X-ray Study on
the Activity History of Sagittarius A* byThree-dimensional View of the Galactic Center

Ryunosuke Gando (= S. G. Ryu)
Department of Physics, Graduate School of Science, Kyoto University
Kitashirakawa Oiwake-cho, Sakyo-ku, Kyoto, 606-8502, Japan
ryu@cr.scphys.kyoto-u.ac.jp

This thesis was submitted to the above department on January 7, 2013in partial fulﬁllment of the requirements
for the degree of Doctor of Philosophy in physics.

Abstract
Context. The most characteristic high-energy phenomena in the Galactic center (GC) arerepresented by strong K-shell emission lines from iron, particularly the 6.7 keV and 6.4 keVlines. The 6.7 keV line is from highly ionized iron in hot plasma (Galactic center plasma
emission: GCPE) of ∼ 107−8 K. The GCPE extends over the entire GC region for about2◦ with monotonic ﬂux decrease as a function of its distance from the central super massive
black hole (SMBH), Sagittarius (Sgr) A*. The 6.4 keV line is from neutral iron in cold giant
molecular clouds (MCs) in the central molecular zone (CMZ), such as Sgr B and Sgr A.
The X-ray spectra of these MCs have non-thermal continuum emission with large equivalent
width at 6.4 keV and exhibit strong absorption below 5 keV. These spectral features strongly
support the X-ray Reﬂection Nebulae (XRNe) scenario; the X-rays are due to ﬂuorescent
and Thomson scattering from MCs illuminated by external X-rays. Since the ﬂuxes of XRNe
require a high X-ray luminosity for irradiation larger than those of any brightest Galactic
X-ray sources, the potential candidate of the illuminating source is the SMBH Sgr A*. The
time variabilities of the 6.4 keV line and the continuum hard X-rays reinforced the idea thatthe emissions from XRNe are due to a big X-ray ﬂare (or ﬂares) from Sgr A* in the past.
According to rough estimations from the light-traveling time of projected distances between
XRNe and Sgr A*, the ﬂare should have occurred at about several hundred of years ago
and the past X-ray luminosity must be about 106 times higher than present. However, theline-of-sight positions of MCs/XRNe were not well known, and hence quantitative estimationof light-traveling time and X-ray luminosity for the past ﬂares was not possible.
Aims. In this thesis, we present an original self-consistent methodology to directly determine the three-dimensional (3-D) positions (projected 2-D plus line-of-sight positions) ofmultiple XRNe and quantitatively investigate the past activity history of Sgr A*.
Methods & Results. We carefully examined the X-ray spectra from the Suzaku observations on the Sgr B MC complex. The spectra consist of emissions from the XRNe and
the GCPE; the GCPE is decomposed into two components, one with large absorptions ofNH ≥ 10 × 1022 H cm−2 (Abs1) and another with absorptions of NH ∼ 6 × 1022 H cm−2
(Abs2). The large Abs1 is positively correlated with the 6.4 keV ﬂux, and hence is dueto the MCs/XRNe. The small Abs2 is nearly constant among all the XRNe, whose values
correspond to the Galactic inter-stellar absorption between the GC and the Sun. We further
testiﬁed and applied this spectral model to the other XRNe in the Sgr C, Sgr D, and Sgr E
regions. Combing the two results of the GCPE spatial distribution as an ellipsoid around
Sgr A* and the GCPE ﬂux ratio between its part with Abs1 (behind XRN) and the otherpart free from Abs1 (in front of XRN), we derived the line-of-sight position for each XRN

inside the GCPE.
We obtained an overall view for the 3-D distribution of XRNe around the GC. We foundthat these XRNe, coincident with MCs in diﬀerent radial velocities, are largely separated
in the line-of-sight position with respect to Sgr A*. XRNe in Sgr B are mostly located in
front; the most massive one in Sgr C is behind, while the other two are in front; those in
Sgr E and Sgr D are located in front and slightly behind, respectively. These 3-D positionsof XRNe are consistent with the MC structures in the bar potential suggested by the radio
observations assuming possible Galactic dynamics (e.g., circular rotations). This enhances
the reliability of our new method which directly determines the 3-D position. Using X-ray
parameters of the 3-D positions, the ﬂuxes, and the XRN column densities (Abs1), we ﬁnallyreconstructed a long-term light curve of Sgr A*. During the past 50 to 600 years, the SMBH
had been in a ﬂare state with luminosities between LX = 1–3 × 1039 erg s−1 , which were far
brighter than the present luminosity by a factor of ∼ 106 . The light curve was not ﬂat top in
the ﬂare epochs, but exhibited signiﬁcant variations within a factor of ∼ 3. The ﬂuctuations
were also conﬁrmed by observations of diﬀerent short-term (∼ 5 years) variabilities in the
6.4 keV ﬂux, of either increase or decrease, from diﬀerent XRNe in Sgr B and Sgr C. Thus,
we concluded that the X-ray activity of Sgr A* in the past 600 years was not a simple single
event, but consisted of multiple short ﬂares occurred continuously and sporadically.

Contents

1 Introduction & Thesis Composition

1

2 Review

3

2.1

The Milky Way Galaxy and the Center Region . . . . . . . . . . . . . . . . .

4

2.1.1

Basic Information of the Galaxy . . . . . . . . . . . . . . . . . . . . .

4

2.1.2

Central Molecular Zone . . . . . . . . . . . . . . . . . . . . . . . . . .

8Super Massive Black Hole, Sagittarius (Sgr) A* . . . . . . . . . . . . . . . .

8

2.2.1

Presence and Basic Properties . . . . . . . . . . . . . . . . . . . . . .

8

2.2.2

X-ray Observations of Sgr A* . . . . . . . . . . . . . . . . . . . . . .

10Thermal X-rays from the Galactic Center . . . . . . . . . . . . . . . . . . . .

14

2.3.1

Early X-ray Observations on the Galactic Plane . . . . . . . . . . . .

14

2.3.2

The Fe XXV Kα (6.7 keV) Diﬀuse Emission and the GC Hot Plasma .

16

2.3.3

The S XV Kα (2.45 keV) Diﬀuse Emission and Possible Presence of
Large-scale Low-temperature Plasma . . . . . . . . . . . . . . . . . .

17

Latest Results. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

18

Fe I Kα Diﬀuse Emissions in the Galactic Center . . . . . . . . . . . . . . . .

19

2.4.1

The 6.4 keV Line and its Origins . . . . . . . . . . . . . . . . . . . .

19

2.4.2

Observational Results of the 6.4 keV Emissions . . . . . . . . . . . .

23

2.4.3

X-ray Reﬂection Nebula and Past Activities of Sgr A* . . . . . . . .

24

Objectives and Approach of the Thesis . . . . . . . . . . . . . . . . . . . . .

27

2.2

2.3

2.3.4
2.4

2.5

i

CONTENTS
3 Instruments

29

3.1

Overview of the Suzaku Satellite . . . . . . . . . . . . . . . . . . . . . . . . .

29

3.2

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

33

3.2.1

Overview of the XRT . . . . . . . . . . . . . . . . . . . . . . . . . . .

33

3.2.2

In-ﬂight Performance . . . . . . . . . . . . . . . . . . . . . . . . . . .

35

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

37

3.3.1

Introduction of X-ray Charge Coupled Device . . . . . . . . . . . . .

37

3.3.2

Overview and Speciﬁcations of XIS . . . . . . . . . . . . . . . . . . .

38

3.3.3

In-ﬂight Performance . . . . . . . . . . . . . . . . . . . . . . . . . . .

41

Suzaku Advantages in the Study of the Galactic Center . . . . . . . . . . . .

44

3.3

3.4

4 Study of Sagittarius B

45

4.1

Observation and Data Reduction . . . . . . . . . . . . . . . . . . . . . . . .

45

4.2

Data Analysis and Results . . . . . . . . . . . . . . . . . . . . . . . . . . . .

46

4.2.1

Overall Features of the X-Ray Images . . . . . . . . . . . . . . . . . .

46

4.2.2

Complex Components in the Spectra . . . . . . . . . . . . . . . . . .

48

4.2.3

Model Constructions and Fitting . . . . . . . . . . . . . . . . . . . .

52

Discussion for Sgr B . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

58

4.3.1

Justiﬁcations of Model Constraints and Assumptions . . . . . . . . .

58

4.3.2

Line-of-Sight Distribution of Sgr B MCs in the GCPE . . . . . . . . .

59

4.3.3

XRN Features in Sgr B . . . . . . . . . . . . . . . . . . . . . . . . . .

61

4.3.4

The Correlation of CS Line and 6.4 keV Line . . . . . . . . . . . . . .

62

4.3

5 Study of Sagittarius C

65

5.1

Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

65

5.2

Data Analysis and Results . . . . . . . . . . . . . . . . . . . . . . . . . . . .

66

5.2.1

66

Images . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

- ii -

CONTENTS

5.3

5.2.2

Spectral Model . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

68

5.2.3

Spectral Fitting . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

69

Discussion for Sgr C . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

73

5.3.1

Association with the Radio MCs . . . . . . . . . . . . . . . . . . . . .

73

5.3.2

Two-temperature Structure of the GC Plasma and the XRN Parameters 74

5.3.3

Time Variability of the 6.4 keV Line and the Flare Property of Sgr A* 75

6 Study of Sgr D & Sgr E

77

6.1

Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

77

6.2

Data Analysis and Results . . . . . . . . . . . . . . . . . . . . . . . . . . . .

78

6.2.1

Images . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

78

6.2.2

Spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

78

Discussion for Sgr D and Sgr E . . . . . . . . . . . . . . . . . . . . . . . . .

80

6.3.1

Association with the Radio MCs . . . . . . . . . . . . . . . . . . . . .

80

6.3.2

Origin of the 6.4 keV Line . . . . . . . . . . . . . . . . . . . . . . . .

81

6.3

7 Overall Results & Discussion

83

7.1

Longitude Distribution of Physical Parameters . . . . . . . . . . . . . . . . .

83

7.2

Face-on Distribution of XRNe in the GC Region . . . . . . . . . . . . . . . .

85

7.2.1

Computation of Line-of-sight Position . . . . . . . . . . . . . . . . . .

85

7.2.2

Association with the CMZ . . . . . . . . . . . . . . . . . . . . . . . .

90

Active History of Sagittarius A* . . . . . . . . . . . . . . . . . . . . . . . . .

93

7.3.1

Derivation of X-ray Light Curve . . . . . . . . . . . . . . . . . . . . .

93

7.3.2

Behaviours of the Past Flare State . . . . . . . . . . . . . . . . . . .

94

7.3.3

Interpretation as Low-Luminosity AGN . . . . . . . . . . . . . . . . .

97

7.3

8 Conclusions

99- iii -

CONTENTS
References

100

List of Figures and Tables

106

Acknowledgment

113- iv -

Chapter 1
Introduction & Thesis Composition
Thanks to its vicinity, the central several-hundred-parsec region of our Galaxy (GC;
Galactic center) provides an unique environment for detailed observational studies on highenergy phenomena associated with the super massive black hole (SMBH). Two major discoveries of the GC region are represented by the iron K-shell lines. One is the hot plasma of
107−8 K emitting ionized iron (Fe XXV Kα) line at 6.7 keV (e.g., Koyama et al. 1986, 2007c).
This hot plasma extends over the entire GC region for about 2◦ with monotonic decrease
from the center (Uchiyama et al. 2011, 2012). Another prominent component is the 6.4 keVline of neutral iron (Fe I Kα) found from cold (typically ∼70 K) giant molecular clouds(MCs) in the central molecular zone (CMZ; Morris & Serabyn 1996), such as Sagittarius(Sgr) B, Sgr C, and Sgr A (e.g., Koyama et al. 1996; Murakami et al. 2001; Muno et al.
2007). The spectra of these 6.4 keV bright MCs exhibit common features of large equivalent
width (EW ≥ 1 keV) at 6.4 keV and strong absorption (NH ≥ 1023 H cm−2 ) at the continuum X-rays below 5 keV. Rapid 6.4 keV time variabilities, consistent with injecting sourceof light speed, are also observed from the MCs (e.g., Inui et al. 2009; Ponti et al. 2010).
These results strongly suggest external X-rays as the primary source producing the 6.4 keVline rather than the charged particles (electron or proton; Yusef-Zadeh et al. 2007; Dogiel etal. 2011).
In the context of the X-ray origin, a luminous external source emitting hard X-rays(over the Fe absorption edge at 7.1 keV) is required to irradiate the MCs and generate
the ﬂuorescent 6.4 keV emission, i.e., the X-ray reﬂection Nebulae (XRNe; Koyama et al.
1996; Sunyaev et al. 1993). Since no present Galactic X-ray source is bright enough to
explain the observed 6.4 keV intensities from XRNe, the SMBH with a mass of 4×106 M⊙(Ghez et al. 2008) sited in the GC (i.e., Sgr A*) is believed to be the most reasonablecandidate (e.g., Koyama et al. 1996; Ponti et al. 2010; Nobukawa et al. 2011). Previous
works proposed past ﬂare from Sgr A*, with X-ray (2–10 keV) luminosity on the order of1

CHAPTER 1. INTRODUCTION & THESIS COMPOSITION
LX ∼ 1039 erg s−1 , occurred several hundreds years ago. However, this picture is not clear
due to poor understanding of the three-dimensional position of the MCs.
In this thesis, we present an original self-consistent methodology to measure the threedimensional (line-of-sight) position of multiple XRNe and quantitatively investigate the past
active history of Sgr A*. Following is the composition. In chapter 2, we review the background knowledge of the GC region to introduce the objectives of this thesis. In chapter 3,
we describe principles and speciﬁcations of the X-ray satellite ”Suzaku” and the onboard
instruments used for the data analysiss. In chapter 4, we construct a spectral model for the
Galactic center diﬀuse X-rays and propose a new methodology to investigate positions of
the Sgr B MCs along the line of sight: a correlation study of the ﬂux of the 6.4 keV linefrom the MCs and the line-of-sight absorption of the Galactic center plasma emissions. In
chapter 5, we study the XRNe in the Sgr C region and testify the spectral model proposed in
Sgr B. We investigate the short-term time variability of the 6.4 keV line. We also establish
the spectral composition for the Galactic center plasma emission. In chapter 6, we apply
the model developed in Sgr B and Sgr C (chapter 4–5) to the Sgr D and Sgr E regions
located near the boundary of the CMZ. We probe and investigate the farthest molecular
clouds from the GC for the 6.4 keV emissions. In chapter 7, we summarize and analysis
the results obtained in chapter 4–6 to quantitatively determine the line-of-sight positions
for the multiple XRNe/MCs over the GC and discuss their three-dimensional distribution inassociation with the CMZ. We ﬁnally derive an X-ray light curve for Sgr A* and discuss its
active history over the past six centuries.
Throughout this thesis, we use the Galactic coordinates, where the horizontal and vertical axes are the Galactic longitude (l) and latitude (b), respectively. The positive and negative directions of l correspond to east and west, respectively. We adopt 8.0 kpc (Ghez et al.
2008) as the distance between the Sun and Sgr A* (l = −0.◦ 056, b = −0.◦ 046; Yusef-Zadeh
et al. 1999), then 1.0 degree corresponds to 140 pc (or 456 ly) at the GC. The parameter
uncertainties are quoted at the 90% statistical conﬁdence-level (1.64 σ) range unless noted
otherwise.

-2-

Chapter 2
Review
In this chapter, we review the important topics of the Galactic center to explain theobjectives of this thesis.

Contents
2.1

2.2

2.3

The Milky Way Galaxy and the Center Region . . . . . . . . .

4

2.1.1

Basic Information of the Galaxy . . . . . . . . . . . . . . . . . . .

4

2.1.2

Central Molecular Zone . . . . . . . . . . . . . . . . . . . . . . . .

8Super Massive Black Hole, Sagittarius (Sgr) A* . . . . . . . . .

8

2.2.1

Presence and Basic Properties. . . . . . . . . . . . . . . . . . . .

8

2.2.2

X-ray Observations of Sgr A* . . . . . . . . . . . . . . . . . . . . .

10Thermal X-rays from the Galactic Center . . . . . . . . . . . . .

14

2.3.1

Early X-ray Observations on the Galactic Plane

14

2.3.2

The Fe XXV Kα (6.7 keV) Diﬀuse Emission and the GC Hot Plasma 16

2.3.3

The S XV Kα (2.45 keV) Diﬀuse Emission and Possible Presence of
Large-scale Low-temperature Plasma . . . . . . . . . . . . . . . . .

17

Latest Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

18

Fe I Kα Diﬀuse Emissions in the Galactic Center . . . . . . . .

19

2.4.1

The 6.4 keV Line and its Origins. . . . . . . . . . . . . . . . . .

19

2.4.2

Observational Results of the 6.4 keV Emissions . . . . . . . . . . .

23

2.4.3

X-ray Reﬂection Nebula and Past Activities of Sgr A* . . . . . . .

24

Objectives and Approach of the Thesis . . . . . . . . . . . . . .

27

2.3.4
2.4

2.5

3

. . . . . . . . . .

CHAPTER 2. REVIEW

2.1

The Milky Way Galaxy and the Center Region

2.1.1

Basic Information of the Galaxy

Composition and Size
The Milky Way is the galaxy hosting our Earth and solar system, which is usually called”the Galaxy” (ﬁgure 2.1). As revealed by early observations in the 20th century (e.g., Oort etal., 1958; Okuda et al., 1977), the Galaxy is a typical spiral galaxy with concentric structure
consisting of three main components: the disk, the bulge, and the halo. The disk has apancake-like shape which is approximately ∼30 kpc in diameter and ∼0.3 kpc in thickness.
The bulge is a bar-like spheroid, with major radius of ∼2 kpc and minor radius of ∼1.0 kpc,
sitting on the disk. The bulge (bar) orientation is about 20◦ with respect to our line of sight.
The very central ∼500 pc region of the bugle is the ”Galactic center (GC)”. The halo is a
big sphere with diameter of ∼50 kpc wrapping up the disk and the halo. The solar system islocated on the disk and is about ∼8 kpc away from the center. A global view of the Galaxy
is shown in ﬁgure 2.2.

Milky Way

Scorpius

Sagittarius

Centaurs

Vela

Fig. 2.1.—: Optical picture of the Milky Way galaxy (Laustsen et al. , 1987).

The Galaxy consists mostly of the stars, which dominate the entire galactic mass. The
disk, including the bulge, contains ∼100 billion stars (corresponding to mass of ∼ 1011 M⊙ )with exponential spatial distribution dropping from the center. The scale length and heightof the stellar-mass density are ∼3.5 kpc and ∼0.3 kpc, respectively. The bulge has ∼10billion stars distributed in its bar potential. The stellar density is ∼1 star pc−3 near thecenter and ∼0.1 star pc−3 near the Sun.

-4-

2.1. THE MILKY WAY GALAXY AND THE CENTER REGION(a) Face-on(b) Edge-on

Halo

Halo

Disk

+30º

Scutu

~4 kpc

m

+180º
Sun

★
GC

l = 0º

Bulge
Bar

West

Galactic longitude

m Ar

East

Bulge
Disk

★

Sun

GC (~0.5 kpc)

~8 kpcA
ter
Ou~30 kpc
Sagittarius Arm

rm

Perse
us A

–30º

rm

Fig. 2.2.—: Schematic global view of the Galaxy. (a) Face-on view. (b) Edge-on view.

Interstellar Medium
Between the stars is the interstellar medium (ISM), whose mass is ∼5% with respect
to the stelar mass. The ISM mainly composed of the hydrogen (H I, H2 , and H II) gases andthe dust. The component ratio between gas and dust is ∼100:1.

Fig. 2.3.—: Optical extinction (AV )map of the Galaxy (Dobashi et al., 2005).

The ISM is concentrated on the disk, as observed in diﬀerent wavelengths. One fundamental parameter is the equivalent hydrogen column density (NH ) along the line of sight. In
the optical/infrared band, lights from stars are absorbed and scattered by the dust, which re-5-

CHAPTER 2. REVIEW
sults in the interstellar extinction (ﬁgure 2.3; Dobashi et al. 2005). The measured extinction(AV ) can predict the NH from the following relation (e.g., Güver & Özel 2009).
NH [cm−2 ] = 2.2 × 1021 AV [mag](2.1)

In the radio observations, H I line (21 cm) and molecular rotation transition lines are powerfultools to trace the neutral hydrogen atom and molecules (e.g., CO) in the ISM. The line
intensity and its doppler shift tell the column density and the line-of-sight velocity (i.e.,radial velocity) of the gas, respectively. For example, we show a CO map in ﬁgure 2.4a(Dame et al., 2001). The NH of H I gas is expressed by integration of the observed line∫
intensities along velocity ranges ( T dv),∫

NH I [cm−2 ] = 1.8 × 1018

T21 cm dv [K km s−1 ](2.2)

where T21 cm is the intensity of the 21-cm line. The H2 molecule is usually investigatedindirectly from the coexisting molecules such as CO and CS; the NH can be estimated as
below.
NH2 [cm−2 ] = 0.7 × 1020
NH2 [cm−2 ] = 4 × 1021∫∫

TCO dv [K km s−1 ]TCS dv [K km s−1 ]

(2.3)
(2.4)

where TCO and TCS is the intensity of the 12 C18 O (J = 1 → 0) line and the 12 C 32 S (J = 1 → 0)∫
line, respectively. The conversion factors ( [NH2 ]/[ T dv]) are values reported by Torii
et al. (2010) for the CO line and Tsuboi et al. (1999) for the CS line. Note that theseparameters contain systematic uncertainties depending on the molecular composition (H, C,
and O) ratio and the critical densities (i.e., optically thin or thick conditions). On the other
hand, the X-ray absorption is a direct way to measure the NH of the ISM, which is about∼6×1022 H cm−2 between the Sun and the Galactic center (e.g., Sakano et al. 2002).
3 Dimensional Dynamics and Structure
The plot of radial velocity (VLOS or VLSR ) against the Galactic longitude (l-V diagram)shows an inclined bar-like structure (ﬁgure 2.4b), which indicates that the entire Galaxy
is rotating around the center basically in a circular motion, in spite of some local random
motions. From the solar point of view, ISM on the east side (l >0) of the center is moving
away (VLOS >0) while ISM of the west side is moving toward. The rotation velocity near thesolar system (LSR; local standard of rest) is V0 ≈ 220 km s−1 , at distance of R0 =8 kpc fromthe GC (e.g., Yusef-Zadeh et al. 1999).
The 3 dimensional (line-of-side) distribution of the Galactic ISM is a basic and very
curious topic, but is not available for direct observation because we are located inside the-6-2.1. THE MILKY WAY GALAXY AND THE CENTER REGION(b)

LSR redial velocity (km s–1)

Galactic Latitude (degree)(a)Galactic longitude (degree)

Galactic longitude (degree)

Fig. 2.4.—: Radio CO (1–0) map of the Galaxy (Dame et al., 2001). (a) Spatial distribution on
the Galactic Plane. (b) Velocity distribution along the longitude (l-V diagram).

(a)

(b)

Fig. 2.5.—: (a) Radial velocity ﬁeld on the Galactic plane calculated with the rotation curve(Nakanishi & Sofue, 2003). (b) Face-on map of molecular (H2 ) gas in the Galaxy (Nakanishi &Sofue, 2006).

-7-

CHAPTER 2. REVIEW
Galaxy. Assuming that the motion of the Galactic ISM is purely circular, the distance fromthe Galactic center (R) is a function of the radial velocity (V ) as below,
VLOS (l) =

[R

0

R

]

v(R) − V0 sin (l)(2.5)

where v(R) is the known rotation curve (e.g., Clemens 1985; Dehnen & Binney 1998). By
the de-convolution of the observed l-V diagram, Nakanishi & Sofue (2006) investigated theface-on view of ISM in the Galaxy (see ﬁgure 2.5).

2.1.2

Central Molecular Zone

The GC region is in complex with many molecular clouds (MCs), emission nebulae (HII
regions), supernova remnants (SNR), non-thermal radio emissions, and so on. The main< 1.5◦ ; Morris & Serabyn 1996) which
component is the central molecular zone (CMZ, |l| ∼contains molecular mass of ∼ 3 × 107 M⊙ (Dahmen et al. 1998), approximately corresponds
to 10% of the whole molecular gas in the Galaxy. Sagittarius (Sgr) B, Sgr A and Sgr C are
the most well-known and massive MCs in the CMZ (see ﬁgure 2.6). The gas temperature anddensity of MCs in the CMZ are typically ∼70 K (ranges from 30 to 200 K) and ≥ 103 cm−3 ,respectively.
The 3D structure of the CMZ has been studied mainly by the radio observations. Usingl-V diagrams of the CO and CS lines and based on the assumption of the circular motions of
the GC region, Sofue (1995) and Tsuboi et al. (1999) predicted the line-of-sight distributionof MCs. On the other hand, Sawada et al. (2004) examined the spatial correlation between
2.6-mm CO line and 18-cm OH absorption line, and investigated the overview of face-on
distribution of MCs without assumption of dynamical motions. However, these results were
not conclusive and could not reach a consistent view of the CMZ (c.f., ﬁgure 2.6b).

2.2

Super Massive Black Hole, Sagittarius (Sgr) A*

2.2.1

Presence and Basic Properties

It is now widely recognized that every galaxy harbors a super massive black hole (SMBH)
at the center. The mass of the SMBHs shows extensive correlation with the mass of the bulge,
with a ratio of Mbh /Mbugle ∼ 10−3 (ﬁgure 2.8; e.g., Magorrian et al. 1998). This suggests
interaction and co-revolution of the SMBH and the bulge in galaxies. The mass of our
Galactic bulge (∼ 1010 M⊙ ) predicts a SMBH existence with mass of ∼ 107 M⊙ at the GC.
-8-2.2. SUPER MASSIVE BLACK HOLE, SAGITTARIUS (SGR) A*

(a)
Sgr C

Sgr B
Sgr D

Sgr A

(b)

Fig. 2.6.—: Radio CS (1–0) map of the central molecular zone (Tsuboi et al., 1999). (a) Spatial
distribution on the Galactic Plane. (b) Velocity distribution along the longitude (l-V diagram).
38(a)

T SU B O I , H A N D A , & U K I T A

(c)

(b)

GCB
F IG. 8.ÈSchematic diagram of the GCB, which is a face-on view from positive latitudeFig. 2.7.—: Diﬀerent predictions on the face-on view of the central molecular zone. (a) Sofue(1995). (b) Sawada et al. (2004). EMR represents the ”expanding molecular ring”. (c) Tsuboi etal. (1999). GCB represents the ”Galactic center bow”.

The SMBH of the Galaxy was early noticed at the continuum radio band as a very-9-

CHAPTER 2. REVIEW

Fig. 2.8.—: Mass correlation between the super mass black hole and its host bulge (Magorrian
et al., 1998).

compact point source inside the Sgr A complex of ∼10 pc (see ﬁgure 2.10a), which coincides
with the dynamical center of the Galaxy. The SMBH is thereafter named Sgr A* and theexistence becomes conclusive after infrared observations of the proper motions of stars (Sstars) within the central region of ∼1′′ ×1′′ (ﬁgure 2.9; e.g., Ghez et al. 2005). From theastrometric and radial velocities (Keplerian orbits) of the S stars , the distance to and themass of Sgr A* is well determined to be R0 = 8.0 ± 0.6 kpc and Mbh∗ = 4.1 ± 0.6 × 106 M⊙ ,respectively (Ghez et al., 2008). The Galactic location of Sgr A* is measured to be atl = −0.◦ 056 and b = −0.◦ 046, by radio VLA observations (Yusef-Zadeh et al., 1999).

2.2.2

X-ray Observations of Sgr A*

After the establishment of presence of the nearest SMBH, Sgr A* is considered as one of
the most precious samples to study the high-energy phenomena of SMBH and its surrounding
environment. Many X-ray observations have been carried out at the GC region. Started
with the Einstein Observatory (Watson et al., 1981), ROSAT (Predehl & Truemper, 1994),
BeppoSAX (Sidoli et al., 1999), and ASCA (Koyama et al., 1996) have detected X-rays in
0.5–10 keV from the inner region of ∼0.◦ 5 near the radio position of Sgr A*. From theseobservations, the X-ray absorption was measured to be NH ≈ 5–8×1022 H cm−2 with thermal
bremsstrahlung model of kB T ≈5 keV. Then, the un-absorbed X-ray luminosity in 2–10 keV- 10 -

2.2. SUPER MASSIVE BLACK HOLE, SAGITTARIUS (SGR) A*

Fig. 2.9.—: Proper motions of S stars around SMBH, Sgr A* (Ghez et al., 2005).

(a)

(b)

898

Galactic plane orientationB A G A N OF F E T A L .

900

V ol. 591(c)

B A G A N OF F E T A L .

Sgr A East
Circumnuclear
Disk

60"

Sgr A*

Northern Arm
of Sgr A West

Sgr A West

2 pc

Fig. 2.10.—: X-ray Image of the vicinity of Sgr A* by Chandra (Baganoﬀ et al., 2003). (a)Schematic view of the main constituents of the Sgr A radio complex. (b) X-ray (0.5–7 keV) imageof the central 8.′ 4×8.′ 4 region around Sgr A*. (c) same as (b) but of the central 1′ ×1′ region.

was estimated to be LX ≤ 1035 –1036 erg s−1 , as the upper limit of Sgr A* because of theseobservations could not exclude contaminations from nearby sources. In the hard X-ray
observations (e.g., Spartan-1, Kawai et al. 1988; Granat, Sunyaev et al. 1993), the presenceof a long-term variable source near the position of Sgr A* was suggested and the average
4–20 keV luminosity was reported to be LX ∼ 1036 erg s−1 .

- 11 -

CHAPTER 2. REVIEW

(a)

(b)

Fig. 2.11.—: X-ray spectra of Sgr A* and ﬁtting with absorbed power-law models. (a) Spectrum
from the inner 1.′′ 5 radius circle by Chandra (Baganoﬀ et al., 2003). (b) Spectrum from the inner
10′′ radius circle by XMM-Newton (Porquet et al., 2003).

The Chandra X-ray observatory (c.f., Weisskopf et al. 1996), with the unprecedented
high-angular resolution of ∼0.′′ 5, resolved the individual sources inside the inner 1′ ×1′ region.
Baganoﬀ et al. (2003) identiﬁed X-rays from the radio position of Sgr A* within ∼0.′′ 3, theprobability of a false match was estimated to be less than 0.5%. The direct observation of
Sgr A* allowed them to extract the pure spectrum without large contamination from nearby
sources. The Sgr A* spectrum was well explained by an absorbed power-law model with22
large NH = 9.8+4.4
H cm−2 and Γ = 2.7+1.3−3.0 ×10−0.9 (see ﬁgure 2.11a). The X-ray luminosity of+3.0
33
Sgr A* in 2–10 keV was LX =2.4−0.6 × 10 erg s−1 in the quiescent phase.
Short X-ray ﬂares have been found from Sgr A*. Baganoﬀ et al. (2001) reported the35
ﬁrst ﬂare detection in 2000 October, with duration of ∼10 ks and LX =1.0+0.1
erg s−1 ,−0.1 × 10
which is about 45 times brighter than that in the quiescent phase. The biggest ﬂare so far
was detected by XMM-Newton in 2002 October (Porquet et al., 2003), which exhibited a35
duration of ∼2.7 ks and LX of 3.6+0.3
erg s−1 , i.e., an increase by a factor of ∼160.
−0.4 × 10
The photon index of Sgr A* spectrum varies in diﬀerent ﬂares, which was Γ = 1.3+0.5−0.6 in 2000
+0.3
and Γ = 2.5−0.3 in 2002.
As reviewed above, direct X-ray observations toward Sgr A* over the past two decades
have shown that the present luminosity is in range of LX ∼1033 –1035 erg s−1 , which is extremely low. The bolometric luminosity of a ∼ 4 × 106 M⊙ SMBH (Sgr A*) radiating at the
Eddington rate1 is LEdd = 5×1044 erg s−1 . Therefore, at present, Sgr A* is about 9–11 ordersof magnitude fainter than its Eddington limit (∼ 10−11 LEdd ), i.e., the quiescent/oﬀ state.
1

Maximum luminosity of the source where the radiation force outwards equals to the gravitational force
inwards: LEdd = 4πGMbh mp c/σT = 1.26 × (Mbh /M⊙ ) 1038 erg s−1 .

- 12 -

2.2. SUPER MASSIVE BLACK HOLE, SAGITTARIUS (SGR) A*

(a)

(b)
~

eposcorr~

Time (s)Time (s)Fig. 2.12.—: Short X-ray ﬂares of Sgr A*. (a) Flare detected in 2000 October by Chandra(Baganoﬀ et al., 2001). (b) Flare detected in 2002 October by XMM-Newton (Porquet et al.,2003).

The present Sgr A* luminosity is even about 6 orders of magnitude lower than the luminosities of other 106−7 -M⊙ class SMBHs sitting in nearby spiral galaxies, which emit X-rays
at LX ∼ 1039 –1040 erg s−1 or ∼ 10−6 LEdd (e.g., M 81, Ishisaki et al. 1996; M 51, Terashima& Wilson 2001). These SMBHs are classiﬁed as the low-luminosity active galactic nucleus(LLAGN), which are described by the model of advection-dominated accretion ﬂow (ADAF;
c.f., Narayan 2005 and references therein). X-ray emission from LLAGN is dominated bythermal Comptonization of hot electrons (kB T ∼100 keV) in optically-thin ADAF gas, which
is distinct from the black-body radiation (optically thick) in the standard thin accretion disk
model (Shakura & Sunyaev 1973).
On the other hand, as indirect observation of Sgr A*, X-ray reﬂections are found fromgiant molecular clouds around Sgr A* and are proposed as echoes of its ﬂares with LX ∼ 1039
erg s−1 occurred in the past (e.g., Koyama et al. 1996). Details will be reviewed in section 2.4.

- 13 -

CHAPTER 2. REVIEW

2.3

Thermal X-rays from the Galactic Center

2.3.1

Early X-ray Observations on the Galactic Plane

Point Sources
X-ray observations on the Galactic plane and central region have found many brightpoint sources, most of which are X-ray binaries related to white dwarfs, neutron stars, andblack holes (e.g., Watson et al. 1981; Kawai et al. 1988; Sakano et al. 2002). The point
sources typically show spectra explained by power-law model with photon index of Γ = 2–3
or featureless thermal bremsstrahlung model with kB T of ∼10 keV. The X-ray (2–10 keV)luminosities of these individual objects are in range from 1035 to 1037 erg s−1 .
Diﬀuse X-rays

(a)

(b)

(c)

Fig. 2.13.—: X-ray observation of the Galactic ridge by the Tenma satellite (Koyama et al.,1986). (a) X-ray spectrum of the Galactic ridge and the detection of the 6.7 keV line. (b) Galactic
distribution of the 6.7 keV line. (c) Galactic distribution of the continuum.

Beside the numerous discrete sources, diﬀuse X-ray emissions (>2 keV) were found to- 14 -

2.3. THERMAL X-RAYS FROM THE GALACTIC CENTER

(a)

(b)

Fig. 2.14.—: X-ray observation of the Galactic center by the Ginga satellite (Yamauchi et al.,1990). (a) Galactic distribution of the 6.7 keV line intensity. (b) Interpretation of the 6.7 keV
Galactic distribution as two-dimensional gaussian.

widely extend from either side of the GC out to the Galactic ridge for several kpc (tens ofdegree) in longitude (e.g., Iwan et al. 1982; Worrall et al. 1982; Warwick et al. 1985; Bleach
et al. 1972). This large-scale emission is called the Galactic ridge X-ray emission (GRXE).
From improved spectra observed with the Tenma satellite, Koyama et al. (1986) and Koyama(1989) discovered thermal emissions with intense 6.7 keV (He-like, Fe XXV) iron line fromthe ridge region at l ∼ 270◦ –330◦ (see ﬁgure 2.13a). The GRXE spectra were described with
optically-thin plasma model with hot temperature of kB T = 5–10 keV.
Subsequently with the Ginga satellite, Yamauchi et al. (1990) performed the 6.7 keVemission-line mapping around the GC, and reported the intense elliptical distribution within
the inner ∼2◦ (ﬁgure 2.14) region. The line intensity at the GC is about 10 times larger
than that from the Galactic ridge (Koyama et al., 1989); this excess can be referred to as
the Galactic center diﬀuse X-rays (GCDX). The 6.7 keV emission from highly-ionized iron
implies the existence of hot plasma with kB T of ∼ 10 keV concentrated in the GC region.
The next breakthrough is by the ASCA satellite which ﬁrstly introduced the photon
counting type of X-ray CCDs (Burke et al., 2001; Koyama et al., 2007a) for the highresolution imaging spectroscopy. From observations of the GCDX from the Sgr A and- 15 -

CHAPTER 2. REVIEW(a)Counts / s/ keV

Sgr B regions, Koyama et al. (1996) revealed a fact that the iron line at ∼6.7 keV is acomplex of three lines: the 6.7 keV (Fe XXV), the 6.9 keV (Fe XXVI), and the 6.4 keV (Fe I)
lines (see ﬁgure 2.15). The 6.7 keV and 6.9 keV lines are thermal X-rays from ionized ironatoms (Galactic center plasma emission), while the 6.4 keV line is emitted from neutral ironatoms by non-thermal processes. These two emissions have completely diﬀerent origins andphysics behind, hence, we explain them separately in the remaining parts of section 2.3 andsection 2.4.

(b)Fig. 2.15.—: (a) X-ray spectra of the GC region by CCDs onboard ASCA (Koyama et al., 1996).
(b) same as (a) but by CCDs onboard Suzaku (Koyama et al., 2007c).

2.3.2

The Fe XXV Kα (6.7 keV) Diﬀuse Emission and the GC Hot
PlasmaThe prominent feature of the GCDX spectrum (see ﬁgure 2.15) is strong K-shell emission lines from highly ionized atoms, such as Fe XXV (6.7 keV) and Fe XXVI (6.97 keV). Twophysical processes have been considered for 6.7 keV and 6.9 keV lines: the collision excitation (CE) and re-emission in an optically thin and high-temperature plasma, or the charge
exchange (CX) recombination during a collision between neutral hydrogen and ionized irons
produced by cosmic-rays (e.g., Tanaka et al. 2000). The 6.7 keV line is mainly composedof the forbidden line (6636 eV) and the resonance line (6700 eV). The CX process produces
stronger forbidden line than the resonance line (Otranto et al. 2006) and therefore a line
center near 6636 eV, while the resonance line is stronger in the CE case and therefore a line
center near 6700 eV (c.f., Koyama et al. 2007c and references therein).
Using the Suzaku satellite, Koyama et al. (2007c) made careful energy measurements- 16 -

2.3. THERMAL X-RAYS FROM THE GALACTIC CENTERof the 6.7 keV and 6.9 keV lines with deep exposure of ∼200 ks near the Sgr A region(ﬁgure 2.15b). The line center at 6.7 keV at was well determined to be 6680±1 eV, whichfavors the CE process in plasma. Moreover, the line-intensity ratios of Fe XXVI Kα (6.97
keV)/Fe XXV Kα (6.7 keV) and Fe XXV Kβ (7.88 keV)/Fe XXV Kα (6.7 keV) are about 0.33
and 0.09, respectively. These values indicate that the plasma is in a collisional ionization
equilibrium (CIE) at the temperature of kB T = 6.5 ± 0.1 keV. The iron abundance of theplasma was determined to be ZFe /Z⊙ ≈ 1. In conclusion, the 6.7 keV emission of the GCDX
is produced during processes in a hot plasma of 7×107 K; we can called it the Galactic centerplasma emission (GCPE), hereafter.
Koyama et al. (1996) made rough estimation for the physical parameters, and then raised
the question on the origin of the GC hot plasma. Adopting a temperature of kB T ≈10 keV
and a size of ∼1◦ (i.e., 140 pc at 8 kpc), the electron (ion) density is ∼1 cm−3 and thetotal thermal energy of the plasma is Eplasma ∼ 1053 –1054 erg. With the sound velocity of1600 km s−1 (at kB T =10 keV), the time scale for the plasma to expand over 140 pc isτplasma ∼ 105 years. The thermal temperature (kB T ∼10 keV) is larger than the energy of
Galactic gravity; the rotation velocity of molecular gases induced from the gravity potential
is only ≤ 250 km s−1 (ﬁgure 2.6) compared to the sound velocity of the hot plasma. Hence,the plasma would have escaped from the GC region if there were not an huge energy supply.
The required energy rate is about Eplasma /τplasma ∼ 1041 –1042 erg s−1 . As a candidate forthe energy origin, one supernova (SN) event every 100–1000 years (i.e., 10−2 –10−3 SN yr−1 )within the vicinity of GC can somehow explain that energy rate numerically. However,the supernova origin seems to be unrealistic because the supernova rate in the GC (∼ 10−4
SN yr−1 ; estimation from Uchiyama et al. 2012, Diehl et al. 2006, and references therein)
is much lower than the above expectation and the reported temperatures of the GC SNRs
(kB T ∼1 keV; e.g., Nobukawa et al. 2008; Tsuru et al. 2009; Mori et al. 2008; Sawada et al.
2009 ) are also diﬃcult to explain that of the GC plasma (kB T ≈6.5 keV). Still, the energy
origin of the GC hot plasma is a mystery if it is truly diﬀuse.

2.3.3

The S XV Kα (2.45 keV) Diﬀuse Emission and Possible
Presence of Large-scale Low-temperature Plasma

Beside the iron lines, the GCDX spectrum (c.f., ﬁgure 2.15a) also displays strong line
emission of ionized sulfur at 2.45 keV (S XV Kα). The 2.45 keV line is mainly produced
from a plasma with lower temperature of kT ∼1 keV (compared with kB T ≈6.5 keV of the
hot plasma represented by the 6.7 keV line), which is close to the temperature of many SNR
candidates found in the GC region (e.g., Nobukawa et al. 2008; Tsuru et al. 2009; Mori et al.
2008). Interestingly, similar to the 6.7 keV emission, the diﬀuse 2.45 keV emission is present
even in the absence of local sources such as SNRs. The 2.45 keV emission probably originates- 17 -

CHAPTER 2. REVIEW
from another large-scale plasma with kB T ∼1 keV. However, the 2–3 keV band is sensitive to
X-ray absorption with NH ∼ 10×1022 H cm−2 from the CMZ (section 2.1.2), which makes theinvestigation of the 2.45 keV line diﬃcult. The properties of the low-temperature plasma,temperature and relation with the high-temperature (6.5 keV) plasma, are still not clear.

2.3.4

Latest Results

Spatial Distribution of the Emission Lines from the GC to the GR
The 6.7 keV emission is known to spread over the GC region for about 1◦ with amonotonic decrease as the distance from Sgr A* increases (Koyama et al. 1989; Yamauchi etal. 1990). Recently, with Suzaku, Uchiyama et al. (2011) performed an intensive investigationon the 6.7 keV (Fe XXV Kα) line proﬁle for the whole GC region (∼5◦ ×2◦ ; see ﬁgure 2.16),
then Uchiyama et al. (2012) extended the work to the other emission lines of highly ionizedsulfur (S), argon (Ar), and calcium (Ca). As shown in ﬁgure 2.17, these works quantitatively
determined the spatial (l) distributions of plasma emissions from the GC to the Galactic ridge(GR). The e-folding lengths of the GC and GR components are measured to be lGC ≈0.6◦ andlGR ≈50◦ , respectively.
Origins as the Source of Thermal Emissions from the GC and the GR
After the discovery of the 6.7 keV and other thermal emissions from the Galactic plane(GR and GC), two origins have been proposed as the source: a large-scale hot plasma of
truly diﬀuse gases or a superposition of unresolved pointed sources.
In a region at (l, b)=(0.◦ 1, −1.◦ 4), Revnivtsev et al. (2009) resolved more than 80%of the 6.7 keV ﬂux into point sources, with the high-angular-resolution and deep-exposure
observation by Chandra. Prime candidates of these point sources are cataclysmic variables(CVs) and active binaries (ABs), whose spectra display irons lines (6.4, 6.7, and 6.9 keV)and indicate high temperature of kB T ∼ 10 keV (e.g., Ezuka & Ishida 1999; Ishida et al. 2009;
Yuasa et al. 2010). Despite some discrepancies of the observed iron-line equivalent width
between CVs/ABs (EW6.7 keV ∼200 eV) and the GR (EW6.7 keV ∼500 eV; e.g., Uchiyama
et al. 2012; Yamauchi et al. 2009), an assembly of numerous discrete faint X-ray stars arepresently proposed as the main origin of the GR emission (Revnivtsev et al. 2009; Yuasa et
al. 2012).
On the other hand, in the GC region, the 6.7 keV proﬁle demonstrates the excess of∼4 times over that of the stellar mass distribution (Uchiyama et al., 2011). Also from thestudy of equivalent width of the 6.4 keV and 6.7 keV lines, the point-source contribution to- 18 -

2.4. FE I Kα DIFFUSE EMISSIONS IN THE GALACTIC CENTER(a) 2–3 keV
Galactic latitude b (degree)

0.2
0.0
-0.2
-0.4
1.0

0.0

359.0

0.0

359.0

(b) 6.7 keV
0.2
0.0
-0.2

-0.4
1.0

Galactic longitude l (degree)
Fig. 2.16.—: X-ray band maps of the GC plasmas by Suzaku (constructed from Uchiyama et al.
2011, 2012). (a) The 6.7 keV (6.5–6.8 keV) band for Fe XXV Kα. (b) The 2–3 keV band for S XVKα and continuum emission.

the 6.7 keV emission near Sgr A is ∼1/6, estimated by Koyama et al. (2008). These results
conﬁrmed the presence of large-scale diﬀuse thermal plasma as the dominating source of the
6.7 keV emission in the GC.

2.4

Fe I Kα Diﬀuse Emissions in the Galactic Center

2.4.1

The 6.4 keV Line and its Origins

The GCDX has another main component, non-thermal emissions of the 6.4 keV linefrom neutral irons (Fe I Kα) and the related continuum X-rays (i.e., 6.4 keV emissions). As
shown in ﬁgure 2.18, the 6.4 keV line has been found from giant molecular clouds (MCs)such as Sagittarius (Sgr) B, Sgr A, and Sgr C (e.g., Koyama et al. 1996; Murakami et al.
- 19 -

CHAPTER 2. REVIEW-6Photons cm-2 s-1 arcmin(a)

Fe XXV K
10-5
Photons cm-2 s-1 arcmin-2

10

-2

10

-5

10-7
10

-8

10-9

0.01

0.1

1

10

100

*

10

(b)

S XV K

-6

10-7
10

-8

10-9

0.01

0.1

1

10

100

*

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

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

Fig. 2.17.—: Intensity proﬁles of the plasma emission lines from the GC by Suzaku (Uchiyama
et al., 2012). (a) Proﬁle for the 6.7 keV line ( Fe XXV Kα). (b) Proﬁle for the 2.45 keV line ( S XV
Kα).

2001; Park et al. 2004; Muno et al. 2007). Origins of X-ray and sub-relativistic charged
particles (electron or proton) have been proposed for producing the 6.4 keV emissions.
X-ray Origin
The 6.4 keV line and the related continuum emission are respectively due to ﬂuorescence and Thomson scattering taking place inside a MC, as consequences of irradiation by
external X-rays. For the ﬂuorescence line, the injecting photons have to be harder thanE ≥7.1 keV, i.e., iron absorption edge, to generate an electron vacancy in the K-shell. Thenthe characteristic X-rays are re-emitted isotropically. Assuming an power-law model withX
photon index of Γ= 1.5 for the irradiating X-ray source, the resultant 6.4 keV ﬂux (F6.4) canbe approximated to the following form (c.f., Sunyaev & Churazov 1998; Capelli et al. 2012).

X≈ 1.2 × ZFe
F6.4Ω
τT I8 keV [photon cm−2 s−1 ]
2
4πRGC

(2.6)

Here Ω is the solid angle of the target cloud (with projected area S) seen from the illuminating
source (at distance D), i.e., Ω ≈ S/D2 . The other parameters of RGC , ZFe , τT , and I8 keV
are the distance to the GC, the iron abundance w.r.t solar, optical depth due to Thomson
scattering, and the source ﬂux at 8 keV (in unit of photon s−1 keV−1 ). On the other hand,the continuum emission depends on the angle of Thomson scattering (θ); its ﬂux at 6.4 keV(C6.4 ) is written in the following approximate expression.

- 20 -

2.4. FE I Kα DIFFUSE EMISSIONS IN THE GALACTIC CENTER

Fig. 2.18.—: Intensity distribution of the 6.4 keV emission in the GC by ASCA (Koyama et al.,1996).

C6.4 ≈ 1.17 × (1 + cos2 (θ))Ω
τT I8 keV [photon cm−2 s−1 keV−1 ]
2
4πRGC

(2.7)

X
Then, the equivalent width of the 6.4 keV line (EW6.4 keV ), deﬁned as F6.4
/C6.4 , is expressedin following equation.

X
EW6.4
keV ≈

1.0 × ZFe[keV]
1 + cos2 (θ)(2.8)

Considering the uncertainty of ZFe near GC (∼1–2 solar; e.g., Koyama et al. 2007c; Nobukawa
et al. 2010) and the range of cosθ (i.e., 0–1), EW6.4 keV in X-ray origin should be larger than
∼0.5 keV. In addition, strong absorption features in the spectrum, bending continuum shape
below 5 keV and iron absorption edge at 7.1 keV, are also expected because X-ray interactions
occur deep inside the MCs (c.f., NH for τT = 0.1 − 1 is 1023 − 1024 H cm−2 ).
- 21 -

CHAPTER 2. REVIEWElectron Origin
Another possible process to produce the 6.4 keV line is the electron (E ∼ 100 keV)
impact on the MCs. Many synchrotron-emitting radio ﬁlaments have been found from the
GC region (e.g., LaRosa et al. 2000; Law et al. 2008), which suggests the presence of electrons
up to ∼ 100 keV. In contrast to the X-ray origin, the electron-MC impact process produces
weak 6.4 keV feature with EW6.4 keV of (250–300)×ZFe eV (e.g., Yusef-Zadeh et al. 2007;
Tatischeﬀ 2003). Since electron with energies below ∼100 keV would be stopped at surfaceof MCs with NH ∼0.4×1022 H cm−2 (Tatischeﬀ 2003), the 6.4 keV ﬂux is expressed by theenergy density of the injecting electron (U ). The 6.4 keV intensity (surface brightness) canbe approximated as below (Yusef-Zadeh et al. 2007 and references therein).

Electron
I6.4
≈
keV

4 × 10−7 ZFe (U) [photon cm−2 s−1 arcmin−2 ]100 eV cm−3

(2.9)

Fig. 2.19.—: Equivalent width of the 6.4 keV line produced by protons and electrons, for thesolar abundance of iron (Dogiel et al., 2011). The horizontal axis shows the range of γ, which isthe proton/electron index of injecting source assuming power-law spectrum (N ∝ E γ ). Einj is themean kinetic energy per proton.

- 22 -

2.4. FE I Kα DIFFUSE EMISSIONS IN THE GALACTIC CENTER
Proton Origin
In principle, the 6.4 keV line can also be produced by bombardment of sub-relativistic
protons (E ∼ 100 MeV). The presence of proton has been implied by discovery of diﬀuse TeV
emission spatially correlated with the CMZ (Aharonian et al. 2006). The primary protons
can be generated by accretion onto the Sgr A* in star capture events (Dogiel et al. 2011;
Chernyshov et al. 2012). The EW6.4 keV has strong dependence (from ∼0.5 keV to ∼10keV) on the proton index of the injecting source assuming power-law spectrum (N ∝ E γ ; seeﬁgure 2.19).
Because of the relatively long lifetime of protons in comparison with the average timeof star capture, the 6.4 keV line emission in the GC is quasi-stationary. In this case, extreme parameters of the proton model (star capture frequency, diﬀusion coeﬃcient, etc) arenecessary (Chernyshov et al., 2012).

2.4.2

Observational Results of the 6.4 keV Emissions

The observed X-ray spectra of many giant MCs (e.g., Sgr B2, Sgr C, Sgr A; see ﬁgure 2.20) exhibit common features of large equivalent width (EW ≥ 1 keV) at the 6.4 keVline and strong absorption (NH ≥ 1023 H cm−2 ) at the continuum below ∼ 5 keV, whichfavor the X-ray origin for the 6.4 keV emissions (e.g., Koyama et al. 1996; Murakami et al.
2000; Nakajima et al. 2009; Park et al. 2004; Nobukawa et al. 2010), rather than the electron
origins (e.g., Yusef-Zadeh et al. 2007).
The X-ray origin has become even more conclusive after the discovery of the rapid
6.4 keV variabilities (see ﬁgure 2.21) in a few years from the small regions of a few lightyears near Sgr B (Koyama et al. 2008; Inui et al. 2009) and Sgr A (Muno et al. 2007; Ponti
et al. 2010). At the same time, the emission by the Thomson scattering is also expected
because the 6.4 keV ﬂuorescence is produced by the hard X-ray excitation above the ironabsorption edge at 7.1 keV. Recently in Sgr B, a-few-years time variabilities of the continuum
hard X-rays (E ≥ 8 keV) in correlation with the 6.4 keV line are reported by Terrier et al.
(2010) and Nobukawa et al. (2011). These fast variabilities indicate interactions in MCs bylight-speed particle, i.e., X-ray.
On the other hand, there has been only one report for 6.4 keV emission consistent
with electron-origin spectrum of low EW and small absorption, which is G 0.16−0.21(EW ∼0.3 keV and NH ∼1×1022 H cm−2 ) found by Fukuoka et al. (2009) near the footof the largest GC radio non-thermal ﬁlament, the radio arc (e.g., LaRosa et al. 2000).

- 23 -

CHAPTER 2. REVIEW(b) M359.47 0.15

0.01

counts/sec/keV(a) Sgr B2

10 3

10 4

2
0
2
5

10

Energy (keV)

Fig. 5.

−0.07 (a) and M 359.47−0.15 (b).

Fig. 2.20.—: X-ray spectra of giant MCs in the GC region after subtraction of the GCPE. (a)spectrum of Sgr B2 (Koyama et al., 2007b). (b) spectrum of M 359.43−0.07 in Sgr C (Nakajima
et al., 2009). The spectral features suggest the X-ray origin (see text).
~~(a) Sgr B2σ=(b) Sgr A, G 0.1–0.1
.

=

± .

×

−

−

−

~

−

. ± .

J =

(c) Sgr A, Bridge

−
−

−
−

±

2

×

−

l b =

=

=

=

α

− −

:ı

:ı

Fig. 2.21.—: Time variabilities of the 6.4 keV ﬂuxes from giant MCs in a few years. (a) Sgr B2(Inui et al., 2009). (b) G0.11−0.11, (c) Bridge (Sgr A; Ponti et al. 2010).

2.4.3

X-ray Reﬂection Nebula and Past Activities of Sgr A*

XRN Model
In the context of the X-ray origin (see ﬁgure 2.22), a luminous external source emitting
hard X-rays is required to irradiate the MCs to produce the ﬂuorescent 6.4 keV emission,i.e., the X-ray reﬂection Nebulae (XRNe; Koyama et al. 1996; Sunyaev et al. 1993). The
estimation of X-ray (2–10 keV) luminosity of the irradiating source (LX ) depends mainlyon the 6.4 keV ﬂux by reﬂection (F6.4 ), the iron column density of the entire MC (NFe =3.3 × 10−5 ZFe NH ), and the distance (D) between the MC/XRN and the source (c.f., Sunyaev& Churazov 1998; Murakami et al. 2000; Muno et al. 2007; Nobukawa et al. 2008; Capelli- 24 -

2.4. FE I Kα DIFFUSE EMISSIONS IN THE GALACTIC CENTER
et al. 2012; also see equation (2.6)). Moreover, there are two cases for LX according√ to√
the relation between the ﬂare duration (∆t) and the cloud size (∼ projected area = S;
typically ∼10 ly).
√
If the ﬂare duration is longer than the light-crossing time of a cloud (∆t > S/c), theentire cloud is illuminated and LX is expressed in the following form.
−1−1
LX ≈ A × F6.4 × D2 × NFe
× S −1 × σFe(2.10)

Here, A is the integration factor (for I8keV in equation(2.6)) depending on the photon index
of the X-ray spectrum of the primary source. S is the projective area of the cloud. σFe is
the known cross section of photoelectric absorption at 7.1 keV.
√
When the ﬂare duration is shorter than the light-crossing time of a cloud (∆t < S/c),
only a part of the cloud (∼ c∆t) is illuminated. Thus, the actual L√X would be higher thanthe estimation in long-ﬂare case (equation (2.10)) by a factor of ∼ S/(c∆t), if assuming auniform density (n) of the cloud. In other words, equation (2.10) corresponds to the averageluminosity of the short ﬂare over the light-crossing time of the cloud.

Face-on View

Y, Line of Sight★

X, Longitude

SourceS

NH ~ n

S

D
Cloud
X-ray
Reflection

T+ t
T

RGC =8 kpc

Equi-delay timecT =(X2+Y2)1/2 +Y<=>
2
Y = –X /2cT + cT/2

ObserverFig. 2.22.—: Schematic face-on view of the geometry for the XRN model.

- 25 -

CHAPTER 2. REVIEW
Unsolved Issues
The observed F6.4 from MCs (e.g., Sgr B2: ∼ 10−4 photons cm−2 s−1 ) require a primarysource with luminosity of Lx > 1037 erg s−1 (Koyama et al. 1996; Nobukawa et al. 2011).
Since nearby Galactic X-ray sources, with masses of ∼ 1M⊙ , are not presently bright enough(LX ≤ 1037 erg s−1 ) to explain the observed 6.4 keV ﬂuxes from XRNe, the SMBH sited in
the GC (i.e., Sgr A*) is believed to be the most reasonable candidate (e.g., Koyama et al.
1996; Ponti et al. 2010; Terrier et al. 2010; Nobukawa et al. 2011; Capelli et al. 2012).
Using the projected distance or poorly known positions between XRNe and Sgr A*,
previous works roughly predicted that Sgr A* was very active several hundred years ago
with the luminosity of Lx ∼ 1039 erg s−1 (Koyama et al. 1996; Nobukawa et al. 2008; Ponti
et al. 2010). This is about 6 order of magnitude higher than the present luminosity in thequiescent state (Lx ∼ 1033 erg s−1 ; Baganoﬀ et al. 2003).
However, the property of the past activity of Sgr A* is still poorly understood. To begin
with, it is not clear whether there was a single ﬂare or multiple ﬂares. Moreover, appropriate
estimation for the luminosity, the occurrence time, and the duration of the ﬂare were notpossible. The distance is the most important factor to calculate the luminosity because otherparameters can be directly determined from the X-ray observation (see equation (2.10)). The
line-of-sight position of XRN deﬁne the light-traveling time from Sgr A*, i.e., the occurrencetime of the ﬂare. So far, these parameters are not well known for MCs in the CMZ (seeﬁgure 2.7 and text in section 2.1.2). Thus, in order to access the past of Sgr A*, we need
the information on the three-dimensional (3-D; projected 2-D plus line-of-sight) geometry of
XRNe.

- 26 -

2.5. OBJECTIVES AND APPROACH OF THE THESIS

2.5

Objectives and Approach of the Thesis

Considering the topics and issues of the GC region as reviewed above, we set the objectives of this thesis as below.
• The nature of un-investigated diﬀuse 6.4 keV emissions and XRNe.
• The nature of the low-temperature plasma emitting the diﬀuse 2.45 keV line.
• The 3-D positions of giant MCs in the CMZ.
• Quantitative information on the past activity (X-ray light curve) of the SMBH, Sgr A*.

Galactic latitude b (degree)

We present an original self-consistent approach based on detailed 3-D spectral modeling andanalysis of various emission components, for multiple regions of giant MCs in the GC region(Sgr B, Sgr C, Sgr D, and Sgr E; see ﬁgure 2.23) with the Suzaku X-ray data.

0.4
0.2

(chapter 6)

6.4 keV Map of the GC– Suzaku/XIS FOV
Sgr B (chapter 4)

Sgr D

Sgr C (chapter 5)
Sgr A*
★

0.0

Sgr E

-0.2

(chapter 6)

-0.4
1.5

1.0

0.5

0.0

359.5

359.0

358.5

Galactic longitude l (degree)

Fig. 2.23.—: The 6.4 keV intensity map of the Galactic center by Suzaku. The black boxes are
the ﬁeld of views (FOV) of the Suzaku/XIS (chapter 3) observations used for the data analyses(chapter 4, 5, and 6) in this thesis. The white circles are the excluded regions for bright point
sources.

- 27 -

CHAPTER 2. REVIEW- 28 -

Chapter 3
Instruments
In this chapter, we review the X-ray satellite ”Suzaku” and the onboard instruments
used for the data analysis in the thesis. We give an overview of the satellite in the ﬁrst
section and describe principles and speciﬁcations of the telescope and the detector in thefollowing sections.

Contents
3.1

Overview of the Suzaku Satellite . . . . . . . . . . . . . . . . . .

29

3.2

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

33

3.2.1

Overview of the XRT . . . . . . . . . . . . . . . . . . . . . . . . .

33

3.2.2

In-ﬂight Performance . . . . . . . . . . . . . . . . . . . . . . . . . .

35

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

37

3.3.1

Introduction of X-ray Charge Coupled Device . . . . . . . . . . . .

37

3.3.2

Overview and Speciﬁcations of XIS . . . . . . . . . . . . . . . . . .

38

3.3.3

In-ﬂight Performance . . . . . . . . . . . . . . . . . . . . . . . . . .

41

Suzaku Advantages in the Study of the Galactic Center . . . .

44

3.3

3.4

3.1

Overview of the Suzaku Satellite

Following the successful Ginga and ASCA satellites, Suzaku (Mitsuda et al., 2007) is theﬁfth Japanese X-ray astronomy satellite developed under a Japan–US collaboration. The
spacecraft was launched by the Japanese M-V-6 rocket from the JAXA Uchinoura Space
Center (USC) on July 10, 2005. It was successfully put into a near-circular orbit with an
29

CHAPTER 3. INSTRUMENTS
apogee of 568 km, an inclination of 31.9 degrees, and an orbital period of about 96 minutes.
Figure 3.1(a) shows a schematic view of Suzaku in orbit. The spacecraft length is 6.5 m
along the telescope axis after extending the extensible optical bench (EOB) and the total
weight is ∼ 1700 kg, which makes it the largest and the heaviest Japanese X-ray satellite atpresent.

Fig. 3.1.—: (a) Schematic view of the Suzaku satellite (Mitsuda et al., 2007). (b) Schematiccross-section of the Suzaku internal structure after the EOB extension (Mitsuda et al., 2007). (c)The orbit of Suzaku around Earth.

The power supply is from solar panels ﬁxed on the satellite, which provide the total- 30 -

3.1. OVERVIEW OF THE SUZAKU SATELLITE
electric power of ∼ 660 W. The Sun-angle constraint to the pointing direction is 65◦ –110◦
and therefore most celestial sources are available for observations for about 50 days every 6months. The normal mode of operations have the spacecraft pointing in a single direction
for at least 1/4 day (i.e., 10 ks net exposure time). Thus, most targets are blocked by the
Earth for about 1/3 of each orbit, while some objects near the orbital poles can be observed
nearly continuously. The observing eﬃciency of the satellite, measured after four years ofoperation, is about 43%.

Fig. 3.2.—: Conﬁguration of the detectors in the Suzaku satellite (Suzaku Technical Description,2011). The deﬁnition of the XIS coordinates is also shown.

The scientiﬁc payload of Suzaku consists of three distinct co-aligned detector systems:
the X-ray Imaging Spectrometer (XIS: Koyama et al. 2007a), the Hard X-ray Detector (HXD:
Takahashi et al. 2007; Kokubun et al. 2007), and the X-ray Spectrometer (XRS: Kelley et al.
2007). Each of XRS and XISs is located at the focal plane of the dedicated X-ray Telescope
modules (XRT: Serlemitsos et al. 2007). Figure 3.2 shows the conﬁguration of detectors on
the base panel.
The XISs are X-ray CCD cameras with a large ﬁeld of view (∼18’×18’), moderate spectral resolution, and low background in the 0.3–12 keV band. The XRS is a micro-calorimeterfor high-resolution spectroscopy, but is no longer operational due to a thermal accident soon
- 31 -

CHAPTER 3. INSTRUMENTS
after the launch1 . It will not be used and discussed further. The HXD (Takahashi et al.
2007; Kokubun et al. 2007) is a non-imaging collimated hard X-ray instrument sensitive inthe 10–600 keV band. A summary of the Suzaku capabilities is given in table 3.1.
For the scientiﬁc objectives of this thesis, we used the Suzaku data obtained with the
XRT+XIS system. We describe their details in the following sections.
Table 3.1—: Summary of the Suzaku capabilities.

Space craft

XRT

Orbit apogee
Orbital period
Observing eﬃciency
Focal length
Field of viewEﬀective area

XIS

HXD

Angular resolution
Field of view
Bandpass
Pixel grid
Pixel size
Energy resolution
Eﬀective area
Time resolution
Field of view
Bandpass
Energy resolution
Eﬀective area
Time resolution

568 km
96 min∼ 43%
4.75 m17′ at 1.5 keV
13′ at 18 keV
440 cm2 at 1.5 keV
250 cm2 at 8 keV
2′ (HPD)
17.′ 8 × 17.′ 8
0.2–12 keV
1024 × 1024
24 µm × 24 µm∼130 eV at 6 keV (FWHM)
330 cm2 (FI), 370 cm2 (BI) at 1.5 keV
8 s (normal mode), 7.8 ms (P-sum mode)′′4.◦ 5 × 4.◦ 5 (≥ 100 keV), 34 × 34 (≤ 100 keV)
10–600 keV (PIN: 10–70 keV, GSO: 40–600 keV)
∼3.0 keV (FWHM; PIN)√
7.6/ EMeV % (FWHM; GSO)
160 cm2 at 20 keV, 260 cm2 at 100 keV
61 µs

1

On August 8, 2005, a thermal short between the helium and neon tanks resulted in the liquid helium
coolant venting to space.

- 32 -

3.2. X-RAY TELESCOPE (XRT)

3.2

X-Ray Telescope (XRT)

3.2.1

Overview of the XRT

Suzaku carries ﬁve light-weight thin-foil XRTs to collect the X-ray photon with large
eﬀect area. Four XRT modules are used for the XIS cameras and are referred to as XRT-I0,
I1, I2, and I3 (ﬁgure 3.3), the remaining one is for the XRS, designated as XRT-S. The focal
length of the XRT-I is 4.75 m. The XRTs are developed jointly by NASA/GSFC, Nagoya
University, Tokyo Metropolitan University, and ISAS/JAXA.

Fig. 3.3.—: (a) Schematic view of the XRT system mounted on the top plate of the Extensible
Optical Bench.(b) Picture of the XRT-I1 module (Serlemitsos et al., 2007).

Here we explain the principles of X-ray reﬂection. X-rays are reﬂected by smoothsurfaces of high-electron-density material at very shallow angles of ≤ 1◦ (grazing incidence).
Since the reﬂection is non-dispersive, circular X-ray mirrors for grazing incidence can focus
X-rays over a broad energy band. In this conﬁguration, only a fraction of the geometricarea of the telescope can be used to collect X-ray photons for imaging, so multiple layers are
combined to increase the eﬀective area.
In Wolter type-I reﬂective optics, X-rays are reﬂected twice sequentially by a parabolic- 33 -

CHAPTER 3. INSTRUMENTS

X-rays
Focal poiFig. 3.4.—: Schematic cross-section of Wolter-I type X-ray mirror. Incident X-rays are reﬂected
twice at very shallow angles and concentrated at the focal point.

surface and a hyperbolic surface (ﬁgure 3.3). The precise Wolter design yields a spatial resolution of ∼ 1 arcsec, such as the X-ray telescope adopted for the Chandra X-ray Observatory(CXO; e.g., Weisskopf et al. 1996). However, this sacriﬁces the eﬀective area, because thesmall surface roughness requires a heavy substrate and limits the number of nesting layers.
On the other hand, it is diﬃcult to construct a precise Wolter geometry with thin foils. In
this case, the primary and secondary reﬂectors are approximated by simple cones, then thealignment inaccuracy and surface roughness limit the spatial resolution to about ∼1′ .
In the tradeoﬀ between the angular resolution and the eﬀective area, Suzaku takes thelatter and adopts the grazing-incidence reﬂective optics consisting of many compactly nested
and thin conical elements, rather than the Wolter type-I reﬂective optics with few and thicksubstrates. The small thickness of the reﬂectors on Suzaku permits high density nesting
and thus provide large collecting eﬃciency with a moderate imaging capability in the energy
range of 0.2–12 keV. The reﬂector surface is coated with Au with a thickness of ≥ 1000 Å forthe high atomic number and therefor the high electron density. The weight of each telescope
units are under 20 kg.
The angular resolutions of the XRTs range from 1.′ 8 to 2.′ 3, expressed in terms of halfpower diameter (HPD), which is the diameter within which half of the focused X-rays (point
source) are enclosed. The angular resolution does not signiﬁcantly depend on the energy of
the incident X-rays in the energy range of Suzaku, 0.2–12 keV.

- 34 -

3.2. X-RAY TELESCOPE (XRT)

3.2.2

In-ﬂight PerformanceFig. 3.5.—: (a) Total eﬀective area of the XRT combined with the quantum eﬃciency of XISs
(Serlemitsos et al., 2007). Those of other X-ray missions are also shown for comparison. (b) Oﬀ-axis
angle dependence (vignetting eﬀect) of the eﬀective area, normalized at the optical axes. The softand hard bands are shown with black and red points, respectively (Serlemitsos et al., 2007). Data
are obtained from the observations of Crab nebula. The solid lines show the simulated model.

As the results of the grazing-incidence reﬂective optics of XRT and four XISs, Suzaku
has moderate eﬀective area in the soft band ( 440 cm2 at 1.5 keV) and large eﬀective area inthe hard band (250 cm2 at 8 keV), in comparison with other X-ray satellites. The energydependent on-axis eﬀective area, including the detector eﬃciency, are shown in the left panelof ﬁgure 3.5. At larger oﬀ-axis angles, incident X-rays are partially obscured by neighboring
reﬂectors and therefore smaller eﬀective area, which is called the vignetting eﬀect. This eﬀect
is prominent at high energy because high-energy photons have smaller critical angles andare more eﬃciently reﬂected at innermost mirrors, where the cone angles is small and thereﬂectors are tightly packed. Figure 3.5 (b) shows the vignetting curve in the soft (3–6 keV)and hard (8–10 keV) bands. Since the eﬀective area is deﬁned by the number of source
photons within the CCD ﬁeld of view (FOV) of a 17.′ 8 square, the model curves show sharp
drops at ∼ 8′ .
The angular resolution was veriﬁed with the SS Cyg data in quiescence state, as atarget of point source (ﬁgure 3.6). The point source image appears to be a cross or abutterﬂy because the reﬂectors are made of four sections (”quadrants”) and therefore thespaces between any two quadrants are dead areas. Figure 3.7a and b show the point spreadfunction (PSF) and the encircled energy function (EEF), respectively. These are typicalindicators of the spatial resolution. Since the primary factor of the XRT spatial resolution- 35 -

CHAPTER 3. INSTRUMENTS
is dominated by the imperfect alignment of the nested reﬂectors, and therefore the PSF isnearly independent on the energy.

(a)

(b)

Fig. 3.6.—: (a) Observed image of SS Cyg (point source) at the optical axis taken with XRT-I0(Suzaku Technical Description, 2011). (b) Simulated image of SS Cyg with XRT-I0. Note that theabrupt drop of the model curves at ∼ 8′ is due to the source approaching the detector edge.

Fig. 3.7.—: Point spread function (PSF; left panel) and the encircle energy fraction (EEF; rightpanel) of the XRT-I0 (for XIS0) at the focal plane (Serlemitsos et al., 2007).

- 36 -

3.3. X-RAY IMAGING SPECTROMETER (XIS)

3.3

X-ray Imaging Spectrometer (XIS)

3.3.1

Introduction of X-ray Charge Coupled Device

Since the ASCA satellite, the charge-coupled device (CCD) has become the standard
detector in the X-ray astronomical missions, i.e., Chandra (Garmire et al., 2003), XMMNewton (Strüder et al., 2001; Turner et al., 2001), and Suzaku (Koyama et al., 2007a). The
CCDs have good performance both in high-resolution imaging (∼20 µm pitch pixel) andhigh energy resolution with fano-limited readout noise (≤ 10 electrons in rms), which enable
scientists to obtain the spatial and spectral information at the same time. This plays a very
important role in studies of most celestial objects including supernova remnants, pulsar wind
nebulae, and cluster of galaxies.
Through the photoelectric absorption in the silicon, CCD converts an incident X-ray
photon to a electron cloud, the amplitude of which is proportional to the energy of the absorbed X-ray photon. The cloud made in the depletion layer in each pixel is then transferredto the gate of the output transistor by three-phase electric potential clocking. In MOS-type
CCDs (ﬁgure 3.10), the gate is made of poli-Si (Metal), the insulator is made of SiO2 (Oxide),and the depletion layer is made of Si (Semiconductor).
In the typical CCD readout architecture, the signal of each pixel is transferred in theform of charge over a long distance to a small number of readout nodes and then converted
into voltage serially for all pixels. The readout noise is substantially reduced by applying
the correlated double sampling (CDS) method at the readout node. The gain is uniform
because of the small number of readout nodes. On the other hand, the time resolution isdetermined by the number of pixels, e.g., the readout time per frame is a few seconds for a
CCD consisting of ∼1 M pixels. Another feature during the transfer is the charge transfer
ineﬃciency (CTI), in which some fractions of the charge are trapped and lost due to thedefects in silicon crystal after radiation damages by the cosmic rays in space. This eﬀect
can be suppressed and well corrected with the charge injection technology such as appliedin Suzaku/XIS (Uchiyama et al., 2009).
After vigorous developments over the past two decades, X-ray CCDs have nearly matured and leaves little room for further essential advancements. In the near future, CCDs
may be replaced by the progressing active pixel sensors (APS; e.g., Lechner et al. 2010;
Prigozhin et al. 2009; Ryu et al. 2011), which access pixels directly and therefore arecapable of more ﬂexible and faster readout (see ﬁgure 3.8).

- 37 -

CHAPTER 3. INSTRUMENTS

Fig. 3.8.—: Readout architectures of the CCD and the active pixel sensor (APS).

3.3.2

Overview and Speciﬁcations of XIS

The Suzaku XIS system (ﬁgure 3.9) is comprised of four MOS-type CCD cameras: threesets of front-illuminated (FI) CCD cameras (XIS 0, 2, and 3) and one set of back-illuminated(BI). Each CCD chip has an imaging area consisting of 1024 × 1024 pixels (1 pixel = 24µm
square) for the 17.′ 8×17.′ 8 ﬁeld of view (FOV). Schematic views are shown in ﬁgure 3.10
and ﬁgure 3.11.

(a)

(c)

(b)
Hood
105 mmBaffleElectricvalve

Bonnet
125 mmDoor
Door calibration
sourceWall calibration
sourceBase
70 mm

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

Flexible print

Fig. 3.9.—: (a) Picture of one of the XIS systems. (b) Cross section of the XIS system (Koyama
et al., 2007a). (c) Picture of the XIS CCD sensor (Koyama et al., 2007a).

The front side of CCD is coated with the gate structure made of Si (0.28 µm) and- 38 -

3.3. X-RAY IMAGING SPECTROMETER (XIS)SiO2 (0.44 µm). The back surface of BI CCD consists of 5 nm HfO2 , 1 nm Ag, and 3 nm
SiO2 , which is made by the chemisorption charging process in order to improves the chargecollection eﬃciency near the back surface. These are the dead layers because X-ray charges
generated there are not collected to the sense node. The dead layers of the FI and BI CCDs
are ∼ 0.7 µm and ∼ 10 nm (see ﬁgure 3.10), respectively. Due to the absence of the gate
structure, BI CCD achieve a high quantum eﬃciency (QE) even below ∼ 1 keV. On theother hand, FI CCDs employ thicker depletion layer than BI CCDs (∼ 76 µm for FI CCDs
and ∼ 42 µm for BI CCD), and therefore have higher QE for high-energy photons. The QEof FI and BI XIS, including the eﬀects described above, are show in ﬁgure 3.12. Each XIS
also has an optical blocking ﬁlter (OBF) to shutoﬀ the optical and UV photons. The OBF
is made of ∼ 100 Å polyimide (C22 H10 N2 O4 ) coated with Al of a ∼ 1200Å thickness.

X-ray

X-ray
~0.7 µm
XIS1
Readout

XIS0, 2, 3

Gate
Insulator~10 nmeDepletion layer
Insulator
Gate

eDepletion layer
Front illumination~76 µmReadout

Back illumination~42 µmFig. 3.10.—: XIS (MOS-type CCDs) in front illumination and back illumination.

When the CCD is operated in the frame transfer mode, the charges are transferred fromthe imaging area (IA) to the frame store (FS) region after an exposure (ﬁgure 3.11). The
charges in the FS region are transferred vertically along the ActY direction and horizontally
along the ActX direction sequentially. In both the IA and FS regions, the CCD chips have aformat of 1024 × 1024 pixels and four read-out nodes. Thus, a CCD chip is composed of four
segments of 256 × 1024 pixels (A–D; ﬁgure 3.11). The pixel size of the IA is 24 µm×24 µm.
In combination with XRT, one pixel corresponds to 1.′′ 04×1.′′ 04 of the sky, and the ﬁeld ofview is 17.′ 8×17.′ 8.

- 39 -

CHAPTER 3. INSTRUMENTS

Quantum efficiency

Fig. 3.11.—: Schematic view of the XIS CCD and the readout system (Koyama et al., 2007a).

1
FI (XIS 0)
BI (XIS 1)

0.5

0

0

5
10
Energy (keV)

Fig. 3.12.—: Quantum eﬃciency of the XIS as a function of incident photon energy. The
black and red curves are for XIS 0 and XIS 1 respectively (Koyama et al., 2007a). The eﬀect ofcontaminant on the OBF is not included.

- 40 -

3.3. X-RAY IMAGING SPECTROMETER (XIS)

3.3.3

In-ﬂight Performance

Operation
The on-orbit operating temperature of XISs is ∼ −90◦ C controlled by Pertier coolers
and always monitored at the JAXA Uchinoura Space Center.
The XISs have two clock modes for the on-board data processing: the Normal modeand the Parallel-Sum (P-sum) mode.
• Normal mode: the exposure time is 8 seconds, and all the pixels on the CCD areread out every 8 seconds, which corresponds to the time resolution. The data of pixels
can be combined with either of the 5×5, 3×3, and 2×2 editing modes.
• P-sum mode: the pixel data from multiple rows are summed in the Y-direction in the
FS region, and the summed value is put in the memory as a single row. The 1024 rowsof data are read out in 8 s, and therefore the time resolution in the is 8/1024 = 7.8 ms.
This mode is useful to avoid the photon pileups in observations of very bright sources,
at the sacriﬁce of the spatial information.
The Galactic center data in this thesis are all obtained in the Normal mode with the
3×3 editing mode.
Energy Calibration
For the calibration of the energy scale (gain tuning), each XIS is equipped with two
radioactive sources (55 Fe), located on the side wall of the housing, to illuminate two far
corners from the read-out nodes. A small fraction of the X-rays from the calibration sources isscattered onto the entire CCD. 55 Fe decays to leave 55 Mn via electron capture mainly (88.5%)
from the K-shell. Subsequently, ﬂuorescent X-ray is emitted, accompanied by an L→Kelectron transition (Kα: 5.895 keV) or an M→K electron transition (Kβ: 6.513 keV). These
two lines are used to calibrate the high-energy range. Suzaku also has regular observations on
1E0102−72, which emits the OVII (0.57 keV) and OVIII (0.65 keV) lines, for the calibrationof the low-energy range.
Charge Injection Capability
The charge transfer ineﬃciency (CTI) has been increasing since the launch due to charge
traps caused by radiation damage (Nakajima et al., 2008; Ozawa et al., 2009). Consequently,- 41 -

CHAPTER 3. INSTRUMENTS

Mn K line FWHM (eV)

the energy resolution has degraded, which was initially of ∼ 130 eV at 5.9 keV in full width at
half maximum. (ﬁgure 3.13). To mitigate the degradation, spaced-row charge injection (SCI)
technique is applied for almost all observations since October 2006. In the SCI technique,electrons are injected from the furthest side from the read-out nodes to ﬁll the traps as”sacriﬁcial charges”. This prevents charges generated by X-rays from being captured by thetraps and actively reduces the CTI. A new calibration method dedicated for SCI observations
has been established (Uchiyama et al., 2009). With this method, the energy resolution stillretains ∼ 155 eV and ∼ 175 eV at 5.9 keV for FI and BI CCDs respectively as of March 2008
(ﬁgure 3.13).
250
200
150

XIS 0

250
200
150

XIS 1

250
200
150

XIS 3

2007

500

Segment A
Segment D

2008

2009

1000
1500
Time after launch (days)

2010

2011

2000

Fig. 3.13.—: Long-term trend of energy resolution (Suzaku Technical Description, 2011).

Contamination
Due to the accumulation of out-gassing from the satellite onto the OBF, the low energy QE has been gradually degraded (Koyama et al., 2007a). The contaminant is mainlycomposed of carbon and oxygen. The time evolution of the contamination has been measured using regular observations of 1E 0102–072. Although the chemical composition of thecontaminant is still uncertain, it is assumed to be C/O=6. In the data analysis, the contamination eﬀect is included in ancillary response ﬁles made by xissimarfgen (Ishisaki etal., 2007).

- 42 -

3.3. X-RAY IMAGING SPECTROMETER (XIS)

Fig. 3.14.—: Non-X-ray background (NXB) spectra of Suzaku XIS 0 (FI) and XIS 1 (BI). The
spectra are constructed from night Earth observations. For comparisons, the background spectraof ASCA, Chandra, and XMM-Newton are also shown.

Non–X-ray background
The XISs are subject to the Non–X-ray background produced by the cosmic-rays onorbit. Figure 3.14 shows the non–X-ray background (NXB) accumulated during night-earth
observations. The data are obtained during the South Atlantic anomaly passages and thecalibration source areas are excluded. The NXB in the 0.5–10 keV band is about 8×10−8
counts s−1 pix−1 for the FI CCDs and 2×10−7 counts s−1 pix−1 for the BI CCD. The ﬂux of
NXB also depends on the geo-magnetic cut-oﬀ rigidity (COR, ﬁgure 3.14), but great concern
on the COR is not necessary for most of the XIS observations. The NXB can be constructed
for any observations with an accuracy of ∼ 5%, which is much smaller than the spatial
ﬂuctuation of the cosmic X-ray background in the 1–7 keV band, and is almost comparable
to that in the 5–12 keV band (Tawa et al., 2008). Suzaku XISs have low backgrounds incomparison with other X-ray missions (see ﬁgure 3.14), thanks to the combination of the
Suzaku orbit and the instrumental design.
Accidents
So far, the Suzaku CCD cameras have experienced two major accidents. On November
9, 2006, XIS 2 suddenly showed an anomaly, in which about 2/3 of the imaging area became
unusable due to a large amount of charge leakage. Although not conclusive, this event was- 43 -

CHAPTER 3. INSTRUMENTS
possibly caused by a micro-meteoroid impact, which is also seen in XMM-Newton and Swift.
Since this anomaly, XIS 2 has been dysfunctional. On June 23, 2009, XIS 0 also showed ananomaly at once. During a Normal clocking operation, a part of the segment A of XIS 0 was
ﬂooded with a large amount of charges, which caused saturation of the analogue electronics.
The anomaly was very similar to that occurred in the XIS 2 in 2006, probably had the same
origin. Data from the segment A of XIS 0 is not available after this event. From the pointof view in statistical expectation, for XISs, two impacts in four years is not too surprising.

3.4

Suzaku Advantages in the Study of the Galactic
CenterAs reviewed in chapter 2, the observation of Fe lines (6.4 keV and 6.7 keV) and hard
X-rays up to 10 keV is essentially to reveal the nature (hot plasmas and X-ray reﬂection
nebulae) and solve the issues in the Galactic center. Suzaku has two big advantages for thisstudy. First, among all the present X-ray satellites and in the band over 5 keV, the Suzaku
instruments (XRT+XIS) have the best overall performances in good energy resolution, low
and stable NXB, and large eﬀective area. Moreover, we have already strategically obtainedrich Suzaku data of the Galactic center; the overall exposure is over 5 Msec and the FOVs
cover a wide region of ∼4◦ × 2◦ . For these reasons, in the following chapters, we analyze theSuzaku data to attain the objectives of this thesis.

- 44 -

Chapter 4
Study of Sagittarius B
In this chapter, we construct a spectral model for the Galactic center diﬀuse X-rays andpropose a new methodology to investigate positions of the Sgr B molecular clouds (MC) along
the line of sight: a correlation study of the 6.4 keV line from the MCs and the line-of-sightabsorption of the Galactic center plasma emissions (GCPE).

Contents
4.1

Observation and Data Reduction . . . . . . . . . . . . . . . . . .

45

4.2

Data Analysis and Results . . . . . . . . . . . . . . . . . . . . . .

46

4.2.1

Overall Features of the X-Ray Images . . . . . . . . . . . . . . . .

46

4.2.2

Complex Components in the Spectra . . . . . . . . . . . . . . . . .

48

4.2.3

Model Constructions and Fitting . . . . . . . . . . . . . . . . . . .

52

Discussion for Sgr B . . . . . . . . . . . . . . . . . . . . . . . . . .

58

4.3.1

Justiﬁcations of Model Constraints and Assumptions . . . . . . . .

58

4.3.2

Line-of-Sight Distribution of Sgr B MCs in the GCPE . . . . . . .

59

4.3.3

XRN Features in Sgr B . . . . . . . . . . . . . . . . . . . . . . . .

61

4.3.4

The Correlation of CS Line and 6.4 keV Line . . . . . . . . . . . .

62

4.3

4.1

Observation and Data Reduction

The Sgr B region, located at east side of the GC (l ∼0.◦ 7), contains the most wellknown and massive MCs (e.g., Sgr B1 and B2) in the CMZ (c.f., section 2.1.2). Deep
multiple pointing observations toward the Sgr B region were performed with Suzaku/XIS
45

CHAPTER 4. STUDY OF SAGITTARIUS B
from October 2005 to September 2007. The observations were made in the normal clocking
mode with read-out cycle of 8 s. We screened the data by excluding the events observed
below the low and day-Earth elevation angles at 5◦ and 20◦ , in addition to the removal ofevents during passage of the South Atlantic Anomaly. The total eﬀective exposure time wasabout 400 ks. The observation log is shown in table 4.1.
We performed the data reduction and analysis using HEADAS software version 6.5. The
calibration database used to calculate the response of XIS and the eﬀective area of XRT was
the version released1 on 2008-06-02. Since the relative gains and response functions of the FI
CCDs are essentially the same, the FI spectra were merged in the spectral ﬁtting. For theimage and spectral analyses described in the following sections, the non-X-ray background(NXB; Tawa et al. 2008) was generated using the night-Earth observation data and was
subtracted from the raw data.
Table 4.1—: Suzaku Observations near the Sgr B Region

Pointing direction∗
α (J2000.0)
δ (J2000.0)hms
GC Sgr B2
100037060 17 47 30. 60 −28◦ 26′ 36.′′ 6
GC Number2
500005010 17h 47m 04.s 63 −28◦ 37′ 46.′′ 2
GC Sgr B East 501039010 17h 48m 04.s 87 −28◦ 21′ 06.′′ 5
Sgr D SNR
502020010 17h 48m 46.s 13 −28◦ 07′ 38.′′ 6
Target name∗

Obs.ID

Obs.date
Exposure(yyyy-mm-dd)(ks)
2005-10-10
76.6
2006-03-27
88.4
2007-03-03
96.4
2007-09-06
139.1

The center of the XIS ﬁeld of view.

4.2

Data Analysis and Results

4.2.1

Overall Features of the X-Ray Images

Figure 4.1 shows XIS images of the Sgr B region in the (a) 6.3–6.5 keV and (b) 2.5–3.5
keV bands. The images were made by co-adding the FI and BI CCDs data, followed bycorrecting the exposure time and the vignetting eﬀects of the XRTs. The data of the CCD
corners illuminated (contaminated) by the calibration sources were excluded. To increase
the visibility, we binned the images with 4×4 pixels and smoothed them with a Gaussian
kernel of σ = 1′ .
1

http://www.astro.isas.ac.jp/suzaku/caldb/- 46 -

4.2. DATA ANALYSIS AND RESULTS

Galactic latitude [degree]

0.10(a) 6. 3 6. 5 keV band
Sgr B2(M0.66)

0.00

M 0.74-0.09
Sgr B1(M0.51)

-0.10

-0.20

Region 1

M 0.74-sub

Region 2

0.75

0.50

0.25

0.75

0.50

0.25

-0.30
1.25

Galactic latitude [degree]

0.10

1.00

(b) 2. 5 3. 5 keV band

0.00

-0.10

-0.20

Profile

-0.30
1.25

1.00

Galactic longitude [degree]
Fig. 4.1.—: XIS band images around the Sgr B region smoothed with a Gaussian kernel of σ = 1′ .
(a) The 6.4 keV band (6.3–6.5 keV), which contains the Fe I Kα line, tracing the dense molecular
clouds. (b) Medium band (2.5–3.5 keV) showing the distribution of the GCDX in low energy. The
CS-line contours (Tsuboi et al. 1999) in velocity range of −50–140 km s−1 are added with green
solid lines in (a) and (b). A hint of anti-correlation between (a) and (b) near l ∼0.◦ 75 would be due
to strong absorption by the molecular clouds. Spectra regions are shown in black solid lines in (a);the spectrum region for M 0.74-sub is the largest circle excluding the small circle of M 0.74-0.09.
The bright point source (SAX J1748.2−2808; Nobukawa et al. 2009) is excluded for Region 1. The
scan proﬁle in ﬁgure 4.2 is made from the large rectangle (0.4 degree × 0.1 degree) in (b).

- 47 -

CHAPTER 4. STUDY OF SAGITTARIUS B
The bright clumps in ﬁgure 4.1a are Sgr B2 (M 0.66−0.02), M 0.74−0.09, and Sgr B1(M 0.51−0.10). These are reported to be candidates of the X-ray reﬂection Nebula (XRN;
Koyama et al. 2007b; Nobukawa et al. 2008). In addition, we found a new diﬀuse structurein the 6.4 keV line at the south of M 0.74−0.09 (hereafter, M 0.74-sub; ﬁgure 4.1a).
The GCDX ﬂux in the 2.5–3.5 keV band (ﬁgure 4.1b) is sensitive to absorption in NH
ranges of 1022 –1024 H cm−2 . This range corresponds to the interstellar absorption toward
the GC region and the intra-cloud absorption of the MCs. Comparing ﬁgure 4.1a and b, wefound a hint of anti-correlation between the 6.3–6.5 keV (6.4 keV line) and the 2.5–3.5 keVband ﬂuxes near l=0.◦ 75: the 2.5–3.5 keV band ﬂux is relatively weak in contrast with thestrong emission of the 6.4 keV line from the Sgr B MCs.
To conﬁrm this anti-correlation, we made one-dimensional ﬂux proﬁles in the 2.5–3.5 keVband and the 6.4 keV-line band along the Galactic longitude of l=1.◦ 0–0.◦ 6 (see ﬁgure 4.2).
For comparison, we also made proﬁle and contours (in ﬁgure 4.1a and b) of the radio CS line
(J = 1 →0) by extracting the velocity-integrated ﬂux in the range of −50–140 km s−1 (Tsuboi
et al. 1999), within which most of the Sgr B MCs is included. The CS-line ﬂux proﬁle isadded in ﬁgure 4.2. For all of the proﬁles, the vertical axis represents the sum-value at the
Galactic latitude b = −0.◦ 15 ± 0.◦ 05 (see the solid rectangle in ﬁgure 4.1b); the ﬂux of each
proﬁle is normalized at the maximum value, and hence the relative ﬂux range is from 0 to 1.
As shown in ﬁgure 4.2, the ﬂux proﬁle of the 6.4 keV line shows a bump at l =0.◦ 65–
0.◦ 85, which is in general agreement with that of the CS line. The 2.5–3.5 keV band ﬂux,
on the other hand, shows an intensity drop at the bump of the 6.4 keV line and the CS line,
which implies absorption due to the MCs. Since the 2.5–3.5 keV band ﬂux is sum of the
GCPE behind the MCs (should be anti-correlated to the 6.4 keV-line ﬂux) and the GCPE
in front of the MCs (no correlation to the 6.4 keV-line ﬂux), the intensity drop would besmeared out to only ∼ 15%. Detailed study of the spectra is given in section 4.2.2. We notethat the angular size of the 6.4 keV bump is slightly narrower than that of the CS line andthat of the intensity drop in the 2.5–3.5 keV band at l ∼ 0.◦ 85 (section 4.3.4).

4.2.2

Complex Components in the Spectra

We obtained X-ray spectra of the Sgr B region in the 0.5–10.0 keV band from the 6solid line regions given in ﬁgure 4.1a. These are Sgr B2, M 0.74−0.09, Sgr B1, M 0.74-sub,Region 1, and Region 2. In ﬁgure 4.3, for brevity, we show 2 typical examples of the spectra.
The spectra have two notable common features (see ﬁgure 4.3). One is that the spectra exhibit many emission lines such as Fe XXV Kα (6.7 keV), Fe XXVI Kα (6.97 keV),Si XIII Kα (1.86 keV), S XV Kα (2.45 keV), and Fe I Kα (6.4 keV). These are K-shell lines- 48 -

Relative flux (normalized)

4.2. DATA ANALYSIS AND RESULTS

1
0.8
0.6
0.4
0.2
0

1.0

0.9

0.8

0.7

0.6

Galactic longitude [degree]
Fig. 4.2.—: One-dimensional proﬁles of the Sgr B region along the Galactic longitude froml = 1.0◦ –0.6◦ at b = −0.15◦ ± 0.05◦ (see the solid rectangle in ﬁgure 1b). Gray open circles and black
ﬁlled squares are data of the 2.5–3.5 keV band and the 6.4 keV band (6.3–6.5 keV) with 1-σ error
bars, respectively. The radio CS line proﬁle (−50–140 km s−1 ) is plotted with the black solid line.
The ﬂux of each data is normalized to its maximum value.

from various atoms in various ionization states and hence the spectra are likely to consist ofmulti-temperature components, but may be reduced to at least three temperatures, which
are responsible for the emission lines of the 6.4 keV (Fe I, cool gas), the 2.45 keV (S XV, hot
plasma), and the 6.7 keV (Fe XXV, very hot plasma). The other feature is that the continuum
shape shows broad local minima near 1–3 keV. Since the absorption of the GCDX spectrum
is in the range of 6×1022 –1024 H cm−2 , the ﬂux should rapidly drop below ∼ 2 keV. The presence of the local minima, therefore indicates that the full-band spectra have an additional
soft X-ray component with no signiﬁcant (small) absorption (the foreground emission). As aconsequence, we decompose the spectra of the Sgr B regions to the following ﬁve components:(1) Very hot plasma responsible for the 6.7 keV line.
(2) Hot plasma responsible for the 2.45 keV line.
(3) Cool gas component that emits the 6.4 keV line.
(4) Low energy (0.5–2 keV) X-rays with small absorption.
(5) The cosmic X-ray background (CXB).
Koyama et al. (2007c) studied component (1) using the iron/nickel line ﬂux ratios of K α- 49 -

-2
-1

5

10

1

5

10

Sgr B2

5

10

5

10

Model B

10 -4

10 -4

-1

-1

10 -3

Model B

10 -3

Counts s keV arcmin-2

M0.74−0.09

12Energy [keV]

-1

10 -3

Model CSgr B2Model C

10 -4

−2 0 2

10 -4

-1

10 -3

Counts s keV arcmin-2-2

M0.74−0.09-1-1Counts s keV arcmin

2
Energy [keV]-2

2
Energy [keV]

-1

-1

10 -4

-1

10 -4

1

0.5

χ

Model A

−2 0 2χ

0.5Counts s keV arcmin

Sgr B210 -3Counts s keV arcmin

Model A

10 -3

M0.74−0.09-1-1Counts s keV arcmin-2

CHAPTER 4. STUDY OF SAGITTARIUS B

0.5

1

2
Energy [keV]

5

10

1

2
Energy [keV]

Fig. 4.3.—: FI spectra of M 0.74-0.09 (left column) and Sgr B2 (right column) as examples fordiﬀerent models. The best-ﬁt spectra and the residuals (χ-distributions) of models A, B, and C
are shown in the top, middle, and bottom rows, respectively. In each spectrum, components of the
NGCE and the XRNE are shown by green and black solid lines; components of GCPE suﬀered by
Abs1 × Abs2 and Abs2 are shown by the red and blue solid lines, respectively.

and K β in the GC region. They concluded that component (1) is a plasma of temperaturekT ≃ 6.5 keV in collisional ionization equilibrium. The iron abundance was determined to be- 50 -

4.2. DATA ANALYSIS AND RESULTS≃ 1 solar2 . We therefore use an APEC model3 for component (1) and express it as APEC1
with ﬁxed temperature and abundances of 6.5 keV and 1 solar, respectively.
With the Suzaku GC survey observations, many SNR candidates were found in the
2.45 keV-line band images. The plasma temperatures are nearly the same at about 1 keV.
(Nobukawa et al. 2008; Mori et al. 2008; Tsuru et al. 2009). The abundances are slightly
diﬀerent in each element, but are consistent with 1 solar. The abundances of S, the mostprominent Kα line emitter, are 0.6–1.3 solar (hereafter, the numbers in the parenthesis are
in 90% error range) for G 0.42−0.04 (Nobukawa et al. 2008), 1.2–1.5 solar for G 359.79−0.26(Mori et al. 2008), and 1.1–2.7 solar for G 359.41−0.12 (Tsuru et al. 2009). Removing
these SNR candidates, we still found the 2.45 keV line prevailing in all of the GC region(component (2)). Thus we use an APEC model for component (2) (APEC2), assuming thatthe temperature and abundances to be ∼ 1 keV (free) and 1 solar (ﬁxed), respectively, thesame values to those of SNR candidates near the GC region.
Component (3) is emission possibly produced by the ﬂuorescence and reﬂection from theMCs (XRNE): the ﬂuorescent Fe I Kα (6.4 keV line) with Thomson scattering continuum.
The spectra of XRNs in the Sgr B and Sgr C complexes were extensively studied with Suzaku(Koyama et al. 2007b; Nobukawa et al. 2008; Nakajima et al. 2009). Following these results,
we use a phenomenological model: a power-law plus 2 Gaussian lines for the Fe I Kα andKβ as,XRNE = A × (E/keV)−Γ + Gaussian1 + Gaussian2[photons cm−2 s−1 arcmin−2 ](4.1)

where the Gaussian center energy and relative intensity ratio of the Fe I Kα, Kβ lines are
ﬁxed to 6.4 keV and 7.05 keV, and 1 : 0.125, according to the theoretical values (Kaastra &Mewe 1993).
Component (4) is low energy (≤ 2 keV) emission extending to nearly 0.5 keV, indicatingsmall or no absorption, and hence the foreground emission, which is unrelated to the GC
region. The origin is unknown but would be either a local Galactic plasma (Ebisawa et
al. 2008) or unresolved faint dM stars (Masui et al. 2009). Apart from the real origin, thespectra of component (4) contains faint line-like structure (see ﬁgure 4.3), we therefore used
an APEC model (APEC3). Since the contribution of component (4) above ∼ 2 keV is small,the ﬁtting results of the spectra above ∼ 2 keV depend very weakly on the model assumptionof component (4).
Component (5) is the cosmic X-ray background (CXB). We applied the CXB modelobtained from the Suzaku data at the north ecliptic pole,2

The solar abundance here is referred to Anders & Grevesse (1989).
Astrophysical Plasma Emission Code: A model of an emission spectrum from collisionally-ionized diﬀuse
gas (Smith et al., 2001).
3- 51 -

CHAPTER 4. STUDY OF SAGITTARIUS B
CXB = 7.4 × 10−7 × (E/keV)−1.486[photons cm−2 s−1 arcmin−2 ](4.2)

As a matter of fact, the surface brightness of the CXB (see also Kushino et al. 2002) in
the Sgr B region is far lower than those of the other components by more than two orders ofmagnitudes, and hence the contribution of the CXB is almost negligible.

4.2.3

Model Constructions and Fitting

For the full-band spectra of the ﬁve components, we tried model ﬁtting with many freeparameters. However, due to large errors, no signiﬁcant constraint on the best-ﬁt parameters was obtained for any further scientiﬁc study and discussion. We therefore set several
constraints on the model based on reasonable assumptions.
We assume that the spectral shape of the GCPE, which comprises the two plasmas(APEC1 and APEC2), is the same in all of the Sgr B regions. We thus combine components (1) and (2),
GCPE = APEC1 + APEC2[photons cm−2 s−1 arcmin−2 ](4.3)

The temperature and abundances of APEC1 are ﬁxed to 6.5 keV and 1 solar, respectively.
We deﬁne a parameter α, the normalization ratio of APEC1 and APEC2, assuming thecommon value in all regions near Sgr B (see ﬁgure 4.6). Thus, free parameters for the
GCPE spectra are the temperature and the normalization factor of APEC2, and α with thesame value among the relevant regions.
The photon indices (Γ) of the continuum of XRNs in Sgr B are scattered. This may be
due to the limited energy band ﬁttings in former studies. On the other hand, the equivalent
width values of the 6.4 keV line are nearly the same (Koyama et al. (2007b), Nobukawa etal. (2008), Nakajima et al. (2009)). Since the emission mechanisms would be the same for
all Sgr B XRNs, we assume that the photon index (Γ) and the equivalent width (EW ) of
the 6.4 keV line are the same in all regions of Sgr B.
Although each of components (1)–(5) should have a diﬀerent absorption, we truncatedthem into three absorptions: Abs1, Abs2, and Abs3. Abs i (i = 1, 2, 3) is the integrated
value along the line of X-ray path into our sight, which is given as exp(−NH × σ(E) ); NH andσ(E) are respectively the hydrogen column density and absorption cross section with solar
abundance. A schematic view of these three absorptions is given in ﬁgure 4.4.
Abs1 is the largest absorption (NH ≥ 1023 H cm−2 ), which takes place in/through theMCs. Abs2 is the interstellar absorption to/beyond the GC of NH ∼ 6×1022 H cm−2 (Sakano- 52 -

4.2. DATA ANALYSIS AND RESULTS
et al. 2002). Abs3 is the smallest absorption applied to the foreground emission (NH <1022 H cm−2 ).

CXB(N H ~ 6)

B
GCPE
l

1 deg

C

Sgr A*

Sgr B (XRNE)(NH >> 6)A
Foreground Emission(N H ~ 6)

Abs3(N H <1)

ObserverFig. 4.4.—: Schematic view of the MCs distribution along the line of sight for Model C. The unitof absorption (NH ) in the parentheses is 1022 H cm−2 .

In all 6 regions, we needed additional common components (4) and (5), which are notdue to the emission from the GC region. We combined these components and deﬁne themas the non GC emission (NGCE):
NGCE = Abs1 × Abs2 × Abs2 × CXB + Abs3 × APEC3[photons cm−2 s−1 arcmin−2 ](4.4)

The term Abs2×Abs2×CXB indicates that the CXB suﬀers from interstellar absorption(Abs2) two times: coming from the front and back sides of the GCPE. Here we assume thatthe absorption due to the back side on the GC is same as that on the front side of the GC.
Referring the schematic view of the geometry of the MCs (XRNE) and the GCPE, we
tried two extreme cases (see ﬁgure 4.4). (A): a MC is located at the near-side edge of the
GCPE, then the XRNE and GCPE would have nearly the same absorption of Abs1×Abs2.
(B): a MC is located at the far-side edge of the GCPE; thus, the XRNE and the GCPE
- 53 -

CHAPTER 4. STUDY OF SAGITTARIUS B
would have independent absorptions, given as Abs1×Abs2 and Abs2, respectively. Model A
and Model B, therefore, are given as,(A) :Abs1 × Abs2 × (XRNE + GCPE) + NGCE

(B) :Abs1 × Abs2 × XRNE + Abs2 × GCPE + NGCE
[photons cm−2 s−1 arcmin−2 ](4.5)

In the ﬁttings, the free parameters are Abs1, Abs2, and the normalization factors of the
6.4 keV-line ﬂux, the GCPE, and the APEC3. While the mixing ratio α of APEC1 andAPEC2 (in the GCPE), Abs3, the temperature and the abundance of APEC3, the equivalent
width of the 6.4 keV line with respect to the power-law continuum (EW6.4 keV ), and thephoton index (Γ) for the power-law spectrum of XRNE are also free, but are assumed to be
the same (commonly linked) in all relevant regions. We list the ﬁtting parameters according
to their properties in table 4.2.
Table 4.2—: Summary of free and ﬁxed parameters in the ﬁttings for Model A, B, and C.

Free parameters for each region
NH (Abs1), NH (Abs2)
normalization (Fe I Kα 6.4 keV line)
normalization (APEC2), position R†
normalization (APEC3)
Free parameters commonly linked in all regions
Absorption: NH (Abs3)
XRNE∗ :
photon index Γ, equivalent width EW6.4 keV
GCPE∗ :kT (APEC2), α (normalization ratio of APEC1/APEC2)
NGCE∗ :kT (APEC3), abundances Z (APEC3)
Fixed parameters
XRNE∗ :Fe I line energy Kα = 6.4 keV, Kβ = 7.05 keV, ﬂux ratio (Kβ/Kα) = 0.125
GCPE∗ :kT (APEC1) = 6.5 keV, Z (APEC1 & APEC2) = 1 solar
NGCE∗ :CXB = equation (4.2)∗
See equation (4.1), equation (4.3), equation (4.4) for the deﬁnitions of XRNE, GCPE, andNGCE, respectively.
†
For Model C only, see equation (4.6) and the text.
Absorption:
XRNE∗ :
GCPE∗ :
NGCE∗ :

We performed a simultaneous ﬁtting for the 2 × 6 spectra (the FI and BI spectra of
6 regions) with Model A and Model B. The best-ﬁt parameters and χ2 /d.o.f are listed intable 4.3. As a results, both models were rejected by the χ2 /d.o.f test. Model A gave alarge χ2 /d.o.f of 1.87, and shows large residuals in the medium range of 1–3 keV and thehard range of 5–7 keV (see ﬁgure 4.3). The best-ﬁt values of Abs1 × Abs2 (with Model A)
is ≃ 10 × 1022 H cm−2 (see table 4.3), which are close to the typical value toward the GC
region. These values may be reasonable for the GCPE but could not be applied to the XRNE,- 54 -

4.2. DATA ANALYSIS AND RESULTS
Table 4.3—: Results of best-ﬁt parameters∗ with Model A and B.

∗

RegionSgr B2
M0.74-0.09
Sgr B1
M0.74-subRegion1
Region2
Region-common parameterskT (keV) (APEC2)
norm
α ( APEC1)
APEC2 norm
Photon index Γ
EW6.4 keV (keV)
Abs3 (1022 H cm−2 )kT (keV) (APEC3)
Z (solar) (APEC3)
χ2 /d.o.f §MODEL A†NH
(Abs1×Abs2)
6.4 keV line‡+0.9
10.9−0.5
3.0+0.1−0.1+0.1
10.7+0.5
2.5
−0.4
−0.1
1.53+0.04
8.4+0.2
−0.2
−0.08
+0.1
7.4−0.4
1.25+0.02−0.02
6.1+0.1
0.24+0.01
−0.1
−0.01
+0.01
7.7+0.1
0.43
−0.3
−0.03

MODEL B†NH
(Abs2)
5.8+0.4−0.1
5.2+1.0−0.8
6.3+0.2−0.1
5.0+0.1−0.1
4.6+0.1−0.1
5.8+0.1−0.2†
NH(Abs1)
54.8+2.8−2.4
40.4+2.0−1.7
26.3+2.0−1.8
21.9+0.6−1.0
18.6+1.9−1.2
31.8+2.2−1.9

0.70 (0.68–0.72)
0.47 (0.46–0.48)
0.52 (0.47–0.59)
1.78 (1.75–1.80)
0.21 (0.21–0.22)
1.02 (0.99–1.05)
0.004 (0.002–0.006)
4749/2537 = 1.872

6.4 keV line‡
6.6+0.2−0.2
4.3+0.1−0.6
2.2+0.1−0.1
1.63+0.03−0.02
0.33+0.01
−0.01
+0.02
0.76−0.04

0.69 (0.68–0.74)
0.59 (0.57–0.60)
2.16 (2.10–2.23)
1.31 (1.28–1.34)
0.20 (0.19–0.21)
0.93 (0.92–0.94)
0.008 (0.007–0.011)
4479/2537 = 1.765

The uncertainties are at 90% conﬁdence level.
The value of NH in the unit of 1022 H cm−2 .
‡Flux in the unit of 10−6 photons cm−2 s−1 arcmin−2 .
§The results obtained by the simultaneous ﬁtting of the 12 spectra (the FI and BI spectraof 6 regions).
†

because the absorption of each XRNE was found to be > 1023 H cm−2 by independent analysis
of XRN (Koyama et al. (2007b); Nobukawa et al. (2008); Nakajima et al. (2009)).
In Model B, the apparent residuals in the medium and hard range found in Model Adisappeared partially, and hence χ2 /d.o.f was improved to 1.77. We still see unnaturalresiduals between 2–4 keV. Although the χ2 /d.o.f of 1.77 is large enough to reject Model B
at more than the 99.9% conﬁdence level, this value is slightly smaller than that of Model A.
At the end of this section, we discuss why we reject Model B from a diﬀerent point of view.
We then constructed Model C, adding a new free parameter R (for each region), in
which a MC/XRN is inside the GCPE with a relative position of R (0–1); the smaller R
indicates the nearer (front) side.
(C) : Abs1 × Abs2 × (XRNE + (1 − R) × GCPE)+Abs2×(R × GCPE) + NGCE
[photons cm−2 s−1 arcmin−2 ](4.6)

We performed a simultaneous ﬁtting for Model C by the same procedure as for Model A
and B. Although Model C has six new free parameters (R for the six regions) more thanModel A and Model B, the degree of freedom (d.o.f) of Model C (2531) is almost equal tothat of Model A and and B (2537). On the other hand, the χ2 /d.o.f of Model C is largely- 55 -

CHAPTER 4. STUDY OF SAGITTARIUS B
improved to 1.34 from those of Model A (1.87) and Model B (1.77).
One may argue that the spectra are contaminated by unresolved point sources in the
Sgr B region. In fact, Koyama et al. (2009) estimated that the point source ﬂuxes provide asigniﬁcant fraction of the GCDX (GCPE + XRNE) emission, up to about 1/6 near Sgr A*,
which may partially contribute to the 6.7 keV-line and the 6.4 keV-line ﬂuxes in the GCPE
and XRNE, respectively. According to the Chandra cataloged sources reported by Muno
et al. (2006), the number of point sources per solid angle near Sgr B2 is ∼ 10% comparedto that near Sgr A*, while Nobukawa et al. (2008) shows that the GCDX ﬂux near Sgr B2
is also ∼ 10% of that near Sgr A*. We can infer that the contribution of point sources to
GCDX would be roughly the same in the Sgr B region. The point source contribution may
cause possible systematic errors in the present spectral analysis. The most serious errors,
therefore, should be the region-to-region ﬂuctuation of the fractions of the point source
ﬂuxes, which may be rather small eﬀects in a relative comparison of the best-ﬁt parameters.
At this moment, we have no information on the exact ﬂux and spectrum of the integrated
point sources in the Sgr B region. We therefore ignore the point source contribution.
All together, taking account of possible systematic errors due to the point source contribution and those due to the simpliﬁed model, we regard Model C as a good approximation of
the overall spectra in the Sgr B region, although the χ2 /d.o.f of 1.34 is statistically marginal
to be acceptable. The best-ﬁt common parameters for all regions are as follows (hereafter,
the numbers in the parenthesis are in 90% error range). The temperature of APEC2 is 0.87
(0.81–0.91) keV and the mixing ratio (α) in the GCPE is determined to be 0.27 (0.26–0.28).
The photon index (Γ) and the equivalent width (EW ) of the 6.4 keV line for the XRNE are
1.72 (1.64–1.80) and 1.59 (1.54–1.63) keV, respectively. The absorption (Abs3), temperature(kT ) and the abundance of the foreground emission (APEC3) are 0.17 (0.16–0.18)×1022 H
cm−2 , 0.85 (0.83–0.87) keV, and 0.011 (0.009–0.014) of solar, respectively. The other best-ﬁt
parameters of Model C for each region are listed in table 4.4.
The 6.4 keV line is likely to be ﬂuorescence in the MC generated by an external X-ray
source. Then the ﬂux should be proportional to the product of the external X-ray ﬂux andthe number of iron atoms along the line-of-sight X-ray path in the MC, in the opticallythin case at the energy near 6.4–7.14 keV (NH ≤ 1024 H cm−2 ). The number of iron atoms
is approximately proportional to the column density NH of the MC. Therefore, in the casethat the external X-ray ﬂuxes are the same among the MCs, Abs1 should be approximately
proportional to the 6.4 keV-line ﬂux.
In ﬁgure 4.5, we plot the best-ﬁt Model C values of Abs1 and Abs2, as a function of the
6.4 keV line ﬂux. For a comparison, we also plot the results of Model B. A notable fact tosupport Model C is that Abs1 is in good proportionality to the 6.4 keV ﬂux. This feature is in
4

E = 7.1 keV is the K absorption-edge of neutral iron.

- 56 -

4.2. DATA ANALYSIS AND RESULTS
Table 4.4—: Detailed results of best-ﬁt parameters with Model C.

Absorption N H [1022-2

H cm ]

RegionNH† (Abs1) NH† (Abs2) 6.4 keV line‡ APEC3 norm§ GCPE norm∥ Position R#
Sgr B236.5+2.4
8.2+0.3
5.1+0.3
0.09+0.01
3.40+0.36
0.44+0.06
−2.2
−0.5
−0.3
−0.01
−0.48
−0.06
+2.0
+0.4
+0.2
+0.01
+0.29
M0.74-0.09
22.9−1.8
6.8−0.4
3.3−0.2
0.11−0.01
2.67−0.28
0.31+0.04−0.04
Sgr B1
12.3+1.0
7.0+0.2
1.8+0.1
0.11+0.01
3.46+0.25
0.41+0.06
−1.0
−0.3
−0.1
−0.01
−0.38
−0.06
+0.1
+0.03
+0.003+0.10+0.04
M0.74-sub
11.5+0.5
5.3
1.42
0.090
1.66
0.26
−0.5
−0.2
−0.03
−0.004
−0.09
−0.04
+0.05
+0.006
+0.01
+0.2
+0.2
0.15+0.01
1.40−0.04
0.105−0.004
0.26−0.01
3.6−0.3
Region17.2−0.2−0.01
0.34+0.04
2.88+0.11
0.130+0.004
0.47+0.03
6.3+0.1
Region2
12.0+0.6
−0.04
−0.18
−0.036
−0.03
−0.1
−0.8
χ2 /d.o.f ∗∗3382/2531 = 1.336∗
The uncertainties are at 90% conﬁdence level.
†
The value of NH in the unit of 1022 H cm−2 .
‡Flux in the unit of 10−6 photons cm−2 s−1 arcmin−2 .
§The normalization factor for the foreground emission (c.f., APEC3 in equation (4.4)) of the APEC
model re-normalized with the region size A [arcmin2 ]; it is expressed as 7.07 × 10−11 /(4πD2 A) EM[cm−5 arcmin−2 ], where D and EM are the distance to the source [cm], and the emission measure [cm−3 ],respectively.
∥
In same expression as §, but the normalization factor of APEC2 in the GCPE (c.f., equation (4.3) andtext).
#
R is the fraction of the GCPE suﬀered by Abs2, which indicates the positions of MCs (see equation (4.6)
and ﬁgure 4.7).
∗∗
The results obtained by the simultaneous ﬁtting of the 12 spectra (the FI and BI spectra of 6 regions).

60
55

Abs1 :
Abs2 :
Model C

50
45
40
35
30
25
20
15
10
5
0

Abs1 :
Abs2 :
Model B

GC typical level ~ 6

0

1
2
3
4
5
6
7
-2
-6
-2 -1
6.4 keV-line flux [10 photons cm s arcmin ]Fig. 4.5.—: Correlations between the 6.4 keV-line ﬂux and absorptions (made from data listed intable 4.3 and table 4.4). Results of Model C and Model B are plotted by the black and gray marks;
ﬁlled squares and open circles represent Abs1 and Abs2, respectively.

- 57 -

CHAPTER 4. STUDY OF SAGITTARIUS B
contrast to the results of Model B, where the proportionality is not very good in the low-ﬂux
range. Most importantly, Model B requires an extra oﬀ-set of about 2×1023 H cm−2 , which is
un-realistic. This leads us to exclude Model B in addition to the χ2 /d.o.f test. We estimatedNH values of the Sgr B MCs from the data of the CO line (Sofue 1995; Dahmen et al. 1998)
and the CS line (Tsuboi et al. 1999). They are in the range of NH = 1023 –1024 H cm−2 , ingood agreement with the Model C prediction.
As shown in ﬁgure 4.5 for Model C, we found that Abs2 is almost constant at NH ∼6 × 1022 H cm−2 . This is consistent with the view that Abs2 is interstellar absorption (seeﬁgure 4.4) toward the GC of about NH ≃ 6 × 1022 H cm−2 (Rieke et al. 1989; Sakano et al.
2002).

4.3

Discussion for Sgr B

4.3.1

Justiﬁcations of Model Constraints and Assumptions

In the previous section, we introduced a complicated model to described the spectra of
the GCDX. In order to obtain meaningful values (no large errors) of physical parameters, we
further added some constraints and assumptions. If, for example, we let the normalization
ratio α (APEC1/APEC2) to be free for all regions, we can not give a signiﬁcant constraint on
the absorptions of GCPE and/or the temperature of APEC1, because the two plasma emission become nearly the same in the band near 2 keV. We summarize the adopted constraints
and assumptions below for further discussions:• The 6.5 keV plasma and ∼ 1 keV plasma, which comprise the GCPE, have nearly thesame angular distribution near the Sgr B region. (see ﬁgure 4.6)
• The abundances of the GCPE are nearly 1 solar, the same value as those of nearby
new SNR candidates.
• The XRNE has a common photon index (Γ) and EW6.4 keV near the Sgr B region.
• The point sources contribution to the GCDX is small and hence could be ignored in
the Sgr B region.
The most fundamental constraint is the ﬁrst item; the factor α is constant from region
to region. This implies that the angular distributions of the very hot (6.5 keV) and hot(∼1 keV) plasmas are nearly the same near Sgr B. In order to judge this constraint, weshow the ﬂux distribution of the 6.7 keV and 2.45 keV lines along the galactic longitude of- 58 -

4.3. DISCUSSION FOR SGR B
l = 1.◦ 2–0.◦ 4 at b = −0.◦ 15 ± 0.◦ 05 (ﬁgure 4.6). The 6.7 keV and 2.45 keV lines are strong
characteristic lines that well represent, respectively, the distributions of the 6.5 keV plasma
and the ∼ 1 keV plasma with the least systematic errors.
1

Flux (arbitrary)--- 6.7 keV Line
2.45 keV Line

0.1

0.01
1.2 1.1

1.0 0.9 0.8 0.7 0.6
Galactic longitude [degree]

0.5

0.4

Fig. 4.6.—: Flux distribution of the 6.7 keV line (6.5 keV plasma; dashed line) and the 2.45 keVline (∼ 1 keV plasma; solid line) along the galactic longitude from l = 1.◦ 2–0.◦ 4 at b = −0.◦ 15±0.◦ 05.
Two parallel dotted lines guide eyes to see the same distribution of the 6.5 keV plasma and the ∼ 1
keV plasma.

As is found in ﬁgure 4.6, the 2.45 keV line and the 6.7 keV line have almost the same
slope (angular distribution). The small drop of the 2.45 keV-line ﬂux near at l = 0.◦ 85–0.◦ 6
would be due to absorption by the M 0.74−0.09 MC and the drop at l = 0.◦ 6–0.◦ 4 is due tothe absorption by the Sgr B1 MC. This ﬁgure, therefore, indicates that the emission of the∼1 keV plasma (APEC2) and the 6.5 keV plasma (APEC1) have a common normalization
ratio (α) for the GCPE near the Sgr B regions. The justiﬁcation of the other constraints andassumptions listed in the last three items are separately discussed in the previous sections.
We thus conclude that Model C with the above listed constrains and assumptions provides
a practical approximation for the overall spectra in the Sgr B region, although the χ2 /d.o.fof 1.34 is statistically marginal to be acceptable.

4.3.2

Line-of-Sight Distribution of Sgr B MCs in the GCPE

From the results of section 4.2.3, we conclude that the MCs of Sgr B should be locatedin the GCPE (Model C), neither in the near-side edge (Model A) nor in the far-side edge
(Model B) of the GCPE. For simplicity, we assume that the volume emissivity of the GCPE
- 59 -

CHAPTER 4. STUDY OF SAGITTARIUS B
is approximately a uniform sphere with a uniform plasma density on the Galactic plane.
Koyama et al. (1989) and Yamauchi et al. (1990) reported that the 6.7 keV line distributedwith a Gaussian shape of 1.8 deg (FWHM) (in parallel to the Galactic plane). We therefore
take the boundary (radius) at |l| = 1.25 deg, within which 90% of the GCPE is included.
The center of the GCPE is approximately Sgr A* (at l= -0.◦ 056; Yusef-Zadeh et al. 1999).
Then the parameter R gives the line-of-sight positions of the MCs within the GCPE5 , whereR = 0 and 0.5 indicate the anterior border line and the center line (Sgr A* is on this line)of the GCPE, respectively. We plot the best-ﬁt position parameter (R) for the 6 regions(indicated in ﬁgure 4.1a) in the schematic view of ﬁgure 4.7. As demonstrated in ﬁgure 4.7,
the MCs in the Sgr B region are generally on the near-side with respect to Sgr A*.

[degree]1.25 1.0 0.75 0.5 0.25 0.0

Sgr A*l
Sgr B1
Sgr B2
Region2
M0.74-0.09
Region1

M0.74-sub

GCPE

R=

ObserverFig. 4.7.—: Relative positions of each MC (ﬁlled circles) along the line of sight, where thehorizontal lines of R = 0.0 and 0.5 imply the anterior border line and the GC line (Sgr A* is on
this line), respectively. The circles ﬁlled with darker colors indicate MCs of the larger NH values.
∫∫
The emission measure EM can be written as ne nH dV = ne nH S dl ∝ ne nH SL, where L is the lengthof the GCPE in the line of the sight, S and ne nH are the projected area and the product of densities of
front
all
electrons and hydrogens, respectively. Then the parameter R is the fraction of EMGCPE/EMGCPE, whichfront
allis equal to LGCPE /LGCPE : the relative line-of-sight position of MCs in the GCPE.
5- 60 -

4.3. DISCUSSION FOR SGR B
It may be suspicious that the best-ﬁt values of R in Model C indicate that the MCs
are all closer to the near-side edge of the GCPE rather than to the far-side edge, while
χ2 /d.o.f. = 1.872 of Model A (MCs at the near-side edge) is rather worse than χ2 /d.o.f.
= 1.765 of Model B (MCs at far-side edge). As shown in equation (4.5), Model B have theextra-part Abs2×GCPE which Model A does not have. This gives a slightly better χ2 /d.o.f.
of Model B (1.765) than Model A (1.872). On the other hand, the improvement of Model C(χ2 /d.o.f.= 1.336) from Models A (and B) is far larger than that between Model A and B.
In this section, we have only shown the line-of-sight positions for MCs from the valuesof R under an assumption of a uniform density for the GCPE. We can furthermore determine the actual positions quantitatively using the R parameter together with the threedimensional information of the GCPE distribution measured by Uchiyama et al. (2012).
This analysis will be performed in chapter 7.

4.3.3

XRN Features in Sgr B

According to former studies (Koyama et al. 2007b; Nobukawa et al. 2008), the equivalent
width of the 6.4 keV line and the continuum photon index (EW , Γ) of the Sgr B XRNs havebeen reported to be (1.1 keV, 3.2) for Sgr B2, (1.55 keV, 1.4) for M 0.74-0.09, and (1.4 keV,
1.8) for Sgr B1. The Γ values are signiﬁcantly scattered among the XRNs, but the variationof EW is not very large. The GCPE ﬂux is position-dependent, which means that a proper
background-selection is diﬃcult. The large scatter of Γ values in the former studies may
be due to the nearby GCPE-subtraction and would also be due to the band limited (e.g.
5–8 keV) spectral ﬁtting. In this work (Model C), the EW and Γ values are estimated byan independent method, a simultaneous ﬁtting of a wide energy band (0.5–10.0 keV), in
which nearby GCPE was not subtracted but directly included as the GCPE model. We
also note that these uncertainties do not largely aﬀect the EW values. In fact, We foundthat the averaged EW ≃ 1.6 keV, which is consistent with the previous studies of the SgrB XRNs. Koyama et al. (2009) studied the GCDX near the Sgr A region (|l| ≤ 0.◦ 2), andfound that the Γ of the continuum ﬂux is 1.9; this reduced to 1.4 if we exclude the point
sources contribution. Nakajima et al. (2009) studied the 6.4 keV clumps in the Sgr C region
and found Γ to be in range of 1.6–1.9. We determined the common Γ in Sgr B to be ≃ 1.7;
this value agrees with the results in Sgr A, B, and C. Thus, we infer that the XRNE may
have a uniﬁed photon index of 1.4–1.9 in the overall GC region.
As discussed by Koyama et al. (2008) and Nobukawa et al. (2008), a possible external
X-ray source is Sgr A*, therefore the X-ray ﬂux at the Sgr B MCs should be proportional to
L0 D−2 , where L0 and D are, respectively, the luminosity of Sgr A* and the distance between
Sgr A* and MCs. Thus the ﬂux of the ﬂuorescent 6.4 keV line should be modiﬁed by this- 61 -

CHAPTER 4. STUDY OF SAGITTARIUS B

40
distance corrected :equation 7 : ---22-2

Abs1 N H [10 H cm ]

35
30
25
20
15
10
5
0

0

1
2
3
4
5
7
8
9 10
6
-6
-2 -1
-2
6.4 keV-line flux [10 photons cm s arcmin ]Fig. 4.8.—: The best-ﬁt absorption (Abs1) vs the distance-corrected 6.4 keV-line ﬂux fromresults of Model C in table 4.4 (open triangles), where the correction factors are estimated fromthe distances between MCs and Sgr A* with ﬁgure 4.7, and are normalized to that of Sgr B1.

distance eﬀect. In ﬁgure 4.8, we plotted distance-corrected 6.4-keV ﬂuxes as a function of
Abs1, where the distances are estimated using ﬁgure 4.7. We then ﬁtted the correlation
between the best-ﬁt absorption (Abs1) and the 6.4 keV line ﬂux (distance-corrected), andobtained,
NH (Abs1) = 3.99 (±0.56) × I6.4 keV [1022 H cm−2 ],I6.4 keV : [10−6 photons s−1 cm−2 arcmin−2 ](4.7)

This relation is generally in agreement with the ﬂuorescent process due to irradiation by
Sgr A* with the past luminosity of 2–3×1039 ergs s−1 . If L0 was time constant, the relation
between the distance-corrected 6.4 keV-line ﬂux vs. Abs1 should be in good proportionality.
However, we see some data scatters from the proportionality relation (ﬁgure 4.8). It can be
explained by assuming a time variable ﬂux of the irradiating source Sgr A* (Koyama et al.
2008; Inui et al. 2009). We will show and discuss its details in chapter 7.

4.3.4

The Correlation of CS Line and 6.4 keV LineAs shown in ﬁgure 4.2, the 2.5–3.5 keV band ﬂux shows an intensity drop at the bumpof the 6.4 keV line and the CS line. The anti-correlation to the 2.5–3.5 keV band ﬂux isclearer in the CS line than that in the 6.4 keV line. This may be because that the CS line
- 62 -

4.3. DISCUSSION FOR SGR B
traces all (GC + foreground) MCs in the line of sight from which the GCPE may suﬀer
absorption, while the 6.4 keV line only traces MCs near the irradiating source in the GC
region. The better anti-correlation of the CS line against the 2.5–3.5 keV band ﬂux indicatesthat a large amount of MCs, identiﬁed by the velocity-integrated CS lines but not bright in
the 6.4 keV line, is located in the near side of the GCPE. The angular size of the 6.4 keVbump is slightly narrower than that of the CS line, as seen at the bump edges near l ∼ 0.◦ 85
and l ∼0.◦ 63. The narrower angular size of the 6.4 keV bump than that of the CS line may
also be explained if Sgr A* is the external X-ray source to produce the 6.4 keV clumps. The
surface brightness of ﬂuorescent X-rays is faint at the near and far sides of the cloud-limb
because of small NH values, but becomes brighter when the external X-rays go into the large
NH (≥ 1023 H cm−2 ) region in the MCs. Another possible reason of the faint 6.4 keV-line ﬂux
at the far-side cloud limb from Sgr A* is due to the time variability of the Sgr A* X-rays;
an X-ray front of possible X-ray outbursts about 300 years ago (Koyama et al. 2008; Inui etal. 2009) does not yet reach at the far-side limb of l ∼ 0.◦ 85.
Most of the CS MCs near Sgr B show positive velocities of 0–100 km s−1 . From spatialcomparison between individual 6.4 keV clumps and the CS (J=1→0) maps in diﬀerent velocity ranges (Tsuboi et al. 1999), we can roughly estimate that Sgr B2, Sgr B1, M 0.74−0.09,
and M 0.74-sub lie in the velocities of +50 (40–80) km s−1 , +30 (20–40) km s−1 , +50 (40–60) km s−1 , and +50 (30–80) km s−1 , respectively.

- 63 -

CHAPTER 4. STUDY OF SAGITTARIUS B- 64 -

Chapter 5
Study of Sagittarius C
In this chapter, we study the X-ray reﬂection nebulae (XRNe) in the Sgr C region andtestify the spectral model proposed in Sgr B (chapter 4). We investigate the short-term time
variability of the 6.4 keV line. We also establish the spectral composition for the Galactic
center plasma emission.

Contents
5.1

Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

65

5.2

Data Analysis and Results . . . . . . . . . . . . . . . . . . . . . .

66

5.2.1

Images . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

66

5.2.2

Spectral Model . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

68

5.2.3

Spectral Fitting . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

69

Discussion for Sgr C . . . . . . . . . . . . . . . . . . . . . . . . . .

73

5.3.1

Association with the Radio MCs . . . . . . . . . . . . . . . . . . .

73

5.3.2

Two-temperature Structure of the GC Plasma and the XRNParameters . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

74

5.3

5.3.3

5.1

Time Variability of the 6.4 keV Line and the Flare Property of Sgr A* 75

Observation

The Sgr C complex is composed of several giant molecular clouds located at l ∼ 0.5◦ , on
the opposite side of Sgr B with respect to Sgr A*. We have made two deep pointing observations on the Sgr C region separated by ∼5 years between February 2006 and September
65CHAPTER 5. STUDY OF SAGITTARIUS C
2010. These observations are conducted using the Suzaku/XIS. We operated the XIS in thenormal clocking mode with a read-out cycle of 8 s. The details of Suzaku and XISs are given
in chapter 2. The observation log of the Sgr C data for spectral analysis is shown in table 5.1.
Information on archive data of other relevant GC regions used for image construction canbe found in Uchiyama et al. (2011).
Table 5.1—: Suzaku Observations at the Sgr C Region.

Target Obs. ID
Obs. point (FOV center)nameα (J2000) δ (J2000)h
Sgr C 500018010 17 44m 37.30s−29◦ 28′ 10.′′ 2
Sgr C 505031010 17h 44m 58.01s−29◦ 22′ 51.′′ 2

5.2

Obs. date(yy-mm-dd)
2006-02-20
2010-09-25

Eﬀective
exposure
106.9 ks
82.6 ks

Data Analysis and Results

We performed the data reduction and analysis using HEADAS software version 6.12.
The calibration database used to calculate the response of XIS and the eﬀective area of XRT
was the version released on 2011-02-10. The screening for the XIS data were performed in
the same way as those of Sgr B (described in section 4.1). The FI spectra were merged in
the spectral ﬁtting, and the NXB was subtracted from the raw data. The XRT vignetting
eﬀects and the exposure-time diﬀerences were corrected.

5.2.1

Images

We constructed the XIS images of the Sgr C region with the 6.4 keV-line (Fe I), the6.7 keV-line (Fe XXV), and the 2–3 keV bands (see ﬁgure 5.1). The images were generated
by adding data from multiple observations together (c.f., table 5.1 and Uchiyama et al. 2011)
to cover a wide region near Sgr C. In addition, the FI and BI data were combined to increase
the statistics. For the 6.4 keV and 6.7 keV line images, the ﬂuxes underlying the continuum
were estimated from the 5–6 keV band and subtracted from the 6.3–6.5 and 6.6–6.8 keVbands, respectively. Details of the subtraction procedures are explained in Nobukawa et al.
(2010).
Figure 5.1a displays many bright clumps in the 6.4 keV-line band near Sgr C. These arespatially associated with giant MCs (see ﬁgure 5.1c) as observed in the radio CS (J=1→0)
line by Tsuboi et al. (1999). We designate the brightest three 6.4 keV clumps: C1, C2, and- 66 -

5.2. DATA ANALYSIS AND RESULTS

Fig. 5.1.—: Added XIS (FI+BI) images near the Sgr C region. The images are binned with
12×12 pixels and smoothened with a Gaussian kernel of 5′′ . (a) 6.4 keV line (Fe I) image obtainedby subtracting the continuum component (see text) from the 6.3–6.5 keV band ﬂux. The XIS ﬁeldof views (FOV) of the 2009 and 2006 observations are represented by the large boxes. The regions
for spectral analyses (C1–3, ref 1–2) are indicated by solid circles. (b) Same as panel (a), but for the
6.7 keV line (Fe XXV) image. (c) 6.4 keV line image (gray) with the superposed radio CS (J=1→0)
contours (Tsuboi et al. (1999)) in the velocity range of −100 ± 20 km s−1 (red), −60 ± 20 km s−1(green), and +50 ± 30 km s−1 (blue). The contour level begins at the intensity of 5 K km s−1 andincreases every 5 K km s−1 . (d) Same as panel (a), but for the 2–3 keV band image (color). The
region of G359.41−0.12 in Tsuru et al. (2009) is represented by the dashed circle.

C3 (solid circles; ﬁgure 5.1a). The clumps C1 and C2 were previously identiﬁed as XRNe by
Nakajima et al. (2009), while C3 is a new XRN candidate revealed in the 2010 observation.
Figure 5.1b (the 6.7 keV line image) traces the high-temperature plasma (HP) of the
GCPE (section 3.2), and shows a more uniform distribution than the 6.4 keV emission shownin ﬁgure 5.1a. Figure 5.1d shows the 2–3 keV band image including the 2.45 keV line (S XV)
and displays the distribution of the low-temperature plasma (LP) and individual thermal
diﬀuse sources, mainly supernova remnants candidates, with temperatures of kT ∼1 keV.
The faint diﬀuse sources near the Sgr C clumps are the supernova remnants G359.41−0.12- 67 -

CHAPTER 5. STUDY OF SAGITTARIUS C
and its outﬂow ”Chimney” (Tsuru et al. 2009). To investigate the nature of the GCPE (HPand LP), we selected two regions with relatively weak 6.4 keV emission for spectral analyses,
and designate them as ref1 and ref2 (solid circles; ﬁgure 5.1a, b, and d).

5.2.2

Spectral Model

The X-ray spectra of the GC region have been extensively studied in chapter 4 andUchiyama et al. (2012). The spectra can be reproduced by the superposition of four components: the GCPE, the XRN emissions (XRNE), the foreground emission (FE), and thecosmic X-ray background (CXB). These components have the following properties. The
GCPE has two-temperature plasmas with kT ∼1 keV and kT ∼7 keV for LP and HP, respectively. The XRNE is expressed as neutral ion lines (Kα and Kβ) associated with apower-law component. Since the XRNE is formed by Thomson scattering of a continuum
component and ﬂuorescence of iron atoms in the MCs, it shows a large absorption (Abs1)of NH ∼ 1023 H cm−2 (e.g., Murakami et al. 2001), which is approximately equal to the NH
through the MC. The GCPE including the HP and LP is extended in the GC region, and the
MC is inside this region. Thus, a fraction (R) of the X-ray from the GCPE is not absorbedby the MC while the other fraction (1−R) is absorbed by the MC (Abs1). Both GCPE andXRNE exhibit an common interstellar absorption (Abs2) of NH ∼ 6 ×1022 H cm−2 between
the GC and the sun (see chapter 4; Sakano et al. 2002). The FE is approximated to athermal plasma of kT ∼1 keV with small absorption. The CXB (e.g., Kushino et al. 2002)
has the absorption of Abs2 by twice (interstellar absorption in front of and behind the GC),in addition to Abs1 (MC). Thus, the spectral model can be expressed by equation (5.1). A
schematic picture of this model is shown in ﬁgure 5.2.
f (E ) = [Abs1 × Abs2 × (1 − R) + Abs2 × R] × (HP + LP)+[Abs1 × Abs2] × XRNE+[Abs1 × Abs2 × Abs2] × CXB + FE

(5.1)

In chapter 4, we successfully applied this model to the X-ray spectra of the 6.4 keVclumps near Sgr B. In addition, Muno et al. (2004) and Nobukawa et al. (2010) reproduced
the X-ray spectra of Sgr A regions with similar models. Therefore, we apply the same modelto the Sgr C region. In ﬁgure 5.3, we show a simulated spectrum (XIS/FI; 0.5–10.0 keV) withthe typical GC parameters (see caption). The plasma emission code used in this chapter is
APEC1 and the absorption is evaluated in solar abundances2 .
1Astrophysical Plasma Emission Code: A model of an emission spectrum from diﬀuse gas in collisionallyionized equilibrium (Smith et al. 2001).
2
The solar abundance in this chapter is referred to the values from Anders & Grevesse (1989).

- 68 -

5.2. DATA ANALYSIS AND RESULTS

ISM

Abs2

Y
Z
1-R

★

X

GCPE

Sgr A*

Y0~300 pc
XRNE

MC Abs1

R
8 kpc

ISM

Abs2

Observer V•
Fig. 5.2.—: A schematic face-on view for line-of-sight alignment. Coordinates X, Z, and Y
correspond to directions of l, b, and line of sight, with respect to Sgr A*.

5.2.3

Spectral Fitting

As shown ﬁgure 5.5, we extract seven spectra from ﬁve regions (solid circles in ﬁgure 5.1).
The spectra of C1 and C2 are obtained twice in the observations of 2006 and 2010, whereas
C3 is only observed in 2010. All spectra are simultaneously ﬁtted in the band of 0.5–10.0 keVwith the model given in equation (5.1). Since the spectral model is highly complicated withmany physical parameters, it is impractical to ﬁt the data with all parameters free. Therefore,
we adopt several reasonable constraints in the ﬁtting. The parameter settings are the sameas those in Sgr B (chapter 4). The FE is ﬁxed to an APEC model with kT =0.85 keV andabundance of Z=0.011, applied by small absorption of NH =0.17×1022 H cm−2 (chapter 4).
The center energy of the Fe Kα and Kβ lines are respectively ﬁxed to 6.4 keV and 7.05 keVwith the ﬂux ratio (Kα/Kβ) of 0.125 (Kaastra & Mewe 1993). The temperature of the HP(kTHP ) is ﬁxed to 6.5 keV (Koyama et al. 2007b), and the abundances of HP and LP are ﬁxed
to one solar (Z = 1). The parameter α, normalization ratio of HP/LP, can be approximatedto a constant near Sgr C. This estimation is inferred from results of Uchiyama et al. (2012)- 69 -

"#"!

B$(.

!!"#'/&0*1(23(%/$0#+-4(5%-4(67(*118!!(239:(5%-4(67(;!!"*"<";!#8((===(;!!"$%"&'()>";!#"$*+","&'() D-(!!"?@A:(!!(B:!!(3?C

B$(--(!"!&!"#$"%&#'()*+,&#+*+'(-"&#.

"#!

CHAPTER 5. STUDY OF SAGITTARIUS C

"#$

!

%

$

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:"$+E'();$F.

Fig. 5.3.—: Simulated XIS/FI spectrum with typical GC parameters reported in chapter 4. Modelcomponents of GCPE (kTLP =1 keV, kTHP =6.5 keV, α=0.27, R=0.5), XRNE (EW6.4 keV =1.6 keV,
Γ=1.7), FE, and CXB are shown in red, blue, green, and cyan lines, respectively. The column
densities of Abs1 and Abs2 are set to 10×1022 H cm−2 and 6×1022 H cm−2 , respectively.

Fig. 5.4.—: Simulated XIS-FI spectra of the GCPE (kTL =1 keV, kTH =6.5 keV, α=0.27,
NH (Abs1)=10×1022 H cm−2 , NH (Abs2)=6×1022 H cm−2 ) for the parameter R from 0.1 to 0.9.

- 70 -

5.2. DATA ANALYSIS AND RESULTS(c.f., ﬁgure 3 and table 2 therein), which indicate that the two plasmas exhibit nearly thesame spatial distribution with an e-folding length of ∼0.◦ 6.
After all, the free parameters are the temperature of LP (kTLP ), the normalization ratioof HP/LP (α= normHP /normLP ), the photon index (Γ) of the power law, and the equivalent
width of the 6.4 keV line with respect to the power-law continuum (EW6.4 keV ); these freeparameters are linked in all regions including C1–3 and ref1–2. These settings are designed
to verify a common and consistent XRN model. Other parameters, such as line intensity(I6.4 keV ) and absorption (NH ), are free and independent for each region. To investigate thetime variability for C1 and C2, the 6.4 keV intensities of the 2006 and 2010 spectra are set
as free parameters, while other parameters are common between the two spectra.
As is seen in ﬁgure 5.1b, C1 is entirely overlapped with the ”Chimney”, which is athermal plasma outﬂowing from G359.41−0.12. In order to include the Chimney contamination, we add the spectral model with ﬁxed parameters (kT , NH , Z) reported in Tsuru
et al. (2009). The ﬂux (normalization of Chimney) is calculated from the average surface
brightness and the region size of C1 (radius=2.′ 0), which is 35% of the GCPE ﬂux.
All seven spectra are generally reproduced by the model of equation (5.1) with a reduced
χ /d.o.f of 3773/3031=1.26. The ﬁtting results are shown in ﬁgure 5.5 and the best-ﬁt
parameters are listed in table 5.2. The large EW6.4 keV of ∼ 1.2 keV and the strong absorption(NH (Abs1)∼ 1023 H cm−2 ) agree with the XRN scenario for all the regions. The NH of Abs2
are in the range of 5.0–6.5 ×1022 H cm−2 . These values are consistent with the interstellar
medium (ISM) absorption toward the GC region (Nobukawa et al. 2010; Sakano et al. 2002).
As for the time variability of the 6.4 keV line between the 2006 and 2010 observations, C1
exhibits a small increase of 8% at 2.9-σ level, while C2 shows no signiﬁcant change (seeﬁgure 5.6).
2

For conﬁrmation, we also apply independent ﬁtting with free EW6.4 keV parameters foreach region. The spectra of C1, C2, and C3 show EW6.4 keV of 1.1–1.5 keV, 1.1–1.6 keV, and
0.7–1.3 keV, respectively. These values are consistent with the ﬁtting result when EW6.4 keV
is a common free parameter (EW6.4 keV =1.15–1.27 keV; table 5.2), and agree with the XRN
scenario of large EW6.4 keV . However, the independent EW6.4 keV ﬁttings result in large
statistical errors in the R parameters (about 2–4 times larger than those in table 5.2), which
is not practical for determination of cloud positions and discussions in chapter 7. Thus, we
adopt the ﬁtting results with EW6.4 keV as a common free parameter between the regions(table 5.2).

- 71 -

CHAPTER 5. STUDY OF SAGITTARIUS C*!+-.""/0-./!+01""23"$!

%&'()*+,*

C1 (2006)
xis 0+2+3

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xis 0+3

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xis 0+3

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xis 0+3

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Fig. 5.5.—: Fitting results for all selected spectra (regions represented by black solid circles inﬁgure 5.1a) near Sgr C. The FI (XIS 0+2+3 for 2006; XIS 0+3 for 2010) spectra with 1-σ error
bars and the best-ﬁt models (c.f., ﬁgure 5.3) are represented with solid crosses and dashed lines,
respectively. The best-ﬁt parameters are summarized in table 5.2.

- 72 -

!!"#$%&''()*+!',-"'./+0'1+)'23./45+06

5.3. DISCUSSION FOR SGR C

563
562

!!"#
"!"$

561
560
567
.//0 .//1 .//2 .//3 .//4 ./5/ ./55 ./5.
!"#$%&'()*+,-$'%

Fig. 5.6.—: Light curves (2006–2010) of the 6.4 keV intensity near Sgr C. Data points from C1and C2 (c.f., table 5.2) are shown with the ﬁlled and open boxes, respectively.

5.3

Discussion for Sgr C

5.3.1

Association with the Radio MCs

The best-ﬁt values of R (position indicator; see ﬁgure 5.2) are 0.23, 0.76, and 0.26 for
C1, C2, and C3, respectively (table 5.2). These results indicate that C1 and C3 are inforeground to whereas C2 is beyond Sgr A*; they are largely separated in the line-of-sight
positions, although proximately located at the projected distance. We search for the radio
counterparts of C1, C2 and C3 from the CS (J=1→0) maps (Tsuboi et al. 1999) in diﬀerent
velocity ranges from −250 km s−1 to 250 km s−1 . As shown in ﬁgure 5.1c, C2 and C3 arerespectively in good coincidence with MCs in the velocity ranges of −60 ± 20 km s−1 and+50±30 km s−1 . On the other hand, C1 is observed in the two separated velocity ranges of−60±20 km s−1 and −100±20 km s−1 . Using the data corresponding C2 and C3, we derived
a conversion factor between the X-ray column density (Abs1; table 5.2) and the integrated
CS intensity (contours; ﬁgure 5.1c), which is [CS/NH ] ≃ 2.5 [K km s−1 /(1022 H cm−2 )]. The
X-ray NH of MC (Abs1) used in this study is determined directly from the absorption of
GCPE continuum behind the cloud (c.f., ﬁgure 5.2). This method is available even if the
MC is not an XRN, and thus does not depend on the properties of the illuminating source.
Assuming the same [CS/NH ] factor near Sgr C, the observed X-ray column density of C1(6.0–6.8×1022 H cm−2 ; table 5.2) indicates a CS intensity of 15–17 K km s−1 . This generally
excludes the −60 km s−1 MC as the counterpart of C1, because the average CS intensity is
only 5–10 K km s−1 . The −100 km s−1 MC has a CS intensity of 10–25 K km s−1 , which- 73 -

CHAPTER 5. STUDY OF SAGITTARIUS C
Table 5.2—: Best-ﬁt spectral parameters of the Sgr C regions∗ .

Local parametersNH† (Abs2) I6.4 keV ‡FE norm§ GCPE normLP ∥ Fraction R#
+0.3
+0.05
5.1−0.3
1.60−0.05 0.060+0.004
3.0+0.1
0.23 +0.03−0.004
−0.1
−0.04
+0.05
1.73−0.05
0.72+0.12
2.0+0.1
0.100+0.005
1.51 +0.03
6.5+0.3
C2
11.4+1.6−0.03−0.1−0.005
−0.03
−0.2
−1.4
+0.03
1.51−0.03
C3
8.7+0.4
5.0+0.3
1.16+0.05
0.082+0.005
3.0+0.1
0.26+0.03
−0.6
−0.2
−0.05
−0.003
−0.1
−0.03
ref1
7.9+0.5
5.3+0.2
0.36+0.03
0.059+0.003
2.3+0.1
0.40+0.04
−0.8
−0.1
−0.04
−0.003
−0.1
−0.03
+0.4
+0.05
+0.006
+0.1
+0.05
ref2
7.6+0.5
6.4
0.78
0.082
3.3
0.31
−0.7
−0.3
−0.05
−0.006−0.1−0.05
Global parametersHPkTLP
α ratio ( norm
EW6.4 keV
Photon index ΓnormLP )
+0.03
+0.01
0.92−0.05 keV
0.28 −0.05
1.19+0.08
keV
1.63+0.08
−0.04
−0.09
2
∗∗
χ /d.o.f = 3773/3031 = 1.26∗
The spectral regions are shown in ﬁgure 5.1. The uncertainties are at 90% conﬁdence level (statistical).
†The value of column density in the unit of 1022 H cm−2 .
‡Intensity in the unit of 10−6 photons cm−2 s−1 arcmin−2 .
§The normalization factor of the APEC model for the foreground emission. It is normalized with the region
size A [arcmin2 ] and expressed as 7.07 × 10−13 /(4πD2 A) EM [cm−5 arcmin−2 ], where D and EM are thedistance to the source [cm], and the emission measure [cm−3 ], respectively.
∥
In the same expression as §, but the normalization factor of LP in the GCPE.
#
R is the fraction of the GCPE not absorbed by MC (see equation (5.1) and ﬁgure 5.2), which is theindicator of the line-of-sight position for XRNe).
∗∗
The results are obtained by the simultaneous ﬁtting of the 14 spectra from the FI and BI data.
RegionC1

Obs.Yr
2006
2010
2006
2010
2010
2006
2010

NH† (Abs1)
6.5+0.3−0.5

is in agreement with the NH estimated by the X-ray data. Thus, C1 is most likely to besituated in the velocity range of −100 ± 20 km s−1 .

5.3.2

Two-temperature Structure of the GC Plasma and the XRNParameters

The best-ﬁt temperature of LP (kTLP ) and the mixing ratio (α = normHP /normLP ) in
the GCPE at Sgr C are 0.87–0.95 keV, and 0.23–0.29, respectively. These values are nearly
equal to those in Sgr B region: kTLP =0.81–0.91 keV and α =0.26–0.28 (chapter 4)). Sgr C
and Sgr B are located at nearly the symmetrical positions with respect to Sgr A* on the
Galactic plane, and hence these results indicate that the GCPE distribution is symmetricalnot only in the ﬂux proﬁle (Uchiyama et al. 2012) but also in the spectral shape.
The photon index and EW6.4 keV for the 6.4 keV clumps in Sgr C are 1.54–1.71 and 1.15–
1.27 keV, respectively (table 5.2). These values are approximately equal to the Sgr B results- 74 -

5.3. DISCUSSION FOR SGR C(chapter 4). The large EW is in favor of the X-ray reﬂection and ﬂuorescent origin (e.g.,
Koyama et al. 1996; Murakami et al. 2001; section ). The photon index is consistent withcanonical low-luminosity active galactic nuclei (AGN; e.g., Ishisaki et al. 1996; Terashima &Wilson 2001), which favors the XRN scenario owing to the past ﬂare of Sgr A*.

5.3.3

Time Variability of the 6.4 keV Line and the Flare Propertyof Sgr A*

Since the discovery of the neutral iron (6.4 keV) line near Sgr B as the X-ray echo of the
SMBH (Koyama et al. 1996), one basic question has been asked that whether it was a single
ﬂare from Sgr A*. As the solid observational facts by the same instrument (Suzaku/XIS),
Nobukawa et al. (2011) detected the decrease of the 6.4 keV intensity in Sgr B (Sgr B2 andM 0.74−0.09) by a factor of 2.0–2.5 from 2005 to 2009. In contrast, our data show thatthe 6.4 keV intensity in Sgr C1 and C2 is stable or slightly increasing in 2006–2010. The
Sgr B results suggest a ﬂare to be shorter than 4 years, while the Sgr C results require a ﬂareduration longer than 4 years. Hence, these are unlikely due to a same ﬂare. In addition,
XMM-Newtoon (Ponti et al. 2010) found the superluminal motion, ﬂux increase and decrease
within a ∼15-ly region, propagating near Sgr A (0.◦ 0< l <0.◦ 2) in 2002–2009. In summary,although observed in the same years, the 6.4 keV emission (X-ray echoes) around the GC
region shows diﬀerent time behaviors, which are diﬃcult to be explained by the Sgr A*
luminosity at a same time in the past. The most direct veriﬁcation for the single-ﬂare (ormultiple-ﬂares) scenario is to see whether all XRNe around the GC are lying near a single
parabola with focus at Sgr A*, from the face-on view (c.f., ﬁgure 2.22). We will discuss itin the chapter 7.

- 75 -

CHAPTER 5. STUDY OF SAGITTARIUS C- 76 -

Chapter 6
Study of Sgr D & Sgr E
In this chapter, we apply the model developed in Sgr B and Sgr C (chapter 4 andchapter 5) to the Sgr D and Sgr E regions located near the edges of the central molecularzone (CMZ; section 2.1.2). We probe and investigate the farthest molecular clouds from the
GC for the 6.4 keV emissions.

Contents
6.1

Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

77

6.2

Data Analysis and Results . . . . . . . . . . . . . . . . . . . . . .

78

6.2.1

Images . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

78

6.2.2

Spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

78

Discussion for Sgr D and Sgr E . . . . . . . . . . . . . . . . . . .

80

6.3.1

Association with the Radio MCs . . . . . . . . . . . . . . . . . . .

80

6.3.2

Origin of the 6.4 keV Line . . . . . . . . . . . . . . . . . . . . . . .

81

6.3

6.1

Observation

The Sgr D and Sgr E regions, located at the contrary GC sides of l ∼1.◦ 2 and −1.◦ 2,
are bright in CO and/or CS emissions (e.g., Oka et al. 1998; Tsuboi et al. 1999), whichsuggest molecular clouds at the boundary of the CMZ. Suzaku had two pointing observations
with XIS in normal mode to Sgr D region on March 2007 and Sgr E on September 2008,
respectively. The observation log is shown in table 6.1.

77

CHAPTER 6. STUDY OF SGR D & SGR E
Table 6.1—: Suzaku Observations near the Sgr D and Sgr E Regions.

Pointing direction∗
α (J2000.0)
δ (J2000.0)hms
GC Sgr D
501059010 17 48 21. 91 −27◦ 56′ 08.′′ 1
GC LP7 (Sgr E) 503013010 17h 42m 41.s 85 −30◦ 04′ 13.′′ 4Target name∗

Obs.ID

Obs.date
Exposure(yyyy-mm-dd)(ks)
2007-03-15
98.3
2008-09-16
50.1

The center of the XIS ﬁeld of view.

6.2

Data Analysis and Results

In the generation of images and spectra for Sgr D and Sgr E, the screening and reduction
for the XIS data were performed in the same way as those of Sgr B (described in section 4.1).
The FI spectra were merged in the spectral ﬁtting, and the NXB was subtracted from the raw
data. In the following analyses, the software was HEADAS version 6.12 and the calibration
database was the version released on 2011-02-10.

6.2.1

ImagesIn ﬁgure 6.1, we show the X-ray band images near the Sgr D and Sgr E regions.
The 6.7 keV maps (ﬁgure 6.1 middle), tracing the large-scale high-temperature (HP:kB T ≈6.5 keV; Koyama et al. 2007c) plasma, basically show uniform distribution in the
XIS FOVs. Similarly, the 2–3 keV maps (ﬁgure 6.1 top), tracing the low-temperature (LP:kB T ∼1 keV) plasma, is featureless except for the local diﬀuse sources reported by Sawada
et al. (2009).
From the 6.4 keV maps (ﬁgure 6.1 bottom), we discovered clumpy structures both in the
Sgr D and Sgr E FOVs; hereafter, we simply call them ”Clump D” and ”Clump E”. Clump D
and Clump E have angular sizes of ∼10′ and ∼5′ , which are much more extended than the
point spread function of XRT/XIS (∼1.8′ ; see ﬁgure 3.7). Thus, these 6.4 keV emissions
should have diﬀuse origin and are probably associated with molecular clouds (MCs), fromour experiences in the Sgr B and Sgr C regions.

6.2.2

SpectraIn order to investigate the details of new 6.4 keV clumps found in the Sgr D and Sgr E
regions, we extract spectra from the two circular regions shown in ﬁgure 6.1 bottom; the- 78 -

6.2. DATA ANALYSIS AND RESULTS(a) Sgr D
0.2

0.2

0.0

XIS/FOV
0.0(b) Sgr E
0.2keV
2–3

0.2
0.2

2–3 keV

XIS/FOV
0.0

0.0
0.0

-0.2

-0.2

1.6
0.2

1.4

-0.2
-0.2

1.2

1.0

Galactic latitude (degree)

Galactic latitude (degree)

DS1, PS1(Sawada+09)

0.8

6.7 keV
XIS/FOV

0.0

-0.2
1.6
0.2
0.2

1.4
0.2
0.2
10

1.2
11

1.0
13

0.8

1.4
1.6
0.2

Great Annihilator
359.0
358.8
358.6
359.2
1.2 359.0
1.0 358.8
0.8
1.4
1.2
1.0

358.6
0.8

6.7 keV
XIS/FOV

0.0

-0.2
359.2

359.0

358.8

358.6

142

275

539

0.2

6.4
keV
17

-0.2

76

6.4 keV

XIS/FOV
0.0
0.0

0.0
0.0

M0.74
-0.2
-0.2

-0.2
-0.2
Clump D

1.6
1.60.8 359.2
1.4

XIS/FOV

0.0

359.2
359.0
1.4
1.2

359.0
358.8
1.2
1.0

358.8
358.6
1.0
0.8

←Sgr C-0.2
359.2

Galactic longitude (degree)

Clump E
359.0

358.8

358.6

Galactic longitude (degree)

Fig. 6.1.—: X-ray band images of the (a) Sgr D and (b) Sgr E regions by Suzaku/XIS. The bin
size of images is ∼ 1′ . The top, middle, and bottom panels represent images of the medium band
(2–3 keV; S XV and continuum), the 6.7 keV (Fe XXV) line, and the 6.4 keV (Fe I) line, respectively.
The 6.4 keV- and 6.7 keV- line images are intensity maps constructed from the 6.3–6.5 keV bandand the 6.5–6.8 keV band, respectively, after continuum subtractions. The XIS ﬁeld of views (FOV)of the data analyzed in this chapter (table 6.1) are represented by the large solid boxes. The regions
for spectral analysis (Clump D and Clump E) are indicated with solid circles; the bright sources in
Sgr D ((a) top; DS1 and PS1; Sawada et al. 2009) are excluded using the region shown with thedashed lines.

region containing local sources in Sgr D (DS1 and PS1; see ﬁgure 6.1a middle) are excluded.
We apply the spectral model developed in Sgr B and Sgr C (c.f., equation (5.1) and ﬁgure 5.2)
to the spectra of Clump D and Clump E. The temperature of LP, temperature of HP, and- 79 -

CHAPTER 6. STUDY OF SGR D & SGR E
the normalization ratio (HP/LP) are ﬁxed to kB T =0.9 keV, kB T =6.5 keV, and α = 0.27,
respectively. The temperature and absorption of the foreground emission (FE) and thephoton index of the continuum (Γ = 1.7) are also ﬁxed to the results obtained in Sgr B
and Sgr C. Other parameters are free in spectral ﬁtting. We set the equivalent width at
6.4 keV (EW6.4 keV ) as free parameters to testify the X-ray reﬂection process for each region.
As shown in ﬁgure 6.2, the spectra of Clump D and Clump E are well reproduced by our< 1.2. The best-ﬁt results are listed in table 6.2.
spectral model with χ2 /d.o.f ∼(a) Clump D(b) Clump E

Counts /s

0.1

0.01

+ xis0&3+ xis1

10−3

5

5

0

0

−5

+ xis0&3+ xis1

0.01

10−3

χ

χ

Counts /s

0.1

−5
0.5

1

2
Energy (keV)

5

10

0.5

1

2
Energy (keV)

5

10Fig. 6.2.—: XIS spectra of (a) Clump D and (b) Clump E (c.f., solid circles in ﬁgure 6.1 bottom).
The FI and BI results are shown in black and red, respectively. The best-ﬁt spectral models (c.f.,ﬁgure 5.3) are shown in the dashed lines.

6.3

Discussion for Sgr D and Sgr E

6.3.1

Association with the Radio MCs

Strong absorption of NH ∼7×1022 H cm−2 (Abs1) is detected both from Clump D andClump E, which should be due to the MCs lying in the line of sight. We look into the radio
CO maps (Oka et al. 1998) in diﬀerent velocity ranges to search for the MC counterparts.
From the comparison between the CO and 6.4 keV maps in morphology, we identiﬁed two
MCs in velocities near +100 km s−1 and −200 km s−1 that display good coincidence withclump D and clump E, respectively (ﬁgure 6.3). According to the CO maps (ﬁgure 6.3) andequation (2.3), the CO intensities of the counterparts in +110∼+120 km s−1 and −220 ∼−210 km s−1 suggest NH ≈2NH2 ∼4–8×1022 H cm−2 for Clump D and ∼2–4×1022 H cm−2 for- 80 -

6.3. DISCUSSION FOR SGR D AND SGR E
Table 6.2—: Spectral ﬁtting results of 6.4 keV clumps in Sgr D and Sgr E.

sourceNH† (Abs1) NH† (Abs2)
I6.4 †FE§
GCPE§
Clump D
6.7+1.5
5.1+0.3
0.27+0.03
0.12+0.04
0.80+0.10
−1.0
−0.3
−0.03
−0.04
−0.10
χ2 /d.o.f=278/258=1.08
Clump E
6.6+0.6
3.5+0.5
0.40+0.04
0.05+0.04
0.90+0.10
−0.6
−0.4
−0.03
−0.04
−0.10
χ2 /d.o.f=348/299=1.16

R#
0.53+0.23
−0.18

†
EW6.4
keV
0.69+0.12−0.11

0.28+0.07−0.08

0.76+0.14−0.11∗

The parameter uncertainties are at 90% conﬁdence level.
The values of absorption (NH ), 6.4 keV intensity (I6.4 ), and 6.4 keV equivalent width (EW6.4 keV ) arein the units of 1022 H cm−2 , 10−6 photons cm−2 s−1 arcmin−2 , and keV, respectively.
§
The normalization factors of APEC models for the LP of the Galactic center plasma emission (GCPE)
and the foreground emission (FE). The values are re-normalized with the region size A [arcmin2 ], expressed as 7.07 × 10−13 /(4πD2 A) EM [cm−5 arcmin−2 ], where D and EM are the distance to the source[cm], and the emission measure [cm−3 ], respectively.
#
R is the fraction of the GCPE suﬀered by Abs2, which indicates the positions of MCs (c.f., equation (5.1)
and ﬁgure 5.2).
†

Clump E, respectively. Although the above estimation contains uncertainties by a factor of∼ 2 due to the velocity range for integration, the NH values by CO observations are basicallyconsistent with the NH values of MC measured from X-ray absorption (Abs1).
The NH values of Abs2 are 3.5×1022 H cm−2 for Clump E and 5.1×1022 H cm−2 for
Clump D, which are similar to pervious reports of Galactic ISM absorption at |l| ∼1◦ –2.◦ 5 (e.g., Sawada et al. 2009, 2011; Uchiyama et al. 2012). Also by taking the high velocities(VLOS >100 km s−1 ) into consideration, we conclude that the Clump D and Clump E are
MCs located in the GC region, most likely as parts of the CMZ (c.f., section 2.1.2). The
obtained R values, the position indicator (c.f., ﬁgure 4.7; section 4.3.2), suggest the lineof-sight placement that Sgr D (R = 0.53) lies slightly behind, while Sgr E (R = 0.28) issigniﬁcantly in foreground to Sgr A*.

6.3.2

Origin of the 6.4 keV LineThe origin of the 6.4 keV emission can be inferred from the observed equivalent width(EW6.4 keV ), which are 0.6–0.8 keV for Clump D and 0.7–0.9 keV for Clump E. In order
to explain these EW6.4 keV values, the required Fe abundance with respect to solar is 2.5–3
times in the electron process, or 0.6–0.9 times in the X-ray process (c.f., section 2.4.1). Local
thermal sources near Sgr D show Fe abundance of less than 1 solar (Sawada et al. 2009) and
the Galactic ridge emission at |l| ∼2◦ shows sub-solar Fe abundance of 0.5–0.8 (Uchiyama
et al. 2012; Sawada et al. 2011), which are consistent with the values of X-ray origin. Thus,- 81 -

CHAPTER 6. STUDY OF SGR D & SGR E

Galactic latitude (degree)(a) V= 110 ~120 km/s(b) V= –220 ~ –210 km/s

0.5

0.5

0.0

0.0

–0.5

–0.5
–1.0
–1.5
Galactic longitude (degree)

1.5
1.0
Galactic longitude (degree)

Fig. 6.3.—: Radio 12 C6 O (J = 1 → 0) maps from Oka et al. (1998). (a) Sgr D region in velocity
range of +110∼+120 km s−1 . (b) Sgr E region in velocity range of −220 ∼ −210 km s−1 . Contours
are set at intervals of 15 K km s−1 . The red solid circles represent the regions used for X-ray
spectral analyses (c.f., ﬁgure 6.1. bottom)

we suggest that the 6.4 keV emissions from Clump D and Clump E are produced by X-ray
reﬂections. These 6.4 keV clumps are the furthest X-ray reﬂection nebulae (XRNe) presently
observed.

- 82 -

Chapter 7
Overall Results & Discussion
In this chapter, we summarise and analyse the results obtained in chapter 4–6. We
quantitatively determine the line-of-sight positions for the XRNe in the GC region anddiscuss their three dimensional distributions in association with MCs in the CMZ. We
ﬁnally derive an X-ray light curve for the activities of Sgr A* over the past six centuries.

Contents
7.1

Longitude Distribution of Physical Parameters . . . . . . . . .

83

7.2

Face-on Distribution of XRNe in the GC Region . . . . . . . .

85

7.2.1

Computation of Line-of-sight Position . . . . . . . . . . . . . . . .

85

7.2.2

Association with the CMZ . . . . . . . . . . . . . . . . . . . . . . .

90

Active History of Sagittarius A* . . . . . . . . . . . . . . . . . .

93

7.3.1

Derivation of X-ray Light Curve . . . . . . . . . . . . . . . . . . .

93

7.3.2

Behaviours of the Past Flare State . . . . . . . . . . . . . . . . . .

94

7.3.3

Interpretation as Low-Luminosity AGN . . . . . . . . . . . . . . .

97

7.3

7.1

Longitude Distribution of Physical Parameters

From the ﬁtting results in chapter 4–6, we investigate the overall l distribution of physical parameters in the GC region. As shown in ﬁgure 7.1a, we can obtain a clear picture thatthe total absorptions toward the GC are separated into two kinds: Abs1 and Abs2. Abs1
displays nearly uniform distribution with average of NH ≈ 6 × 1022 H cm−2 , which suggests
its origin as the Galactic ISM between the GC and the Sun. The NH values of Abs2 are
83

CHAPTER 7. OVERALL RESULTS & DISCUSSION
comparable or higher than those of Abs1, which implies that Abs2 comes from the giant
MCs (e.g., Sgr B2, Sgr C) in the CMZ. As shown in ﬁgure 7.1b, the l distribution of the
GCPE basically follows the exponential model by Uchiyama et al. (2012), which is decaying
from Sgr A* with an e-folding scale height of ∼0.◦ 6. Thus, the overall distributions of the
GCPE and its absorptions are self-consistent with the schematic picture we constructed inchapter 4 and 5 (see ﬁgure 4.4 and ﬁgure 5.2).

50

(a)

45

Abs1 (Galactic ISM)
Abs2 (GC-MC) + Abs1--- 6e22 H/ cm2

NH (10 22 H cm-2)

40
35

B

30
25

C

20
D

15

E

10
5
0

1.5

Emission Measure (a.u.)

10

1

0.5
0
-0.5
Galactic longitude l (degree)

(b)

-1

-1.5

GCPE
Foreground--- exp(–l/ 0.6º)

1

10-1

1.5

1

0.5
0
-0.5
Galactic longitude l (degree)

-1

-1.5

Fig. 7.1.—: Longitude (l) distribution of physical parameters in the GC. (a) Absorptions of the
GCPE. (b) Emission measures of the GCPE and the foreground emission (FE). The values in thevertical axis are shown in arbitrary units (a.u.).

- 84 -

7.2. FACE-ON DISTRIBUTION OF XRNE IN THE GC REGION

7.2

Face-on Distribution of XRNe in the GC Region

7.2.1

Computation of Line-of-sight Position

As illustrated in ﬁgure 5.2 and ﬁgure 4.4, the parameter R represents the line-of-sight
position of XRN inside the GCPE; R is the fraction of the GCPE in front of the XRN
with respect to the entire GCPE along the line-of-sight. Depending on the distance fromSgr A*, the GCPE is more or less contaminated by the Galactic ridge plasma emission(GRPE; section 2.3.1), which has a similar two-temperature spectral structure but a more
extended spatial distribution (Uchiyama et al. 2012). In order to estimate the XRN positions
numerically, we use the spatial distributions of these plasmas.

(a)

(b)
Sgr A*

1

GCPE Intensity (a.u.)

F (X), GCPE Intensity (a.u.)

I (X,Y)rGC=55 pc

10-1
10-2
-3

10

1

[rGC (pc), rGR (pc), A]–– [60, 6000, 800]--- [60, 6000, 900]–– [55, 5000, 900]--- [50, 4000, 900]–– [50, 4000, 1000]

10-1

10-4
300
200
100

To Sun
600
0
500
400
Y (pc) -100
300
-200
200
X (pc)-300 0 100

10-2
1

10

102

103

104

X (pc)

Fig. 7.2.—: (a) X-Y distribution of the GCPE, I(X,Y ) (see deﬁnition in text). (b) Proﬁle of the∫∞
GCPE along longitude (X-axis) calculated from integration of F (X) = −∞I(X, Y ) dY . The redsolid line represents the proﬁle adopted for calculations of the line-of-sight positions. Other proﬁleswith diﬀerent parameters, corresponding to 10%, 20%, and 10% error ranges for rGC , rGR , and A
in I(X, Y ), respectively, are also plotted to show model uncertainties estimated form the results of
Uchiyama et al. (2012) (c.f., ﬁgure 2.17).

We introduce new coordinates (X,Y,Z) with origins at Sgr A* (l = −0.◦ 056, b = −0.◦ 046;
Yusef-Zadeh et al. 1999), in which X, Z, and Y correspond to the directions of l, b, andline of sight, respectively (see ﬁgure 5.2). Because the relevant XRNe are all on or near
the Galactic plane (|bXRN − bSgr A∗ | ≤0.◦ 1), we ignore the oﬀset of Z, for simplicity, in thefollowing calculations. Uchiyama et al. (2011) and Uchiyama et al. (2012) derived the X-axis
ﬂux distribution for many emission lines of ionized atoms (e.g., Fe and S) in addition to the- 85 -

CHAPTER 7. OVERALL RESULTS & DISCUSSION
soft and hard X-ray continuum bands. The data were ﬁtted with a two-exponential function
along the Galactic plane (X-axis): F (X) = IC exp(−X/lC ) + IR exp(−X/lR ). The values of
e-folding length for HP and LP are similar to each other both in the GCPE and GRPE. In
this paper, we adopted lC =0.◦ 6, lR = 50◦ , and IC /IR =10 as the mean values (see table 2 inUchiyama et al. 2012).
We assume that the X-Y plane distribution of the plasmas (see ﬁgure 7.2a) is anoblate spheroid expanded with two-exponential function as I(X, Y ) = A × exp(−r/rGC ) +√exp(−r/rGR ), where r = X 2 + Y 2 . A is the ﬂux ratio (GCPE/GRPE), and rGC and rGR are
the e-folding scale lengths of GCPE and GRPE, respectively. Integrating I(X, Y ) along the∫∞
Y -axis as F (X) = −8
kpc I(X, Y ) dY , we ﬁnd that the resultant projection of X-distribution
is nearly the same exponential shape as those observed by Uchiyama et al. (2012). We
then compare F (X) with the results of Uchiyama et al. (2012) (ﬁgure 2.17) to match theshapes. The parameters are determined to be rGC =55 pc (or 0.◦ 4), rGR =5000 pc (or 36◦ ),and A=900. The uncertainties of these parameters, including statistical errors (at 90% C.L.)and systematical ﬂuctuations between proﬁles in diﬀerent bands from Uchiyama et al. (2012),
are approximately 10%, 20%, and 10% for rGC , rGR , and A, respectively. In ﬁgure 7.2b, weshow the results of numerical calculation for the GCPE proﬁle F (X) and its uncertainty
ranges.
Considering that an XRN locating at (X0 ,Y0 ) and X0 as its l oﬀset from Sgr A*, the lineof-sight position (Y0 ; c.f., ﬁgure 5.2) is a function of parameter R expressed in the following
equation.
∫ Y0

I(X0 , Y ) dY
−8 kpc I(X0 , Y ) dY

kpc
R(Y0 ) = ∫−8
∞

(7.1)

In ﬁgure 7.3, we show examples of numerical calculation for the Y -R relation (equation (7.1))
at diﬀerent X0 positions, using the adopted spatial parameters of the GCPE and the GRPE
(rGC , rGR , A; c.f., ﬁgure 7.2).
From the R values obtained in chapter 4–6 (table 4.4, table 5.2, table 6.2) and simulated
Y -R curves, we determined the line-of-sight position (Y0 ) for each XRN in Sgr B, Sgr C,
Sgr D, and Sgr E. Here, the Y errors due to the uncertainties of the GCPE spatial parametersare ≤ 20 ly at a given R value in range of 0.2–0.8 (see black lines in ﬁgure 7.3), which are
ignorable compared to the Y errors originated from statistical uncertainties of the best-ﬁt R> 50 ly; table 7.1). Figure 7.4 illustrates the resultant face-on view of XRNe
values (errors ∼around Sgr A* on the X-Y plane. The units of (X, Y ) are converted into the actual lengthof light year (ly). Detailed results are summarized in table 7.1.

- 86 -

7.2. FACE-ON DISTRIBUTION OF XRNE IN THE GC REGION

500(a) X0 = 100 pc

500[rGC (pc), rGR (pc), A]–– [60, 6000, 800]–– [55, 5000, 900]–– [50, 4000, 1000]

400
300

300

100

Y0 (pc)100

Y0 (pc)

200

0

0

-100

-100

-200

-200

-300

-300

-400

-400
0

0.1 0.2

0.3 0.4[rGC (pc), rGR (pc), A]–– [60, 6000, 800]–– [55, 5000, 900]–– [50, 4000, 1000]

400

200

-500

(b) X0 = 170 pc

0.5 0.6 0.7

0.8 0.9

1

-500

R

0

0.1 0.2

0.3 0.4

0.5 0.6 0.7

0.8 0.9

1

R

Fig. 7.3.—: Line-of-sight position (Y0 ) as a function of the parameter R (see equation (7.1)). The
cases for (a) Sgr B at X0 =100 pc and (b) for Sgr E at X0 =170 pc are shown. The red solid lines
represent the Y -R curves adopted for calculations of the line-of-sight positions. Other Y -R curves(black lines), calculated from diﬀerent parameters (rGC , rGR , A) in I(X, Y ), are also plotted toshow model uncertainties estimated from the results of Uchiyama et al. (2012)(c.f., ﬁgure 7.2 andﬁgure 2.17).

- 87 -

CHAPTER 7. OVERALL RESULTS & DISCUSSION

Table 7.1—: Three-dimensional positions∗ of XRNe.

XRNID
Clump D
M0.74−0.09
M0.74-sub
Sgr B2
Sgr B1
Sgr C3
Sgr C2
Sgr C1
Clump E∗l[deg.]
1.17
0.74
0.74
0.66
0.51
−0.41
−0.53
−0.56
−1.24

b
[deg.]
−0.04
−0.09
−0.15
−0.05
−0.11
−0.14
−0.15
−0.06
−0.02

X
[ly]
560
363
363
327
258
-162
-216
-230
-540

Z
[ly]
3
−25
−47
−2
−30
−43
−47
−7
13

Y†
[ly]
55+580−310-202+50−60-280+80−60-55+50−60-75+50−50−173+30−30
189+200−30−241+40−60+60-408−570

Values of l and b are the Galactic coordinates at the centerof spectral regions used for each XRN. Values of X, Z, and Y
respectively correspond to the positions along l, b, and line ofsight, with origins at Sgr A*.
†
The Y errors are calculated from statistical ranges of R at 90%
conﬁdence level.

- 88 -

7.2. FACE-ON DISTRIBUTION OF XRNE IN THE GC REGION

3-D Distribution of XRNe/MCs around Sgr A*

Z (ly)

200

Edge onD

Sgr A*

B2

0

C1

B1

C3 C2

M0.74

40

E
b
l

35

Face-on View of XRNe around Sgr A*-200
Face on

30

400
25

Y (ly)

200

C2

D

20

Sgr A*

0

B2

B1
C3

-200

NH (1022 cm-2)

600

M0.74

-400

15
C1

To
Sun

-600

E

10

↓
Y= -978

5

600

400

200

0

X (ly)-200 -400 -600

Fig. 7.4.—: (a) Three dimensional (face-on and edge-on) distribution of the XRNe/MCs around
Sgr A*. Colors and sizes of the circles represent column densities (NH ) and spatial sizes (fromspectral regions) of the XRNe, respectively. Coordinates X, Z, and Y correspond to positions in
the Galactic longitude (l), latitude (b), and the line of sight from Sgr A*, respectively. Position
parameters for each XRN are given in table 7.1.

- 89 -

CHAPTER 7. OVERALL RESULTS & DISCUSSION

7.2.2

Association with the CMZ

We have obtained a face-on view of XRNe around the GC, directly and solely fromthe X-ray data and analyses. The l-b positions and large column densities (NH ∼ 5 − 40 ×1022 H cm−2 ) of the XRNe are consistent with those of the giant MCs (e.g., Sgr B2 andB1) in the CMZ (Morris & Serabyn 1996). Thus, these XRNe should be major componentsof the CMZ and therefore represent its basic structure. As shown in ﬁgure 7.4b, the Sgr B
XRNe/MCs are mostly located in foreground to Sgr A*. Although proximately (∼40 ly)
located in the projection, Sgr C1 and C3 is in front of the X-axis (Sgr A*), while Sgr C2(the most massive one) is largely behind C1 by ∼400 ly beyond the X-axis. Clump E andClump D lie before and slightly behind Sgr A*, respectively.
In a general perspective of the CMZ in radio studies, the theory of the central bar potential hypothesized closed elliptical orbits called the X1 and X2 families, which are originatedfrom the Lindblad resonance between the gas and the bar (e.g., Contopoulos & Mertzanides
1977; Binney et al. 1991; Morris & Serabyn 1996). The X1 orbits are elongated along thebar, while the X2 orbits are rounder and perpendicular to the bar. In the l-V map, theX1 and X2 orbits should be projected as parallelogram structures, which is consistent withthe radio observational results (see ﬁgure 2.6). Thus, following the X-orbits models, theradio observations on velocities of MCs can predict the line-of-sight positions and their large
structures such as Arm I/II (Sofue, 1995) and the Galactic center bow (GCB; Tsuboi et al.
1999). These are, however, depending on the models and hence not conclusive.
The X-ray face-on view of MCs in the CMZ (ﬁgure 7.4b) is obtained without informationof velocities, which can therefore verify the X-orbits model in the bar potential. In ﬁgure 7.5,
we superpose the CMZ structures (Arm I/II) predicted by radio l-V observations and possible
X1/2 orbits onto our X-ray face-on view. We ﬁnd good coincidences between them. The
details of their relations are discussed in the following paragraphs.
Arm I, Arm II, and GCB
The X-Y distribution of XRNe suggests that Sgr B2, B1, M 0.74, Sgr C1, and Clump D
may belong to Arm I, while Sgr C2 and Clump E is in Arm II, because Sofue (1995) predictedthat Arm I and Arm II are respectively located in the foreground and background withrespect to Sgr A* by assuming a uniform circular rotation (c.f., ﬁgure 2.7a). These two arms
were identiﬁed in the CO observation of l-V diagrams (Sofue 1995; see ﬁgure 7.6): Arm I is
stretching from (−1.◦ 0, −200 km s−1 ) through (−0.◦ 7, −150 km s−1 ) to (0.◦ 9, +80 km s−1 )and Arm II is stretching from (−0.◦ 6, −80 km s−1 ) through (−0.◦ 1, +60 km s−1 ) to (1.◦ 0,+100 km s−1 ). Although Sofue (1995) only deﬁned Arm structures within |l| ≤1◦ (i.e.,|X| ≤400 ly), Arm I and Arm II respectively extend to east and west more than 1.◦ 0,- 90 -

7.2. FACE-ON DISTRIBUTION OF XRNE IN THE GC REGION

1000

(a)
X2 Orb.

25

A rm

II

D
B2

B1

C2

C3

M
0.74

Arm I

-500

X1 Orb.

Sgr A*
★

20

15

C1

NH (10222 cm–22)

Y (ly)
(ly)
Y

500

0

(b)

30

E

10
To
X
majo 1 ~1 Sun
r
5º
axis

1000
1000

500

0

X (ly)

5

-500

-1000

Fig. 7.5.—: (a) Face-on distribution of the XRNe around Sgr A* superposed with Arm I/II
(dashed lines) and possible orbits of X1 and X2 (grey lines). The orbit parameters (major-axis
length, eccentricity, inclination angle) used for inner X2, outer X2, and X1 are (310 ly, 0.2, 105◦ ),(620 ly, 0.5, 105◦ ), and (2400 ly, 0.9, 15◦ ), respectively. (b) Schematic view of X1/2 orbits andArm I/II structures (Morris & Serabyn 1996; Binney et al. 1991; Sofue 1995).

according to Morris & Serabyn (1996) (see ﬁgure 7.5b).
The face-on XRN alignments of M0.74, Sgr B1-2, Sgr C1, and Clump E (ﬁgure 7.4)
also agree with the GCB (ﬁgure 2.7c). The GCB is found from the l-b-V observation of
the CS line by Tsuboi et al. (1999), which shows continuous stream-like appearances startfrom (−0.◦ 80, −0.◦ 08, −150 km s−1 ), bends at (−0.◦ 7, −0.◦ 08, +50 km s−1 ), and end at(1.◦ 50, −0.◦ 15, +50–100 km s−1 ). The GCB corresponds to a more general picture of Arm I.
The absorption features (shadows at l < 0) are observed in the near-IR and H I maps, whichindicates that the eastern part of the GCB is closer to Sgr A* than the western part (Tsuboi
et al. 1999 and references therein; see ﬁgure 2.7c).
In ﬁgure 7.6, we show the l-V distributions of Arm I, Arm II, GCB, and the corresponding XRNe. The line-of-sight velocity for each XRN (c.f., table 7.2) had been estimated bycomparing their positions and morphologies on the l-b plane between X-ray and Radio data
(c.f., chapter 4–6).

- 91 -

CHAPTER 7. OVERALL RESULTS & DISCUSSION

300–– Arm I–– Arm II–– GCB

200

30

25

100
M0.74

0

20

C3
B2

Sgr A*
B1
C2

-100

15

C1

NH (1022 cm–2)

VLOS (km s-1)

D

10

-200
-300

E
l = –1ºl = 1º|600 400 200

5|

0 -200 -400 -600
X (ly)

Fig. 7.6.—: Distribution of the line-of-sight velocity of XRNe near the GC (c.f., table 7.1). The
vertical errors correspond to the typical velocity width of 40 km s−1 . The known l-V structures of
Arm I, Arm II, and GCB are shown with red, blue, and green solid lines, respectively; the dashedlines extending from Arm I and Arm II are interpretations in this thesis (see details in the text).

X1 and X2 Orbits
As shown in ﬁgure 7.5, we can draw two X2 and one X1 orbits in order to match theface-on distribution of XRNe. The Sgr B1–2 and C1–3 are clearly located on a inner X2
orbit, while Clump D, M0.74, and Clump E are on a outer X2 orbit. Clump D and E arealso near the intersections between the outer X2 and X1 orbits. The inclination angle of theX1 major axis with respect to our line of sight is about θ ∼ 15◦ − 30◦ , which is consistent
with model predictions (θ ∼15◦ to 45◦ ; e.g., Morris & Serabyn 1996).
In summary, as an original result, the face-on view of XRNe basically proves the radio
predictions for the CMZ structures such as Arm I and Arm II. The alignments of Sgr B
and Sgr C XRNe suggest the existence of X2 orbits and hence support the bar-potential
theory. This enhances the reliability of our new X-ray method which directly determines theline-of-sight position. Details of the X-ray and radio associations are listed in table 7.2.

- 92 -

7.3. ACTIVE HISTORY OF SAGITTARIUS A*
Table 7.2—: Association of XRNe with radio structures in the CMZ.

XRN
Clump D
M0.74−0.09
M0.74-sub
Sgr B2
Sgr B1
Sgr C3
Sgr C2
Sgr C1
Clump E
†

‡
l†
VLOS
Arm I (GCB)∗
1.17 +110
0.74
+50
y
0.74
+50
y
0.66
+50
y
0.51
+30
y
−0.41 +50
−0.53 −60
−0.56 −100y−1.24 −210y

Arm II∗y

X1∗
y

X2-in∗

X2-out∗
y
y
y

y
y
y
y
y

y
y

y

The Galactic longitude coordinates (in units of degree) at the center of spectral regions used
for each XRN/MC.
‡
Line-of-sight velocity (in units of km s−1 ) measured from radio observations of doppler shift.
Typical velocity width is 40 km s−1 .
∗Possible associations with the radio structures and orbits described in text and ﬁgure 7.5.
The mark ”y” represents yes.

7.3

Active History of Sagittarius A*

7.3.1

Derivation of X-ray Light Curve

On the basis of the XRN model (c.f., ﬁgure 2.22) of the past ﬂares from Sgr A*, we draw
the parabolas in ﬁgure 7.7, with their common foci at Sgr A*, to illustrate the equi-delay-time
contours for X-ray echoes. The time delay (T∗ ) of an echo, light-passage (Sgr A*-XRN-Sun)divided by light speed (c), is given as below.
√
T∗ = ( X 2 + Y 2 + Y )/c
[yr]
(7.2)
As shown in ﬁgure 7.7, we ﬁnd that the XRNe are not located on a single parabola(i.e., same T∗ ), which directly indicates multiple ﬂares from Sgr A* at diﬀerent epochs in thepast. This result is also supported by observations of diﬀerent a-few-years time variabilities,of either ﬂux decrease or increase, from 6.4 keV clumps in the Sgr B, Sgr C, and Sgr A regions(ﬁgure 5.6; ﬁgure 2.21; also see discussions in section 5.3.3). Then, we can quantitatively
investigate the activity history for Sgr A*.
Using results (table 4.4; table 5.2; table 6.2; table 7.1) of the Fe I intensity by reﬂection√(I6.4 keV ), the column density of XRN (Abs1), and the distance (D∗ = X 2 + Y 2 ), the required- 93 -

CHAPTER 7. OVERALL RESULTS & DISCUSSION
X-ray (2–10 keV) luminosity of Sgr A* (LX∗ ) can be expressed in the following equation (fromNobukawa et al. 2008; Sunyaev & Churazov 1998; also c.f., equation (2.20) and ﬁgure 2.22).

LX∗(I6.4 keV /10−6 photons cm−2 s−1 arcmin−2 ) × (D∗ /300 ly)2
≈ 1.0 ×(NH /1022 H cm−2 )[1040 erg s−1 ] (7.3)

Here, isotropic radiation with a photon index of Γ = 1.6 from Sgr A*, sphere structure(with sizes of the analyzed spectral regions) and the solar iron abundance for XRNe are theunderlying assumptions. Also note that equation (7.3) gives the average luminosity of a ﬂare
over the light-crossing time of each XRN (c.f., section 2.4.3).
Using equation (7.2) and (7.3), we compute T∗ and LX∗ for nine XRNe in the Sgr C,
Sgr B, Sgr D, and Sgr E regions. The resultant long-term light curve of Sgr A* is plotted
in ﬁgure 7.8 and its details are summarized in table 7.3. Sgr A* was active with many short
ﬂares between LX∗ = 1–3 × 1039 erg s−1 during the past ∼50 to ∼600 years (i.e., past ﬂarestate). Thanks to the measurement of line-of-sight position of Sgr C2 and Clump D (seeﬁgure 7.7), we ﬁnd that Sgr A* was experiencing periods of high luminosity already 600
years ago, which extends the light curve of Inui et al. (2009) and Ponti et al. (2010) back by
200–500 years.

7.3.2

Behaviours of the Past Flare State

During the past ﬂare state, Sgr A* exhibited sporadic ﬂuctuations (short ﬂares) at
200–300 ago and ∼100 years ago, as represented in ﬁgure 7.8. The short ﬂare at 200–300
years ago would be responsible for the short-term time variability (2005–2009) of Sgr B2 andM 0.74−0.09 (ﬁgure 2.21; Nobukawa et al. 2011; Koyama et al. 2008; Terrier et al. 2010),
and hence the time scale would be ∼5 years. The other short ﬂare at ∼100 years ago hassimilar time scale as inferred by the 6.4 keV time variability (increase) of Sgr C1 in 5 years(2006–2010) at 3-σ level (table 5.2). This ﬂare may correspond to the time variability of the
Sgr A XRNe. Muno et al. (2007) and Koyama et al. (2009) reported variability of the 6.4
keV line from the Sgr A cloud in a short time-scale of 3–5 years at a ∼ 5σ sigma level. Then
Ponti et al. (2010) and Capelli et al. (2012) predicted the ﬂare events occurred between ∼100and ∼400 years ago, which is basically consistent with our results.
The detailed behavior of the activity at ∼ 100 years ago is not completely clear. From theillumination pattern in the Sgr A complex, Ponti et al. (2010) proposed that the ﬂare shouldhave started ∼400 years ago, while Capelli et al. (2012) suggested that an active period ended∼150 years ago. However, the estimated ﬂare luminosities were diﬀerent: LX∗ ∼ 1039 erg s−1- 94 -

7.3. ACTIVE HISTORY OF SAGITTARIUS A*

5

600

4.5
Y=–X2/2cT*+cT* /2

Y (ly)

200
0

-200

T*
(yr)

4

800

3.5

C2

600

D
B2

3
2.5

B1
M
0.74

Sgr A*

400

C3
C1

-400

200

2
1.5

E

To
Sun

-600
600 400 200

501

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

400

0.5

0 -200 -400 -600
X (ly)

Fig. 7.7.—: Face-on distribution of XRNe superposed with parabolas (dashed lines) of equal-timedelay (T∗ ; see deﬁnition in text) contours; diﬀerent parabolas represent echoes due to Sgr A* ﬂares
in diﬀerent past epochs. Coordinates X and Y correspond to positions in the Galactic longitude(l) and in the line of sight from Sgr A*, respectively.

was suggested by Ponti et al. (2010), while LX∗ ∼ 1038 erg s−1 was suggested by Capelli et al.
(2012). These inconsistencies are mainly attributable to unclear line-of-sight positions andpartly due to the diﬃcult estimation of NH for the Sgr A XRNe. The luminosity determinedby Capelli et al. (2012) is ∼ 1037 –1038 erg s−1 around 70 and 130 years ago, which is about
10 times lower than that obtained here: LX∗ ∼ 1039 erg s−1 at ∼90 years ago. The highluminosity indicated by C1 may be attributed to a short-term ﬂare that occurred in thefading period reported by Capelli et al. (2012). A uniﬁed analysis for the Sgr A XRNe by
the same method performed in Sgr B, C, D, and E in this thesis would be able to solve these
problems.
Apart from the 6.4 keV clumps with compact angular distribution such as Sgr B2,
Sgr C1, and Sgr E, we also detected relatively weak 6.4 keV line (I6.4keV ∼ 5×10−7 photons
cm−2 s−1 arcmin−2 ) all over the GC region (e.g., Sgr C ref 1–2, ﬁgure 4.1, table 4.4; Sgr B
region 1, ﬁgure 5.1, table 5.2). The spectra of these regions exhibit large EW6.4 keV above
1 keV and high GC absorption of NH =7–12×1022 H cm−2 , which are consistent with the
XRN model. Thus, the spatially-featureless 6.4 keV emissions are probably originated fromthe X-ray reﬂections by diﬀuse molecular gas ﬁlling the CMZ rather than compact MCs(e.g., Sgr B2). In this case, the observed 6.4 keV intensity is the line-of-sight integration of- 95 -

CHAPTER 7. OVERALL RESULTS & DISCUSSION

LX* (erg s-1)

1040

E

M0.74

C1

B2

D
C2

1039

C3

B1

10 33–35
-700

-600

-500

-400
-300
-200
T * (years from 2010)

-100

0

Fig. 7.8.—: X-ray light curve of Sgr A* for the past 600 years (see table 7.3 for parameterdetails). Red arrows show the results from observations of short-term 6.4 keV variabilities between
5 years. The present luminosity of Sgr A* is quoted from Baganoﬀ et al. (2001) and Porquet et al.
(2003).

reﬂections by continuous (or multiple) X-ray ﬂares over the past 500 years (see details in
Chernyshov et al. 2012). The average luminosity of Sgr A* over past several centuries (LaveX∗ )can be estimated from equation (7.3) using the best-ﬁt parameters (I6.4keV and NH ) of Sgr C
ref 1–2 and Sgr B region 1. Although the gas is most likely distributed continuously along
the line of sight, the lower limit of LaveX∗ is still determined by the minimum distance, whichis the project distance at a given X position (D∗ =|X|). The LaveX∗ estimated from Sgr C ref 1,
ref 2, and Sgr B region 1 are respectively, >0.47, >0.55, and >0.84 in unit of 1039 erg s−1 .
These results are consistent with the picture that Sgr A* had been in a ﬂare state with highluminosity (LX∗ ∼1039 erg s−1 ) over centuries in the past.

- 96 -

7.3. ACTIVE HISTORY OF SAGITTARIUS A*
Table 7.3—: Parameters† of XRNe and past luminosity of Sgr A*.

XRNID
Clump D
Sgr C2‡
Sgr B2‡
Clump E
M0.74−0.09‡
Sgr B1
M0.74-sub
Sgr C1‡
Sgr C3
†

Obs.
D∗§
Year
[ly]
2007
562+290−2
2006 (2010) 287+150−202005 (2009) 331+20−5
2008
677+440−902005 (2009) 416+30−20
2006
269+20−10
2007
459+40−40
2006 (2010) 333+50−30
2010
237+20−20∥

T∗[year]
618+860−260
476+350−50
281+50−40
269+90−130
219+30−30
198+40−30
182+40
−20
92+10
−10
64+10
−10

L#
X∗
[1039 erg s−1 ]1.4+0.3−0.0
1.2+1.6−0.2 (1.2)+0.2
1.7−0.1 (0.65)
2.9+5.1−0.7
2.7+0.4−0.3 (1.4)
1.0+0.2−0.1
2.8+0.5−0.5
2.9+1.1−0.5 (3.1)
0.8+0.2−0.1

Errors of D∗ , T∗ , and LX∗ are estimated using the error ranges of line-ofsight positions Y (see table 7.1).
√§
Distance from Sgr A* calculated from D∗ = X 2 + Y 2 .
∥Time delay of echo from Sgr A* calculated from equation (7.2) and counted
from 2010.
#
Required past X-ray luminosity of Sgr A* calculated from equation (7.3),
also see text.
‡
Values in brackets are the results from observations of short-term 6.4 keVvariabilities between 5 years (C1 and C2, section 5.2.3, this work; B2 andM,0.74, Nobukawa et al. (2011)). Errors due to position uncertainties of Y
are not included.

7.3.3

Interpretation as Low-Luminosity AGN

39
The average X-ray luminosity of Sgr A* in the past ﬂare state was Laveerg s−1X∗ ≃ 2×10(table 7.3), which corresponds to ∼ 4 × 10−6 of the Eddington limit (LE ≃ 5×1044 erg s−1 at
4×106 M⊙ ). As shown in ﬁgure 7.9a, comparing the luminosities, we may state that Sgr A*
had been acting as a low-luminosity AGN, such as M 81* (Ishisaki et al. 1996; Devereux etal. 2003) and M 51* (Terashima & Wilson, 2001), during the past ∼500 years.

The present X-ray luminosity of Sgr A* is LX∗ ∼ 1033−35 erg s−1 (Baganoﬀ et al. 2001;
Porquet et al. 2003), and hence a sudden luminosity drop, by 4–6 orders of magnitude fromthe ﬂare state, should have occurred within ∼100 years. Schawinski et al. (2010) reported
that the quasar IC 2497 with mass of ∼ 109 M⊙ experienced an dramatic luminosity dropby more than 4 orders of magnitude in less than 70000 years. On the other hand, for
Galactic X-ray binaries with a black hole mass of ∼10 M⊙ (e.g., GRS 1915+105; Greiner etal. 1996), a state change from high/soft to low/hard was detected on time scales of ∼1 hour.
- 97 -

CHAPTER 7. OVERALL RESULTS & DISCUSSION

10-2

(b)

GRS 1915+105

IC 2497

104

Cyg X-1

Transition Time (year)(a) 1

LX / LEdd

10-4
Sgr A* (past) ★
M51 M81
IC 2497

10-6
10-8

M31

10-10
10-12
1

Sgr A* (present) ★
2

10

104

106

108

1010

MBH / MSolar

Sgr A* ★

102
1
Cyg X-1

10-2
10-4
10-6
1

GRS 1915+105

102

104

106

108

1010

MBH / MSolar

Fig. 7.9.—: Comparison of observational results between Sgr A*, Galactic black holes, and
LLAGNs. (a) Plot of present X-ray luminosities (normalized with corresponding Eddington limits)against black-hole masses. The past luminosity of Sgr A* by this work is added. (b) Plot of statetransition time against black-hole masses. The diagrams include data points of Sgr A* (Baganoﬀ
et al. 2003), M 81* (Ishisaki et al. 1996; Devereux et al. 2003), M 31* (Garcia et al. 2000), M 51*(Terashima & Wilson 2001; Kohno et al. 1996), IC 2497* (Schawinski et al., 2010), Cyg X-1 (Zhang
et al. 1997), and GRS 1915+105 (Greiner et al. 1996).

Assuming that the state-transition time scale is proportional to black-hole mass, at a massof 4×106 M⊙ , 70000 years (for 109 M⊙ ) and 1 hour (for 10 M⊙ ) are scaled down to 280 years
and up to 46 years, respectively. Then, the time scale of the luminosity drop for Sgr A* is< 100
estimated to be as fast as about 50–300 year, which is consistent with the prediction (∼year) in our light curve (ﬁgure 7.8).

- 98 -

Chapter 8
Conclusions
In this thesis, using the high-quality and high-statistic data by Suzaku, we carried out
an uniﬁed study on various diﬀuse X-ray emissions from the major MC (Sgr B, C, D, andE) regions around the GC. The analyses were performed by careful spectral modeling inconsideration of the three-dimensional alignments. Followings are the highlights.
• The diﬀuse 6.4 keV emissions over |l| ∼ 0.◦ 5–1.◦ 3, including clumpy regions associatedwith giant MCs and spatially-featureless dim regions, are consistent with the spectral
model of XRN.
• The diﬀuse thermal emission in 2–10 keV over |l| ∼ 0.◦ 5–1.◦ 3 can be explained by thespectral model of a two-temperature (kB TLP ≈ 0.9 keV and kB T HP ≈ 6.5 keV) CIE
plasma with the solar abundance; the mixing ratio between emission measures (EM )of the two plasmas is EMHP /EMLP ≈ 0.27.
• The XRNe of Sgr B, C, D, and E, with diﬀerent radial velocities, are largely separated
in the line-of-sight position with respect to Sgr A*. The Sgr B is mostly locatedin foreground. The most massive one of Sgr C is beyond, while another two are in
front. Sgr E and Sgr D are located before and slightly behind, respectively. The faceon distribution of XRNe basically proves the presence of dynamical structures in bar
potential hypothesized by radio observations.
• The past X-ray activity from Sgr A* was not a single event, but consisted of multiple
events. The SMBH had been in a ﬂare state with sporadic ﬂuctuations between LX =1–
3 ×1039 erg s−1 ; the short (∼5 years) ﬂares occurred continuously during the past 50
to 600 years.

99

CHAPTER 8. CONCLUSIONS- 100 -

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List of Figures
2.1

Optical picture of the Milky Way galaxy. . . . . . . . . . . . . . . . . . . . .

4

2.2

Global view of the Galaxy. . . . . . . . . . . . . . . . . . . . . . . . . . . . .

5

2.3

Optical extinction map of the Galaxy. . . . . . . . . . . . . . . . . . . . . . .

5

2.4Radio CO map of the Galaxy. . . . . . . . . . . . . . . . . . . . . . . . . . .

7

2.5

Face-on view of the molecular gas in the Galaxy. . . . . . . . . . . . . . . . .

7

2.6Radio CS map of the central molecular zone. . . . . . . . . . . . . . . . . . .

9

2.7

Diﬀerent predictions on the face-on view of the central molecular zone. . . .

9

2.8Mass correlation between the super mass black hole and its host bulge. . . .

10

2.9

Proper motions of S stars around SMBH, Sgr A*. . . . . . . . . . . . . . . .

11

2.10 X-ray Image of the vicinity of Sgr A* by Chandra. . . . . . . . . . . . . . . .

11

2.11 X-ray spectra of Sgr A*. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

12

2.12 Short X-ray ﬂares of Sgr A*.

. . . . . . . . . . . . . . . . . . . . . . . . . .

13

2.13 X-ray observation of the Galactic ridge by Tenma. . . . . . . . . . . . . . . .

14

2.14 6.7 keV mapping of the Galactic center by Ginga. . . . . . . . . . . . . . . .

15

2.15 CCD X-ray spectra of the GC region by ASCA and Suzaku. . . . . . . . . .

16

2.16 X-ray band maps of the GC plasmas by Suzaku.

19. . . . . . . . . . . . . . .

2.17 Intensity proﬁles of the plasma emission lines from the GC by Suzaku.

. . .

20

2.18 Intensity distribution of the 6.4 keV emission in the GC by ASCA. . . . . . .

21

2.19 Equivalent width of the 6.4 keV line produced by protons and electrons. . . .

22

2.20 X-ray spectra of giant MCs in the GC region. . . . . . . . . . . . . . . . . .

24

107

LIST OF FIGURES
2.21 Time variabilities of the 6.4 keV ﬂux from giant MCs in a few years. . . . . .

24

2.22 Schematic face-on view of the XRN geometry. . . . . . . . . . . . . . . . . .

25

2.23 The 6.4 keV map of the Galactic center by Suzaku. . . . . . . . . . . . . . .

27

3.1

Schematic view of the Suzaku satellite in orbit. . . . . . . . . . . . . . . . .

30

3.2

Conﬁguration of the detectors in the Suzaku satellite. . . . . . . . . . . . . .

31

3.3

Schematic view and picture of the XRT. . . . . . . . . . . . . . . . . . . . .

33

3.4

Schematic cross-section of Wolter-I type X-ray mirror. . . . . . . . . . . . . .

34

3.5

Total eﬀective area and vignetting of the XRT. . . . . . . . . . . . . . . . . .

35

3.6

The Image of SS Cyg at the optical axis taken with XRT-I0. . . . . . . . . .

36

3.7

PSF and EEF of the XRT. . . . . . . . . . . . . . . . . . . . . . . . . . . . .

36

3.8

Readout architectures of CCD and APS. . . . . . . . . . . . . . . . . . . . .

38

3.9

Picture and cross-section of XIS sensor. . . . . . . . . . . . . . . . . . . . . .

38

3.10 XIS/CCDs in front illumination and back illumination. . . . . . . . . . . . .

39

3.11 Schematic view of the XIS CCD. . . . . . . . . . . . . . . . . . . . . . . . .

40

3.12 QE of the XIS. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

40

3.13 Long-term trend of the XIS energy resolution. . . . . . . . . . . . . . . . . .

42

3.14 XIS spectra of the NXB. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

43

4.1

XIS band images around the Sgr B region. . . . . . . . . . . . . . . . . . . .

47

4.2

One-dimensional proﬁles of the Sgr B region. . . . . . . . . . . . . . . . . . .

49

4.3

FI spectra of M 0.74-0.09 and Sgr B2. . . . . . . . . . . . . . . . . . . . . . .

50

4.4

Schematic view of the MCs distribution along the line of sight. . . . . . . . .

53

4.5

Correlations between the 6.4 keV-line ﬂux and absorptions. . . . . . . . . . .

57

4.6

Flux distribution of the 6.7 keV and 2.45 keV lines. . . . . . . . . . . . . . .

59

4.7

Relative positions of each Sgr B MC along the line of sight. . . . . . . . . . .

60

4.8

The absorption (Abs1) vs the distance-corrected 6.4 keV-line ﬂux. . . . . . .

62- 108 -

LIST OF FIGURES
5.1

XIS (FI+BI) images near the Sgr C region. . . . . . . . . . . . . . . . . . . .

67

5.2

A schematic face-on view for line-of-sight alignment for Sgr C. . . . . . . . .

69

5.3

Simulated XIS/FI spectrum of GCPE. . . . . . . . . . . . . . . . . . . . . .

70

5.4

Simulated XIS-FI spectra of the GCPE for the parameter R. . . . . . . . . .

70

5.5

Fitting results for Sgr C spectra.

. . . . . . . . . . . . . . . . . . . . . . . .

72

5.6

Light curves (2006–2010) of the 6.4 keV intensity near Sgr C. . . . . . . . . .

73

6.1

XIS band images near the Sgr D and E regions. . . . . . . . . . . . . . . . .

79

6.2

XIS spectra of 6.4 keV clumps in Sgr D and Sgr E. . . . . . . . . . . . . . .

80

6.3

Radio CO maps of the Sgr D and Sgr E regions. . . . . . . . . . . . . . . . .

82

7.1 l distribution of absorption, GCPE, and FE in the GC region. . . . . . . . .

84

7.2

Spatial distribution of the GCPE. . . . . . . . . . . . . . . . . . . . . . . . .

85

7.3

Line-of-sight position (Y0 ) as a function of the parameter R. . . . . . . . . .

87

7.4

3-D distribution of the XRNe around Sgr A*. . . . . . . . . . . . . . . . . .

89

7.5

Face-on distribution of the XRNe around Sgr A* and X1/2 orbits. . . . . . .

91

7.6

Velocity distribution of XRNe near the GC. . . . . . . . . . . . . . . . . . .

92

7.7

Face-on view of XRNe around Sgr A*. . . . . . . . . . . . . . . . . . . . . .

95

7.8

X-ray light curve of Sgr A* in the past six centuries.

96

7.9

Comparison between Sgr A*, Galactic black holes, and LLAGNs.

- 109 -

. . . . . . . . . . . . .
. . . . . .

98

LIST OF FIGURES- 110 -

List of Tables
3.1

Summary of the Suzaku capabilities. . . . . . . . . . . . . . . . . . . . . . .

32

4.1

Suzaku Observations near the Sgr B Region . . . . . . . . . . . . . . . . . .

46

4.2

Summary of free and ﬁxed parameters in the ﬁttings for Model A, B, and C.

54

4.3

Results of best-ﬁt parameters∗ with Model A and B.

. . . . . . . . . . . . .

55

4.4

Detailed results of best-ﬁt parameters with Model C. . . . . . . . . . . . . .

57

5.1

Suzaku Observations at the Sgr C Region. . . . . . . . . . . . . . . . . . . .

66

5.2

Best-ﬁt spectral parameters of the Sgr C regions∗ .

. . . . . . . . . . . . . .

74

6.1

Suzaku Observations near the Sgr D and Sgr E Regions. . . . . . . . . . . .

78

6.2

Spectral ﬁtting results of 6.4 keV clumps in Sgr D and Sgr E. . . . . . . . .

81

7.1

Three-dimensional positions∗ of XRNe. . . . . . . . . . . . . . . . . . . . . .

88

7.2

Association of XRNe with radio structures in the CMZ. . . . . . . . . . . . .

93

7.3

Parameters† of XRNe and past luminosity of Sgr A*. . . . . . . . . . . . . .

97

111

LIST OF TABLES- 112 -

Acknowledgments
I express my deepest thanks to my supervisors Prof. T.G. Tsuru and Prof.
K. Koyama for giving me the chance, instruction, advice, and encouragement
to carry out this challenging and interesting study.
This work would also not be completed without fruitful discussions with
Dr. M. Nobukawa and Dr. H. Uchiyama.
I really enjoyed the days with the other instructors and colleagues in (orhad been in) the Kyoto-U/cosmic-ray (CR) group, which was indispensable
for a productive and healthy academic life.
I want to thank people at KEK and MIT with whom I have worked on the
SOIPIX detector.
Finally, I am so grateful to my family for the selfless supports throughout
my student career.

113