High Resolution X-Ray Observationsof the Galactic Center Region
Hiroshi Murakami
Department of Physics. Graduate School of Science,
Kyoto University, Sakyo-ku, Kyoto 606-8502, Japan

Submitted to the Department of Physics,
Kyoto University on January 7, 2002,
for the degree of Doctor of Science (Physics)

1

Abstract
We present the ASCA and Chandra observations of diﬀuse X-ray structures in our Galactic
center region. We concentrate on giant molecular clouds, supernova remnants, and other faintstructures which are newly discovered with Chandra.
We observed four giant molecular clouds: Sgr B2, Sgr C, M0.11−0.08 (at the Radio Arc
region), and Clump2. All clouds emit strong 6.4-keV line with the equivalent width of >∼ 1 keV,and exhibit larger column densities than interstellar absorption. The distribution of the 6.4keV line is coincident with molecular cloud. The distribution of Sgr B2 and Sgr C are shiftedto the Galactic Center side, in particular. These features imply that heavy molecular clouds
are irradiated by an external X-ray source in the direction of the Galactic center and emit
fluorescent and scattered X-rays. We call this new category of X-ray source “X-ray Reflection
Nebula”.
Numerical simulations of X-ray reflection nebula model successfully reproduce the observed
image and spectrum when the mass of the cloud is smaller than 0.3 times of the mass estimatedby the radio observations of molecular lines. This result supports that the molecular clouds in
the Galactic center are in pressure equilibrium, not in gravity equilibrium. The abundances of
the clouds are also estimated by numerical simulations. The iron abundance and other lighter
elements of the Sgr B2 cloud are estimated to be 1.5–3.0, and 2.0–5.0, respectively in case ofcloud mass is 0.2 times of the radio estimated mass. The abundances of other clouds are notconstrained in significant region. The lower limit of the iron abundance is 1.0 for all of the threeclouds. The abundance of other lighter elements are constrained to be <∼ 1 for M0.11−0.08,
and <∼ 3 for Sgr C and Clump 2.
The X-ray reflection nebula model are used to investigate the past activities of our Galactic

center. The observed fluorescent and scattered X-rays from the molecular clouds reflect thepast luminosity of the primary source, which would be the massive black hole Sgr A ∗. It was
as luminous as 1039 erg s−1 ∼ 300 – 2000 years ago. Then the luminosity decrease gradually
to present value. This time variability agrees with the implication obtained by the Chandra

observation of Sgr A East. Sgr A East is the 104 -year-old supernova remnant, and the shock
front passed by Sgr A∗ about 103 years ago. Large amount of matters accreted to Sgr A∗ in
the passing time, and then surrounding materials are swept away.
At least three supernova remnants are discovered in the GC region with much higher angular resolution of Chandra than past X-ray observatories. The number of detected supernova
remnants are insuﬃcient to explain the total energy of the high temperature plasma in the GC

2region.
Chandra discovered faint diﬀuse structures in the GC region. They are candidates of supernova remnants, X-ray reflection nebulae, or something else. Some of them exhibit non-thermal
spectrum, and thought to be acceleration sites of high energy cosmic rays.

Contents
1 Introduction

13

2 Review

15

2.1

2.2

Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15
2.1.1

Abundance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 16

2.1.2

Magnetic Field . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 16

The Sgr A Region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17
2.2.1

2.3

2.4

Galactic Nucleus Sgr A∗ . . . . . . . . . . . . . . . . . . . . . . . . . . . 18

The Distribution of Giant Molecular Clouds . . . . . . . . . . . . . . . . . . . . 18
2.3.1

Sgr B2 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19

2.3.2

Sgr C

2.3.3

Radio Arc . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 21

2.3.4

Clump 2 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 22

. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 20

X-ray Sources in the Galactic Center Region . . . . . . . . . . . . . . . . . . . . 24
2.4.1

X-ray Luminosity of Sgr A∗ . . . . . . . . . . . . . . . . . . . . . . . . . 24

2.4.2

High Temperature Plasma . . . . . . . . . . . . . . . . . . . . . . . . . . 24

2.4.3

Neutral Iron Line . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 25

3 Instruments
3.1

3.2

29

ASCA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 29
3.1.1

Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 29

3.1.2

Attitude Control System . . . . . . . . . . . . . . . . . . . . . . . . . . . 29

3.1.3

XRT (X-Ray Telescope) . . . . . . . . . . . . . . . . . . . . . . . . . . . 31

3.1.4

GIS (Gas Imaging Spectrometer) . . . . . . . . . . . . . . . . . . . . . . 32

3.1.5

SIS (Solid-state Imaging Spectrometer). . . . . . . . . . . . . . . . . . 37

Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 41
3

4

CONTENTS
3.2.1

Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 41

3.2.2

PCAD (Pointing Control and Aspect Determination) System . . . . . . . 42

3.2.3

HRMA (High Resolution Mirror Assembly) . . . . . . . . . . . . . . . . . 45

3.2.4

ACIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 47

3.2.5

Other Instruments — HRC, HETG & LETG . . . . . . . . . . . . . . . . 49

4 Observations

53

4.1

ASCA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 53

4.2

Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 55

5 Results I. — 6.4-keV line emissions
5.1

5.2

5.3

5.4

5.5

Sgr B2 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59
5.1.1

ASCA Observations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59

5.1.2

Chandra Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 62

5.1.3

Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 66

Sgr C . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 69
5.2.1

ASCA Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 69

5.2.2

Chandra Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 72

5.2.3

Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 74

M0.11−0.08 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 76
5.3.1

ASCA Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 76

5.3.2

Chandra Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 80

5.3.3

Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 84

Clump 2 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 84
5.4.1

ASCA Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 84

5.4.2

Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 89

Numerical Simulations of XRN Model . . . . . . . . . . . . . . . . . . . . . . . . 90
5.5.1

Conditions of Simulations . . . . . . . . . . . . . . . . . . . . . . . . . . 90

5.5.2

Simulated Images . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 94

5.5.3

Simulated Spectra. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 97

6 Results II. — Supernova Remnants
6.1

59

G 359.92−0.09

105. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 105

6.1.1

Radio Observations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 105

6.1.2

Chandra Observation. . . . . . . . . . . . . . . . . . . . . . . . . . . . 107

CONTENTS

5

6.1.3

Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110

6.2

G0.570−0.018 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 112
6.2.1

ASCA Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 112

6.2.2

Chandra Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 115

6.2.3

Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 117

7 Results III. — Other Structures
7.1

7.2

7.3

Non-Thermal Clumps . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 121
7.1.1

No. 2 — G359.96−0.05 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 121

7.1.2

No. 11 — G0.13−0.01 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 125

6.4-keV Clumps . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 127
7.2.1

No. 10 — G0.09−0.09 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 127

7.2.2

No. 12 — G359.42−0.13 . . . . . . . . . . . . . . . . . . . . . . . . . . . 127

6.7-keV Clumps . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 130
7.3.1

8 Discussion
8.1

8.2

121

No. 1 — G359.88−0.08, a part of G0.570−0.018 . . . . . . . . . . . . . . 130
131

Verification and Consequence of the X-ray Reflection Nebula . . . . . . . . . . . 131
8.1.1

Required Luminosity . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 131

8.1.2

The Past Activities of the Galactic Nucleus . . . . . . . . . . . . . . . . 133

8.1.3

Mass and Abundance of the Cloud . . . . . . . . . . . . . . . . . . . . . 134

8.1.4

Further Verifications . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 136

Origin of the High Temperature Plasma . . . . . . . . . . . . . . . . . . . . . . 137
8.2.1

Energy Injection . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 137

8.2.2

Temperature . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 138

6

CONTENTS

List of Figures
2.1

Radio continuum image of the Sgr A region . . . . . . . . . . . . . . . . . . . . 17

2.2

12

2.3

Radio image of the Sgr B2 region . . . . . . . . . . . . . . . . . . . . . . . . . . 19

2.4Radio image of the Sgr C region . . . . . . . . . . . . . . . . . . . . . . . . . . . 20

2.5Radio image of 12 CO (J =1–0) and 20 cm continuum at the Radio Arc region . 21

2.6

Longitude-velocity diagram in the GC region . . . . . . . . . . . . . . . . . . . . 22

2.7Radio image of 12 CO (J =1–0) at the Clump 2 . . . . . . . . . . . . . . . . . . . 23

2.8

The distribution of the high temperature plasma in the GC region . . . . . . . . 26

2.9

The distribution of the 6.4-keV line intensity in the GC region . . . . . . . . . . 27

3.1

The schematic images of the ASCA satellite . . . . . . . . . . . . . . . . . . . . 30

3.2

Energy and position dependence of XRT eﬀective area. . . . . . . . . . . . . . . 31

3.3

ASCA PSFs of XRT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 33

3.4

The encircled energy function of the XRT . . . . . . . . . . . . . . . . . . . . . 34

3.5

Schematic views for incorrect light paths with the stray light photons . . . . . . 34

3.6

Observed stray light structures due to Crab . . . . . . . . . . . . . . . . . . . . 35

3.7

Schematic diagram of the principle of GIS. . . . . . . . . . . . . . . . . . . . . . 35

3.8

GIS spectra of day earth, blank sky and night earth. . . . . . . . . . . . . . . . 36

3.9

Change of GIS NXB since the launch . . . . . . . . . . . . . . . . . . . . . . . . 37

CO map of the GC region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 18

3.10 The layout of the SIS chips . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 39
3.11 SIS Grade definition. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 40
3.12 The schematic view of the Chandra satellite . . . . . . . . . . . . . . . . . . . . 41
3.13 Aspect determination subsystem of Chandra . . . . . . . . . . . . . . . . . . . . 43
3.14 HRMA mirrors . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 45
3.15 Eﬀective area vs. X-ray energy and oﬀ-axis angle . . . . . . . . . . . . . . . . . 46
3.16 Simulated on-axis images . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 46
3.17 Encircled energy radius of HRMA . . . . . . . . . . . . . . . . . . . . . . . . . . 47
7

8

LIST OF FIGURES
3.18 The schematic drawing of the ACIS layout (overhead view) . . . . . . . . . . . . 48
3.19 The schematic drawing of the ACIS layout (perspective view) . . . . . . . . . . 49
3.20 The energy resolution of ACIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . 50
3.21 The quantum eﬃciency of the ACIS CCDs . . . . . . . . . . . . . . . . . . . . . 50
4.1

Exposure map of the ASCA observations . . . . . . . . . . . . . . . . . . . . . . 54

4.2

The FOVs of Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 57

4.3

The Exposure map of Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . 58

5.1

(Sgr B2) ASCA GIS and SIS images

5.2(Sgr B2) ASCA GIS Spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . 61

5.3

(Sgr B2) Chandra ACIS image . . . . . . . . . . . . . . . . . . . . . . . . . . . . 64

5.4

(Sgr B2) Chandra ACIS Spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . 65

5.5

Equivalent width of fluorescent iron 6.4-keV line . . . . . . . . . . . . . . . . . . 68

5.6

(Sgr C) ASCA GIS image . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70

5.7

(Sgr C) ASCA GIS spectrum

5.8(Sgr C) Chandra ACIS image . . . . . . . . . . . . . . . . . . . . . . . . . . . . 73

5.9

(Sgr C) Chandra ACIS spectrum. . . . . . . . . . . . . . . . . . . . . . . . 60

. . . . . . . . . . . . . . . . . . . . . . . . . . . . 71
. . . . . . . . . . . . . . . . . . . . . . . . . . 75

5.10 (M0.11−0.08) ASCA GIS image . . . . . . . . . . . . . . . . . . . . . . . . . . . 77
5.11 (M0.11−0.08) ASCA GIS spectrum . . . . . . . . . . . . . . . . . . . . . . . . . 79
5.12 (M0.11−0.08) Chandra ACIS image . . . . . . . . . . . . . . . . . . . . . . . . . 81
5.13 (M0.11−0.08) Chandra ACIS spectrum . . . . . . . . . . . . . . . . . . . . . . . 83
5.14 Intensity distribution of the iron neutral line in the GC region . . . . . . . . . . 85
5.15 (Clump 2) ASCA GIS image. . . . . . . . . . . . . . . . . . . . . . . . . . . . 86

5.16 (Clump 2) Equivalent width distribution of iron line obtained with ASCA GIS . 87
5.17 (Clump 2) ASCA GIS spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . 88
5.18 The schematic view of the XRN simulation . . . . . . . . . . . . . . . . . . . . . 92
5.19 Simulated image of the XRN model . . . . . . . . . . . . . . . . . . . . . . . . . 96
5.20 The radial profile of the 6.4-keV line emission . . . . . . . . . . . . . . . . . . . 98
5.21 The allowable region of ηZothers and ηZFe . . . . . . . . . . . . . . . . . . . . . . 99
5.22 A simulated spectrum with XRN model . . . . . . . . . . . . . . . . . . . . . . . 100
5.23 Chandra spectrum of Sgr B2 fitted with XRN Model. . . . . . . . . . . . . . . . 101
5.24 Acceptable regions of spectrum fittings with XRN model (η = 0.2) . . . . . . . . 102
5.25 Confidence contours of spectrum fittings with XRN model (η = 0.1) . . . . . . . 103

LIST OF FIGURES

9

6.1

The radio image of the Sgr A region . . . . . . . . . . . . . . . . . . . . . . . . . 106

6.2(G359.92−0.09) Chandra ACIS image . . . . . . . . . . . . . . . . . . . . . . . . 107

6.3(G359.92−0.09) Radial profiles of the X-ray emission . . . . . . . . . . . . . . . 108

6.4(G359.92−0.09) Chandra ACIS spectrum of the eastern half region . . . . . . . . 109

6.5(G359.92−0.09) X-ray spectrum of the “wisp” . . . . . . . . . . . . . . . . . . . 110

6.6(G359.92−0.09) X-ray spectrum of the southwestern quadrant region . . . . . . 111

6.7(G0.570−0.018) ASCA GIS image . . . . . . . . . . . . . . . . . . . . . . . . . . 113

6.8(G0.570−0.018) ASCA GIS spectrum . . . . . . . . . . . . . . . . . . . . . . . . 114

6.9(G0.570−0.018) Chandra ACIS image . . . . . . . . . . . . . . . . . . . . . . . . 115

6.10 (G0.570−0.018) Radial profile of X-ray emission . . . . . . . . . . . . . . . . . . 116
6.11 (G0.570−0.018) Chandra ACIS spectrum . . . . . . . . . . . . . . . . . . . . . . 117
6.12 (G0.570−0.018) The simultaneously fitting of the ASCA and the Chandra spectra118
7.1

Merged image of the Chandra observations of the Sgr A and the Radio Arc region122

7.2(G359.96−0.05) X-ray and radio images

7.3(G359.96−0.05) Chandra ACIS spectrum . . . . . . . . . . . . . . . . . . . . . . 124

7.4(G0.13−0.01) Chandra ACIS image and spectrum . . . . . . . . . . . . . . . . . 126

7.5(G0.09−0.09) Chandra ACIS image and spectrum . . . . . . . . . . . . . . . . . 128

7.6(G359.42−0.13) Chandra ACIS image and spectrum . . . . . . . . . . . . . . . . 129

8.1

The past X-ray luminosity of Sgr A∗ derived by simulation . . . . . . . . . . . . 135. . . . . . . . . . . . . . . . . . . . . . 123

10

LIST OF FIGURESList of Tables
3.1

PCAD Modes . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 42

3.2

Aspect System Performance of Chandra . . . . . . . . . . . . . . . . . . . . . . . 44

4.1

Observational log of ASCA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 55

4.2

Observational log of Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . 56

5.1

(Sgr B2) Best-Fit Results of the X-ray Spectrum with ASCA . . . . . . . . . . . 62

5.2

(Sgr B2) Best-Fit Results of the Combined Spectrum of Point Sources Obtainedwith Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 63

5.3

(Sgr B2) Best-Fit Results of the X-ray Spectrum with Chandra . . . . . . . . . . 66

5.4

(Sgr C) Best-Fit Results of the X-ray Spectrum with ASCA . . . . . . . . . . . 72

5.5

(Sgr C) Best-Fit Results of the X-ray Spectrum with Chandra . . . . . . . . . . 74

5.6

(M0.11−0.08) Best-Fit Results of the X-ray Spectrum with ASCA . . . . . . . . 80

5.7

(M0.11−0.08) Best-Fit Results of the X-ray Spectrum with Chandra . . . . . . . 82

5.8

(Clump 2) Best-Fit Results of the X-ray spectrum with ASCA fitted with a

5.9

The Parameters of the Molecular Clouds . . . . . . . . . . . . . . . . . . . . . . 94

6.1

(G359.92−0.09) Best-Fit Results of the X-ray Spectrum with Chandra . . . . . . 110

6.2(G0.570−0.018) Best-Fit Results of the X-ray Spectrum with ASCA . . . . . . . 114

6.3(G0.570−0.018) Best-Fit Results of the X-ray Spectrum with Chandra . . . . . . 118

6.4(G0.570−0.018) Best-Fit Results of the X-ray Spectrum with the Simultaneously. . 88Fitting of ASCA and Chandra with a NEI Model . . . . . . . . . . . . . . . . . 119
7.1(G359.96−0.05) Best-Fit Results of the X-ray Spectrum with Chandra . . . . . . 123

7.2(G0.13−0.01) Best-Fit Results of the X-ray Spectrum with Chandra . . . . . . . 126

7.3(G0.09−0.09) Best-Fit Results of the X-ray Spectrum with Chandra . . . . . . . 128

7.4(G359.42−0.13) Best-Fit Results of the X-ray Spectrum with Chandra . . . . . . 130

8.1

Required X-ray Luminosities and Observed Ones
11

. . . . . . . . . . . . . . . . . 132

12

LIST OF TABLES
8.2

Required Luminosities of Sgr A∗ . . . . . . . . . . . . . . . . . . . . . . . . . . . 133

Chapter 1
Introduction
The center of our Galaxy, the Galactic center (GC, hereafter), is one of the most interestingtargets in astronomy. A large amount of matter concentrates in this region: stars, moleculargas, HII regions, neutron stars, supernova remnants, and so on. The GC is a peculiar region
with these various objects. Although many observations have been made to understand thenature of the GC, there are many questions till now. The most puzzling problem is the massive
black hole at the Galactic nucleus. The luminosity is anomalously low for the black hole withthe mass of 2.6×106 M⊙ .
In X-ray band, Japanese X-ray satellites Ginga and ASCA had made important discoveries
which imply the past activities in the GC region: a large scale high temperature plasma extending in a 1.◦ 8 × 1.◦ 0 region, and diﬀuse 6.4-keV line emission. The plasma contains the extremely

large energy of 1054 erg, and 6.4-keV line suggests that there was a strong X-ray emitter in the
GC region. They seem to be contradictory to the present quiescent nucleus. Following theseresults, Chandra is increasing the number of X-ray sources in the GC region with the extremely
high angular resolution. These discoveries give us further information on the GC region.
In this thesis, we attempt to reveal the past activities of the GC region with the ASCA andChandra observations. For the ASCA results, Maeda (1998) summarized the global distributionof X-ray emissions, in particular of emission lines, in the GC region. Then Sakano (2000)
investigate the properties of X-ray point sources. Following these results, we mainly study
about the discrete diﬀuse emissions in this thesis.
The main objectives are:
1. to reveal the emission mechanism of 6.4-keV line emissions,
2. to search an origin of the high temperature plasma,13

14

CHAPTER 1. INTRODUCTION
3. to find new discrete diﬀuse sources.
The contents of this thesis are as follows. In chapter 2, we review the nature of our Galac-

tic center region briefly. Then instruments and the details for observations are explained inchapter 3 and 4, respectively. In chapter 5–7, we report the results of the observations which
are related to each item mentioned above. Finally, we discuss and summarize the results inchapter 8.
In this thesis, we assume the distance to the Galactic center to be 8.5 kpc. All the X-ray
sources in the GC region are assumed to be at the same distance.

Chapter 2
Review
In this chapter, we review the nature of our Galactic center (GC) region. Because the GC
region is highly complex with the various phenomena, we do not cover all the objects in thisreview. We give here only brief descriptions of the radio and X-ray sources which are closely
related to our work. More detailed explanations are shown in already published reviews: e.g.
Morris & Serabyn (1996), Mezger, Ducshl, & Zylka (1996), and Melia & Falcke (2001).

2.1

Overview

Optical and ultraviolet wavelengths are heavily absorbed by the interstellar matters along the
Galactic disk, and can not penetrate from the GC region to us. We can observe the GC region
with radio, infrared, X-ray and γ-ray wavelengths.
With the early radio observations, bright structures are detected, which are named Sgr A–E.
The most luminous region, Sgr A, is the Galactic center (section 2.2). Then numerous radio
observations have been made, and many objects are found; e.g. giant molecular clouds (GMC),
non-thermal filaments, HII regions, and so on. The distribution and the properties of molecular
clouds are stated in section 2.3.
In the infrared band, the GC region is filled with many bright stars. The mass density ofstars rises toward the Galactic center. There are many high mass stars at the close vicinity of
the Galactic nucleus. For example, IRS 16 star cluster, which is at the central parsec of the
Galaxy, is comprised of several O type stars (Simons, Hodapp, & Becklin 1990). These stars
indicate that the star burst activity was occurred at the GC region within the past 107 years.
Also in X-ray band, the distinct features are X-ray point sources, the most of which areneutron star binaries or black hole candidates. Diﬀuse X-ray emissions also exist. Details for15

16

CHAPTER 2. REVIEW

X-ray observations are shown in section 2.4.
Some γ-ray observations revealed that the large amount of positrons exists in the GC region.
The 511 keV line from positron annihilation is observed in many times (e.g. Johnson, Harnden,& Haymes 1972), and the distribution is comprised of three components: a central bulge, theGalactic plane, and an enhancement or extension of emission at positive latitudes above the
Galactic center (Purcell et al. 1997; Kinzer et al. 2001). The origin of the positrons is not
explained yet.
1.8 MeV line produced by the 26 Al decay is also detected from the GC region. The line
emission is revealed to concentrate along the Galactic plane with the observation of COMPTEL
onboard CGRO, and the origin of 26 Al is thought to be supernovae (type II) and/or Wolf-Rayetstars (Prantzos 1998). Since the half-life of 26 Al is only about 106 yr, there should have beenstar burst activities in recent past at the GC region.
High energy cosmic rays with energies of about 1018 eV are also found to concentrate in the
GC region and the Cygnus region (Hayashida et al. 1999). This result indicates that cosmic
rays up to 1018 eV, which have been thought to be an extragalactic origin, are also accelerated
in our Galaxy or perhaps at the GC region.

2.1.1

Abundance

It has been debatable whether the GC region is overabundant or not. The observations of the
Galactic disk SNRs within 3 kpc from the Sun indicate that the abundance gradually becomeslarger toward the GC. From the gradient of abundance increase, the abundance of the GC
region is estimated to be about 3 times of the solar value (Binette et al. 1982). While, direct
observations of the stars in the GC region indicated that the abundance is consistent withthe solar value (Carr, Sellgren, & Balachandran 2000). Due to these contradicted results, theabundance of the GC region is not determined.

2.1.2

Magnetic Field

There are many filamentary structures in the GC region: e.g. Radio Arc (section 2.3.3), Sgr C(section 2.3.2). The radio spectra are non-thermal, and thought to be synchrotron emission withthe strong magnetic field. By the measurement of the linear polarization, the magnetic field isdetermined to be mG order, which is much higher than the averaged interstellar magnetic fieldstrength on the Galactic disk of µG order (e.g. Yusef-Zadeh & Morris 1987).

2.2. THE SGR A REGION

17

Thus the GC region is peculiar region with various energetic phenomena. Below we introduce the brief characteristics of individual sources.

2.2

The Sgr A Region

Sgr A is resolved into three parts: Sgr A East, Sgr A West, and a point source Sgr A∗ (Figure 2.1). Sgr A East has shell-like structure, and thought to be a young supernova remnant.
Recently, Maeda et al. (2001) detected diﬀuse X-ray emission from Sgr A East, and verified
the hypothesis (also see section 8.1.2). Sgr A West is the shape of spiral, which is surrounding
Sgr A∗. Sgr A∗ is a dynamical center of our Galaxy, or Galactic nucleus (see next subsection).

Figure 2.1: VLA radio continuum image of the Sgr A region (Fig. 4 in Yusef-Zadeh, Melia, & Wardle
2000). Shell-like structure Sgr A East (light gray) and the spiral-shaped structure Sgr E West (darkgray) are shown. Note that the coordinates are in equinox 1950.

18

CHAPTER 2. REVIEW

Galactic Nucleus Sgr A ∗2.2.1

Sgr A∗ was discovered about a quarter of a century ago (Balick & Brown 1974), and is a strong
candidate of a massive black hole. Many observations support the presence of the massive black
hole. By the observations of the stellar proper motions in the vicinity of the Galactic nucleus
Sgr A∗ , Genzel et al. (1997) revealed a mass of 2.6 ×106 M⊙ is confined within about one

light-week at Sgr A∗ . Because there is no stable configuration of normal stars, stellar remnantsor substellar entities in such a high concentration of the mass, Sgr A∗ must be a massive blackhole at the Galactic nucleus.
There are many high mass stars around Sgr A∗. Based on stellar wind measurements, thegas supply rate is at least Ṁ ∼ 10−6 M⊙ yr−1 (Genzel et al. 1994). With the standard accretion
model, in which the luminosity is calculated to be ∼ 0.1 Ṁ c2 , the luminosity of Sgr A∗ should be
about 1040 erg s−1 . However, observed luminosity is several orders smaller than expected value.

To solve the contradiction, many theories are presented: e.g. advection-dominated accretion
flow (ADAF; Narayan et al. 1998). However, there is no decisive explanation of the quiescent
Galactic nucleus. It is one of the most important problems till now.

2.3

The Distribution of Giant Molecular Clouds

In the radio millimeter band, many molecular clouds are found in the GC region. Figure 2.2
shows the global distribution of the molecular clouds in the GC region. These clouds generally
exhibit large velocity width of several tens km s −1 , and high temperature of 30–60 K. There
are some giant molecular clouds: Sgr B2, Sgr C, M0.11−0.08 (at the Radio Arc region), andClump 2.

Figure 2.2: The intensity distribution of the 12 CO J =1–0 line in the GC region (Oka et al. 1998).
The positions of giant molecular clouds are indicated by arrows

2.3. THE DISTRIBUTION OF GIANT MOLECULAR CLOUDS

2.3.1

19

Sgr B2The Sgr B2 cloud is located at (l, b) ≃ (0.◦ 7, 0.◦ 0), which is ∼ 100 pc from the Sgr A∗ in the

projected distance. The mass of the cloud is estimated to be ∼ 6–7 ×106 M⊙ with the radio

observations of molecular lines (e.g. Oka et al. 1998; Lis & Goldsmith 1989), and one of themost heaviest cloud in the GC region.

Figure 2.3: Radio images of the Sgr B2 region (Sato et al. 2000). (a) Contour map of the 5 GHz
continuum emission imaged with the VLA (J. Whiteoak & F. Gardner 1985, unpublished). The radio
peak of Sgr B2 (N), (M), (S) are shown. (b) Contour map of the 13 CO emission integrated between
VLSR = 35 and 100 km s−1 . (c) Contour map of the C 18 O emission integrated between VLSR = 30 and
100 km s−1 . The positions of masers are shown in (b) and (c) (Gaume & Claussen 1990; Mehringer,
Goss, & Palmer 1994; Houghton & Whiteoak 1995; Mehringer & Menten 1997; Hasegawa et al. 1985;
Morita et al. 1992). The labels are in units of T A∗ × ∆v in K km s−1 . The coordinates of the maps

are relative to (α, δ) 1950 = (17h 44m 10.s0, −18◦ 22′ 00′′). Half-power beamwidths are shown by a filled
ellipse or circle in the lower left-hand corner of each panel.

With the radio observations, there found to be many ultra compact HII regions and masers in
the molecular cloud (Benson, & Johnston 1984; De Pree, Goss, & Gaume 1998; Shiki, Ohishi, &Deguchi 1997; Mehringer, Goss, & Palmer 1994). Most of them concentrate in radio continuum
peak of Sgr B2 North (N), Main (M), and South (S) (Figure 2.3). The ages of of Sgr B2 (N)and Sgr B2 (M) are estimated to be ∼ 104 years, using the photon flux of the exciting stars andthe ambient gas density (De Pree et al. 1995, 1996). In addition, powerful bipolar outflows are

20

CHAPTER 2. REVIEW

found from Sgr B2 (N) and Sgr B2 (M), with the dynamical age of ∼ 104 years (Vogel, Genzel,& Palmer 1987; Lis, et al. 1993; Mehringer 1995; Kuan & Snyder 1996; Liu et al. 1998). All
these results indicate that there are many clusters of high-mass young stellar objects (YSOs)
in the Sgr B2 cloud.

2.3.2

Sgr C

The prominent objects at the Sgr C region are a shell-like H II region (Sgr C itself) and afilamentary structure perpendicular to the Galactic plane (gray-scale image in Figure 2.4). The
filament is non-thermal, and resembles Radio Arc filament (see next section).
There is also a molecular cloud at the Sgr C region as well as these radio continuum structures. In Figure 2.4, 13 CO line emission is shown by contours. Near the basis of the filament,
and above the HII region, the peak of the cloud is shown (CO 359.4+0.0 in Oka et al. 1998).
We call this cloud the Sgr C cloud in this thesis. The mass of the cloud is about 106 M⊙ .

Figure 2.4: Radio image of the Sgr C region. Contours show the 13 CO line emission at 10, 20, · · ·<−1
160 K km s −1 integrated over the range −84 <∼ v ∼ − 24 km s . The intensity distribution of 1616
MHz radio continuum emission is superposed with gray scale. (Liszt & Spiker 1995)

2.3. THE DISTRIBUTION OF GIANT MOLECULAR CLOUDS

2.3.3

21

Radio Arc

There are notable features of arched filaments and connected straight filaments at the positivelongitude side of the GC (Figure 2.5f). This filed is called the Radio Arc region. The straight
filaments are non-thermal, and perpendicular to the Galactic plane with strong linear polarization. The origin of the radio emission is thought to be synchrotron radiation at strong magnetic
field.

Figure 2.5: The velocity sliced contour maps of 12 CO (J =1–0) emissivity (Oka et al. 2001). Contours
are drawn at a interval of!

TR∗ (V )dV = 20 K km s −1 . The dotted ellipses show a kinematical modelfor the expanding cavity. (f) A VLA radio continuum image at 20 cm (Yusef-Zadeh et al. 1984) witha schematic view of the molecular features (Oka et al. 2001).

22

CHAPTER 2. REVIEW
By the observations of radio emission of molecular lines, a molecular cloud is found in theRadio Arc region (M0.11−0.08 in Lindqvist et al. 1995; CO 0.13−0.13 in Oka et al. 2001;
see Figure 2.5). The mass of the cloud is estimated to be about 106 M⊙ (Oka et al. 2001).
Oka et al. (2001) further discovered the expanding cavity at the opposite side of the molecular
cloud with respect to the filaments. They suggested that magnetic tubes were swept up by the
expanding cavity, and compressed by interacting with the molecular cloud. Then an interaction
between the cloud and magnetic field may accelerate electrons to relativistic energies throughshock acceleration or magnetic field reconnection.

2.3.4

Clump 2

Clump 2 is a giant molecular cloud located at (l, b) ≃ (3 ◦ , 0.◦ 2) found by Bania (1977). As seenin Figure 2.6, the cloud is peculiar with the large velocity width of 50–150 km s −1 . The cloudconsists of 16 cores, each having densities in excess of 2×104 cm−3 (Figure 2.7). The total massis estimated to be 3 × 106 M⊙ (Stark et al. 1991; Bally et al. 1988).

Figure 2.6: Schematic longitude-velocity diagram illustrating the major large-scale kinematic featuresin the GC region (Bania 1977).

2.3. THE DISTRIBUTION OF GIANT MOLECULAR CLOUDS

23

Figure 2.7: Velocity integrated map of 12 CO (J =1–0) emission at the Clump 2 (Stark & Bania
1986). The velocity range is from +40 km s −1 to +240 km s−1 . Contours are drawn at 50 K km s −1interval.

24

CHAPTER 2. REVIEW

2.4

X-ray Sources in the Galactic Center Region

In X-ray band, some important discoveries are made by Japanese X-ray satellites. Near the endof 1980s, Ginga discovered that the central 1.◦ 0 × 1.◦ 8 region is filled with the high temperatureplasma (Koyama et al. 1989, Yamauchi et al. 1990). Then ASCA found the clumpy emissionsof neutral iron line (Koyama et al. 1996). The details are mentioned in following sections.

2.4.1

X-ray Luminosity of Sgr A∗

As mentioned in section 2.2.1, the most important problem in the Galactic center region is alack of the activity of the massive black hole Sgr A ∗ . The X-ray emission from the Sgr A region
was reported by several X-ray satellites (Skinner et al. 1987; Kawai et al. 1988; Syunyaev et al.
1991) with the intensity variation. However, ASCA revealed that the X-ray variability should
be attributable to an eclipsing X-ray burster AX J1745.6−2901 near Sgr A region (Maeda
1996), and the X-ray emission from the Sgr A region is extended. The upper limit of the X-ray
intensity from the Sgr A∗ is constrained to be ≤ 1035 erg s−1 . Chandra resolved Sgr A ∗ for the

first time in X-ray band, and the X-ray luminosity is estimated to be 2 × 10 33 erg s−1 (Baganoﬀ
et al. 2001a).

The X-ray luminosity of 2 × 1033 erg s−1 is very small for the massive black hole with themass of 2.6 × 106 M⊙ . This value is about 10−11 of the Eddington limit.

Recently, Chandra detected an X-ray burst from Sgr A ∗ with the luminosity of 1035 erg s−1(Baganoﬀ et al. 2001b). However, it is still a small value, which can be radiated by a ordinal
neutron star binary.

2.4.2

High Temperature Plasma

Ginga found strong emission of highly ionized iron line from the GC region (Koyama et al.
1989). Since the ionized line emission is typical of the optically thin plasma, the discovery of
the iron line indicated the presence of a large amount of high temperature gas at the GC region.
The temperature is estimated to be about 10 keV. Then Yamauchi et al. (1990) investigated the
distribution of the iron line, and revealed that the plasma is an elliptical structure with the sizeof 1.◦ 8 × 1.◦ 0 (FWHM), which is tilted 21 ◦ ± 3◦ with respect to the Galactic plane (Figure 2.8).

The total energy and the total luminosity of the plasma are (4–8) ×1053 erg and (0.9–3)
×1037 erg s−1 in the 2–10 keV band, respectively. Note that such high temperature plasma

cannot be confined with the gravity of the GC region. The typical time scale of the plasma2.4. X-RAY SOURCES IN THE GALACTIC CENTER REGION

25(obtained by dividing its diameter by the sound velocity) is ∼ 50000 yr. Later observation withthe higher energy resolution of ASCA revealed that the line width of the iron line is slightly
broad with the σ of ∼ 75 eV (Koyama et al. 1996).

Following the discovery of the 1.◦ 8 × 1.◦ 0 plasma withGinga, ASCA found another hightemperature plasma at the close vicinity of Sgr A ∗ by the higher angular resolution than Ginga.
The size and the total energy are ∼ 2′ × 3′ and 3 ×1050 erg, respectively. Recently, with themuch higher angular resolution of Chandra (∼ 0.′′ 5), the plasma was found to fill the radio shellstructure, Sgr A East. Maeda et al. (2001) suggested that supernova explosion about 10000 yr
ago might produce the plasm. However, the origin of the larger scale plasma is an open question
till now.

2.4.3

Neutral Iron LineASCA resolved iron lines from the GC region to three characteristic emission lines from theatoms in respective ionized states; 6.4-keV line from neutral iron, 6.7-keV line from He-like
iron, and 6.9-keV line from H-like iron. 6.7-keV and 6.9-keV lines should be attributable tothe high temperature plasma in the GC region, but 6.4-keV line can not be radiated by theplasma.
The distribution of the 6.4-keV iron line is shown in Figure 2.9 (Koyama et al. 1996). There
are two distinct peaks at the Sgr B2 region and the Radio Arc region, which roughly agree withthe position of the giant molecular clouds. In addition, Martı́n-Pintado et al. (2000) found thecorrelation between the 6.4-keV line emission and the SiO (J = 1–0) line intensity. Because
6.4-keV line is a characteristic radiation from neutral irons, it is natural to think that X-rays
are emitted by molecular clouds. However, clouds are very cold, and can not emit high energy
X-rays in themselves. Koyama et al. (1996) implied that clouds are irradiated by an external
X-rays, and emit fluorescent X-rays as 6.4-keV line. It is the most important theme in thisthesis to reveal the origin and the radiation mechanism of the 6.4-keV line.

26

CHAPTER 2. REVIEW

Figure 2.8: The schematic distribution of the GC plasma obtained with Ginga (Yamauchi et al. 1990)

2.4. X-RAY SOURCES IN THE GALACTIC CENTER REGION

Figure 2.9: The distribution of the 6.4-keV line intensity in the GC region.

27

28

CHAPTER 2. REVIEWChapter 3
Instruments
We use the ASCA and the Chandra data in this thesis. In this chapter, we introduce theproperties of both satellites.

3.1

ASCA

3.1.1

Overview

ASCA was the fourth Japanese astronomical X-ray satellite, which was launched into orbit on
February 20, 1993 (Tanaka et al. 1994). The orbital heights at perigee and apogee were about
520 km and 620 km, respectively, whereas the inclination was 31.◦1. The orbital period wasapproximately 96 minutes. The satellite have been reentry into the atmosphere on March 2,
2000.
ASCA was the first satellite which can obtain X-ray images in hard X-ray band above 2 keV(to 10 keV). The X-Ray Telescopes (XRTs) enable to gather high energy X-rays with the large
number of thin-foil layers. ASCA was equipped with two sets of the Gas Imaging Spectrometers(GISs) and the Solid-state Imaging Spectrometers (SISs) at the foci of four identical XRTs.
Figure 3.1 shows the appearance and the schematic configuration of instruments aboard ASCA.

3.1.2

Attitude Control System

The satellite attitude was controlled with four reaction wheels (RW) and three-axis magnetic
torquers (MTQ). RWs exchange the angular momentum with the satellite to control the attitude, and MTQs are used to adjust the total angular momentum by interacting with an
external magnetic force. As an attitude sensor, ASCA had five inertial reference units (IRUs),
29

30

CHAPTER 3. INSTRUMENTS

S-band antennauzuki

X-band antenna
Shunt regulator

Extendable
optical bench

Sun shade

Solar paddle
NSAS

Y

SSAS

X-ray telescope

X

Z
Satellite coordinates(a) The appearance of the ASCA satellite

Sun

Focal Length 3500mm

Sun Shade
ThermalShield

XRT

SIS

GIS
345mm dia
220mm(b) The configuration Φ of XRTs and the detectors aboard ASCA
Figure 3.1: The schematic images of the ASCA satellite. (a) The appearance of the satellite. Thereis four XRTs pointing to the lower left. (b) Configuration of XRTs and the detectors aboard ASCA.
The XRTs are pointing to the left. Two GISs and two SISs are located at the focal plane.

3.1. ASCA

31

two star trackers (STTs; optical CCD camera), one non-spinning solar aspect sensor (NSAS),
one spinning solar aspect sensor (SSAS), and 3-axis geomagnetic aspect sensors (GAS). IRUs
monitor the relative change of the attitude, whereas STTs are used to determine the absolute
attitude of the satellite. The stability of the attitude was normally better than 10 ′′ , whereas
the absolute pointing accuracy was ≃ 40′′1 (Gotthelf 1996), except just after maneuver. The

accuracy can be improved by correction

3.1.3

XRT (X-Ray Telescope)

The ASCA telescopes are configured according to the Wolter type-1 system: concentric parabolaand hyperbola. The four X-ray telescopes aboard ASCA are identical. The mirrors are formed
by a large number of thin-foils (120 foils) with the diameters from 120 mm to 345 mm, and the
focal length of 3.5 m. On the focal plane, 1 mm corresponds to 0.9822 arcmin. Each of XRTs
consists of four identical quarter mirrors, called quadrants.
Figure 3.2a shows the eﬀective area of XRT in relation to an X-ray energy. In general, theeﬀective area becomes smaller with the higher energy X-rays: about 300 cm 2 below 2.0 keV, and200 cm2 at 5.0 keV. The eﬀective area also decrease with the larger oﬀ-axis angles2 (Figure 3.2b).
At the oﬀ-axis angle of 20′ , the eﬀective area becomes 1/4–1/5 of on-axis value.

(a) Eﬀective area vs. X-ray energy(b) Eﬀective area vs. oﬀ-axis angle

Figure 3.2: Energy and position dependence of ASCA XRT eﬀective area.
A point source image, which is called Point Spread Function (PSF), is shown in Figure 3.3.
1

The accuracy can be improved to ≃ 24′′ for GIS by correcting temperature dependent deviation of theattitude (Gotthelf 2000)
2

This dependence of eﬀective area on incident angle is called vignetting.

32

CHAPTER 3. INSTRUMENTS

Since XRTs are made with 4 parts of quadrants, the image seems to be a so-called 4-leaves
pattern. The 4-leaves image becomes flatter in larger oﬀ-axis angle.
A useful parameter for representing the performance of the mirror is “encircled energyfunction” shown in Figure 3.4. This function shows how many photons are included in acircular region of each diameter with the point spread function of mirrors. When we observe a
point source, about a half of X-rays are gathered in 3′ -diameter region, and 80% of X-rays arein 6′ -diameter.
X-ray photons are designed to reflect twice on the surface of the mirror with the Wolter
type-1 system, but some of X-ray photons from out of the field of view (FOV) reach the focal
plane through incorrect paths (Figure 3.5). These photons are called stray light. The intensityof stray light depends on the spectral shape of the source, as well as the source intensity andstructure, because softer photons are more easily reflected with large angle than harder photons.
The stray light does not focus on the focal plane as the usual image, but produces a diﬀuse
image; for example, the stray light from a point source out of FOV makes a shell-like structure
depending on the source oﬀset distance and the angle. Fig. 3.6 shows the observed stray light
structures on the GIS detector, where the stray light source is Crab. The stray light fromlargely extended sources like the cosmic X-ray background, on the other hand, produces no
distinct structure, but smooth extended structure over the focal plane, because of the largesolid angles.
Details for XRTs are found in Serlemitsos et al. (1995).

3.1.4

GIS (Gas Imaging Spectrometer)

The Gas Imaging Spectrometer (GIS) is an imaging gas scintillation proportional counter(Ohashi et al. 1996; Makishima et al. 1996). ASCA has two GISs called GIS2 and GIS3,
which are operated in parallel. Fig. 3.7 shows the schematic diagram of the principle of photon
detection with GIS. GISs are superior to SISs in following points: a wide FOV with radius of∼25 mm (∼ 25′ ), good time resolution, high detection eﬃciency in high energy band, and alittle degradation for years in orbit.

GIS has sensitivity from 0.7 to 15 keV, whereas the nominal energy range is actually limited
below 10 keV, due to lack of sensitivity with XRT above 10 keV. Although GIS has threeobservational modes: PH, PCAL and MPC modes, we here concentrate on the performance in
only the PH mode because all of the observations that we use in this thesis were operated withPH mode.

3.1. ASCA

33

Figure 3.3: Point spread functions of ASCA XRT at four diﬀerent oﬀ-axis angles. Circle and squareshows the field of view for GIS and SIS, respectively.

34

CHAPTER 3. INSTRUMENTS

Figure 3.4: The encircled energy functions of the XRT, which are normalized at 12 ′.

(a)

(b)

(c)

(d)

Primary
mirrors

Secondarymirrors

Figure 3.5: Schematic views for incorrect light paths with the stray light photons: (a) normal path,(b) reflection only by primary mirrors, and (c) only by secondly mirrors, (d) reflection by back surfaceof mirrors (Ishisaki 1997).

3.1. ASCA

35

60arcmin30arcminCrab

Figure 3.6: Observed stray light structures due to Crab taken with GIS (Ishisaki 1997).

HV

X-r a yGND
52mm

support grids- 6kV

Be entrance window

photo electrons
cloud(10um)

Gas cellDrift Region

Xe + He
96%

4%

10mm
- 5kV

intermediate mesh

Fluorescence
UV light

Scintillation
Region15mm

ground mesh

2.5mm
3.0mm

QuartzImaging PhotoMultiplier Tube

Bialkali photo-electrodeX-multi anode x 16
Y-multi anode x 16
Last Dynode

GND

1400V

PH
RT

Y
01 2 3 4 56 7 8 9 A B C D E F

Pre-Amplifier

XP / YP

Figure 3.7: Schematic diagram of the principle of the GIS detector.

36

CHAPTER 3. INSTRUMENTS
GIS achieves low background count rate with eﬀective discriminations of non-X-ray eventswith pulse height, rise time, spread of light, and the position information. Most of non-X-ray
events are rejected in orbit, whereas some of them are rejected in the data analysis at ground 3 .
Fig. 3.8 shows the GIS spectra of day earth4 , blank sky5 and night earth6 . The night earth
data (image and spectrum) are thought to represent the non-X-ray events.

Ne
Mg

0.85

1.26

O
0.52

day Earth(4,620 ks)

Si
1.74

CXB
(649 ks)night Earth(1,406 ks)

Ar
2.96

Mg Al

1.26 1.49

Sn-L
3.4

Xe-L
4.8

Cu
8.04

Figure 3.8: GIS spectra of day earth, blank sky and night earth (Ishisaki 1997).
The long-term monitoring of the night earth data shows that the GIS non-X-ray background (NXB) has been gradually increasing since the launch, whereas the spectral shape
has not changed significantly (Ishisaki et al. 1997; see also Fig. 3.9). The reason is still not
clear, although it is generally thought to be caused by a gradual build-up of long-term decay
radioisotopes.
The bit assignment for the GIS telemetry data represents the quantity of information on(sensor-ID), pulse height, X-ray detected position for X- and Y-axes, rise time information,
spread of light, and time stamp for each event. The standard bit assignment, with which mostof the observations are done, are (1-)10-8-8-5-0-0 for the above items, respectively; it means that
pulse height information is expressed with 10 bits, namely 1024 ch, and the X-axis information
3

Sometimes the important data (especially, rise time information) to reject the non-X-ray events are notincluded in the telemetry data, depending of the selected GIS bit assignment; in that case, the rejection of
non-X-ray events at ground is not so eﬀective.
4

The earth surface which is illuminated by the sun (at day).

5

The high Galactic latitude fields where no bright X-ray sources are seen.

6

The same as day earth, but at night.

37

GIS2+3 NXB (10-3 counts s-1 cm-2 keV-1 sensor-1)

3.1. ASCA

0.6 - 1.0 keV

1
0.9
0.8

1.0 - 2.0 keV

0.7
0.6

2.0 - 4.0 keV

0.5
4.0 - 8.0 keV

0.4
0

5

10

15

20

25

Month from June 1993

Figure 3.9: Change of GIS NXB fluxes with several energy bands since the launch (Ishisaki 1997)
NXB shows gradual increase since the launch.

is expressed with 8 bits, namely 256 ch, for example. For the standard bit assignment, 1 ch for
X- or Y-axis corresponds to 0.250 mm, and 1 ch for pulse height corresponds to ∼0.012 keV.
The time resolution of GIS depends on bit-rate and GIS bit assignment that the observer
selects. In high bit-rate, the best time resolution is achieved, whereas in medium and low bitrates, the resolution is worse by a factor of 8 and 32, respectively, than that in high bit-rate.
For standard bit assignment, the time resolution in high bit-rate is 65 msec. We can increase
the time resolution by changing the bit assignment, sacrificing other information, such as, risetime information.
The energy gain of GIS is the function of the temperature, the detected position, and thetime from launch, etc. Each of GIS sensor (GIS2 and GIS3) is equipped with the isotope 55 Fe
at the edge of the sensor; we can refer the peak of the spectrum of the calibration isotope tocorrect the gain. Thus we correct the raw gain to the linear scale one.
Details for GIS are found in Ohashi et al. (1996), Makishima et al. (1996).

3.1.5

SIS (Solid-state Imaging Spectrometer)

The Solid-state Imaging Spectrometer (SIS) is an X-ray CCD camera (Burke et al. 1991, 1994;
Yamashita et al. 1997; Yamashita 1998). ASCA has two SISs called SIS0 and SIS1, which areoperated in parallel. SISs exhibit better energy resolution (130 keV at 5.9 keV, at launch),

38

CHAPTER 3. INSTRUMENTS

higher detection eﬃciency in low energy band, higher rejection rate of the particle events, andbetter spatial resolution7 than GISs.
Each of SISs has 4 chips, each of which has 420×422 pixel with a pixel size of 27 µm.
Geometrical area of each chip is 11 mm×11 mm, corresponding to about 11′ × 11′ . All of the

SIS chips can be operated in parallel, although only the selected chips are used in most of theobservations to avoid telemetry saturation. The energy range of SIS is from 0.4 to 10 keV.
Although many kinds of operation modes for SIS are available, we used only the Faint and theBright modes in the observations of the GC region (see Burke et al. 1991, 1994). Hence wehere concentrate on these two modes.
The time resolution of SIS is 4, 8, or 16 seconds, depending on how many chips are used
for the observation. The events detected with SIS are often extended for two or more than
two pixels. We classify the shape of extension into eight categories (Grade 0–7: Fig. 3.11).
Nominally we regard Grade 0, 2, 3 and 4 events as X-ray events, whereas we regard others as
non-X-ray events.
The SISs are known to have the serious degradation since the launch due to the accumulated
radiation damage (Dotani et al. 1997; Yamashita et al. 1997; Yamashita 1998). It makes thelarge calibration uncertainty as well as the degradation of energy resolution, the detection
eﬃciency, and the increase of the hot/flickering pixels.

7

The pixel size of SIS, 27µm, corresponding to 0.′ 027, is the angular resolution of SIS, whereas the half power

diameter of XRTs is 3′ .

3.1. ASCA

39

~22.5 arcmin

Gap: ~0.5mm~18.5 pixels~29.4 arcsec

S0C0

S0C3

S1C2

S1C1

x

x

Satellite coordinates
Y

1
y

S0C1
420 pixels~11.3 mm
~11.1 arcmin

X

y

S0C2

4 0

S1C31P

1

2
2P

S1C0

~22.8 arcmin

y
x

DETX

1 pixel= ( 27 µ m )2
= ( 1.59 arcsec )2x
DETY
Detector coordinatesy422 lines~11.4 mm
~11.2 arcmin

Gap: ~0.13 mm~4.8 pixels~7.6 arcsec

1 1 CCD mode default position
2 2 CCD mode default position
4 4 CCD mode default position

Sun

1

P

2 P Old 1,2 CCD mode default position
0

SIS0 Optical Axis

1

SIS1 Optical Axis

Figure 3.10: The layout of the SIS chips.

40

CHAPTER 3. INSTRUMENTS[Definition][Examples]

Grade 0= perfect single
Grade 1= single+ detouched corners
Grade 2= vertical single-sided split+ detouched corners
Grade 3= left single-sided split+ detouched corners
Grade 4= right single-sided split+ detouched corners
Grade 5= single-sided split
with touched corners
Grade 6
= L-shape or square-shape+ detouched corners

The center pixel.
A pixel whose PH level is larger than the split threshold andwhich is included when summing up the PHs.
A pixel whose PH level is larger than the split threshold andwhich is not included when summing up the PHs.

Figure 3.11: Grade definition and the rule to add pulse heights. The events out of these criteria aredefined as Grade 7.

3.2. CHANDRA

3.2

41

ChandraIn this section, we introduce the property of the Chandra satellite. The large portion of descriptions, figures, and tables are referred to the Chandra Observatory Guide v.3(http://asc.harvard.edu/udocs/docs/docs.html).

3.2.1

Overview

Chandra was launched on the Space Shuttle Columbia on July 23, 1999. The orbit of thesatellite is elliptical with the perigee altitude of 10,000 km, the apogee altitude of 140,000 km.
The orbital period is about 64 hours, which is allowing uninterrupted observing intervals of
more than 48 hours.

Figure 3.12: The schematic view of the Chandra satellite
The most vivid characteristic of Chandra is higher angular resolution than all the previous
X-ray satellites. The Chandra telescope, “High Resolution Mirror Assembly (HRMA)”, consistsof 4 mirrors of Wolter type 1 geometry formed by glass, and achieved the angular resolutionof 0.′′ 5. The satellite attitude is controlled and determined by “Pointing Control and Aspect
Determination (PCAD)” system with the higher accuracy than angular resolution.
Chandra has two focal plane instruments: a microchannel plate, named “High Resolution
Camera (HRC)”, and a X-ray CCD camera, “Advanced CCD Imaging Spectrometer (ACIS)”.
Chandra also has two gratings: “High-Energy Transmission Grating (HETG)”, and “LowEnergy Transmission Grating”. To protect the instruments from the particle radiation damage,42

CHAPTER 3. INSTRUMENTS

Chandra is equipped with a charged-particle detector, the “Electron, Proton, Helium INstrument (EPHIN)”. The schematic view of the satellites is shown in Figure 3.12. The detailed
properties of the instruments are mentioned in following subsections.

3.2.2

PCAD (Pointing Control and Aspect Determination) System

Aspect control
The satellite attitude is controlled by the Reaction Wheel Assembly (RWA). Six wheels aremounted in a pyramidal configuration, and adjust the angular momentum. To unload accumulated momentum by external torques (e.g. gravity gradient, magnetic force), thrusters areused, which is named the Momentum Unloading Propulsion System (MUPS). Reaction Control
System (RCS) also consists of thrusters which change the spacecraft attitude in powered flight,
RCS maneuver, and RCS safe sun modes. PCAD modes and concerned instruments used ineach mode is summarized in Table 3.1.

Table 3.1: PCAD Modes and Concerned Instruments Used in each Mode
Mode
Standby

Sensors—

Control Description—
On-board computer commands to RWA,
RCS, and Solar Array Drive Assembly(SADA) disabled, for initial deployment,
subsystem checkout, etc.
Normal Pointing
IRU, ACARWA
Point at science target, with optional
ditherNormal Maneuver
IRU
RWA
Slew between targets at peak rate of 2 ◦per minute
Normal Sun
IRU, CSS, FSS RWA
Acquire sun and hold spacecraft −Z axis
and solar arrays to the sun
Powered FlightIRU
RCS
Control Chandra during Liquid Apogee
Engine burns
RCS Maneuver
IRU
RCS
Control Chandra using the RCS
Safe Sun
IRU, CSS, FSS RWA
Safe mode: acquire sun and holdspacecraft −Z axis and solar arrays tothe sunDerived Rate Safe Sun IRU, CSS, FSS RWA
Similar to Safe Sun Mode, but using
only one gyro (two axes) plus sun sensor
data
RCS Safe Sun
IRU, CSS, FSS RCSSame as Safe Sun Mode, but using RCS
instead of RWA for control

3.2. CHANDRA

43

Aspect monitor
To determine the accurate pointing of the satellite, Chandra has two kinds of aspect monitor
system. One is Inertial Reference Units (IRU), and the other is constructed with Aspect Camera
Assembly (ACA), Fiducial Light Assemblies (FLA), and Fiducial Transfer System (FTS).
IRU consists of two 2-axis gyroscopes, and can measure angular rates about 4 axes. Since
two IRUs are boarded, 8 axes can be measured in total. All 8 axes are in diﬀerent directions
and no three axes lie in the same plane. However, only four axes data can be read at one time.
ACA is a CCD camera, which is placed on the side of the HRMA, to track optical stars.
Simultaneously, fiducial lights irradiated from the LEDs at the science instrument (SI) module
passes through FTS, and is also imaged onto the ACA CCD (Figure 3.13). ACA can track
up to 8 images, which are typically 5 stars and 3 fiducial lights. Only star images are used tocontrol the attitude. Fiducial lights are used in ground processing to determine the accurate
position by comparison with the star images.

Figure 3.13: Aspect determination subsystem of ACA, FLA, and FTS. Optical stars and the fiducial
light from the SI module are monitored by the CCD camera.

In addition to aspect monitors, there are two sun sensors: Coarse Sun Sensor (CSS) to cover
4π steradian, and Fine Sun Sensor (FSS) for higher resolution of 0.◦ 02. Earth Sensor Assembly

44

CHAPTER 3. INSTRUMENTS(ESA) is used during the orbital insertion phase. These sensors are not used in processing
X-ray images.

Dithering
HRMA exhibits very high spatial resolution, which is comparable to the pixel size of focal
instruments (section 3.2.3). To avoid X-rays from a point source falling on a limited number ofpixels, the line-of-sight is dithered continuously in Normal Point Mode. Dithering distributes
photons over many detector elements (microchannel pores or CCD pixels), and serves several
purposes: reduces uncertainty due to pixel to pixel variation in quantum eﬃciency (QE); allows
sub-sampling of the image, and; provides some exposure in the gaps between the CCD.
The dither pattern is a Lissajous figure with an peak-to-peak amplitude of 16′′ for ACIS,
and 40′′ for HRC. This pattern is corrected in ground processing by using the data of aspectmonitors. The default dither periods are 706.96 sec (pitch) and 999.96 sec (yaw). Due to thedithering, point sources near the edge or the gap of the CCD chips observed intermittently.
These sources could exhibit fake periodicity at these periods.

Performance
The uncertainties of the pointing and the reconstructed image are shown in Table 3.2. Absolute
celestial pointing and PCAD 10-second pointing stability show the accuracy of the satellite attitude (the direction of the optical axis of HRMA) in the observation time. Then the images arereconstructed in ground processing, and each X-ray position is determined with the uncertaintyof less than 0.′′ 3 (RMS diameter). Celestial location accuracy measures the absolute accuracyof the point sources with accurately known coordinates.

Table 3.2: Aspect System Performance of Chandra
Description
Absolute celestial pointing
3.′′ 0 (99.0%, radial)
PCAD 10 sec pointing stability pitch 0.′′ 043 (RMS, 1-axis)
yaw 0.′′ 030′′
Image reconstruction
0. 3 (RMS diameter)
Celestial location
0.′′ 6 (RMS radius)3.2. CHANDRA

3.2.3

45

HRMA (High Resolution Mirror Assembly)

The Chandra telescope has four pairs of Wolter Type-I mirrors, which were fabricated fromZerodur glass, polished, and coated with iridium on a binding layer of chromium (Figure 3.14).
The outer diameters of mirrors are 1.23, 0.99, 0.87, and 0.65 meter, and the focal length is
10.066 meter.

Figure 3.14: HRMA mirrors

Figure 3.15a,b show the eﬀective area of HRMA in relation to X-ray energy and oﬀ-axis
angle, respectively. On-axis eﬀective area is about 800 cm 2 at 0.25 keV, and declines to 400 cm2
at 5.0 keV, 100 cm2 at 8.0 keV. The eﬀective area also decreases as source position departs
from on axis.
The point-spread function (PSF) is very sharp, and most of X-ray photons are collected
within 1′′ -radius region (Figure 3.16). The encircled energy function of HRMA is shown in
Figure 3.17a. At 4.51 keV, the 50% photons are included in about 0.′′ 35-radius, which is called
the encircled energy radius. The encircled energy radius increase, or PSF becomes broad, atthe larger oﬀ-axis angle. Encircled energy radius is shown in Figure 3.17b. The diameter, which
encircles 50% energy, is called Half Power Diameter (HPD).
The central 29′ diameter region of the field of view is free from stray light, because baﬄes
prevent non-reflected or singly reflected X-rays. Then, there is no ghost image.

46

CHAPTER 3. INSTRUMENTS(a) Eﬀective area vs. X-ray energy(b) Eﬀective area vs. oﬀ-axis angle

Figure 3.15: The HRMA eﬀective area versus X-ray energy (left; a) and oﬀ-axis angle (right; b). In
the left panel, shell 1, 3, 4, 6 indicate the mirror layers from outside to inside.

Figure 3.16: Simulated on-axis images with the energies of 1.49, 4.51, 6.40, and 8.63 keV. Surface
brightness contours are logarithmic and spaced in factors of 3. The innermost contour is at 90% of
the peak brightness. The 8.6 keV image core is oﬀ-center due to the shell misalignment.

3.2. CHANDRA

47(a) Encircled energy function(b) Encircled energy radiusFigure 3.17: Encircled energy radius of HRMA

3.2.4

ACIS

The Advanced CCD Imaging Spectrometer (ACIS) is comprised of ten 1024 × 1024 pixel CCDs;

2 × 2 array of ACIS-I (I0–I3) for imaging and 1 × 6 array of ACIS-S (S0–S5) for imaging andgrating spectroscopy. The pixels are 24 µm (∼ 0.′′ 5) square. Two CCDs (S1 and S3) are back-illuminated (BI) and the others are front-illuminated (FI). The layout of CCD chips is shownin Figure 3.18 and Figure 3.19. The CCD chips of ACIS-I and ACIS-S are tilted to follow the
HRMA focal surface and the HETG focal surface, respectively.
ACIS has many observing modes to apply variety observations, but we use only TimedExposure mode with the readout time of 3.24 s. X-ray events are extracted by the same way as
ASCA or grade method. Though grade identification number is diﬀerent from that of ASCA,
X-ray events which are equivalent to ASCA grade 0, 2, 3, 4, 6 are used in standard analysis.
Figure 3.20(a) shows the energy resolution of CCDs before launch. The energy resolutionof FI chips are near the theoretical limit, while BI chips exhibit poorer resolution. After thelaunch, the CCD chips are damaged with low energy protons, reflected by X-ray telescope
onto the focal plane during radiation belt passages. Low energy protons deposit there energy
in the buried channels at the HRMA side of the FI chips, and increase the Charge Transfer
Ineﬃciency (CTI), which leads the degeneration of the energy resolution. Hence the energy
resolution becomes a function of the row number shown in Figure 3.20(b). Since the gate
structure of BI chips are at opposite side to the HRMA, the energy resolution of BI CCDs

48

CHAPTER 3. INSTRUMENTS

ACIS FLIGHT FOCAL PLANE
not constant with Z~22 pixels ~11"22 pixels
= 11"

I0

I1

w203c4r
0

w193c2
1x

I2
w158c4r

w215c2r

2

3

S0

S1

S2

w168c4r
4

w140c4r
5

w182c4r
6}

I3

ACIS-I(aimpoint on I3 = (962, 964))

330 pixels = 163"

S3+
w134c4r
7

S4S5

w457c48

w201c3r
9}

ACIS-S(aimpoint on S3 = (252, 510))+∆Y
18 pixels = 8".8

Image Region

Bottom.

column

three

two

one

Top

BI chip indicator

node zero

+Z

.

row

TargetOffsetCoordinates+Z
Pointing+∆Z

Coordinates-Z

Pixel (0,0)+Y

Sim Motion

Frame StoreCCD Key

NodeDefinitions

Row/Column
Definition

Coordinate
Orientations

Figure 3.18: A schematic overhead view of the ACIS focal plane; the legend of the terminology is
given in the lower left. Nominal aimpoints of ACIS-I and ACIS-S are shown by ’x’ and ’+’ mark,respectively.

3.2. CHANDRA

49(a) ACIS-I(b) ACIS-S

Figure 3.19: A schematic perspective view of the layout of (a) ACIS-I and (b) ACIS-S; note thatvertical axes are not to scale. The aimpoints are indicated in ’+’ marks.

remains at their prelaunch values. In the present operation, the ACIS is not left at the focal
position during radiation belt passages, and no further degradation in performance has beenencountered. By decreasing the CCD temperature from −90 ◦ C (before damaged) to −120◦ C,

the CTI eﬀect is reduced as possible.

The quantum eﬃciency of ACIS is shown in Figure 3.21. Low-energy X-rays are largely
absorbed by Optical Blocking Filter (OBF) and the gate structure of the CCD chips. Since
BI chips have the gate at the opposite side to HRMA, the quantum eﬃciency of BI CCDs
are larger than those of FI CCDs in low-energy band. By the increase of CTI, the quantum
eﬃciency becomes smaller at the farther side from the readout.

3.2.5

Other Instruments — HRC, HETG & LETG

HRC is a microchannel plate comprised of HRC-I for imaging and HRC-S for grating spectroscopy with LETG. The features are a large field of view (30′ × 30′ ; HRC-I), faster time

resolution (∼ 16µs), and higher angular resolution (FWHM ∼ 20µm). However, energy resolu-

tion is poor with ∆E/E ∼ 1 at 1 keV.

HETG and LETG are designed for grating spectroscopy with ACIS-S or HRC-S. HETG is

used for higher energy range of 0.4–10.0 keV, and LETG is for lower energy of 0.07–7.29 keVwith HRC-S and 0.2–8.86 keV with ACIS-S. Their characteristics are high energy resolving

50

CHAPTER 3. INSTRUMENTS(a) prelaunch energy resolution(b) energy resolution after damaged

Figure 3.20: (a) The prelaunch energy resolution of ACIS FI chips (solid lines) and BI chips (dashedand dotted lines). (b) The energy resolution of proton damaged ACIS CCDs (I3 and S3) as a functionof row number. The energy resolution of FI chip (I3) is shown by data points (diamonds, open circles,
filled circles), and that of BI chip is shown by lines (dashed, dash-dotted, solid). These data weretaken at −120 ◦ C.

Figure 3.21: The quantum eﬃciency of the ACIS CCDs. Note that values are including the transmission rate of the OBF.

3.2. CHANDRA

51>
power of ∆E/E <∼ 1000 (1000 at 1 keV) for HETG, and ∆E/E ∼ 1000 between 0.07–0.15 keV
for LETG. Since the resolution of wavelength is constant in each instrument, the resolving
power becomes lower at higher energy band.
These instruments are not used in this thesis. For our objective of investigating diﬀuse
structures in the GC region, ACIS-I is the best instrument with the moderate energy resolution
and suﬃciently good spatial resolution.

52

CHAPTER 3. INSTRUMENTS

Chapter 4
Observations
4.1

ASCA

With ASCA, we observed the GC region over 100 times. Beginning with the mapping observations of central 1◦ × 1◦ region in PV (Performance Verification) phase, we almost cover 4 ◦ × 4◦of the GC region till AO-7 (7th Announcement of Opportunity) phase (see exposure map in
Fig 4.1; also refer to Sakano 2001). In this thesis, we analyze the data of the regions which
include the giant molecular clouds: Sgr B2, Sgr C, Radio Arc (M0.11−0.08), and Clump2.
Most of these regions are observed in many times, and the exposure time of each region is
generally large with about 100 ksec in total. However, the Sgr C region is observed only about
20 ksec in PV phase. The exposure time, the observing date, and the pointing coordinate ofeach observation are shown in Table 4.1.
GISs are operated in the PH mode with the standard bit-assignment (10-8-8-5-0-0; see theprevious chapter). While SIS mode varies from 4-CCD mode in the early observations to 1CCD mode in the later observations, because SISs are largely damaged with charged particle
radiations. In addition, the SISs are smaller FOVs and poorer quantum eﬃciency than GISs.
We mainly use GIS data in this thesis.
The data were post-processed to correct for spatial gain non-linearity. We removed datataken at low Earth elevation angles, at low geomagnetic cutoﬀ rigidities, or during the passage
through the South Atlantic Anomaly. For the SIS, we also excluded the data taken at elevation
angles from the bright Earth rim less than 25◦ .

53

54

CHAPTER 4. OBSERVATIONSExposure map(b (deg))+2°00′+1°00′+0°00′-1°00′-2°00′(l (deg))+2°00′+1°00′+0°00′-1°00′-2°00′

Figure 4.1: Exposure map of the ASCA observations in the GC 5◦ × 5◦ region with the galactic

coordinates (Sakano 2001). The gray levels are log(EXP[sec]) of 3.0, 3.5, 4.0, 4.5, 5.0, 5.5. The
position of the calibration source for GISs is removed from each pointing. Sgr B2 and RadioArc is in the darkest region near the center of the image, and Sgr C is in the second darkest
region at the right of the center. Clump2 (l II ∼ 2.◦ 6) is out of the figure.

4.2. CHANDRA

55

Table 4.1: Observational log of ASCA
ID

1
2
3
4
5
6

Date
Start
(UT)‡
1993/10/01-21:14
1993/10/04-02:11
1994/09/15-22:01
1994/09/22-03:44
1994/09/24-02:03
1999/09/06-02:19

End
(UT)‡
1993/10/02-10:25
1993/10/04-16:30
1994/09/17-14:00
1994/09/23-12:10
1994/09/24-14:41
1999/09/08-18:12

CoordinateslIIbII(deg.) (deg.)
0.49 −0.09
359.440.09
359.99 −0.06
0.54 −0.02
0.54 −0.02
2.61
0.01

Exp. †

Name(ksec)
17.2
19.8
81.9
58.1
19.7
77.1

GC(0.37,−0.37) §GC(−0.37,0.37) §SGR A §
SGR B N1
SGR B N2
Clump2†: The average of the exposures with GIS2 and GIS3 after the screening.
‡: (Year/Month/Date-Hour:Minute).

§ : Observations ID 1–3 include Sgr B2, Sgr C, and the Radio Arc region, respectively.

4.2

ChandraChandra observed the GC region 34 times with ACIS-I: 4 pointing observations for particular
objects (Sgr A∗, Sgr B2, Radio Arc, 1E1740.7−2942) and 30 mapping observations of the
GC survey. Exposure time of a each observation for the GC survey is about 10 ksec, while
pointing observations have larger exposure time of about 50–100 ksec except for 1E 1740.7−2942(∼ 9 ksec). The observation log is shown in Table 4.2.
Figure 4.2 shows the schematic figure of the FOVs of Chandra, and Figure 4.3 shows the
exposure map. GC survey covers −1◦ ≤ lII ≤ 1◦ , −0.◦ 4 ≤ bII ≤ 0.◦ 4 with about 10 ksec exposure.

FOVs of GC survey are overlapping on the Galactic plane (−0.◦2 ≤ bII ≤ 0.◦ 2), hence 20 ksec

exposure is made in these region. All observations are made in Timed-Exposure mode.

Observed data are screened in a standard criteria to remove flaring pixels and artificial
stripes caused probably by hot pixels in the frame-store region and by particles which hit on
the CCD node boundaries. The standard Good Time interval (GTI) filter is also applied tothe events. Then we use a standard ACIS grade set of ASCA grade 0, 2, 3, 4, 6 to distinguish
X-ray events from non X-ray background (NXB) events.

56

CHAPTER 4. OBSERVATIONS

Table 4.2: Observational log of Chandra
ID

1
2
3
4
5
6
7
8
9
10
11
12
13
14
15
16
17
18
19
20
21
22
23
24
25
26
27
28
29
30
31
32
33
34

Date
Start
(TT)‡
1999/09/21-02:52
2000/03/29-09:44
2000/07/07-19:08
2000/08/30-16:59
2001/07/16-02:15
2001/07/16-05:35
2001/07/16-08:44
2001/07/16-11:52
2001/07/16-15:01
2001/07/16-18:09
2001/07/17-14:11
2001/07/17-17:51
2001/07/17-21:19
2001/07/18-00:48
2001/07/18-04:16
2001/07/18-07:45
2001/07/18-11:13
2001/07/18-14:25
2001/07/18-17:37
2001/07/18-20:49
2001/07/19-00:01
2001/07/19-03:21
2001/07/19-06:41
2001/07/19-10:04
2001/07/20-04:37
2001/07/20-08:04
2001/07/20-11:15
2001/07/20-14:44
2001/07/20-18:12
2001/07/20-21:41
2001/07/21-01:09
2001/07/21-04:35
2001/07/21-08:03
2001/07/21-11:32

End
(TT)‡
1999/09/21-16:53
2000/03/30-13:53
2000/07/08-09:14
2000/08/30-20:00
2001/07/16-05:35
2001/07/16-08:44
2001/07/16-11:52
2001/07/16-15:01
2001/07/16-18:09
2001/07/16-21:21
2001/07/17-17:51
2001/07/17-21:19
2001/07/18-00:48
2001/07/18-04:16
2001/07/18-07:45
2001/07/18-11:13
2001/07/18-14:25
2001/07/18-17:37
2001/07/18-20:49
2001/07/19-00:01
2001/07/19-03:21
2001/07/19-06:41
2001/07/19-10:01
2001/07/19-13:25
2001/07/20-08:00
2001/07/20-11:10
2001/07/20-14:38
2001/07/20-18:07
2001/07/20-21:35
2001/07/21-01:04
2001/07/21-04:32
2001/07/21-08:03
2001/07/21-11:32
2001/07/21-15:03

CoordinateslIIbII(deg.) (deg.)
359.94 −0.05
0.59 −0.02
0.14 −0.10
359.11 −0.11
0.77 −0.20
0.78
0.00
0.77
0.19
0.58 −0.20
0.58
0.00
0.58
0.19
0.39 −0.20
0.39
0.00
0.39
0.19
0.19 −0.20
0.19
0.00
0.19
0.19
0.00 −0.20
0.00
0.00
359.990.19
359.80 −0.20
359.800.00
359.800.19
359.61 −0.20
359.610.00
359.610.19
359.41 −0.20
359.410.00
359.410.19
359.21 −0.20
359.220.00
359.220.19
359.02 −0.20
359.020.00
359.02
0.19†: The exposure time of ACIS

Exp. †

Name(ksec)
48.72098.98949.3799.357
10.746
10.55010.56010.56010.56010.54411.93911.75411.76011.76011.76311.76010.76210.76210.76210.762
11.261
11.261
11.261
11.63210.95010.75811.76011.74711.76311.76311.76311.76311.76311.763

SGR A ∗
SGR B2GALACTIC CENTER ARC
1E 1740.7−2942
GCS 1
GCS 2
GCS 3
GCS 4
GCS 5
GCS 6
GCS 7
GCS 8
GCS 9
GCS 10
GCS 11
GCS 12
GCS 13
GCS 14
GCS 15
GCS 16
GCS 17
GCS 18
GCS 19GCS 20
GCS 21
GCS 22
GCS 23
GCS 24
GCS 25
GCS 26
GCS 27
GCS 28
GCS 29GCS 30‡: (Year/Month/Date-Hour:Minute). The unit is Terrestrial Time(see http://tycho.usno.navy.mil/systime.html).

4.2. CHANDRA

57

1+1°00′

0.75
0.5+0°30′

0.25
0

-0.25

3

6

9

12

15

2

5
SgrB2

8

11

14

SgrA*

18

21

24

27

30

17

20

23

26

29

GC Arc1

4

7

10+0°00′

1E1740

13

16

19

22

28

25

-0.5

-0°30′-0.75-1+0°30′+0°00′-0°30′-1°00′-1°00′-1-0.75 -0.5 -0.25 0 0.25 0.5 0.75 1
Figure 4.2: The FOVs of Chandra at the GC region. Each FOV is drawn in 12 -diameter circle.
′

Note that the real field of ACIS-I is a inscribed square of the circle. Label 1–30 shows theobserved field of GCS 1–30, and Sgr A∗ , Sgr B2, Galactic center Arc (GC Arc), 1E1740.7−2942(1E1740) are indicated in their own names.

58

CHAPTER 4. OBSERVATIONSExposure map(b (deg))

2000+1°00′

1750
1500+0°30′

1250
1000+0°00′

750-0°30′

500
250-1°00′+1°00′(l (deg)) 250+0°30′+0°00′-0°30′-1°00′

500 750 1000 1250 1500 1750 2000

Figure 4.3: The exposure map of the Chandra observations in the GC region. The observed
time of each observation is shown in gray scale with log of 103.5 –105.5 sec.

Chapter 5
Results I. — 6.4-keV line emissions
At first, we study about the origin of the 6.4-keV line in the GC region. Koyama et al. (1996)
already indicated that 6.4-keV line emissions are coincident with the giant molecular clouds(Figure 2.9). Then we investigate the X-ray spectra and images of the heavy molecular clouds:
Sgr B2, Sgr C, M0.11−0.08 (at the Radio Arc region), and Clump 2. Most of the clouds areobserved with both ASCA and Chandra, but Clump 2 is observed only with ASCA.

5.1

Sgr B2

Sgr B2 is one of the heaviest giant molecular clouds in the Galactic center region, and the mostluminous in the 6.4-keV line band (see section 2.3, 2.4).

5.1.1

ASCA Observations

Two observations of Sgr B2 were made with ASCA on October 1 1993, and on September 22–24
1994. In these observations, we analyze both GIS and SIS data. The net observed time is about
95 ksec for GIS and 85 ksec for SIS (section 4.1). Since no time variability is found between
the two observations, we analyzed the combined data. The results are already published in
Murakami et al. (2000).
6.4-keV line image
Figure 5.1 shows the X-ray images in narrow energy bands with a central energy of 6.4 keV
and width of twice the energy resolution (FWHM): 5.8–7.0 keV for the GIS and 6.2–6.6 keV
for the SIS. We found a distinct peak near the center of the molecular cloud. The position
59

60

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONSof the X-ray peak is determined as R.A. (2000) = 17h 47m 20.s 5, Dec (2000) = −28◦ 24′ 20′′(l = 0.◦ 6496, b = −0.◦0470) and R.A. (2000) = 17h 47m 20.s 3, Dec (2000) = −28◦ 24′ 13′′(l = 0.6509, b = −0.0454), for GIS and SIS respectively. The statistical errors are 10′′ for the

GIS and 15′′ for the SIS.

The SIS image is laid over the radio intensity contours of the CH 3CN line (Bally et al.
1988). The X-ray peak is shifted from that of the radio distribution by ∼ 1′ .2 to the Galactic

center side.

Galactic Longitude (degree)
0.8
0.6
0.4
0.2

0.1

Galactic Latitude (degree)

Galactic Latitude (degree)

0.2

0.0
-0.1

-0.2
0.8

0.5
0.4
0.6
0.7
Galactic Longitude (degree)

0.2

0.0

-0.2

0.3

(a): SIS image(b) GIS image

Figure 5.1: 6.4-keV line image of Sgr B2 obtained with ASCA. (a) SIS image laid over the CH 3CN line
contours (Bally et al. 1988). The 6.4-keV brightness distribution is shifted from the radio distributionby ∼ 1′ .2 to the Galactic center side (to the right in the figure). (b) The 6.4-keV line image with

GIS. The source and the background regions are shown by the solid circle and the dotted ellipse,
respectively. The dotted circle encloses the other X-ray bright spot (G0.570−0.018; see section 6.2),
which is excluded from the background region. The bright source at the boundary is a galactic binary
X-ray source (1E 1743.1−2843).

Spectrum
The GIS spectrum given in Figure 5.2 is obtained by summing the X-ray photons in 3′ -radius
circle around the X-ray peak of the GIS image. For the background spectrum, we used anelliptical region with the major axis parallel to the Galactic plane, excluding the region of
Sgr B2 (a 3′ -radius circle) and the other X-ray bright spot (the other 3′ -radius circle) at thewest of Sgr B2 (G 0.570−0.018; see section 6.2). The source and the background regions are5.1. SGR B2

61

Figure 5.2: The GIS spectrum of Sgr B2. The best-fit model is shown in solid line.
shown in Figure 5.1b with the solid circle and dotted ellipse, respectively.
In order to derive quantitative feature of the Sgr B2 X-rays, we fit the spectrum to aphenomenological model, a power-law model with a Gaussian line. We used the Morrison &McCammon (1983) cross section for the absorption. Due to large absorption at low energy, theavailable data to be fitted are in the 4.0–10.0 keV band. However the limited energy band andrather poor statistics do not allow us to constrain all parameters simultaneously. Therefore weassumed a power-law of (fixed) photon index 2.0 (same as Koyama et al. 1996). We obtained
the acceptable fittings with the best-fit χ2 of 30.9 (34 degrees of freedom; hereafter d.o.f). The
best-fit parameters are given in Table 5.1. The 6.4-keV line, as we expected, is very strong
with an equivalent width of 2.9 keV. The hydrogen column density is N H ∼ 8 × 1023 H cm−2 ,and the absorption corrected luminosity is ∼ 1035 erg s−1 .

We also made a SIS spectrum and fit it with the same model of the GIS. However, since
the limited photon statistics prevented us from determining the absorption depth, we fixed
the hydrogen column to the best fit value for the GIS. The results are given in Table 5.1.
The SIS results were found to be consistent with the GIS. The narrow field of view of SIS
severely constrains the selection of the background region for the SIS spectrum and increases
background-subtraction uncertainties. Here and after, we thus analyzed only the GIS spectrum

62

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

for more detailed studies.

Table 5.1: Best-Fit Results of the ASCA Spectrum of Sgr B2.
GIS
SIS

NH a
8.3+2.5−2.0
8.3 (fixed)

ContinuumΓb
FX c
2.0 (fixed) 15+3−2
2.0 (fixed) 18 ± 6

ECd6.38 ± 0.03
6.43 +0.02−0.03

Iron line
FX c
9.7 ± 1.0
9.7 ± 1.7a: Hydrogen column density [1023 H cm−2 ].

b: Photon index of a power-law model for the continuum spectrum.

EW e
2.9 +0.3−0.9
2.1 +2.1−0.8

LX f
1.1+0.7−0.5
1.4+0.7−0.9

χ2 /d.o.f.
30.9/34
35.1/45c : Flux (no correction of absorption) in the 4–10 keV band [10 −5 photons cm−2 s−1 ].
d : Center energy of iron line [keV].
e: Equivalent width of iron line [keV].

f: Absorption corrected luminosity in the 4–10 keV band [10 35 erg s−1 ].

To estimate the contamination of the 6.7-keV line from diﬀuse hot plasma in the Galactic
center region, we fit the spectrum to the power-law and two narrow Gaussian lines at 6.4 and6.7 keV. We found no significant line at 6.7 keV, with the flux less than 10% of the 6.4-keV
line. Thus the observed line profile is reproduced by the 6.4-keV line alone.
The observed hydrogen column of 8 × 1023 H cm−2 is at least 5 (nearly 10) times larger

than that of interstellar gas to the GC region (Sakano et al. 2000; see also section 5.5.1). Thismeans that the large absorption column is due to local gas near or at the Sgr B2 cloud.
We found a deep iron edge in the spectrum. We fit the spectrum allowing the iron column
density to be free. Then the iron column density is estimated to be 4×10 19 Fe cm−2 .

5.1.2

Chandra Observation

With Chandra, we can obtain highly resolved image, which arrows us to exclude the contributionof the point sources, and study detailed structure of the diﬀuse emission. The results are alreadypublished in Murakami, Koyama, & Maeda (2001).
6.4-keV line image
We make the 6.4 keV-line (6.15–6.55 keV) image to highlight the most detailed structure.
Using the procedure given in Baganoﬀ et al. (2001) and the algorithm of Ebeling, White &Rangarajan (2000), the image is flat-fielded, and is smoothed adaptively with circular Gaussian5.1. SGR B2

63

kernels. Figure 5.3 is the 6.4 keV line image overlaid on the radio intensity contours of the13

CO molecule (Sato et al. 2000). The global morphology of the 6.4-keV emission is a concave-

shape located at the south-west side (right hand side) of the molecular cloud pointing to the
Galactic center direction (right hand side). The diﬀuse emission is distributed at R.A. (2000)= 17h 47m 13s –27s , Dec (2000) = −28◦ 21.′3–24.′ 8 (3′ × 3.′ 5 rectangle shown by solid line in
Figure 5.3).

We search the point sources in this region by the “wavdetect” method with the 3–8 keVimage. 17 point sources are resolved in 3′ ×3.′5 region. The positions of the resolved point sources

are indicated by solid circles with diameters of six times of the HPD of PSF. No significant

excess of the point sources is found in this narrow band (6.15-6.55 keV) image, which indicatesthat the 6.4 keV line of normal point sources, unlike the diﬀuse emission, is not prominent.
The local excess near the center and upper-left of the molecular cloud are possibly due to H II
complexes, Sgr B2 Main at (l, b) ≃ (0.67, −0.04) and North at (l, b) ≃ (0.68, −0.03). These two
local excess occupy only 5 % flux in the total diﬀuse flux.
Spectrum
We first fit the combined spectrum of the resolved point sources. The spectrum can be reproduced by a thin thermal plasma model with the temperature of ∼ 8.5 keV. The best-fit

parameters are shown in Table 5.2. The large column density of hydrogen indicates that mostof these sources are located in the cloud. A detailed study of the point sources is made by
Takagi, Murakami, & Koyama (2001).

Table 5.2: Best-Fit Results of the Combined Spectrum of Point Sources Obtained with Chandra
NHa
kT b
FXc
LdX
χ2/d.o.f.
1.9+0.6
8.5+9.3
1.5+0.5
2.9+1.1
12.2/12
−0.4
−3.7
−0.2
−0.8
a : Hydrogen column density [10 23 H cm−2 ].
b : Temperature of a thin thermal plasma [keV].
c: Flux (no correction of absorption) in the 2–10 keV band [10 −5photons cm−2 s−1 ].

d: Absorption corrected luminosity in the 2–10 keV band [10 33erg s−1 ].

Figure 5.4 shows the spectrum of the diﬀuse X-ray emission by using the data from thesolid square and subtracting background taken from the dotted circle in Figure 5.3. We finda faint peak at 7.1 keV, which is the center energy of Kβ line of neutral iron. We then fit the64

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

Figure 5.3: Chandra ACIS image of the Sgr B2 region in the 6.4-keV line band (6.15–6.55 keV). The
contours show the radio intensity of the 13 CO molecule (Sato et al. 2000). The X-ray emission is
located at the south-west (right hand) side of the molecular cloud. The source and the backgroundregions are shown by solid and dotted lines, respectively.

5.1. SGR B2

65

spectrum with a phenomenological model, a power-law function of photon index 2.0, adding Kαand Kβ lines of neutral iron and absorption with the Ba&lucińska-Church & McCammon (1992)
cross-sections for the solar abundances (Feldman 1992). Iron abundance, which is sensitive tothe edge depth, is treated separately as a free parameter. The center energy and flux of ironKα line are free parameters, but those of Kβ are linked to the Kα parameters by fixing theratio to the laboratory values. Point source contamination is taken into account by adding thebest-fit model of the combined point-source spectrum.
The best-fit parameters are given in Table 5.3. We find that the iron Kα line energy is
6.38+0.02−0.01 keV, which is consistent with neutral irons of 6.40 keV. In order to examine whether
highly ionized irons are present or not, we add 6.70 keV line, Kα of He-like iron, and constrain
that the equivalent width is less than 0.07 keV, or at most 3 % of that of neutral irons. Column
density of neutral iron (NFe ) determined from the K-edge structure is ∼ 3.4 × 10 19 cm−2, while

NH determined from the low energy cutoﬀ is ≃ 8.8 × 1023 cm−2. Both are nearly ten times

larger than those of the Galactic interstellar absorption (Sakano 2000; see 5.5.1), hence are
attributable to the gas in Sgr B2. The ASCA results, in spite of its poor spatial resolution, areroughly consistent with the Chandra best-fit parameters, supporting that the contribution of
the point sources is almost negligible.

Figure 5.4: X-ray spectrum of Sgr B2 obtained with Chandra ACIS. The best-fit spectrum is shownin solid line.

66

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

Table 5.3: Best-Fit Results of the X-ray Spectrum of Sgr B2 obtained with Chandra.
NH a
8.8+2.0−1.5

NFeb
3.4+3.6−2.2

ContinuumΓc
FX d
2.0 (fixed) 12 +2−1

Ee
6.38 +0.02
−0.01

FX d
5.6 ± 0.5

Iron line
e
EW f
EC
+0.02
1.7 +0.2
7.04
−0.1
−0.01

FX d
1.1 ± 0.1

EW f
0.28 +0.02
−0.03

LX g
1.0+0.4−0.2

χ2 /d.o.f.
20.6/19a : Hydrogen column density [10 23 H cm−2 ].
b : Iron column density [10 19H cm−2 ].
c : Photon index of a power-law model for the continuum spectrum.
d : Flux (no correction of absorption) in the 4–10 keV band [10 −5 photons cm −2 s−1 ].
e : Center energy of iron line [keV].
f : Equivalent width of iron line [keV].
g : Absorption corrected luminosity in the 4–10 keV band [10 35erg s −1 ].

5.1.3

Discussion

The ASCA observations discovered a 6.4-keV line emission, and the Chandra observation clearly
detect diﬀuse X-rays from Sgr B2 by resolving underlying point sources. The diﬀuse X-rays are
emitted from the south-west half (right side in Figure 5.3) of the cloud with a concave-shapemorphology pointing to the Galactic center direction. The spectrum is characterized with thepronounced iron Kα line at 6.38+0.02−0.01 keV, heavy absorption at iron K-edge and at low energy.

Furthermore, Kβ line at 7.04 +0.02−0.01 keV is also discovered with Chandra. We examine the originof the diﬀuse X-rays based on these results.
Young stellar objects

The giant molecular cloud Sgr B2 is known to be one of the richest regions of ongoing starformation in the Galaxy (e.g., de Vicente et al. 2000). In fact, many compact H II regions aswell as maser sources have been found from the cloud. These sources trace the star forming site
in the cloud, and are found from the ridge running north to south (Figure 2.3). Since youngstellar objects (YSO), regardless low mass or high mass stars, are rather strong X-ray sources
with luminosity ranging from 1032−33 (massive YSO) to 1028−29 erg s−1 (low mass YSO) (e.g.,
Feigelson 1999), numerous unresolved YSOs in Sgr B2, if any, may mimic the diﬀuse X-rays.
Sidoli et al. (2001) also proposed the point source origin from the observations by BeppoSAX.
The spatial distribution of the diﬀuse X-rays, a concave-shape, however is largely diﬀerent

5.1. SGR B2

67

from that of the north-south ridge of star formation. The X-ray spectra of YSOs are generally
consistent with optically thin thermal plasma of temperature ranging from 1 keV to a few keV,with the 6.7 keV line of He-like irons present in some cases. The spectrum of the Sgr B2 diﬀuse
is completely diﬀerent from that of typical YSOs, hence a YSO origin is ruled out for the bulkof the diﬀuse X-rays. Partial contribution of YSOs may still be conceivable in the two local
enhancement at the cloud center and north, Sgr B2 Main and North (see Figure 5.3). We hence
examine the X-ray spectrum from these two regions and find no essential diﬀerence from thetotal diﬀuse spectrum. We only see a hint of a weak 6.7 keV line with the equivalent width of
0.10±0.07 keV. Consequently, we can safely conclude that no significant contribution of YSOs
is found in the spectrum and morphology of the diﬀuse X-rays of Sgr B2.
X-Rey reflection due to internal point sources
It may be conceivable that the resolved point sources in Sgr B2 irradiate the surrounding gas
and produce the reflected X-rays. These point sources, however, are uniformly distributed in
the cloud (see Figure 5.3), which disagrees with the concave-shape of the diﬀuse X-rays. The
integrated luminosity of all the resolved point sources is 3× 1033 erg s−1 , which is only 3% of
the luminosity of the diﬀuse X-rays. Therefore internal point sources can be ignored.
X-Ray Reflection Nebula model
The most plausible scenario is “X-ray Reflection Nebula” (XRN) model, as proposed in Koyama
et al. (1996).In this model, the molecular cloud is irradiated by an external X-ray source, andemit fluorescent and scattered X-rays. The large absorption, the strong neutral iron line, andthe shifted X-ray peak from the cloud core agree with the X-ray features irradiated heavy
molecular cloud. The 6.4-keV line is strong evidence of reflection, in particular. Since X-ray
7
emitting gas is very hot with the temperature of >∼ 10 K, iron atoms are immediately ionized
to He- or H-like state. Neutral irons are too cold to emit thermal X-rays by themselves, andcan emit only fluorescent X-rays.
Inoue (1985) calculated the equivalent width of fluorescent 6.4-keV line in three cases (Figure 5.5). If we observe only scattered and fluorescent X-rays, the equivalent width should be24−2
about 2 keV at the range of column density <∼ 10 H cm . The observed equivalent width of
1.71 keV from Sgr B2 is consistent to this case, and other two cases (fluorescent and primary
X-rays are observed simultaneously) cannot produce such a strong 6.4-keV line.
We examine this scenario in detail with numerical simulations in section 5.5.

68

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

Figure 5.5: The simulated equivalent width of fluorescent iron 6.4-keV line as a function of column
density of scattering matter. The simulation is made in three cases as follows. I. Primary source
can be directly observed. Both fluorescent X-rays and incident (not absorbed) X-rays are seen. II.
X-rays from primary source are also absorbed by surrounding matter. III. Primary source is hidden
or scattering matter is irradiated by an external source. Only fluorescent X-rays are observed.

5.2. SGR C

5.2

69

Sgr C

In the opposite side of the Sgr B2, there is the other giant molecular cloud at the Sgr C region.
The cloud is also heavy of about 106 M⊙ , and we call the cloud the Sgr C cloud in this thesis(section 2.3).

5.2.1

ASCA Observation

ASCA observed the Sgr C region on October 4 1993 in PV phase. The net observed time isabout 20 ksec (section 4.1). The results are already published in Murakami et al. (2001).
6.4-keV line image
We made a narrow energy band image with a central energy of 6.4 keV and a width of twice theenergy resolution (FWHM): 5.8–7.0 keV for the GIS. Figure 5.6 shows the GIS image overlaid
with the contour map of the density distribution of the cold cloud by the radio observation of
CS with the radial velocity of −120 km s−1 to −110 km s−1 (Figure 3 in Tsuboi, Handa, &Ukita 1999). The position accuracy is about 5′′ for the radio map, and 24′′ for the GIS image(Gotthelf et al. 2000) The 6.4 keV band image shows an X-ray peak (inner region of the solid
circle) near the molecular cloud core CO 359.4 + 0.0 (Oka et al. 1998), which is one of the giant
molecular clouds in the Sgr C complex, with a mass of ∼ 106 M⊙ . We therefore suspect that

the X-ray emission is associated with the molecular cloud Sgr C.
Spectrum

We made X-ray spectrum using the X-ray photons in a circle of 2′ .7 radius around the X-ray
peak (see Figure 5.6). The background region is selected as a 2′ .7 radius region with the same
Galactic latitude (dotted circle in Figure 5.6).
Figure 5.7 shows the background subtracted GIS spectrum. We fit the spectrum with apower-law continuum and a Gaussian line model. Due to the limited statistics, we fixed thepower-law photon index Γ to be 2.0, the same value as assumed for Sgr B2. The fit is acceptable,with the best-fit parameters shown in Table 5.4. The central energy of the line appears at
6.28+0.15−0.21 keV, in agreement with the fluorescent line from neutral irons. The line equivalent

23
H cm−2 ,width and absorption column density are both very large: ∼ 0.8 keV and 1.3+0.3−0.4 × 10

respectively.

Since the Galactic diﬀuse background may not be uniform, we checked the ambiguity of

70

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

Figure 5.6: 6.4-keV line image of Sgr C. The contour shows the radio intensity distribution of CS
with the radial velocity of −120 km s −1 to −110 km s−1 (Tsuboi, Handa, & Ukita 1999). The source

and the background regions are shown by solid and dotted lines, respectively.

5.2. SGR C

71

the background selection. We selected three other background regions with the same Galactic
latitude as the source region, and fitted them with the same model. The best-fit parameters of
the line central energy, line intensity, and the absorption are in the range of statistical errors.
However, X-ray flux of the continuum component diﬀers larger than the error. The 2.0–10.0 keV
flux varies 12–38 ×10−5 photons s−1 cm−2 . Thus the continuum flux has large uncertainty by

the background selection, however, since we are generally interested in the iron line flux, thisflux variation is not serious.

Adding two narrow Gaussian lines at 6.4 and 6.7 keV on a power-law continuum, we foundthat the 6.7-keV line from highly ionized iron is very weak; the upper limit of 9 ×10 −14

erg s−1 cm−2 is about 1/4 of the best-fit value of the 6.4-keV line. We thus conclude thatthe emission line from the Sgr C cloud is dominated by the fluorescent line from cold irons.
We also made an SIS spectrum from the same region of GIS, and fitted it with the samemodel as with the GIS. The best-fit spectral parameters are shown in 5.4, and are consistent
with those obtained from the GIS spectrum. Due to the limited statistics, the SIS spectrum
does not further constrain the spectral parameters.

Figure 5.7: The GIS spectrum of Sgr C. The best-fit model is shown by solid line

72

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

Table 5.4: Best-Fit Results of the X-ray Spectrum of Sgr C obtained with ASCA.
GIS
SIS
a

NH a
12.6+3.5−3.3
9.7+3.9−2.5

ContinuumΓb
FX c
2.0 (fixed) 30+9−6
2.0 (fixed) 20+7−5ECd
6.28+0.15−0.21
6.38+0.10−0.31

Iron line
FX c
3.5+1.4−2.2
2.7+2.3−1.5

EW e
0.8+0.4−0.5
1.1+1.1−0.7

LX f
5.0+1.3−1.0
3.0+0.9−0.8

χ2 /d.o.f.
16.1/20
8.9/14: Hydrogen column density [1022 H cm−2 ].

b : Photon index of a power-law model for the continuum spectrum.
c : Flux (no correction of absorption) in the 4–10 keV band [10 −5photons cm−2 s−1 ].
d: Center energy of iron line [keV].

e: Equivalent width of iron line [keV].

f : Absorption corrected luminosity in the 4–10 keV band [10 34 erg s−1 ].

5.2.2

Chandra Observation

Sgr C region is observed twice with Chandra as a part of the Galactic center survey (section
4.2). The net observing time is about 20 ksec.

6.4-keV line image
We first applied the “wavdetect” method to the 3–8 keV image to remove a contaminationof point sources, and detected 37 point sources in the 17′ × 17′ field. Since no point sources
are found around the molecular cloud, we simply exclude X-rays from point sources within

the diameter of twice of the PSF size, and smoothed with Gaussian function of σ = 16′′ toinvestigate a distribution of diﬀuse emissions. Figure 5.8 shows the smoothed image superposedwith radio intensity of the CS line (Tsuboi, Handa, Ukita 1999).
The diﬀuse X-ray emission is detected at (l =359.◦ 4–359.◦ 5, b = −0.◦ 1–0.◦ 0), which is almost

the same position as ASCA observation. The distribution of X-ray intensity has a crescentshape which is symmetric to Sgr B2 with respect to the GC, and at the Galactic center side of
the cloud core. Note that the radio image is sliced in radial velocity, and we cannot estimate
accurate separation between the X-ray peak and the cloud core.
There is another diﬀuse filamentary structure at (359.◦ 42, −0.◦ 13), which is near the faintemission detected by BeppoSAX (Sidoli et al. 2001). The nature of this emission is discussed

in section 7.2.2. Other three bright spots at (l, b) = (359.◦ 32, 0.◦ 06), (359.◦ 39, 0.◦ 01) and (359.◦ 55,
0.◦ 01) are contamination of bright point sources.

5.2. SGR C

73

Figure 5.8: Chandra image of Sgr C obtained with ACIS. The contour shows the radio intensity
distribution of CS with the radial velocity of −120 km s −1 to −110 km s−1 (Tsuboi, Handa, & Ukita

1999). The source and the background regions are shown by solid and dotted lines, respectively.

74

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

spectrum
We accumulated X-rays from the circular region indicated by the solid line in Figure 5.8, andmake a spectrum. No point source is detected in this region with the 3.0–8.0 keV band image.
The background spectrum is obtained from the other 2′ -radius circle at the same Galactic
latitude (dotted circle in Figure 5.8). The spectrum is shown in Figure 5.9. There are large
excesses at ∼ 6.4 and ∼ 7.1 keV, which are the center energies of Kα and Kβ lines of neutral

iron. We then fit the spectrum with a power-law and two narrow Gaussian lines with the fixedratio of the center energy and the intensity. We fixed a photon index of the power-law model
to be 2.0 as Sgr B2. The best-fit parameters are shown in Table 5.5.
The strong neutral iron is detected with the equivalent width of ∼ 3.2 keV. The absorptionof 13.8 × 1022 H cm−2 is slightly larger than the Galactic interstellar absorption (Sakano 2000;

see section 5.5.1). X-ray flux is smaller than ASCA results. It would be attributable to smaller
size of source region (∼ 1/4) and diﬀerence in background region.
We further fit the spectrum adding 6.7-keV line. The equivalent width is constrained to be< 260 eV.

Table 5.5: Best-Fit Results of the X-ray Spectrum of Sgr C Obtained with Chandra.

ACIS

NH a
13.8 +14.2−5.9

ContinuumΓb
FX c
2.0 (fixed) 29 ± 13dEC6.38 +0.06
−0.29

FX c
14+6
−7

Iron linedEW e
EC
3.2 ± 1.5 7.04 +0.06
−0.32

FX c
2.0 ± 0.9

EW e
0.5 +0.3−0.2

LX f
6.4+7.8−3.6

χ2 /d.o.f.
10.4/13a : Hydrogen column density [10 22 H cm−2 ].
b : Photon index of a power-law model for the continuum spectrum.
c : Flux (no correction of absorption) in the 2–10 keV band [10 −6 photons cm −2 s−1 ].
d : Center energy of iron line [keV].
e : Equivalent width of iron line [keV].
f : Absorption corrected luminosity in the 2–10 keV band [10 33 erg s −1 ].

5.2.3

Discussion

Comparison with radio observations
In the radio continuum band, the most distinct features are H II regions and non-thermal filaments. However, no X-ray peak agrees with the radio filament. The filamentary structure in5.2. SGR C

Figure 5.9: Chandra spectrum of Sgr C with the best-fit spectrum in solid line.

75

76

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

Figure 5.8 is slightly shifted from the radio filament. While the distribution of the molecular
cloud agree with the 6.4-keV line emission.
Emission mechanism
We detected a diﬀuse X-ray emission from the molecular cloud near Sgr C. Strong neutral ironlines and large absorption column are confirmed. The X-ray image resembles Sgr B2; the peakis shifted to GC side. No point source is detected in the cloud, which clearly rejects the internal
point source origin. With the similarity of Sgr B2, the molecular cloud is a candidate for an
XRN.
The detailed verification of the XRN model is mentioned in section 5.5.

5.3

M0.11−0.08

M0.11−0.08 is a molecular cloud near the Radio Arc, and the nearest giant molecular cloudto the Galactic center. Koyama et al. (1996) already reported the detection of strong neutral
iron line around here with ASCA. In this section, we further study the X-ray spectrum anddistribution by long exposure observation with ASCA and Chandra.

5.3.1

ASCA Observation

After the discovery of the 6.4-keV line emission in the PV phase (Koyama et al. 1996), we
made follow-up observation of the Radio Arc region ∼ 80 ksec with ASCA (section 4.1).
6.4-keV line image
Figure 5.10 shows the 5.8–7.0 keV image in the Radio Arc region with 80 ksec exposure. Thereis a faint peak at (l, b) ≃ (0.◦ 1, −0.◦1) as shown by solid white circle, which agrees with theposition of the bright spot seen in the brightness distribution map of the 6.4-keV line with earlyobservation in Koyama (1996) (see Figure 2.9).This position is also coincident with M0.11−0.08.
However, the detailed comparison with the distribution of molecular cloud is made by using
the Chandra image in the next subsection.
spectrum
We make a spectrum with the source region of 3′ -radius circle. To exclude the contaminationfrom the Sgr A region, the background region is determined to be another 3′ -radius circle which

5.3. M0.11−0.08

77+0°10′
180

160+0°00′
140

120-0°10′
100

80-0°20′
60+0°20′
40+0°10′
60

80+0°00′
100

120-0°10′
140

160

180

Figure 5.10: ASCA GIS image of the Radio Arc region in 5.8–7.0 keV band. Bright spot is shown bysolid circle. The dashed line shows the background region. Two white regions are saturated with thebright X-ray sources: 1E 1743.1−2843 (left) and Sgr A region (right).

78

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

is symmetric in regard to the Galactic latitude of Sgr A ∗. The source and the background regions
are shown in solid and dashed lines respectively, in Figure 5.10. The spectrum exhibits strong
iron line at ∼ 6.4 keV (Figure 5.11).
We first fit the spectrum with a power-law and a Gaussian line with the free center energy
and width. Though the spectrum is roughly fitted with the best-fit values of photon indexΓ ∼ 1.9, the center energy of Gaussian EC ∼ 6.43 keV, and the equivalent width of ∼ 2.3 keV,
it is not acceptable with the χ 2 of 35.4 (24 d.o.f.). There are residuals at ∼ 1.9 keV (He-like Si),

∼ 2.5 keV (He-like S), and ∼ 3.1 keV (He-like Ar). These features indicate that the spectrum

contains thermal component. The width of Gaussian is also large with the σ of ∼ 300 eV,
which implies that the iron line is a blend of neutral and ionized states.

A single thin-thermal plasma model also cannot reproduce the spectrum with a large residual
at ∼ 6.4 keV and the χ2 of 50.9 (27 d.o.f.). Then we fit the spectrum with a thin-thermal plasma

model and a simple XRN model used for Sgr B2 and Sgr C (a power-law and two Gaussianlines of iron Kα and Kβ). To convergence the fitting, we fixed the abundance of the thermal
plasma to be 1.0, the photon index of the power-law to be 2.0 (same as Sgr B2), and equivalent
width of lines to be 2.0 for Kα as the expected value in Inoue (1985). The center energies andthe intensity ratio of iron lines are fixed to be the laboratory values.
At first we determine the parameters of thermal component with the spectrum of 1.0–3.5 keVband. Since XRN ordinary exhibits large absorption column and very weak X-ray emission in
low energy band, the thermal component is dominant in this band. We obtain acceptable
22
H cm−2 ,fittings with best-fit parameters and χ2 of kT = 2.8 (> 1.2) keV, NH = 5.1+2.9−2.1 × 10

and χ2/d.o.f = 15.4/10. Although there still left residuals at ∼ 1.9 keV and ∼ 3.1 keV, reduced
χ2 becomes worse if we fit the spectrum with the abundance variable.

We fixed the temperature and the absorption column to be these values, and fit the 1.5–
10 keV band adding the simple XRN model. The normalization of the thermal component istreated as variable. We then obtain acceptable fittings. Best fit values are shown in Table 5.6.
We confirmed the strong neutral iron line. Though the center energy of the line is slightly
smaller than 6.4 keV, it should be due to systematic uncertainty of detector or fitting method(aﬀected by the thermal component).
We thus conclude that M0.11−0.08 also emits strong 6.4-keV line with the equivalent widthof ∼ 2.0 keV, though there is another component of thin-thermal plasma.

5.3. M0.11−0.08

79

Figure 5.11: The X-ray spectrum of M0.11−0.08 obtained with ASCA. The solid line shows the bestfit spectrum, of which dotted line shows the thermal component and dashed lines show the simple
XRN model (power-law and Kα, Kβ lines of neutral iron).

80

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

Table 5.6: Best-Fit Results of the X-ray Spectrum of M0.11−0.08 obtained with ASCA GIS(thermal component)
NH a
5.1 (fixed)

kT b
2.8 (fixed)

Zc
1.0 (fixed)

FX d
34 +3
−4

LX e
3.1±0.4(simple XRN model)
Continuum
NH a54.7 +20.3
−17.3

Γf
2.0 (fixed)

FX d
8.4±2.1

Iron lines
g
EC
+0.08
6.28 −0.14KαFX d3.7+1.0−0.9

EW h
2.0 (fixed)

g
EC
+0.08
6.93 −0.16

total
Kβ
FX d
0.7+0.1−0.2
χ2 /d.o.f.

EW h
0.3 (fixed)
···

LX e
5.4 +1.4−1.5
33.6/26

a : Hydrogen column density [10 22 H cm−2 ].
b : The temperature of the plasma [keV].
c : Chemical abundance with ratio to solar abundance.
d : Flux (no correction of absorption) in the 2–10 keV band [10 −5 photons cm −2 s−1 ]. e : Absorption corrected luminosity in the

2–10 keV band [10 34 erg s −1 ].
f : Photon index of a power-law model for the continuum spectrum.
g : Center energy of iron line [keV].
h : Equivalent width of iron line [keV].

5.3.2

Chandra Observation

Chandra observed the Radio Arc region about 50 ksec in August 2000 (Section 4.2), which isincluding the molecular cloud M0.11−0.08.
Image
Figure 5.12 shows the 6.0–7.0 keV image obtained with Chandra, which is adaptively smoothedwith the same method as Sgr B2. We confirm the diﬀuse X-ray emission at (l, b) ≃ (0.◦ 1, −0.◦ 1).

In addition, several point sources and filamentary structures are also found. We extract the

point sources by the “wavdetect” method in the 3–8 keV image, and detect 186 point sourcesin 17′ × 17′ region.

The intensity distribution of CS line with the radial velocity of 20–30 km s −1 is laid on

the X-ray image. The diﬀuse emission is roughly coincident with the dense cloud. A larger
filamentary structure is at the edge of the GC side of the cloud (see section 7.2.1). Anotherfilament is located at the opposite side of the cloud (see section 7.1.2). These filaments areshown by dashed elliptical region in Figure 5.12. They are distinct in 3.0–8.0 keV raw (notsmoothed) image (Figure 7.1, 7.4b, 7.5b).

5.3. M0.11−0.08

81

Figure 5.12: 6–7 keV image around M0.11−0.08 obtained with Chandra. Point sources areshown by circles, and two filamentary structures are indicated by ellipses. Contour shows the
CS line intensity with the radial velocity of 20 km s −1 to 30 km s−1 (Tsuboi, Handa, & Ukita
1999). The source and the background regions are shown by solid and dashed lines.

82

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

spectrum
Since we revealed that the diﬀuse emission is coincident with the molecular cloud, we extract
X-rays from the rectangle region of 6′ × 4.′7 enclosing the cloud. The contamination of point

sources and filaments are excluded. Background region is surrounding rectangle of the source

region. The source and background regions are shown in Figure 5.12. We fit the spectrumwith the same model as ASCA. Since the energy resolution of ACIS is better than GIS, we
fit the spectrum with the parameters of a thermal and a simple XRN components variable.
We obtain acceptable fittings with χ2/d.o.f. of 26.7/25. The best-fit parameters are shown in
Table 5.7. Figure 5.13 shows the spectrum and best-fit model in solid line. The thermal and thesimple XRN components are shown in dotted and dashed lines respectively. The parametersof thermal component are almost the same as ASCA results. We verified that the assumption
in fitting analysis of ASCA spectrum is valid. We also confirmed that the simple XRN model
can reproduce the spectrum. The absorption column of XRN component is also large as otherclouds.
We also fit the combined spectrum of resolved point sources. The total luminosity of point
sources is ∼ 2 × 1034 erg s−1 , which is smaller than that of XRN component.
Table 5.7: Best-Fit Results of the X-ray Spectrum of M0.11−0.08 Obtained with Chandra.
(thermal component)
NH a
6.9+1.0−0.7

kT b
2.8+0.7
−0.5

FX d
30 +3
−2

Zc
1.0 (fixed)

LX e
3.4+0.4−0.3(simple XRN model)
Continuum
NH a23.5 +10.3−7.5

Γf
2.0 (fixed)

FX d
21±2

Iron linesgEC
6.36 +0.01
−0.03

Kα
FX d
7.7±0.8EW h
2.0 (fixed)

total
g

EC
7.01 +0.02
−0.03

Kβ
FX d
1.2+0.1−0.2
χ2 /d.o.f.

EW h
0.3 (fixed)
···

LX e
6.1 +0.7−0.6
26.7/25a : Hydrogen column density [10 22 H cm−2 ].
b : The temperature of the plasma [keV].
c : Chemical abundance with ratio to solar abundance.
d : Flux (no correction of absorption) in the 2–10 keV band [10 −5 photons cm −2 s−1 ]. e : Absorption corrected luminosity in the

2–10 keV band [10 34 erg s −1 ].
f : Photon index of a power-law model for the continuum spectrum.
g : Center energy of iron line [keV].
h : Equivalent width of iron line [keV].

5.3. M0.11−0.08

83

Figure 5.13: The X-ray spectrum of M0.11−0.08 obtained with Chandra. The best-fit spectrum isshown by solid line. The dotted line shows the thermal component, and the dashed lines show the
simple XRN spectrum (a power-law with two narrow lines of neutral iron Kα and Kβ).

84

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

5.3.3

Discussion

By the observations of ASCA and Chandra, we revealed that the X-ray spectrum of M0.11−0.08
exhibits strong 6.4-keV line. Though there is also a thermal component, fluorescent X-rays
should be emitted from the cloud. By using the simple XRN model, we can reproduce thespectrum. It is possible that M0.11−0.08 is also an XRN. This hypothesis is further verified
with numerical simulations in section 5.5.

5.4

Clump 2

Clump 2 is a large molecular cloud at (l, b) = (2.◦ 7–3.◦ 4, 0.◦ 0–0.◦ 8) (Stark&Bania 1986). The
molecular cloud consists of a lot of cores with the total mass of 3 × 106 M⊙ (section 2.3).

5.4.1

ASCA Observation

Maeda (1998) revealed the excess of 6.4-keV line in this field with ASCA Galactic plane survey.
The Galactic longitudinal distribution of 6.4-keV line decreases with the larger longitude, but
exhibits a bump at l ∼ 3◦ (Figure 5.14). This region includes the Clump 2 cloud.
The 6.4-keV line and the large molecular cloud are common features in XRN, then we madea follow-up observation with ASCA in 1999 July. We concentrate the results of the follow-up
observation in this section. The observed time is about 80 ksec (section 4.1).
6.4-keV line image
Figure 5.15 shows the 5.8–7.0 keV image. This field is located near GX 3+1, and the stray light◦
contaminates the upper-right part of the image (l >∼ 0. 0). There is no other distinct emission.
We then made the equivalent width map by the same method in Sakano (2000). Figure 5.16
shows the distribution of the equivalent width of the iron line. The excess is detected at (l, b) ≃(0.◦ 0, 2.◦ 7). The X-ray peak is located at the GC side of the main body of Clump 2, and rather
agrees with the envelope gas than a dense core (Figure 2.7). Photon statistics and the angularresolution is insuﬃcient to make a detailed comparison between X-ray and radio image.
spectrum
We extract X-rays from the 4′ -radius circular region, and subtracted a background spectrum
obtained from the annulus with the radius of 4′ –6′ (solid and dashed circles in Figure 5.16).

5.4. CLUMP 2

85

Clump2

Sgr B2

Sgr A*{
Figure 5.14: Intensity distribution of the iron neutral line in the GC region (Maeda 1998).

Figure 5.17 shows the spectrum. We fit the spectrum with a phenomenological model of apower-law and a Gaussian line. This model cannot reproduce the spectrum with the large
reduced χ2 of 1.85 (6 d.o.f.). Though we confirmed the strong iron line, we cannot limit theionization state with the large error of the center energy of 6.21–6.95 keV. The equivalent widthof the line is ∼ 5.9 keV. There is a residual at >∼ 7.0 keV. We also fit the spectrum with a thinthermal plasma model. The abundances of elements are varied with the fixed ratio to the solar
values. This model is rejected with the reduced χ 2 of 1.87 (7 d.o.f.). There is a large residual
at 6–7.5 keV, or the strong iron line can not be reproduced with the thermal model.
We then fit the spectrum with a power-law model and two narrow Gaussian lines withthe fixed center energies of iron Kα and Kβ (simple XRN model; see previous sections). The
center energies and intensity ratio of iron lines are fixed to the laboratory values. We obtainbetter fitting with the reduced χ 2 of 1.73 (7 d.o.f). This χ2 value is slightly larger than 90%
accepted region with the probability of 9.7%. If we fix the equivalent width of the iron line
to be 2.0 keV, we can obtain acceptable fittings with the best-fit values shown in Table 5.8.
Absorption column of ∼ 5 ×1022 H cm−2 is consistent with interstellar absorption (Sakano
2000; see section 5.5.1).

86

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

250+0°20′
200

150+0°00′

100

50-0°20′+3°00′
50+2°40′
100+2°20′
150

200

250

Figure 5.15: The 5.8–7.0 keV image of the Clump 2 region. The stray light of GX 3+1 is found inupper right.

.

5.4. CLUMP 2

87

200

180

+0°10′
160

140+0°00′
120

100-0°10′
80+2°50′+2°40′+2°30′

60

60

80

100

120

140

160

180

Figure 5.16: The distribution of the equivalent width of iron line on Clump 2 obtained with ASCA.

.

Gray image shows the equivalent width of 0.0–2.0 keV in linear scale.

88

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

Figure 5.17: The X-ray spectrum of the diﬀuse X-ray emission at Clump 2 region. The solid line
shows the best-fit spectrum with the simple XRN model.

Table 5.8: Best-Fit Results of the X-ray spectrum of Clump 2 obtained with ASCA GIS.

GIS

NH a
5.2 +22.9−3.7

ContinuumΓb
FX c
2.0 (fixed) 22±11Ed6.40 (fixed)

FX c
4.3 +2.2−2.1

Continuum
EW eEd
2.0 (fixed) 7.06 (fixed)a : Hydrogen column density [10 22 H cm−2 ].
b : Photon index of a power-law model for the continuum spectrum.
c : Flux (no correction of absorption) in the 2–10 keV band [10 −6 photons cm −2 s−1 ].
d : Center energy of iron line [keV].
e : Equivalent width of iron line [keV].
f : Absorption corrected luminosity in the 2–10 keV band [10 33 erg s −1 ].

FX c
0.6±0.3EW e
0.3 (fixed)

LX d
2.5 +2.1−1.4

χ2 /d.o.f.
12.4/8

5.4. CLUMP 2

5.4.2

89

Discussion

With 80-ksec observation with ASCA, we confirmed the diﬀuse X-ray emission which exhibitsa strong iron line. Due to the limited photon statistics, we cannot constrain the physicalparameters, such as ionization state of iron atoms. However, the equivalent width is too large
to be explained by a thermal plasma model in ordinal abundance. If it is neutral iron line,
large equivalent width is evidence for reflection. While, the absorption column is not larger
than interstellar absorption. Hence it is diﬃcult to conclude the origin of X-rays from Clump 2.
However, since the simple XRN model can reproduce the spectrum, we regard Clump 2 tobe a candidate of XRN. The verification of XRN model by numerical simulation is seen in thenext section.

90

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

5.5

Numerical Simulations of XRN Model

In section 5.1–5.4, we observed the molecular clouds in the GC region, and revealed the common
features as follows;
1. we confirmed the presence of the very strong emission lines of neutral iron,
2. we found larger column densities than interstellar absorption (except Clump 2),
3. we found that the X-ray emitting region is extended and coincident with the molecular
cloud. X-ray images of Sgr B2 and Sgr C are shifted to the GC side, in particular.
These peculiar X-ray spectra and distributions imply that the X-rays from these clouds are
attributable to Thomson scattering (continuum emissions), photo-electric absorption of neutral
iron atoms (low-energy cutoﬀ and iron K-edge), and fluorescence (6.4-keV line), produced byan irradiation of an external X-ray source in the direction of the GC. The clouds should be
the “X-ray Reflection Nebulae (XRNe)”. In this section, we simulate the images and spectraof XRNe, and verify the hypothesis by comparing with the observations. In this analysis, we
mainly use the Chandra results. However, we use the ASCA results for Clump 2, which isobserved only with ASCA.

5.5.1

Conditions of Simulations

The geometry and configuration for XRN simulations are given in Figure 5.18, with the following
assumptions.
(1) For the mass distribution, we adopt the result of 13 CO and the C18 O observations by
Lis & Goldsmith (1989). The model cloud is cylinder-shape with its axis parallel to our line ofsight.
The density distribution is shown as follows (in case of Sgr B2);
"

nH2
1cm−3

#

4

= 5.5 × 10$r1.25pc%−2+ 2.2 × 103 ,(5.1)

where nH2 is the number density of hydrogen molecules, and r is the projected distance fromthe center of the cloud. The thickness of the cloud is same as the radius. This formula gives
the total molecular mass to be 6.3 × 106 M⊙ in the 45-pc diameter region. For the other clouds

than Sgr B2, we assume the same density gradient as Sgr B2, and normalize the size and thedensity of the model cloud by using the observed mass and size of the cloud.

5.5. NUMERICAL SIMULATIONS OF XRN MODEL

91

Since the density distribution is too simplified, we modify the shape of the model cloud for
Sgr B2 to an elliptical cylinder to agree with the distribution of the radio intensity (e.g. Sato etal. 2000; see Figure 2.3). The ellipticity is 2, and its major axis is parallel to the line of longitude(−30◦ to the Galactic plane; see Figure 5.19). Other clouds are assumed to be cylindricallysymmetry, because X-ray images of these clouds are not simulated (see section 5.5.2).
(2) A primary source (irradiating source) is located in the GC-side with the normal angle
to the cloud (cylinder) axis. The spectrum of the primary source is a power-law with photon
index of 2 (Koyama et al. 1996).
For numerical simulations, we divided the model cloud into cells with same size, and calculated absorbed, reflected, and fluorescent X-rays in each cell as a function of incident X-ray
energy (E) from the primary source. The size (number) of the cell is determined by the required fineness of the simulation. Figure 5.18 shows an example of the simulation of the 45-pc
diameter cloud. When the cell size is 0.2 pc, the cloud is resolved into 225 × 225 × 113.
For simplicity, we adopted further assumptions.
(3) The reflection and fluorescence are isotropic, and the absorption is given by the analytical
function of the photoelectric cross section for the standard interstellar matter (Morrison &McCammon, 1983). The cross section of an iron atom is taken from Henke et al. (1982).
(4) Since the Thomson scattering optical depth is much smaller than that of the photoelectric absorption in the relevant energy band below 10 keV, we neglect the multiple Thomson
scattering.
Then the probability Pscat (E) that a primary X-ray photon arrives at a cell, is scattered,and comes into our line of sight, is given as a function of photoelectric absorption cross section(σ(E)), the column density from the primary source to the cell (N Ha ), the sell size (l), Thomson
scattering cross section (σT ), the Hydrogen number density in the cell (n cell ), and the column
density from the cell to the surface of the cloud (N Hb ) (see Figure 5.18).
Pscat = lncell σT exp (−NHa σ(E ) − NHb σ(E )).

(5.2)

Some fractions of the primary X-rays with energies above 7.1 keV, the K-edge of neutral
iron, are absorbed with the cross section of an iron atom (σFe (E)). The absorbed X-rays
are re-emitted as the 6.4-keV photons, with the probability of the fluorescence yield of 0.34
(Bambynek et al. 1972).
Thus the probability P6.4 , that a primary X-ray (E ≥ 7.1 keV) is converted to the 6.4 keVline and comes into our line of sight, is similarly given;

92

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

P6.4 = 0.34lncell ZFe σFe(E) exp (−NHa σ(E ) − NHb σ(6.4)),(5.3)

where ZFe is abundance of iron.
In both the cases, the interstellar absorption from the primary source to the cloud is neglected. Adding the Galactic absorption from each cloud to the observer, we finally obtained
simulated spectra.

113 cellsbNH

φ = 225 cells
Line of SightaNH

SgrB2 Cloud

The primary X-ray

Figure 5.18: The schematic view of the XRN simulation . The cloud (45 pc diameter, 22.5 pc
thickness) is divided into 225 × 225 × 113 cells, and the penetrating and reflecting (scattering)rate at each cell is calculated.

5.5. NUMERICAL SIMULATIONS OF XRN MODEL

93

Mass and Size of each Cloud
We refer to radio observations of molecular lines for the size and the mass of each cloud. Since
values vary with the definition of the parameters (particularly in size), we use the ratio to
Sgr B2 cloud in each parameter to determine model cloud as possible.
In Table 5.9, we show the observed parameters of the clouds obtained with GC survey by
Oka et al. (1998) and Stark et al. (1991). In these observations, the clouds are analyzed inunified manner, so we can obtain the accumulate ratio of parameters. For the Sgr C cloud,
both survey data gives almost the same ratio to Sgr B2. We hence adopt the averaged ratioof 0.13 for the mass and 0.76 for the size, and estimate those parameters of model cloud to be
8.5 ×105 M⊙ and 17.2 pc. The mass and size of Clump 2 is not mentioned in Oka et al. (1998).

We used the data in Stark et al (1991), only.

M0.11−0.08 is not mentioned in Oka et al. (1998) nor Stark et al. (1991), and we cannot
derive the ratio to Sgr B2. This cloud is investigated in detail by Oka et al. (2001) and Tsuboi
et al. (1997) with CO and CS emission lines, respectively. The size of the cloud is almost thesame in both observations, but the mass in Oka et al. (2001) is about an order of magnitude
higher than the Tsuboi et al. (1997). It is because CS line is radiated from central dense core,
while CO intensity may be due to the contribution of low-density gas. Since the model cloudof Lis & Goldsmith (1989) is based on CO data and the survey observations are made with COlines, we refer to Oka et al. (2001). However, the uncertainty of mass and size may be large
for M0.11−0.08. The estimated parameters of model clouds of four clouds are summarized inTable 5.9.
We must consider another factor in determining the cloud mass. Molecular clouds have two
possible solutions in mass estimation: gravity equilibrium mass and pressure equilibrium mass.
Oka et al. (1998) observed 12 CO and 13 CO line emissions, and revealed that if the clouds are
assumed to be in gravity equilibrium, virial masses are an order of magnitude larger than masses
estimated from the CO luminosity. To remove the discrepancy, they suggest that molecular
clouds in the GC region should be in pressure equilibrium with the hot gas and/or the magnetic
field near the cloud region. In this case, the masses become smaller than estimated values inTable 5.9. We then multiply a parameter η(< 1) by each mass of the cloud.
Interstellar Absorption to each Cloud
The interstellar absorption to the GC region is systematically investigated by Sakano (2000).
◦
The column density is almost constant in the Galactic longitude direction (|l II | <∼ 1 ), but largely

94

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

Table 5.9: The Parameters of the Molecular Clouds Used in XRN Simulation

Sgr B2
Sgr C
Clump 2
M0.11−0.08

Mass (×106 M⊙ )Oka98∗
Stark‡
Model
7.1
6
6.3
0.96 (0.14)§ 0.8 (0.13)§ 0.85 (0.13)§—
3 (0.5)§3.2 (0.5)§
—
—
1.2¶

Oka98∗
36
28 (0.78) §——Size† (pc)
Stark‡
Model
100
22.5
75 (0.75)§ 17.2 (0.76)§
480 (4.8)§ 108 (4.8)§—
4.3¶† : The radius of the cloud∗: The values are from Oka et al. (1998)‡: The values are from Stark et al. (1991)§ : Values in parentheses show the ratio to the Sgr B2 cloud in each parameter.
¶ : The values are from Oka et al. (2001)

decreases with the larger Galactic latitude. The 2-dimensional distribution can be expressedas:
NH = 4.1 × e− 12& ′ '2|l |2.1

× |b′ |−0.43(1022 H cm−2 ),(5.4)

where |l ′| ≡ |lII + 0.◦ 056| and |b′| ≡ |bII + 0.◦ 046| in the units of degree1.

The center positions of the X-ray emissions are (l ′ , b′) = (0.724, 0.012), (−0.508, −0.001),(0.173, −0.062), (2.728, 0.049) for Sgr B2, Sgr C, M0.11−0.08, and Clump 2, respectively. The
absorption column of each cloud is calculated to be 6–77 ×1022 H cm−2 . However, considering

the extent of the clouds (∼ 0.◦ 1) and the rapid increase in absorption near b′ = 0, these values
are not significant. We then substitute the conservative limit of b ′ = 0.◦ 1, and obtain the columnof ∼ 1023 H cm−2 for Sgr B2, Sgr C, M0.11−0.08, and ∼ 5 × 1022 H cm−2 for Clump 2. We
then adopted fixed values of 1023 H cm−2 (Sgr B2, Sgr C, M0.11−0.08) and 5 × 1022 H cm−2
(Clump 2) for the fittings of simulated spectra.

5.5.2

Simulated Images

For image analysis, we use only Chandra image of Sgr B2. Other clouds have poor photon
statistics, and rather complicated density distributions, which are not valid for image analysis.
While Sgr B2 is the most luminous in 6.4-keV line band and distributed in simple oval shape(Figure 2.3b, c).
1

Our Galactic nucleus Sgr A∗ is located at (lII , bII) = (−0.056, −0.046).

5.5. NUMERICAL SIMULATIONS OF XRN MODEL

95

Figure 5.19 is a simulated fluorescent line image of Sgr B2.The whole distribution of thefluorescent X-rays are concave shape pointing toward the irradiating source. These features
resemble the Chandra image (Figure 5.3).
With the XRN model, the molecular cloud is irradiated by an external X-ray source. The
X-rays penetrate into the molecular cloud, suﬀering with absorption, and excite iron atoms.
Then fluorescent X-rays are emitted. If the cloud is optically thick, the X-rays cannot penetrate
to the center of the cloud. Thus fluorescent X-ray distribution is shifted toward the irradiating
source.
Hence the separation between the X-ray peak and the center of the cloud depends on theabsorption column. We then compare the separation between the simulated and the observed
fluorescent line image. Though the assumed shape of cylindrical configuration may not bevalid for image simulation, the separation is a good indicator to estimate the cloud mass andabundances.
To compare the separation, we concentrate on the illuminated side of the molecular cloud.
We make radial profiles of the observed and simulated flux in the opening angle of ±60 degrees

above and below the horizontal line pointing to the irradiating source located in the Galactic
center direction.
The radial profile of the observed data in the 6.15–6.55 keV band is shown by the crosses
in Figure 5.20, where the background level estimated from background region in Figure 5.3 isindicated by the dashed line. The central region within 10 ′′ radius are excluded because they

are contaminated by the local excess of Sgr B2 Main. We also note that the other local excess,Sgr B2 North, is outside the opening angle of ±60 degrees. We can see that the observed
separation is about ∼ 40′′ . While, simulated image with the solar abundance and η = 1 shows
the separation of ∼ 90′′ , which is significantly larger than observed value.

We then simulated the images with various abundances and masses. In the XRN model,
X-rays above the 7.11 keV energy are absorbed mainly by iron atoms, while 6.4-keV X-rays
are absorbed mainly by lighter elements such as carbon, nitrogen and oxygen2. Thus simulated
separations depend on the number of iron atoms and other lighter elements, which are represented by two independent parameters3 of ηZFe and ηZothers , where ZFe shows the abundanceof iron, and Zothers shows the abundance of lighter elements. η is a coeﬃcient of the mass
2

X-rays above 7.11 keV also absorbed by lighter elements, but the eﬃciency is lower than iron atoms. The

simulation includes cross sections of lighter elements even in higher energy
3
Since the absorption cross sections of hydrogen and helium atoms are much smaller, we ignore the dependence on these atoms (total number is nearly proportional to η).

96

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

Figure 5.19: Simulated image of the iron fluorescent line with the XRN model. Contours show
the density distribution of the model cloud. Primary X-ray source is to the right.

5.5. NUMERICAL SIMULATIONS OF XRN MODEL

97(section 5.5.1). We then simulate images with the various sets of (ηZFe , ηZothers ).
The simulated profile fits the observed data well at the allowable region shown in Figure 5.21;
the separation between X-ray peak and the cloud center is 1.5±0.4 pc (or 36±10′′ ). We show
the best-fit profiles of (ηZothers , ηZFe ) = (0.3, 0.3) in Figure 5.20 in solid line.
Thus we revealed that the Chandra image can be reproduced by the simulated image. However, we need more complicated assumption of gas distribution for more accurate quantitative
comparison.

5.5.3

Simulated Spectra

Figure 5.22 shows an simulated spectrum with the solar abundance and η = 1, in case of Sgr B2.
There are strong Kα (6.40 keV) and Kβ (7.06 keV) lines of neutral iron and a deep iron K-edge(7.11 keV). In this subsection, we fit the simulated spectrum to the observed spectrum of thecloud. Spectral simulation is made for all the clouds.
Sgr B2The simulated spectra are convolved with the response function of ACIS and are comparedwith the observed spectrum. We first fit the observed spectrum of Sgr B2 to the simulated one
with the XRN model of solar abundances, where a free parameter is only the normalization of
the flux.
This simulated spectrum, however, is rejected with a reduced χ 2 of 1.48 (23 d.o.f.). Residuals
are found in the flux of the 6.4 keV line, depth of the K-edge and low-energy absorption.
As already mentioned in image analysis, X-rays above the 7.11 keV energy are absorbed
mainly by iron atoms, while the lower energy X-rays are absorbed mainly by lighter elements
such as carbon, nitrogen and oxygen. While continuum X-ray flux by the Thomson scattering is proportional to the number of electrons (∼hydrogen atoms). Accordingly, the ratio ofabsorption depth at 7.11 keV and low energy band is proportional to the abundance ratio of
iron and the other elements. The equivalent width of the 6.40 keV line constrains the ironabundance. We thus simulate the spectrum with various abundance sets of iron and the lighter
elements and compare with the observed spectrum.
Since the number of hydrogen atoms independently depend on η, there are three independent
variables of η, ZFe , and Zothers . We then simulate the spectra with various sets of (ZFe , Zothers ),and η.
We obtained acceptable fittings. The 90% confidence region is laid on the Figure 5.21 in

98

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

X10 -2

counts/pixel

3

2

1

0

0

20
40
60
80
Separation from cloud center (arcsec)

Figure 5.20: The radial profile of the surface brightness of the 6.4 keV line (crosses). The solid line
shows the simulation of the XRN model, and the dashed straight line indicates the background level
estimated from the background region shown in Figure 5.3.

5.5. NUMERICAL SIMULATIONS OF XRN MODEL

99

Figure 5.21: The allowable region of ηZ others and ηZFe constrained by image analysis (between two
solid lines). Spectral constraints are also indicated by dashed (η = 0.4), dash-dotted (η = 0.3), dotted(η = 0.2), and dash-dot-dot-doted (η = 0.1) lines.

100

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS

Figure 5.22: A simulated spectrum with XRN model.
(6.40 keV), Kβ (7.06 keV) lines, and K-edge (7.11 keV).

The spectrum exhibits iron Kα

5.5. NUMERICAL SIMULATIONS OF XRN MODEL

101

case of η = 0.1, 0.2, 0.3, and 0.4. We can find that the allowable regions of the image and thespectral analysis overlap each other at η < 0.3. No significant lower limit of η is determined. Incase of η = 0.2, the best-fit parameters within the overlapping region are (ZFe , Zothers ) = (2.5,2.0). The best-fit model spectrum is shown in Figure 5.23.
Hence we revealed that observed image and spectrum of Sgr B2 can be explained by the
XRN model.

Figure 5.23: X-ray spectrum of Sgr B2 obtained with Chandra ACIS and best-fit spectrum of
XRN model in case η = 0.2 (dashed line). The dotted line indicates the contribution of all the
resolved point sources in Sgr B2.

Sgr C, M0.11−0.08, Clump 2
With the same procedure as Sgr B2, we fit the spectra of the rest three clouds. We have alreadyknown that η < 0.3 in case of Sgr B2. Assuming the coeﬃcient η to be same in all clouds, we
simulate the spectra in two cases: η = 0.1, 0.2. The acceptable regions of each case is shownin Figure 5.24, 5.25.
All of the spectra can be reproduced by simulated spectra. Hence XRN hypothesis is verified
with numerical simulations.

102

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONS(a) Sgr B2(b) Sgr C(c) M0.11−0.08(d) Clump 2

Figure 5.24: Acceptable regions of spectral fittings with XRN model (η = 0.2). 80%, 90%, and
95% acceptance level is shown by dotted, solid, and dashed lines, respectively. The acceptableregions of Sgr C, M0.11−0.08, and Clump 2 are left side of the lines.

5.5. NUMERICAL SIMULATIONS OF XRN MODEL

103(a) Sgr B2(b) Sgr C(c) M0.11−0.08(d) Clump 2

Figure 5.25: Confidence contours of spectral fittings with XRN model (η = 0.1). 80%, 90%, and
95% acceptance level is shown by dotted, solid, and dashed lines, respectively. The acceptableregions of Sgr C, M0.11−0.08, and Clump 2 are left side of the lines.

104

CHAPTER 5. RESULTS I. — 6.4-KEV LINE EMISSIONSChapter 6
Results II. — Supernova Remnants
Supernova remnants (SNRs) are the plausible candidate to produce the high-temperatureplasma extended in the GC region. However, the number of discovered SNRs is insuﬃcient
to explain the large total energy of the plasma. Recently, high angular resolution of Chandra
enables us to search very small SNRs because of the short evolution time after the explosion, orvery faint SNRs which are buried in a surrounding X-ray emissions. In this chapter, I introduce
the newly discovered SNR candidates: G 359.92−0.09 and G0.570−0.018.

6.1

G 359.92−0.09

G 359.92−0.09 is newly discovered in the X-ray band with Chandra. This SNR is at the southof Sgr A East.

6.1.1

Radio Observations

Ho et al. (1985) detected an elongated future, called “wisp”, at ∼ 4′ south of Sgr A∗ with theradio continuum observations (Figure 6.1). The radio emission from “wisp” is non-thermal andpolarized 10–20 %. Combined with the distribution of NH3 molecules, the authors indicatedthat the “wisp” is a part of the SNR, which is interacting with a molecular cloud. The inward
curve of the Sgr A East and the other condensations aligned in a circular shape (shown by
dotted line in Figure 6.1a) imply a shell-like structure, which is attributable to SNR. The later
observation of the dynamics of molecular material supported the interaction with the molecular
cloud and Sgr A East, then this SNR candidate is named G 359.92−0.09 (Coil & Ho 2000).
However, this object is detected only by radio band. In addition, only a few clumps areobserved. The radio observations are insuﬃcient to conclude that G 359.92−0.09 is SNR.
105

106

CHAPTER 6. RESULTS II. — SUPERNOVA REMNANTS(a) Radio 6 cm image(b) NH3 image

Figure 6.1: The radio image of the Sgr A region (Ho et al. 1985). (a) 6 cm continuum map. A
elongated future “wisp” (E) continues northwest into another condensation (F). Both condensations
and other three features at R.A. = 17h 42m 40.s0, Dec = -29◦ 01′ 00′′ (epoch 1950) may be part of a
SNR with shell-type morphology (dotted circle). (b) The 6 cm continuum map superposed on an NH 3
emission map. Two molecular condensations of M−0.02 − 0.07 and M−0.13 − 0.08 are shown. The

shell-like structure seems to interact with M0.13 − 0.08 at the position of “wisp”.

6.1. G 359.92−0.09

6.1.2

107

Chandra Observation

Image
Figure 6.2 shows the Chandra 3–8 keV image around the Sgr A region. A sharp structure is
detected at the same position as “wisp”. In addition, diﬀuse X-ray emission is detected fromthe eastern half (left in the figure) in the circular structure indicated in radio band.
Figure 6.3 shows the radial profiles of the diﬀuse X-ray emission in the each quadrant of
the circular structure (divided by the dashed lines in Figure 6.2). There are excesses in theshell at the northeast and the southeast part compared with the outside region. The boundary
shape is similar to a step-function. The northwest part shows no clear boundary because of
the contamination of the intense X-ray emission from Sgr A East. The southwest part shows
no excess in the shell, but has very sharp boundary, which is attributable to the “wisp”.

-0.1

0.0

0.0 0.1

359.9

359.8

Figure 6.2: The X-ray image of G 359.92−0.09 in 3.0–8.0 keV band. Up is north, and left is east.
There is diﬀuse emission in the circular region suggested by radio observations (solid line). Radial
profiles are taken from the each quadrant of the circular region divided by dashed lines.

108

CHAPTER 6. RESULTS II. — SUPERNOVA REMNANTS

0.5

0.5

0.45

0.45

0.4

0.4

0.35

0.35

0.3

0.3

0.25

0.25

0.2

0.2

0.15

0.15

0.1

0.1

0.05

0.05

0

0

50

100 150 200 250 300 350 400
radius (pixel)

0

0

50

(a) Northeast(b) Northwest

0.5

0.5

0.45

0.45

0.4

0.4

0.35

0.35

0.3

0.3

0.25

0.25

0.2

0.2

0.15

0.15

0.1

0.1

0.05

0.05

0

0

50

100 150 200 250 300 350 400
radius (pixel)(c) Southeast

100 150 200 250 300 350 400
radius (pixel)

0

0

50

100 150 200 250 300 350 400
radius (pixel)(d) Southwest

Figure 6.3: Radial profiles of the X-ray emission from (a) northeast, (b) northwest, (c) southeast,and (d) southwest part of the circular structure. In each figure, the horizontal line shows the radiusfrom the shell center, and the vertical line shows the X-ray counts per pixel. Dashed lines show theboundary of circular region shown in Figure 6.2.

6.1. G 359.92−0.09

109SpectrumAt first, we extract the X-rays from the eastern half of the shell structure, and make thespectrum (Figure 6.4). The source region is inside of the circle of r < 2 ′ , and the backgroundregion is the annulus of 2′ < r < 2.′ 8, where r is the distance from the center of the shell. The
spectrum exhibits the strong iron K-line. We fit the spectrum with a phenomenological modelof the thermal bremsstrahlung and an Gaussian line. The center energy of the iron is 6.62 keV,which is slightly lower than that of the He-like iron. We then fit the spectrum with the nonequilibrium ionization (NEI) model, and obtain the best-fit values shown in Table 6.1. The
Hydrogen column density of about 6 × 1022 H cm−2 is consistent with the Galactic absorption

to the GC (Sakano 2000). The ionization parameter nt, where n and t are plasma density and11s cm−3. In this fitting, we fixed theelapsed time after the plasma is heated-up, is 1.4 +5.4
−0.6 × 10

metal abundances to be the solar values.

Figure 6.5 and Figure 6.6 show the X-ray spectrum of the “wisp” and the southwestern
quadrant, respectively. Both spectra show the iron K-line. We fit the spectra with the samemodel as eastern half. The best-fit values are also shown in Table 6.1. Both region exhibit
larger absorption than eastern half region.

Figure 6.4: The X-ray spectrum of the eastern half region of the G359.92−0.09. The best-fit spectrumwith NEI model is shown by solid line.

110

CHAPTER 6. RESULTS II. — SUPERNOVA REMNANTS

Table 6.1: Table
NH aeastern half
6.0+2.3
−2.3
wisp
37+7
−5
southwestern quadrant 18+15
−11

kT b
11.4+68.5−8.0
12.7 (> 7.9)
2.6+74.7−0.7

ntc
13.5+54.4−5.4
884 (> 0.01)
20.6 (> 0.1)

FX d
15+1
−2
4.4 ± 0.5
4.3+0.9−0.8

LX e
χ2/d.o.f.
+1.7
1.8−0.4
44.37/38
1.2 ± 0.2
6.0/5
+220
1.4−1.1
9.8/9

a : Hydrogen column density [10 22 H cm−2 ].
b: The temperature of the plasma [keV].

c: The ionization equilibrium parameter [10 10 s cm−3 ]d : Flux (no correction of absorption) in the 2–10 keV band [10 −13 erg cm −2 s−1 ].
e : Absorption corrected luminosity in the 2–10 keV band [10 34erg s−1 ].

Figure 6.5: The X-ray spectrum of the “wisp”.

6.1.3

Discussion

X-ray emission from the western partIn section 6.1.2, we confirmed the presence of the diﬀuse X-ray emission at the eastern partof G359.92−0.09, but there is no clear excess in the western part. In addition, the spectrum

6.1. G 359.92−0.09

111

Figure 6.6: The X-ray spectrum of the southwestern quadrant of the G359.92−0.09.
analysis indicates that there is a larger absorption column in the west. These results imply
that the western part of the SNR is behind the molecular cloud: M−0.13−0.08 in Figure 6.1b.
The column density of the cloud is estimated to be (3–8) ×1023 H cm−2 by Zylka et al. (1990)

and (8–14) ×1023 H cm−2 by Coil & Ho (1999). These values are larger than, or suﬃcient to
explain the X-ray absorption.

Coil & Ho (2000) investigated the dynamics of the molecular clouds within 15 pc of the
Galactic center, and concluded that the cloud and molecular flow, called southern stream, areinteracting with G359.92−0.09. Though they suggested that the molecular cloud and southern
stream are more behind the SNR, it is marginal and diﬃcult to decide the position in line ofsight only with the radio data (e.g. Coil & Ho 1999). Thus we suspect that at least a part ofthe molecular cloud is in front of the SNR, and the lack of X-ray intensity in the west is only
due to the absorption.

Physical parameters
In this section, we use only eastern half data to estimate the physical parameters of the SNR.
We assumed the distance to the SNR to be 8.5 kpc. With the spectral fitting, we obtained thebest-fit values of emission measure (E.M. = n2 V ) of 1.1×1057 cm−3 , and ionization parameter
nt of 1.4×1011 cm−3 s, where V is the plasma volume. Assuming the X-rays are emitted from112

CHAPTER 6. RESULTS II. — SUPERNOVA REMNANTS

the uniform hemisphere with the size of the spectrum region of 2 ′ (∼ 5 pc) radius, the volume
is calculated to be ∼ 7 × 1057 cm−3 . Hence the plasma density (n) and ionization age (t) are

estimated to be about 0.4 cm−3 and 104 year, respectively.

6.2

G0.570−0.018

This X-ray source is located at the 7′ southwest from Sgr B2. ASCA detected a faint X-ray
peak from the same observed region of Sgr B2. Then Chandra revealed a shell-like structurewith the radius of about 10′′ . The detailed description about G0.570−0.018 will be publishedsoon (Senda et al. 2001).

6.2.1

ASCA Observation

Because G0.570−0.018 is near the Sgr B2 region, it is observed simultaneously with Sgr B2(Section 5.1.1). The net observed time is about 95 ksec for GIS.

Image
Figure 6.7 shows the 6.0–7.0 keV image near the Sgr B2 region. The brightest X-ray peak is
Sgr B2, but there is another peak near the center of the image (solid circle). This source isalready mentioned in Sakano, Yokogawa, & Murakami (1999) and Sakano (2000), and named
AX J1747.0−2828. We can not determine that X-rays are from a point source or an extended
structure with the angular resolution of ASCA. The position of the peak is 0.◦55 ≤ l ≤ 0.◦ 57,−0.◦ 02 ≤ b ≤ 0.◦ 00.
SpectrumThe spectrum is obtained from the 3′ radius region around the X-ray peak, and the backgroundregion is annulus with the radius of 3′ –4′ , which are indicated by solid and dashed lines in
Figure 6.7, respectively. Figure 6.8 shows the spectrum.
We fit the spectrum with the phenomenological model of a thermal bremsstrahlung and a
Gaussian line, and obtain acceptable fittings. The best-fit parameters are shown in Table 6.2.
Strong Kα line of iron is detected with the center energy of 6.5–6.7 keV and the equivalent
with of about 3.7 keV.

6.2. G0.570−0.018

113

60.5

+0°40′+0°05′+0°35′-0°10′+0°30′ -0°05′+0°00′

04.5
94.5
150.5
Figure 6.7: The 6.0–7.0 keV image around G0.570−0.018 observed with ASCA. The X-ray image is
BLOW BLOW IMSMO(Gauss) g3_6-7.img0

same as figure 5.1, but zoomed in G0.570−0.018

114

CHAPTER 6. RESULTS II. — SUPERNOVA REMNANTS

Figure 6.8: The X-ray spectrum of G0.570−0.018 obtained with ASCA. The solid line shows thebest-fit spectrum with a thermal bremsstrahlung and a Gaussian line model.

Table 6.2: Best-Fit Results of the X-ray spectrum of G0.570−0.018 obtained with ASCA.
a

GIS

NH
8.1+18.8−5.3

ContinuumkT b
FX c
13.6(>1.3) 4.7 ± 1.4

ECd6.60 +0.12−0.09

Iron line
FX c
2.6 ± 0.8

a : Hydrogen column density [10 22 H cm−2 ].
b : The temperature of the plasma [keV].
c : Flux (no correction of absorption) in the 2–10 keV band [10 −13 erg cm −2 s−1 ].
d : Center energy of iron line [keV].
e : Equivalent width of iron line [keV].
f : Absorption corrected luminosity in the 2–10 keV band [10 33 erg s −1 ].

EW e
3.7

LX f
9 +3
−2

χ2 /d.o.f.
6.4/12

6.2. G0.570−0.018

6.2.2

115

Chandra Observation

G0.570−0.018 is included in the 100 ksec observation field of Sgr B2 with Chandra.
Image
Figure 6.9 shows the Chandra image of G0.570−0.018 in 4.0–7.0 keV band. X-ray emission
shows clear shell-like structure of about 10′′ radius. The center of the shell is R.A.= 17 h 47m 02.s 6,
Dec= −28◦ 27′ 33′′ (epoch 2000), which corresponds to (l, b) = (0.◦ 570, −0.◦ 018), the same positionof the ASCA diﬀuse source within the error.

To investigate the shape of X-ray emission, we make a radial profile of X-ray emission fromthe center of the shell. Figure 6.10 is the radial profile in the 4.0–7.0 keV band. There is theclear excess around 10′′ with the width of about 6′′ (FWHM). Hence it is revealed that thediﬀuse emission is distributed as circular shape.

47.5-0°00′20″+0°34′20″-0°01′20″+0°34′-0°01′+0°33′40″-0°00′40″

10.5
339.5
412.5
Figure 6.9: Chandra image of G0.570−0.018 in 4–7 keV band. A faint shell-like structure is shown.
MO(Gauss) IMSMO(Gauss) IMSMO(Gauss) BLOW SN
Spectrum
The Chandra X-ray spectrum is made from the same area as with ASCA, while the background
spectrum is taken from a source-free region with the same CTI eﬀect as the source region. We
fit the spectrum with a thermal bremsstrahlung and a Gaussian line, and obtained the best-fit
parameters shown in Table 6.3. The parameters are consistent with the ASCA results withinnormalized photon number (counts/bin)

116

CHAPTER 6. RESULTS II. — SUPERNOVA REMNANTS

0.7
0.6
0.5
0.4
0.3
0.2
0.1
0
0

5

10

15

20

25

30

/ REBIN PROR:CTS:SNR_acisf944_filter2_4-7keV_2bin.img0

distance (arcsec)
Figure 6.10: The radial profile of X-ray emissions from G0.570−0.018. The horizontal line shows the
distance from the center of the shell-like structure.

6.2. G0.570−0.018

117

90% errors. The center energy and the equivalent width of the iron line is 6.47–6.53 keV and4.1 keV, respectively.

Figure 6.11: The X-ray spectrum of G0.570−0.018 obtained with Chandra.

6.2.3

Discussion

The faint X-ray emission discovered with ASCA, lately confirmed to have a shell-like structure.
The spectrum can be reproduced with the thin thermal plasma model. These results imply
that the X-rays are radiated from a SNR. In this section, we further investigate the propertiesof the G0.570−0.018.

Spectrum fitting with NEI model
Because the center energy of the line is smaller than that of the He-like iron, we fit the ASCA
and the Chandra spectra simultaneously with NEI model. The best-fit parameters and spectra
are shown in table 6.4 and figure 6.12, respectively. The ionization parameter nt is very small as
1.7×1010 cm−3 s−1 . The X-ray emitting plasma is far from a collisional ionization equilibrium
state. The hydrogen column density of NH ∼ 1.4 × 1023 cm−2 is a consistent value with theinterstellar absorption to the GC (Sakano 2000).

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CHAPTER 6. RESULTS II. — SUPERNOVA REMNANTS

Table 6.3: Best-Fit Results of the X-ray spectrum of G0.570−0.018 obtained with Chandra.
GIS

NH a
11.9+3.4−4.2

ContinuumkT b
FX c
4.8(>3.4) 5.0 ± 0.4

Iron line
ECd
FX c
6.50±0.03 3.2 +0.4
−0.3

EW e
4.1

LX f
13±1

χ2 /d.o.f.
26.2/22a : Hydrogen column density [10 22 H cm−2 ].
b : The temperature of the plasma [keV].
c: Flux (no correction of absorption) in the 2–10 keV band [10 −13erg cm−2 s−1 ].

d: Center energy of iron line [keV].

e : Equivalent width of iron line [keV].
f : Absorption corrected luminosity in the 2–10 keV band [10 33 erg s−1 ].

Figure 6.12: The simultaneously fitting of the ASCA (gray) and the Chandra (black) spectra. The
solid lines show the best-fit spectra of NEI model.

6.2. G0.570−0.018

119

Table 6.4: Best-Fit Results of the X-ray Spectrum with the simultaneously fitting of ASCA andChandra with a NEI model
GIS&ACISaNH
13.9+3.3−3.2

kT b
6.1+19.7−3.0

Continuum
nt
Zc
+1.0
1.7−0.4 4.5+5.4
−2.9

FX d
8.2 ± 0.5

LX e
16 +2
−1

Reduced χ2 /d.o.f.
41.3/40a: Hydrogen column density [1023 H cm−2 ].

b: Plasma temperature of a thin thermal plasma model for the continuum spectrum.

c : Chemical abundance with ratio to solar abundance.
d : Flux (no correction of absorption) in the 2–10 keV band [10 −6photons cm−2 s−1 ].
e: Absorption corrected luminosity in the 2–10 keV band [10 33erg s−1 ].

Physical parameters
With the same procedure as G359.92−0.09, we estimate the density and the age of the plasma
separately. From the image and the radial profile (section 6.2.2), we assumed the X-ray emitting
region to be a spherical shell with the radius of (10 ± 3)′′ or 0.3–0.5 pc. Then the volume of theplasma is 1.2 ×1055 cm3. The E.M. is derived from a new spectrum region with the smaller

radius to exclude X-rays from the surrounding region. The E.M. is 5.5 ×1056 cm−3, which is

67% of that of the whole area of G359.92−0.09. Then using the best-fit ionization parameter(nt) of 1.7 ×1010 cm−3 s, the plasma density (n) and the ionization age (t) are estimated to be

6.7 cm−3 and 80 year, respectively. The expanding velocity is about 4900 km s −1 . The age ofthe plasma is very young, and the SNR should be in a free expansion phase. In fact, the Sedov
self-similar model leads the inconsistent results (Senda 2001).

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CHAPTER 6. RESULTS II. — SUPERNOVA REMNANTS

Chapter 7
Results III. — Other Structures
With the high angular resolution of Chandra, we discovered many futures in the Galactic
center region as well as SNRs or XRNs. In this section, we concentrate on new discoveries with
Chandra.
Figure 7.1 shows the combined image of the Chandra observations of the Sgr A and the RadioArc region. There are many clumps in the field (labeled 1–11). The filamentary structure isalso found in the Sgr C region (Figure 5.8). This clump is labeled No. 12. As mentioned later,
the physical properties are diﬀerent from clump to clump; some show the 6.7 keV line, others
are the 6.4 keV line emitters, and the others have non-thermal spectra. We classify the X-ray
clumps based on the presence (absence) of iron K-line and its center energy (E C ); 6.4-keVclumps, 6.7-keV clumps, and non-thermal clumps. The detailed study of these clumps is madeby Bamba et al. (2001).
In this chapter, we investigate the representative samples of these clumps. No. 2, 11 fornon-thermal, No. 10, 12 for 6.4 keV, and No. 1 for 6.7-keV clumps.

7.1

Non-Thermal Clumps

7.1.1

No. 2 — G359.96−0.05

Image
Figure 7.2 shows the X-ray image and the radio continuum image of Sgr A region. The faintfilamentary structure is found at the northeastern edge of Sgr A East. The X-ray structure
is coincident with a faint radio filament. We named this structure G359.96−0.05 from theposition in the Galactic coordinates.
121

122

CHAPTER 7. RESULTS III. — OTHER STRUCTURES

Figure 7.1: Merged image of the Chandra observations of the Sgr A and the Radio Arc region. Energy
band is 3.0–8.0 keV. Many diﬀuse structures are shown, and indicated by number 1–11.

7.1. NON-THERMAL CLUMPS

123

Spectrum
Figure 7.3 shows the X-ray spectrum of G359.96−0.05. The spectrum region is 3.′′ 5×7′′ rectangle
surrounding the filament. The spectrum exhibits no line structure. We fit the spectrum witha power-law model, and obtain the best-fit parameters shown in Table 7.1. Due to the limited
photon statistics, we cannot give significant constraint on the upper-limit of iron line intensity.

(b) Chandra 2-6keV
X-ray filament

30"

30"
(a) Radio image(b) Chandra image

Figure 7.2: The X-ray and radio images of an X-ray filament G359.96−0.05. (a) and (b) show
radio and Chandra images of the identical field around Sgr A ∗ (Koyama 2001). The X-ray and radio
filaments are located near at each other, implying physical correlation between them.

Table 7.1: Best-Fit Results of the X-ray spectrum of G359.96−0.05.
Power-law
a
NH
Γb
FX c
LX d
χ2/d.o.f.
+7.2
+1.4
+0.2
+35.1
ACIS 7.3−5.5 1.7−0.6 2.4−0.3 3.0−0.6
9.5/15a : Hydrogen column density [10 22 H cm−2 ].
b: Photon index of a power-law model for the continuum spectrum.

c : Flux (no correction of absorption) in the 2.0–10.0 keV band [10 −13 erg cm−2 s−1 ].
d : Absorption corrected luminosity in the 2.0–10.0 keV band [10 33erg s−1 ].

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CHAPTER 7. RESULTS III. — OTHER STRUCTURES

Figure 7.3: Chandra ACIS spectrum of G359.96−0.05. No line future is found.

7.1. NON-THERMAL CLUMPS

125

Discussion
We found a filamentary structure at the Sgr A region. The spectrum shows no line structure,and can be reproduced with power-law model. The shape and the spectrum indicates that the
X-rays from G359.96−0.05 is non-thermal origin, which may be attributable to a synchrotron
emission by high energy electrons. Further verifications of the hypothesis requires more accurate
observations in both radio and X-ray band.

7.1.2

No. 11 — G0.13−0.01

Image
Figure 7.4 shows the 3.0–8.0 keV band image of G0.13−0.01. It shows thin filamentary structurewith the size of about 60′′ × 10′′ . It is located at the edge of the molecular cloud M0.11−0.08(section 5.3). Tsuboi, Ukita, & Handa (1998) and Oka et al. (2001) found an expanding shellat this side of the cloud.
Spectrum
We make the X-ray spectrum of G0.13−0.01 from the oval region surrounding the filament.
Since G0.13−0.01 is buried in the diﬀuse X-ray emission from the cloud (section 5.3), we
subtract the background spectrum extracted from the neighbor oval region. The spectrum andthe background regions are shown in Figure 7.4a in solid and dashed lines, respectively.
The spectrum can be reproduced with a single power-law model. The best-fit spectrum andparameters are shown in Figure 7.4b and Table 7.2, respectively. To estimate the upper-limitof the iron line emission, we add a narrow Gaussian line with the fixed center energy of 6.4 or
6.7 keV. In both cases, the equivalent width of the line is <∼ 0.2 keV. Thus we conclude that
G0.13−0.01 is a non-thermal filament.
Discussion
We detected the non-thermal emission from G0.13−0.01. Tsuboi, Ukita, & Handa (1997) andOka et al. (2001) suggested that magnetic field is compressed by the interaction between
the expanding shell and the molecular cloud, and electrons may be accelerated to relativisticenergies through shock acceleration (section 2.3.3).
The X-rays from G0.13−0.01 should be synchrotron emission from high energy electrons.
Chandra observation strongly support the hypothesis that this region is the cite of the electron

126

CHAPTER 7. RESULTS III. — OTHER STRUCTURES

-0.10

0.15

0.10

(a) Chandra image(b) Chandra spectrum

Figure 7.4: (a) 3.0–8.0 keV band image of G0.13−0.01 obtained with Chandra. Filamentary structureis shown in solid ellipse region. The dashed ellipse shows the background region. (b) Chandra spectrumof G0.13−0.01 fitted with a power-law model.

Table 7.2: Best-Fit Results of the X-ray Spectrum of G0.13−0.01.
Power-lawaΓb
FX c
LX d χ2 /d.o.f.
NH
ACIS 7.8+6.8
2.1+1.2
3.0+0.4
4+39
23.0/25
−3.3
−0.8
−0.3
−3
a

: Hydrogen column density [1022 H cm−2 ].

b: Photon index of a power-law model for the continuum spectrum.

c : Flux (no correction of absorption) in the 2.0–10.0 keV band [10 −13 erg cm−2 s−1 ].
d : Absorption corrected luminosity in the 2.0–10.0 keV band [10 33erg s−1 ].

7.2. 6.4-KEV CLUMPS

127

acceleration. Note that the strong radio emission is outside of the cloud, and not coincident
with G0.13−0.01 (see Figrue 2.5).

7.2

6.4-keV Clumps

7.2.1

No. 10 — G0.09−0.09

Image
Figure 7.5a shows the 3.0–8.0 keV image around G0.09−0.09. There is a filamentary structurewith the size of ∼ 160′′ × 16′′ . This filament is located at the edge of the M0.11−0.08 cloud on

the GC side (see Figure 5.12).
Spectrum

Figure 7.5b shows the spectrum of G0.09−0.09. The spectrum and the background regions
are shown in Figure 7.5a. We first fit the spectrum with a power-law and a narrow Gaussianline model. Though this model allows acceptable fittings, the photon index of the power-law
model cannot be constrained in significant range. We then fix the photon index Γ to be 2.0,
and obtain the best-fit parameters as shown in Table 7.3. The center energy is ∼ 6.4 keV, andconsistent with that of neutral iron. The absorption column of ∼ 3 ×1023 H cm−2 is larger

than interstellar absorption to the GC region (Sakano 2000).
Discussion

As seen in Figure 5.12, this filament is located at the GC side of M0.11−0.08. The spectrum
exhibits strong iron line with the equivalent width of ∼ 1 keV, and large absorption. Thesefutures suggest that G0.09−0.09 is an XRN, which emits fluorescent and scattered X-rays. A
small condensation of molecular material may be located at this position, and emit fluorescent
X-rays. More detailed radio observations of molecular lines are required to reveal the origin of
the X-ray emission.

7.2.2

No. 12 — G359.42−0.13

Image
We found the faint elongated structure in iron line image at the Sgr C region (Figure 5.8a).
It is also found in 3.0–8.0 keV image (Figure 7.6a), which is perpendicular to the Galactic

128

CHAPTER 7. RESULTS III. — OTHER STRUCTURES

-0.10

-0.05
0.10

0.05

Figure 7.5: (a) 3.0–8.0 keV band image of G0.09−0.09 obtained with Chandra. Very faint filamentary
structure is detected (solid ellipse). The dashed line shows the background region. (b) Chandra
spectrum of G0.09−0.09. A strong iron line is detected.

Table 7.3: Best-Fit Results of the X-ray Spectrum of G0.09−0.09.

NH a
ACIS 27.7+9.3−8.4

Power-lawΓb
FX c
2.0 (fixed) 4.5±0.6Ed6.42+0.04−0.05

Iron line
FX c
EW e
0.9+0.3
1.1+0.2
−0.2
−0.3

LX f
13+5
−4

Reduced χ2 (d.o.f.)
21.2/20a : Hydrogen column density [10 22 H cm−2 ].
b: Photon index of a power-law model for the continuum spectrum.

c: Flux (no correction of absorption) in the 2.0–10.0 keV band [10 −13erg cm−2 s−1 ].

d : Absorption corrected luminosity in the 2.0–10.0 keV band [10 33erg s−1 ].

7.2. 6.4-KEV CLUMPS

129

plane. The size is about 5′ × 1.′5. It is not coincident with radio filaments or molecular clouds(Figure 2.4).
SpectrumThe spectrum of G359.42−0.13 is shown in Figure 7.6b. The spectrum and background regions
are shown in Figure 7.6a by solid and dashed lines, respectively. The spectrum shows a line
emission at 6–7 keV. We then fit the spectrum with a power-law and a narrow line model.
The best-fit parameters are shown in Table 7.4. The center energy of the line agrees with theneutral iron, and the column density of hydrogen is ∼ 5 × 1022 H cm−2 .
Discussion
We detected strong iron line from from G359.42−0.13, but the absorption is rather small. It
is consistent or smaller than interstellar absorption to the GC region. The small absorption iscontrary to the XRN mechanism. One possibility is that the emission line is from low-ionized
iron, or X-rays are from NEI plasma. In fact, NEI plasma model can reproduce the spectrumwith the best-fit parameters as follows: NH ∼ 9 × 1022 , kT ∼ 2.6 keV, and the the ionization

parameter nt ∼ 3 × 1010 (χ2 /d.o.f = 5.0/9). X-ray flux is 9 × 10−13 erg s−1 cm−2 in 2–10 keVband. Though the limited photon statistics prevents us from further study, G359.42−0.13 may
be a young SNR like G0.570−0.018.
-0.15

-0.10

359.45-0.05

359.40359.35(a) Chandra image of G359.42−0.13(b) Chandra spectrum of G359.42−0.13

Figure 7.6: (a) 3.0–8.0 keV band image of G359.42−0.13 obtained with Chandra. The concentrationof the X-ray emission is shown in solid ellipse region. Dashed line shows the background region. (b)Chandra spectrum of G359.42−0.13. Iron line is detected at ∼ 6.4 keV.

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CHAPTER 7. RESULTS III. — OTHER STRUCTURES

Table 7.4: Best-Fit Results of the X-ray Spectrum of G359.42−0.13.

NH a
ACIS 5.3+1.4−0.9

Power-lawΓb
FX c
2.0 (fixed) 9.2±1.0

Ed6.38+0.19−0.21

Iron line
FX c
EW e
1.3+0.5
1.0±0.5
−0.6

LX f
12+3
−2

Reduced χ2 (d.o.f.)
7.8/9

a : Hydrogen column density [10 22 H cm−2 ].
b: Photon index of a power-law model for the continuum spectrum.

c: Flux (no correction of absorption) in the 2.0–10.0 keV band [10 −13erg cm−2 s−1 ].

d : Absorption corrected luminosity in the 2.0–10.0 keV band [10 33erg s−1 ].

7.3

6.7-keV Clumps

7.3.1

No. 1 — G359.88−0.08, a part of G0.570−0.018

We detected filamentary structure at the south of the Sgr A region. The spectrum exhibits
6.7-keV line, and the distribution is coincident with a radio structure of “wisp”.
The nature of this clump is already studied in section 6.1, and G359.88−0.08 is revealed to
be a part of SNR. The spectrum and the image of this structure are shown in section 6.1.
Thus 6.7-keV clumps are possible to be SNRs. We can find new SNRs in the Galactic centerby the X-ray observations.

Chapter 8
Discussion
We revealed that;• giant molecular clouds emit 6.4-keV line
• there are new faint SNRs in the GC region• there are numerous new X-ray features in the GC region
in the previous chapters. We discuss the implication of these results in this chapter.

8.1

Verification and Consequence of the X-ray Reflection Nebula

In chapter 5, we investigate the giant molecular clouds in the GC region, and revealed that
all clouds emit strong 6.4-keV line. Since the 6.4-keV line is a characteristic line from neutral
iron atoms, we suspect that the cloud is irradiated by an external luminous X-ray source. Thishypothesis is confirmed by the simulated spectra and image (only Sgr B2) of the XRN model.
However, there remains an important problem; what is the primary X-ray source (or primary
X-ray sources) to radiate molecular clouds? In this section, we investigate what is the external
primary source.

8.1.1

Required Luminosity

From the equation (5.2), (5.3), we calculate the fractions of the reflected and fluorescent fluxes,
hence estimate the required luminosity of the primary source to give the observed fluxes ofcontinuum X-rays and the 6.4-keV photons.
131

132

CHAPTER 8. DISCUSSION

At first, we estimate the luminosity of primary source of Sgr B2. We adopt the best-fit
values of η = 0.2, and (ZFe , Zothers ) = (2.5, 2.0) for this analysis. The required luminosity is,39

L2−10keV ∼ 10

$

d
100pc

%2

erg s−1 ,(8.1)

where d is the distance from the primary X-ray source to Sgr B2. Since the reflecting region isextended largely over the Sgr B2 cloud, the required X-ray flux should be the value averaged
over a time of about 100 yr. This may exclude galactic binary transient sources as the primary
X-ray, whose duty cycle of flaring may be very short. We hence compare the required flux withthe observed luminosities of rather stable sources found in the catalog near the GC in Table 8.1(Sakano et al. 2001).
The observed luminosities of cataloged bright X-ray sources near the GC are 2–46 ×1036 erg s−1 ,which are much lower than that required to explain the reflected luminosity of Sgr B2. The

present X-ray sources cannot produce the strong fluorescent X-rays observed from the molecular clouds. As already mentioned, the luminosity of internal point sources are also insuﬃcient.
In fact, required luminosity exceeds the Eddington luminosity of neutron star.
This “no luminous primary source” problem leads to an exciting consequence; there was avery luminous X-ray source in the GC region in the past. The plausible candidate of the pastluminous X-ray source is the massive black hole Sgr A ∗, as already proposed in Koyama et al.
(1996). Sgr A∗ was bright some hundreds of years ago, the light travel time between Sgr A∗
and Sgr B2, and is dim at present. The shifted X-ray images of Sgr B2 and Sgr C support thatthe primary X-ray source is in the direction of GC.

Table 8.1: Comparison between Required Luminosities and the Observed Ones
X-ray sourcesda Lbobs Lcreq
1E 1740.7−2942
233
3 5400A 1742−294
175
3 30001E 1743.1−2843
622
380
SLX 1744−299/300 250
4 6300GX 3+1275
46 7600
a : The projected distance to Sgr B2 [pc].
b : The observed luminosity in the 2–10 keV band (Sakano et al. 2001) [10 36 erg s−1 ].
c: The required luminosity to account for the observed X-ray fluxes of Sgr B2 [10 36 erg s−1 ].

If the primary source of Sgr B2 is the Galactic nucleus (not a neighborhood of Sgr B2), theother clouds are also irradiated by the same primary source. We further verify the primary8.1. VERIFICATION AND CONSEQUENCE OF THE X-RAY REFLECTION NEBULA133
source by the required luminosities for other clouds.
Table 8.2 shows the required luminosities for other clouds in case the primary source is
Sgr A∗. The errors are derived only by the errors in observed luminosities. We use the same
parameters as Sgr B2: η = 0.2, and (ZFe , Zothers ) = (2.5, 2.0) for Sgr C and Clump 2. However,since these values are not in the 90% acceptance region of M0.11−0.08, we adopt the best-fit
values of (ZFe , Zothers ) = (1.5, 0.5) with η = 0.2 for M0.11−0.08.
SgrC, which is at almost the same distance as Sgr B2 from Sgr A∗ , requires smaller luminosity. However, considering the systematic uncertainties in spectral analysis 1 , these values are
consistent. In fact, ASCA results requires about 7 ± 2 × 1038 erg s−1 , which is consistent withthat of Sgr B2.

Required luminosity for Clump 2 is also consistent. While, M0.11−0.08 requires an orderof magnitude lower than other three clouds. It may be attributable to long time variability of
Sgr A∗, because M0.11−0.08 is closest to Sgr A∗. We conclude that fluorescent X-rays from allof the cloud can be explained by past activities of the Galactic nucleus Sgr A ∗ .

Table 8.2: Required Luminosities of Sgr A ∗ to explain the observed intensity of fluorescent andscattered X-rays.
blII
bII daI6.4[deg][deg] [pc] [×10 −5 ph cm−2]
SgrB2
0.67 −0.03 109
5.6
SgrC−0.56 −0.05 76
1.4
M0.11−0.08
0.12 −0.11 28
3.7
Clump22.67
0.00 409
0.4

LcSgrA[1038 erg s−1 ]
9+3
−2
3+3
−2
0.7±0.1
5+4−3

a : The distance from the Sgr A∗ to the cloud.
b: The observed intensity of 6.4-keV line.

c: The required luminosity of Sgr A ∗ to produce the observed 6.4-keV line emission. The errors arederived only by the errors in observed luminosities.

8.1.2

The Past Activities of the Galactic Nucleus

In the previous subsection, we estimated the required luminosities to explain the X-ray intensityof four molecular clouds. These luminosities should indicate the past luminosity of Sgr A ∗corresponding to the light travel time from Sgr A∗ to each cloud. We thus convert the distance1

There is another uncertainty in the radial distance to us.

134

CHAPTER 8. DISCUSSION

from Sgr A∗ to elapsed time, and make a light curve for the past 2000 years (Figure 8.1). The
luminosity is almost constant until a few hundreds years ago, and seems to decrease graduallyto the present value. However, there is only five data points. We cannot discuss the detailed
variability. The time span of each data point corresponds to the size of the cloud. Our method
is also insensitive to time variability in shorter time scale.
The past activity of Sgr A∗ would be generated by a surge of accretion on the MBH due,for example, to the passage of dense shell of a young supernova remnant Sgr A East as is
discussed by Maeda et al. (2001) based on the new Chandra results of the GC. They estimatedthe age of SNR to be ∼ 104 years, and the dense shell reached Sgr A ∗ ∼ 103 years ago. Thenthe shell swept the surrounding matter, and the luminosity became anomalously low. This
time scale agrees with the derived light curve by the XRN model. The luminosity could be

as bright as the Eddington luminosity of 3 ×1044 erg s−1 , which is also enough to explain theintensity of fluorescent X-rays. Coker (2001) further confirmed this results, and predicted that3
the luminosity of Sgr A∗ should increase over the next >∼ 10 years as the accreting matter isaccumulated again by cluster’s winds.
Hence our results of past activities of Sgr A∗ suggest that Sgr A∗ is accidentally in a low
luminosity phase at present.

8.1.3

Mass and Abundance of the CloudOka et al. (1998) proposed that the molecular clouds in the GC region is in pressure equilibrium.
In this case, the masses of the clouds become about an order of magnitude smaller, and thetotal mass within the central 400 pc of the Galaxy is (2–6) ×107 M⊙ .
The total mass of the Galactic center is estimated by several methods. Diﬀuse γ-ray in the
Galaxy is mainly produced with interaction of cosmic ray and matter. The total mass can be
estimated by diﬀuse γ-ray flux. Hence the upper limit of the total mass was given as 6 × 107

M⊙ within 500 pc (Blitz et al. 1985). Sakano (2000) also estimated the total mass within lessthan 380 pc to be 6.4 × 107 M⊙ from the absorption column to bright X-ray sources in the

GC region. This value is thought to be a upper limit since it is derived by assuming the solar
abundance. If the abundance is higher than solar value, it is reduced by a factor.

Thus all observations favor pressure equilibrium. We estimated the mass of the cloud to be< 0.3 times gravity equilibrium mass. Our result also supports the pressure equilibrium.
Chemical abundances of the GC region is also an important problem. It has been debatable
whether the GC region is overabundant or not (section 2.1.1). Our results of Sgr B2 shows8.1. VERIFICATION AND CONSEQUENCE OF THE X-RAY REFLECTION NEBULA135

Figure 8.1: The time variability of the Sgr A ∗ for the past 2000 years derived by XRN simulation.
The luminosities are estimated by fittings of observed spectra with simulations of XRN model for−100 – −2000 years ago. The time span corresponds to the size of each cloud. The luminosity ofrecent 10 years is from the direct observations with ASCA and Chandra (Maeda 1996; Baganoﬀ et al.
2001a, 2001b).

136

CHAPTER 8. DISCUSSION<<<
overabundance with 1.5 <∼ ZFe ∼ 3.0 and 2.0 ∼ Zothers ∼ 5.0 (η = 0.2, here and hereafter). While

we only obtained lower limit of ZFe , and upper limit of Zothers in case of Sgr C, M0.11−0.08,
and Clump 2. M0.11−0.08, in particular, shows the upper limit of Zothers <∼ 1.0, which isinconsistent with other clouds. If it is a real feature, the abundances are diﬀerent from cloud
to cloud. However, the spectrum of M0.11−0.08 includes thermal component. It may arise
systematic uncertainty to reduce the abundance. If we underestimate the X-ray flux of thermal
component, absorption in lower energy band becomes smaller, and lower abundance is obtained.
ZFe is constrained to be >∼ 1.0 for all the clouds. The iron abundance is mainly determinedby equivalent width of the fluorescent iron line. The line structures generally have better photon
statistics than continuum emission, and the ratio of fluorescent and scattered X-rays are notaﬀected by the cloud shape. Thus the abundance of iron has relatively good confidence thanthat of other lighter elements.
The XRN scenario has a potentiality to determine the total gas mass and chemical abundances. However, since the assumed density distribution may not be valid for image analysis,
further observations, especially in radio band, are essential to obtain accurate parameters.

8.1.4

Further Verifications

Our hypothesis of XRN is confirmed by numerical simulations of image and spectrum. Future
observations will further confirm the hypothesis of XRN.
The most convincing evidence is a detection of polarized X-rays. The scattered X-rays arepolarized perpendicular to the direction of primary X-rays. Churazov, Sunyaev, & Sazonov(2001) estimated the degree of polarization to be about 33%. Polarized X-rays also can be agood probe to determine the direction of primary X-ray source. This method is valid even forthe small cloud, which is optically thin to the center, and the X-ray image is not shifted fromthe center.
We can investigate the short time variability of the primary source by the evolution of Xray image. Sunyaev & Churazov (1998) simulated the morphology of the surface brightness
distribution, the equivalent width and the shape of the fluorescent line from the Sgr B2 cloud
as a function of time and relative position of the XRN and the primary (irradiating) source.
They examined both the cases of a flare-like and a step-function (X-ray flux is constant but
suddenly decreases to zero) irradiation by an X-ray source. As they demonstrated, a shorttime variability is the key to investigating the nature of the irradiating source, whether it is
the Galactic nucleus or a binary X-ray source. For comparison with data, some observations

8.2. ORIGIN OF THE HIGH TEMPERATURE PLASMA

137

with high angular resolution are necessary. The Chandra resolution of 0.5 arcsec, for example,
can attain the time resolution of about 1 month, which is typical flare duration of a transientbinary.

8.2

Origin of the High Temperature Plasma

With the high angular resolution of Chandra, we discovered new SNRs and SNR candidates
from the GC region (chapter 6, 7). SNRs are one of the most plausible candidates to produce
the high temperature plasma filling the central 1 ◦ × 2◦ region of our Galaxy. In this section,

we investigate the contribution of SNRs to high temperature plasma.

8.2.1

Energy Injection

The temperature of the GC plasma is ∼ 10 keV. Since the Galactic gravity cannot confine

such a high temperature plasma, the plasma is expanding. If we assume the expanding speed
to be the sound velocity (v ∼ 1600 km s−1 ), it takes about 50000 yr to reach the present size

(Koyama 1996). It means that the energy of 1054 erg is injected into the GC region within thepast 50000 yr.

One SNR releases about 1051 erg. Then a thousand of SNRs are required to produce thetotal energy. If supernova explosion rate is constant, it corresponds to one supernova per 50
yr. We discovered two SNRs, which is younger than 104 yr. In addition, Maeda et al. (2001)
revealed that the shell-like radio structure Sgr A East is also a SNR with the age of ∼ 104 yr.

We then assume that we can observe SNRs which are younger than 10000 yr by ∼ 50 ksec
observations.

With this assumption, we estimate the expected number of SNRs in the GC region. Long
exposure observations are made at three regions (Sgr A, Radio Arc, and Sgr B2). The net
observed area is about 800 arcmin2. While, the size of the GC plasma is 1.◦ 8 × 1.◦ 0 ≃ 5100
arcmin2, which is about 6.4 times of the long exposed area. Thus we simply estimate thatabout 20 SNRs younger than 10000 yr are hidden in the short exposed or not observed area in
the GC region. The number of SNRs which is younger than 50000 yr becomes 100. It is anorder of magnitude lower than required number. Thus SNRs discovered till now explain only
10% of the total energy of the GC plasma.
However, there may be additional SNRs in the GC region. Some fractions of newly discovered clumps may be SNRs. Radio observations discovered numerous shell-like structures from

138

CHAPTER 8. DISCUSSION

the GC region (e.g. Oka et al. 2001), which may also be SNRs. The discovery of G0.570−0.018
may be a piece of evidence that there are many supernovae in the GC region. The age of thisSNR is ∼ 80 yr, which is near the expected supernova rate. Further observations in X-ray andradio wavelength would reveal the accurate contribution of SNRs.

8.2.2

Temperature

The temperature is also a diﬃcult problem to explain the high temperature plasm by SNRs.
The temperature of GC plasma is ∼ 10 keV, which is much higher than that of the Galactic
SNR. This fact cause doubt on a thermal origin of the GC plasma (Tanaka et al. 2000)

Newly discovered SNRs in the GC region tend to show high temperature compared with
Galactic SNRs out of the GC region. For example, the temperature of G0.570−0.018 is ∼6 keV, and that of G359.92−0.09 is ∼ 11 keV. If it is common to GC SNRs, high temperatureof the GC plasma can be explained. However, observed temperatures have very large errors. It
is not certain whether these SNRs really exhibit high temperature comparable to the plasma
temperature. The lower limit of these temperatures are ∼ 3 keV. Further detail observations
and studies for the evolution of SNR in peculiar environment in the GC are required.

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Acknowledgments
I am deeply thankful to Prof. K. Koyama for his 5-year guidance of X-ray astronomy. He gave
me many precious advises and told me the attitude to research works. I would like to express
my gratitude to Dr. Y. Maeda and Dr. M. Sakano. They always give me valuable suggestions.
The large portion of this work is based on their pioneer works. I am indebted to Ms. A. Bambaand Mr. A. Senda for the analysis of a SNR and faint structures in the Galactic center.
I thank all the members of ASCA and Chandra team. These wonderful X-ray satellites are
essential for my study for the Galactic center. Some of the ASCA data is obtained by a part of
ASCA Galactic Plane/Center survey project. I am grateful to the principal investigator, Dr.
S. Yamauchi. He also gave me many useful suggestions. I also use the data of the Galactic
center survey with Chandra, which are contributed by Prof. Q. D. Wang.
I am deeply grateful to Prof. G. Garmire and Dr. Y. Tsuboi for their kind treatment in mystay at Pennsylvania State University to analyze Chandra data. I also had good discussions with
Prof. E. Feigelson. I am helped by Dr. K. Ebisawa, Y. Terashima in my visit to NASA/GSFC.
Many people give me useful suggestions. I thank to Dr. T. Oka, Prof. T. Hasegawa, Prof.
F. Sato, Mr. T. Sawada. I am grateful to Drs. H. Matsumoto, H. Tomida, K. Hamaguchi, andM. Nishiuchi for their kind instruction in X-ray astrophysics and analysis software.
I acknowledge all the people in Cosmic Ray laboratory in Kyoto University for their supportand encouragement. I spent pleasant days with them. This work is financially supported by
the Japan Society for the Promotion of Science.
I especially thanks to my family and my friends, finally.

145