Study of X-ray Emission from the Giant Molecular Clouds in
the Galactic Center Region

Masayoshi Nobukawa
Department of Physics, Graduate School of Science, Kyoto University
Kitashirakawa Oiwake-cho, Sakyo-ku, Kyoto, 606-8502, Japan
nobukawa@cr.scphys.kyoto-u.ac.jp
January 5 2011

Abstract
There are K-shell emission lines of neutral Fe (Fe I Kα) correlated with giant molecular
clouds in the Galactic center (GC) region. The Fe I Kα emission is one of the key forthe structure and activity of the GC, including the super-massive black hole Sagittarius
A* (hereafter Sgr A*) with a huge mass 4 × 106 M¯ . The origin of the Fe I K is still
open question, but possibly due either to X-ray radiation from an external object (X-ray
reflection nebula: XRN) or to bombardment of low energy cosmic-ray electrons (LECRe)to neutral atoms in dense molecular clouds. In order to resolve the problem, we madedeep observations of molecular clouds emitting the Fe I Kα line in the Sagittarius (Sgr)A, B1, B2, and C regions with the Japanese X-ray scientific satellite Suzaku.
We discovered Kα lines of neutral Ar, Ca, Cr, Mn atoms in addition to those of Fe andNi from a cloud in the Sgr A region. The equivalent widths of these emission lines requirethat the metal abundances in the GC region are ∼1.6 and ≥ 4 solar values in the XRN
and LECRe scenarios, respectively. On the other hand, the X-ray spectrum of the hot
plasma extended over the GC region exhibits the metal abundances of 1–2 solar values.
The XRN model agrees with these results.
We detected X-ray emissions with strong Fe I Kα lines from molecular clouds in the
Sgr B1 and C regions. The X-ray spectra consist of the Fe I Kα lines with large equivalent
widths of 1–2 keV and deeply-absorbed hard continua (column density NH ∼ 1023 cm−2 ,photon index Γ ∼ 1.7), and hence these clouds are also likely to be XRNe.
We found that the Fe I Kα fluxes of the two clouds in the Sgr B2 region had decreased
in correlation to the hard-continuum fluxes by factor of 0.4–0.5 from 2005 to 2009. The
time interval of 4 years is nearly equal to the light-traveling across the clouds. The rapid
and correlated time-variability in addition to the large equivalent width of Fe I Kα are
consistently explained by the XRN scenario due to an X-ray flare from an external sourcewith the duration time of shorter than 4 years. The peak-position of the Fe I-Kα emission
shifted slightly to a larger distance from the GC, Sgr A*, and hence the X-ray flare source
should be located in the Sgr A*-side with respect to Sgr B. Based on these facts, theluminosity of the X-ray flare is needed to be larger than 1040 erg s−1 , which is 50 times
larger than the Eddington limit of an object with the sun’s mass. Therefore Sgr A* is anunique candidate to illuminate the clouds in X-ray.
When the Sgr A* flare is applied to the origin of the Fe I K emission from the otherclouds, the required luminosities are similar values of ∼ 1041 erg s−1 . Thus, we reached aconclusion that the super-massive black hole Sgr A* made a single X-ray flare 300 years
ago and illuminated the molecular clouds in the vicinity.
3Contents
1 Introduction
1.1 Overview of the central region in Milky Way Galaxy . . . . . . . . . . . . .
1.2 SMBH, Sgr A* . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
1.2.1 X-ray Observations of Sgr A* . . . . . . . . . . . . . . . . . . . . . .
1.2.2 Gamma-rays from the direction near Sgr A* . . . . . . . . . . . . . .
1.3 Diffuse X-Ray Emission from the GC Region . . . . . . . . . . . . . . . . .
1.3.1 Discovery of Fe Lines from the GC Region . . . . . . . . . . . . . . .
1.3.2 Origin of Fe XXV and Fe XXVI . . . . . . . . . . . . . . . . . . . . .
1.3.3 GC Hot Plasma . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
1.4 Fe I Kα Emission from Giant Molecular Clouds . . . . . . . . . . . . . . . .
1.4.1 Possible Origins for the Fe I Kα Emission . . . . . . . . . . . . . . .
1.4.2 Recent Results—time variability of X-ray emissions— . . . . . . . .
1.5 Aim of this thesis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

1
1
1
2
4
5
5
7
8
10
10
11
13

2 Suzaku X-ray Satellite
2.1 Suzaku Satellite . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2 X-Ray Telescope . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3 X-Ray Imaging Spectrometer . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.1 Introduction of X-ray Charge Coupled Device Cameras . . . . . . .
2.3.2 Overview of XIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.3 Charge Injection Capability . . . . . . . . . . . . . . . . . . . . . . .
2.3.4 XIS Performance in orbit . . . . . . . . . . . . . . . . . . . . . . . .

15
16
16
23
23
24
27
30

3 K-shell Lines of Neutral Atoms
3.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2 Observations and Data Reduction . . . . . . . . . . . . . . . . . . . . . . .
3.3 Analysis and Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.1 6.4 keV line Image . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.2 Gain Tuning . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.3 Model Construction . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.4 Spectral Fitting . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.4 Discussion for the X-ray Emission from the Sgr A region . . . . . . . . . . .
3.4.1 The GC hot plasma . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.4.2 Origin of the Neutral Clump . . . . . . . . . . . . . . . . . . . . . .

37
37
38
39
39
40
41
43
46
46
47

i

ii

CONTENTS

3.5

Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 49

4 Other 6.4 keV Clouds
4.1 Sgr B1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.1 Introduction for Sgr B1 . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.2 Observation and Data Reduction . . . . . . . . . . . . . . . . . . . .
4.1.3 Analysis and Results . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.4 Discussion for M 0.51−0.10 . . . . . . . . . . . . . . . . . . . . . . .
4.1.5 Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2 Sgr C . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.1 Introduction for Sgr C . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.2 Observation and Data Reduction . . . . . . . . . . . . . . . . . . . .
4.2.3 Analysis and Results . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.4 Discussion for the 6.4 keV Clouds in Sgr C . . . . . . . . . . . . . .
4.2.5 Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

53
54
54
54
55
60
63
65
65
65
65
71
76

5 Time-Variability of X-rays from Sgr B2
5.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.2 Observation and Data Reduction . . . . . . . . . . . . . . . . . . . . . . . .
5.3 Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3.1 X-ray Images . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3.2 Spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.4 Discussion for the 6.4 keV Clouds in Sgr B2 . . . . . . . . . . . . . . . . . .

77
77
78
78
78
81
82

6 Discussion
85
6.1 Features of the 6.4 keV Clouds . . . . . . . . . . . . . . . . . . . . . . . . . 85
6.2 Irradiating Source of the 6.4 keV Clouds . . . . . . . . . . . . . . . . . . . . 87

List of Figures
1.1
1.2
1.3
1.4
1.5
1.6
1.7
1.8

Surrounding environment of Sgr A* . . . . . . . . . . . . . . . . . . . . . . 2
X-ray flare from Sgr A* . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4
X-ray spectra and 6.7 keV line intensity distribution (Ginga) . . . . . . . . 6
X-ray spectrum of the GCDX obtained by ASCA . . . . . . . . . . . . . . . 7
X-ray spectrum of the GCDX obtained by ASCA . . . . . . . . . . . . . . . 8
X-ray map in the 6.4 keV line obtained by ASCA . . . . . . . . . . . . . . . 10
Time variabilities of the 6.4 keV fluxes from the Sgr B2 clouds . . . . . . . 11
Fe I Kα image of the different XMM-Newton observations . . . . . . . . . . 12

2.1
2.2
2.3
2.4

15
17
17

Schematic view of the Suzaku satellite . . . . . . . . . . . . . . . . . . . . .
X-ray telescope onboard the Suzaku satellite . . . . . . . . . . . . . . . . .
Wolter type-I grazing incident optics in the XRT . . . . . . . . . . . . . . .
Comparison of effective areas of the X-ray telescopes onboard the X-ray
satellites . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.5 Point spread function of the XRT . . . . . . . . . . . . . . . . . . . . . . . .
2.6 Vignetting of the XRT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.7 X-ray path through the XRT . . . . . . . . . . . . . . . . . . . . . . . . . .
2.8 Wide-range vignetting function of the Suzaku XRT . . . . . . . . . . . . . .
2.9 Schematic view of a MOS structure . . . . . . . . . . . . . . . . . . . . . . .
2.10 Schematic view of charge transfer in a CCD . . . . . . . . . . . . . . . . . .
2.11 Picture and cross section of XIS . . . . . . . . . . . . . . . . . . . . . . . . .
2.12 Picture of the CCD installed in the base . . . . . . . . . . . . . . . . . . . .
2.13 Schematic view of the XIS CCD . . . . . . . . . . . . . . . . . . . . . . . .
2.14 Schematic view of the CI technique . . . . . . . . . . . . . . . . . . . . . . .
2.15 Schematic view of measuring procedure of CTI with the CFCI technique . .
2.16 Pulse height of the Mn I Kα line from the calibration source. . . . . . . . .
2.17 Time histories of the Mn I Kα line energy. . . . . . . . . . . . . . . . . . . .
2.18 Time histories of the OXII Kα line energy . . . . . . . . . . . . . . . . . . .
2.19 Schematic response function of the XIS . . . . . . . . . . . . . . . . . . . .
2.20 Time histories of the Mn I Kα line width . . . . . . . . . . . . . . . . . . .
2.21 Time history of the OVII Kα line width . . . . . . . . . . . . . . . . . . . .
2.22 Quantum efficiency of the XIS . . . . . . . . . . . . . . . . . . . . . . . . . .
2.23 Spectra of the non-X-ray background . . . . . . . . . . . . . . . . . . . . . .
2.24 Dependence of the NXB on the COR and ACTY . . . . . . . . . . . . . . .
iii

19
20
20
21
22
23
24
25
25
26
28
29
30
31
31
32
33
33
34
35
36

iv

LIST OF FIGURES

2.25 XIS background counting rate . . . . . . . . . . . . . . . . . . . . . . . . . . 36
3.1
3.2
3.3
3.4
3.5
3.6
3.7
3.8

Sgr A region in the 6.4 keV image . . . . . . . . . . . . . . . . . . . . . . . 37
X-ray image in the 6.4 keV neutral iron line at the east region of Sgr A* . . 40
X-ray spectra in the Source and the Reference regions . . . . . . . . . . . . 41
Spectral fitting with a model consisting of plasma and neutral components. 42
Line intensity ratios of He- and H-like ions. . . . . . . . . . . . . . . . . . . 45
Relationship between the electron spectral index and the X-ray photon index 48
Equivalent widths of Kα lines of various neutral atoms. . . . . . . . . . . . 49
Equivalent width of various elements expected in the XRN and LECRe models 50

4.1
4.2
4.3
4.4
4.5

6.4 keV image in the GC region . . . . . . . . . . . . . . . . . . . . . . . . . 53
Full spectrum of the Sgr B1 region . . . . . . . . . . . . . . . . . . . . . . . 55
X-ray images in the Sgr B1 region. . . . . . . . . . . . . . . . . . . . . . . . 57
The 6.7 keV band image between the Sgr A and Sgr B2 regions . . . . . . . 58
Distributions of the 6.7 keV line flux and the photon flux ratio between the
6.7 and 6.9 keV lines . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59
4.6 Background-subtracted spectrum of M 0.51−0.10. . . . . . . . . . . . . . . . 60
4.7 Distribution of the molecular clouds and H II regions in the GC region . . . 61
4.8 Full spectrum of the Sgr C region . . . . . . . . . . . . . . . . . . . . . . . . 66
4.9 X-ray images of the Sgr C region . . . . . . . . . . . . . . . . . . . . . . . . 67
4.10 NXB-subtracted spectra of M 359.43−0.07, M 359.47−0.15, and the outerregion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 68
4.11 Correlation between the equivalent widths of the Fe Kα and Fe XXV Kα lines 71
4.12 Background-subtracted spectra of M 359.43−0.07 and M 359.47−0.15 . . . . 72
4.13 X-ray and radio features in the Sgr C region . . . . . . . . . . . . . . . . . . 74
4.14 6.4 keV image toward the Sgr C region with radio contours of the CS emission line . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 75
5.1
5.2
5.3
5.4
5.5

Sgr B region in the 6.4 keV image . . . . . . . . . . . . . . . . . . . . . . . 77
X-ray images in Fe I Kα and 8-10 keV near Sgr B2 and M 0.74−0.09. . . . . 79
Projected profile of the 6.4 keV line intensity in the 6.4 keV clouds . . . . . 80
X-ray spectra of molecular clouds, Sgr B2 and M 0.74−0.09 on 2005 and 2009. 81
Schematic face-on view of movement of ”X-ray echoes” in the molecular
clouds . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83

6.1
6.2

Required luminosity of the irradiating source . . . . . . . . . . . . . . . . . 88
Distribution of the 6.4 keV Clouds . . . . . . . . . . . . . . . . . . . . . . . 88

List of Tables
3.1
3.2
3.3

Observation Data List. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 39
Fitting results of P lasma. . . . . . . . . . . . . . . . . . . . . . . . . . . . . 46
Fitting Results of N eutral and those of calculated values.∗ . . . . . . . . . 47

4.1
4.2
4.3
4.4
4.5

Result of spectral fitting in Sgr B1 . . . . . . . . . . . . . . . . . . . . . . .
Comparison between required luminosities and observed luminosities . . . .
Center energies of the Fe, Ni line from the GCDX . . . . . . . . . . . . . . .
Best-fit parameters for the NXB-subtracted spectra . . . . . . . . . . . . . .
The best-fit parameters of the spectral fittings to the background (GCDX)
subtracted spectra of M 359.43−0.07 and M 359.47−0.15. . . . . . . . . . .

61
62
68
69
73

5.1

Best-fit parameters.∗ . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 82

6.1

Parameters of 6.4 keV clouds . . . . . . . . . . . . . . . . . . . . . . . . . . 86v

Chapter 1

Introduction
1.1

Overview of the central region in Milky Way Galaxy

The central region of Milky Way Galaxy (Galactic center, GC) is the closest among
those of all galaxies in the universe. A super-massive black hole (SMBH) Sagittarius A*(hereafter, Sgr A*) has been found, which sits at a dynamical center of the Galactic gravity(e.g., Ghez et al. 2008). Many molecular gases, stars, and supernova remnants concentratein the surroundings of Sgr A*. Interactions between these objects induce star-formations.
Many molecular clouds, stellar clusters, and H II regions have been frequently producedin the GC. Strong magnetic field (several tens of µG) has been observed and the energy
density of cosmic-rays is also significantly higher than that in the vicinity of the Earth.
The GC is an excellent laboratory to research high energy phenomena in astrophysics.
The GC are difficult to be observed in the optical and ultraviolet wavebands becauseof large extinction and absorption due to inter-stellar dust and gas (AV ∼ 30 mag) (e.g.
Rieke et al. 1989). The 30 mag of visual extinction corresponds to a column density
NH ∼ 6 × 1022 cm−2 (Predehl & Schmitt 1995), and hence the obscuring medium becomes
partially transparent to X-rays from the GC region at energies > 2 keV. The GC can bealso observed in radio and infrared band with long wavelengths.

1.2Super-Massive Black Hole, Sagittarius A*

Nuclei of most galaxies are now believed to harbour super-massive black holes with
106 –109 M¯ (solar mass). The motions of stars in the central few pc of Milky Way Galaxy
indicate the presence of a super-massive black hole with a mass of ∼ 4 × 106 M¯ (Genzel
et al. 2000; Ghez et al. 2000; Ghez et al. 2008). The black hole coincides with a compact
non-thermal radio source, Sgr A*, and the radio emission is thought to be powered by
the gravitational potential energy released by matter as it accretes onto the super-massive
black hole. However, no counterpart had been found in other wavelength. Sgr A* wasmuch fainter especially in X-rays, which has cast some doubt on adoption of the accretion
model.
A cluster of luminous young stars lies in the immediate vicinity of Sgr A*. Many ofthem are windy emission-line stars that presumably supply the matter which the black1

2

CHAPTER 1. INTRODUCTION

Figure 1.1: Left panel: Schematic diagram of the principle constituents of the Sgr A radio
complex (Baganoff et al. 2003). The outer ellipse depicts the location of the radio shellof Sgr A East, which is a filled-center structure in X-rays. The circumnuclear disk may
affect the morphology of the super-massive black hole, Sgr A*, lies at the center of a cavity
surrounded by the circumnuclear disk. The ionized gas features of Sgr A West and thecentral cluster of luminous, hot stars also lie within the cavity. Right panel: Chandra Xray image in 0.5–7 keV with logarithmic scaling of the central 10 × 10 of the GC (Baganoff
et al. 2003). The black cross marks the radio position of Sgr A*. The cross lies superposedon the X-ray source that we associate with Sgr A* based on the extremely close positional
coincidence. Other sources are a counter part of infra-red stars or new X-ray binaries.
hole is presently accreting (e.g., Krabbe et al. 1995; Paumard et al. 2001). These stars
and their colliding winds are themselves potential X-ray sources (Ozernoy et al. 1997).
This cluster excites a parsec-scale H II region, Sgr A West, that is well studied at radio
and infrared wavelengths. The surrounding of Sgr A West is the dense, predominantly
neutral circumnuclear disk (Morris & Serabyn 1996 and references therein). Beyond that,on a scale of ∼10 pc, a supernova remnant, Sgr A East, with a nonthermal radio shellsurrounds the circumnuclear disk in projection.

1.2.1

X-ray Observations of Sgr A*

X-ray observations were carried out with various instruments. The observations with
an X-ray telescope were started by the Einstein satellite. 1E 1742.5-2859, the strongest
among detected sources was centered only 2000 from the position of Sgr A* (Watson etal. 1981). Assuming an absorbed thermal bremsstrahlung model with kT =5 keV and
NH = 6 × 1022 cm−2 , Watson et al. (1981) estimated the absorption-corrected 0.5–4.5 keVluminosity of this source to be 9.6 × 1034 erg s−1 . Hard X-ray observations were madewith 30 –50 using Spacelab-2/XRT (Skinner et al. 1987), Spartan-1 (Kawai et al. 1988),
and Granat/ART-P (Sunyaev et al. 1993; Pavlinsky et al. 1994). These observations
suggested the presence of a long-term variable source near the position of Sgr A* with an1.2. SMBH, SGR A*

3

average 4–20 keV luminosity of ∼ 1036 erg s−1 .
ROSAT observed the Galactic center and detected 14 sources within the central 300 ×300
region. With the relatively high spatial resolution of 1000 –2000 , it resolved 1E 1742.5-2859
into three sources, of which RX J1745.6-2900 was coincident within 1000 with the radio
position of Sgr A* (Predehl & Trümper 1994; Predehl & Zinnecker 1996). Predehl &Trümper (1994) adopted a thermal bremsstrahlung model with kT = 5 keV, but with
NH = 1.5 × 1023 cm−2 to obtain broadband agreement with the hard X-ray data describedabove, and derived an unabsorbed 0.8–2.5 keV luminosity of ∼ 1036 erg s−1 for the source.
The first X-ray image with an X-ray CCD of the GC was made in 1993 with ASCA
(Koyama et al. 1996). The angular resolution of the X-ray telescope onboard ASCA was∼10 . ASCA detected diffuse thermal emission (kT ∼ 10 keV) covering the central square
degree of the GC (Galactic center diffuse X-rays, GCDX: see section 1.3). A 20 ×30 ellipticalregion filling a shell of a supernova remnant, Sgr A East showed bright diffuse emission.
After correction for a measured absorption of NH ∼ 7×1022 cm−2 , the unabsorbed 2–10 keV
luminosity of this gas was found to be ∼ 1036 erg s−1 . No subtraction was performed forthe spatially variable local background, and consequently ASCA could only place an upper
limit of ∼ 1036 erg s−1 on the X-ray luminosity of Sgr A*.
Koyama et al. (1996) found a hard X-ray source, AX J1745.6−2901 located 1.30 awayfrom Sgr A*. During their second observation made in 1994, Maeda et al. (1996) discovered an X-ray burst and eclipses with a period of 8.4 hr from the hard source, establishingthat it was an eclipsing low-mass X-ray binary (LMXB: see also Hyodo et al. 2009). Only
one Ariel V cataloged transient source, A 1742-289 (Eyles et al. 1975), which appeared in
1975, positionally coincides within the error region. However, Kennea & Skinner (1996)
reanalyzed the Ariel V data taken in 1975 and found no eclipses from A 1742-289. Hence,
the hard source was identified as a newly discovered LMXB and given the name AX
J1745.6-2901 (Maeda et al. 1996). Maeda et al. (1996) reported that the absorbed flux
from this source varied from 1 × 10−11 to 4 × 10−11 erg cm−2 s−1 , which was similar tothe variations reported previously by the lower resolution hard X-ray instruments. Hence,
the hard X-ray fluxes attributed to Sgr A* may have been contaminated significantly by
AX J1745.6-2901 and A 1742-289 (see also Beckert et al. 1996).
A BeppoSAX observation, with on-axis angular resolution of ∼1.30 and an energy range
similar to ASCA, was performed in 1997 (Sidoli et al. 1999b). BeppoSAX detected the
diffuse emission near Sgr A*, measured the absorption column to be NH ∼ 8 × 1022 cm−2 ,and set a tighter upper limit on the 2–10 keV luminosity of Sgr A* of ∼ 1035 erg s−1 .
Recently, the Chandra satellite with the highest resolution 0.500 found a compact Xray source, GXOGC J174540.0−290027, coincides with Sgr A* within 0.300 (Baganoff et al.
2001; 2003). The spectrum is represented with an absorbed power-law with photon index
Γ ∼ 2.7 and column density NH ∼ 1023 cm−2 or with an absorbed optically thin thermal
plasma with kT ∼ 2 keV and NH ∼ 1023 cm−2 . The detected X-ray flux in 2–10 keV of
10−13 erg s−1 cm−2 is too faint to be detect with the previous X-ray instruments. The
33 erg s−1 .
absorption-corrected luminosity in 2–10 keV was estimated to be 2.2+0.4−0.3 × 10
Baganoff et al. (2001) also discovered that Sgr A* flared up in 2000. The luminosity in theflare state was 1 ± 0.1 × 1035 erg s−1 , about fifty times higher than that in the quiescent

4

CHAPTER 1. INTRODUCTION

Figure 1.2: X-ray light curve taken from Sgr A* (Baganoff et al. 2001). (a) Light curvein hard band (4.5–8 keV). (b) Same as (a) but in soft band (2–4.5 keV) counts. The
transverse and vertical axes show the time offset from the start of the observation andcount rates normalized to 600 s. During the quiescent intervals at the beginning and endof this observation, the mean count rate in the 2–8 keV band is 6.4 ± 0.6 × 10−3 counts s−1 .
The detected count rate at the peak of the flare is 0.16 ± 0.01 counts s−1 , which is larger
than that in the quiescent level by more than one order of magnitude.
state. The photon index of the spectra seemed to vary from ∼ 2.7 in quiescence to ∼ 1.0
in flaring. The luminosity is 108 –1010 times fainter than the Eddington luminosity of asource with the mass of 4×106 M¯ . Baganoff et al. (2003) discussed the radiation model of
Sgr A*. In the model, low energy photons (λ ∼ sub-mm) are generated by synchrotron ofrelativistic electrons, and then X-rays are produced via up-scattering of the sub-millimeterphotons off the relativistic electrons (synchrotron self-Compton model).

1.2.2

Gamma-rays from the direction near Sgr A*

Gamma-rays were also detected from the position coincide with Sgr A* by some instruments. The EGRET team found that the 100 MeV to 10 GeV gamma-ray emission
comes from the GC region at (l, b) = (0.11◦ , −0.04◦ ) with the errors 0.13◦ in the 95%
confidence level (Mayer-Hasselwander et al. 1998). An Imaging Atmospheric Cerenkov
Telescope (IACT), Cangaroo-II detected very high energy (VHE) TeV gamma-rays froma point-like source at a position consistent with that of the Sgr A* (Tsuchiya et al. 2004).
The IACT, H.E.S.S. also detected VHE gamma-rays in 0.1–10 TeV and decided the peak
position of (l, b) = (0.0◦ ± 0.1◦ , 0.0◦ ± 0.1◦ ) (Aharonian et al. 2006). Although highenergy gamma-rays from Sgr A* suggested that super-massive black hole is a site wherecosmic-rays are accelerated, the less position uncertainty than that in the radio and X-ray
results has remained some origin for the gamma-ray emission detected from the GC.
Some galactic nuclei have high luminosity comparable to the Eddington limit (Active

1.3. DIFFUSE X-RAY EMISSION FROM THE GC REGION

5

galactic nuclei: AGN). The reason why the nucleus of Milky Way Galaxy is much fainterthan the AGN is under intense debate.

1.3

Diffuse X-Ray Emission from the GC Region

There are a number of types of astrophysical objects which emit X-rays, from galaxy
clusters, through black holes in active galactic nuclei to galactic objects such as supernova
remnants, stars, and binary systems containing white dwarfs, neutron stars or black holes(X-ray binaries). A combination of many unresolved X-ray sources is thought to produce
the observed X-ray background. These are divided into discrete (point-like) and extended
objects.
Early X-ray observations discovered the presence of several discrete objects concentrating on the Galactic plane, especially in the Galactic center region (e.g., Watson et al.
1981; Skinner et al. 1987; Kawai et al. 1988). The objects are bright X-ray binaries of
black holes or neutron stars with high luminosity of LX = 1036 –1038 erg s−1 .
On the other hand, diffuse X-ray emission which were not able to be resolved into
discrete sources was found from the Galactic plane by various instruments, rockets, balloons, and satellites (Cooke et al. 1969; Hudson et al. 1971; Bleach et al. 1972; Worrall
et al. 1982; Iwan et al. 1982). This emission has been called as the Galactic ridge X-ray
emission (GRXE). The hard spectrum of the GRXE was explained by an optically thinthermal plasma with a temperature kT ∼ 10 keV (Iwan et al. 1982) or by a power-law
model with a photon index Γ = 2–3 (Bleach et al. 1972).
Many possible origins have been proposed by many researches so far. Worrall et al.
(1982) proposed that the seemingly extended emission of the GRXE comes from integration of discrete sources, accreting white dwarfs (cataclysmic variables) and binary stars
with strong coronal activity. On the other hand, Koyama et al. (1989) suggested that
the GRXE is truly diffuse, and originates from thermal hot plasma gas with temperature
kT =1–10 keV. Ginzburg & Syrovatskii (1964) and Blumenthal & Gould (1970) attempted
to explain by another scenario that the nonthermal X-ray spectrum is generated by inverse Compton scattering of the cosmic microwave background (T ∼ 3 K) on cosmic-ray
electrons of a few GeV. Protheroe et al. (1980) interpreted the X-ray emission by thesynchrotron radiation of the cosmic-ray electrons of 1014 eV interacting with a interstellar
magnetic field of several µG. However, no critical information from the observations had
been found and the origin had been under discussion for a long time. It is not understoodwhether the GRXE is integration of discrete stars or truly diffuse.

1.3.1

Discovery of Fe Lines from the GC Region

The Tenma and Ginga satellites made a breakthrough discovery of a 6.7 keV emission
line of Fe XXV Kα in the GRXE spectrum (see left panel in figure 1.3). Moreover, thedistribution of the 6.7 keV emission line intensity has two components widely extending
over the Galactic plane and concentrating within ∼ 1◦ of the Galactic center (GC) at
l = 0◦ (Koyama et al. 1986a; Koyama et al. 1990; Yamauchi et al. 1990). This diffuse

6

CHAPTER 1. INTRODUCTION

Figure 1.3: Left: X-ray spectra of the GCDX at (l, b) = (1.2◦ , 0.5◦ ) and (358.0◦ , −0.6◦ )near the Galactic center taken by the Ginga satellite (Yamauchi et al. 1990). A pronounced Fe line at 6.7 keV in the spectra is modeled with a Gaussian. Right: Galactic
longitude distribution of the Fe line intensity (Yamauchi et al. 1990). Each panel shows
the distribution at a different latitude b =1.7◦ , 0.5◦ , −0.6◦ , −1.1◦ , and −2.1◦ . The arrows
indicates the intensity peak of each profile.
X-ray emission is called as the GCDX. The existence of the highly ionized Fe suggested
that the diffuse X-ray emission possibly comes from the optically thin thermal plasma
with the temperature of kT ∼ 10 keV.
Detailed image and spectrum in the hard X-ray band up to 10 keV were obtained by
the X-ray telescopes aboard the ASCA satellite (Koyama et al. 1996). Superior energy
resolution of the X-ray CCD cameras (SIS) demonstrated that the Fe emission line is due
to complex of three different ionization states, and consists of emission lines at 6.4, 6.7,
and 7.0 keV (figure 1.4). The 6.7 and 7.0 keV lines are from highly-ionized ions of He-like(Fe XXV Kα) and H-like (Fe XXVI Kα) Fe, respectively. On the other hand, the energy
6.4 keV of the other line indicated that Fe is a neutral or lowly-ionized state. Thus, the
GCDX consists of at least two components. One associates with 6.7 and 7.0 keV Fe lines
and the other is a neutral component with the 6.4 keV Fe line.

1.3. DIFFUSE X-RAY EMISSION FROM THE GC REGION

7

Figure 1.4: X-ray spectrum of the GCDX obtained by ASCA. The Fe lines of neutral(6.4 keV), He-like (6.7 keV), and H-like (7.0 keV) are resolved. There are many emission
lines of the other heavy elements, Si, S, Ar, and Ca in highly-ionized states.
We review the previous researches of the 6.7 and 7.0 keV lines in the next section,
while those of the 6.4 keV line is mentioned in section 1.4.

1.3.2

Origin of Fe XXV and Fe XXVI

Two plausible origins for the highly-ionized Fe lines (Fe XXV and Fe XXVI) havebeen proposed. One is collisional excitation in the thin thermal plasma with the hightemperature as above mentioned (Koyama et al. 1989; Yamauchi et al. 1990), and theother is charge exchange (CX) recombination due to interaction between bare Fe in cosmicrays and a hydrogen atom in an interstellar medium (Tanaka et al. 2000).
A broadened Fe line width of ∼ 70 eV was found in the X-ray spectrum obtained by
ASCA (Tanaka et al. 2000). This suggested that the velocity of the Fe ion was 3000–5000
km s−1 , which was greatly exceeded the sound speed ∼ 1000 km s−1 in the ∼10 keV
plasma. Although intensities of the two Fe lines are expected to be the same in the CX
scenario, the observed intensity ratio between Fe XXV Kα and Fe XXVI Kα is ∼ 1 : 2
across all the GC region (Koyama et al. 1996). Moreover, no excesses were seen even atregions with dense molecular gases, i.e. molecular clouds (Koyama et al. 1996). These
facts deprecate the CX scenario and favored the hot plasma one.
The Fe XXV Kα line has fine structures consisting of different transitions of resonance,intercombination, and forbidden. The center energies of the resonance, intercombination,and forbidden lines are 6.70, 6.67–6.68, and 6.64 keV, respectively. The forbidden line isstronger than the resonance line in the CX process, whereas the resonance line is stronger
in the hot plasma (Koyama et al. 2007c, and reference therein). Although the X-ray
CCDs, which have the highest energy resolution among the instruments launched in orbit,
cannot resolve the fine structures, the two origins are distinguishable by the center energyof the Fe XXV Kα line. The hot plasma should provides the center energy 6.680–6.685 keV
for the Fe XXV Kα (Koyama et al. 2007c). On the other hand, Wargelin et al. (2005)
measured by a laboratory experiment that the Fe XXV Kα line in the CX process has

8

CHAPTER 1. INTRODUCTION

Figure 1.5: Left: X-ray spectrum of the GCDX obtained by Suzaku (Koyama et al. 2007c).
The model is a thin thermal plasma of kT = 6.5 keV plus a power-law of the photon index
Γ = 1.4 with 3 Gaussian lines at 6.40, 6.68, and 6.97 keV and an iron absorption edge at
7.11 keV. Right: line fluxes of Fe XXV Kα are given by the squares, while the integrated
point source fluxes in the 4.7–8 keV band are plotted by the crosses (Koyama et al. 2007c).
The horizontal axis is the Galactic longitude, but the vertical axis is a logarithmic scale.
lower energy 6.666 ± 0.005 keV.
Suzaku, which is described in chapter 2, obtained the most excellent X-ray spectrumof the GCDX in the past, and decided the center energy 6.680 ± 0.001 keV for the Fe
XXV Kα line successfully (left panel in figure 1.5). The value is inconsistent with that in
the CX origin proposed by Tanaka et al. (2000), but coincides with that expected in the
hot plasma.
Suzaku also gave a severe constraint for the Fe XXV Kα line broadening of ≤ 30 eV.
The intrinsic line broadening of Fe XXVI Kα is nearly zero. On the assumption of the
CX process, the velocity of the bare Fe is driven to be nearly zero. In such a case,emission lines at the Fe XXV Rydberg series of large principle quantum numbers (n ≤ 10)
are predicted, but no enhancement at the energy was observed in the GCDX spectrum.
These observational facts indicate that the highly-ionized Fe lines come from the hightemperature hot plasma with kT ∼ 6.5 keV covering the ∼ 1◦ centered region of the GC,
hereafter the GC hot plasma (Koyama et al. 2007c).

1.3.3

GC Hot Plasma

Whether the GC hot plasma is truly diffuse or the integration of faint point sources
have been under debate. White dwarf binaries (cataclysmic variables: CV) and active
binaries (e.g. RS CVn) have thin thermal spectra with kT = several keV, which is similar
to that of the GC hot plasma. However, characteristic properties of X-ray spectra anddistributions of the binary sources measured at the vicinity of the solar system does not
explain the GC hot plasma (Yamauchi et al. 1990).
The 100 central region of the GC was observed with a excellent spacial resolution ∼ 0.500
of Chandra (Muno et al. 2003). More than two thousands point sources are resolved. The

1.3. DIFFUSE X-RAY EMISSION FROM THE GC REGION

9lowest flux is 10−17 erg s−1 cm−2 , which corresponds to that of a faint CV. The summed
flux of the resolved point sources is only 10% of the flux of the total diffuse emission.
However, Revnivtsev et al. (2007) re-analyzed the same Chandra data correcting themeasured source counts at low fluxes for bias associated with Poisson noise and resolved
more number of point sources, whose integrated flux contributes 40% of the total emissionof the GC hot plasma. Revnivtsev et al. (2007) proposed that the GC hot plasma iscompletely divided into discrete sources including fainter CVs and active binaries thanthe detected ones.
Recently, Suzaku obtained a detailed intensity map in Fe XXV Kα around the GC.
Koyama et al. (2007c) found a difference between distributions of the Fe XXV Kα flux
and the total flux of the resolved discrete sources at l = −0.4◦ –+0.2◦ (see right panel infigure 1.5), and hence they suspected that the truly diffuse hot plasma should accountat least some part of the GC hot plasma. Moreover, Koyama et al. (2009) separated
the Suzaku data of the GCDX into 4.50 × 4.50 sections to obtained positional variation of
the Fe lines and continuum emission in the 5–10 keV band. The sum of the 6.4 keV and6.7 keV line fluxes with a ratio of 1 : 2 showed good proportionality to the continuum flux,
and hence the continuum flux of the GCDX was phenomenologically decomposed into the
6.4 keV- and 6.7 keV-associated continua with a flux ratio of 1 : 2. The results indicatedthat the contribution of integrated flux of point sources to the GCDX was ∼ 1/6 (Koyama
et al. 2009).
Each scenario has a problem. If the plasma is the integration of discrete sources, theFe XXV Kα equivalent width ∼ 600 eV (e.g. Muno et al. 2003) of the GCDX cannot be
explained only by known X-ray sources, such as CVs with a rather small value ∼ 200 eV(Ezuka & Ishida 1999). This needs larger metal abundances than the solar value in the
GC region or faint X-ray sources that have been invisible so far.
If the GC hot plasma is truly diffuse, it has an enormous amount of energy 1053 –
1054 erg, which corresponds to 100–1000 supernova explosions. In addition, the kT =6.5 keV plasma is not bounded by gravitational potential of Milky Way Galaxy and runs
away within about 105 years. A huge energy injection 1041 erg s−1 is needed to sustain
the diffuse hot plasma. Although the X-ray and radio observations have found about
10 supernova explosions for 104 years, the required energy injection is not explained.
The results obtained by the past observations required that many undetected supernova
remnants are buried in a heap of the hot plasma, or the super-massive black hole Sgr A*
heated up the surrounding gases.

10

CHAPTER 1. INTRODUCTION

Figure 1.6: X-ray map in the 6.4 keV line obtained by ASCA (Koyama et al. 1996). There
are two bright spots in the Sgr B2 and Radio Arc regions at the upper left and the centerof this figure.

1.4

Fe I Kα Emission from Giant Molecular Clouds

ASCA found the Fe line at 6.4 keV in the GCDX spectrum (section 1.3). The distribution of the 6.4 keV emission is not uniform like the 6.7 keV emission, i.e. the GC hot
plasma, but showed clumpy structures corresponding to giant molecular clouds in the Sgr
B2 and the Radio Arc regions (figure 1.6: Koyama et al. 1996). Hence, the 6.4 keV line
should be produced from large amount of neutral Fe atoms in the clouds. Hereafter, wecall the molecular clouds emitting the 6.4 keV (Fe I Kα) line as ”6.4 keV clouds”.
In order to emit the 6.4 keV fluorescent X-ray, the neutral Fe atom should be ionizedby an external particle with higher energy than the Fe I-K edge (7.1 keV). However, theorigin of the 6.4 keV line from the clouds is still one of the open questions among highenergy phenomena in the GC region.

1.4.1

Possible Origins for the Fe I Kα Emission

Koyama et al. (1996) proposed that the 6.4 keV emission from the Sgr B2 cloud isdue to the K-shell ionization of neutral iron atoms by external X-rays (an X-ray reflection
nebula; XRN). The characteristic features of the large equivalent width (1–2 keV) of
the Fe I Kα line and the deep absorption edge at the Fe I-K edge corresponding to NH ∼1024 cm−2 were successfully fitted with the XRN model (Murakami et al. 2000). The peak
position of the emitting region was found to be shifted from the dense core of the molecular
cloud toward the GC by 10 –20 (Murakami et al. 2000). Murakami et al. (2001b) also
discovered a 6.4 keV cloud in the Sgr C region. Based on these results of the 6.4 keV clouds,
the strong X-ray source with high luminosity ≥ 1038−39 erg s−1 is necessary (Koyama et1.4. FE I Kα EMISSION FROM GIANT MOLECULAR CLOUDS

11

al. 1996; Murakami et al. 2001a, 2001b). The super-massive black hole, Sgr A* has beensuspected to be the irradiating source (e.g. Sunyaev et al. 1993; Koyama et al. 1996).
However, the present X-ray luminosity of Sgr A* is 1033 –1034 erg s−1 (Baganoff et al.
2001, 2003). Therefore, Sgr A* in a few hundred years ago would have been 104−6 times
brighter than now (e.g. Koyama et al. 1996).
On the other hand, Yusef-Zadeh et al. (2007) proposed that the origin of the 6.4 keVemission would be low energy cosmic-ray electrons (LECRe: Ee =10–100 keV), because
they found that the X-rays correlated with synchrotron-emitting radio filaments. Energy
densities 20–103 eV cm−3 of the electrons were estimated for the 6.4 keV emission in the
Sgr B1, B2 and C regions. Yusef-Zadeh et al. (2007) also explained the cosmic-ray heatingof molecular gas and TeV gamma-ray emission in the GC region. In particular, inverse
Compton scattering of the sub-millimeter radiation from dust by the relativistic electrons
may contribute substantially to the TeV gamma-ray emission.
In the XRN scenario, hard continuum X-rays should be generated due to the Thomson
scattering of the incident external X-rays. The LECRe scenario also predicts a hard
continuum of the bremsstrahlung by the relativistic electrons. In fact, the INTEGRAL
satellite detected hard X-rays upto 100 keV from the Sgr B2 cloud. The hard X-ray
spectrum agreed with the XRN model when the photon index of the incident X-rays
and the optical depth of the cloud are Γ = 1.8 ± 0.2 and τ = 0.4 ± 0.1, respectively(Revnivtsev et al. 2004). The X-ray luminosity in the hard X-ray band of the Sgr B2
cloud is LX ∼ 1035 erg s−1 (Revnivtsev et al. 2004). On the other hand, since only ∼ 10−5of the energy of the relativistic electrons can go into hard X-ray radiation at ∼ 50 keV,
the total energy release of cosmic-ray electrons should be > 1040 erg s−1 . The value iscomparable to the bolometric luminosity of Sgr B2 mostly observed in infrared (Gordon
et al. 1993), and hence the LECRe scenario seems to be unlikely (Revnivtsev et al. 2004).

1.4.2

Recent Results—time variability of X-ray emissions—Figure 1.7: Left: Time variability of the 6.4 keV flux from the Sgr B2 clouds (M 0.66−0.02,
Inui et al. 2009). The vertical axis means the surface brightness in the 6.4 keV linein the unit of photons s−1 cm−2 arcmin−2 . Right: Same as the left panel, but from
G 0.570−0.018.

12

CHAPTER 1. INTRODUCTION

Thanks to the high sensitivity for extended faint sources of Suzaku, many new 6.4 keV
clouds have been discovered in a variety of regions (Koyama et al. 2007b; Nobukawa et al.
2008; Fukuoka et al. 2009; Nakajima et al. 2009; Nakashima et al. 2010), and their X-ray
spectra, in particular the Fe line, were well studied. Most of the 6.4 keV clouds have thestrong 6.4 keV lines with large equivalent widths 1–2 keV and hard continua absorbed bydense media (photon index Γ ∼ 2, absorption column density NH ≥ 1023 ).
Koyama et al. (2008) and Inui et al. (2009) discovered time variability of the 6.4 keVflux in the Sgr B2 region from the archival data for 10 years obtained with Suzaku,
XMM-Newton, Chandra, and ASCA. The 6.4 keV flux of the most brightest cloud Sgr B2
was ∼ 3 × 10−6 photons s−1 cm−2 arcmin−2 in 1994–2001, while that decreased to ∼ 2 ×
10−6 photons s−1 cm−2 arcmin−2 in 2004–2005 (figure 1.7). The other spot, G 0.570−0.018,
is found to be conspicuous only in the Chandra observation in 2000 (see figure 1.7).
Muno et al.(2007) also discovered the changes in the 4–8 keV fluxes and morphologiesof the 6.4 keV clouds in the Radio Arc region from the Chandra observations in 2002–2005.
The flux change of the 6.4 keV line of the 6.4 keV clouds was detected at the 5σ confidence
level by the combination of the Chandra and Suzaku data (Koyama et al. 2009). Ponti etal. (2010) also studied the 6.4 keV clouds in the Radio Arc region using the XMM-Newton
data in 2001–2009 with total exposure time ∼ 1.2 Ms. Some of the clouds flickered in the
6.4 keV line as is shown in figure 1.8, while some of them showed constant luminosities.

Figure 1.8: Fe I Kα images obtained with the different XMM-Newton observations (Ponti
et al. 2010). A brightening of the regions marked with numbers of 1, 2, 3, and 4 is clear.
An X-ray flux decrease from the Sgr B2 cloud also in the hard X-ray (20–60 keV) band
was detected from the seven years monitoring data taken by INTEGRAL (Terrier et al.
2010). The characteristic time scale of the decay was ∼ 8 years, comparable with the1.5. AIM OF THIS THESIS

13

light-travel time across the cloud core.
Since the variability over several years from a pc-scale object requires light-speed communication, these facts support the XRN scenario: X-rays from external sources irradiate
molecular clouds and produce the fluorescence iron line at 6.4 keV and Thompson scattering continuum emission in the same manner. However, low-energy cosmic-ray electrons
are not formally ruled out since the lifetime of ∼100 keV electrons in the dense clouds issmaller than the measured decay time of several years, so that if the cosmic-ray injection
had ceased with a similar or smaller timescale, we might obtain a similar flux decrease.
We have left room for debating the origin of the 6.4 keV line from the GC region.

1.5

Aim of this thesis

In this thesis, we aim to examine the detailed features of the 6.4 keV emission and toreveal the origin.
1. Is the X-ray emission from the 6.4 keV clouds explained only by the X-ray reflection
model, not by the low energy electron model?
2. What is the irradiating source of the 6.4 keV clouds?
3. How long and how bright was the Sgr A* flare in the past? (if it is the irradiating
source.)The detailed analyses and results for discrete sources are written in chapter 3, 4, and
5. Since the two processes should produced different X-ray spectra in principle, we canclearly distinguish the origin by the more detailed X-ray spectra. Although only theneutral Fe line has been detected so far, neutral K-shell lines of the other atoms shouldbe emitted from the clouds. In chapter 3, we state a new constrain for the origin (XRNor LECRe) based on K-shell lines of several heavy elements that are newly discoveredin this thesis. In chapter 4, we examined 6.4 keV clouds found in other regions anddiscussed the origins in detail. Previous observations indicated that the most plausible
source irradiating the 6.4 keV clouds is the central super-massive black hole, Sgr A*.
According to time variabilities of the 6.4 keV and hard X-rays from the two clouds in the
Sgr B2 region, an unique constrain on the irradiating source is made in chapter 5. At theend of this thesis, we draw a conclusion for the 6.4 keV emission in the GC region on thebasis of the the obtained results.
We use the Galactic coordinates in this thesis, where the transverse and vertical axes
are the Galactic longitude and latitude. We also consider the positive and negative directions of the Galactic longitude as ”east” and ”west”, respectively. In the same way, thepositive and negative directions in the Galactic latitude are ”north” and ”south”. Errors
are estimated at the 90% confidence levels in this thesis, unless otherwise stated.

Chapter 2

Instrument:
Suzaku X-ray Satellite

XRT−S
Sun shade
XRT−I(4 units)

EOB
Star
tracker

Side panel(8 panels in total)

Gyro−
scopes
XRS

Base
panel

HXD
XIS
(4 units)

Figure 2.1: Left: Schematic view of the Suzaku satellite in orbit (Mitsuda et al. 2007).
Right: Side view of the internal structure (Mitsuda et al. 2007).

For this work, we used data obtained with the X-ray Imaging Spectrometers (XIS:Koyama et al. 2007a) equipped on the focal plane of the X-Ray telescope (Serlemitsos
et al. 2007) aboard the Suzaku X-ray satellite (Mitsuda et al. 2007). We introduce theoverview and performances of the Suzaku satellite and the onboard instruments in thischapter.
15

16

2.1

CHAPTER 2. SUZAKU X-RAY SATELLITE

Overview of the Suzaku Satellite

The Suzaku satellite, which is the fifth X-ray science satellite in Japan, was launchedby Japan Aerospace Exploration Agency (JAXA) with the M-V Launch Vehicle from the
JAXA’s Uchinoura Space Center (USC) on 2005 July 10. The mission name before thelaunch was Astro-E2. The Suzaku satellite was developed by the Institute of Space andAstronautical Science of JAXA (ISAS/JAXA) in cooperation with the National Aeronautics and Space Administration’s Goddard Space Flight Center (NASA/GSFC) and manyother institutions.
The Suzaku satellite was carried into the ∼570 km altitude orbit and is orbiting the
Earth with the period of ∼96 minutes and the inclination angle relative to the equator of
31◦ . The low altitude makes the significant advantage of the satellite of lower non-X-ray
background (NXB) caused by cosmic-rays than the other X-ray satellite, Chandra and
XMM-Newton, although efficiency of observation becomes relatively small (∼ 43%) since
the Earth occultation. Operations for observation, health check, and data acquisition are
performed by communication with large parabola antennas with the diameters of 20 and
34 m in the USC 4 or 5 times per a day. Total communication time is about 50 minutes
per one day.
The key feature of the Suzaku satellite is high sensitivity wide-band (0.2–600 keV)
X-ray spectroscopy. The Suzaku satellite is equipped on the base plate with the three
instruments of the high resolution X-ray spectrometer (X-Ray Spectrometer; XRS, 0.5–
10 keV: Kelley et al. 2007), the X-ray imaging spectrometer (XIS, 0.2–12 keV), and thenon-imaging collimated hard X-ray detector (Hard X-ray Detector; HXD, 10–600 keV:
Takahashi et al. 2007; Kokubun et al. 2007). The arrangement of the instruments isshown in the right panel of figure 2.1. We concentrate the data taken by the XIS with theenergy band of 0.2–12 keV in this thesis. The details of the XRT and the XIS are shownin the following sections.

2.2

X-Ray Telescope (XRT)In order to collect X-rays and to take X-ray images in the < 10 keV band, five setsof the XRT are mounted at the top of the satellite (Serlemitsos et al. 2007). One unit
(XRT-S) is used for the XRS and the others (XRT-I0 to XRT-I3) are for the XIS (seefigure 2.1).

DesignSince refractive indices of matter to X-rays is less than 1, X-rays cannot be focusedby a refracting telescope. X-rays are full-reflected if the incident grazing angle is shallow
enough (≤ 1◦ ), and thus X-rays are focused using this full-reflection.
An approximate the Wolter type-I optics is used in the XRT. The XRT is the thin-foilnested X-ray telescope consisting of many concentric thin reflecting mirrors of thin aluminum foils coated with gold. The shape of the reflecting foils are not complete paraboloid
but are approximated by cones in the Wolter type I grazing incident optics (figure 2.3).

2.2. X-RAY TELESCOPE

17

Figure 2.2: X-ray telescope onboard the Suzaku satellite (Serlemitsos et al. 2007).

Figure 2.3: Schematic view of an X-ray telescope composed of nested mirror shells with a
Wolter type-I grazing incident optics. This figures are adopted from Mori et al. (2005).

18

CHAPTER 2. SUZAKU X-RAY SATELLITE

X-rays are reflected twice on the surface of the primary and secondary reflectors and arecollected on the focal plane. The XRT is basically the same as the X-ray telescope onboard
the ASCA satellite (Serlemitsos et al. 1995), while the following major improvements were
made.
The incident grazing angle has been reduced to be ≤ 0.6◦ and the diameter of the XRT
has increased from ∼ 350 mm to ∼ 400 mm. As a result, the effective area especially inthe higher energy band than that of Fe-K lines (6–7 keV) has been significantly improved
since X-rays with the higher energy can be collected than by the ASCA/XRT. The focal
length also changed to be from 3.5 m to 4.5 (for XRS) and 4.75 m (for XIS).
The ASCA reflectors were made by coating acrylic lacquer on the aluminum foils
followed by vacuum deposition of a gold layer on the top. The surface of the foils was
not smooth sufficiently, resulting in the image broadening. For the Suzaku XRT, the gold
layer was sputtered onto the glass tube of a highly smooth surface, and was transferred to
the foil reflectors. A half power diameter (HPD) is the diameter of a circle which contains
50% of flux of a point source. The change of the fabricating method made the HPD from∼ 3.60 (ASCA) to ∼ 2.00 (Suzaku).
X-rays from the field of view (FOV) are properly focused when they are reflected by
the two surfaces. If X-rays from the outside of the FOV are reflected only by one surface
or by the backside of the mirrors, an incorrect X-ray image appears on the focal plane(stray light). To avoid the stray light, a pre-collimator is equipped on the top of the XRT.
The entrance of the XRT is covered with a thermal shield in order to isolate the
XRT thermally from the space and to keep the temperature within the specified range of
20 ± 7.5 ◦C. The thermal shield also works to block optical light from the sky and fromthe surface of the Earth illuminated by the Sun.

Performance
The XRT has light-weight of ∼20 kg and large effective area (450 cm2 at 1.5 keV,
250 cm2 at 7 keV) per one unit. Since the four sets of the XRT-I modules point to thesame direction, the total effective area is calculated by summing those of all the XRT.
Figure 2.4 shows the total effective area including the detector efficiency, compared withthose of the Chandra and XMM-Newton satellites. Although the total weight of the XRT
is at least 80 kg, the effective area at 7 keV is comparable to that of the XMM-Newton
mirrors, which has the total weight of ∼ 1300 kg. On the other hand, the spacial resolution(HPD) of ∼ 20 is larger than the other X-ray telescopes onboard the other active satellites,
Chandra (High resolution mirror array: ∼1500 kg) and XMM-Newton by about one orderof magnitude since the design philosophy of the XRT is to achieve the the better spectral
power, especially in the 6–10 keV band including Fe-K lines in the confines of a severe
weight limit due to the launch vehicle.
The HPDs of the XRTs are 10 .8–20 .3. Point spread function (PSF) is an intensity
profile of a point source as a function of radius from the point source. Enclosed energyfunction (EEF) is a integration of PSF and shows the ratio of the flux of the enclosed
circle to the total flux of a point source. PSF and EEF of one of the XRTs are shown infigure 2.5.

2.2. X-RAY TELESCOPE

19

Figure 2.4: Effective areas of the Suzaku XRT, compared with the other telescopes onboard
the Chandra and XMM-Newton satellites (Serlemitsos et al. 2007).
Tightly-nested structures increases the possibility of reflection other than the normal
double reflection within the telescope structure. Because of these unexpected incidences,
X-ray of bright sources outside the field of view reaches the focal plane (stray light). Thishas been a large problem for nested X-ray telescopes (e.g. ASCA/XRT).
The effective area for off-axis incident X-rays is smaller than that for on-axis X-rays
because off-axis X-rays are partially obscured by neighbor mirror shells. This ”vignetting”
effect is prominent at high energy. A maximum incident angle for full-reflection is smaller
for higher energy X-rays than for lower ones. High energy X-rays are reflected only at
inner mirror shells, where the cone angles is small and the mirror shells are tightly packed(see figure 2.3), and highly subject to shadowing. The off-axis angle dependencies of theeffective area for the soft (3–6 keV) and the hard (8–10 keV) bands are shown in figure
2.6.
Figure2.8 shows the wide-band vignetting function of the XRT. The stray light (thesecondary reflection) dominants at the off-axis angle of 200 –700 . The pre-collimator successfully reduces the stray light although some portion in the off-axis angle of ∼ 200 and
600 –700 is not completely diminished.
From real data taken by observations of the Crab Nebula, optical axes of the fourXRT-I modules were measured to scatter within ∼ 1.30 . The observation efficiency of the
XRT is more than 97% even at the high energy limit of 8–10 keV.

20

CHAPTER 2. SUZAKU X-RAY SATELLITE

Figure 2.5: Left: Point spread function (PSF) of the XRT of XIS 0 (Serlemitsos et al.
2007). This is obtained with the observation of SS Cyg, which is a bright point-likesource. Right: Enclosed energy function XRT of XIS 0 (Serlemitsos et al. 2007). This isobtained with the integration of the PSF in the left panel.

Figure 2.6: Off-axis angle dependence of the effective area normalized at the optical axes
(Serlemitsos et al. 2007). The black and red points shows those in the soft and hardbands obtained with the Crab nebula observations, respectively (vignetting). The solidlines show the simulated models.

2.2. X-RAY TELESCOPE

21

Figure 2.7: Major reflection paths occurring in the XRT. τ represents the oblique angleof the primary reflectors measured from the optical axis (Serlemitsos et al. 2007). (a)Normal double reflection of the X-rays arriving from the on-axis direction. Incident Xrays are bent by an angle of 4τ in total. (b) Secondary reflection: an X-ray arrives at
the focal plane only if the incident angle θ measured from the optical axis is in the rangeof τ < θ < 2τ . (c) Stray-light path that gives rise to the brightest ghost among various
backside reflections. The pattern occurs then the X-ray incident angle is in the range of2τ < θ < 3τ .

22

CHAPTER 2. SUZAKU X-RAY SATELLITE

Figure 2.8: Wide-range vignetting function of the XRT at the three energies of Al-Kα(1.49 keV), Ti-Kα (4.51 keV), and Cu-Kα (8.04 keV) for off-axis angles up to 800 (Mori etal. 2005). The circles and triangles are the measured stray effective area at 1.49 keV before
and after installation of the pre-collimator, respectively. The continuous and dashed curves
are the effective areas of the secondary reflection with and without the pre-collimator,
respectively, while the dash-dotted curves are those of the normal double reflection. Allof the curves were simulated by the ray-tracing program.

2.3. X-RAY IMAGING SPECTROMETER

2.3

X-Ray Imaging Spectrometer (XIS)

2.3.1

Introduction of X-ray Charge Coupled Device Cameras

23

The XIS consists of MOS-type charge coupled devices (CCDs). An X-ray CCD isa silicon semiconductor detector with two-dimensional imaging capability. Each pixel
consists of Metal (poly-silicon) and Oxide (SiO2 ) as an insulator and p-type semiconductor(silicon). This is called a MOS structure (figure 2.9).

Metal
Oxide
Semiconductor

Figure 2.9: Schematic view of a MOS structure.
The principle of an X-ray CCD is briefly explained in the following. In a p-typesemiconductor, holes are the major carriers. A positive voltage applied to the metalelectrodes repels holes in the semiconductor and leaves behind a depleted layer where anelectric field exists. When an X-ray photon enters into a pixel, photo-electrons are created
and then electron-hole pairs are generated. The number of the generated electrons is
proportional to the energy of the incident photon. The average energy required to form
an electron-hole pair is 3.65 eV at −90◦ C. The holes and electrons are separated by theelectronic field in the depletion layer before they recombine. The electrons are collected
into the potential wells beneath the electrodes as signal charges. After the exposure
period, the signal charges are transferred to the readout node by clocking the voltages ofthe electrodes.
The CCDs of the XIS employ the three-phase transfer type electrode structure. As
shown in left panel of figure 2.10, in this type, three electrodes are deposited on each pixel.
By applying the three-phase clock voltage to the electrodes, the potentials beneath theelectrodes change periodically and the signal charges are transferred.
The CCD of the XIS is classified as a frame transfer (FT) CCD in the scheme of thecharge transfer and readout. A FT CCD consists of an imaging area and a frame-store
region shown in right panel of figure 2.10. The imaging area is exposed but the framestored region is shielded from X-rays. After the exposure period, the signal charges aretransferred in a short time from the imaging area to the frame-store region. The signal
charges in the frame-store region are transferred vertically to the horizontal shift registerby each row. The charges in the horizontal shift register are transferred to the readoutnode by each pixel. This process is repeated until the charges in all the pixels of theframe-store are read out. When the charges are transferred to the frame-store area, new

24

CHAPTER 2. SUZAKU X-RAY SATELLITE

1 pixelClock voltage

1
Electrode

Imaging area
Potential

Frame−store region

Time

Horizontal shift register

Electrons

Readout
node

2

Figure 2.10: Schematic view of three phase transfer (left panel) and frame transfer (rightpanel) in a CCD (Takagi et al. 2006)
exposure starts in the imaging area. The exposures and readouts are made continuously.
In this way, amounts of charges in each pixel are measured and an image and a spectrum
are taken with an X-ray CCD camera. An X-ray CCD camera has also time resolution
depending on the exposure cycle.

2.3.2

Overview of XIS

The XIS is an X-ray CCD camera system which can take image, spectral and photometrical data in the 0.2–12 keV band. The XIS has four sensors (XIS 0, 1, 2, and 3). Eachsensor consists of one CCD chip and a camera body. A photograph and the cross sectionof one of the XIS sensors (CCD + camera body) are shown in figure 2.11.
The CCD chip of the XIS is a MOS-type three-phase FT CCD. A photograph andthe schematic view of the CCD is shown in figure 2.12 and 2.13. The CCD chips of
the XIS are fabricated by the Lincoln Laboratory of MIT and basically same as those of
ACIS/Chandra (Burke et al. 2007) except for a charge injection structure. The details of
the charge injection structure are found in the following section 2.3.3.
The CCD has imaging and frame-store areas with 1024×1024 pixels (figure 2.13). The
signals are transferred down in figure 2.13. The pixel sizes of the imaging area and theframe-store region are 24 µm × 24 µm and 21 µm × 13.5 µm, respectively. In the imaging
area, one pixel and the entire field of view correspond to 100 .04 × 100 .04 and 170 .8 × 170 .8,
respectively, combined with the XRT whose focal length is 4.75 m. Note that the spatial
resolution of the XIS for celestial objects is determined by not the pixel size of the XIS
but the imaging capability of the XRT.
Each CCD chip is divided into four segments (from A to D) and each segment has one
readout node. One of the CCD chips (XIS 1) is a back-illuminated type (BI) and the others
are a front-illuminated type (FI). The front side of a CCD has a structure of electrodeand insulator (“gate” structure) made of poly-Si and SiO2 layers. The thicknesses of the2.3. X-RAY IMAGING SPECTROMETER

25

Electric
valve

Hood
105 mmBaffleDoor calibration
sourceWall calibration
sourceBonnet
125 mm

Door

Base
70 mm

Optical blocking
filter (OBF)
Cover shield
Thermo-electric
CCD
cooler (TEC)
Heat sink

Flexible print

Figure 2.11: X-ray Imaging Spectrometer (Koyama et al. 2007a). Left panel: Picture of
one of the XIS sensors. Right panel: Cross section of the XIS sensor.

Figure 2.12: Picture of the CCD installed in the base (Koyama et al. 2007a).

26

CHAPTER 2. SUZAKU X-RAY SATELLITE

Figure 2.13: Schematic view of the XIS CCD (Koyama et al. 2007a). The CCD consists offour segments (A, B, C, and D), each with a dedicated read-out node. An incident X-ray
is received on the top half of the CCD (imaging area). The signal charge corresponding
to the incident X-ray is transferred to the bottom half (frame-store area). The attachedcalibration source irradiates the top corners with Mn I K lines (5.895 and 6.49 keV X-rays).
The charge injection register is equipped at the top of the imaging area.

2.3. X-RAY IMAGING SPECTROMETER

27

poly-Si and SiO2 layers are ∼ 0.28 µm and ∼ 0.44 µm in the case of the XIS. The FI CCD
is less sensitive than the BI CCD to X-rays in the <2 keV band since the gate structureabsorbs soft incident X-rays. On the other hand, the depletion layer of the FI CCD is
thicker (∼ 65 µm) than that of the BI CCD (∼ 42 µm) and thus the FI CCD has highquantum efficiency to hard X-rays than the BI CCD.
For the in-orbit calibration, radiative isotopes 55 Fe are equipped (“Wall calibration
source” in figure 2.11). They emit the Mn I Kα (5.895 keV) and Kβ (6.492 keV) lines tothe upper edge of the segment A and D constantly (see figure 2.13).
The CCDs of the XIS are operated at a temperature of −90◦ C cooled by radiators andPertier coolers. Each XIS has an optical blocking filter (OBF) to block optical and UV
photons. The OBF is made of 100 nm polyimide coated with Al of a 120 nm thickness.
XIS 2 suddenly showed an anomaly on November 9, 2006, and it has not been operated
since then. Although there is no direct evidence, the micro-meteoroid impact might havecaused the anomaly 1 .

2.3.3

Charge Injection Capability

The XIS is equipped with the CI structure for the first time among the satellite-onboard
CCD cameras. A charge injection (CI) structure is installed adjacent to the topmost row
of the imaging area. The CI structure allows us to inject an arbitrary amount of charge
in a nearly arbitrary spatial pattern. The CI has been used in two ways to mitigate theeffect of the radiation damage in orbit; checker-flag CI (CFCI) and spaced-row CI (SCI)
techniques. The CI function is important to analyze the data used in this thesis because
the difference of the observation dates are up to three and half years. Generally, systematic
errors from a time degradation of a CCD become serious problems for such data sets. In
the case of the XIS data used in this thesis, however, such systematic errors are small dueto these CI techniques.
The function of the CI is described in figure 2.14. A serial CI register of 1024 pixels
long is attached next to the upper edge of the imaging area. An input gate is equippedleft of the CI register (figure 2.14). Pulling down the potential for electrons at the input
gate and the next electrode (S3 in right panel of figure 2.14), the potential well is filledwith charges with an amount of Q. Then, pulling up the potential makes the charges to be
spilled. The amount of charges is controlled by the offset voltage between the input gate
and the next electrode (S3). It can inject charges from ∼ 50 e− to ∼ 4000 e− per pixel;
the equivalent X-ray energies range from ∼300 eV to ∼15 keV. The deposited charges aretransferred horizontally through the serial CI resister, vertically through transferred into
the imaging area, and to an arbitrary position of the CCD chip. Prigozhin et al. (2008)
have reported details concerning the CI structure.
X-ray CCDs suffer from radiation damage in orbit. The damage causes the increaseof the charge transfer inefficiency (CTI) which is defined as the ratio of lost charges
to transferred charges in one pixel by one transfer. This causes decrease of gain anddegradation of energy resolution.
1

See http://www.astro.isas.jaxa.jp/suzaku/proposal/ao3/suzaku td/ and similar micro-meteoroid impact events of XMM-Newton are reported by Strüder et al. (2001)

28

CHAPTER 2. SUZAKU X-RAY SATELLITE

Figure 2.14: Schematic view of the charge injection (CI) technique (Bautz et al. 2004).
Left: CI structure. Charge packets are injected to the charge-injection register from an
input gate located at the edge of the register. After depositing the packets over the register,a vertical clock runs to inject the packets into the imaging area. Right: Schematic viewfor the injection of charges at the input gate. The offset voltage between the input gate
and the S3 electrode can control the amount of injected charge (Q).

The main origin of the CTI is thought to be the increase of defects in the crystal latticegenerated by cosmic-rays in orbit. These defects form charge traps that can capture signal
electrons. The CTI results in the degradation of the determination accuracy of the energy
and the energy resolution for two reasons: (1) the pulse height strongly depends on theposition of an X-ray event, since the X-ray event loses more electric charges as the numberof transfer increases, and (2) the loss of charge is a stochastic process, and thus there is afluctuation in the amount of lost charge.
Checker-Flag Charge Injection
One strategy to mitigate the effect of the CTI is to measure the CTI of each column
precisely and correct the lost charges. Here the ”column” is defined as the vertical columnof the pixels along the ACTY in figure 2.13. The checker-flag CI (CFCI) technique makes
the column-to-column CTI measurement possible. Here the outline of the CFCI technique
is described referring figure 2.15. First, a “test” charge packet is injected into the top CCD
row. Then, after a gap of a few rows, five continuous packets are injected with the sameamount of charge of the test packet. The former four packets are called “sacrificial” charge
packets, while the last one is called a “reference” charge packet. The test packet loses its
charge by charge traps. On the other hand, the reference packet does not suffer from any
charge loss, because the traps are already filled by the preceding sacrificial charges. Thus,
we can measure the CTI of each column by comparing the pulse heights of the reference
charge and the test charge.
Using the obtained CTI values, the column-to-column CTI correction is applied for the2.3. X-RAY IMAGING SPECTROMETER

29

Figure 2.15: Schematic view of measuring procedure of CTI with the CFCI technique(Nakajima et al. 2008). The amount of lost charge (δQ) can be estimated by comparing
the PHA of the test packet, which suffers from CTI, to that of the reference packet, which
is not affected by the traps.
data observed between the launch and September 2006. The energy resolution of the FI
CCD at 5.9 keV was ∼140 eV (FWHM) at August 2005, and had degraded to ∼200 eV at
August 2006 without the column-to-column CTI correction. The column-to-column CTI
correction improves the energy resolution to ∼170 eV (FWHM) at August 2006. Nakajima
et al. (2008) and Ozawa et al. (2009) reported the details of the CFCI in orbit and thecolumn-to-column correction.
However, we cannot correct the fluctuation by the readout-to-readout in a specificcolumn, which still degrades energy resolution even after the column-to-column CTI correction in principle.
Spaced-row Charge Injection
The spaced-row charge injection (SCI) technique can reduce the CTI actively andimprove the energy resolution. In the SCI mode, a charge packet is injected into CCD
rows periodically during the observation. The injected charge fills the radiation-induced
traps as a sacrificial charge packet, and thus prevents some of the traps from capturing
signal charges produced by X-rays. Results based on ground experiments using the SCI
technique with radiation-damaged CCDs have been reported by Tomida et al. (1997) andBautz et al. (2004), but no in-orbit experiment had been done.
The Suzaku XIS operated with the SCI technique in orbit for the first time from August
2006. In the SCI mode of the XIS, a charge is injected into every 54th row. The amountsof the injected charge into each pixel are equivalent to the X-ray energy of ∼6 keV andof ∼2 keV, for the FI and BI, respectively. The energy resolution was improved from∼200 eV to ∼140 eV at 5.9 keV (Bautz et al. 2007). The SCI has been applied in anormal observation since October 2006 (the SCI mode). In the SCI mode, since the CTI

30

CHAPTER 2. SUZAKU X-RAY SATELLITE

depends on the distance from a charge injected row, the pulse heights show a sawtooth
distribution depending on the transfer number shown in figure 2.16. This non-uniformityof the pulse heights degrades the determination accuracy of the energy and the energy
resolution. Thus we developed a new CTI correction method for the SCI mode, which iscalled a “sawtooth” CTI correction. This method corrects the non-uniformity like shown infigure 2.16. All the XIS data of the SCI mode are processed with the sawtooth correction.
The details of the sawtooth correction and CTI measurement method for the SCI modeare described in Uchiyama et al. (2009).

1640

Pulse height (ch)

1630
1620
1610
1600

Readout
Corrected

1590
800

830

860
890
920
i (transfer number)

950

980

Figure 2.16: Relation between pulse height of the Mn I Kα line from the onboard calibration source and transfer number i: This is the data of February 2008. Black and red
marks are data before and after our new CTI correction, respectively (Uchiyama et al.
2009).

2.3.4

XIS Performance in orbit

Determination Accuracy of the Energy Scale
Figure 2.17 shows the time histories of the center energy of the Mn I Kα line obtainedwith the onboard calibration sources. These data are processed with the column-to-column
CTI correction for the SCI-off data or the saw-tooth CTI correction for the SCI-on data.
The center energy of the Mn I Kα line have been constantly 5.895 keV. The uncertaintyof the energy-scale determination is less than ∼ 20 eV (0.3%) at 5.9 keV.
Figure 2.18 shows the time histories of the center energy of the O VII Kα line (0.570
keV) obtained from 1E 0102-7219, which is a well-studied supernovae remnant and commonly used as a celestial calibration source for X-ray astronomical satellites recently.
Narrow lines from 1E 0102−7219 were resolved with the grating spectrometer on XMMNewton (see Rasmussen et al. (2001) and Plucinsky et al. (2008) in detail). The uncertainty of energy-scale determination of the XIS is less than ∼ 10 eV (1–2%) at 0.5–1 keV.
There is no good celestial calibration source for the 1.5–3 keV band, and the gain
and response functions may jump at the Si-K edge (1.839 keV). Therefore, the relatively

2.3. X-RAY IMAGING SPECTROMETER

31

Figure 2.17: Time histories of the CTI-corrected central energy of the Mn I Kα line fromthe 55 Fe calibration sources for XIS 0 (FI) and XIS 1 (BI). The left and right figures show
data taken in the SCI-off (Ozawa et al. 2009) and SCI-on (Uchiyama et al. 2009) modes,
respectively.

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

170

165

~ 4 eV

160
400

Time since 2005-7-10 [day]

500

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

900

1000

Figure 2.18: Time histories of the center energy of the O VII Kα of 1E 0102.2-7219. Left:
SCI-off mode (Ozawa et al. 2009). Right: SCI-on mode, where the center energy is
described in the unit of pulse height channel and the result of XIS 1 is shown (Uchiyama
et al. 2009).

32

CHAPTER 2. SUZAKU X-RAY SATELLITE

Figure 2.19: Schematic response function of the XIS. 1. main peak, 2. sub peak, 3.
triangle component, 4. Si escape, 5. Si Kα line, and 6. constant component. The physicalinterpretation of the each component is shown in Matsumoto et al. (2006).
large uncertainty remains in the band. Recent studies of SNRs using the XIS show theenergy-scale determination in the 1.5–3 keV can have the systematic error of 20–30 eV(1–2%).
Response and Energy Resolution
The pulse height distribution of the signals for a monochromatic incident X-ray energy
is called a response function of the detector. The schematic response function of the XIS
is shown in figure 2.19. The main-peak component is made by complete-gathered signal
charges. The sub-peak, triangle and constant components come from events some of whichcharges are lost. The Si escape line happens when a fluorescent Si X-ray photon generated
by a photoelectric absorption escapes. The Si Kα line also occurs when the fluorescent
Si photon is absorbed in an other pixel. The details of physical interpretations of the
response function are reported by Matsumoto et al. (2006).
The energy resolution of the XIS is defined as the width of the main peak. The
time history of the widths of Mn I Kα from the onboard calibration sources is shown infigure 2.20. The intrinsic width of the Mn I Kα line is almost 0 eV, although, strictly
speaking, since this line is a blend of Mn K I α1 5.888 keV and α2 5.899 keV, the widthof the merged Mn I Kα is a few eV. Figure 2.20 shows the time histories of the energy
resolution at 5.9 keV. The energy resolution of the BI has degraded more rapidly thanthat of the FI in the SCI mode. This may be caused by difference of the amount of theinjected charges.
The time history of the energy resolution at 0.57 keV (O VII Kα) is shown in figure 2.21.
This figure is obtained with the observations of 1E 0102.2-7219. The calibration databasefor the XIS includes these time history of the energy resolution properly.

2.3. X-RAY IMAGING SPECTROMETER

33

Figure 2.20: Time histories of the width of the Mn I Kα line obtained with the onboard
calibration sources. Left: in the SCI-off mode (Ozawa et al. 2009). Right: in the SCI-on
mode (Uchiyama et al. 2009).

O VIII Kα FWHM (eV)70

XIS 1(BI)60

XIS 0(FI)
50

400

500

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

900

1000

Figure 2.21: Time history of the line widths of the O VII Kα obtained with the onboard
calibration sources in the SCI-on mode (Uchiyama et al. 2009).

34

CHAPTER 2. SUZAKU X-RAY SATELLITE

Figure 2.22: Quantum efficiency of the XIS sensor as a function of incident energy (Koyama
et al. 2007a). The solid and dashed lines indicate the FI and BI CCDs, respectively.
Quantum Efficiency and Effective Area
Conversion efficiency of an incident X-ray photon to an detected X-ray event is called
quantum efficiency (QE), as we shown in figure 2.22. The difference of the QEs between
the FI and the BI comes from the illumination type and the thickness of the depletion
layers.
The QE in the low energy band has degraded by unexpected contamination (due toout-gassing from the satellite) accumulating on the OBF (Koyama et al. 2007a; Mori etal. 2007). The contaminant is mainly composed of carbon and oxygen. Although thechemical composition of the contaminant is still uncertain, it is assumed in the calibration
to be C/O=6 in number. The time evolution of the contaminant has been measured using
regular observations of 1E 0102.2-7219. Note that, in this work, Fe Kα lines (∼ 7 keV) aremain targets and the contamination effects almost nothing to the energy band.
Non-X-ray Background
One of the characteristic advantages of the XIS is a low and stable non-X-ray background (NXB), which is caused by cosmic-ray particles in the Suzaku orbit. Figure 2.23
shows the NXB spectra of the FI and BI sensors when the XIS observed the dark (night)
Earth (no celestial X-ray is incident). The data during the passage through the south
Atlantic anomaly (SAA) or from the calibration source area is excluded. The lines at 5.9
and 6.5 keV are due to scattered X-rays from the calibration sources (Mn Kα and Kβ).
Other than these, many lines from materials in the XIS and the XRT, Kα lines of Al, Si,
Ni, Kβ of Ni, and Lα, Kβ, and Mα lines of Au are detected.
The NXB rates in the 0.4–12 keV band is 0.1–0.2 counts s−1 in the FI sensors and2.3. X-RAY IMAGING SPECTROMETER

35

Figure 2.23: Spectra of the non-X-ray background for XIS 0 (black), XIS 1 (red), and XIS 3(blue). The spectra are constructed from night Earth observations from the SCI start to2010 May.
0.3–0.6 counts s−1 in the BI sensor. The NXB is not uniform over the chip. Since somefraction of the NXB is produced in the frame-store region and the NXB events made in
pixels of the larger ACTY take the longer time to be read out (figure 2.13; Yamauchi
et al. 2006), the flux tends to be slightly higher at larger ACTY as shown in the left
panel of figure 2.24. The NXB also depends on the cut-off rigidity (COR) as shown in
the right panel of figure 2.24. The NXB flux is higher at lower COR, where many cosmicray particles encounter the satellite. Thus, we need to select the NXB for an appropriate
range of the COR in order to subtract the NXB from an observed data. Tawa et al. (2008)
modeled the NXB spectra as a function of the COR and the position on the CCD. The
model reconstructs the NXB spectra within the accuracy of 5%.
Thanks to the good energy resolution, large collecting area, and the low/stable NXB
background, the XIS has a good spectral performance especially in the 5–10 keV band,
including Fe-K lines. These capabilities also enable us to make line-resolved imaging forlow surface brightness sources.

36

CHAPTER 2. SUZAKU X-RAY SATELLITE

Figure 2.24: Left: NXB rate (average intensity for 5–10 keV) depends on the cut-off rigidity(COR) for each XIS sensor (black: XIS 0, red: XIS 1, blue: XIS 3). The NXB fluxes varyby a factor of ∼ 2 depending on the COR (Tawa et al. 2008). Right: Dependence of the
NXB on ACTY. Black and red lines indicate the continuum (2.5–5.5 keV) and Ni Kα line
(7.2–7.8 keV) components, respectively (Tawa et al. 2008).

Figure 2.25: XIS background counting rate as a function of energy. They are normalized
by the effective areas and the fields of view, which is a good measure of the sensitivity
determined by the background for spatially extended sources. The background spectra of
ASCA, Chandra, and XMM-Newton are shown for comparisons (Mitsuda et al. 2007).

Chapter 3

Discovery of K-shell Lines of
Neutral Atoms from the Galactic
Center Region
0.5

Sgr B2 (Chapter 5)0.0Arches cluster

M359.43-0.07
Sgr A(Chapter 3) M359.47-0.15

M0.74-0.09

Sgr A*
Sgr B1(Section. 4.1)

Sgr C (Section 4.2)

-0.5
1.0

0.5

0.0

359.5

359.0

Figure 3.1: X-ray image in the 6.4 keV line taken by Suzaku/XIS. The transverse andvertical axes are the Galactic longitude and latitude in the unit of degree (◦ ). The solidboxes indicate the field of view of the observations used in this work. We see the brightest
6.4 keV cloud among the GC region in the Sgr A region indicated with the yellow square.
The Arches cluster is in the northeast beside of the cloud (e.g. Tsujimoto et al. 2007).
The super-massive black hole, Sgr A* inhabits at a position of the small circle. The otherregions are referred in the next chapters.

3.1

Introduction

ASCA found clumpy structures in the 6.4 keV line from the Sgr B2 and the RadioArc regions (Koyama et al. 1996). The clumps correspond to giant molecular clouds
37

38

CHAPTER 3. K-SHELL LINES OF NEUTRAL ATOMS

found by radio observations (e.g. Tsuboi et al. 1999). The X-ray spectrum of the Sgr B2
cloud has a prominent 6.4 keV line on the continuum emission with deep absorption of
NH ∼ 1024 cm−2 (e.g. Murakami et al. 2000). Recently, many new 6.4 keV clouds havebeen discovered in a variety of regions (Koyama et al. 2007b; Yusef-Zadeh et al. 2007;
Nobukawa et al. 2008; Fukuoka et al. 2009; Nakajima et al. 2009).
Koyama et al. (1996) and Murakami et al. (2000) suggested that the 6.4 keV emission
from the 6.4 keV clouds is due to the K-shell ionization of iron atoms by external X-rays,possibly from the super-massive black hole, Sgr A* (an X-ray reflection nebula; XRN).
The present X-ray luminosity of Sgr A* is LX ∼ 1033 erg s−1 , although some flares with
LX = 1034 –1035 erg s−1 have been found by Chandra and XMM-Newton (Baganoff et
al. 2001, 2003; Porquet et al. 2003). However, the luminosity required for the 6.4 keVemission is 1038−39 erg s−1 cm−2 . Therefore, Sgr A* in a few hundred years ago would
have been 103−6 times brighter than now (e. g. Koyama et al. 1996). On the other
hand, Yusef-Zadeh et al. (2007) proposed that the origin of the 6.4 keV emission would
be low energy cosmic-ray electrons (Ee =10–100 keV), because they found that the X-rays
correlated with non-thermal radio filaments. The origin of the 6.4 keV line has been an
open question so far.
Neutral atoms of lighter elements than Fe, such as Ar and Ca, should also exist in the
6.4 keV clouds, and hence would emit K-shell emission lines. These would provide new
information to constrain the origin of the 6.4 keV clouds in the GC region since there is adifference between the X-ray and electron origins, for example, between the cross sectionsof K-shell photo-absorption by X-rays and K-shell ionization by electrons. However, only
the K-shell lines of neutral Fe and Ni atoms have been discovered so far.
The Sgr A region is the best place for the search of neutral K-shell lines of various elements because many bright clouds have been found at (l, b)∼(0.1◦ , −0.1◦ ) (see figure 3.1)
and the absorption column density is relatively small among the 6.4 keV clouds in the GC
region. Moreover, many observations with the total exposure time ∼ 200 ks have been
performed with Suzaku (Koyama et al. 2007c; Hyodo et al. 2009).
We analyzed the X-ray data obtained with the Suzaku/XIS and found emission lines
from neutral Ar, Ca, Cr, and Mn atoms for the first time. This chapter reports on detailed
analyses, results, and the severe constraint on the origin of the 6.4 keV emission.

3.2

Observations and Data Reduction

The eastern vicinity of Sgr A* was observed 4 times in 2005, 2006, and 2007 with
Suzaku/XIS. Each pointing position and roll angle were almost the same. The observation
log is listed in table 3.1.
The XIS observations were made in the normal clocking and the full-window mode.
XIS 2 has been out of use since 2006 November, hence we did not use the XIS 2 data in 2007
September. The data in 2007 September were taken with the SCI technique. For the data
reduction, we used the xispi software version 2008-04-10, and the calibration database
version 2008-08-25. We removed the data during the epoch of low-Earth elevation anglesof less than 5◦ (ELV< 5◦ ), day Earth elevation angles of less than 10◦ (DYE ELV< 10◦ ),3.3. ANALYSIS AND RESULTS

39

Table 3.1: Observation Data List.
Observation ID
100027010
100037040
100048010
102013010

R.A.
266.5146
266.5133
266.5135
266.5129Decl.
-28.9267-28.9266-28.9269-28.9278

XIS
0123
0123
0123
013

SCI
off
off
off
on

obs start
2005-09-23
2005-09-30
2006-09-08
2007-09-03

exposure time*
44.8 ks
43.0 ks
63.0 ks
51.4 ks*After the data screening described in the text.

and the South Atlantic Anomaly. The good exposure times are listed in table 3.1. The
overall spectral resolutions (FI/BI) at 5.9 keV were 150/150, 175/185, and 140/175 eV(FWHM) for the 2005, 2006, and 2007 observations, respectively.
We sorted the NXB with the COR values using xisnxbgen from the night-Earth data
released by the Suzaku XIS team (see section 2.3.4). The COR-sorted NXB was subtracted
from the raw data with the same COR.
Since we checked that the relative gain of the FI sensors were well calibrated and the
response functions were essentially the same, we co-added the FI data. We also summed
the spectra of the four observations to increase the statistics.
We analyzed the data using the software package HEASoft 6.6.1. For spectral fittings,
we made XIS response files using xisrmfgen, and auxiliary files using xissimarfgen.
Errors quoted in this chapter are estimated at the 90% confidence levels unless otherwise
stated.

3.3

Analysis and Results

3.3.1

6.4 keV line Image

In order to depict 6.4 keV clouds from the GC region, we made an X-ray image inthe 6.4 keV line by the following procedures. We first made X-ray images in the 5–6 keV
and in the 6.3–6.5 keV bands, after subtracting the NXB and correcting the vignetting
effect. We extracted the X-ray spectrum from a 80 circular region near the field center,and fitted the 5–6 keV band spectrum with a power-law model. The best-fit photon index
in the 5.0–6.0 keV band is 1.25 ± 0.07. The ratio of the photon flux in the 6.3–6.5 keVband to that in the 5.0–6.0 keV band was estimated to be 0.165, by extrapolation of thepower-law index. We multiplied the 5.0–6.0 keV band image by 0.165, and subtracted
from the 6.3–6.5 keV band image. The result is shown in figure 3.2.
We can see a bright region near the center of the 6.4 keV line image, as is shownwith the solid line of a gourd-like shape (hereafter, the Source region). Around the Source
region, we can see fainter emissions, which may be due to unresolved miner 6.4 keV clouds.
The X-ray emission from the GC hot plasma contributes to the spectrum of the Source
region. In order to examine the contribution of the GC hot plasma accurately, we selected
the dashed region, excluding the supernova remnant, Sgr A East (Maeda et al. 2002).,

40

CHAPTER 3. K-SHELL LINES OF NEUTRAL ATOMS

6.4 keV line map

Arches cluster

0.1

Galactic latitude (deg)

Source0.0-0.1
Sgr A*-0.2
Reference
0.3

0.2

0.1

0.0

359.9

Galactic longitude (deg)Figure 3.2: X-ray image in the 6.4 keV neutral iron line (the underlying continuum fluxes
are subtracted). The solid square is the FOV of the XIS. The small circle indicates the
position of Sgr A*. We extracted the source spectrum from the solid region (Source). In
order to examine the contribution of the GC hot plasma, we also selected the surrounding
region (Reference) marked with the dashed line. The areas of the Source and Reference
regions are 43.3 and 40.5 arcmin2 , respectively. The bright source in the northeast edgeof the FOV is the Arches cluster (e.g., Tsujimoto et al. 2007).
whose size and position are ∼ 30 and (l, b)∼(−359.5◦ , −0.5◦ ). Hereafter, we call this region
as Reference.

3.3.2

Gain Tuning

The nominal uncertainty of the absolute energy of XIS is ∼ 10 eV (the Suzaku XIS
team1 ). To be more accurate, we checked the energy of the Kα lines of H-like S (S XV) of
the spectra and the calibration source 55 Fe (emits the Mn I Kα line at 5895 eV) attachedat the corner of the XIS. Compared with the theoretical value obtained from Atomic &Molecular Database in the Institute of Applied Physics and Computational Mathematics(CAMDB)2 , we found that the energy differences of S XV Kα and Mn I Kα were 5 ± 2
and 0 ± 1 eV for FI, while those for BI were consistent with the theoretical energies. We,12

http://www.astro.isas.ac.jp/suzaku/doc/suzaku td
http:www.camdb.ac.cnespectraspectra search.asp

3.3. ANALYSIS AND RESULTS

41

CaXVIII Kα

SXIV KβArXVI KαSXIV KαSXV KαSiXIV Kβ

0.1
0.01

Counts s-1 keV-1

SourceReference

FeI KαFeXXV KαFeXXVI Kα+ FeI Kβ
Ni I Kα25Energy (keV)

10

Figure 3.3: X-ray Spectra (FI) in the Source (black) and the Reference (red) regions,where the NXB was already subtracted. Identified emission lines are indicated by thedashed lines with their corresponding elements. Errors of the data are estimated at the1σ confidence level.
therefore, fine-tuned the energy scale for FI as a liner function;∆E = −0.0015 × (E0 − 5895) (eV),(3.1)

where E0 and ∆E are the original and energy shift, respectively.

3.3.3

Model Construction

Figure 3.3 shows the X-ray spectra of the FI sensor in the Source and Reference
regions. The spectra have several ionized emission lines of heavy elements, such as Si, S,
Ar, Ca, and Fe. These lines would come from the GC hot plasma with the temperatureof kT ∼ 7 keV (Koyama et al. 2007c) and possibly kT ∼1 keV (Ryu et al. 2009).
On the other hand, we can see prominent Kα lines of neutral Fe and Ni at 6.4 keV
and 7.5 keV in the spectra. The neutral lines indicate the existence of a large amount of
neutral Fe and Ni atoms in the regions. The neutral lines from the Reference region may
be due to many faint unresolved 6.4 keV clouds.
The X-ray spectra of the two regions (Source and Reference) would be a complexof the GC hot plasma and the 6.4 keV clouds (Koyama et al. 2009; Ryu et al. 2009).
We, therefore, made a fitting model composed of the GC hot plasma (here P lasma), the6.4 keV clouds (hear N eutral) and the Cosmic X-ray background (here CXB). Thus, themodel is given by;
M odel = Abs1 × [P lasma + N eutral] + CXB,

(3.2)

42

CHAPTER 3. K-SHELL LINES OF NEUTRAL ATOMS

0.1

Fe

0.01

Ar

Ca
Cr

Mn

5
5 −5 0

(c)

(d)

0

(b)

5 −5 0

Ni

−5

χ

χ

χ

Counts s−1 keV−1

(a)

2

5
Energy (keV)

10

Figure 3.4: (a): FI Spectrum of the Source region fitted with the final model (step 3)
in subsection 3.3.4 (see text). The red and blue histograms are the best-fit P lasma andN eutral components, respectively. The Gaussian line at 3.2 keV is added to compensate
the uncertainty of the instrumental response. The blue and red dashed lines indicate theenergies of the neutral Kα of Ar, Ca, Cr, Mn and the He-like Kα of Cr and Mn. (b):Residual of the fitting with the models of 2-kT AP EC, a power-law, and neutral lines of
Fe I Kα, Kβ, Ni I Kα (step 1). (c): Same as (b), but the ionized Cr and Mn lines wereadded (step 2). (d): Same as (c), but the neutral Ar, Ca, Cr, and Mn Kα lines were added(step 3). Errors of the data were estimated at the 1σ confidence level.

3.3. ANALYSIS AND RESULTS

43

where Abs1 is absorption of the inter-stellar medium toward the GC region. The absorption is determinated by the spectral shape at a lower energy band where photo-absorption
is more effective than at a high energy band.
Ryu et al. (2009) found that the GC hot plasma (P lasma) has two-temperaturecomponents, and hence the P lasma model can be described as
P lasma = AP EC1 + AP EC2 (photons s−1 cm−2 ),(3.3)

where AP EC1 and AP EC2 are thin thermal plasma code in the XSPEC package for thelow (kT1 ) and high (kT2 ) temperature components, respectively.
N eutral is composed of neutral lines and associated continuum emission, and hence
can be given by
N eutral = Abs2 × [A × (E/keV)−Γ + Gaussians](photons s−1 cm−2 ),(3.4)

where Abs2 is self-absorption in the 6.4 keV clouds. Neutral lines are expressed by Gaussians, while the continuum is represented by a power-law model with the photon index Γand the normalization A.
The cosmic X-ray background (CXB) flux is lower than the whole spectra shown infigure 3.3 by more than two orders of magnitude. We nevertheless add the CXB model
made by Kushino et al. (2002),
CXB = Abs1 × Abs1 × 7.4 × 10−7 × (E/keV)−1.41(photons keV−1 s−1 cm−2 arcmin−2 ),(3.5)

where absorption column density (Abs1) is applied twice because the origin of the CXB
is extragalactic.

3.3.4

Spectral Fitting

Since the surface brightness of P lasma is not constant, the spectrum of N eutral cannot
be directly obtained by the subtracting the surrounding (Reference) region. We therefore
tried to simultaneously fit these two regions and to determine the spectra of the N eutral
and P lasma models, separately.
The X-ray photons below ∼2 keV from the GC region are heavily absorbed by aninter-stellar medium of NH ∼ 6 × 1022 cm−2 , and the X-ray spectrum below ∼2 keV
is dominated by the local background (Ryu et al. 2009). We therefore used only the
2–10 keV band for the spectral analysis.
In the fitting of P lasma, we assumed that the temperatures (kT1 and kT2 ) and elemental abundances (Z) of Si, S, Fe, and Ni were common in the two regions (Source andReference). Abundances of Ar and Ca were set to be the same as that of S. Also, Abs1 wastreated to be common in the two regions. On the other hand, we set the normalizationsof the hot plasma (AP EC1, AP EC2) to be independent free parameters.
For N eutral, common free parameters among the two regions are the power-law index(Γ), the equivalent widths (flux ratio) of the neutral lines to the continuum, and line center

44

CHAPTER 3. K-SHELL LINES OF NEUTRAL ATOMS

energies. The normalization of the power-law (A) is an independent free parameter. With
these constraints, we proceeded the model fitting with equation (3.2) along the following
3 steps:

Step 1: Gaussians for K-shell Lines from Fe and Ni
For the N eutral model of equation (3.4), we first tried a fitting with three Gaussians
for Fe I Kα (6.40 keV), Fe I Kβ (7.06 keV), and Ni I Kα (7.49 keV) lines (e.g. Koyama etal. 1996; Koyama et al. 2007c).
This fit left many residuals, including that at ∼ 3.2 keV. The energy is near the M-edgeof Au (3.2–3.4 keV), and hence the residuals are likely due to the calibration uncertainty of
the response of the XRT. In fact, similar residuals were reported by Kubota et al. (2007).
We, thus, added a broad Gaussian line at ∼ 3.2 keV according to Kubota et al. (2007).
The best-fit center energy and the line width (1 sigma) of the Gaussian are 3.24 ± 0.02 keV
and 0.07 ± 0.02 keV, respectively.
The fitting residuals are displayed simply for the FI spectra of the Source region infigure 3.4b, although the fitting was made simultaneously for the both sensors (FI and BI)
and both the two regions. As shown in figure 3.4b, this model was not able to reproduce
the observed spectra (χ2 /d.o.f is 1447/906).

Step 2: K-shell Lines from Highly Ionized Cr and Mn
One reason of the unacceptable fit is that the APEC model does not contain Cr and
Mn K-lines at 5–6 keV. We therefore added two Gaussian lines in the P lasma model of
equation (3.3) at ∼ 5.7 and 6.2 keV for the Cr XXIII and Mn XXIV Kα lines, then, thecenter energies were obtained to be 5.68 ± 0.02 and 6.19 ± 0.02 keV, respectively. These
are consistent with theoretical energies of 5.67 and 6.17 keV (CAMDB), the He-like Cr(Cr XXIII) and Mn (Mn XXIV).
The hot plasma can also emit Kα lines of H-like Cr (Cr XXIV) and Mn (Mn XXV) at
5.94 keV and 6.44 keV, respectively. Using the two-APEC model with the temperaturesof 1 keV and 7 keV and the flux ratios given in table 3.2, we plotted the intensity ratio of
H-like to He-like Kα lines as a function of atomic number (Ar, Ca, Fe) in figure 3.5. Byinterpolation, we estimated the line intensity ratios of Cr and Mn to be 0.22 and 0.28. We
added these four lines in the P lasma model:
P lasma = AP EC1 + AP EC2 + 4 Gaussians (Cr, Mn)(photons s−1 cm−2 ),(3.6)

where the intensity ratios of H-like to He-like Cr and Mn were fixed to be 0.22 and 0.28.
We also set the equivalent widths of Cr XXIII Kα and Mn XXIV Kα to the continuum of
the P lasma model to be the same in the two regions.
This fitting gave χ2 /d.o.f. of 1367/902. The residuals from this model in the Source
spectrum are shown in figure 3.4c. In the next step (step 3), we used the P lasma modelof equation (3.6).

3.3. ANALYSIS AND RESULTS

45[H-like]/[He-like]

0.5
0.4

Fe

Mn

0.3

Cr
0.2
0.1
0

Ar
18

Ca

20

22

24

26

28

Atomic Number
Figure 3.5: Line-intensity ratios of He- and H-like ions. The filled circles are calculated
values using the APEC model. The values of Ar, Ca, and Fe are 0.057, 0.085, and 0.35,
respectively. By interpolating these data, the line intensity ratios of Cr and Mn are
estimated to be 0.22 and 0.28 (open circles).

Step 3: K-shell Lines from Neutral Ar, Ca, Cr, and Mn
Since we still found line-like residuals at the energies of ∼3.0, 3.7, 5.4, and 5.9 keV,
we added four Gaussian lines to the N eutral model of equation (3.4) at these energies(in figure 3.4c, also see dashed lines in figure 3.4). Then the line energies were found tobe 2.94 ± 0.02, 3.69 ± 0.02, 5.41 ± 0.04, and 5.94 ± 0.03 keV, respectively. These values
correspond to theoretical energies of 2.96, 3.69, 5.41, and 5.90 keV for the Kα lines of Ar I,
Ca I, Cr I, and Mn I. The fitting residuals in the spectrum of the Source region are shownin figure 3.4d.
As demonstrated in figure 3.4d, the χ2 /d.o.f. is largely improved from 1367/902 (step2) to 1203/894 (step 3), but is not acceptable even at the 99% confidence level. Although
the energy resolution in the hard X-ray band is well calibrated with the calibration source
55 Fe (emits the Mn I Kα line), that in the soft band, especially about 2 keV, may havesome uncertainty due to the absence of available calibration sources. Thus the deviation
at ∼2.5 keV would come from the incomplete response function in the soft energy band.
Residuals are also found at ∼6.5–7.0 keV, with a slight mismatch of the center energyof Fe XXV Kα. This suggests that the GC hot plasma has more than two temperature components. We therefore tried a fitting with a three-kT plasma model. Although the fit was
improved from χ2 /d.o.f.=1203/894 to 1123/891, the best-fit parameters for the N eutral
and P lasma models do not change from those in the two-kT model within the statisticalerrors. We, thus, regard that the two-kT model after step 3 is a good approximation toderive the physical parameters.

46

CHAPTER 3. K-SHELL LINES OF NEUTRAL ATOMS

Table 3.2: Fitting results of P lasma.
Component
Abs1
AP EC1AP EC2Abundances

Cr XXIII Kα
Mn XXIV KαParameter
Unit
22
NH
10 cm−2
kT1
keV
Norm
†
kT2
keV
Norm
†
Si
solar
S, Ar, Ca
solar
Fe
solar
Ni
solar
Gaussian Linesenergy
(keV)
EW
(eV)
energy
(keV)
EW
(eV)

Value∗
6.75 ± 0.13
1.01+0.01−0.02
12 ± 17.0 ± 0.1
1.3 ± 0.1
2.52+0.09−0.17
1.87 ± 0.07
1.16+0.07−0.04
1.64 ± 0.37
5.68 ± 0.02
22 ± 6
6.19 ± 0.02
39 ± 6∗ The uncertainties indicate the 90% confidence levels.
R† Emission measure 10−12 /(4πD2 ) ne nH dV , where D is the
distance to the source (cm), ne and nH are the electron andhydrogen density (cm−3 ), respectively.

For the discussion on the 6.4 keV cloud in the Source region, we list the best-fit parameters for the P lasma model in table 3.2, while those for the N eutral model are listed
in table 3.3. For comparison, the best-fit parameters of the Reference region are: normalizations of AP EC1 and AP EC2 are 6.8 ± 0.5 and 0.90+0.06−0.04 in the same unit shown in ta22
ble 2, respectively. Abs2 and normalization of the power-law component are 12.7+1.2−1.8 ×10−3 photons keV−1 s−1 cm−2 , respectively. The other parameters
H cm−2 and 3.0+0.5−0.4 × 10
are the same as the Source region.

3.4

Discussion for the X-ray Emission from the Sgr A region

3.4.1

The GC hot plasma

The absorption column density toward the GC region (Abs1) is ∼ 6.7 × 1022 H cm−2 ,a typical value to the GC region (e.g., Muno et al. 2004). This is the first constraint thatthe low temperature kT = 1.01+0.01−0.02 keV plasma is also in the GC region.
Koyama et al. (2007c), using the same data set of this paper, reported the temperatureof the GC hot plasma is kT = 6.5 ± 0.1 keV. The 1 keV plasma also emits Fe XXV Kα,
but no significant Fe XXVI Kα. Since Koyama et al. (2007c) ignored the Fe lines from the
1.0 keV plasma, they under-estimated the flux ratio of Fe XXVI Kα to Fe XXV Kα, i.e.,
the plasma temperature. Indeed, about 20% of the Fe XXV Kα line may come from the
1.0 keV plasma. The flux ratio of the Fe lines is estimated to be 0.35 and 0.42 about the3.4. DISCUSSION FOR THE X-RAY EMISSION FROM THE SGR A REGION

47

Table 3.3: Fitting Results of N eutral and those of calculated values.∗

Energy (keV)
2.94 ± 0.02
3.69 ± 0.02
5.41 ± 0.04
5.94 ± 0.03
6.404 ± 0.002
7.06 (fixed)
7.48 ± 0.02

Identification
Ar I Kα
Ca I Kα
Cr I Kα
Mn I KαFe I KαFe I Kβ
Ni I Kα

Photon Index
Normalization
Abs2(NH )Neutral Lines
Intensity† EW(eV)‡170+60140 ± 40−40
54+14
83 ± 13−99.5 ± 2.5
24 ± 77.4 ± 2.2
22 ± 7340 ± 10 1150 ± 90
40 ± 3160 ± 2018 ± 383 ± 13
ContinuumΓk22
10 cm−2

EW(XRN)§
45
35
10
7.6
730
120
53

EW(LECRe)§
12
10
3.5
2.8
270
38
18

1.87 ± 0.04
9.6+1.6−1.3
12.0 ± 1.1∗ The uncertainties indicate the 90% confidence levels.
† Absorption-corrected line intensity in the unit of 10−6 photon s−1 cm−2 .
‡ Observed equivalent widths of the neutral lines to the power-law continuum in equation (3.4).
§ Calculated equivalent widths in the XRN and LECRe scenarios when the metal abundances are the solar values
and the photon index of the incident X-rays is 1.87.
kNormalization at 1 keV in the unit of 10−3 photons keV−1 s−1 cm−2 .

6.5 keV and 7.0 keV plasmas by the APEC model, respectively. Taking the contribution of
the low temperature plasma into account, the result in Koyama et al. (2007c) is consistent
with our work.
The Fe and Ni abundances of 1.1–1.2 and 1.3–2.0 solar are more precise than theprevious work (Koyama et al. 2007c). The abundances of lighter elements, Si and S, (Ar,Ca), were respectively determined to be 2.3–2.6 and 1.8–2.0 solar, for the first time. The
highly ionized Cr and Mn lines in the GC region were discovered for the first time, but
details are beyond the scope of this thesis.

3.4.2

Origin of the Neutral Clump

We discovered the K-shell lines of neutral Ar, Ca, Cr, and Mn from the bright 6.4 keV
cloud toward the Sgr A region at the significance levels of 6.8, 9.6, 6.1, and 5.4 σ, respectively. The absorption column densities, NH (Abs2) in the Source and Reference regions
22 cm−1 , respectively.
are 12.0(±1.1) × 1022 and 12.7+1.2−1.8 × 10
The photon index, Γ of the continuum (power-law) is 1.87 ± 0.04, which is consistent
with the result of Koyama et al. (2009) in the same region. The equivalent widths of theneutral Kα lines to the power-law continuum are ∼140, 83, 24, 22, 1150, 83 eV in Ar, Ca,
Cr, Mn, Fe, and Ni, respectively.
The major possibility for the origin of the 6.4 keV clouds is the ionization of neutral
atoms by either low energy cosmic-ray electrons (LECRe: Yusef-Zadeh et al. 2007) with an

CHAPTER 3. K-SHELL LINES OF NEUTRAL ATOMS

2.5
2
1.5

Γ (Photon Index)

3

3.5

48

0

1
2
α (electron distribution)

3

4

Figure 3.6: Photon index Γ of the X-ray spectrum emitted from a molecular cloud whenelectrons with the spectral index α are injected (solid line).
energy of 10–100 keV, or X-rays of external sources (XRN: Koyama et al. 1996; Murakami
et al. 2000). The ionization cross-sections for these processes are very different. They also
produce a continuum emission: the bremsstrahlung in the LECRe scenario and Thomson
scattering in the XRN scenario. As a result, the two scenarios make different X-ray spectra.
In particular, sharp contrasts are the photon index of the continuum and the equivalent
widths of the neutral lines.
The continuum emission in the X-ray spectrum produced by the LECRe scenario is
an integration of the bremsstrahlung by electrons with various energies. According toTatischeff (2003), we calculated the X-ray spectrum with various indexes of their energy
distributions into a molecular cloud. Figure 3.6 shows a relationship between the electron
spectral index α and the X-ray photon index Γ. As a result, the photon index of Γ = 1.9
corresponds to an index of the LECRe source spectrum of α ∼ 3. On the other hand, theThomson scattering do not change the photon index from the incident spectrum of theexternal source.
Tatischeff (2003) and Yusef-Zadeh et al. calculated the observed X-ray spectrum produced in the LECRe model. On the other hand, Murakami et al. (2000) estimated the
XRN spectrum by a numerical simulation. We improved the calculations to the otherrelevant elements. Figure 3.8 shows the equivalent widths of the K-shell lines for Ar, Ca,
Cr, Mn, Fe, and Ni calculated in the XRN (black) and LECRe models on the assumptionthat the elemental abundances in the cloud is the solar values. We found that abundantelements such as Fe have a large equivalent width. Although the equivalent width depends
on the incident spectral index, we found a large difference between the absolute values inthe two scenarios. Thus, the equivalent width of the line is a good indicator for constraintof the origin.
The equivalent widths produced by the LECRe with the index α = 3 were estimatedas shown with the red solid line in figure 3.7 as well as in figure 3.8. These are listed intable 3.3 and plotted in figure 3.7 together with the best-fit results.
From figure 3.7, the equivalent width of each element in the solar abundance is larger

3.5. SUMMARY

49

10

100

1000

X−ray
Electron1

Equivalent Width (eV)

Equivalent Width

Ar

Ca

Cr
Mn
Element

Fe

Ni

Figure 3.7: Equivalent widths of Kα line of various neutral atoms. Black and red lines
are the calculated values for the X-ray (XRN) and electron (LECRe) scenarios, respectively. The data points marked with the open circles are the observed value in our work.
Errors were estimated at the 90% confidence level. The black and red dashed lines are to
guide eyes, which are XRN and LECRe scenarios in 1.6 solar and 4.0 solar abundances,respectively.
than that expected in the both of the LECRe and XRN scenarios. For the LECRe scenario,
the abundances of the molecular cloud must be ∼4-times larger than the solar value, while
the XRN scenario requires ∼1.6 solar abundance.
Since the molecular cloud may be formed by condensation of the ambient materials,
the abundances should be similar to those of 1–2 solar in the GC hot plasma. Accordingly,
the neutral lines from the GC region likely come from the fluorescence by external X-rays.

3.5

Summary

The summary of the results in this chapter is as follows:• We found that a kT = 1.0 keV plasma exists in the GC region as well as thekT = 7.0 keV plasma. The elemental abundances of Si, S (Ar, Ca), Fe, Ni were
measured to be 2.3–2.6, 1.8–1.9, 1.1–1.2, 1.3–2.0 times larger than the solar value.
We also discovered He-like Cr and Mn Kα lines from the GC region.
• K-shell lines of neutral Ar, Ca, Cr, and Mn were firstly discovered in addition to
those of Fe and Ni. Equivalent widths of these atoms are ∼140, 83, 24, 22, 1150,
and 83 eV, respectively.

50

CHAPTER 3. K-SHELL LINES OF NEUTRAL ATOMS

1

1.5
2
2.5
Observed photon index Γ

100

X−ray
Electron10100

Equivalent width (eV)X−rayElectron10

Equivalent width (eV)

1000

Ca

1000

Ar

3

1

1000
1.5
2
2.5
Observed photon index Γ

100

X−ray
Electron10

Equivalent width (eV)100
10Equivalent width (eV)X−rayElectron

1

3

1

1.5
2
2.5
Observed photon index Γ

3

1000

Ni

10
1

1.5
2
2.5
Observed photon index Γ

3

X−ray
Electron100

100

Equivalent width (eV)X−rayElectron

10

1000

Fe

Equivalent width (eV)

3

Mn

1000

Cr

1.5
2
2.5
Observed photon index Γ

1

1.5
2
2.5
Observed photon index Γ3

Figure 3.8: Equivalent width of various elements respective to a photon index of an observed X-ray spectrum in the XRN (black) and LECRe (red) models.

3.5. SUMMARY

51• The observed equivalent widths of Ar, Ca, Cr, Mn, Fe and Ni favor the X-ray
radiation origin to the molecular cloud with 1–2 solar abundance.

Chapter 4

Other 6.4 keV Clouds in the Sgr
B1 and C Regions
In figure 4.1, we see many molecular clouds emitting Fe I Kα X-rays other than in the
Sgr A region, which was analyzed in chapter 3). As is shown by solid squares in figure 4.1,
the Sgr B1 and Sgr C regions are at the east and west side relative to the GC, i.e., SgrA*, respectively. In this chapter, I investigate detailed X-ray features of 6.4 keV clouds
found in the Sgr B1 and Sgr C regions (Murakami et al. 2001; Yusef-Zadeh et al. 2007).
Analyses and discussions for Sgr B1 and C are made separately in section 4.1 and 4.2,
respectively.

0.5

Sgr B2 (Chapter 5)0.0Arches cluster

M359.43-0.07
Sgr A(Chapter 3) M359.47-0.15

M0.74-0.09

Sgr A*
Sgr B1(Section. 4.1)

Sgr C (Section 4.2)

-0.5
1.0

0.5

0.0

359.5

359.0

Figure 4.1: X-ray image in the 6.4 keV line same as figure 3.1. In this chapter, 6.4 keV
clouds, Sgr B1, M 359.43−0.07, and M 359.47−0.15, in the yellow boxed regions are focusedin this chapter.

53

54

CHAPTER 4. OTHER 6.4 KEV CLOUDS

4.1

Sgr B1

4.1.1

Introduction for Sgr B1

The radio complex Sgr B1 located to the west, or to the negative Galactic longitude of
the Sgr B2 giant molecular cloud consists of a H II region surrounded by a molecular loop
(Sofue 1990). The expansion of the molecular loop is explained by the past activity of the
H II region (Sofue 1990). A 6.4 keV cloud was discovered in this region by Yusef-Zadeh
et al. (2007). OH and H2 O masers have been found and a lot of intermediate and/or
low-mass stars are formed (Mehringer et al. 1993).
Yusef-Zadeh et al. (2007) suggested that the 6.4 keV X-rays are due to collision of
low-energy cosmic ray electrons with molecular gas because of a correlation between nonthermal radio filaments and the X-ray features taken by Chandra. The scenario explains
not only the Fe I Kα emission, but also cosmic-ray heating of molecular gas and diffuse
TeV emission from the Galactic molecular clouds. However, Yusef-Zadeh et al. (2007) did
not exhibit detailed features in the X-ray spectrum.
Thanks to the characteristic features of the large effective area and low/stable background, the XIS are instruments suitable to observe diffuse and faint objects. We had
observed the Sgr B1 region for about 100 ks with Suzaku in order to obtain a good spectrum of the 6.4 keV cloud detected by Yusef-Zadeh et al. (2007). We report on the detailed
analysis and result of the Sgr B1 cloud in this section.

4.1.2

Observation and Data Reduction

We observed the Sgr B1 region with Suzaku/XIS on 2006 March 27–29. The pointing
angle was (R.A., Decl.)=(266.7693◦ , -28.6295◦ ). The XIS data were taken in the normal
mode. The XIS pulse-height data for each X-ray event were converted to Pulse Invariant (PI) channels using the xispi software and the calibration database released by theSuzaku/XIS team. After removing the epoch of low earth elevation angle less than 5
degree (ELV< 5◦ ), the day earth elevation angle less than 10 degree (DYE ELV< 10◦ )and the South Atlantic Anomaly (SAA), the effective exposure time was about 95 ks. We
made the data reduction and analysis using HEADAS software 6.1.2 and XSPEC 11.3.2
for Sgr B1.
Since relative gain among FI sensors are well calibrated and the response functions are
essentially the same, we co-added their data. We also applied the non-X-ray background(NXB) data from the night Earth data released by the Suzaku XIS team1 . The NXB data
were sorted with the geomagnetic cut-off rigidity (COR) because of the variation of the
NXB corresponding to the COR values.

4.1. SGR B1

55

0.1

Ar

Ca

S

Fe

0.01

Counts s−1 keV−1

Sgr B1

2

5
Energy (keV)

10

Figure 4.2: X-ray spectrum of the Sgr B1 region (FI: XIS 0+2+3). The spectrum wascollected from the full FOV, but excluding the corners where the calibration sources (55 Fe)
illuminate Mn I K lines. The non-X-ray backgrounds were already subtracted.

4.1.3

Analysis and Results

The Overall Features
We show the X-ray spectrum of the Sgr B1 region in figure 4.2. It is the FI (the averageof XIS0, 2 and 3) spectrum extracted from the full FOV excluding the corners where thecalibration sources of 55 Fe illuminate. The NXB was already subtracted from the spectra.
The spectrum shows K emission lines from He-like and/or H-like ions of S, Ar, Ca and
Fe. In addition, Kα emission line at 6.4 keV and K-edge absorption at 7.1 keV of neutralor low ionized Fe are cleanly detected. Three Fe lines at 6.4, 6.7 and 6.9 keV are wellresolved thanks to the excellent energy resolution of the XIS.
We examined the characteristic lines of Fe I, Fe XXV, and Fe XXVI. We show narrow
band images of 6.4 keV (6.30 to 6.50 keV; Fe I) and 6.7 keV (6.57 to 6.77 keV; Fe XXV) in
figure 4.3a, b and c, respectively. They are co-added images of the four XIS, in which the
NXB was subtracted and the correction of exposure and vignetting effects were alreadydone. The exposure maps were made with xissim (Ishisaki et al. 2007) in HEADAS
software 6.1.2.
We see clearly a diffuse source in the elliptic solid region shown in figures 4.3a. We
hereafter call the source M 0.51−0.10. M 0.51−0.10 is identified with the 6.4 keV cloud
Sgr B1 discovered by Yusef-Zadeh et al. (2007) whose position of l, b ∼ (0.5◦ , −0.1◦ ) is
the same.
We here note the excesses near the western edge of the FOV in each image. A known
bright X-ray source with the flux of a few 10−10 ergs s−1 cm−2 , 1E 1743.1−2843, is located1

http://www.astro.isas.jaxa.jp/suzaku/analysis/xis/nte/

56

CHAPTER 4. OTHER 6.4 KEV CLOUDS

outside the FOV in the northwest direction (Del Santo et al. 2006) and its position ismarked with a small green circle in figure 4.3. The northwest excess is due to the XRTPSF (point spread function) tail of 1E 1743.1−2843. It has a hard X-ray spectrum, whichis consistent with the result that the northwest excess is seen in all the panels in figure 4.3.
Spatial Distribution of the Fe XXV Kα emission
Figure 4.3b and figure 4.4 show the narrow band image of 6.7 keV (Fe XXV), namely
distribution of the GCDX. Figures 4, 5 and 6 of Koyama et al. (2007c) gave the profileof the fluxes of the 6.7 keV line and the flux ratios of 6.7 and 6.9 (Fe XXVI) keV lines asa function of the galactic longitude between l = −0.4◦ –+0.2◦ . We made the same profiles
for the Sgr B1 region of l = +0.3◦ –+0.6◦ and compared them with the results of Koyama
et al. (2007c).
First, we divided the Sgr B1 region into 4×4 small regions excluding the four corners
illuminated by the built-in calibration sources as shown in figure 4.3c. Next, we obtained
the co-added spectrum of the 3 FIs (XIS0, 2 and 3) for each small region. We finallyobtained 6.7 keV and 6.9 keV line fluxes and their ratio for each small region by fitting
the spectrum with a continuum (thermal bremsstrahlung), Fe lines, which are Fe I Kα andKβ (6.4 and 7.06 keV), Fe XXV Kα (6.7 keV), and Fe XXVI Lyα (6.9 keV), and the cosmic
X-ray background (CXB). Koyama et al. (2006) suggested the 6.7 keV emission smoothly
extends to the Sgr B2 region (l = +0.6◦ –+0.8◦ ) with almost the same temperature. Thus,
we made follow-up analyses on the Fe emission lines for the Sgr B2 region in the same
method as done for Sgr B1.
In figure 4.5 we show the 6.7 keV line flux (top) and the ratio between the photon
fluxes of 6.7 keV and 6.9 keV lines (bottom) as a function of the Galactic longitude.
Excluding the data of the region at G 0.61+0.01 (will be mentioned later on), we foundthat the 6.7 keV line flux exponentially decreases from the peak at the GC (l ∼ −0.05◦ )and smoothly connects the Sgr B1 and B2 regions. The line flux ratios of [Fe XXVI]/[Fe
XXV] in the Sgr B1 and B2 regions are almost the same as the ones in l = −0.4◦ –+0.2◦ .
We note that the three data points marked with filled circles in the Sgr B2 region
having high 6.7 keV fluxes and low flux ratios between the 6.7 keV and 6.9 keV lines
correspond to SNR G 0.61+0.01, which Koyama et al. (2007b) recently discovered with
Suzaku. The plasma temperature of G 0.61+0.01 is kT ∼ 3 keV (Koyama et al. 2007b),
which is significantly lower than that of the GCDX (kT ∼6–7 keV; Koyama et al. 2007c).
We see that the 6.7 keV intensity of the small region [3-2] is rather higher than expectationfrom the whole tendency and the ratio of [Fe XXVI]/[Fe XXV] is lower than the average ofthe GCDX. G 0.61+0.01 is located to the east of the small region [3-2] (see figure 4.4). This
result suggests that the SNR would have a larger size than the one reported by Koyama
et al. (2007b) and the small region [3-2] is a part of the G 0.61+0.01.
Spectrum of a 6.4 keV Clump, M 0.51−0.10
We examine M 0.51−0.10, a diffuse source in the 6.4 keV narrow band image of figure
4.3a. We identified this diffuse source with the 6.4 keV cloud, Sgr B1 detected by Yusef-

4.1. SGR B1

57(b) 6.7 keV(a) 6.4 keV
1E1743.1-2843

0.0

0.0
M0.51-0.10

b

b

-0.2

-0.2

0.5

0.3

0.5

0.3

l (degree)l (degree)(c) regions

0.0

b
-0.2

3-1 2-0
3-2 2-1 1-0
2-2 1-1
2-3 1-2 0-1
1-3 0-2

0.5

0.3

l (degree)

Figure 4.3: The narrow band images in 6.30–6.50 keV (a) and 6.57–6.77 keV (b). They are
the co-added images of the four XIS. The subtraction of the NXB was already done andthe effects of the exposure and vignetting are taken into account. The solid and dashedlines in the panel (a) show the regions for the spectra of M 0.51−0.10 and its background,respectively. The minor and major axes of each elliptical region are 2.30 × 3.10 (source),2.80 × 3.80 (background), respectively. (c) shows the small square regions and their ID
numbers used in section 4.1.3.

58

CHAPTER 4. OTHER 6.4 KEV CLOUDS

Suzaku 6.7 keV image
0.2

CXOGC J174645.3-281546

Arches cluster

G0.61+0.01

b

0.0

3-1
3-2

2-0
2-1

2-2

-0.2

Sgr B2 region

2-3

1-0
1-1

1-2
1-3

Sgr A

0-1
0-2

Sgr B1 region

1.0

0.5
l (degree)

0.0

Figure 4.4: The 6.7 keV band image of the Sgr A to Sgr B2 regions with Suzaku. Excepting
Sgr B1, the image is adopted from Koyama et al. (2007b), Koyama et al. (2007c), Hyodo
et al. (2008), and Mori et al. (2008). The small square regions and the numbers are thesame as figure 4.3d. G 0.61+0.01 is located to the east of the Sgr B1 region (the solid
ellipse; Koyama et al. 2007b). We can think that G 0.61+0.01 is a part of the dashedshell. There are the other bright sources, CXOGC J174645.3–281546 (Muno et al. 2006;
Hyodo et al. 2008) and Arches cluster (Tsujimoto et al. 2007).

Zadeh et al. (2007). But we will call the 6.4 keV cloud M 0.51−0.10 in order to avoid theconfusion with the H II region “Sgr B1” at (l, b) = (0.5◦ , −0.05◦ ) near the 6.4 keV cloud.
We obtained the background-subtracted spectra shown in figure 4.6, in which the sourceand background spectra are extracted from the solid and dashed regions in figure 4.3a,
respectively. We selected the background region referring to the top panel of figure 4.5, sothat the 6.7 keV line fluxes, namely the GCDX, of the two regions are almost the samelevel within the statistical errors. We assume that since the NXB and CXB are the samein the source and background regions, their contribution is successfully subtracted.
We checked a residual contribution of point sources after the background subtraction
with the point source catalog by Muno et al. (2006). There are 14 and 11 point sourcesin the source and background regions, respectively. Those fluxes per a square arcminutein the 2 to 10 keV band are 5.0 × 10−15 erg cm−2 s−1 arcmin−2 in the source region, and6.4×10−15 erg cm−2 s−1 arcmin−2 in the background region. Though the flux in the background region is somehow larger than that in the source region, the difference (1.4 × 10−15
erg cm−2 s−1 arcmin−2 ) is only 3% of the flux of M 0.51−0.10 (5.4 × 10−14 erg cm−2 s−1
arcmin−2 ). Thus, the contribution from the point sources after the background subtraction
can be ignored.

59

10−6 photons s−1cm−2arcmin−2

4.1. SGR B1

2
1
0.5

2−0 1−0
0−1
3−2 2−1
2−2 1−1
3−1 1−2
0−2
2−3 1−3

Sgr B1
Sgr B20.2[H−like Fe Kα]/[He−like Fe Kα]

0.5
0Galactic longitude (degree)

−0.5

0.6
2−1

0.4

0.2

2−3 2−21−2
1−1 1−0
1−3 2−0
3−1
0−1
0−2
3−2

Sgr B1
Sgr B2
0

0.5

0

−0.5

Galactic longitude (degree)
Figure 4.5: Top panel: The 6.7 keV line fluxes of the GC region at l = −0.4◦ ∼ +0.8◦
including the Sgr B1 and B2 regions. The data of the Sgr B1 region obtained in thisobservation are colored with green. Data marked with open circles are adopted from
Koyama et al. (2007c). The data with filled circles correspond to SNR G 0.61+0.01(Koyama et al. 2007b). The ID number of the data correspond to the regions with thesame ID number in figure 4.3c. Bottom panel: The same as in the top panel but for thephoton flux ratios between the 6.7 and 6.9 keV lines.

60

CHAPTER 4. OTHER 6.4 KEV CLOUDS

0.01
10−3
−2 0 2χCounts s−1 keV−1

FIs (XIS0+2+3)

2

5
Energy (keV)

10

Figure 4.6: Background-subtracted FI spectrum extracted from M 0.51−0.10. The sourceand background spectra are extracted from the solid and dashed region in figure 4.3a,
respectively.
The XIS spectra of M 0.51−0.10 contain the prominent 6.4 keV emission line. We
fitted the spectra with the model consisting of an absorbed power-law component plus
two gaussian lines of Fe I Kα (6.4 keV) and Kβ (7.06 keV). We obtained an acceptable
result and show the best-fit parameters in table 4.1. Yusef-Zadeh et al. (2007) reported
the equivalent width of the Fe I Kα line is ∼ 0.6 keV, smaller than our results of ∼ 1.4
keV. Yusef-Zadeh et al. (2007) subtracted only the NXB from the source spectrum. We
examined that the equivalent width of the Fe I Kα line is obtained to be ∼ 0.5 keV fromour spectra when we subtracted only the NXB from the source spectra. We believe thatthe difference between the result of Yusef-Zadeh et al. (2007) and ours is due to thedifferent assumption for the background emission.

4.1.4

Discussion for M 0.51−0.10

The absorption column of 1.5 × 1023 cm−2 is 2 times or more larger than the typicalvalue of 6 × 1022 cm−2 for sources in the GC region. The strong 6.4 keV emission line
implies that a large amount of iron in the neutral state exists at M 0.51−0.10. Thus,
these results suggest M 0.51−0.10 is a local cool dense cloud in the GC region. Indeed, amolecular cloud lies at the same position of M 0.51−0.10, which is a part of the Sgr B1
molecular shell proposed by Sofue (1990) (see figure 4.7).
Two possible scenarios for the the origin of the 6.4 keV emission line have been proposed. One is photo-ionization by X-rays, namely the XRN (X-ray reflection nebula)
scenario. The other is inner shell ionization by collision of low energy cosmic ray electrons. The observed large equivalent width of 1.4 keV and absorption column reaching∼ 1023 cm−2 are consistent with those of the XRN scenario. On the other hand, the4.1. SGR B1

61

Table 4.1: Result of spectral fitting of M 0.51−0.10. The model is an absorbed power-law
plus two gaussian lines.
Model component
Absorption (NH ) (1023 cm−2 )Continuum (power-law):
Photon index (Γ)
Gaussian1 (Fe I Kα):Line center energy (eV)Intensity (10−5 photons cm−2 s−1 )Equivalent Width (keV)
Gaussian2 (Fe I Kβ):Line center energy (eV)†Intensity (10−5 photons cm−2 s−1 )Equivalent Width (keV)
Observed Flux (10−12 ergs cm−2 s−1 )‡Luminosity (1034 ergs s−1 )§χ2 /d.o.f.

value
1.5+0.2−0.1
1.8+0.4−0.5
6402+6−7
2.8+0.2−0.4
1.4+0.3−0.3
7061
0.44+0.17−0.21
0.23+0.22−0.21
1.2+0.1−0.1
1.9+0.1−0.1
50/37

note: The uncertainties indicate the 90% confidence limits.
† Fixed to 1.103 × E(Fe I Kα).
‡ In the 2 to 10 keV band.
§ Absorption corrected in the 2 to 10 keV band.

Figure 4.7: Distribution of the molecular gas in the GC region showed with the contours(13 CO). A structure around at (l, b) ∼ (0.5◦ , 0.0◦ ) is the molecular shell in the Sgr B1
region (Sofue et al. 1990). The arrow indicates the small molecular cloud at (l, b) =(0.5◦ , −0.1◦ ), which may be a counterpart of the 6.4 keV cloud, M 0.51−0.10. The grayscale indicates radio continuum structures (H II region).

62

CHAPTER 4. OTHER 6.4 KEV CLOUDS

Table 4.2: Comparison between required luminosities and observed luminosities
the irradiating source
candidates
1E 1743.1−2843
Sgr A*

distance∗
D(pc)
40
100

photon index
Γ
1.9†
2.0‡

Lobs †(ergs s−1 )3 × 1036
2 × 1033

Lreq §(ergs s−1 )2 × 1038
2 × 1039

note: Those source are assumed to be at the same distance as M 0.51−0.10, 8.5 kpc.
The bright source 1E 1743.1−2843 is a neutron star or black hole LMXB. Sgr A* is asuper-massive black hole at the center of our galaxy.
∗ the distance to M 0.51−0.10.
† the observed photon index and luminosity (Porquet et al. 2003; Baganoff et al. 2003).
‡ The assumed value by Koyama et al. (1996) and Murakami et al. (2000).
§ the required luminosity to account for the observed 6.4 keV flux of M 0.51−0.10. It is
obtained from equation (4.1).

scenario of electron collision expects rather small equivalent width of ∼ 300 eV and an absorption column less than ∼ 1021−22 cm−2 assuming the elemental composition is the sameas the solar system (see table 3.3 in section 3.4.2, Tatischeff 2003). The electron collision
scenario requires the iron over-abundance by a factor of 4 – 5, while the photo-ionization
model does not require the iron over-abundance. Thus, we conclude that M 0.51−0.10 islikely to be an X-ray reflection nebula.
In order to examine the irradiating source, we estimate its luminosity to require theM 0.51−0.10 fluorescence. The cloud absorbs X-rays with the energies higher than 7.1 keVfrom the external source and emits the fluorescent line of 6.4 keV X-rays. The photon fluxof the 6.4 keV line, C6.4keV (photons s−1 cm−2 ), is described as the following equation,
C6.4keV = ²

Ω
4π

Z ∞
7.1keV³´

1 − e−NFe σFe(E) AE −Γ dE.

(4.1)

Here, we assume that the irradiating source has a power-law spectrum (AE −Γ ) and the ironabundance of the molecular cloud of M 0.51−0.10 is solar ([Fe]/[H] = 3 × 10−5 ). Since thetypical absorption column for sources in the GC region is NH = 0.6×1023 cm−2 , the one of
M 0.51−0.10 itself is estimated to be NH = 0.9×1023 cm−2 , i.e. NFe = 2.7×101 8 cm−2 fromthe results of the spectral fitting. Ω = π(r/D)2 is the solid angle covered by M 0.51−0.10
from the view point of the irradiating source, where r is the radius of M 0.51−0.10 and D
is the distance between M 0.51−0.10 and the irradiating source. ² is the fluorescent yield of
0.34 for an iron atom. The K-shell photo-ionization cross section per an iron atom against
an X-ray with the energy above 7.1 keV is σFe = 6.0 × 10−18 (E/1keV)−2.58 (cm2 ) (Henke
et al. 1982). C6.4keV is the observed Fe I Kα line flux of 2.8 × 10−5 photons cm−2 s−1 .
We discuss candidates for the irradiating source by estimating its luminosity with theequation (4.1). One is a bright X-ray source 1E 1743.1−2843, the PSF tail of which isseen at the northwestern edge of the XIS FOV in figure 4.3. 1E 1743.1−2843 is reported
to be a neutron star or a black hole low mass X-ray binary in the GC region, or toward
the GC region but beyond there (Porquet et al. 2003; Del Santo et al. 2006).
In the case that 1E 1743.1−2843 is located in the GC region, the projected distance4.1. SGR B1

63

to the M 0.51−0.10 cloud is D = 40 pc. Since the power-law slope for the X-ray spectrumof 1E 1743.1−2843 is reported to be Γ = 1.9 by Porquet et al. (2003), we obtain theluminosity of 2 × 1038 ergs s−1 required for the 6.4 keV line flux of M 0.51−0.10. However,
1E 1743.1−2843 persistently has the luminosity of (1 − 4) × 1036 ergs s−1 (e.g. Cremonesi
et al 1999; Porquet et al. 2003), which is 2 orders lower than the required luminosity. In
addition, no bursts have been observed from 1E 1743.1−2843 in extensive observations
over the last 20 years. In the case that 1E 1743.1−2843 would be beyond the GC region,
it is too far to account for the 6.4 keV line flux of M 0.51−0.10. Thus, 1E 1743.1−2843 isunlikely for the irradiating source of M 0.51−0.10.
Total luminosity of all the cataloged bright point sources within 50 pc (200 ) from
M 0.51−0.10 is ∼ 3 × 1035 ergs s−1 , which is far lower than the required luminosity of1038−39 ergs s−1 .
Koyama et al. (1996) and Murakami et al. (2000) proposed that the GC supermassive black hole, Sgr A* irradiated molecular clouds in the Sgr B2 region. Sgr A*
is then thought to have been 106 or more times brighter about 300 years ago, the lighttraveling time between Sgr B2 and Sgr A*.
We apply the XRN scenario by past activity of Sgr A* to M 0.51−0.10. The Sgr B
molecular complex, consisting of Sgr B1 (l = 0.5◦ ) and B2 (l = 0.8◦ ), belongs to the“Galactic-Center molecular Arm” surrounding Sgr A* (Sofue 1995). Sgr B1 and Sgr B2
are located at the same three dimensional distance from Sgr A* (∼ 100 pc) though the
projected distances are different. The photon index for Sgr A* is reported to be Γ = 2.7
by Baganoff et al. (2003), which is not consistent with that of M 0.51−0.10 (Γ = 1.3−2.2).
However, Baganoff et al. (2001) reported that the photon index became hard (Γ ∼ 1.0)
when Sgr A* is in flare-up state.
Koyama et al. (1996) and Murakami et al. (2000) assumed that the photon index
is 2.0 for the power-law spectrum of Sgr A*, and reported that the required luminosity
for Sgr B2 cloud is ∼ 3 × 1039 ergs s−1 . We also estimated the required luminosity forM 0.51−0.10 to be ∼ 2 × 1039 ergs s−1 . Since we do not take the limited duration timeof the X-ray irradiation from the past Sgr A* flare into account in this calculation, it isconsistent with the case in Sgr B2.

4.1.5

Summary

In this section, we analyzed the data of the Sgr B1 region observed with Suzaku withthe long exposure ∼100 ks.
• The line flux ratio [Fe XXVI]/[Fe XXV] of the Sgr B1 and B2 regions (l = +0.3◦ −− + 0.8◦ ) is consistent with that of the region at l = −0.4◦ − − + 0.2◦ . The Fe XXVline flux exponentially decreases from the GC center to the Sgr B1 and B2 regions.
This results suggest that the GCDX extends at least up to the Sgr B1 and B2 regionswith a constant temperature of kT =6–7 keV.
• M 0.51−0.10 is identified with the 6.4 keV cloud “Sgr B1” detected by Yusef-Zadeh
et al. (2007). The X-ray spectra of M 0.51−0.10 exhibit a absorption column ofNH ∼ 1.5 × 1023 cm−2 and a large equivalent width of ∼ 1.4 keV for Fe I Kα

64

CHAPTER 4. OTHER 6.4 KEV CLOUDS

fluorescence emission line. It suggests that the 6.4 keV X-ray emission is unlikely
due to collision of low energy cosmic ray electrons but due to reflection of external
hard X-rays, and hence M 0.51−0.10 is thought to be an X-ray reflection nebula. The
bright X-ray point source 1E 1743.1−2843 or the collection of fainter point sources
near M 0.51−0.10 does not explain the luminosity of ≥ 1038−39 ergs s−1 required forthe 6.4 keV X-rays from M 0.51−0.10. The XRN scenario by past activity of Sgr A*,
which was successfully applied to Sgr B2, may also be applied to M 0.51−0.10.

4.2. SGR C

4.2

Sgr C

4.2.1

Introduction for Sgr C

65

The Sgr C region located at l ∼ −359.5◦ consists of giant molecular clouds, large H II
regions, and radio non-thermal filaments (NTF). From the 13 CO observations, Liszt andSpiker (1995) estimated 6.1 × 105 solar mass for the molecular gas.
Murakami et al. (2001b) discovered a 6.4 keV cloud with ASCA, and interpreted thatthe cloud is an X-ray reflection nebula (XRN) because the X-ray spectrum exhibited alarge equivalent width (EW ) of the 6.4 keV emission line with a strong absorption like thatin Sgr B2. On the other hand, Yusef-Zadeh et al. (2007) discussed possible associationof the 6.4 keV emission lines with some radio NTFs based on the Chandra observation,
and argued that the 6.4 keV line is due to the impact of low energy cosmic-ray electrons(LECRe).
These results and arguments are based on the short exposure observations (20 ks and
22 ks by the ASCA and Chandra, respectively). In order to fix the above debates, we,therefore, made a long Suzaku observation on the Sgr C region.

4.2.2

Observation and Data Reduction

The Sgr C region was observed with the Suzaku/XIS from 2006 February 20 to 23. The
telescope optical axis position was R.A.= 17h 44m 37.30s , Decl.= −29◦ 280 10.200 (J2000.0).
A cleaned event list was obtained from the processed data with the version of 2.0.6.132
by removing events taken during the passage of the South Atlantic Anomaly, the elevation angles from the night Earth rim of < 5◦ and from the sun-lit Earth rim of < 20◦ ,and the telemetry saturation. After these data screenings, the net exposure was 107 ks.
We analyzed the screened data using the HEADAS software version 6.4, XSPEC version
11.3.2g3 . We utilized the calibration databases released on 2008 February 1 4 .
We also made COR-sorted NXB data sets using the night-Earth data released by theSuzaku XIS team 5 . For the following imaging and spectral studies in the Sgr C region(section 4.2.3 and 4.2.3), we used the X-ray data after subtracting the NXB that werecompiled to have the same COR distribution as that during the Sgr C observation.
To increase statistics, hereafter, we co-added the four XIS data for imaging study, inwhich exposure and vignetting correction were made. For the spectral study, we co-added
the three FI CCD data, and treated BI data separately.

4.2.3

Analysis and Results

X-ray Image in the Emission Line
Figure 4.8 shows the energy spectrum of the whole Sgr C region, which was extracted
from the circular region with a radius of 6.950 from the FOV center. The spectrum shows2

http://www.astro.isas.jaxa.jp/suzaku/process/
http://heasarc.gsfc.nasa.gov/docs/software/lheasoft/4
http://www.astro.isas.jaxa.jp/suzaku/caldb/
5
http://www.astro.isas.jaxa.jp/suzaku/analysis/xis/nte/3

66

CHAPTER 4. OTHER 6.4 KEV CLOUDS

normalized counts/sec/keV

1

Ar
0.1

S

Ca
Fe

0.01

10−3

2

5
channel energy (keV)

10

Figure 4.8: The co-added GCDX spectrum of the FIs (XIS0+2+3) in the Sgr C region
extracted from the 6.950 radius circle in the center of the XIS FOV. The NXB component
was subtracted. The spectrum of BI (XIS1) was simultaneously analyzed, but the figure
is not shown here, for brevity.
K-shell emission lines from He-like and/or H-like ions of S, Ar, Ca, and Fe. In addition,
the 6.4 keV emission line and the K-edge absorption at 7.1 keV due to Fe I are detected.
The three Fe K-shell emission lines at 6.4, 6.7 and 6.9 keV in the Galactic center diffuse
X-ray emission (GCDX: Koyama et al. 1989; Yamauchi et al. 1990; Koyama et al. 2007c)
are clearly resolved.
Since the Suzaku coordinate has 90% confident error of ∼ 20 00 (Uchiyama et al. 2008),
we made a coordinate correction using a Chandra observation. In the energy band of
0.7–1.5 keV, a point source with the flux of ∼ 6 × 10−15 erg cm−2 s−1 was detected at(R.A., Decl.)J2000.0 =(17h 44m 52.1s , −29◦ 310 56.700 ) in the XIS image. In the error circle ofSuzaku, there was only one Chandra point source at (R.A., Decl.)J2000.0 = (17h 44m 53.1s ,−29◦ 310 45.500 ). Since the flux was almost the same as that of the Suzaku source, we safely
concluded these point sources are identical. We, then fine-tuned the Suzaku coordinate
by shifting ∆(R.A., Decl.)=(13.700 , 11.200 ).
Referring figure 4.8, we made the narrow band image of the 2.45 keV-line (2.35–
2.50 keV; Kα of S XV) and the 6.4 keV-line (6.28–6.42 keV; Kα of Fe I). The images after
the smoothing with the Gaussian kernel of 48 pixels (0.830 ) are shown in figure 4.9. In
figure 4.9a, a bright 2.45 keV-line clump is found at (l, b) = (359.407◦ , −0.119◦ ), hence
designated as G 359.41−0.12. Likely, four 6.4 keV-line clumps found in figure 4.9b aredesignated as M 359.43−0.07, M 359.47−0.15, M 359.43−0.12, and M 359.38−0.00, where
M 359.43−0.12 is in the region of G 359.41−0.12.
The Emission Line Feature in the Sgr C Region
The photon statistics are limited to make measurable spectra for M 359.38−0.00 and
M 359.43−0.12. Furthermore, M 359.43−0.12 is in the 2.45 keV-source, G 359.41−0.12.
Therefore we concentrated on the two bright sources, M 359.43−0.07 and M 359.47−0.15,

4.2. SGR C

67(b)6.28-6.42keV(a)2.35-2.50keV
05:00.0

05:00.0

0:00:00.0

0:00:00.0

05:00.0-0:10:00.0

G359.41-0.12

M359.43-0.07
05:00.0

M359.43-0.12-0:10:00.0

15:00.0

15:00.0

40:00.0

Galactic latitude

Galactic latitude

M359.38-0.00

359:30:00.0
20:00.0
Galactic longitude

40:00.0

M359.47-0.15

359:30:00.0
20:00.0
Galactic longitude

Figure 4.9: (a) The 2.45 keV-line (S XV Kα) image in the energy band of 2.35–2.50 keV
and (b) the 6.4 keV-line (Fe I Kα) image in the 6.28–6.42 keV band. The source andbackground regions for the spectral analysis are shown with the solid and dashed magenta
ellipses, where the data within the blue ellipse (G 359.41−0.12) were excluded to minimize
the contamination from the strong 2.45 keV emission line and its associated continuum.
and the surrounding background regions.
As is seen in figure 4.8, the GCDX with strong 6.7 keV (Fe XXV Kα) and 6.9 keV (FeXXVI Lyα) lines is the largest background for the local enhancements such as the 6.4 keVclumps. In addition, the flux of the GCDX is not uniform but variable from position toposition. To see the flux distribution and nature of the GCDX near Sgr C, we examined
the NXB-subtracted (GCDX is not subtracted) spectra in the 6.4-keV clumps and thesurrounding areas including candidate background regions (see figure 4.9).
The spectra from the whole Sgr C region is already made in section 4.2.3 (figure 4.8).
The spectra of M 359.43−0.07 and M 359.47−0.15 are extracted from the regions shownwith the solid magenta ellipses in figure 4.9b. The background spectra for these 6.4 keVclumps were obtained from the annuli between the inner solid and outer dashed ellipses.
Since G 359.41−0.12, the solid blue ellipse in figure 4.9a, may have strong 2.45 keV-linewith an associated continuum emission, we excluded G 359.41−0.12 from the spectra of
the 6.4-keV sources and their backgrounds. As the surrounding area, we made additional
spectrum from a circle at the FOV center with a radius of 7.80 but excluding the above
cited source and background regions (hereafter, the outer region).
The spectral fittings were made in the same energy band with the same phenomenological model and free parameters as those used in Koyama et al. (2007c; 2009): a single
power-law and 10 Gaussian emission lines (table 4.3). The spectra of FIs and BI were
simultaneously fitted, and acceptable χ2 were obtained from all the spectra. The FI spectra and the best-fit model of M 359.43−0.07 and M 359.47−0.15, and the outer region of
Sgr C are shown in figure 4.10 as examples. The best-fit parameters are listed in table 4.4.

CHAPTER 4. OTHER 6.4 KEV CLOUDS(a)M359.43−0.07
0.01

10−3

10−4

normalized counts/sec/keV

normalized counts/sec/keV

68

10−3

10−4
2

2
0

χ

χ

(b)M359.47−0.15
0.01

0
−2

−2

−4
5

10

5

10

channel energy (keV)channel energy (keV)χ

normalized counts/sec/keV1(c)the outer region
0.1

0.01

10−3
2
0
−2
−4
5

10
channel energy (keV)

Figure 4.10: Same as figure 4.8, but for M 359.43−0.07 (a), M 359.47−0.15 (b), and theouter region (c). The phenomenological model with single power-law, 10 Gaussian lines,
and the cosmic X-ray background component is applied.

Table 4.3: Center energies of the Fe and Ni lines from the GCDX obtained by Suzaku(Koyama et al. 2007c).
line
Fe I KαFe XXV Kα
Fe XXVI Kα
Fe I Kβ
Ni I Kβ
Ni XXVII Kα
Fe XXV Kβ
Fe XXVI Kβ
Fe XXV Kγ
Fe XXVI Kγ

observed energy (eV)6409+1−16680+1−1
6969+6−37069 = 1.103 × E(Fe I Kα)
7490+12−147781+24−31
7891 = 110 + E(Ni XXVII Kα)
8220+31−228264 = 44 + E(Fe XXVI Kβ)
8681+33−32

Area(arcmin2 )
152
121
7.09
11.7
6.91
22.4

1.72(1.65 − 1.78)
1.70(1.63 − 1.82)
1.82(1.65 − 2.06)
1.64(1.47 − 1.93)
1.96(1.59 − 2.24)
1.75(1.59 − 1.97)

ΓF6.4
0.79(0.77 − 0.81)
0.55(0.53 − 0.57)
1.91(1.79 − 2.06)
1.00(0.92 − 1.08)
2.16(2.03 − 2.31)
1.02(0.97 − 1.08)

F6.7
0.68(0.65 − 0.70)
0.68(0.66 − 0.71)
1.08(0.96 − 1.22)
0.93(0.86 − 1.03)
0.63(0.53 − 0.73)
0.67(0.62 − 0.72)

Line Flux†Line Equivalent width (eV)
EW6.4
EW6.7
458(448 − 470) 420(401 − 430)
346(333 − 363) 454(441 − 471)
670(625 − 720) 400(357 − 454)
397(365 − 428) 388(360 − 431)
966(909 − 1030) 298(253 − 348)
587(554 − 620) 404(374 − 437)

(1) the whole Sgr C region, (2) the outer region, (3) M 359.43−0.07, (4) background region for M 359.43−0.07, (5) M 359.47−0.15,(6) background region for M 359.47−0.15.
∗ 90% confidence limits are in parentheses.
† Absorption-uncorrected line flux in the unit of 10−6 ph cm−2 s−1 arcmin−2 .

(1)
(2)
(3)
(4)
(5)
(6)

RegionTable 4.4: Best-fit parameters for the NXB-subtracted spectra∗.

4.2. SGR C
69

70

CHAPTER 4. OTHER 6.4 KEV CLOUDS

Koyama et al. (2009) found that the GCDX in the Sgr A region is phenomenologicallydecomposed into the 6.7-keV line plus associated continuum (6.7-component) and the 6.4keV line plus associated continuum (6.4-component); the EWs of the 6.4 keV (EW6.4 )and the 6.7 keV (EW6.7 ) lines are given by the relation EW6.7 + 0.50(±0.06) × EW6.4 =0.62(±0.07)[keV]. Then, in the limit of EW6.4 → 0, EW6.7 was estimated to be 0.62 ±0.07 keV. On the other hand, in the limit of EW6.7 → 0, EW6.4 is 1.2 ± 0.2 keV. They
also made the GCDX-subtracted spectra of the two 6.4 keV clumps in the Sgr A region(source 1 and 2 in Koyama et al. 2009) and obtained consistent EW6.4 values as thoseestimated from the phenomenological relation.
Stimulated by the successful approach of Koyama et al. (2009), we used the same
method to decompose the GCDX in the Sgr C region into the 6.4-component and 6.7component. Using the data in table 4.4, we plot the EW6.4 and EW6.7 relation of the six
regions (figure 4.11). The outer region, the whole Sgr C and the background region forM 359.43−0.07 follow to, but the M 359.43−0.07 and M 359.47−0.15 regions come above
the best-fit correlation line of the Sgr A region; these are systematically larger EW6.4 thanthe Sgr A region. The best-fit linear relations between EW6.7 and EW6.4 for all cited
regions except the whole Sgr C is;
EW6.7 + 0.22(±0.12) × EW6.4 = 0.53(±0.06)[keV],(4.2)

where the errors mean 90% confidence levels.
This relation indicates that EW6.7 in the 6.7-component is 0.53±0.06 keV (at EW6.4 →0) for the Sgr C region, consistent with that of the Sgr A region of 0.62 ± 0.07 keV(Koyama et al. 2009). On the other hand, EW6.4 in the 6.4-component is 2.4+2.3−0.7 keV (at
EW6.7 → 0), larger than those in the Sgr A region (Koyama et al. 2009).
Spectra of M 359.43−0.07 and M 359.47−0.15
From table 4.4, we see that the 6.7 keV line flux in M 359.43−0.07 and M 359.47−0.15
are almost identical to those of the respective background regions. The 6.7 keV flux
differences are 16% and 6% of those in M 359.43−0.07 and M 359.47−0.15, respectively.
Since the 6.7 keV line flux is a good indicator of the flux of the GCDX (Koyama etal. 2007b; 2009), we conclude that the background regions are properly selected withinpossible systematic error of 6–16%, which is smaller than the statistical errors.
Therefore we made the GCDX-subtracted spectra of these 6.4 keV clumps (figure 4.12).
The energy-dependent vignetting were corrected by multiplying the effective-area ratio of
the source region to corresponding background region for each energy bin of the GCDX
spectra.
We first applied a model of a power-law and a narrow Gaussian line at 6.4 keV withabsorption. The metal abundances of the inter stellar medium were fixed to the solar
value (Anders & Grevesse 1989), while the center energy of the 6.4 keV line and its flux
were free parameters. We simultaneously fit this model to the FIs and BI spectra, andfound a line-like residual at 7.0–7.1 keV, with the flux of ∼10–20% of the 6.4 keV (Fe I
Kα) line. From the energy and flux, the residual is likely to be the Kβ of neutral Fe. We
then added second narrow Gaussian for the Fe I Kβ line, fixing its center energy to 1.103

4.2. SGR C

71

EW(6.7keV) (eV)

600

400

200

0

0

500

1000
1500
EW(6.4keV) (eV)

2000

2500

Figure 4.11: The correlation between the EWs of the Fe I-Kα line (EW6.4 ) and Fe XXV-Kαline (EW6.7 ). The dashed line shows the best fit relation in the Sgr A region (adopted fromfigure 2d of Koyama et al. 2009). The open and filled circle show the whole Sgr C region
and the outer region, respectively. The source regions of M 359.43−0.07 and M 359.47−0.15
are plotted with the open square and triangle, while their background regions are givenby the filled marks. The solid line is the best-fit relation for all regions except the whole
Sgr C of EW6.7 + 0.22 × EW6.4 = 0.53 [keV].
times of that of Kα. This model was accepted as is shown in figure 4.12, while the best-fit
parameters are listed in table 4.5.

4.2.4

Discussion for the 6.4 keV Clouds in Sgr C

We detected four 6.4 keV clumps, M 359.43−0.07, M 359.47−0.15, M 359.43−0.12,
and M 359.38−0.00 with Suzaku. From the former bright clumps, M 359.43−0.07 andM 359.47−0.15, we found Kα and Kβ-lines from Fe I. The most important fact is thatEWs of the Kα-lines are extremely large as ∼2 keV (figure 4.12, table 4.5). These large
EWs are independently supported by the analysis of the GCDX in the Sgr C regions (figure 4.11), and hence are highly reliable. We note that Yusef-Zadeh et al. (2007) reported
small EW6.4 from all the clumps, but their estimation was based on the spectra where no
GCDX is subtracted.
Comparison with the Past X-Ray Observations
Murakami et al. (2001b) found excess emission at (l, b) = (359.43◦ , −0.04◦ ) in the
X-ray image of the 5.8–7.0 keV band using the ASCA satellite. Although the energy
resolution was limited, an emission line at the energy consistent with Fe I-Kα line wasfound. Hence Murakami et al. (2001b) regarded this clump as an XRN.

CHAPTER 4. OTHER 6.4 KEV CLOUDS

normalized counts/sec/keV

72

0.01

(a)M359.43−0.07

10−3

10−4

χ

2
0
−2

normalized counts/sec/keV

5channel energy (keV)

0.01

10

(b)M359.47−0.15

10−3

10−4

χ

2
0
−2
5
channel energy (keV)

10

Figure 4.12: Same as figure 4.10, but for GCDX-subtracted spectra of M 359.43−0.07 (a)and M 359.47−0.15 (b). The Kα and Kβ emission lines of Fe I and power-law are shownby dot-dashed, dotted, and dashed lines, respectively.

The Suzaku source M 359.43−0.07 is located in the ASCA clump, but the peak position is systematically shifted. For comparison, we extracted the Suzaku spectrum fromthe same source and background regions as those of the ASCA clump (Murakami etal. 2001b). The background subtracted spectrum has the 6.4 keV line flux of 1.8+0.3−0.2 ×−5−2−1
10 ph cm s (absorption uncorrected), which agrees with that of the ASCA clump−5 ph cm−2 s−1 ). With Suzaku, we found that the surface brightness ra(3.5+1.4−2.2 × 10
tio of the 6.4 keV emission lines between the M 359.43−0.07, the ASCA clump and theouter background region are approximately 3 : 2 : 1 (1.9, 1.4, and 0.6 in the unit of
10−6 ph cm−2 s−1 arcmin−2 ). Hence no significant detection of the full region of the
ASCA clump in figure 4.9b may be due to its relatively lower surface brightness thanM 359.43−0.07. On the other hand, no ASCA detection of M 359.43−0.07 is puzzling,although possible time variability can not be excluded. ASCA also detected no flux from

4.2. SGR C

73

Table 4.5: The best-fit parameters of the spectral fittings to the background (GCDX)
subtracted spectra of M 359.43−0.07 and M 359.47−0.15.
Parameter
Absorbed power-law model:
Column density NH (1023 cm−2 )Photon index
Gaussian 1 (Fe I Kα):Line energy (keV)Line flux (10−6 ph cm−2 s−1 ) ∗Equivalent width (keV)
Gaussian 2 (Fe I Kβ):Line energy (keV)†Line flux (10−6 ph cm−2 s−1 ) ∗Flux (10−13 erg cm−2 s−1 )‡χ2 /d.o.f.

M 359.43−0.07

M 359.47−0.15

0.92(0.48 − 1.41)
1.67(1.51 − 1.82)

0.82(0.65 − 1.18)
1.61(1.52 − 1.86)

6.41(6.39 − 6.42)
6.43(5.27 − 7.39)
2.18(1.79 − 2.51)

6.41(6.40 − 6.42)
8.82(7.87 − 10.0)
1.96(1.75 − 2.23)

7.07
1.05(0.43 − 2.14)
2.70
27.92/41

7.07
1.91(1.22 − 3.00)
4.11
49.67/46

90% confidence limits are in parentheses.
∗ Line fluxes are not corrected for the absorption.
† The line center is fixed so that the energy ratio is 1.103 (Kβ/Kα=7.06keV/6.40keV).
‡ Observed flux of the diffuse emission in the energy band of 3–10 keV. Absorption
is not corrected.

the other bright Suzaku source M 359.47−0.15. However, this source is near the edge of
the ASCA field.
The Suzaku spectrum of the whole Sgr C region (figure 4.8) shows the EWs of 458+12−10 eVand 420+10
eV
for
the
6.4
keV
and
6.7
keV
emissionlines,
respectively.
The
surfacebright−19
ness in the 5–8 keV energy band is 7.1 × 10−14 erg cm−2 s−1 arcmin−2 . The Chandra
analysis of the same region, in which point sources cataloged in the list of Muno et al.
(2006) are excluded, shows that the EWs of the 6.4 keV and 6.7 keV emission lines are
460 ± 100 eV and ∼ 400 eV, respectively. The surface brightness in the 5–8 keV band is∼ 7.5 × 10−14 erg cm−2 s−1 arcmin−2 (figure 1b in Yusef-Zadeh et al. 2007). These arevery similar to those of Suzaku, and hence point source contribution with the flux level
larger than 1 × 10−13 erg cm−2 s−1 (Muno et al. 2006) may not be large, at least in thewhole Sgr C region.
Detailed comparison in the smaller scale emission, however, shows significant differences. M 359.47−0.15 and a part of M 359.43−0.07 can be seen in the narrow band image
at 6.4 keV with Chandra (left panel in figure 4.13; Yusef-Zadeh et al. 2007). However,M 359.43−0.07 is not found in the EW map of Chandra (right panel in figure 4.13; YusefZadeh et al. 2007). This is a puzzle, because the Suzaku EW of M 359.43−0.07 is one of
the largest among cited regions in the NXB-subtracted spectra. One possibility is a time
variability as was found in Sgr B2 and Radio Arc regions (chapter 5 in this thesis; Muno etal. 2007; Koyama et al. 2008; Inui et al. 2009). In the Chandra image of the 2–6 keV band
(figure 1a of Yusef-Zadeh et al. 2007), a clump named as G 359.45−0.07 is found. How-

74

CHAPTER 4. OTHER 6.4 KEV CLOUDS

Figure 4.13: Left panel: Contours of velocity-integrated 13 CO emission between −77.2
and −22 km s−1 superimposed on a gray-scale 6.4 keV Kα line emission with a 3000
resolution taken by the Chandra satellite (figure 5a in Yusef-Zadeh et al. 2007). Rightpanel: Contours of X-ray continuum emission superimposed on a gray-scale equivalent
width map of 6.4 keV line emission.
ever Suzaku found no excess in this region with the 6.4 keV line, and hence G 359.45−0.07
would not be 6.4 keV line emitter. Conversely no excess from M 359.38−0.00 is found witheither Chandra or ASCA, possibly due to limited flux, and hence this is a new 6.4 keVclump found with the Suzaku satellite.
Ionization Mechanism
Two scenarios have been proposed for the origin of the 6.4 keV emission line from
MCs in the GC region. One is X-ray photo ionization by external X-ray sources (the XRN
scenario: Koyama et al. 1996; Sunyaev & Churazov 1998; Park et al. 2004). The other isthe inner shell ionization by the impact of LECRe (the electron bombardment scenario:
Valinia et al. 2000; Yusef-Zadeh et al. 2002). The LECRe scenario expects EW6.4 of∼ 300 eV for the solar abundance Fe (table 3.3, see also Tatischeff 2003; Yusef-Zadeh etal. 2007), and hence is difficult to explain the observed large EW6.4 of ∼ 2 keV, unless
Fe is extremely over-abundant by a factor of 6–7. We found no observational evidence
for such over-abundance. On the other hand, XRN scenario naturally explain the large
EW6.4 of the M 359.43−0.07 and M 359.47−0.15, and hence is more likely.
Photo-Ionization Source
Since the typical absorption toward GC is ∼ 0.6 × 1023 H cm−2 (Rieke et al. 1989;
Sakano et al. 2002), intrinsic absorption depth of M 359.43−0.07 and M 359.47−0.15 are≤ 0.8 × 1023 H cm−2 and ≤ 0.6 × 1023 H cm−2 , respectively. The areas of M 359.43−0.07
and M 359.47−0.15 shown in figure 4.9b are 7 arcmin2 , then their masses are estimatedto be ≤ 4 × 104 , ≤ 3 × 104 solar mass, respectively. From these masses and fluxes of the4.2. SGR C

75

05:00.0

Galactic latitude

0:00:00.0

05:00.0-0:10:00.0

15:00.0

40:00.0

359:30:00.0
20:00.0
Galactic longitude

Figure 4.14: The contours of channel map of CS J = 1 − 0 emission line overlaid on the
narrow band image at 6.4 keV shown in figure 4.9b. The first contour level and contour
interval are 14.3 and 4.8 K km s−1 , respectively. Emission in the velocity ranges of −130 to−90 km s−1 is shown with green and that of −80 to −30 km s−1 is magenta, respectively.
6.4 keV lines, the luminosity of photo-ionizing source can be estimated following Sunyaev
and Churazov (1998) as ≥ 1 × 1038 erg s−1 and ≥ 2 × 1038 erg s−1 .
There is no X-ray object bright enough inside or nearby Sgr C region. Even thebrightest near-by sources, KS 1741−293 and the Great Annihilator (1E 1740.7–2943) are
impossible to explain this luminosity. Thus we arrive to the same conclusion for the
6.4 keV clumps origin near Sgr B2, a past Sgr A∗ activity of ≥ 1 × 1038 erg s−1 at 240 yr
ago.
Comparison with the Molecular Line Observations
The absorption column derived from the background-subtracted spectra of M 359.43−0.07
and M 359.47−0.15 are NH ∼ 1 × 1023 H cm−2 . We require similar amounts of column in
the MCs that corresponds to M 359.38−0.00 and M 359.43−0.12. Referring the CS molecular line map over the velocity range of VLSR = −200 to +200 km s−1 (Tsuboi et al. 1999,
and the electric data), we searched for possible counterpart of the 6.4 keV clumps withNH ≥ 1023 H cm−2 (figure 4.14).
M 359.43−0.07 : MCs of NH ∼ 1023 H cm−2 are seen both in the velocity ranges of −50
to −70 km s−1 (also Yusef-Zadeh et al. 2007) and −90 to −110 km s−1 (also Murakami
et al. 2001b).

76

CHAPTER 4. OTHER 6.4 KEV CLOUDS

M 359.47−0.15 : An MC with NH ≥ 1023 H cm−2 in the velocity range of −50 to−70 km s−1 exists in this region. M 359.47−0.15 is located in the western edge of−65 km s−1 MC as is also noticed by Yusef-Zadeh et al. (2007). There is no other
MC of ∼ 1023 H cm−2 in the other velocity range.
M 359.38−0.00 : No MC is seen in the −50 to −70 km s−1 band but an MC in thevelocity range of −110 to −130 km s−1 is associated to the 6.4-keV clump; The column
density of this radio cloud is roughly half of those in the directions of M 359.43−0.07 andM 359.47−0.15.
M 359.43−0.12 : There is a weaker peak in the −130 to −120 km s−1 band than thatat M 359.38−0.00.
Sofue (1995) and Sawada et al. (2004) proposed that molecular gases in the GC
region consists of two arms, Arm I and Arm II. The −130 to −90 km s−1 MC is physicallyassociated with Arm I, while the −80 to −30 km s−1 MC is a part of Arm II. M 359.38−0.00
and M 359.47−0.15 belong to Arm I and Arm II, respectively, while M 359.43−0.07 has
both possibilities. Since Arm II is located far side of Arm I (Sofue 1995; Sawada et al.
2004), X-rays from Sgr A∗ of several hundred years ago arrived at M 359.47−0.15 probably
earlier than M 359.38−0.00.

4.2.5

Summary

1. We found four diffuse 6.4 keV clumps, M 359.43−0.07, M 359.47−0.15, M 359.43−0.12
and M 359.38−0.00 in the Sgr C region. The last one, M 359.38−0.00 is newly discovered.
2. The spectra of the two bright clumps, M 359.43−0.07 and M 359.47−0.15, havepower-law of photon index of 1.6–1.7 with a very large EW6.4 of 2.0–2.2 keV, thelargest EW6.4 among the 6.4 keV clumps in the GC.
3. The large EW6.4 supports that the origin of 6.4 keV clumps is due to X-ray reflection
irradiated by hard X-ray sources.
4. M 359.38−0.00 and M 359.47−0.15 possibly associate with MCs in different velocity
ranges, indicating these are located in the different molecular arms.

Chapter 5

Discovery of Correlated
Time-Variability between the
Fe I K and Hard X-rays from
Sgr B20.5

Sgr B2 (Chapter 5)0.0Arches cluster

M359.43-0.07
Sgr A(Chapter 3) M359.47-0.15

M0.74-0.09

Sgr A*
Sgr B1(Section. 4.1)

Sgr C (Section 4.2)

-0.5
1.0

0.5

0.0

359.5

359.0

Figure 5.1: 6.4 keV image in the GC region, same as figure 3.1. In this chapter, we focus
on two bright clumps in the Sgr B2 region indicated with the yellow boxed region.

5.1

Introduction

In the previous chapters 3 and 4, the external X-ray origin for the Fe I K line in the GC
region was established by the large equivalent widths of the K-shell lines. Some molecular
clouds ”reflects” the X-rays from an external bright object (X-ray reflection nebula; XRN).
Accordingly, we call X-rays from the MCs as ”X-ray echo”.
The more serious unresolved problem is ”What and where the external X-ray source77

78

CHAPTER 5. TIME-VARIABILITY OF X-RAYS FROM SGR B2

is ?”. One proposed scenario is that Sgr A*, the super-massive black hole with around 4
million solar masses (Ghez et al. 2008) in the GC located at a several hundred light-years
distance from the molecular clouds (MCs) (Koyama et al. 1996; Murakami et al. 2000),
was bright several hundred years ago with the X-ray luminosity millions times higher thanthe present one of 1033 –1034 erg s−1 (Baganoff et al. 2001).
Although the required luminosity of the external X-ray source depends on the distanceto the MCs, the duration time of irradiating X-rays, and size and density of the MCs,
current data were unable to exclude a possibility of intense X-ray flares of near-by transient
sources. Also no clear evidence that Sgr A* was so bright in the past has been presented.
Recently rapid time-variability of Fe I-K was discovered from some molecular clouds in
the Sgr A and B2 regions (Muno et al. 2007; Koyama et al. 2008; Inui et al. 2009; Ponti
et al. 2010). However, those observations were not enough to distinguish the verify theorigin because the positions of the clouds are unknown and/or systematical uncertainties
due to a combination of different instruments remain.
In order to identify the external X-ray source for the ”X-ray echo”, Sgr B2 has uniqueproperties as, (i) Sgr B2 is one of the brightest 6.4 keV clouds, (ii) the 3-dimensional
position is well determined (Ryu et al. 2009), and (iii) the Galactic center diffuse X-ray
emission (GCDX: Koyama et al. 1996) is relatively small and uniform, and hence possible
contamination from the GCDX can be reliably subtracted.
This chapter provides clear evidence for the scenario of a past big-flare of Sgr A*. Errors
in this chapter are estimated at the 90% confidence level unless otherwise mentioned.

5.2

Observation and Data Reduction

We have observed the Sgr B region with the Suzaku/XIS. The observations were madetwo times in October 2005 and September 2009 at the identical pointing positions within
400 . Sequence numbers for the two observations are 100037060 and 504004020. The effective
exposures in 2005 and 2009 were 21.3 hours and 56.1 hours, respectively. Four XIS (XIS0,
1, 2, and 3) and three XIS (XIS0, 1, and 3) cameras were available in 2005 and 2009,
respectively. We discard the data obtained in a one fourth field (segment A) of XIS0 for
both observations, which has become nonfunctional since June 2008.
Although the effective energy range of XIS is 0.2–10 keV, we use the 4–10 keV banddata for the spectral analyses, because soft X-rays from the target (Sgr B) are absorbed by
the interstellar medium toward the GC region. We construct the non-X-ray background(NXB) data caused by the cosmic-rays using xisnxbgen (Tawa et al. 2008) and subtracted
the NXB from the observed data. We then corrected the local differences of the effective
area for the data analyses (the vignetting correction).

5.3

Analysis

5.3.1

X-ray Images

Figure 5.2a shows X-ray images in the Fe I-Kα (6.4 keV) line of the Sgr B region in
2005, where bright diffuse sources are clearly detected (Koyama et al. 2007b). Following

5.3. ANALYSIS

0.1

79

(a)
Sgr B2

0.1

FeI Kα2005

0.0

FeI Kα
2009

(b)

0.0

b

b
-0.1

-0.1
M0.74-0.09

-0.2

-0.2

Background
0.8

0.6

0.7

0.5

0.8

0.1

0.1

0.0

Hard continuum
2009

(d)

0.0

b

b
-0.1

-0.1

-0.2

-0.2
0.8

0.6

0.7

0.5

0.8

l
0

0.5

l

Hard continuum
2005

(c)

0.6

0.7

l

0.6

0.7

0.5

l
0.5

1.0

Figure 5.2: X-ray images near Sgr B2 and M 0.74−0.09 in the field of view of 17.80 ×13.40 .
The vertical and horizontal axes are the Galactic longitude (l) and latitude (b) in degree(◦ ), respectively. (a): The Fe I-Kα (6.4 keV) band images in 2005. The bin size is 8.300
and is smoothed with a Gaussian kernel of 1σ=4200 . Profiles of the 6.4 keV line intensity
in the Sgr B2 and M 0.74−0.09 are taken from the yellow boxes (see text and figure 5.3).
(b): The same as (a) but in 2009. (c): The hard continuum (8-10 keV) band images in2005. The bin size is 33.300 and is smoothed with a Gaussian kernel of 1σ =10000 . (d)The same as (c) but in 2009. The source regions for the Sgr B2 and M 0.74−0.09 spectra
are shown in the white solid circles, centered at (l, b) = (0.6676◦ , −0.0226◦ ) with radiusof 1.70 and (l, b) = (0.7394◦ , −0.087◦ ) with radius =1.50 , respectively. The backgroundregion is shown in the white dashed ellipse centered at (l, b) = (0.6410◦ , −0.1523◦ ) withthe major and minor radii of 5.50 and 3.50 , respectively. The inside white dashed circle at(l, b) = (0.6885◦ , −0.1523◦ ) with radius = 1.50 is excluded for the background region (seetext). In this figure, 0.1◦ (60 ) corresponds to 50 light-years. The color scale is arbitrarily
adjusted for comparison.

80

CHAPTER 5. TIME-VARIABILITY OF X-RAYS FROM SGR B2(a) Sgr B2(b) M0.74−0.09
2005
2009

0.5

2005
2009

0.5

Ћ

Ћ

0

1
Intensity (A.U.)

Intensity (A.U.)

1

0.65
0.7
Galactic Longitude (degree)

0

0.7
0.72
0.74
0.76
Galactic Longitude (degree)

Figure 5.3: (a): Projected profile of the 6.4 keV line intensity in the Sgr B2 cloud (l =0.63◦ –0.72◦ ) extracted from the yellow box in figure 5.2. The integral interval is b =−0.05◦ –0.00◦ . The peak positions l = 0.6610◦ ± 0.0027◦ in 2005 and l = 0.6692◦ ± 0.0056◦
in 2009 are indicated with the black and red arrows, respectively. The peak position has
shifted by ∆l = 0.0082◦ ±0.0062◦ at a 2.1σ confidence level. (b): Same as (a), but in theM 0.74−0.09 cloud. The integral interval is b = −0.11◦ –−0.07◦ . We do not find the peakshift (∆l = 0.0045◦ ±0.0110◦ ).

the naming by the previous work (Koyama et al. 2007b), we hereafter refer to the twosources as Sgr B2 and M 0.74−0.09. Figure 5.2b shows the same image but taken 4
years later in 2009. Comparing the two images, we find that the two bright sources
have darkened in the 4 years, whereas no significant change of the brightness is in thesurrounding region. We also find that the peak position (the Galactic longitude) of Sgr
B2 in 2005 is l = 0.6610◦ ±0.0027◦ , while that in 2009 is l = 0.6692◦ ±0.0056◦ . Thus,
the peak position has shifted by ∆l = 0.0082◦ ±0.0062◦ at a 2.1σ confidence level in 4
years. No significant shift is found in the Galactic latitude (b) direction. Then the shiftcorresponds to 4.0±3.0 light-years toward the east direction, which indicates that theshift speed is consistent with the light-velocity. For M 0.74−0.09, a peak shift is not found(∆l = 0.0045◦ ±0.0110◦ ).
We also make the hard X-ray images of the Sgr B region in the 8–10 keV band for the2005 and 2009 observations in figure 5.2c and 5.2d, respectively. As shown in figure 5.2c,
we find hard X-ray excesses at the positions of Sgr B2 and M 0.74−0.09 in 2005. This is
the first resolved hard X-ray image of Sgr B2 and M 0.74−0.09, although the hard X-ray
excess near Sgr B2 has been found with insufficient spacial resolution of 120 by the past Xray observation (Terrier et al. 2010). On the other hand, the hard X-ray fluxes from theseMCs are almost comparable to those of the surrounding areas and hence no significant
hard X-ray excesses are found in the 2009 data (see figure 5.2d). These facts suggest thatthe hard X-rays from Sgr B2 and M 0.74−0.09 decreased from 2005 to 2009, in correlation
to the decrease of the Fe I-K fluxes.

5.3. ANALYSIS

5.3.2

81

SpectraCounts s−1 keV−1

0.01(a) Sgr B2+ 2005+ 2009

10−3
10−4

χ

5
0
−5
5

Counts s−1 keV−1

0.01

10

Energy (keV)(b) M0.74-0.09+ 2005+ 2009

10−3
10−4

χ

5
0
−5
5

Energy (keV)

10

Figure 5.4: X-ray spectra of Sgr B2 (a) and M 0.74−0.09 (b). Spectra of the 2005 observation and those of the 2009 observation are indicated with the black and red crosses,
respectively. The best-fit models are shown with the solid histograms. Residuals between
the observed data and the models are shown in the bottom panels of (a) and (b).
We investigate the X-ray spectra to study the time variability more quantitatively.
The regions selected for the Sgr B2 and M 0.74−0.09 spectra are indicated with the whitesolid lines in figure 5.2. The background region is also shown with the white dashed lines,
where the dashed small circle in which a MC might be emitting the faint Fe I-K emission
is excluded. Figure 5.4a and 5.4b are the background-subtracted spectra of Sgr B2 andM 0.74−0.09, respectively. The spectra in 2005 and 2009 are shown by black and red
crosses, respectively. As seen in the both spectra, the fluxes of the Fe I-Kα line and thehard continuum have simultaneously decreased from 2005 to 2009.

82

CHAPTER 5. TIME-VARIABILITY OF X-RAYS FROM SGR B2Table 5.1: Best-fit parameters.∗
Source (Year)
unit
Sgr B2 (2005)
Sgr B2 (2009)
M 0.74−0.09 (2005)
M 0.74−0.09 (2009)

NH
1023 cm−2
8.4±1.4
8.8±2.2
5.7±0.6
6.5±1.9

Fe I Kα
10−5 †
13.6±1.1
5.2±0.8
5.1±0.6
2.7±0.5

Fe I Kβ10−5 †
1.3±0.5<0.6
0.4±0.3
0.5±0.3

5–10 keV
10−12 ‡
4.3±0.8
2.1±0.5
1.7±0.3
0.9±0.3∗ The error ranges in this table are calculated at the 90% confidence level.
† Absorption-corrected value in the unit of photons s−1 cm−2 .
‡ Absorption-corrected flux of the continuum excluding the FeI K lines. The
unit is ergs s−1 cm−2 .

We fit the spectra with a phenomenological model of an absorbed power-law plus Fe
I-Kα and Fe I-Kβ lines at 6.4 keV and 7.06 keV, respectively. A supernova remnants(SNR) candidate G 0.61+0.01 (Koyama et al. 2007b) is located 40 away from Sgr B2,
and may slightly contaminate the Sgr B2 spectrum. We therefore take into account thecontamination of the thermal plasma of G 0.61+0.01, by adding a plasma model (Mazzotta
et al. 1998) with fixed temperature of kT = 3.2 keV and iron abundance ZFe =5.1 solar(Koyama et al. 2007b). Then the contamination is found to be less than 10%, and hence
the error due to possible uncertainty of the contaminated flux can be ignored.
The four spectra (Sgr B2 and M 0.74−0.09 for 2005 and 2009) are simultaneously fittedwith free parameters of the absorption column density NH , the iron abundance ZFe , thephoton index Γ, and the fluxes of the continuum and the Fe I-K lines. Among them, only
ZFe and Γ are set to be common for all the spectra. The fit is acceptable with chi2 /d.o.f.
of 216/217 (null hypothesis probability=0.6). The best-fit common parameters of ZFe andΓ are 1.3±0.5 solar and 2.5±0.9, respectively. Those of the other free parameters are listed
in table 5.1.
We find that the fluxes of Fe I-Kα and the hard continuum band for Sgr B2 decreaseby a factor of 0.39±0.06 and 0.49±0.15 from 2005 to 2009, respectively. These variation
factors are consistent within the 1σ confidence level with the mean value of 0.41±0.06. The
variation factors of 0.53±0.10 (Fe I-Kα) and 0.53±0.21 (continuum band) for M 0.74−0.09
are also consistent with the mean value of 0.53±0.09.

5.4

Discussion for the 6.4 keV Clouds in Sgr B2

We confirmed with the 2005 data that the two MCs (Sgr B2 and M 0.74−0.09) exhibit
intense Fe I-K lines. Clear excesses of the hard continuum (8–10 keV) from the MCs arefound for the first time. The most important discovery is the correlated decrease in 4 years
between the Fe I-K lines and the hard continuum in the two MCs. Since the data for theseresults are obtained with the same instrument (Suzaku/XIS) and the identical pointing
field, possible systematic errors would be minimized. The correlated fluxes change in the
two MCs and the two components (Fe I-K and hard continuum) immediately indicate that

5.4. DISCUSSION FOR THE 6.4 KEV CLOUDS IN SGR B2

83

they are in a common origin.
As we have already discussed, the origin is likely to be an X-ray irradiation by an
external source (”X-ray echo”). Then the equivalent width of the Fe I-Kα line (EW6.4keV )is expressed as 1.0×(ZFe /1 solar) keV (chapter 3 in this thesis; Tatischeff et al. 2003;
Nobukawa et al. 2010). Since the observed EW6.4keV is 1.0–1.5 keV, we derive the ironabundance as ZFe =1.0–1.5 solar. The absorption column density NH of about 1024 cm−2
is significantly higher than the value of 1023 cm−2 for typical sources in the GC (Muno etal. 2004), and hence the major fraction would be due to the absorption inside the MCs.
The depth of the iron absorption edge at 7.1 keV gives the iron abundance in the MCs of
ZFe = 1.3 ± 0.5 solar, which is consistent with that derived from the EW6.4keV . These facts
firmly support the ”X-ray echo” for the X-ray emission from Sgr B2 and M 0.74−0.09: thefluorescent Fe I-K line and the Thomson-scattered hard-continuum.
We find that the position of the ”X-ray echo” in Sgr B2 had shifted for 4 light-years
during 4 years toward the east direction, and the fluxes of Sgr B2 and M 0.74−0.09 had
dropped by factors of 0.41±0.6 and 0.53±0.09, respectively. These results indicate that
the X-ray irradiating source had flared up in duration of shorter than 4 years, and was
located in the GC direction with respective to Sgr B2.

Figure 5.5: Schematic face-on view of movement of ”X-ray echoes” in the molecular clouds.
The irradiating source, i.e. the super-massive black hole Sgr A* is indicated with the yellow
star. The line-of-sight positions of Sgr B2 and M 0.74-0.09 are adopted from results by theprevious work (Ryu et al. 2009). The solid parabola guides eyes to see the X-ray echoesfrom Sgr B2 and M 0.74-0.09 arriving at the observer at the same time (equal-time delayline).The unit of distance is in light-years (lys).
The schematic view of the ”X-ray echo” scenario is shown in figure 5.5. As the incident
X-ray flare moves at light speed in the MCs, the bright spot also moves. Since the84

CHAPTER 5. TIME-VARIABILITY OF X-RAYS FROM SGR B2

density of the MCs is not uniform, brightness of the X-ray echoes would be higher inhigh-density regions than those in low-density regions. The X-ray echoes in 2005 would
be produced in high-density regions, and then moved into low-density regions in 2009.
In fact, the radio observation reported a dense region at l = 0.664◦ ±0.001◦ (Sato et al.
2000), which is almost coincident with the position of the bright spot of Sgr B2 in 2005
(l = 0.661◦ ±0.003◦ ). The darker spot at l = 0.669◦ in 2009 means that probably the echo
had already passed the dense region.
The required luminosity (LX ) of the X-ray flare to produce the ”X-ray echo” of theMCs depends on the duration time (t) of the flare, the distances (d) to the MCs, and theoptical depths (τ ) along the X-ray pass in the MCs. To be more exact, LX also depends onthe density profiles and the sizes of MCs. We here approximate these fine-tuning factors
by a single parameter NH . Using the observed NH of about 1024 cm−2 , τ is estimated to
be about 0.1. Then we can estimate LX from Sunyaev et al. (1998); Ryu et al. (2009) as,
µ

¶

µ

¶

2
−1
t
d
erg s−1 .
(5.1)
year
300 lyWe have already noted that the X-ray flare source should be located in the west of Sgr
B2. The minimum value of the required luminosity is therefore given when the flare-sourceis near the west-rim of Sgr B2. In this case, the distance d to the other MC, M 0.74−0.09 isabout 0.1◦ , or 50 light-years (see figure 5.2). Since the observed duration time t is shorter
than 4 years, the minimum value of the required luminosity is LX = 4 × 1038 erg s−1 .
Equation (5.1) is derived from the simple assumptions and approximations, and hencesome uncertainties exist in LX = 4 × 1038 erg s−1 but the errors should not be larger than
one order of magnitude.
The most conservative value of LX is 4 × 1037 erg s−1 , which is nearly equal to or
higher than those of the brightest X-ray sources in our Galaxy. Since we found no suchstable source in the GC, one possible candidate is a transient X-ray binary system. The
X-ray monitoring in the GC has been continuing for nearly 10 years with the INTEGRAL
satellite (Jensen et al. 2003). All the bright transient sources had the peak luminosities
lower than a few times of 1037 erg s−1 (Sakano et al. 2002; Werner et al. 2003). Moreover,the flare durations (t) are shorter than 0.1 years (Sakano et al. 2002; Werner et al. 2003).
The averaged luminosity of transient sources for 4 years are not larger than 1036 erg s−1 ,which is 40 times lower than the minimum required luminosity by the most conservative
estimation.
These facts relegate the transient binary to an implausible candidate for the irradiating
source. We therefore reach the unique conclusion of a big-flare of Sgr A*, the super-massive
black hole. Since d (Sgr B2 to Sgr A*) is 300 light-years (figure 5.5), the big-flare withthe luminosity of larger than 1.5 × 1040 erg s−1 had occurred about 300 years ago. The
flare duration was shorter than 4 years. At present, the X-ray luminosity of Sgr A* is
about 1034 erg s−1 and exhibits many flares, but the flare scales are relatively small andat most 50 times with respect to the quiescent level (Baganoff et al. 2001). Our predictedbig-flare, which is more than 104 times bigger than the biggest flare observed so far, would
be new evidence for the past high activity of the super-massive black hole at the center of
our Galaxy.

LX = 6 × 1040

Chapter 6

Discussion
6.1

Features of the 6.4 keV Clouds

We observed the clouds emitting the strong Fe I K (6.4 keV) lines in the GC region
with the Suzaku/XIS, and studied the detailed features of the clouds. Two possible originsof the Fe I K lines were proposed so far: ionization of neutral atoms due to low energy
cosmic-ray electrons (LECRe) with 10–100 keV (Valinia et al. 2000; Yusef-Zadeh et al.
2007), and photo-ionization due to external X-rays (X-ray reflection nebula: XRN) (e.g.
Koyama et al. 1996; Sunyaev et al. 1998).
We found that the X-ray spectra produced by the two scenarios have large difference inequivalent widths of emission lines. In chapter 3, we predicted that the equivalent widthof Fe I Kα to be ∼0.3 keV and ∼1 keV in the LECRe and XRN scenarios, respectively.
Those values for the other atoms were shown in figure 3.8 and table 3.3. From the Suzaku
data, we obtained that equivalent widths of the Fe I Kα line were large (1–2 keV) for all
clouds. The fact indicates that the X-ray emission from the 6.4 keV clouds is due to the
X-ray origin (chapter 3, 5, and 4).
Best-fit parameters of the 6.4 keV clouds analyzed in this thesis are shown in table 6.1.
The other 6.4 keV clouds, G 0.174−0.233 and M 359.23−0.04 were found by Fukuoka
et al. (2009) and Nakashima et al. (2010), respectively. These sources are shown infigure 6.1. G 0.174−0.233 is at (l, b) ∼ (0.2◦ , 0.2◦ ), ∼ 200 southeast apart from Sgr A*.
M 359.23−0.04 is at (l, b) ∼ (359.2◦ , 0.0◦ ), to the west side from Sgr C. These clouds havealso large equivalent widths ∼ 1 keV of Fe I Kα, and hence are suspected to to be XRNe
(Fukuoka et al. 2009; Nakashima et al. 2010) The obtained parameters are also listed intable 6.1 for this discussion.
The observed equivalent width may suggest that Fe is abundant in each cloud comparedwith the solar value. In particular, since M 359.43−0.07 and M 359.47−0.15 have the large
equivalent widths of the 6.4 keV lines, the Sgr C region is rich in Fe, whose abundance islarger than 2 solar.

85‡
I6.4keV10−5 ph s−1 cm−2
34.0 ± 1.0
13.6 ± 0.7
5.1 ± 0.4
2.8+0.2−0.4
6.4+1.0−1.1
8.8+1.2−0.9
0.6 ± 0.1
3.4+0.6
−0.8

EW
keV
1.15 ± 0.09
1.3 ± 0.2
1.3 ± 0.3
1.4 ± 0.3
2.2+0.3−0.4
2.0+0.2−0.2
1.0 ± 0.2
1.2+0.2−0.4

NH10−23 cm−2
1.9 ± 0.1
8.4 ± 0.9
5.7 ± 0.4
1.5+0.2−0.1
0.9+0.5−0.4
0.8+0.4−0.1
0.8 ± 0.2
3.1+0.4−1.2
1.87 ± 0.04
2.5 ± 0.6
2.5† †
1.8+0.4−0.5
1.7+0.1−0.2
1.6+0.3−0.1
1.7+0.1−0.2
2.9+0.6−0.7

ΓH cm−3
4 × 103
4 × 104
3 × 104
3 × 103∼ 5 × 103
∼ 5 × 103
∼ 1 × 103
7 × 103

n§

k

dproj
pc
30
110
120
90
75
75
50
100

d#
pc
50
110
170
90
75
75
75
100

Lreq ∗∗
erg s−1
2 × 1040
1.5 × 1040
2 × 1040
3 × 1040
3 × 1040
3 × 1040
3 × 1040
2 × 1040∗ Errors are given at the 90% confidence levels.
† The parameters of G 0.174−0.233 and M 359.23−0.04 are referred to Fukuoka et al. (2009) and Nakashima et al. (2010), respectively.
‡ Absorption-corrected intensity of the 6.4 keV line.
§ Hydrogen densities in the clouds.
k Projected distance to the super-massive black hole, Sgr A*.
# Actual distance to Sgr A* on the assumption that the clouds align on the paraboloidal surface (see text).
∗ Required luminosity for the 6.4 keV line if the irradiating source is Sgr A*.
†† Linked to the same value with that of Sgr B2.

Value
Unit
Sgr A
Sgr B2
M 0.74−0.09
Sgr B1
M 359.43−0.07
M 359.47−0.15
G 0.174−0.233†
M 359.23−0.04†

Table 6.1: Best-fit parameters of various 6.4 keV clouds studied in this thesis.

86
CHAPTER 6. DISCUSSION

6.2. IRRADIATING SOURCE OF THE 6.4 KEV CLOUDS

87

The observed absorption column densities NH are 1023 –1024 cm−2 for the clouds. The
NH values are larger than that of the interstellar matter in the GC region (0.6×1023 cm−2 )(Predehl & Schmitt 1995; Muno et al. 2006). The large NH should reflects the dense
matter in the clouds. Number densities n in the clouds can be roughly estimated from theobserved NH and surface areas S if the clouds are like a sphere. The density is described√
as n = NH / S, and is shown in table 6.1.

6.2

Irradiating Source of the 6.4 keV Clouds

We do not know whether a single object illuminate all the 6.4 keV clouds or not.
Bright X-ray binaries are found in the GC region, We may consider the possibility of a
scenario that the 6.4 keV clouds are illuminated by different objects. However, only one
cloud, M 359.23−0.04, accompanies with a bright source (Great Annihilator) with enoughluminosity to explain the observed Fe I Kα flux (Nakashima et al. 2010). Thus, this
scenario seems to be unlikely.
On the other hand, in the case of the single origin, the irradiating source should exist
at the central region since the clouds are distributed on the east (Sgr B1, B2) and west
(Sgr C) regions with respect to the GC. Moreover, the peak shift of the Sgr B2 cloud alsofavors this idea that the source is near the GC (chapter 5). We estimated the required
luminosity Lreq > 1040 erg s−1 for the Sgr B2 cloud in chapter 5. The source with thehigh luminosity should have a large mass of > 50 M¯ because the Eddington luminosity is
LEdd = 2 × 1038 M¯ erg s−1 . Hence the super-massive black hole, Sgr A*, with the hugemass of ∼ 4 × 106 M¯ is the only plausible candidate.
From the two Suzaku observations with 4 years interval, we discovered the time variability of the X-ray emission from the Sgr B2 clouds (chapter 5). The fact suggests thatthe external irradiation was not made persistently but instantaneously, and hence Sgr A*
had flared up for the duration time < 4 years in the past. We examine whether the single
flare of Sgr A* simultaneously illuminated all the 6.4 keV clouds including the others thanthe Sgr B2 cloud.
According to Sunyaev et al. (1998), the required luminosity of the irradiating source
is describe as,µ40

Lreq = 5 × 10

I6.4 keV
10−5

¶µ

n
3
10 cm−3¶−1 µd100 pc

¶2 µ

∆t
4 yr

¶−1

erg s−1 .

(6.1)

I6.4 keV are the 6.4 keV line intensity in the unit of photons s−1 cm−2 . A hydrogen number
density n in the clouds. d is a distance between the 6.4 keV cloud and the irradiating
source. ∆t is the duration time for which the external source irradiates the cloud.
The distances d to the irradiating source for 6.4 keV clouds are not all known, while
Ryu et al. (2009) estimated the relative position of the Sgr B2 cloud to Sgr A* (seefigure 5.5). On the assumption of the single flare origin, total passes in which X-rays fromSgr A* reach the clouds and then the generated 6.4 keV X-rays arrive at an observer (theSuzaku satellite) should be the same. Then, the 6.4 keV clouds should align on the same
paraboloidal surface plotted in figure 6.2, on the basis of the position of the Sgr B2 cloud.

88

CHAPTER 6. DISCUSSION

0.5

M359.43-0.07

Sgr B2
Sgr A*
0.0

M359.23-0.04

M359.47-0.15

M0.74-0.09

Sgr B1

Sgr A
Sgr C

G0.174-0.233
-0.5

359.0

359.5

0.0

0.5

1.0

Figure 6.1: X-ray image of the GC region in the 6.4 keV line, same as figure 3.1. Dashedregions indicated molecular clouds found by the radio observations but with no 6.4 keVemission.

G0.174-0.233

Face on Viewof the GC region

Sgr A

Sgr B1[ly]

300

Sgr B2100
Sgr A*

M359.23-0.04

Line of Sight

M0.74-0.09

Sgr C

ObserverFigure 6.2: Schematic face-on view of distribution of the molecular clouds emitting the
6.4 keV line if the irradiating source is a single big flare of the super-massive black hole,
Sgr A*, which is marked with the yellow star. The solid curve is a parabola on which
the light passes of all molecular clouds (open circles) are equal to that of Sgr A*—Sgr B2(filled circle)—observer.

6.2. IRRADIATING SOURCE OF THE 6.4 KEV CLOUDS

89

The actual distances d are estimated from the distribution. According to the result in
the Sgr B2 observations (chapter 5), the duration time of the past flare is assumed to be
4 years in the following discussion. As stated in the discussion of chapter 5, we note that
there are some uncertainties in the estimation but the errors should not be larger than
one order of magnitude.
We estimated the required luminosity Lreq for the 6.4 keV clouds by the above equation (6.1) and described the results in table 6.1. The Lreq of all the clouds are ∼ 1–
3 × 1040 erg s−1 . Thus, the single origin of the Sgr A* flare successfully explains theobserved results.
Sofue (1990) mentioned a giant molecular cloud with 2 × 106 M¯ at (l, b) = (0.52◦ ,−0.05◦ ) to the north of the Sgr B1 cloud in a radio observation of 13 CO emission line
(Bally et al. 1987). The large mass is comparable to that of the Sgr B2 cloud, but we havefound no X-ray emission at the corresponding position. By a radio observation, Tsuboi
et al. (1999) found a giant molecular cloud at (l, b) = (359.55◦ , −0.05◦ ) in the Sgr C
region. However, the radio cloud is ∼ 0.1◦ apart from the 6.4 keV clouds, M 359.43−0.07
or M 359.47−0.15 and is dim in X-ray. A giant molecular cloud without 6.4 keV emission
has been also found at (l, b) ∼ (−359.85◦ , −0.10◦ ), whose projected distance to Sgr A* is
about 0.1◦ (Sofue 1990; Tsuboi et al. 1999). Velocities obtained from the Doppler shift ofthe molecule lines suggest that these molecular clouds exist in the GC region. The X-ray
flare of Sgr A* might have already pass over those molecular clouds or have not reachedthem.

Acknowledgement
I would like to express my sincere gratitude to my supervisor, Prof. Koyama forproviding me this precious study opportunity as a Ph.D student in his laboratory. My
deepest appreciation goes to my supervisor, Prof. Tsuru for his elaborated guidance,
considerable encouragement and invaluable discussion that make my research of great
achievement and my study life unforgettable.
I am very grateful to members of the present and past cosmic-ray laboratory.
I would also like to express my gratitude to my family for their moral support andwarm encouragements.

91References[1] Aharonian, F., et al. 2006, Nature, 439, 695[2] Anders, E., & Grevesse, N. 1989, Geochim. Cosmochim. Acta, 53, 197[3] Baganoff, F. K., et al. 2001, Nature, 413, 45[4] Baganoff, F. K., et al. 2003, ApJ, 591, 891[5] Balick, B., & Brown, R. L. 1974, ApJ, 194, 265[6] BaÃlucińska-Church, M., & McCammon, D. 1992, ApJ, 400, 699[7] Bautz, M. W., et al. 1998, proc. SPIE, 3444, 210[8] Bautz, M. W., Kissel, S. E., Prigozhin, G. Y., LaMarr, B., Burke, B. E., & Gregory,J. A. 2004, Proc. SPIE, 5501, 111[9] Bautz, M. W., et al. 2007, Proc. SPIE, 6686, 66860Q[10] Beckert, T., Duschl, W. J., Mezger, P. G., & Zylka, R. 1996, A&A, 307, 450[11] Bleach, R. D., Boldt, E. A., Holt, S. S., Schwartz, D. A., & Serlemitsos, P. J. 1972,
ApJ, 174, L101
[12] Blumenthal, G. R., & Gould, R. J. 1970, Rev. Mod. Phys., 42, 237[13] Brinkmann, W., Kawai, N., Matsuoka, M., & Fink, H. H. 1991, A&A, 241, 112[14] Burke, B. E., Gregory, J. A., Bautz, M. W., Prigozhin, G. Y., Kissel, S. E., Kosicki,
B. B., Loomis, A. H., & Young, D. J. 1997, IEEE Transactions on Electron Devices,44, 1633[15] Burke, B. E., Mountain, R. W., Doniels, P. J., Cooper, M. J., Dolat, V. S. 1993,
Proc. SPIE Vol. 2006, p. 272-285, EUV, X-Ray, and Gamma-Ray Instrumentation for
Astronomy IV, Oswald H. Siegmund; Ed.
[16] Cooke, B. A., Griffiths, R. E., & Pounds, K. A. 1969, Nature, 224, 134[17] Cremonesi, D. I., Mereghetti, S., Sidoli, L., & Israel, G. L. 1999, A&A, 345, 826[18] Del Santo, M., Sidoli, L., Bazzano, A., Cocchi, M., De Cesare, G., Paizis, A., &Ubertini, P. 2006, A&A, 456, 1105
93

94

REFERENCES[19] Dutra, C.M., & Bica, E. 2000, A&A, 359, L9[20] Eisenhauer, F., et al. 2003, ApJL, 597, L121
[21] Eyles, C. J., Skinner, G. K., Willmore, A. P., & Rosenberg, F. D. 1975, Nature, 257,
291
[22] Fukuoka, R., Koyama, K., Ryu, S. G., & Tsuru, T. G. 2009, PASJ, 61, 593[23] Garmire, G. P., et al. 1992, in AIAA, Space Programs and Technologies Conference(Paper 92-1473) (New York: AIAA)[24] Genzel, R., Pichon, C., Eckart, A., Gerhard, O. E., & Ott, T. 2000, MNRAS, 317,
348[25] Ghez, A. M., Morris, M., Becklin, E. E., Tanner, A., & Kremenek, T. 2000, Nature,
407, 349[26] Ghez, A., M., et al. 2003, ApJ, 586, 127L
[27] Ghez, A., M., Salim, , S., Hornstein, S. D., Tanner, A., Lu, J. R., Morris, M., Becklin,
E. E., & Duchene, G. 2005, ApJ, 620, 744[28] Ghez, A. M., et al. 2008, ApJ, 689, 1044
[29] Ginzburg, V. L., & Syrovatskii, S. I. 1964, The Origin of Cosmic Rays (New York:
Pergamon)[30] Henke, B. L., Lee, P., Tanaka, T. J., Shimabukuro, R. L., & Fujikawa, B. K. 1982,
At. Data Nucl. Data Tables, 27, 1[31] Hudson, H. S., Peterson, L. E., & Schwartz, D. A. 1971, Nature, 230, 177[32] Hyodo, Y, Tsujimoto, M., Koyama, K., Nishiyama, S., Nagata, T., Sakon, I., Murakami, H., & Matsumoto, H. 2008, PASJ, 60, 173[33] Hyodo, Y., Ueda, Y., Yuasa, T., Maeda, Y., Makishima, K., & Koyama, K. 2009,
PASJ, 61, S99[34] Inui, T., Koyama, K., Matsumoto, H., & Tsuru, T. G. 2009, PASJ, 61, S241
[35] Ishisaki, Y., et al. 2007, PASJ, 59, S113
[36] Iwan, D., Shafer, R. A., Marshall, F. E., Boldt, E. A., Mushotzky, R. F., & Stottlemyer, A. 1982, /apj, 260, 111[37] Jansen, F., et al. 2001, A&A, 365, L1[38] Jensen, P. L., et al. 2003, A&A, 411, L7[39] Kawai, N., Fenimore, E. E., Middleditch, J., Cruddace, R. G., Fritz, G. G., Snyder,
W. A., & Ulmer, M. P. 1988, ApJ, 330, 130

REFERENCES

95[40] Kelley, L. R., et al. 2007, PASJ, 59, S77[41] Kennea, J. A., & Skinner G. K. 1996, PASJ, 48, L117
[42] Kokubun, M., et al. 2007, PASJ, 59, S53[43] Kotani, T., Kawai, N., Matsuoka, M., & Brinkmann, W. 1996, PASJ, 48, 619[44] Koyama, K., Makishima, K., Tanaka, Y., & Tsunemi, H. 1986a, PASJ, 38, 121[45] Koyama, K., Ikeuchi, S., & Tomisaka, K. 1986b, PASJ, 38, 503[46] Koyama, K., Awaki, H., Kunieda, H., Takano, S., Tawara, Y., Yamauchi, S., Hatsukade, I., & Nagase, F. 1989, Nature, 339, 603[47] Koyama, K., Maeda, Y., Sonobe, T., Takeshima, T., Tanaka, Y., & Yamauchi, S.
1996, PASJ, 48, 249[48] Koyama, K., Hyodo, Y., & Inui, T. 2006, Journal of Physics Conference Series, 54,
95[49] Koyama, K., et al. 2007a, PASJ, 59, S23[50] Koyama, K., et al. 2007b, PASJ, 59, S221
[51] Koyama, K., et al. 2007c, PASJ, 59, S245
[52] Koyama, K., Inui, T., Matsumoto, H., & Tsuru, T. G. 2008, PASJ, 60, S201
[53] Koyama, K., Takikawa, Y., Hyodo, Y., Inui, T., Nobukawa, M., Matsumoto, H., &Tsuru, G. T. 2009, PASJ, 61, S255
[54] Krabbe, A., et al. 1995, ApJ, 447, L95[55] Kubota, A. et al. 2007, PASJ, 59, S185
[56] Kushino, A., Ishisaki, Y., Morita, U., Yamasaki, N. Y., Ishida, M., Ohashi, T., &Ueda, Y., 2002, PASJ, 54, 327[57] LaRosa, T. N., Kassim, N. E., Lazio, T. J. W., & Hyman, S. D. 2000, A&A, 119, 207[58] Liszt, H. S., & Spiker, R. W. 1995, ApJS, 98, 259[59] Law, C., & Yusef-Zadeh, F. 2004, ApJ, 611, 858[60] Maeda, Y., Koyama, K., Sakano, M., Takeshima, T., & Yamauchi, S. 1996, PASJ,48, 417[61] Maeda, Y., et al. 2002, ApJ, 570, 671[62] Markevitch, M., Sunyaev, R. A., & Pavlinsky, M. 1993, Nature, 364, 40[63] Matsumoto, H., et al. 2006, Proc. SPIE, 6266

96

REFERENCES[64] Mazzotta, P., Mazzitelli, G., Colafrancesco, S., & Vittorio, N. 1998, A&AS, 133, 403[65] Mehringer, D. M., Palmer, P., & Goss, W. M. 1993, ApJ, 402L, 69M[66] Mendoza, C., Kallman, T. R., Bautista, M. A., & Palmeri, P. 2004, A&A, 414, 377[67] Mirabel, I. F. 1994, ApJS, 92, 369[68] Mitsuda, K., et al. 2007, PASJ, 59, S1[69] Mori, H., et al. 2005, PASJ, 57, 245[70] Mori, H., et al. 2007, Proc. SPIE, 6686
[71] Mori, H., Tsuru, T. Go, Hyodo, Y., Koyama, K., & Senda, A. 2008, PASJ, 60, 183[72] Morris, M., & Serabyn, E. 1996, ARA&A, 34, 645[73] Muno, M. P., et al. 2004, ApJ, 613, 326[74] Muno, M. P., Bauer, F. E., Bandyopadhyay, R. M., & Wang, Q. D. 2006, ApJS, 165,173[75] Muno, M. P., Baganoff, F. K., Brandt, W. N., Park, S., & Morris, M. R. 2007, ApJL,
656, L69[76] Muno, M. P., et al. 2008, ApJS, 181, 110[77] Murakami, H., Koyama, K., Sakano, M., Tsujimoto, M., & Maeda, Y. 2000, ApJ,
534, 283[78] Murakami, H., Koyama, K., & Maeda, Y. 2001, ApJ, 558, 687[79] Murakami, H., Koyama, M., Tsujimoto, M., Maeda, Y., & Sakano, M. 2001, ApJ,
550, 297[80] Nakajima, H., et al. 2008, PASJ, 60, S1[81] Nakajima, H. et al. 2009, PASJ, 61, S233
[82] Nakashima, S., Nobukawa, M., Tsuru, T. Go, Koyama, K., & Uchiyama, H. 2010,
PASJ, 62, 971[83] Nobukawa, M., et al. 2008, PASJ, 60, S191
[84] Nobukawa, M., Koyama, K., Ryu, S. G., Tsuru, T. G., & Tatischeff, V. 2010, PASJ,62, 423[85] Nobukawa, M., Ryu, S. G., Tsuru, T. G., & Koyama, K. submitted to ApJL[86] Oka, T., Hasegawa, T., Sato, F., Tsuboi, M., & Miyazaki, A. 1999 ApJS, 118, 455[87] Ozawa, M., et al. PASJ, 61, S1

REFERENCES

97[88] Ozernoy, L. M., Genzel, R., & Usov, V. V. 1997, MNRAS, 288, 237[89] Park, S., Muno, M. P., Baganoff, F. K., Maeda, Y., Morris, M., Howard, C.,
Bautz, M. W., & Garmire, G. P. 2004, ApJ, 603, 548[90] Paumard, T., Maillard, J. P., Morris, M., & Rigaut, F. 2001, A&A, 366, 466[91] Pavlinsky, M. N., Grebenev, S. A., & Sunyaev, R. A. 1994, ApJ, 425, 110[92] Ponti, G., Terrier, R., Goldwurm, A., Belanger, G., & Trap, G. 2010, ApJ, 714, 732[93] Porquet, D., Rodrigues, J., Corbel, S., Goldoni, P., Warwick, R. S., Goldwurm, A.,& Decourchelle, A. 2003, A&A, 406, 299[94] Predehl, P., & Trümper, J. 1994, A&A, 290, L29[95] Predehl, P., & Schmitt, J. H. M. M. 1995, A&A, 293, 889[96] Predehl, P., & Zinnecker, H. 1996, in ASP Conf. Ser. 102, The Galactic Center, ed.
R. Gredel (San Francisco: ASP), 415[97] Predehl, P., Costantini, E., Hasinger, G., & Tanaka, Y. 2003, Astronomische
Nachrichten, 324, 73[98] Prigozhin, G., Burke, B., Bautz, M., Kissel, S., & Lamarr, B. 2008, IEEE Transactions
on Electron Devices, 55, 2111
[99] Protheroe, R. J., Wolfendale, A. W., & Wdowczyk, J., 1980, MNRAS, 192, 445[100] Rasmussen, A. P., Behar, E., Kahn, S. M., den Herder, J. W., & van der Heyden,
K. 2001, A&A, 365, L231
[101] Revnivtsev, M. G., et al. 2004, A&A, 425, L49[102] Revnivtsev, M., Vikhlinin, A., & Sazonov, S. 2007, A&A, 473, 857[103] Rieke. G. H., Rieke, M. J., & Paul, A. E. 1989, ApJ, 336, 752[104] Plucinsky, P. P., et al. 2008, Proc. SPIE, 7011
[105] Ryu, G. S., Koyama, K., Nobukawa, M., Fukuoka, R., & Tsuru, G. T., 2009, PASJ,61, 751[106] Sakano, M., Imanishi, K., Tsujimoto, M., Koyama, K., & Maeda, Y. 1999, ApJ, 520,
316
[107] Sakano, M., Koyama, K., Murakami, H., Maeda, Y., & Yamauchi, S. 2002, ApJS,
138, 19[108] Sato, F., Hasegawa, T., Whiteoak, J. B., & Miyawaki, R. 2000, ApJ, 535, 857[109] Sawada, T., Hasegawa, T., Handa, T., & Cohen, R. J. 2004, MNRAS, 349, 1167

98

REFERENCES[110] Serlemitsos, P. J., et al. 1995, PASJ, 47, 105[111] Serlemitsos, P. J., et al. 2007, PASJ, 59, S9[112] Sidoli, L., Mereghetti, S., Israel, G. L., Chiappetti, L., Treves, A., & Orlandini, M.
1999, ApJ, 525, 215[113] Sidoli, L., Mereghetti, S., Favata, F., Oosterbroek, T., & Parmar, A. N. 2006, ApJ,
456, 287[114] Skinner et al.(1987) Skinner, G. K., Willmore, A. P., Eyles, C. J., Bertram, D., &Church, M. J. 1987, Nature, 330, 544[115] Smith, D. M., Heindl, W. A., & Swank, J. H. 2002, ApJ, 578, 89[116] Sofue, Y. 1990, PASJ, 42, 827[117] Sofue, Y. 1995, PASJ, 47, 527[118] Strüder, L., et al. 2001, A&A, 365, L18[119] Sunyaev, R. A., Markevitch, M., & Pavlinsky, M. 1993, ApJ, 407, 606[120] Sunyaev, R., & Churazov, E. 1998, MNRAS, 297, 1279
[121] Takahashi, T., et al. 2007, PASJ, 59, S35[122] Tanaka, Y., Koyama, K., Maeda, Y., Sonobe, T. 2000, PASJ, 52, L25[123] Tatischeff, V. 2003, in Final Stage of Stellar Evolution, ed. C. Motch &J. M. Hameury (Les Ulis: EDP Sciences), 79[124] Tawa, N., et al. 2008, PASJ, 60, S11[125] Terrier, R., et al. 2010, ApJ, 719, 143[126] Tomida, H., et al. 1997, PASJ, 49, 405[127] Tsuboi, M., Handa, T., & Ukita, N. 1999, ApJS, 120, 1[128] Tsujimoto, M., Hyodo, Y., & Koyama, K. 2007, PASJ, 59, S229
[129] Turner, M. J. L., et al. 2001, A&A, 365, L27[130] Uchiyama, H., et al. 2009, PASJ, 61, S9[131] Uchiyama, Y., et al. 2008, PASJ, 60, S35[132] Valinia, A., Tatischeff, V., Arnaud, K., Ebisawa, K., & Ramaty, R. 2000, ApJ, 543,
733
[133] Watson et al.(1981) Watson, M. G., Willingale, R., Hertz, P., & Grindlay, J. E.
1981, ApJ, 250, 142

REFERENCES

99[134] Weisskopf, M. C., Brinkman, B., Canizares, C., Garmire, G., Murray, S., & Van
Speybroeck, L. P. 2002, PASP, 114, 1[135] Werner, N., et al. 2003, A&A, 416, 311[136] Worrall, D. M., Marshall, F. E., Boldt, E. A., & Swank, J. H. 1982. ApJ, 255, 111[137] Yamaguchi, H., et al. 2006, Proc. SPIE, 6266
[138] Yamauchi, S., Kawada, M., Koyama, K., Kunieda, H., Tawara, Y., & Hatsukade, I.
1990, ApJ, 365, 532[139] Yusef-Zadeh, F., Law, C., Wardle, M., Wang, Q. D., Fruscione, A., Lang, C. C., &Cotera, A. 2002, ApJ, 570, 665[140] Yusef-Zadeh, F., Muno, M., Wardle, M., & Lis, D. C. 2007, ApJ, 656, 847