ionization is marginally increase as higher the Galactic latitude.
–61–

CHAPTER 5. ANALYSIS ON GC SOUTH

0.1

Counts sï1 keVï1

ï5

0

r

5

0.01

0.1
0.01
0
ï5

r

5

Counts sï1 keVï11SïSrc21SïSrc1

1

2
Energy (keV)

5

10

10

1

2
Energy (keV)

5

10

1

2
Energy (keV)

5

10

1
0.1

Counts sï1 keVï1

5
0

r

ï5
2
Energy (keV)

5

10

0.1
0.01
5
0
ï5

ï5

0

r

5

0.01

0.1

Counts sï1 keVï11SïSrc61

SïSrc5Counts sï1 keVï1

5

0.01

0.1

Counts sï1 keVï1

0.01
5
0

r

ï5

1

r

2
Energy (keV)SïSrc41SïSrc3

1

1

2
Energy (keV)

5

10

Figure 5.6. Spectra of the S-Src 1–10 regions fitted with the modified model (see text). Only the
co-added FI spectra and models are displayed for visibility, although the FI and BI spectra were
simultaneously fitted.

–62–

5.3. SPECTRUM ANALYSIS

0.1

Counts sï1 keVï1

ï5

0

r

5

0.01

0.1
0.01
0
ï5

r

5

Counts sï1 keVï11SïSrc81SïSrc7

1

2
Energy (keV)

5

10

5

10

1

2
Energy (keV)

5

10

1
0.1

Counts sï1 keVï1

ï5

0

r

5

0.01

0.1

Counts sï1 keVï1

0.01
5
0
ï5

r

2
Energy (keV)SïSrc101SïSrc91

12
Energy (keV)

5

10

Figure 5.6. Continued.
Table 5.5. Fitting results of the spatially resolved spectra with the NEIJ model.
Region∗NH
GC(1022 cm−2 )NH(1022 cm−2 )

kTe
(keV)

ne t
(1011 s cm−3 )EM †(1058 cm−3 )SB‡FE ratio§

χ2 /dof

S-Src 1
S-Src 2
S-Src 3S-Src 4
S-Src 5
S-Src 6
S-Src 7
S-Src 8
S-Src 9S-Src 10

2.31
1.54
1.25
0.84
0.76
0.65
0.56
0.66
0.46
0.41

3.27+0.12−0.16
2.30+0.14−0.18
1.53+0.17−0.26
0.86+0.03−0.02
0.73 ± 0.02
0.60 ± 0.02
0.48 ± 0.02
0.36+0.04−0.02
0.24 ± 0.02
0.36 ± 0.03

0.63 ± 0.03
0.75 ± 0.04
0.61+0.28−0.12
0.50+0.02−0.03
0.46 ± 0.02
0.48 ± 0.02
0.44 ± 0.02
0.41+0.03−0.06
0.43+0.03−0.02
0.38 ± 0.03

> 21
> 13
> 4.7
6.2+0.4−0.5
5.8 ± 0.3
5.5 ± 0.3
4.6+0.3−0.2
5.0+0.4−0.6
4.9 ± 0.4
5.1±+0.6−0.5

5.3+1.4−0.7
2.3+0.4−0.3
2.8+1.3−1.1
1.9+0.3−0.1
2.6 ± 0.2
4.1 ± 0.3
2.5 ± 0.2
1.1+0.36−0.1
1.2 ± 0.1
0.39 ± 0.05

9.0+2.4−1.2
6.9+1.2−0.9
5.4+2.5−2.1
8.3+1.3−0.4
8.2+0.6−0.6
5.3+0.4−0.4
4.1+0.3−0.3
1.7+0.6−0.1
2.0+0.2−0.2
1.4+0.2−0.2

2.1 ± 0.2
2.6 ± 0.1
2.3 ± 0.4
1 (fix)
1 (fix)
1 (fix)
1 (fix)
1 (fix)
1 (fix)
1 (fix)525/328
399/328
381/328
415/328
408/328495/328
417/328
541/328570/328552/328

Notes:
kTinit and elemental abundances for the NEIJ model were fixed to the values listed in Table 5.3.
* Absorption column density for the BG emissions that were estimated from the near-infrared extinction investigated by
Gonzalez et al. (2012). They were fixed to the listed values.
† Emission measure assuming the distance of 8 kpc.
‡ Absorption-corrected surface brightness in the 0.7–5.0 keV band with a unit of 10−14 erg s−1 cm−2 arcmin−2§ Ratio of the FE flux between in the best-fit model and in the original BG model shown in Table 5.2.

–63–

Hydrogen Column Density (1022 cm−2)

CHAPTER 5. ANALYSIS ON GC SOUTH

1

10

-1

-2.0

-1.8

-1.6

-1.4

-1.2

-1.0

-0.8

-0.6

Galactic Latitude (degree)

Figure 5.7. Absorption column densities along the Galactic latitude. The black and red data
points show the NH GC estimated from the near-infrared extinction and the best-fit NH values
obtained from the spatially resolved analysis, respectively.

5.4

Physical Properties of GC SouthAs shown in Figure 5.7, the absorption column densities for GC south (red) are consistent
with those toward the GC region (black). Therefore, GC south is most likely to be located in
the GC region (8 kpc distance). Then, the physical size of the plasma is ∼ 200 pc × 70 pc at
this distance. Such a large diffuse plasma have not been discovered in the GC region except
for the GCDX. Since the sound velocity of the plasma with kT = 1.6 keV is 510 km s−1 , thedynamical timescale, which takes for the plasma to travel the 200 pc length, is estimated tobe ∼4 × 105 yr.
Assuming the depths of the emitting regions are 70 pc, we estimated the volume of eachregion (S-Src 1–10) as the apparent size × the depth (see Table 5.6). Then, we calculated
the electron density (ne ), the pressure (p), the mass (M ), and the thermal energy (Eth ). We
also estimated the recombination timescales from the best-fit ne t and the derived ne . These
results are summarized in Table 5.6 and plotted in Figure 5.9. We discovered that ne and p
decreases with as higher the Galactic latitude. The recombination timescale is slightly larger
as further from the GC and ∼3 × 105 yr, which is consistent with the dynamical timescale.
The obtained M and Eth reach ∼ 3000 M and ∼ 8 × 1051 erg, respectively.

–64–5.4. PHYSICAL PROPERTIES OF GC SOUTH

1(b) Electron temperaturekTe (keV)

SB (erg cm−2 s−1 arcmin−2)(a) Surface brightness

10

1
-2.0

-1.8

-1.6

-1.4

-1.2

-1.0

-0.8

-1
10

-0.6

-2.0

-1.8

(c) Relaxation timescale

Diff. average charge statent (1011 s cm−3)

10
8
6
4
2

-2

-1.8

-1.6

-1.4

-1.2

-1

-1.6

-1.4

-1.2

-1.0

-0.8

-0.6

Galactic Latitude (degree)

Galactic Latitude (degree)

-0.8

-0.6

Galactic Latitude (degree)1(d) Difference of average charge state

0.8
0.6
0.4
0.2
0
-0.2

CIE state
-2

-1.8

-1.6

-1.4

-1.2

-1

-0.8

-0.6

Galactic Lattitude (degree)

Figure 5.8. Best-fit parameters along the Galactic latitude. (a) surface brightness, (b) kTe , (c)ne t, and (d) difference of the average charge state form the CIE model.

–65–CHAPTER 5. ANALYSIS ON GC SOUTH

Table 5.6. Physical parameters of each region derived from the spatially resolved spectrum fitting.
RegionS-Src 1
S-Src 2
S-Src 3S-Src 4
S-Src 5
S-Src 6
S-Src 7
S-Src 8
S-Src 9S-Src 10V∗
(1060 cm3 )ne(10 cm−3 )

4.5
2.7
3.6
1.4
1.8
4.5
3.4
3.3
3.0
1.2

+1.4
10.9−0.7+0.8
9.2−0.6
8.8+2.0−1.7+0.9
11.5−0.3
12.0 ± 0.5
9.6 ± 0.4
8.6 ± 0.3
+0.9
5.8−0.2
6.3 ± 0.3
5.6 ± 0.4

−2

−10

(10

P
erg cm−3 )1.6+0.2−0.1
1.6+0.2−0.1
1.3+0.7−0.4
1.4 ± 0.1
1.3 ± 0.1
1.1 ± 0.1
0.91 ± 0.06
0.57+0.10−0.09
0.65+0.05−0.04
0.51 ± 0.05

Notes:
Quoted errors include only statistical errors in the spectrum fitting.
∗Assumed the line of sight length of 70 pc.

–66–

M(10 solar)

Eth(1051 erg)

Recombination time(105 year)

5.7+0.8−0.4
3.0+0.3−0.3
3.7+0.9−0.7
1.9+0.2−0.1
2.5+0.1−0.1
5.0+0.2−0.2
3.4+0.1−0.1
2.2+0.4−0.1
2.2+0.1−0.1
0.8+0.1−0.1

1.7+0.2−0.1
1.0+0.1−0.1
1.1+0.6−0.3
0.46+0.04−0.03
0.55+0.03−0.03
1.1+0.1−0.1
0.71+0.04−0.04
0.43+0.08−0.07
0.46+0.04−0.03 1
0.15+0.02
−0.02

—
—
> 1.6
1.7 ± 0.1
1.5 ± 0.1
1.8 ± 0.1
1.7 ± 0.1
2.7+0.5−0.3
2.5+0.2−0.2
2.9+0.4−0.3

2

5.4. PHYSICAL PROPERTIES OF GC SOUTH(b) Pressure profile

-1

-10

10

Pressure (10-3Density (cm )-3erg cm )(a) Electron density profile

-2.0

-1.8

-1.6
-1.4
-1.2
-1.0
Galactic latitude (degree)

-0.8

-0.6

1

-2.0

-1.8

-1.6
-1.4
-1.2
-1.0
Galactic latitude (degree)

-0.8

-0.6

(c) Timescale profile
6

45Time 10 year

5

3
2
1
-2

-1.8

-1.6
-1.4
-1.2
-1
Galactic latitude (degree)

-0.8

-0.6

Figure 5.9. Spatial profiles of (a) ne , (b) P , and (c) recombination timescale of the plasma.

–67–

CHAPTER 5. ANALYSIS ON GC SOUTH–68–

Chapter 6
Analysis on GC North
In this chapter, we present the analysis results of GC North, a possible counterpart to
GC South.

Contents
6.1

Analysis Method . . . . . . . . . . . . . . . . . . . . . . . . . . . .

70

6.2

Background Estimation . . . . . . . . . . . . . . . . . . . . . . . .

71

6.3

Spectrum Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . .

73

6.3.1

Merged spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . .

73

6.3.2

Spatially resolved spectra . . . . . . . . . . . . . . . . . . . . . . .

73

Physical Properties of GC North . . . . . . . . . . . . . . . . . .

75

6.4

69

Galactic latitude (degree)

CHAPTER 6. ANALYSIS ON GC NORTH

1.0

N-Src 1

Reference region

N-Src 20.5

Sgr A*

0.0
1.0

0.5
0.0
−0.5
Galactic longitude (degree)−1.0

Figure 6.1. Close-up view of the GC North region in the 1.0–3.0 keV band take by Suzaku. The
green and white squares indicate the source and reference regions, respectively. The cross shows
the location of Sgr A*.

6.1

Analysis Method

We found GC North extending ∼1◦ north of Sgr A* from the ASCA image (Figure4.3). However, photon statistics and energy resolution of the GIS data are not sufficient forspectral investigation. Therefore, we used the Suzaku data, although the entire GC North isnot covered. Figure 6.1 shows the close-up view of the GC North region in the middle band
taken by Suzaku. GC North is not seen in the soft and the hard band (see Figure 4.1). We
chose the source and reference regions as shown in the green and white rectangles in Figure
6.1, respectively. Observations used in this analysis are listed in Table 6.1. The criteria ofthe data calibration and screening is the same as described in Section 4.1.1. The RMFs and
ARFs were also compiled in the same manner.
–70–6.2. BACKGROUND ESTIMATION
Table 6.1. Observation logs of the GC North region.
Sequence #α2000.0

Aim Pointδ2000.0lbObs. Date(yyyy-mm-dd)

Exposure(ks)

Obs. Name

22
22

GC19
GC22

24

GC20

Source regions
502007010
502010010

265.◦ 86266.◦ 38−28.◦ 45−28.◦ 340.◦ 160.◦ 500.◦ 660.◦ 33

2007-10-11
2007-10-13

Reference regions
502008010

6.2

265.◦ 66−28.◦ 73359.◦ 83

0.◦ 66

2007-10-12

Background Estimation

As same as the case of GC South, precise Galactic BG estimation is crucial for the
analysis of GC North. Therefore, we investigated a spectrum in the reference region (white
rectangle in Figure 6.1). The obtained spectrum is shown in Figure 6.2. Its spectrum shapeis similar to that in the Ref 2 region used in the GC South analysis (see Figure 5.2)
We applied the BG model constructed in Section 5.2. At first, all the parameters were
fixed to the values shown in Table 5.2 except for NH GC and f actor, because the interstellar
column density and the intensity of the GRXE are different from position to position. The
parameter of f actor was treated as a free parameter while the absorption column density was
fixed to a value estimated from near-infrared extinction (Gonzalez et al. 2012). The extinction in the reference region (l ∼ 0.◦ 17, b ∼ 0.◦ 66) is AKs = 1.21, so that NH GC was estimatedto be 2.2 × 1022 H cm−2 using the conversion factor of NH /AKs = 1.8×1022 cm−2 mag−1 (seeSection 5.2). The fitting result is plotted on Figure 6.2(a). This model roughly represents
the spectrum, but large residuals remain with χ2 /d.o.f = 541/237.
We then allowed to vary kTHP , kTLP , and EMFE for the fine-tuning of the model. As
a result, the fit is considerably improved with χ2 /d.o.f = 344/234, but small residuals still
remain at the energy of S XIII Kα line (2.4 keV) and the S XXV Kα line (6.7 keV). We
finally varied the abundances of Si and Fe, and obtained the best-fit model shown in Figure
6.2(b) with χ2 /d.o.f = 289/232. The best-fit parameters are listed in Table 6.2.
The obtained kTHP , kTLP , and EMFE are consistent with the original BG model. On theother hand, the abundances of Si and Fe are larger than 70% and 30% from the original BG
model, respectively. We confirmed that the flux in the 5–8 keV band is consistent with thatof the model composed by Uchiyama et al. (2013). Therefore this model is plausible.
–71–

CHAPTER 6. ANALYSIS ON GC NORTH

0.1

Counts s−1 keV−1
12
Energy (keV)

5

10

0
−5

χ

5

0.01

0.1
0.01
0
−5

χ

5

Counts s−1 keV−1

1

(b)

1

(a)

1

2
Energy (keV)

5

10

Figure 6.2. Spectra of the reference region. (a) the model which is the same as GC South BG.
(b) the model fine-tuned.

Table 6.2. Fitting results of the spectra in the reference regions for GC North.
— HP + LP + Gaussian —— FE —

Free parameters
f actorkTHP (keV)
kTLP (keV)
Si (solar)
Fe = Ni (solar)

1.38 ± 0.07
6.7 ± 0.6
0.81 ± 0.06
1.6 ± 0.2
0.69 ± 0.09∗
(1058 cm−3 )EMFE

1.26 ± 0.06

Fixed parameters
NH GC (1022 cm−2 )∗EMHP(1058 cm−3 )∗EMLP (1058 cm−3 )Mg = Ne (solar)
S (solar)
Ar (solar)
Ca (solar)
Other Elements (solar)
Fe I Kα flux (10−6 ph cm−2 s−1 )NH FE (1022 cm−2 )kTFE (keV)
Ne (solar)
Mg (solar)
Other Elements (solar)

2.2
0.38
0.45
0.43
1.6
3.5
2.8
1
6.4

0.42
0.33
0.06
0.19
0.022

Notes:R
Emission measure assuming the distance of 8 kpc. EM = ne np dV , where ne , np , and V , are theelectron density, the proton density, and the emitting volume, respectively.
∗–72–

6.3. SPECTRUM ANALYSIS

6.3

Spectrum Analysis

6.3.1

Merged spectrum

We extracted the spectra of GC North from the two rectangle regions shown in Figure 6.2(N-Src 1 and N-Src 2). Firstly, these spectra were co-added to increase the photon statistics(hereafter the merged spectrum). Figure 6.3 shows the merged spectrum. Based on theimage analysis in Section 6.1, the emissions below 1 keV and above 5 keV are likely to be
the FE and the GRXE, respectively. In order to subtract these component, we used the BG
model constructed in the previous section. The absorption column density of the GRXE isevaluated from the infrared extinction in this region, and found to be 2.5× 1022 H cm−2 .
The normalization factor is allowed to vary. Applying this BG model, we found the excessof the data with emission lines of Mg, Si, and S as shown in Figure 6.3.
These emission lines suggest an optically-thin thermal plasma origin, and thus we added
a CIE plasma model. Free parameters were NH , kTe , EM , and the abundances of Mg, Si
and S. The abundances of the other elements were fixed to the values in GC South (seeTable 5.3), because they are not determined due to poor photon statistics and/or strong
absorption of GC North. After fitting with this model, we found a Gaussian like residual at
1.2 keV, which is likely the missing Fe-L shell lines (see Section 5.3.1). Therefore, we also
added a Gaussian at 1.2 keV in the same manner as done in the GC South. The spectrum isrepresented by this model with χ2 /d.o.f = 236/185. The best-fit parameters are tabulated
in Table 6.3. We also attempted RP and IP models, but no significant improvement was
obtained.

6.3.2

Spatially resolved spectra

From the analysis of the merged spectrum, we found that the CIE model is favor in
GC North. In order to investigate the spatial variation of the plasma parameters, we thenanalyzed the spectra of N-Src 1 and N-Src 2 separately. For the BG model, the absorption
column density of NH GC in each region is re-estimated from the near-infrared extinction,and listed in Table 6.3.
Their spectra were fitted with the CIE models as same as the merged-spectrum fitting
but elemental abundances were fixed to the best-fit value of the merged spectrum. Figure
6.4 and Table 6.3 shows the fitting results and obtained plasma parameters, respectively.
We obtained the different electron temperature for each region–73–CHAPTER 6. ANALYSIS ON GC NORTHCounts s-1 keV-1

100

10-1

10-2
5

χ

0
-5

1

2
Energy (keV)

5

10

Figure 6.3. Merged Spectra of GC North fitted with the CIE model. The gray, red, and purple
lines indicate the BG model, the best-fit CIE model, and the missing Fe L-shell line respectively.
Only the co-added FI spectrum and model are displayed for visibility, although the FI and BI
spectra were simultaneously fitted.

Table 6.3. Fitting results of the merged spectrum.
Parameter (Unit)

Merged spectrum

N-Src 1

N-Src 2NH∗ GC (1022 cm−2 )NH (1022 cm−2 )EM † (1058 cm−3 )kTe (keV)
Mg (solar)
Si (solar)
S (solar)SB‡χ2 /dof

2.5 (fix)
2.6+0.3−0.2
3.4+0.8−0.6
0.73 ± 0.04
0.9+0.3−0.2
0.9 ± 0.1
1.2 ± 0.2
2.2+0.5−0.4
236/185

2.0 (fix)
2.5 ± 0.2
2.5+0.6−0.5
0.61 ± 0.04
0.9 (fixed)
0.9 (fixed)
1.2 (fixed)
2.3+0.6−0.5
216/185

2.8 (fix)
3.3 ± 0.3
2.2+0.8−0.5
0.76 ± 0.06
0.9 (fixed)
0.9 (fixed)
1.2 (fixed)
2.7+1.0−0.6315/220

Notes:* Absorption column density for the BG emissions that were estimated from thenear-infrared extinction investigated by Gonzalez et al. (2012).
† Emission measure assuming the distance of 8 kpc.
‡ Absorption-corrected surface brightness in the 1.5–3.0 keV band with a unit of
10−14 erg s−1 cm−2 arcmin−2–74–

6.4. PHYSICAL PROPERTIES OF GC NORTH(a)(b)
Counts s-1 keV-1

100

10-1

10-2

10-1

10-2

5

5

0

0

-5

χ

χ

Counts s-1 keV-1

100

1

2
Energy (keV)

5

10

-5

1

2
Energy (keV)

5

10

Figure 6.4. Spectra extracted from (a) N-Src 1 and (b) N-Src 2. The gray, red, and purple lines
indicate the BG models, the best-fit CIE models, and the missing Fe L-shell lines respectively. Only
the co-added FI spectra and models are displayed for visibility, although the FI and BI spectra
were simultaneously fitted.

6.4

Physical Properties of GC NorthAs shown in Table 6.3, the obtained absorption column densities for the N-Src 1 and
N-Src 2 regions are consistent with estimated NH GC , whose typical errors are ∼ 15%. We
therefore concluded that GC North is truly located at the GC distance (8 kpc). At this distance, apparent size of the N-Src 1 and N-Src 2 regions are 900 pc2 and 860 pc2 , respectively.
Assuming the length along the line of sight to be 70 pc, we estimated the volume ofthe plasma, and calculated physical parameters. The results are summarized in Table 6.4.
The density and pressure in N-Src 1 is slightly smaller than that in N-Src 2. The total massand thermal energy are ∼460 M and ∼1.5 × 1051 erg, respectively. Since the entire sizeof GC North seen in the ASCA image is about 2 times larger than that of analysis region,
total mass and thermal energy probably reach ∼900 M and ∼3 × 1051 erg, respectively.
The sound velocity of the plasma (kT ∼ 0.7 keV) is 400 km s−1 . Therefore we have thedynamical timescale of ∼4 × 105 yr for crossing the plasma size of ∼ 150 pc.

–75–CHAPTER 6. ANALYSIS ON GC NORTH

Table 6.4. Physical parameters of the N-Src 1 and N-Src 2.
RegionN-Src 1
N-Src 2V∗
(1060 cm3 )ne(10 cm−3 )

1.7
1.6

1.2 ± 0.2
1.2 ± 0.2

−1

−10

(10

p
erg cm−3 )1.8 ± 0.2
2.1 ± 0.4

M
(10 solar)

Eth(1050 erg)

2.4 ± 0.3
2.2 ± 0.4

7.0 ± 1.0
7.8 ± 1.6

2

Notes: Quoted errors include only statistical errors in the spectrum fitting.
* Assumed the line of sight length of 70 pc.

–76–

Chapter 7
Discussion
In this chapter, we summarize the analysis results in the preceding chapters, and discusspossible origins of GC South and GC North.

Contents
7.1

Summary of the Obtained Parameters . . . . . . . . . . . . . . .

78

7.2

Origin of GC South . . . . . . . . . . . . . . . . . . . . . . . . . .

79

7.2.1

Energy source . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

79

7.2.2

Mechanism to form the recombining plasma . . . . . . . . . . . . .

82

7.3

Origin of GC North . . . . . . . . . . . . . . . . . . . . . . . . . .

85

7.4

Connection between GC South and GC North . . . . . . . . . .

85

77

CHAPTER 7. DISCUSSION

Table 7.1. Typical physical parameters of GC South and GC North
Region

kTe
(keV)

ne
(cm−3 )

Ionization state

Abundance∗(solar)

Eth
(erg)

M
(M )

t†rec
(year)t‡dyn(year)GC South
GC Nouth

0.4–0.8
0.7

0.06–0.12
0.12

RP or CIE
CIE

0.7–0.8
0.9–1.2

8 × 1051
3 × 1051

3000
900

3 × 105—

4 × 105
4 × 105

Notes:* Abundances of Mg, Si, and S.
† Recombination timescale of the plasma‡ Dynamical timescale (size/sound velocity) of the plasma.

7.1

Summary of the Obtained Parameters

In the preceding chapters, we derived the physical parameters of GC South and GC
North. Typical parameters are tabulated in Table 7.1. We also summarized the spatial
variation of the plasma parameters along the Galactic latitude in Figure 7.1 and Figure 7.2.
Figure 7.1 shows the profile of the absorption column density for the GC South/North
compared with the near-infrared extinction (Gonzalez et al. 2012). We adopted the conversion factor of NH /AKs = 1.8 × 1022 H cm−2 mag−1 derived in Section 5.2. The absorption
column density for the GCDX reported by Uchiyama et al. (2013) is also potted as a reference. There is a ∼40% discrepancy between the X-ray absorption and the near-infrared
extinction at b = 0◦ . Gonzalez et al. (2012) noted that extinction in the GC region (AKs > 2
or NH > 3.6 × 1022 H cm−2 ) is underestimated by a few tens of percent because of incompleteness of the sample. Thus, this discrepancy is mainly due to the observational bias. The
gradient of the X-ray absorption well agrees with that of the near-infrared extinction, whichindicates that GC South and GC North are located at the GC distance.
Figure 7.2 shows the spatial profiles of the surface brightness, temperature, density andpressure. We also plotted those of LP and HP in the GCDX as a reference. We referred to
the result reported by Uchiyama et al. (2013) for the surface brightness. Other parametersof the GCDX are derived from Muno et al. (2004), in which the central 170 × 170 filed wasinvestigated. Since the parameters of GC North are determined predominantly in the 1.5–3.0
keV band, absorption-corrected surface brightnesses of all the regions, including GC Southand the GCDX, were calculated in this band. We see that all of the parameters decline
smoothly from the GC to the high Galactic latitude and are nearly symmetrical about theGC, although the data in the positive latitude is insufficient.
At first, we consider a possible origin of GC South using these information. Then, we
discuss the origin of GC North and revealed the connection between GC South and GC
North.
–78–

7.2. ORIGIN OF GC SOUTH

Column density (1022-2

H cm )

10

1

10

-1

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

Figure 7.1. Absorption column densities along the Galactic latitude obtained from the GC
South and GC North (the red circles), the GCDX (the bule squares; Uchiyama et al. 2013),
and the infrared extinction (the black triangles; Gonzalez et al. 2012) converted by the factorof 1.8 × 1022 H cm−2 mag−1 .

7.2

Origin of GC South

7.2.1

Energy source

Typical thermal energies observed in Galactic X-ray SNRs are ∼1049−50 erg, which is
1–10% of the total energy released by a single SN (1051 erg). We found the thermal energyof GC South to be 8 × 1051 erg. Therefore, GC South is about two orders of magnitude moreenergetic than a single SNR, and hence a single SN origin is rejected. In fact, non-thermal
radio shell, which is evidence for an SNR, is not found near GC south (e.g., LaRosa et al.
2000).
As seen in Galactic superbubbles, multiple SNe and stellar winds in an active starforming region can provide sufficient energy for the formation of GC South. However, there
is no active star-forming region at the location of GC South except for the S-Src 1–3 filed(see Section 5.3.2). Since some of the H II regions found in the S-Src 1–3 field are locatedat the foreground of the GC, they are not associated with the GC South. Moreover, the–79–CHAPTER 7. DISCUSSION

10

(b)

Temperature (keV)

101

SB (10

-14

-1

-2

2

erg s cm arcmin )

(a)

10

-1

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

1

10

0.5

-1

(c)

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

(d)

-3

Density (cm )-3erg cm )

10

-10

-1

Pressure (10

10

-2

-1.5

-1
-0.5
0
Galactic latitude (degree)

0.5

1

Figure 7.2. Spatial profiles of (a) surface brightness, (b) electron temperature, (c) electron density,and (d) thermal pressure along the Galactic latitude. The red circles indicates the parameters of
GC South and GC North. The blue triangles and squares are the parameters of LP and HP in the
GCDX, respectively (Uchiyama et al. 2013; Muno et al. 2004).

temperature and density of GC South are higher than those of typical Galactic superbubbles(kT < 0.3 keV and n ∼ 0.01 cm−3 ). Therefore, the superbubble origin is not favored for GC
South.

Remaining energy sources are GC activities. We found that the recombining timescaleof GC South is ∼3 × 105 yr, and is consistent with the dynamical timescale (see Table
7.1). Therefore, required energy-injection rate is estimated as 8 × 1051 ergs/3 × 105 yr =9 × 1038 erg s−1 . We examine whether GC activities can provide this energy in the following
section.
–80–

7.2. ORIGIN OF GC SOUTH
GC activities at present
The current SN rate in the nuclear region (r < 50 pc) is estimated to be only one per
10 yr, corresponding to 1038 erg s−1 (assumed 10% efficiency). It is not sufficient to form
GC South. Moreover, the current Sgr A* luminosity of 1033 –1035 erg s−1 in the X-ray bandis also 4–6 orders of magnitude lower than the required energy. Therefore, the GC activities
at present are not able to form GC South.
4

Past starburst
One possible energy source is a past nuclear starburst activity. In starburst galaxies, weusually see outflows of plasmas with temperature of 0.3–2.0 keV, the so-called superwinds.
The superwind in NGC 253 has the size and the energy of ∼1 kpc and ∼1053 erg, respectively(Strickland et al. 2000). Although they are about one orders of magnitude higher than thoseof GC South, other physical features (temperature, density, and morphology) are similar.
Yusef-Zadeh et al. (2009) claimed that the GC SFR 105 –107 yr ago is one order ofmagnitude higher than that at present. The past starburst scenario is also proposed as theorigin of the Fermi bubbles, although their timesscale of 107 yr is different from that of GC
South (105 yr). If the nuclear SN rate was one order of magnitude higher than the current
rate 105 yr ago, the energy of ∼1039 erg s−1 was provided and GC South was formed like asuperwind.
We then check the consistency between the superwind origin and the obtained plasma
parameters. In order to blow out the plasma from the GC, the pressure of the ISM in the
Galactic bulge should be lower than that of the plasma. According to an model of the
ISM distribution (Ferrière et al. 2007), we found that the density of H I and H+ in the GC
South region are < 5 cm−3 and ∼0.15 cm−3 , respectively. The contribution of H2 clouds is
negligible in this region (see Figure 2.3). Then, the thermal pressure of H I (∼103 K) andH+ (∼104 K) is < 10−12 erg s−1 and ∼10−13 erg s−1 , respectively. They are two orders ofmagnitude lower than the thermal pressure of GC South, and hence plasma can be blownout from the GC. In Figure 7.2(d), we can also see that the thermal pressure decreases fromthe GC to the southern edge of GC South. These facts support the superwind origin.
Past Sgr A* jet
Another possible energy source is a jet from the central SMBH. Some AGNs exhibit
thermal plasma emissions, which are produced by jet-cloud interactions (e.g., NGC 4258 andNGC 4151). These plasmas have the size of a few kpc and the temperature of ∼0.5 keV, andresemble GC South.
–81–CHAPTER 7. DISCUSSIONAs reviewed in the Section 2.3.3, Sgr A* was active 102 –103 yr ago with the luminosity of1039 erg s−1 (Ryu et al. 2013). If Sgr A* was more active 105 yr ago, it probably emitted a jetwhose energy is comparable to the luminosity of the central engine. Then, the jet possibly
generated a shock-heated hot plasma in the GC South region. Such a jet-like activity 106 yr
ago is also proposed as the origin of the Fermi bubbles although it is still controversial.
Thus, the past Sgr A* jet origin is not ruled out.

7.2.2

Mechanism to form the recombining plasma

We discovered the RP in GC South (S-Src 4–10), which is not explained in the standard
evolution of a shock-heated plasma; a shock-heated low-density (∼1 cm−3 ) plasma is inan ionizing phase for a long period because the time scale of electron heating (∼100 yr)
is much shorter than that of collisional ionization (∼104 yr). Therefore, to realize an RP,
some specific events should occur: either a rapid decrease in electron temperature or anenhancement of only the ionization. We discuss possible scenarios in the following.
Adiabatic expansion scenario
In Galactic RP SNRs , two mechanisms for a decrease in electron temperature have been
proposed (e.g., Sawada & Koyama 2012; Uchida et al. 2012), and supported by numericalsimulations (e.g., Zhou et al. 2011; Shimizu et al. 2012). We attempt to apply these scenarios
to GC south. One is thermal conduction via interactions with ambient cold molecular clouds.
However, no such molecular clouds have been found in this region. Moreover, electron
temperature is lower in the higher Galactic latitude, where cold gas density become lower(Ferrière et al. 2007). Therefore this scenario is unlikely.
The other scenario is adiabatic expansion of the plasma. In this scenario, the CIE
plasma was formed in a dense circumstellar medium at the starburst site. Then, the plasma
were blown out rapidly to a thin circumstellar space, and the electron temperature rapidly
decrased as a result of adiabatic expansion. The electron temperature is expected to decrease
with distance away from the GC, and the ionization state should be CIE near the GC while
it becomes an over-ionized state at the far side of the GC. The observational results are
consistent with this pictures (see Figure 5.8). In adiabatic condition, T ∝ nγ−1 is also
expected. For ideal monatomic gas (γ = 5/3), the index of the above equation is 2/3.
Figure 7.3(a) shows the correlation between the observed kTe and ne . We can see a positive
correlation between them, and fitting with a power-law function gave us the best-fit index
of 0.4 ± 0.3. It is consistent with the adiabatic equation within the 90% confidence interval
as shown in the red bow tie in Figure 7.3(a). Although kTe and ne (or emission measure)are coupled with each other in the spectral fitting, their correlation is negative as shown in–82–

7.2. ORIGIN OF GC SOUTH(a) Temperature v.s. Density
2(b) Error contours
0.48

0.47
kTe (keV)

kTe (keV)

1

0.5

ba

Adia

0.2

nsion

xpa
tic e

0.05

x=
(inde

2σ
2σ

0.46

3σ

1σ
3σ

0.45

2/3)

-3

Density (cm )

0.1

0.2

0.44
2.4

2.5

2.6

2.7

2.8

Emission measure (1058 cm-3)

Figure 7.3. (a) Correlation between the temperature and density of GC South derived from thespatially resolved fitting. A red bot tie shows the best-fit power-law function. A dotted blue line
indicates the adiabatic process of T ∝ n2/3 (see text). (b) Typical error contours between thetemperature and the emission measure in the spectral fitting. The contours show 1σ, 2σ, and 3σ
levels.

Figure 7.3(b). Therefore, the positive correlation between kTe and ne is reliable. Then, theadiabatic expansion scenario is favored for GC South.
We further examine whether plasmas located in the current GC region satisfy the initialcondition of GC South. The initial plasma should have the temperature of 1.6 keV (seeSection 4.3.1). From a simple extrapolation of the adiabatic equation obtained from the kTe –ne relation (Figure 6.3), we obtained the initial plasma density of 0.6 cm−3 . The densityof the GCDX, both the LP and the HP, is smaller than the expected value by a factorof 3–4. Mori et al. (2009) discovered the GC superbubble, which has kTe = 0.91 keV andne = 0.26 cm−3 , but they does not agree with the initial plasma parameters of GC South. The
Sgr A East SNR is only a possible candidate. It consists of the low temperature (kTe = 1.2
keV, ne = 4.6 cm−3 ) and high temperature (kTe = 6 keV, ne = 0.6 cm−3 ) plasma (Koyama
et al. 2007a). These parameters are close to or larger than the required parameters. The
past starburst activity might form a large amount of plasmas like the Sgr A East SNR.
Sgr A* photoionization scenario
Not rapid electron cooling but selective ionization also forms an RP. One possibility
is photoionization; a plasma is illuminated by a strong X-ray source and is initially over–83–

CHAPTER 7. DISCUSSION
ionized. After the termination of photoionization, the plasma enters a recombining phase.
The balance between photoionization and recombination is described by a single parameter,ξ = L/(nR2 ), where L, n, and R are the luminosity of the source, the gas density, and the
distance between the irradiating source and the plasma, respectively. To achieve nearly thesame ionization state as the 1.6 keV CIE plasma, ξ should be ∼1000 (Kallman & Bautista2001). Then, the required luminosity is estimated as

43

L = 3.6 × 10

R
200 pc2 nerg s−1 ,0.1 cm−3(7.1)

where R is assumed to be the size of GC South. Only a past Sgr A* flare with the nearly
Eddington luminosity (∼1044 erg s−1 ) can achieve such a extremely high luminosity. The flareshould not be isotropic radiation, because the GCDX, which is in the vicinity of Sgr A*, is not
an RP but a CIE plasma (Koyama et al. 2007b). Therefore, GC south must be illuminatedby collimated radiation like a jet, if the photoionization scenario is the case.
We discussed a Sgr A* jet as the possible energy source of GC South in the previous
section. Assuming this jet also emitted collimated X-rays such as a blazar activity, theplasma was formed and photoionized simultaneously. In this scenario, entire GC Southregion should became an RP. Moreover, a higher ionization state is expected as the plasma
is closer to Sgr A*, because the radiation field is stronger near Sgr A*. However, the observed
spatial variation of the ionization states is opposite to this prediction. As shown in Figure
5.8(d), the plasmas in S-Src 1–3 are not RPs but CIE plasmas, and there is no significant
variation of the ionization states in the S-Src 4–10 region. Therefore this scenario is notfavored.
Alternatively, we consider another scenario, where a plasma was generated by a nuclear
starburst activity and was then photoionized by a Sgr A* jet in the course of an outflow.
Rapid electron cooling due to adiabatic expansion is not necessary in this scenario. Assuming
the starburst-region size of 50 pc, required luminosity is estimated to be ∼1042 erg s−1 from
Eq.7.1. If Sgr A* became quiescent in the middle of the starburst period, no more RP wasformed. Thus, this scenario is able to explain the composite ionization state of GC South;
plasmas generated in early phase (S-Src 4-10) were photoionized by a Sgr A* jet, but Sgr A*
became quiescent after photoionizing the S-Src 4 plasma. Then, the plasma formed later(S-Src 1-3) were no more photionized and kept the CIE states. Since the boundary between
S-Src 3 and S-Src 4 is 150 pc south of Sgr A*, phase transition of Sgr A* should be occur
3 × 105 year ago when the outflow speed is 510 km s−1 . It is consistent with the recombining
timescale of S-Src 4 within a factor of 2. Therefore, this scenario is cannot be excluded.
–84–

7.3. ORIGIN OF GC NORTH

Galactic latitude (degree)

1.0

(a)

(b)
GC North

GC North

0.5-35 km s-1 Ridge

0 km s-1 Ridge

CMZ

CMZ

0.0
−0.5 CMZ

−1.0
1.0

CMZ

0.0

359.0 1.0
Galactic longitude (degree)

0.0

359.0

Figure 7.4. Images of CO (J=2–1) emission in the velocity range of (a) −3–0 km s−1 and (b)−39–−31 km s−1 (Enokiya et al. 2013). The overlaid green contours are the X-ray emission derived
from the ASCA image in the 1.5–3.0 keV band (Figure 4.3).

7.3

Origin of GC North

The thermal energy of GC North is about one order of magnitude higher than that oftypical SNRs. In addition, star forming region is not found near GC North. Therefore, asingle or a multiple SN origin is unlikely.
The most plausible energy source is a GC activity as proposed for GC South. Since therequired energy-injection rate estimated from the dynamical timescale is 3 × 1038 erg s−1 ,which is comparable to that of GC South, the current GC activities cannot provide sufficientenergy as discussed in Section 7.2.1. Therefore, a past starburst or a jet from Sgr A* arepossible origins for GC North.

7.4

Connection between GC South and GC North

We show that both GC South and GC North were formed by past GC activities ∼105 yr
ago. Since they have nearly the same timescales, we suspect that the identical past activity
formed both GC South and North. Indeed, parameters of kTe , ne , abundance, Eth , andM are similar between GC South and GC North as shown in Table 7.1, and the identical–85–

CHAPTER 7. DISCUSSION
origin is supported. If this is the case, a past nuclear starburst is the energy source of GC
Sout/North from the discussion of Section 7.2.1. However, morphology and ionization states
are different from each other. We consider the cause of these differences, although spatial
variation of plasma parameters for GC North is poorly constrained,
The emission of GC South extends to b = −2.◦ 0 along l = 0.◦ 0, while that of GC Northextends to only b = 1.◦ 2 and shifted to the east (l ∼ 0.◦ 25). One possible cause of the morphological difference is asymmetrical distribution of the ISM. Enokiya et al. (2013) discovered
the molecular ridges (the 0 km s−1 ridge and the −35 km s−1 ridge) to the north of the GC,
which have the mass of ∼105 M . Figure 7.4 shows the images of CO (J=2–1) overlaid with
ASCA X-ray contours. We found that the X-ray emission of GC North shows anticorrelation
with the molecular ridges. Assuming the its density of ∼103 cm−3 (critical density of CO
J=2–1) and the temperature of ∼70 K (typical value of the CMZ), the thermal pressure of
the molecular ridge is ∼1010 erg cm−3 and comparable to that of GC North. Therefore, we
consider that both or either molecular ridge has prevented the plasma from outflowing anddistorted the plasma shape.
We found no RP in GC North contrary to in GC South. One simple explanation is that
the Suzaku observations have not covered the RP region in GC North. Since we discovered
the RP in b < −1.◦ 0, the plasma in b > 1.◦ 0 could be an RP. Another explanation is thatan RP was not formed in GC North. As shown in Figure 7.4, the GC North plasma may
be prevented from expanding by the molecular ridges. If adiabatic expansion scenario isthe case for GC South (see Section 7.2.2), the GC North plasma was not able to expand
adiabatically and therefore did not become an RP. In any case, feature Suzaku observationsof the GC North region will reveal the true. We finally summarized the activity history on
Figure 7.5.

–86–

7.4. CONNECTION BETWEEN GC SOUTH AND GC NORTHFB

Relased energy rate (erg s-1)

1042

(Sgr A* jet origin)

FB
1040 (Starburst origin)

Jet?

GC South/North

XRN

1038

10% SN rate(r < 50 pc)
Current
Sgr A*

1036

Flare

1034
Quiescent

1032

107

106

105

104

103

102

101

Look back time (years)

Figure 7.5. Schematic diagram of the revealed GC activities.

–87–

CHAPTER 7. DISCUSSION–88–

Chapter 8
Conclusion
In this thesis, we performed the survey of the south and north of the GC with Suzaku,
and derived the following results.
• We found the large diffuse plasmas extending south and north perpendicular to the
Galactic plane: GC South and GC North.
• GC South is a thin thermal plasma in a recombining phase, and is located at the GC
distance. This is the first sample of an RP in the GC region. Fitting with a simple
RP model gave us the plasma parameters of kTe = 0.5 keV, kTinit = 1.6 keV, andne t = 5 × 1011 s cm−3 . The size of GC South is about 200 pc. Then, the density, total
thermal energy and mass are 0.1 cm−3 , 8 × 1051 erg and 3000 M , respectively.
• The electron temperature, density, and pressure of GC South smoothly decline withincreasing the distance from the GC. The plasma near GC is in a CIE state while thatof far side is in a recombining phase.
• GC North is also located at the GC distance, but its ionization state is different fromthat of GC South: an CIE plasma with kTe = 0.8 keV. The size of GC North is about
100 pc, and the its density is found to be 0.1 cm−3 . The total thermal energy and massreach 3 × 1051 erg and 900 M , respectively. The dynamical timescale of the plasmais ∼ 105 yr.
• Both GC South and North is not a single SNR or a superbubble. The mini-superwind
due to nuclear starburst activity ∼ 105 is a possible scenario. An RP in GC Southis likely to be formed by the rapid electron cooling due to adiabatic expansion of theoutflowing plasma.

89

CHAPTER 8. CONCLUSION–90–

Appendix A
Emission from a Thermal Plasma
A.1

Ionization State in a Plasma

A thermal plasma is one in which free-electrons distributed in the Maxwell-Boltzman
distribution:


3/2

me
me ve22
f (ve ) = 4πve
exp −
,
(A.1)
2πkTe
2kTe
where ve , me , k, and Te are electron mass, electron velocity, the Boltzman constant, andelectron temperature, respectively. In this chapter, we focus on the plasma under the condition of high electron temperature (kTe ≈ 1 keV) and low electron density (ne < 1 cm−3 ),which is typical for the X-ray emitting plasma in the universe.
In a thermal plasma, collisional ionization occur when a free electron has a kinetic energy
higher than binding energy and collides with an ion. It is also possible for a free electron
to be captured by an ion into a bound state. This process is called recombination. Ionization degree of ions are deterred by the competition between ionization and recombination
rates; the plasma becomes ionization equilibrium state when an ionization rate is equal to a
recombination rate.
Since density of hot plasmas is low, non-ionization equilibrium state is usual in theuniverse. If the ionization degree is low compared with the electron temperature (underionized state), the plasma is dominated by an ionization process and is referred to as an“ionizing plasma”. Such a plasma is widely observed in young SNRs, where the plasmas
were recently shock-heated. In the opposite condition (over-ionized state), a recombining
process is enhanced, and thus the plasma is referred to as an “recombining plasma”. Sucha condition can be seen in several mixed-morphology SNRs or photoionized plasmas aroundbright X-ray sources.
91

CHAPTER A. EMISSION FROM A THERMAL PLASMA

A.2

Radiation Process in a Thermal pPasma

Bremsstrahlung
When a free electron is scattered by an ion, photon is emitted, the so-called bremsstrahlungor free-free emission. An emissivity of bremsstrahlung produced by the interaction between
thermal electrons and ion is represented by
dW
32πe6
=
dV dtdν3me c2



2π
3kme

1/2
T

−1/2



hν
Z ne ni exp −kTe
2

gf f (Te , ν),(A.2)

where c, Z, ni , h, are gf f the speed of light, the atomic number of a target ion, ion density,
the Planck constant, and the velocity averaged Gaunt factor. For a typical thermal plasma
emitting X-rays, the velocity averaged Gaunt factor is approximated as
gf f (Te , ν) =

3kTe
πhν

−1/2
.

(A.3)

Taking the solar abundance of elements into account, hydrogen is the major element leading
to the bremsstrahlung in a plasma.

Line emission
A bounded electron is excited by a collision of free electrons in a thermal plasma. Subsequently, de-excitation of the electron often leads to photon emission called line emission
or bound-bound emission. In the low density regime, the so-called coronal limit, an excitedelectron immediately de-excited. Therefore, line flux is determined by a collisional excitation
rate and is proportional to ne ni .
Since flux of the bremsstrahlung is almost proportional to ne nH , we can derive the relative
abundance of a heavy element to hydrogen (ni /nH ) from the ratio of the line flux and thebremsstrahlung flux. Line emission is useful for plasma diagnostics. We are also able toobtain the ion population in an element by using the flux ratio of two lines which come froman identical element but different ionization states.

Radiative recombination continuum
Photon emission associated with recombination is called radiative recombination continuum or free-bound emission. The energy of the photon equals the kinetic energy of theelectron plus the binding energy where the electron recombines. Therefore, the major feature–92–A.3. SPECTRUM OF A RECOMBINING PLASMAof RRC spectrum is a edge, which corresponds to the binding energy. The emissivity of the
RRC with a thermal electron distribution is given bydW= 4ne ni EdV dtdE



E−I
kTe



1
2πme kTe

1/2



E−I
exp −
kTe



σnrec ,

(A.4)

where E, I, and σnrec are emitted photon energy, binding energy, and the cross section of recombination. According to Nakayama & Masai (2001), σnrec for an H-like ion is approximatedas
−1
1 3kTe1recσn ∝ 3
(A.5)
+ 2
n
2I
n
Therefore, recombination to the lowest levels is the major process. If a free electron recombines to excited state, de-excitation of the electron subsequently occurs. The line emission
emitted by this process is especially called the cascade line.

A.3

Spectrum of a Recombining Plasma

In an ionization-equilibrium plasma with high electron temperature (kTe > 0.1 keV),
RRC is weaker than the bremsstrahlung, and thus there is no apparent RRC features in an
X-ray spectrum. However, in a recombining plasma, a spectrum shows strong RRC features
since the majority of the ions are recombining (Figure A.1).

–93–

CHAPTER A. EMISSION FROM A THERMAL PLASMA

100

Equillibirum plasma
Recombining plasmaphtons s-1 cm-2 keV

10-1
10-2

Edge of
S RRC

10-3
10-4

1

Edges of
Si RRC
2
Energy (keV)

3

4

Figure A.1. Spectrum models of thermal plasmas with kTe = 0.5 keV. The ionization equilibrium
and recombining plasmas are shown in black and red, respectively.

–94–

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Acknownledgments
My cordial gratitude goes to my supervisors Prof. Takeshi Go Tsuru and Prof. Katsuji
Koyama who gave me the precious opportunity to conduct this interesting study and has
been encouraging me throughout my graduate course.
I am also appreciate to Dr. Masayoshi Nobukawa, Dr. Takaaki Tanaka, Dr. Hiroyuki
Uchida, and Dr. Hideki Uchiyama who gave me a number of invaluable comments to mystudy, including English proofreading.
I would like to thank the members of the Cosmic-ray laboratory in Kyoto University forenjoying the academic life together.
Finally, I express supreme thanks to all of my family and friends for their warm supportand encouragements.