X-ray Study of the Variable Neutral Iron Line Emission of
the Sgr B2 Complex in the Galactic Center Region
Tatsuya Inui
Department of Physics, Faculty of Science
Graduate School of Kyoto University
March 13, 2008

Abstract
We report the ﬁrst discovery of the neutral iron line variability in the Galactic
molecular complex Sgr B2. This discovery is a crucial hint not only for the originof the neutral iron line emission in the Galactic center but also for the past activityof the quiescent Galactic center black hole Sgr A∗ in the present.
We operate a comparison analysis with the past observations from 1994 to 2005.
We make the cross calibration of the absolute eﬃciency among the four diﬀerent
observatories using the invariable 6.67 keV line and ﬁnd a clear variability of the
6.40 keV line emission. The variability is primarily originated from the two distinct sources, M 0.66−0.02 (Sgr B2 cloud) and G 0.570−0.018. The intensity of the
6.40 keV line has a peak in 2000 and decreases to 66% in ﬁve years. The ten year
time scale of variability requires that the origin invoking the 6.40 keV line emission
has nearly a light speed, strongly supporting the X-ray reﬂection scenerio. At thesame time, the 6.40 keV light curve indicates the past activity of the irradiating
source. The most possible source is the Galactic center black hole Sgr A∗ , which
was thought to be in an active phase about 300 years ago. We now see a decay of
the giant ﬂare through the 6.40 keV line of the molecular clouds.
The high energy resolving power of the Suzaku XIS provides the separate maps of
the K-shell transition lines from Fe I (6.40 keV) and Fe XXV (6.67 keV) and has discovered new X-ray sources from each map. Whereas the 6.67 keV lines are smoothly
distributed over the Sgr B2 region, which is due to the Galactic hot plasma, a local
excess is found near at (l, b) = (0.◦61, 0.◦ 01), possibly due to a new young SNR. The
plasma temperature is kT = 3 keV and the age is estimated to be younger than
7 × 103 years. The 6.40 keV image is clumpy with local excesses near at the Sgr B2
cloud and at (l, b) = (0.◦ 74, −0.◦ 09). Since the spectral features are quite similar to
those of the Sgr B2 cloud, the other excess can be a new X-ray reﬂection nebulae(XRN).

Contents
1 Introduction

12 Review
2.1 The Galactic Center Region . . . . . . . . . . . . . . . . . . . . . . .
2.1.1 Galactic Nucleus Sgr A∗ . . . . . . . . . . . . . . . . . . . . .
2.1.2 The Galactic Center Diﬀuse X-ray Emission . . . . . . . . . .
2.1.3 Supernova Remnants . . . . . . . . . . . . . . . . . . . . . . .
2.1.4 Giant Molecular Complexes . . . . . . . . . . . . . . . . . . .
2.2 A Giant Molecular Complex — Sgr B . . . . . . . . . . . . . . . . . .
2.2.1 Sgr B2 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2.2 Sgr B1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3 X-ray Emission Processes from Neutral Elements . . . . . . . . . . .
2.3.1 Photoionizaion . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.2 Collisional Ionization by Electrons . . . . . . . . . . . . . . . .
2.3.3 Collisional Ionization by Fast Ions . . . . . . . . . . . . . . . .

3
3
3
5
8
8
12
12
14
14
14
14
18

3 Instruments
3.1 X-ray Satellites and their Mirrors . . . . . . . . . . . . . . . . . . . .
3.1.1 ASCA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.1.2 Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.1.3 XMM-Newton . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.1.4 Suzaku . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2 CCD cameras . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.1 SIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.2 ACIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.3 EPIC . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.4 XIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

20
20
20
22
25
28
29
30
33
36
39

4 Observations
4.1 The Suzaku Observation . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.1 Data Collection . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.2 The Gain Tuning . . . . . . . . . . . . . . . . . . . . . . . . .
4.1.3 The Position Tuning . . . . . . . . . . . . . . . . . . . . . . .
4.2 The XMM-Newton Observations . . . . . . . . . . . . . . . . . . . . .
4.2.1 Data Collection . . . . . . . . . . . . . . . . . . . . . . . . . .

49
49
49
49
50
50
50

i

CONTENTS

ii

4.2.2 The Position Tuning . . . . . . . . . . . . . . . . . . . . . . .
The Chandra Observations . . . . . . . . . . . . . . . . . . . . . . . .
4.3.1 Data Collection . . . . . . . . . . . . . . . . . . . . . . . . . .
4.3.2 Position Accuracy . . . . . . . . . . . . . . . . . . . . . . . . .
The ASCA Observations . . . . . . . . . . . . . . . . . . . . . . . . .
4.4.1 Data Collection . . . . . . . . . . . . . . . . . . . . . . . . . .
4.4.2 The Position Tuning . . . . . . . . . . . . . . . . . . . . . . .

53
53
53
54
55
55
55

5 Result I — Suzaku Discoveries of New Diﬀuse Sources
5.1 Suzaku Overview of the Sgr B2 Region . . . . . . . . . . . . . . . . .
5.2 New Diﬀuse Sources . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.2.1 SNR Candidate — G 0.61+0.01 . . . . . . . . . . . . . . . . .
5.2.2 XRN Candidate — M 0.79−0.09 . . . . . . . . . . . . . . . . .

56
56
60
60
66

6 Result II — Time Variabilities of Neutral Iron Line
6.1 Spectral Analysis and Error Estimation . . . . . . . . . . . . . . . . .
6.1.1 Suzaku . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.1.2 XMM-Newton . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.1.3 Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.1.4 ASCA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.1.5 Comparison Analysis . . . . . . . . . . . . . . . . . . . . . . .
6.2 Variability Veriﬁcations on Individual Sources . . . . . . . . . . . . .
6.2.1 Surface Brightness Mapping . . . . . . . . . . . . . . . . . . .
6.2.2 M 0.66−0.02 . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.2.3 G0.570−0.018 . . . . . . . . . . . . . . . . . . . . . . . . . . .

74
74
75
77
81
81
82
86
86
86
93

7 Summary and Discussion
7.1 Ionized Iron Line Sources . . . . . . . . . . . . . . . . . . . . . . . . .
7.1.1 Galactic Center Diﬀuse X-ray Emission . . . . . . . . . . . . .
7.1.2 Supernova Remnants . . . . . . . . . . . . . . . . . . . . . . .
7.2 Neutral Iron Line Sources . . . . . . . . . . . . . . . . . . . . . . . .
7.2.1 The Origin . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.2.2 Relation to the Past Activity of Sgr A∗ . . . . . . . . . . . . .

96
96
96
96
97
97
98

4.3

4.4

A Tables for Spectral Results of M 0.66−0.02
104
A.1 Continuum Models . . . . . . . . . . . . . . . . . . . . . . . . . . . . 104
A.2 Background Analysis of the MOS and the PN detector . . . . . . . . 104
B Tables for Spectral Results of G 0.570−0.018
110
B.1 Continuum Models . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110
B.2 Background Analysis of the MOS and the PN detector . . . . . . . . 110

List of Figures
2.1
2.2

2.3
2.4
2.5

2.6
2.7

2.8
2.9

Adaptively smoothed 0.5–7.0 keV image of Sgr A∗ overlaid contoursof VLA 6 cm (Baganoﬀ et al. 2003). . . . . . . . . . . . . . . . . . . . 4
Wideband spectrum of Sgr A∗ . The spectrumof Sgr A∗ is based onmeasurements at 7 mm, 450 μm, 850 μm, 3.6 μm, 1.6 μm, and Xrays at 1 keV. The quiescent X-ray and near-IR ﬂuxes are taken from
Baganoﬀ et al. (2001), Genzel et al. (2003), and Ghez et al. (2005).
The quiescent ﬂux is shown in blue color and the ﬂux associated withﬂaring activity including the X-ray ﬂare is shown in red color. The
hard gamma ray emission from the Galactic center at 1 TeV based on
HESS observations (Aharonian et al. 2004) and the high ﬂuxes fromINTEGRAL observations are not connected and are only shown by
themselves in black. Further descriptions are shown in Yusef-Zadeh
et al. (2006). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 6
6.7 keV narrow band image of the Galactic Center obtained with theSuzaku XIS. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 7
2.45 keV narrow band image of the Galactic Center obtained with
the Suzaku XIS. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 9
Velocity-integrated CS J=1–0 emission in the Galactic center region(Tsuboi et al. 1999). The velocity integrated range is from −200 to
200 km s1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10
6.4 keV narrow band image of the Galactic Center with the Suzaku
XIS. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 11
2–10 keV band image of Sgr B2 obtained with the Chandra ACIS
(Takagi et al. 2002), overlaid on coutours of H II regions (Gaumeet al. 1995). The solid circles indicate X-ray point sources. . . . . . . 12
Adaptively smoothed Chandra image in the 6.15–6.55 keV band (Murakami
et al. 2001), overlaid on coutours of 13 CO line ﬂux (Sato et al. 2000) . 13(a) Cross sections for K-shell ionization of various atoms. Solid curvesindicate the calcurated result by electron impacts from the semiempirical formula. Dashed curves indicate cross sections for K-shell
photoionization of N and Fe. (b) Stopping range of low-energy electrons in hydrogen. See Tatischeﬀ (2003) for detail descriptions. . . . 15

iii

LIST OF FIGURES

iv

2.10 Simulated equivalent width of ﬂuorescent Fe K line as a functionof the thickness NH of scattering matter. The simulation is madein the following cases; I: primary source is directly observed, II: Xrays from primary source are absorbed by a surrounding matter, III:
primary source is totally hidden or scattering matter is irradiated by
an external source. . . . . . . . . . . . . . . . . . . . . . . . . . . . . 16
2.11 X-ray emission produced by suprathermal electrons with the source
spectrum of ∝ Ee−2 interacting in a thick target composed of a neutral
atomic gas of the solar abundance. The dashed line shows the eﬀectof photoelectric absorption, with a H column density of 4 × 1021 cm−2 . 17
2.12 X-ray emission produced by fast ions with cosmic-ray source (CRS)
composition. IB: inverse bremsstrahlung; SEB: secondary electron
bremsstrahlung; RRC: radiative recombination continuum. See Tatischeﬀ(2003) for detailed descriptions. . . . . . . . . . . . . . . . . . . . . . 19
3.1

Schematic pictures of the ASCA (upper left), Chandra (upper right),
XMM-Newton (lower left) and Suzaku satellites (lower right). . . . .
3.2 ASCA XRT eﬀective area (one telescope) for diﬀerent oﬀ-axis angle(θ) and azimuthal angle (φ). . . . . . . . . . . . . . . . . . . . . . . .
3.3 The four nested HRMA mirror pairs and associated structures. See
http://cxc.harvard.edu/proposer/POG/html/HRMA.html. . . . . . .
3.4 The HRMA/ACIS and HRMA/HRC eﬀective areas versus X-ray energy in linear (top) and log (bottom) scale. . . . . . . . . . . . . . . .
3.5 The net eﬀective area of all XMM-Newton X-ray telescopes, EPIC
and RGS. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.6 Eﬀective area of one XRT + XIS system, for both the FI and BI
chips. (No contamination included) . . . . . . . . . . . . . . . . . . .
3.7 Cross section of SIS (Gendreau 1995). . . . . . . . . . . . . . . . . . .
3.8 Left:Four individual CCDs are combined in a single package to makea single SIS focal plane (Gendreau 1995). Right:Conﬁguration of the
two SIS on the ASCA focal plane. . . . . . . . . . . . . . . . . . . . .
3.9 Energy resolution of the GIS and the SIS. For the SIS, measured or
estimated energy resolutions of SIS1 in the 1CCD mode are plotted
for at launch and 1, 2, 3, 4 and 5 years after the launch. . . . . . . .
3.10 The SIS spectrum for the dark earth data (Gendreau 1995). . . . . .
3.11 A schematic drawing of the ACIS focal plane; insight to the terminology is given in the lower left. Note the nominal aimpoints: on S3(the ‘+’) and on I3 (the ‘x’) (Weisskopf et al. 2003). . . . . . . . . . .
3.12 The ACIS background spectra of I0 (black) and S3 (red) extracted
from the Chandra blank-sky database. . . . . . . . . . . . . . . . . .
3.13 An example of the ACIS background counting rate versus time: S3(the BI chip; top curve) and I2 (the FI chip; bottom curve). . . . . .
3.14 A sketch of the ﬁeld of view of the two types of EPIC camera; MOS(left) and PN (right). . . . . . . . . . . . . . . . . . . . . . . . . . . .

21
23
24
26
27
29
31

31

32
33

34
35
37
38

LIST OF FIGURES
3.15 Energy resolution of MOS1-CCD1 (red asterisk), MOS2-CCD1 (green
square) and PN-CCD4 (black triangle). . . . . . . . . . . . . . . . . .
3.16 Upper: Background spectra for the MOS1 camera (left) and the PN
camera (right) during an observation with the ﬁlter wheel in the closed
position. Lower: Images of the MOS camera of Al K line at 1.48 keV(left) and the PN of Cu K at 8.04 keV (right) (Carter & Read 2007).
3.17 Layout of the XRTs on the Suzaku spacecraft. . . . . . . . . . . . . .
3.18 One XIS instrument. . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.19 The time history of the energy resolution for Mn Kα from the 55 Fe
calibration sources for XIS. . . . . . . . . . . . . . . . . . . . . . . . .
3.20 The XIS background rate for each of the four XIS detectors, withprominent ﬂuorescent lines marked. These spectra are based on theobservations of the dark Earth for ∼ 800ks. . . . . . . . . . . . . . . .
3.21 ACTY dependence of the NXB for XIS1 and XIS2. . . . . . . . . . .
3.22 Cut-oﬀ rigidity dependence of the NXB (average intensity in 5–10keV) for each sensor. . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.1

4.2

4.3

4.4
5.1

5.2

v

39

40
41
42
44

45
47
48

The Suzaku image of the Sgr B region in the 2–10 keV band. All thefour XIS data are co-added. The dotted square is the XIS ﬁeld ofview (FOV). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51
Background light curves of 2001 (left) and 2004 (right) observations
with the MOS and the PN detectors. Events with over 10 keV andpattern 0 are accumulated. Holizontal lines indicate the screening
criteria (0.4 counts s−1 for MOS and 1.0 counts s−1 for PN). . . . . . 52
The XMM-Newton image of the Sgr B region in the 0.3–10 keV band.
All MOS and PN data in 2004 are co-added. Circles indicate point
sources identiﬁed with the Chandra point source catalogue (Munoet al. 2003). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 53
The distribution of the Chandra aspect oﬀsets (Weisskopf et al. 2003). 54
The X-ray spectra from the full FOV of the XIS, except the CCD
corners irradiated by the build-in calibration sources. All four XIS
data sets are co-added. . . . . . . . . . . . . . . . . . . . . . . . . . . 57
Fluxes of the 6.40 keV (a) and the 6.67 keV lines (b), the ﬂux ratio of
6.97 keV/6.67 keV (c), and the continuum ﬂux (normalization of the
7-keV bremsstrahlung) map (d). The embedded numbers are ﬂuxes
in units of 10−6 photons cm−2 s−1 arcmin−2 or ﬂux ratios. The color
means that the average value is white, and the increasing (decreasing)excess above (below) the average is given with increasing deepness ofred (green), where the maximum (minimum) is 2.0 (0.1) times of theaverage value. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59

LIST OF FIGURES
5.3

5.4

5.5

5.6

5.7

5.8

5.9
6.1

6.2

6.3

6.4

6.5

The 6.67 keV line map (the 6.58–6.74 keV band map) showing brightspot at the northwest corner with the 2–5 keV continuum map as
green contours. The ﬂuxes are normalized by the 6.7 keV ﬂat-ﬁeld
image. The source and background regions are shown by the whitesolid and dotted ellipses, respectively. . . . . . . . . . . . . . . . . . .
The exposure corrected 6.67 keV line maps of (a) Chandra 2000 and(b) XMM-Newton 2004. The source regions are shown by the solid
circles. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Top: the X-ray spectrum of the sum of 3 FI CCDs (XIS0, 2 and 3)
for the new SNR (G 0.61+0.01) with the best-ﬁt VPSHOCK model.
Bottom: same as the top, but of the BI CCD (XIS1). . . . . . . . . .
The 6.40 keV line map (the 6.33–6.46 keV band map) showing brightspots at the Sgr B2 cloud and M 0.74-0.09. The ﬂuxes are normalized
by the 6.40 keV ﬂat-ﬁeld image. The sources and background regions
are shown by the solid circles and dotted ellipse, respectively. . . . . .
The exposure corrected 6.40 keV line maps of (a) Chandra in 2000
and (b) XMM-Newton in 2004. The source regions are shown by thesolid circles. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Top: the X-ray spectrum of the sum of the 3 FI CCDs (XIS0, 2 and
3) for M 0.66−0.02 (Sgr B2 cloud) with an absorbed power-law modeland two Gaussian lines. Bottom: same as the top but of the BI CCD(XIS1). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Same as ﬁgure 5.8, but for a new source M 0.74-0.09. . . . . . . . . .

vi

60

62

63

67

69

70
71

The XMM-Newton image in 6–7 keV band. The source region for
the Sgr B2 region is described in a solid line. Dashed line circles areindividual sources discovered in this region. . . . . . . . . . . . . . . . 75
The blank-sky-subtracted spectra of the Sgr B2 region, with theSuzaku XIS FI sensors (a) and BI sensor (b) in 2005, XMM-Newton
MOS (c) and PN (d) in 2004, (e) and (f) in 2001, Chandra ACIS-I(g) in 2001 and (h) in 2000, and ASCA SIS (i) in 1994. The solidlines indicate the best-ﬁt models (see text). . . . . . . . . . . . . . . . 76
The 6.67 keV line surface brightness obtained with each observationin the Sgr B2 region. In this plot, each error bar indicates the 90%
conﬁdence. The values are marked with red circles for Model I, green
x marks for Model II, blue squares for Model III and aqua triangles
for Model IV. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83
The 6.67 keV line surface brightness with Model I but each error bar
in this plot indicates one sigma statistical error plus the systematic
error of the detector. The data are marked with a triangle mark for
Suzaku, a circle for XMM-Newton, an asterisk for Chandra and asquare for ASCA. The average ﬂux is shown in a solid horizontal line. 84
The light curve of the corrected 6.40 keV line surface brightness in
the Sgr B2 region. The XMM-Newton values are the average of MOS
and PN. In this plot, each error bar indicates the 90% conﬁdence limit. 85

LIST OF FIGURES

vii

6.6

The surface brightness maps of iron lines obtained with (a) Suzaku
XIS 2005, (b) XMM-Newton MOS and PN 2004, (c) Chandra ACIS-I
2000, and (d) ASCA SIS 1994. Each pixel size is 50 × 50 . . . . . . 87
6.7 The spectra for M0.66−0.02 obtained with Suzaku XIS FI sensors(a) and BI sensor (b) in 2005, XMM-Newton MOS (c) and PN (d) in
2004, (e) and (f) in 2001, Chandra ACIS-I (g) in 2000, and ASCA SIS(h) in 1994. The solid lines indicate the best-ﬁt models with Model I. 88
6.8 The 6.67 keV line surface brightness obtained with each observationin M 0.66−0.02. In this plot, each error bar indicates the 90% conﬁdence. The values are marked with red circles for Model I, green x
marks for Model II, blue squares for Model III and aqua triangles for
Model IV. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 90
6.9 The light curve of the 6.40 keV line surface brightness from M0.66−0.02(Sgr B2 cloud). In this plot, each error bar indicates the 90% conﬁdence limit. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 91
6.10 The spectra for G0.570−0.018 obtained with Suzaku XIS FI sensors(a) and BI sensor (b) in 2005, XMM-Newton MOS (c) and PN (d)in 2004, (e) and (f) in 2001, Chandra ACIS-I (g) in 2000, and ASCASIS (h) in 1994. The solid lines indicate the best-ﬁt models. . . . . . 92
6.11 The light curve of the 6.40 keV line surface brightness from G 0.570−0.018.
In this plot, each error bar indicates the 90% conﬁdence limit. . . . . 95
7.1

The time trend of the continuum ﬂuxes in M 0.66−0.02. XMM values
in 2001 and 2004 are the average of MOS and PN. In this plot, eacherror bar indicates the 90% conﬁdence limit. . . . . . . . . . . . . . . 98

List of Tables
3.1
3.2

Comparison of X-ray telescopes . . . . . . . . . . . . . . . . . . . . . 21
Comparison of X-ray CCDs . . . . . . . . . . . . . . . . . . . . . . . 30

4.1

Observation data list . . . . . . . . . . . . . . . . . . . . . . . . . . . 49

5.1

The best-ﬁt parameters for G 0.61+0.01 with the VPSHOCK modelplus two emission lines . . . . . . . . . . . . . . . . . . . . . . . . . . 64
The physical parameters of G 0.61+0.01 . . . . . . . . . . . . . . . . 66
The result of spectral ﬁttings of M 0.66−0.02 (Sgr B2 cloud) andM0.74−0.09 with a power-law and two Gaussian models. . . . . . . . 73

5.2
5.3
6.1
6.2
6.3
6.4
6.5
6.6
6.7
6.8
6.9

The Suzaku result of spectral ﬁtting for the Sgr B2 region . . . . . . 78
The XMM-Newton 2004 result of spectral ﬁtting for the Sgr B2 region 79
The XMM-Newton 2001 result of spectral ﬁtting for the Sgr B2 region 79
The XMM-Newton 2004 results of spectral ﬁtting for the Sgr B2
region with multiplied blank-sky backgrounds . . . . . . . . . . . . . 80
The XMM-Newton 2001 results of spectral ﬁtting for the Sgr B2
region with multiplied blank-sky backgrounds . . . . . . . . . . . . . 81
The Chandra 2000 and 2001 results of spectral ﬁtting for the Sgr B2
region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 82
The ASCA 1994 result of spectral ﬁtting for the Sgr B2 region . . . . 83
The result of spectral ﬁtting for M 0.66−0.02 with Model I . . . . . . 89
The result of spectral ﬁtting for G 0.570−0.018 with Model I plus one
Gaussian function . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 94

A.1
A.2
A.3
A.4

The result of spectral ﬁtting for M 0.66−0.02 with Model II . . . . . . 105
The result of spectral ﬁtting for M 0.66−0.02 with Model III . . . . . 106
The result of spectral ﬁtting for M 0.66−0.02 with Model IV . . . . . 107
The XMM-Newton 2004 results of spectral ﬁtting for M 0.66−0.02
with multiplied blank-sky backgrounds . . . . . . . . . . . . . . . . . 108
A.5 The XMM-Newton 2001 results of spectral ﬁtting for M 0.66−0.02
with multiplied blank-sky backgrounds . . . . . . . . . . . . . . . . . 109

B.1 The result of spectral ﬁtting for G 0.570−0.018 with Model II plus
one Gaussian function . . . . . . . . . . . . . . . . . . . . . . . . . . 111
B.2 The result of spectral ﬁtting for G 0.570−0.018 with Model III plus
one Gaussian function . . . . . . . . . . . . . . . . . . . . . . . . . . 112
viii

LIST OF TABLES

ix

B.3 The result of spectral ﬁtting for G 0.570−0.018 with Model IV plus
one Gaussian function . . . . . . . . . . . . . . . . . . . . . . . . . . 113
B.4 The XMM-Newton 2004 results of spectral ﬁtting for G 0.570−0.018
with multiplied blank-sky backgrounds . . . . . . . . . . . . . . . . . 114
B.5 The XMM-Newton 2001 results of spectral ﬁtting for G 0.570−0.018
with multiplied blank-sky backgrounds . . . . . . . . . . . . . . . . . 115

Chapter 1
Introduction
The Milky Way galaxy, our Galaxy, harbors the closest massive black hole at thecenter. The Galactic center (GC), crowded with various objects, is one of the mostinteresting targets in astronomy. The most puzzling problem is the central supermassive black hole with anomalously low luminosity.
One key to solve this problem was found with the past X-ray observations of the
Galactic molecular clouds whose temperature is too low to emit X-ray photons bythemselves. The discovery of neutral iron ﬂuorescense line emission with the ASCA
and Chandra X-ray satellites suggested that the possible irradiating source is thecentral suparmassive black hole Sgr A∗ and that 300 years ago the luminosity was
six order larger than that at the present. Murakami (2002) summarized these resultsof the past observations. However, this hypothesis still has two problems. Firstly,
the neutral iron line emission can be also produced by energetic cosmic particles.
Secondly, no time variability of the line emission had been found expected in thishypothesis.
In this thesis, we shed light on the largest Galactic molecular complex Sgr B2
with the strongest neutral iron line emission and aim to solve these problems. Noted
above, the GC region is the complex area with hot plasma, suparnova remnants(SNRs), H II regions and so on. The ﬁrst objective is to reveal the emission structure
all over the Sgr B2 region with the high performance of detecting a faint diﬀuse
emission of the Suzaku X-ray satellite. The second objective is to ﬁnd the time
variability of neutral iron line emission from molecular clouds by comparing thepast observations.
In chapter 2, we review the observational results of our Galactic center andthe Sgr B2 region, and X-ray line emission processes from neutral elements in acool molecular cloud. We use X-ray data obtained with CCD cameras onboard
four diﬀerent satellites. We introduce the charasteristics of the X-ray satellites,1

CHAPTER 1. INTRODUCTION

2

particularly their eﬃciencies and backgrounds, in chapter 3. We explain the data
reduction processes for the observations in chapter 4. In chapter 5, we report theresults of the Suzaku observation for the ﬁrst objective. In chapter 6, we report thediscovery of time variability of the neutral iron line by a comparison analysis using
the past observations with the precise systematic error study of each instrument.
Finally we discuss and summarize these results in chapter 7.

Chapter 2
Review
2.1

The Galactic Center Region

The GC region, shrouded by the intervening gas and dust, is actively exploredat radio, infrared, X-ray and gamma ray bands. The environment of the GC is socomplex that many components coexist within a central deep gravitational potential
well. The principal components are the supermassive black hole Sagittarius A∗ (SgrA∗ ), surrounding clusters of stars, ionized gas streamers, diﬀuse hot gas, supernova
remnants and molecular rings. Hereafter, we assume 8.5 kpc as the distance to the
GC (Reid et al. 1988) and 6×1022 cm−2 as the Galactic absorption hydrogen column
density.

2.1.1

Galactic Nucleus Sgr A∗

Sgr A∗ was found as a compact radio source (Balick & Brown 1974). The motionsof stars in the central part of our Galaxy indicate the presence of a supermassive
black hole with a mass of about 4 × 106 M , coincident with Sgr A∗ (Eisenhauer et
al. 2005; Ghez et al. 2005). In spite of the huge gravitational energy, the radiationis quite weak, particular in the X-ray band. The direct observation of Sgr A∗ inthe X-ray band was ﬁrst made with Chandra as shown in ﬁgure 2.1, which foundthat its quiescent luminosity is only 2 × 1033 erg s−1 (Baganoﬀ et al. 2001). Thisluminosity is far fainter than that of a typical active galactic nucleus (∼ 1041–46 erg
s−1 ). Flare activities with a duration time of 10 ks were also found (Baganoﬀ et al.
2001; Goldwurm et al. 2003), but their amplitudes were just a factor of ten.
Observations with other wavelengths have been extensively done. Recently, thelong-sought near-IR counterpart to Sgr A∗ was discovered (Genzel et al. 2003), andthe point-like detection of TeV gamma ray source coincident within 1 arcmin of Sgr
A∗ is reported (Aharonian et al. 2004; Tsuchiya et al. 2004; Kosack et al. 2004). The
3

CHAPTER 2. REVIEW

4

Figure 2.1: Adaptively smoothed 0.5–7.0 keV image of Sgr A∗ overlaid contours of
VLA 6 cm (Baganoﬀ et al. 2003).

CHAPTER 2. REVIEW

5

wideband spectrum of Sgr A∗ is displayed in ﬁgure 2.2 (Yusef-Zadeh et al. 2006). A
number of theoretical models are proposed to explain its faint quiescent luminosity(e.g., Liu et al. 2004; Yuan et al. 2004; Atoyan & Dermer 2004; Goldston et al. 2005;
Markoﬀ 2005; Eckart et al. 2006).

2.1.2

The Galactic Center Diﬀuse X-ray Emission

The Ginga X-ray satellite ﬁrst detected the iron K-shell lines near the Galactic
Center (Koyama et al. 1989,Yamauchi et al. 1990). The line energy is about 6.7 keVand the emission extends about 1◦ –2◦ over the GC region, which is called as the
Galactic Center Diﬀuse X-ray Emission (GCDX). ASCA resolved the iron K-shellline detected with Ginga into three distinct lines at 6.4, 6.7, and 6.9 keV (Koyama
et al. 1996). The 6.7 and 6.9 keV lines are likely due to He-like (Fe XXV) and hydrogenic (Fe XXVI) ions of iron. With respect to the 6.7 keV and 6.9 keV lines, collisional excitation in optically-thin, high-temperature plasma, and charge exchange
recombination are both considered. Suzaku obtained the global distribution of the
6.7 keV line emission shown in ﬁgure 2.3 and the GCDX spectrum with the highest
quality determining the line centroid energy of Fe XXV within 1 eV (Koyama et al.
2007d). The line center energies and widths of Fe XXV Kα and Fe XXVI Lyα favor acollisional excitation plasma for the origin of the GCDX. The electron temperature
determined from the line ﬂux ratio of Fe XXV Kα / Fe XXV Kβ is similar to the ionization temperature determined from that of Fe XXV Kα / Fe XXVI Lyα. Thus the
GCDX plasma is close to ionization equilibrium. The temperature of the thin thermal plasma determined from the GCDX spectrum is ∼ 10 keV. The total thermalenergy is estimated to be 1053–54 erg, or equivalent to 102 –103 supernovae explosions.
The temperature of ∼10 keV is higher than that bounded by the Galactic gravity,
hence the plasma should escape from the GC region. The time scale estimated by
the plasma size (1◦ or ∼ 150 pc at 8.5kpc) divided by the sound velocity of the∼ 10 keV plasma is ∼ 105 years. Therefore the energy of ∼ 1053–54 erg should be
supplied in the past ∼ 105 years; 1 suparnova every 100–1000 years is required in
the GC region.
The 6.4 keV line is likely to be Kα from the neutral and/or a blend of low charge
states of iron (Fe I). In fact, the 6.4 keV line is localized near molecular clouds suchas the radio complex Sgr B2 and the Radio Arc regions. Detailed descriptions about
6.4 keV are made in section 2.1.4.

CHAPTER 2. REVIEW

6

Figure 2.2: Wideband spectrum of Sgr A∗ . The spectrumof Sgr A∗ is based onmeasurements at 7 mm, 450 μm, 850 μm, 3.6 μm, 1.6 μm, and X-rays at 1 keV. The
quiescent X-ray and near-IR ﬂuxes are taken from Baganoﬀ et al. (2001), Genzel etal. (2003), and Ghez et al. (2005). The quiescent ﬂux is shown in blue color andthe ﬂux associated with ﬂaring activity including the X-ray ﬂare is shown in red
color. The hard gamma ray emission from the Galactic center at 1 TeV based on
HESS observations (Aharonian et al. 2004) and the high ﬂuxes from INTEGRAL
observations are not connected and are only shown by themselves in black. Further
descriptions are shown in Yusef-Zadeh et al. (2006).

CHAPTER 2. REVIEW

Galactic longitude (degree)

2.0
-1.0

0.0

1.0

6.7 keV

1.0

0.0

359.0

358.0

7

Galactic latitude (degree)
Figure 2.3: 6.7 keV narrow band image of the Galactic Center obtained with theSuzaku XIS.

CHAPTER 2. REVIEW

2.1.3

8

Supernova Remnants

With radio studies of our Galaxy, over two hundred SNRs are known (Green 2004).
Nine of them are located in the GC region (1◦ –2◦ ) and three of the SNRs associate Xray emissions. These X-ray emissions are seen as globular diﬀuse morphologies while
the radio ones are shell-like morphologies which are so called “Mixed-Morphology
SNRs” (Rho & Petre 1998). One interpretation of the MM SNR is that a hot(∼ 107 K) central ejecta is seen with X-ray and cool (≤ 106 K) thermal shell are
obscured with Galactic high density intersteller medium. The thermal plasma witha temperature of 1 keV exhibits a strong S XV Kα (2.45 keV) line mission. Therefore
the narrow line imaging of the S XV Kα line is the best probe to search a Galactic
SNR. The Suzaku satellite has observed many GC regions and composed the S XVKα line map (shown in ﬁgure 2.4). Some structures have been already reported as
SNR sources (Koyama et al. 2007c; Nobukawa et al. 2008; Mori et al. 2008).

2.1.4

Giant Molecular Complexes

Cold molecular gas is strongly concentrated in the inner few hundred parsecs of the
Galaxy. It seems to form a huge cloud complex with a size of Δl × Δb = 3◦ × 1◦ ,which is called as the Galactic center molecular cloud complex. In the Galactic
center molecular cloud complex, gas temperature is known to be relatively high
(30–60 K) and roughly uniform over several hundred parsecs. The total molecularmass is ≥ 2 × 107 M (Oka et al. 1998).
Figure 2.5 shows the spatial distribution of dense molecular clouds in the Galactic
center region with velocity-integrated CS J=1–0 emission in the range v=−200 to
200 km s−1 , which corresponds to the central molecular zone (CMZ) (Morris &Serabyn 1996). The molecular clouds, which consist of long and curved streams
mainly along the Galactic plane, are roughly associated with strong continuum
sources, the Sgr A radio arc complex, Sgr B, Sgr C, and Sgr D.
Although these molecular clouds are so cool that they cannot emit X-ray photons
themselves, 6.4 keV neutral iron line (Fe I) emissions were found (Koyama et al.
1996). Figure 2.6 is the Suzaku view of the GC region in the 6.4 keV band. The X-ray
emissions show clumpy structure associating some molecular cloud complexes. They
need external sources, such as X-rays or particles, to cause inner-shell ionizations.
Further details of Sgr B and the emission mechanism of neutral lines, the main
theme of this thesis, are described in the following sections.

CHAPTER 2. REVIEW

Galactic longitude (degree)

1.0
2.0
-1.0

0.0

1.0

Suzaku 36 observations
2.45 keV

0.0

359.0

358.0

9

Galactic latitude (degree)
Figure 2.4: 2.45 keV narrow band image of the Galactic Center obtained with theSuzaku XIS.

CHAPTER 2. REVIEW

10

Figure 2.5: Velocity-integrated CS J=1–0 emission in the Galactic center region(Tsuboi et al. 1999). The velocity integrated range is from −200 to 200 km s1 .

CHAPTER 2. REVIEW

Galactic longitude (degree)

2.0
-1.0

0.0

1.0

6.4 keV

1.0

0.0

359.0

358.0

11

Galactic latitude (degree)
Figure 2.6: 6.4 keV narrow band image of the Galactic Center with the Suzaku XIS.

CHAPTER 2. REVIEW

12

Figure 2.7: 2–10 keV band image of Sgr B2 obtained with the Chandra ACIS (Takagi
et al. 2002), overlaid on coutours of H II regions (Gaume et al. 1995). The solid circles
indicate X-ray point sources.

2.2

A Giant Molecular Complex — Sgr B

Sgr B is one of a number of Giant Molecular Complexes (GMCs) bound in relatively
tight orbits at a projected distance of ∼100 pc around Sgr A∗ . The total mass
is 0.4–1.0×107 M (Tsuboi et al. 1999). The Sgr B complex contains two bright
subregions, Sgr B1 and Sgr B2, which were named in radio continuum observations(Mehringer et al. 1995).

2.2.1

Sgr B2

Sgr B2 is the largest molecular cloud in the Galaxy and contains one of its largest
complexes of H II regions, being dominated by numerous ultracompact and hypercompact H II regions which inhabit three dense cores (labeled ”north”, ”main”, and”south”). These are, themselves, located inside a structure labeled ”the Envelope”(Gordon et al. 1993). The cores are small (0.5 pc), light (103 –104 M ; correspond-

CHAPTER 2. REVIEW

13

Figure 2.8: Adaptively smoothed Chandra image in the 6.15–6.55 keV band
(Murakami et al. 2001), overlaid on coutours of 13 CO line ﬂux (Sato et al. 2000)
ing to 5% of the Sgr B2 mass), and dense (106 –107 cm−3 ). On the other hand,the Envelope is massive (7.6 × 105 M ), and less dense (105 cm−3 ). The averagehydrogen density across Sgr B2 is 106 cm−3 . There have been many ultracompact(UC) H II sources identiﬁed in the Sgr B2 main alone (Gaume et al. 1995). Takagi
et al. (2002) have discovered many compact X-ray sources in the cloud center with
Chandra and the X-ray ﬂuxes and spectra indicate that these are likely high mass
young stellar objects.
Sgr B2 is also a strong 6.4 keV line source. ASCA ﬁrst discovered the 6.4 keV line
emission with a strong Fe K absorption edge (Koyama et al. 1996), and Chandra has
resolved its morphology (ﬁgure 2.8, Murakami et al. 2001). The equivalent widthof this line is about 1–2 keV and the depth of the absorption edge is equivalent
to NH ∼ 1024 cm−2 . The X-ray morphology is a concave shape at the GC side.
With these observational facts, they proposed that Sgr B2 is an X-ray reﬂection
nebula (XRN) irradiated by the GC source Sgr A∗ , which was thought to be X-ray
bright about 300 years ago, the light traveling time between Sgr B2 and Sgr A∗ .
On the other hand, the counter scenario is proposed in Yusef-Zadeh et al. (2007).
The nonthermal electrons cause the 6.4 keV line emission by colliding with ambient
matter to produce the inner-shell ionizations of iron atoms. The determination of
the 6.4 keV line origin is one of the main theme of this thesis.

CHAPTER 2. REVIEW

2.2.2

14

Sgr B1

Sgr B1 is comparatively less studied than Sgr B2. It is an H II region of opticaldepth much less than cm wavelengths, suggesting that it is an older, evolved structure no longer containing the dense, hot star forming UC H II regions found in Sgr
B2 (Mehringer et al. 1995). In the recent study with the Suzaku satellite, two diﬀuse
sources are found (Nobukawa et al. 2008). G 0.42−0.04 is a candidate for a supernova remnant which exhibits a strong S XV Kα emission line. M 0.51−0.10 has aprominent Fe I Kα emission line and a strongly absorbed continuum which is similar
to the X-ray spectrum of Sgr B2.

2.3

X-ray Emission Processes from Neutral Elements

Here we review the X-ray line emission processes from neutral elements, largely
owing to Inoue (1985) and Tatischeﬀ (2003). We introduce three processes to cause
inner-shell ionization; X-ray photoionization and collisional ionization by electronsand by ions.

2.3.1

Photoionizaion

Figure 2.9a shows a compilation of cross sections for K-shell ionization. The photoionization cross sections are steep functions of energy varying approximately as
EX−3 from ionization thresholds. Inoue (1985) operated a numerical simulation of
the ﬂuorescence iron line emission from neutral matter. The incident X-rays whichhave a spectrum with an index of 1.1 invoke the K-shell ionization of iron and produce the Fe K line and the continuum emission by Thomson scattering. Figure 2.10
shows the expected equivalent width of the Fe K line as a function of the thicknessof the matter NH . The line I represents the ﬂuorescence eﬃciency for the incident
X-rays. The eﬃciency ﬁsrt increases with the thickness NH of the gas but turns
to decreases due to the absorption of the line photons in a certain NH . The line
II indicates the equivalent width for the continuum including the absorption of theincident Xrays and the scattered component and the line III only considers only thescattered continuum.

2.3.2

Collisional Ionization by Electrons

The cross sections for the collisional ionization, shown in ﬁgure 2.9a, have a much
weaker energy dependence than that for photoionization, in particular, arise in theCHAPTER 2. REVIEW

15

Figure 2.9: (a) Cross sections for K-shell ionization of various atoms. Solid curvesindicate the calcurated result by electron impacts from the semi-empirical formula.
Dashed curves indicate cross sections for K-shell photoionization of N and Fe. (b)Stopping range of low-energy electrons in hydrogen. See Tatischeﬀ (2003) for detail
descriptions.

CHAPTER 2. REVIEW

16

Figure 2.10: Simulated equivalent width of ﬂuorescent Fe K line as a function of the
thickness NH of scattering matter. The simulation is made in the following cases;
I: primary source is directly observed, II: X-rays from primary source are absorbed
by a surrounding matter, III: primary source is totally hidden or scattering matter
is irradiated by an external source.

CHAPTER 2. REVIEW

17

Figure 2.11: X-ray emission produced by suprathermal electrons with the source
spectrum of ∝ Ee−2 interacting in a thick target composed of a neutral atomic gas ofthe solar abundance. The dashed line shows the eﬀect of photoelectric absorption,with a H column density of 4 × 1021 cm−2 .
MeV region caused by relativistic eﬀects (e.g. Hoﬀmann et al. 1979). The mosteﬃcient energy for an interaction of iron is 10–100 keV. Figure 2.9b shows the
stopping range of low-energy electrons in a neutral H gas. Suprathermal electrons
with an energy of 10–100 keV would stop in a neutral atomic gas of solar composition
with a thickness of 4 × 1021 cm−2 . Taking these basic conﬁgurations into account,Tatischeﬀ (2003) calculated the spectrum (shown in ﬁgure 2.11) for suprathermal
electrons with an index of 2. The equivalent width of Kα lines from O, Ne, Si and Fe
is 355, 107, 73 and 290 eV, respectively. A signiﬁcant line emission from elementsof relatively low atomic number could allow to distinguish X-ray line production
by suprathermal electrons from ﬂuorescence emission, since the strong decrease of
the ﬂuorescence yields with decreasing atomic number z is compensated by theapproximate z −4.3 dependence of the K-shell ionization cross section for electrons
while z −2.3 for photons (ﬁgure 2.9a).

CHAPTER 2. REVIEW

2.3.3

18

Collisional Ionization by Fast Ions

Tatischeﬀ (2003) also calculated an X-ray spectrum from accelerated ion interactions shown in ﬁgure 2.12. The energetic particles are assumed to be the same composition as the Galactic cosmic-ray ions. In this calculation,photoelectric absorption
in the ambient medium is not taken into account for both clarity and simplicity.
The narrow lines are due to K-shell vacancy production in the ambient atoms
by the fast ions. In comparison with the similar lines produced by photoionization
or by electron impact, the lines produced by heavy ion collisions could be shiftedby several tens of electron-volts, signiﬁcantly broadened and split up into several
components, owing to multiple simultaneous ionizations (Garcia et al. 1973), which
is not taken into account in the calculation shown in ﬁgure 2.12. For example,
the neutral Fe Kα line produced by 1.9 MeV/nucleon O impacts is blueshifted by∼50 eV in comparison with that produced by proton impacts, and has a width of∼100 eV (FWHM). The prominent, broad line features in ﬁgure 2.12 are due toatomic de-excitaions in the fast ions following electron capture by charge-transfer
reactions and collisional excitation.
Continuum X-rays are essentially produced by inverse bremsstrahlung (IB), whichis the kinematic inverse of normal bremsstrahlung, i.e. the radiation of an electron
at rest in the moving Coulomb ﬁeld of a fast ion. Thus, for nonrelativistic protonsof kinetic energy Ep interacting in stationary H, the inverse bremsstrahlung crosssection is almost identical to the bremsstrahlung cross section for electrons of energy(me /mp )Ep also interacting in stationary H. The recent study in shocked astrophysical plasmas shows that IB from accelerated ions is generally expected to be much
lower than bremsstrahlung from suprathermal electrons. There are the other two
continuum emission from the radiative recombinations of ambient electrons to theK-shell of the fast ions (RRC; Dogiel et al. 1998) and from the secondary, knocked-on
electrons that subsequently produce bremsstrahlung by interacting with the ambient atoms (SEB). These two emissions should generally be dominated by the IB
production.

CHAPTER 2. REVIEW

19

Figure 2.12: X-ray emission produced by fast ions with cosmic-ray source (CRS)
composition. IB: inverse bremsstrahlung; SEB: secondary electron bremsstrahlung;
RRC: radiative recombination continuum. See Tatischeﬀ (2003) for detailed descriptions.

Chapter 3
Instruments
In this study, the X-ray detector with the suﬃcient energy resolution to resolve the
Fe K lines and the spatial resolution to distinguish diﬀuse sources with a scale ofan arcsecond is necessary. A Charge-Coupled Device (CCD) is the most powerfulimaging device for the present X-ray astronomy with the energy resolution of ∼130 eV at 5.9 keV (FWHM), and suﬃcient detection eﬃciency in the 0.1–10 keVband. Thus, the X-ray CCD detectors are the most suitable devices for this study.
We use observation data from four X-ray satellites boarding the CCD cameras,
ASCA, Chandra, XMM-Newton and Suzaku. We introduce these space observatories in the point of the eﬃciency of their mirrors, the space environments, and thebackgrounds of the CCD cameras.

3.1

X-ray Satellites and their Mirrors

The characteristics of four satellites are diﬀerent among each other, which specify imaging or spectroscopy reﬂecting on the X-ray mirror system. Figure 3.1 isschematic views of the four X-ray observatories, ASCA, Chandra, XMM-Newton
and Suzaku. The typical properties of thier spacecrafts and mirrors are listed intable 3.1.

3.1.1

ASCA

The Advanced Satellite for Cosmology and Astrophysics, ASCA, is the Japanesefourth X-ray satellite lanunched on February 20th, 1993 (Tanaka et al. 1994). The
orbital heights at perigee and apogee are about 520 km and 620 km, respectively,
whereas the inclination is 31.◦ 1. The orbital period is approximately 96 minutes. The
satellite re-entered into the atmosphere on March 2, 2000. ASCA has four identical
X-ray telescope with nested thin foil mirrors (XRT; Serlemitsos et al. 1995). The
20

CHAPTER 3. INSTRUMENTS

21

Figure 3.1: Schematic pictures of the ASCA (upper left), Chandra (upper right),
XMM-Newton (lower left) and Suzaku satellites (lower right).

Table 3.1: Comparison of X-ray telescopes

Observatory
CountryYear
Orbit [km]Orbital period
Mirror system
Focal length [m]Eﬀective area@1 keV [cm2 ]Field of view
Angular Resolution

ASCA
Chandra
XMM-Newton
Japan
USA
ESA
1993–2000
1999–
1999–
520–620 16,000–133,000 7,000–114,000
96 mins
48 hrs64 hrs
XRT
HRMA
—
3.5
10
7.5
325 × 4800 × 11500 × 3
24
30
30


2.9
0. 5
15

Suzaku
Japan
2005–
568
96 mins
XRT
4.75
450 × 417
2

CHAPTER 3. INSTRUMENTS

22

focal-plane instruments consist of two kinds of detectors: two CCD cameras and two
gas counters. The former are called Solid-state Imaging Spectrometers (SISs: Burkeet al. 1994), and the latter are called Gas Imaging Spectrometers (GISs: Ohashi etal. 1996; Makishima et al. 1996).
The XRT of ASCA utilizes total X-ray reﬂection under the Wolter type I conﬁguration, which employs paraboloid and hyperboloid surfaces as the primary andthe secondary mirrors. The basic design of the telescope approximates the conventional opticals paired hyperbolic and parabolic surfaces with conical surfaces realised
as a set of closely nested concentric thin foils (120 per telescope) held together inquadrants by metal ﬁttings. The foils are made of aluminium and coated with 10
to 20 microns of acryllic and 60 nm of gold, which provided the reﬂecting surface.
Although the use of foils, which cannot be smoothed by polishing, reduces the sharpness of the image compared with conventional X-ray optics, it does provide a larger
collecting area and broader pass band. The angular resolution is so poor as ∼ 3
arcmin due to the shape errors of the mirror surfaces. The eﬀective area of XRT isshown in ﬁgure 3.2.

3.1.2

ChandraChandra (Weisskopf et al. 2003) was launched on the Space Shuttle Columbua on
July 23, 1999. The orbit of the satellite is elliptical with the perogee altitude of
16,000 km, and the apogee altitude of 133,000 km. The orbital period is about 64
hours, which is allowing uninterrupted observing intervals of more than 48 hours.
The most vivid characteristic of Chandra is the highest angular resolution in
all the previous X-ray satellites. The Chandra telescope, “High Resolution Mirror
Assembly (HRMA)”, consists of four mirrors of Wolter type I geometry formedby glass, and achieved the angular resolution of 0. 5. The mirrors are coated withiridium on a binding layer of chromium. In order to keep high accuracy of the mirror
surface, the glass is very thick and heavy with the total mass of about 1,500 kg.
Chandra has two X-ray detectors on the focal plane: the Advanced CCD Imaging
Spectrometer (ACIS) and the High Resolution Camera (HRC). The HRC is amultiple-channel plate optimized for imaging without the ability of spectroscopy.
Chandra can operate two types of grating: the High-Energy Transmission Grating(HETG) and the Low-Energy Transmission Grating (LETG).
The HRMA, shown schematically in ﬁgure 3.3, contains the nested mirrors, center, forward and aft aperture plates, baﬄes, inner and outer cylinders, mounts, preand post-collimators, ﬁducial light transfer components, mirror support sleeves, for-

CHAPTER 3. INSTRUMENTS

23

Figure
3.2:
ASCA
XRT
eﬀectivearea(onetelescope)for
diﬀerent
oﬀ-axis angle(θ) andazimuthal
angle
(φ).
See
http://heasarc.gsfc.nasa.gov/docs/asca/asca nra06/appendix e/appendix e.html.

CHAPTER 3. INSTRUMENTS

24

Figure 3.3: The four nested HRMA mirror pairs and associated structures. See
http://cxc.harvard.edu/proposer/POG/html/HRMA.html.

CHAPTER 3. INSTRUMENTS

25

ward and aft contamination covers, ﬂux contamination monitors, and thermal control hardware. The outer mirror pair is number 1, and, progressing inwards, 3, 4,
and 6. The original design had six mirror pairs; numbers 2 and 5 were eliminated.
The pair diameters range from about 0.65 to 1.23 meters. The distance from thecenter of the Central Aperture Plate (CAP) separating the paraboloid and hyperboloid mirrors to the HRMA focus is 10.0548 meters, with each mirror pair varying
slightly about this value. Note that this distance is close to, but not exactly, thefocal length. An annular on-axis beam enters each mirror pair, is reﬂected fromparaboloids and hyperboloids and exits to converge to a focus.
Figure 3.4 shows the HRMA eﬀective areas with focal plane detectors as a function of X-ray energy. The structure near 2 keV is due to the iridium M-edge.

3.1.3

XMM-Newton

XMM-Newton is a ESA’s X-ray observatory launched on December 10th, 1999. It
carries two distinct types of telescopes: three X-ray telescopes with diﬀerent X-ray
detectoers in their foci, and a 30-cm optical/UV telescope with a microchannel-platepreampliﬁed CCD detector in its focal plane. XMM-Newton provides three scientiﬁc
instruments: European Photon Imaging Camera (EPIC) comprised of three X-ray
CCD detectors, Reﬂection Grating Spectrometer (RGS) with high resolution X-ray
spectroscopy and Optical Monitor (OM).
Each of the three X-ray telescopes on board XMM-Newton consists of 58 Wolter
I grazing-incidence mirrors which are nested in a coaxial and cofocal conﬁguration.
The design of the optics was driven by the requirement of obtaining the highest
possible eﬀective area over a wide range of energies, with particular emphasis in
the region around 7 keV. Thus, the mirror system had to utilize a very shallow
grazing angle of 30 in order to provide suﬃcient reﬂectivity at high energies. The
focal length of the telescopes is 7.5 meters and the diameter of the largest mirrors is
70 cm, to be compatible with the shroud of the launcher. Each telescope includes,
apart from the mirror modules, baﬄes for visible and X-ray stray-light suppression
and an electron deﬂector for diverting soft electrons. Figure 3.5 shows the eﬀective
area of all XMM-Newton X-ray telescopes, EPIC and RGS. The XMM-Newton
mirros are most eﬃcient in the energy range from 0.1 to 10 keV, with a maximum
at about 5 keV and a pronounced edge near 2 keV (Au M edge).

CHAPTER 3. INSTRUMENTS

26

Figure 3.4:
The HRMA/ACIS and HRMA/HRC eﬀective areas
sus X-ray energy in linear (top) and log (bottom) scale.
http://cxc.harvard.edu/proposer/POG/html/HRMA.html for details.

verSeeCHAPTER 3. INSTRUMENTS

27

Figure3.5:
Thenet
eﬀective
area
of
all
XMMNewton
X-ray
telescopes,
EPIC
and
RGS.
See
http://xmm.vilspa.esa.es/external/xmm user support/documentation/uhb 2.5/node20.html.

CHAPTER 3. INSTRUMENTS

3.1.4

28

SuzakuSuzaku is the ﬁfth Japanese X-ray astronomy satellite. It was developed underJapan-US international collaboration and was launched on July 10, 2005, fromJAXA Uchinoura Space Center. The scientiﬁc payload of Suzaku initially consisted of three distinct co-aligned scientiﬁc instruments. Suzaku is in many ways
similar to ASCA in terms of orbit, pointing, and tracking capabilities. As a result,the operational constraints for Suzaku are also similar to those of ASCA. Suzaku isplaced in a near-circular orbit with an apogee of 568 km, an inclination of 31.◦9, andan orbital period of about 96 minutes. The maximum slew rate of the spacecraftis 6 degrees/min, and settling to the ﬁnal attitude takes ∼ 10 minutes, using thestar trackers. The normal mode of operations will have the spacecraft pointing ina single direction for at least 1/4 day. With this constraint, most targets will beocculted by the Earth for about one third of each orbit, but some objects near theorbital poles can be observed nearly continuously. The observing eﬃciency of thesatellite as measured after a year of operation is about 45%.
There are four X-ray sensitive imaging CCD cameras (X-ray Imaging Spectrometers;
XISs), three front-illuminated (FI; energy range 0.4–12 keV) and one back-illuminated(BI; energy range 0.2–12 keV), capable of moderate energy resolution. Each XIS is
located in the focal plane of a dedicated X-ray telescope (XRT). The second instrument is the non-imaging, collimated Hard X-ray Detector (HXD), which extends the
bandpass of the observatory to much higher energies with its 10–600 keV bandpass.
Since, a thermal short between the helium and neon tanks resulted in the liquid
helium coolant venting to space, which was happened on August 8, 2005, the highresolution spectrometer, the X-Ray Spectrometer (XRS), is no longer operational.
The Suzaku XRTs consist of closely nested thin-foil reﬂectors, reﬂecting X-ray atsmall grazing angles. An XRT is a cylindrical structure, having the following layered
components: 1. a thermal shield at the entrance aperture to help maintain a uniform
temperature; 2. a pre-collimator mounted on metal rings for stray light elimination;
3. a primary stage for the ﬁrst X-ray reﬂection; 4. a secondary stage for the second
X-ray reﬂection; 5. a base ring for structural integrity and interface with the EOBof the spacecraft. All these components, except the base rings, are constructed in 90◦

segments. Four of these quadrants are coupled together by interconnect-couplers
and also by the top and base rings. The telescope housings are made of aluminumfor an optimal strength to mass ratio. Each reﬂector consists of a substrate also
made of aluminum and an epoxy layer that couples the reﬂecting gold surface to thesubstrate.

CHAPTER 3. INSTRUMENTS

29

Figure 3.6: Eﬀective area of one XRT + XIS system, for both the FI and BI chips.
(No contamination included)
Including the alignment bars, collimating pieces, screws and washers, couplers,
retaining plates, housing panels and rings, each XRT-I consists of over 4112 mechanically separated parts. In total, nearly 7000 qualiﬁed reﬂectors were used andover 1 million cm2 of gold surface was coated.
Each of the co-aligned XRTs features an X-ray mirror with an angular resolutionof ∼ 2 (HPD). Figure 3.6 shows the total eﬀective area of the XIS+XRT, which
includes features due to the elemental composition of the XIS and XRT. K-shell absorption edges from the oxygen (0.54 keV) and aluminum (1.56 keV) in the blocking
ﬁlters are present, as well as a number of weak M-shell features between 2–3 keVarising from the gold in the XRT.

3.2

CCD cameras

Table 3.2 shows typical parameters of the CCD cameras onboard ASCA, Chandra,
XMM-Newton and Suzaku.
Every CCD initially had a good energy resolution of ∼ 130 eV. However, Charge

CHAPTER 3. INSTRUMENTS

30

Table 3.2: Comparison of X-ray CCDs

Observatory
CountryYear
Detector
No. of chips
No. of pixels/chip
Pixel size [μm]Depletion depth [μm]Readout noise [e− ]ASCA
ChandraJapanUSA
1993–2000
1999–
SIS
ACIS
8
8 (FI) / 2 (BI)420 × 420
1024 × 1024
27
24
35
70
4–6
2

XMM-Newton
Suzaku
ESA
Japan
1999–
2005–
MOS
PN
XIS
7
12
3 (FI) / 1 (BI)600 × 600 378 × 64
1024 × 1024
40
150
24
35
280
70(FI) / 45(BI)——2transfer eﬃciency (CTE), a measure of how eﬀectively the CCD transfers the charge
from one pixel to the next, gradually decreases due to the radiation damage in space.
The decrease of the CTE, or in other words, the increase of the charge transfer
ineﬃciency (CTI), causes the loss of charge during the charge transfer. This reduces
the apparent X-ray energy and hence brings about a reduction of gain. Although
the gain can be recovered by the measurement of the CTI, the energy resolution
remains degraded.
Another problem in space is a background. The in-ﬂight background spectrumof a CCD depends on the readout electronics in the soft band and the particle
environment in the high energy bands. The particle environments depend on theorbits of the satellites. We describe the details in the following.

3.2.1

SIS

The SIS (Burke et al. 1994) is the ﬁrst satellite-bourne X-ray experiment that utilizes
CCDs in the photon counting mode. Designed focal plane and conﬁguration of the
SIS are shown in ﬁgure 3.8.
Energy resolution
Figure 3.9 shows the energy resolution of SIS and GIS. The energy resolution of
SIS is degraded because of the increase of the CTI. Additinally, the residual dark
distribution (RDD) eﬀect caused additional degradation of the spectral response.
CCDs always have read-out noise: pixels without photon or particle events will havea mean of ”dark level” and some distribution around this mean level. This is usually

CHAPTER 3. INSTRUMENTS

31

Figure 3.7: Cross section of SIS (Gendreau 1995).

Figure 3.8: Left:Four individual CCDs are combined in a single package to make asingle SIS focal plane (Gendreau 1995). Right:Conﬁguration of the two SIS on the
ASCA focal plane.

CHAPTER 3. INSTRUMENTS

32

Figure 3.9: Energy resolution of the GIS and the SIS. For the SIS, measured or estimated energy resolutions of SIS1 in the 1CCD mode are plotted for at launch and 1, 2, 3, 4 and 5 years after the launch. SIS0 resolution is degrading somewhat more slowly than that of SIS1. In 2 and 4
CCD modes, the RDD eﬀect further degrades the resolution somewhat. See
http://heasarc.gsfc.nasa.gov/docs/asca/asca nra06/appendix e/node9.html for details.
treated as the Gaussian distribution. In the case of ASCA, the on-board processor
maintains a coarse map of the dark levels (each chip is divided into 16 by 16 regions
by default) and subtracts these from individual pixel values. The dark frame map iscalculated assuming a symmetric noise around the mean dark level. Unfortunately,
the cumulative eﬀects of the radiation damage have created a population of active
pixels which skews this distribution, leading to imperfect dark level subtraction
on-board. This eﬀect is called the Residual Dark Distribution (RDD). The RDD
depends on the clocking mode, and therefore, it is most pronounced in 4-CCD mode
(an integration time of 16 s), and less important in 1-CCD mode (an integration
time of 4 s).

CHAPTER 3. INSTRUMENTS

33

Figure 3.10: The SIS spectrum for the dark earth data (Gendreau 1995).
Background
The SIS background is quite stable and its count rate correlates with the local cutoﬀ
rigidity value, indicating that the background is mainly induced by high energy
particles, except at the South Atlantic Anomaly (SAA). Figure 3.10 shows the SIS
spectrum of the dark earth data. Since ASCA took an almost circular orbit withan average altitude of 520–620 km, inside the Earth’s magnetosphere, the SIS was
little aﬀected by soft protons, which seriously damaged ACIS and EPIC describedin the following sections.

3.2.2

ACIS

The Advanced CCD Imaging Spectrometer (ACIS) is comprised of ten CCDs with
1024 × 1024 pixels; 2 × 2 array of ACIS-I(I0–I3) for imaging and 1 × 6 array of ACISS (S0–S5) for imaging and grating spectroscopy. The pixel size is 24 μm and theangler resplution with the HRMA is 0.5 arcmin. Two CCDs (S1 and S3) are backilluminated (BI) and the others are front-illuminated (FI). The layout of CCD chips

CHAPTER 3. INSTRUMENTS

34

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

I0

I1

w203c4r
0

w193c2
1
x

I2

}

I3

w158c4r

w215c2r

2

3

ACIS-I(aimpoint on I3 = (949, 978))330 pixels = 163"
S0

S1

S2

w168c4r
4

w140c4r
5

w182c4r
6

S3+
w134c4r
7

S4S5

w457c48

w201c3r
9}

ACIS-S(aimpoint on S3 = (235, 497))+ΔY
18 pixels = 8".8

Image Region

Bottom.

columnthree

two

BI chip indicator

one

Top

node zero

+Z

.

row

TargetOffsetCoordinates+Z
Pointing+ΔZ

Coordinates-Z

Pixel (0,0)+Y

Sim Motion

Frame StoreCCD Key

NodeDefinitions

Row/Column
Definition

Coordinate
Orientations

Figure 3.11: A schematic drawing of the ACIS focal plane; insight to the terminology
is given in the lower left. Note the nominal aimpoints: on S3 (the ‘+’) and on I3(the ‘x’) (Weisskopf et al. 2003).
is shown in ﬁgure 3.11.
Energy resolution
After the launch, the CCD chips were seriously damaged with low energy protons
reﬂected by the X-ray telescope onto the focal plane during radiation belt passage.
Low energy protons deposited thier energy in the buried channels at the HRMA
side of the FI chips and increased the CTI. BI chips, on the other hand, were
not impacted although they exhibited poorer resolution in pre-launch. Since thechannels are near the gates and the gates face in the direction opposite to the HRMA,
the energy resolution for the BI chips remains nearly at pre-launch values. In thepresent operation, the ACIS is not left at the focal position during the radiation
belt passage, and no serious degradation in performance has been encountered. By
decreasing the CCD temperature from −90◦ C (before damaged) to −120◦ C, the CTI

CHAPTER 3. INSTRUMENTS

0.5

35

0.1

Au M

Ni K

0.05

Al K

0.02

counts/sec/keV

0.2

Si K

0.5

1
2
channel energy (keV)

5

Figure 3.12: The ACIS background spectra of I0 (black) and S3 (red) extracted
from the Chandra blank-sky database.
eﬀect is reduced as possible.
BackgroundFigure 3.12 shows the ACIS background spectra measured in a stowed position(moving ACIS further away from the external calibration source than in the HRC-I
position). In general, the background counting rates are stable during an observation. Furthermore, the spectral shape of the non-X-ray background has beenremarkably constant during 2000–2005 for both BI and FI chips, even though theoverall background rate showed secular changes by 20–30%. When the quiescentbackground spectra from diﬀerent observations are normalized to the same rate in
the 10–12 keV interval, they match each other to within ±3% across the whole
Chandra energy band. Occasionally, however, there are large variations (ﬂares).
Figure 3.13 shows the frequency of such variations when compared to the quiescentbackground. An average fraction of the exposure aﬀected by ﬂares above the ﬁltering threshold used for the blank-sky background datasets (a factor of 1.2 above theCHAPTER 3. INSTRUMENTS

36

nominal rate) was about 6% for the FI chips and up to 1/3 for the BI chips during
the ﬁrst few years of the mission.

3.2.3

EPIC

The European Photon Imaging Camera (EPIC) consists of two MOS (Metal Oxide
Semiconductor) CCD cameras (Turner et al. 2001) and one pn-type CCD camera(PN) (Strüder et al. 2001). The PN chip array is slightly oﬀset with respect to theoptical axis of its X-ray telescope so that the nominal, on-axis observing position
does not fall on the central chip boundary. This ensures that ≥ 90% of the energyof an on-axis point source are collected on one PN chip. The two EPIC MOS
cameras are rotated by 90◦ with respect to each other. The dead spaces between the
MOS chips are not gaps, but unusable areas due to detector edges. The MOS chips
physically overlap each other, the central one being located slightly behind the onesin the outer ring.
Energy resolution
The energy resolution of EPIC, shown in ﬁgure 3.15, is monitored by the onboard
calibration source, which produces strong Al Kα and Mn Kα lines. The resolutionof the MOS cameras degraded signiﬁcantly with time up to the epoch where thecameras were cooled from −100◦ C to −120◦ C. After this epoch the degradation has
been minima. In the same description of the ACIS CCD, the PN CCDs are lessaﬀected since they are back-illumination type.
Background
XMM-Newton takes a highly elliptical orbit with an apogee of about 115,000 km anda perigee of about 6000 km and thus mostly ﬂy outside of the Earth’s magnetosphere.
This is thought to be responsible for the violent variability of the background of the
EPIC cameras as large as two orders of magnitudes. This variation is caused byprotons scattered through the mirror system.
The background of the EPIC cameras consists of mainly two components: thecosmic X-ray background (CXB) and the instrumental non-Xray background (NXB).
The latter component is mainly due to the interactions of charged particles with thestructure surrounding the detectors and the detectors themselves. This component
can be further divided into two components: a transient component and a quiescentcomponent. Since spatial and temporal characteristics are diﬀerent between thesecomponents, it is important to study them separately.

CHAPTER 3. INSTRUMENTS

37

Figure 3.13: An example of the ACIS background counting rate versus
time: S3 (the BI chip; top curve) and I2 (the FI chip; bottom curve).
These are for the standard grades and the band from 0.3–10 keV. See
http://cxc.harvard.edu/proposer/POG/html/ACIS.html for details.

CHAPTER 3. INSTRUMENTS

38

Figure 3.14: A sketch of the ﬁeld of view of the two types of EPIC camera; MOS (left) and PN (right).
The shaded circle depicts a 30 arcmin
diameter area.
For the alignment of the diﬀerent cameras with respect
to each other in the XMM-Newton focal plane refer to the text.
See
http://xmm.vilspa.esa.es/external/xmm user support/documentation/uhb 2.5/node24.html
for details.

CHAPTER 3. INSTRUMENTS

39

Figure 3.15: Energy resolution of MOS1-CCD1 (red asterisk), MOS2-CCD1 (green
square) and PN-CCD4 (black triangle). The drop after rev. 532 is due to the coolingof the cameras. See http://xmm.vilspa.esa.es/docs/documents/CAL-TN-0018.pdf
for details.
The transient component is thought to be induced by soft protons with energies
less than a few hundred keV. Since the soft protons are scattered through the mirror
system, they are vignetted. The spectra of soft proton ﬂares are variable and noclear correlation has been found between intensity and spectral shape. Therefore,
observation intervals aﬀected by this component have to be screened oﬀ.
The quiescenet component is induced by high-energy cosmic rays with energies larger than several hundred MeV which penetrate the satellite. The quiescentbackground consists of a continuum component and lines. While the continuum
component is induced by high energy particles directly in the detector, the lines
correspond to the ﬂuorescent lines of the surrounding structures. Whereas the continuum component is spatially ﬂat, the line component spatially changes dependingon the detector structure. Figure 3.16 shows background spectra for the MOS and
the PN cameras and images of the typical intrinsic lines. Further details are described in Carter & Read (2007).

3.2.4

XIS

The four Suzaku X-ray Imaging Spectrometers (XISs) are named XIS0, 1, 2 and 3,
and each CCD is located in the focal plane of an X-ray Telescope; those telescopes
are known respectively as XRT-I0, XRT-I1, XRT-I2, and XRT-I3 (ﬁgure 3.17). One

CHAPTER 3. INSTRUMENTS

40

Figure 3.16: Upper: Background spectra for the MOS1 camera (left) and the PN
camera (right) during an observation with the ﬁlter wheel in the closed position.
Lower: Images of the MOS camera of Al K line at 1.48 keV (left) and the PN of Cu
K at 8.04 keV (right) (Carter & Read 2007).

CHAPTER 3. INSTRUMENTS

41

Figure 3.17: Layout of the XRTs on the Suzaku spacecraft.
of the XISs, XIS1, uses a back illuminated CCDs, while the other three use front
illuminated CCDs.
A single XIS CCD chip consists of four segments (marked A, B, C and D inﬁgure 3.18). The Analog Electronics (AE) drives the CCD and processes its data.
Charge clouds produced in the exposure area in the CCD are transferred to the
Frame Store Area (FSA) after the exposure according to the clocks supplied by the
AE. To minimize the thermal noise, the sensors need to be kept at ∼ −90◦ C duringobservations. This is accomplished by thermo-electric coolers (TECs), controlled
by TEC Control Electronics, or TCE. The AE and TCE are located in the same
housing. The digital electronics system for the XISs consists of two Pixel Processing
Units (PPU) and one Main Processing Unit (MPU). The PPUs access the raw CCD
data in the Pixel RAM, carry out event detection, and send event data to the MPU.
The MPU edits and packets the event data, and sends them to the satellite’s main
digital processor.
To reduce contamination of the X-ray signal by optical and UV light, each XIS

CHAPTER 3. INSTRUMENTS

Figure 3.18: One XIS instrument.

42

CHAPTER 3. INSTRUMENTS

43

has an Optical Blocking Filter (OBF) located in front of it. The OBF is made ofpolyimide with a thickness of 1000 Å, coated with a total of 1200 Å of aluminum(400 Å on one side and 800 Å on the other side). To facilitate the in-ﬂight
calibration of the XISs, each CCD sensor has two 55 Fe calibration sources. Thesesources are located on the side wall of the housing and are collimated in order toilluminate two corners of the CCD. They can easily be seen in two corners of each
CCD. A small number of these X-rays scatter onto the entire CCD.
After the launch of Suzaku, the four XISs had been operated successfully formore than a year. However, a large amount of leaked charge suddenly appeared in
one of the sensors, XIS2, in November 2006. Because the leak of charge was unable
to be suppressed, we stopped using XIS2 for scientiﬁc observations.
Energy Resolution
Performance of the XIS is gradually degrading since the launch of Suzaku due to theradiation damage. Figure 3.19 shows the time history of the energy resolution for
Mn Kα from the 55 Fe calibration sources for XIS. Energy resolution at Sept. 2006
measured by the 55 Fe is about 200 eV (FWHM) at 5.9 keV, while it was ∼130–135
eV just after the launch.
In principle, a precision charge injection capability can mitigate the eﬀects of inﬂight radiation damage (Bautz et al. 2004). A suﬃcient quantity of charge injectedperiodically during the array readout process (spaced-row charge injection; SCI) willtend to ﬁll radiation-induced traps, and in so doing reduce the eﬀects of these traps
on the devices charge transfer eﬃciency (CTE). Because the traps need to be ﬁlled
quite often, we inject charge every 54 rows in every exposure. The charge injectedrows are not usable to detect X-ray photons, and become dead area.
With this SCI technique, XIS achieve to keep the high energy resolution (∼ 140–
150 eV at 5.9 keV) in a long period. However, Suzaku data used in this thesis aretaken in an early time from the launch and not operated with SCI.
BackgroundAll four XISs have low backgrounds, due to a combination of the Suzaku orbit andthe instrumental design. In the XIS, the background originates from the cosmic Xray background (CXB) combined with charged particles (NXB). Currently, ﬂickering
pixels are a negligible component of the background.
There are also ﬂuorescence features arising from the calibration source as wellas material in the XIS and the XRTs. In XIS0, the Mn line due to the scatteredCHAPTER 3. INSTRUMENTS

44

Figure 3.19: The time history of the energy resolution for Mn Kα from the 55 Fe
calibration sources for XIS.

CHAPTER 3. INSTRUMENTS

45

Figure 3.20: The XIS background rate for each of the four XIS detectors, withprominent ﬂuorescent lines marked. These spectra are based on the observations of
the dark Earth for ∼ 800ks.

CHAPTER 3. INSTRUMENTS

46

X-rays from the calibration sources contaminates the whole chip. The other lines
are ﬂuorescent lines from the material used for the sensor. The NXB is not uniform
over the chip. Its ﬂux tends to be slightly higher at larger ACTY as shown in ﬁgure
3.21. This is because some fraction of the NXB is produced in the frame-store region.
The fraction may be diﬀerent between the ﬂuorescent lines and the continuum. This
causes slight diﬀerence in the ACTY dependence of the NXB.
The total ﬂux of the NXB depends on the cut-oﬀ rigidity (COR) as shown inﬁgure 3.22. This may be natural as the NXB is produced by the charged particles,
whose ﬂux is higher at lower COR. This means that, when we use the NXB databaseto subtract the background, we need to select the NXB with an appropriate rangeof the COR, which match with the COR distribution of the on-source observations.
For this purpose, the NXB database available from the Suzaku web page is sortedwith the COR.
The count rate of the PIN upper discriminator (PIN-UD) in the HXD also correlates the ﬂux of the NXB. Recent studies show the PIN-UD provides slightly better
reproducibility of the NXB. The reproducibility of the NXB in the 5–12 keV bandis evaluated to be 3–4% of the NXB, when we use the PIN-UD as the NXB sorting
parameter (Tawa et al. 2008).

CHAPTER 3. INSTRUMENTS

Figure 3.21: ACTY dependence of the NXB for XIS1 and XIS2.

47

CHAPTER 3. INSTRUMENTS

48

Figure 3.22: Cut-oﬀ rigidity dependence of the NXB (average intensity in 5–10 keV)for each sensor.

Chapter 4
Observations
We use four deep observations and two short surveys listed in table 4.1. For imaging
analysis, we only use the deep obsevation data.

4.1

The Suzaku Observation

4.1.1

Data Collection

The Sgr B region was observed with the XIS on 10-12 October 2005. The XIS
observation was made with the normal mode. The eﬀective exposure time after
removing the epoch of low Earth elevation angle ( ELV ≤5◦ ) and the South Atlantic
Anomaly was about 89 ksec.

4.1.2

The Gain Tuning

In a quick look of the spectrum, we found strong lines at ∼6.7 keV and ∼6.4 keV, inthe whole CCD imaging area (IA). These lines may be due to the largely extended

Table 4.1: Observation data list
Observatory

Obs. ID

ASCA
ASCA
Chandra
XMM-Newton
Chandra
XMM-Newton
Suzaku52006000
52006001
944
0112971501
2280
0203930101
100037060

Date(yyyy/mm/dd)
1994/09/22
1994/09/24
2000/03/29
2001/04/01
2001/07/16
2004/09/04
2005/10/10

49

Exposure Time(sec)
58,39521,50098,62912,64910,42342,282
89,072

CHAPTER 4. OBSERVATIONS

50

Galactic center diﬀuse X-ray emission (GCDX). These lines are most likely Kαlines of Fe XXV (6.67 keV) and Fe I (6.40 keV). Using the center energies of the twostrong lines, we made ﬁne correction of the CTI (Charge Transfer Ineﬃciency), andﬁne gain tuning in XIS-to-XIS and segment-to-segment levels (relative gain tuning).
Then the absolute gain tuning is made using the 55 Fe calibration sources irradiating
the CCD corners. Details of this procedure and its usefulness are demonstrated in
Koyama et al. (2007d).

4.1.3

The Position Tuning

After the CTI correction and ﬁne gain tuning, we added all the XIS data and made acomposite image of the Sgr B region in the 2–10 keV band (ﬁgure 4.1). The diﬀuse
enhancement in the northwest corresponds to the Sgr B2 complex. In addition,
we found two point sources at the southeast edge of the XIS ﬁeld of view (FOV).
We made radial proﬁles of these two sources and determined the peak positions tobe (l, b) = (0.◦ 5762, − 0.◦ 1736) and (0.◦6626, − 0.◦ 2225) in the nominal Suzaku
coordinate. We searched for the Chandra Galactic center survey map, and foundpossible counterparts, CXO J174741−283213 and CXO J174804.8−282919 (Munoet al. 2003). The Galactic coordinates of these sources are (l,b) = (0.◦ 5762,−0.◦ 1796)and (0.◦6625, −0.◦ 2289). Therefore the Suzaku nominal coordinate is systematicallyshifted by (Δl, Δb) = (0.◦0001, 0.◦ 0062) from the Chandra coordinate. Since theaspect solution of Chandra is very accurate, within sub-arcsec, we made ﬁne-tunedof the Suzaku position by shifting (−0.◦ 0001, − 0.◦ 0062) in the (l, b) coordinate.
Hereafter we use this re-registered coordinate.

4.2

The XMM-Newton Observations

4.2.1

Data Collection

The X-ray data have been obtained with the EPIC cameras onboard XMM-Newton
on 1 April 2001 and 4 September 2004. The observations were operated in theextended full frame mode. The data have been analysed using the Science Analysis
Software (sas 7.1.0). Event ﬁles for both the PN and the MOS detectors have been
produced using the epproc and emproc tasks of sas, respectively. The event ﬁles werescreened for the high particle background periods. The background light curves andthe screening criteria are shown in ﬁgure 4.2.
In our analysis, we have dealt only with events corresponding to patterns 0–4
for the PN and 0–12 for the MOS instruments.

CHAPTER 4. OBSERVATIONS

51

15m00.0s

XIS FOV

Declination-28d20m00.0s

25m00.0s

PSRC 130m00.0s

PSRC 235m00.0s

GC
30.0s

17h48m00.0s

30.0s

47m00.0s

46m30.0s

Right ascension
Figure 4.1: The Suzaku image of the Sgr B region in the 2–10 keV band. All thefour XIS data are co-added. The dotted square is the XIS ﬁeld of view (FOV).

CHAPTER 4. OBSERVATIONS

52

MOS1 (>10 keV)

40

Counts per 50 sec

0

20

400
200
0

Counts per 50 sec

600

60

MOS1(>10 keV)

1.02475×108 1.0248×108 1.02485×108 1.0249×108 1.02495×108
TIME s

2.1066×108 2.1067×108 2.1068×108 2.1069×108
TIME sMOS2 (>10 keV)

40

Counts per 50 sec

0

20

400
200
0

Counts per 50 sec

600

60

MOS2(>10 keV)

2.107×108

1.02475×108 1.0248×108 1.02485×108 1.0249×108 1.02495×108
TIME s

2.1066×108 2.1067×108 2.1068×108 2.1069×108
TIME s

PN (>10 keV)

200

Counts per 50 sec

0

100

1000
0

500

Counts per 50 sec

1500

300

PN(>10 keV)

2.107×108

1.02475×108 1.0248×108 1.02485×108 1.0249×108 1.02495×108
TIME s

2.1066×108 2.1067×108 2.1068×108 2.1069×108
TIME s

2.107×108

Figure 4.2: Background light curves of 2001 (left) and 2004 (right) observations
with the MOS and the PN detectors. Events with over 10 keV and pattern 0 are
accumulated. Holizontal lines indicate the screening criteria (0.4 counts s−1 for MOS
and 1.0 counts s−1 for PN).

CHAPTER 4. OBSERVATIONS

53

Figure 4.3: The XMM-Newton image of the Sgr B region in the 0.3–10 keV band.
All MOS and PN data in 2004 are co-added. Circles indicate point sources identiﬁed
with the Chandra point source catalogue (Muno et al. 2003).

4.2.2

The Position Tuning

For the observation in 2004, we make the total band images (0.3–10 keV) of eachdetector and detect point sources; 93 from MOS1, 97 from MOS2, 111 from PN.
There are 27 sources detected with all the detectors within positional uncertainties
of 10 arcsec which is comparable to the size of PSF of EPIC at 6 keV, and 19 sources
have possible counterparts with the point sources catalogued with Chandra (Munoet al. 2003). The composite image and distribution of identiﬁed point sources areshown in ﬁgure 4.3. Since the mean of their diﬀerences againt the Chandra positionsis (−0.◦ 0006, 0.◦ 0003) in the celestial coordinate, we tune the XMM-Newton frameby shifting (ΔRA, ΔDEC) = (0.◦0006, − 0.◦ 0003).

4.3

The Chandra Observations

4.3.1

Data Collection

Sgr B2 was observed using the ACIS-I onboard the Chandra observatory on 29
March 2000 and 16 July 2001. We use event ﬁles provided by the standard pipeline

CHAPTER 4. OBSERVATIONS

54

Figure 4.4: The distribution of the Chandra aspect oﬀsets (Weisskopf et al. 2003).
processing. Only Grade 0, 2, 3, 4 and 6 events are used in the analysis. Each CCD
chip subtends an 8.3-arcmin square on the sky while the pixel size is 0.5 arcsec. The
on-axis spatial resolution is 0.5-arcsec with a full width at half-maximum (FWHM).
Images, spectra, ancillary ﬁles and response matrices have been created using theciao v3.4 software.

4.3.2

Position Accuracy

Taking the advantage of the highest positional resolution, the Chandra aspect oﬀsets have been determined by numerous observations with X-ray sources having
well-determined radio or optical/IR counterparts. Figure 4.4 shows the distributionof aspect oﬀsets for Chandra. The absolute positional accuracy is determined to be
0.6 arcsec (Weisskopf et al. 2003), which is the best quality among the four observatories. In this thesis, we regard the Chandra frame as the standard astrometry.

CHAPTER 4. OBSERVATIONS

4.4

The ASCA Observations

4.4.1

Data Collection

55

The ASCA observations of Sgr B2 were performed on 22 and 24 September 1994 in
4-CCD mode for ∼ 80ks in total. We use the event ﬁles provided by the standard
pipeline processing. Since the observation dates and aim points are almost the same,
we merge the two observation data. Since the GIS data do not have enough energy
resolution (∼ 470 eV at 5.9 keV) to separate the iron lines, we use the SIS dataalone for the analysis. We use the software package HEASoft 6.2 in the analysis.

4.4.2

The Position Tuning

Since no available point sources for the absolute astrometry correction are found in
the SIS ﬁeld of view, the astrometric uncertainty of the SIS image remains about
40 arcsec1 .

1

See http://heasarc.gsfc.nasa.gov/docs/asca/cal probs.html

Chapter 5
Result I — Suzaku Discoveries of
New Diﬀuse Sources
5.1

Suzaku Overview of the Sgr B2 Region

The X-ray spectrum of all the Sgr B region is given in ﬁgure 5.1. The spectra of
the four XISs (XIS0–XIS3) were co-added and the night earth spectrum (non X-ray
background; here, NXB) was subtracted. With the excellent energy resolution of the
XIS for diﬀuse sources, we can clearly resolve the 6.40 keV, 6.67 keV and 6.97 keVlines. These are Kα lines from neutral Fe I, He-like Fe XXV and hydrogenic Fe XXVI.
The 6.97 keV line may contain a small fraction of Kβ line of Fe I (7.06 keV). Weaklines seen above 7 keV are Kα of Ni I (at ∼ 7.5 keV), Kα of Ni XXVII + Kβ of
Fe XXV (at ∼7.8–7.9 keV) and Kβ of Fe XXVI + Kγ of Fe XXV (at ∼8.2–8.3 keV).
In order to study more quantitatively, we divided the XIS FOV into 16 (4 × 4)
areas as is shown in ﬁgure 5.2 by the solid grids. For brevity, we assign coordinates
for each square position such as (0, 0) and (3, 3), where the former and latterare lower-left and upper-right positions, respectively. Since the roll angle of theobservation was nearly 0◦ , the north of this coordinate is also the north of the sky
coordinate. We made X-ray spectra from the 16 squares. Since we already made ﬁne
adjustment among the four XISs, we added the four XIS data to make a composite
spectrum. The NXB spectra were constructed from the night earth data in the samedetector area of the 16-square regions. The NXB spectra from the four XIS werealso co-added and were subtracted from the 16 squares of the Sgr B2 region. In theseadding procedure, the four corners, (0, 0), (0, 3), (3, 0) and (3, 3) contain the Mn I
Kα and Kβ line from the built-in calibration sources, and we therefore excluded
these areas in the following analysis. The NXB subtracted spectra were ﬁtted witha thermal bremsstrahlung model, plus the Fe absorption edge at 7.1 keV and four56

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES57

Figure 5.1: The X-ray spectra from the full FOV of the XIS, except the CCD corners
irradiated by the build-in calibration sources. All four XIS data sets are co-added.

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES58
emission lines, which are the Kα lines from Fe I (6.40 keV), Fe XXV (6.67 keV) andFe XXVI (6.97 keV) and the Kβ line from Fe I (7.06 keV). Since the energies of thelatter two lines have a small separation of only about 90 eV, we constrain the center
energy and ﬂux such that the separation energy between the Kα and Kβ of Fe I is
ﬁxed to the theoretical value of 658 eV (7058 − 6400) and the ﬂux of Kβ is ﬁxed tobe 12.5% of that of Kα.
Since we ﬁnd that the best-ﬁt bremsstrahlung temperature is the same (7 keV)in all the regions within errors, we ﬁx the temperature to be 7 keV, and extract
the line and continuum ﬂuxes from the 12 squares. Using the best-ﬁt ﬂuxes of the
6.40 keV, 6.67 keV, and 6.97 keV lines, we make color code maps (ﬁgure 5.2), which
are the ﬂuxes of the 6.40 keV (ﬁgure 5.2a) and the 6.67 keV lines (ﬁgure 5.2b) andthe ﬂux ratio of the 6.97 keV line to the 6.67 keV line (ﬁgure 5.2c). The continuum
ﬂux (normalization of the 7-keV bremsstrahlung) map is also given in ﬁgure 5.2d.
Using the ﬂat ﬁeld simulation data provided by the XRT team, we corrected thetelescope vignetting eﬀect.
Both the 6.40 and 6.67 keV line maps show a large contrast in red and green
colors, which means that the ﬂuxes have large variation from position to position.
On the other hand, the 6.97 keV/6.67 keV ratio is rather smooth (the contrast ofcolors is weak) except the position of (3, 2). If we exclude this position, the meanvalue of the other 11 squares is 0.39. Koyama et al. (2007d) reported that thebulk of the GC diﬀuse emission is due to hot plasma in nearly collisional ionization
equilibrium (CIE) with temperature of about 6–7 keV. Since the mean value of the
6.97 keV/6.67 keV line ratio in the Sgr B2 region is almost identical to that of theGC, we safely assume that the origin of these lines is the same: hot plasma in the
CIE with a temperature of 7 keV. The absolute surface brightness is about a halfof the GC. The continuum ﬂux is also smooth with a small but a systematic excess
at the north position (2, 2). A signiﬁcant excess of the 6.67 keV ﬂux and deﬁcitof the 6.97 keV/6.67 keV ﬂux ratio at the position of (3, 2) indicates a presenceof an extra component with a strong 6.67 keV line but no strong 6.97 keV line.
The rather smooth distribution of the continuum X-ray ﬂuxes (see ﬁgure 5.2d) alsosuggests that the excess emission is mainly due to the 6.67 keV line. Since the 6.67
keV line is most likely the Kα of Fe XXV, the provable origin of the local 6.67 keVexcess is due to hot plasma of the temperature lower than that of the large scale GC
plasma. The 6.97 keV/6.67 keV ﬂux ratio at the position (3, 2) is 0.14, which canbe converted to the ionization temperature to be 4 keV. This temperature is typicalto a young SNR like Sgr A East.

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES59

Figure 5.2: Fluxes of the 6.40 keV (a) and the 6.67 keV lines (b), the ﬂux ratio of 6.97 keV/6.67 keV (c), and the continuum ﬂux (normalization of the 7-keV
bremsstrahlung) map (d). The embedded numbers are ﬂuxes in units of 10−6 photons cm−2 s−1 arcmin−2 or ﬂux ratios. The color means that the average value iswhite, and the increasing (decreasing) excess above (below) the average is given
with increasing deepness of red (green), where the maximum (minimum) is 2.0 (0.1)
times of the average value.

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES60

Figure 5.3: The 6.67 keV line map (the 6.58–6.74 keV band map) showing bright spotat the northwest corner with the 2–5 keV continuum map as green contours. The
ﬂuxes are normalized by the 6.7 keV ﬂat-ﬁeld image. The source and backgroundregions are shown by the white solid and dotted ellipses, respectively.
The 6.40 keV ﬂuxes (ﬁgure 5.2a) shows local excesses at the position of Sgr B2(2, 2) and the east of Sgr B2 (1, 3). The presence of 6.40 keV is strong evidence fordense and cool molecular clouds.

5.2

New Diﬀuse Sources

5.2.1

SNR Candidate — G 0.61+0.01

We made a narrow band image of 6.67 keV (the 6.58–6.74 keV band) in ﬁgure 5.3.
We see a clear 6.67 keV ﬂux excess at the northwest corner. To conﬁrm the 6.67
keV excess, we referred to the Chandra and XMM-Newton data with long exposure
time (shown in ﬁgure 5.4). Since the energy resolution of the Chandra ACIS islimiting in the separation of the 6.40 and 6.67 keV lines, it is unclear whether the
6.67 keV source is present or not. The XMM image in the 6.67 keV band shows

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES61
a clear elongated structure near the same position. Accordingly the presence ofthe 6.67 keV line source is certain. On the other hand, the continuum band image(shown as green contours in ﬁgure 5.3) shows only a hint of enhancement, andshows no clear structure. This source is therefore very peculiar, being dominant
only in the 6.67 keV line. Since the dominance of the 6.67 keV line suggests thatthe excess is a new SNR (see also the discussion below), we designate this source
Suzaku J1747.0−2824.5 (G 0.61+0.01) from its center position.
We made the NXB-subtracted spectra of the SNR from the region indicated by
the solid ellipse in ﬁgure 5.3 for both the three FI CCDs (XIS0, XIS2 and XIS3 areco-added) and one BI CCD (XIS1). In these spectra, the cosmic X-ray background(CXB) and GCDX are still included. We therefore made the NXB-subtracted spectra for the region indicated by the dotted ellipse, and subtracted this local background (CXB + GCDX) from the source spectrum. All the spectra have beencorrected for the vignetting at 6.67 keV. The results are given in ﬁgure 5.5, for theFIs and BI, separately. We see a pronounced peak at 6.67 keV, but not at 6.97 keV.
The 6.67 keV line shape is asymmetric with a tail at lower energy. In order toverify the line structure, we derive ﬂuxes of the 6.40, 6.67 and 6.97 keV lines (Kαline of Fe I, XXV and XXVI) in the source and the background regions, by applying
a phenomenological model (a bremsstrahlung continuum and many Gaussian lines)
in the raw data (no background subtraction). The resulting 6.40, 6.67 and 6.97 keVline ﬂuxes are 2.22 ± 0.22, 5.17 ± 0.52 and 0.48 ± 0.05 for the G0.61+0.01 (source)
region, and 0.61 ± 0.06, 0.68 ± 0.07 and 0.30 ± 0.03 for the background region, where
the ﬂux unit is 10−6 photons cm−2 s−1 arcmin−2 . In contrast to the 6.67 keV line,
we see no signiﬁcant excess in the 6.97 keV line from the source region comparedto the background region. Thus we conﬁrm that G 0.61+0.01 emits strong 6.67 keVline, but very weak 6.97 keV line. The small excess of the 6.40 keV line makes thelow energy tail in the 6.67 keV line.
The FI and BI spectra are simultaneously ﬁtted with a plane parallel shockmodel (VPSHOCK in the XSPEC package) adding two Gaussian lines at 6.40 keV
and 7.06 keV. These two lines represent the Kα and Kβ lines of Fe I, where theﬂux of latter line is ﬁxed at 12.5% of the ﬂux of the former line (Kaastra & Mewe
1993). The best-ﬁt results and parameters are shown in ﬁgure 5.5 and table 5.1.
Although we detected the 6.40 keV line from G 0.61+0.01, it is very diﬃcult to judge
whether this line is really attributable to G 0.61+0.01, due to spilled-over ﬂux fromthe adjacent source Sgr B2, or due to a ﬂuctuation of a larger scale structure in the
6.40 keV line. As for the last possibility, we see a large scale 6.40 keV enhancement

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES62

Figure 5.4: The exposure corrected 6.67 keV line maps of (a) Chandra 2000 and (b)XMM-Newton 2004. The source regions are shown by the solid circles.

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES63

Figure 5.5: Top: the X-ray spectrum of the sum of 3 FI CCDs (XIS0, 2 and 3) for
the new SNR (G 0.61+0.01) with the best-ﬁt VPSHOCK model. Bottom: same as
the top, but of the BI CCD (XIS1).

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES64

Table 5.1: The best-ﬁt parameters for G 0.61+0.01 with the VPSHOCK model plus
two emission lines
Parameter
NH (1023 H cm−2 )kT (keV)ne t (1011 cm−3 s)Abundances∗
Ca
Fe
Neutral iron linesa
I6.40 (10−6 photons cm−2 s−1 )I7.06 (10−6 photons cm−2 s−1 )Flux and Luminosity
F2−10 † (10−13 ergs cm−2 s−1 )L2−10 ‡ (1034 ergs s−1 )χ2 /dofValue
1.6+0.7−0.4
3.2+2.3−0.9
1.9+4.7−0.8
3.5+3.1−2.4
5.1+1.2−1.1
5.1+2.4−2.5
0.6
7.5+1.1−2.2
1.5+0.1−0.2
99.0/78

Note—The uncertainties indicate the 90% conﬁdence limit.
∗
The elements which are not listed are ﬁxed at 1.0 (solar ratio).
a
The line energy of Kα and Kβ is ﬁxed at the theoretical value(6.40 and 7.06 keV, respectively; Kaastra and Mewe 1993) andthe intensity of Kβ is ﬁxed at 12.5% (Kaastra and Mewe 1993) ofthat of Kα.
†
Observed ﬂux in the 2.0–10.0 keV band.
‡
Absorption corrected luminosity in the 2.0–10.0 keV band.

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES65
in the northwest compared to the background region in the southeast (see ﬁgure5.6). In any case, we ignore this line in the discussion of G 0.61+0.01 because the6.40 keV line ﬂux is only 3% of that in Sgr B2 (see tables 5.1 and 5.3). Since theSuzaku spatial resolution is not good enough, there may be possible contaminationof unresolved point sources. To check this problem, we searched for point sources
using the Chandra archival data and found no point source in the source region. On
the other hand, in the background region, there are 48 point sources. The total ﬂux
in the 2–10 keV band is 5×10−13 ergs cm−2 s−1 , which is only ∼2% of the CXB +GCDX ﬂux of 2.6×10−11 ergs cm−2 s−1 , and hence can be ignored in the present
data analysis and discussion.
The best-ﬁt temperature of ∼3 keV and the over-abundance of Fe are consistent
with an ejecta of a SNR and are similar to those found in the central region of Sgr A
East, a young SNR near the GC. The high temperature component of Sgr A East iskT ∼4–6 keV (Sakano et al. 2004; Park et al. 2005; Koyama et al. 2007c), and iron isover-abundant by a factor of 4–5 (Maeda et al. 2002; Sakano et al. 2004; Park et al.
2005; Koyama et al. 2007c). Thus G 0.61+0.01 is likely an ejecta dominant central
region of an SNR. We note that Sgr A East has a low-temperature componentof about 1 keV, while G 0.61+0.01 has not. The absence of the low-temperatureplasma may be due to the large absorption. The NH value of 1.6 × 1023 H cm−2 islarger than that of the typical value to the GC (6 × 1022 H cm−2 ) (Sakano et al.
2002). Therefore, G 0.61+0.01 could be located behind or at the rim of the Sgr B2
cloud. Since G 0.61+0.01 is located in the south of an expanding radio shell (Okaet al. 1998), which is probably interacting with the Sgr B2 cloud rim, we regard
that the distance of G 0.61+0.01 is the same as that of Sgr B2 and is assumed tobe 8.5 kpc (Reid et al. 1988). Then the 2–10 keV band luminosity is estimated tobe 1.5 × 1034 ergs s−1 , which is typical for an ejecta plasma of an SNR. The sizeof G 0.61+0.01 (the solid ellipse in ﬁgure 5.3) is 2.2 ×4.8 , which corresponds to
5.5 pc × 12 pc at a distance of 8.5 kpc. Assuming the plasma emission is due to a
uniform density ellipsoid with the 3-axis radii of 2.7 pc, 2.7 pc and 6 pc, we estimate
physical parameters of G 0.61+0.01 (table 5.2). Although the iron abundance is 3–4
times the solar value, the total mass of protons are much larger than those of irons.
We therefore assume that electron density (ne ) is equal to the proton density (np ),and that protons carry most of the plasma mass (mp ). Dividing the radius of the
major axis (6 pc) by the sound velocity of the 3.2 keV plasma (v = 1.4×108 cm s−1 ),we obtain the dynamical time scale (tdyn ) of ∼4×103 years. If, instead, we use
the ionization parameter (ne t) and electron density (ne ), then the ionization time

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES66

Table 5.2: The physical parameters of G 0.61+0.01
Parameter
Emission measure (EM)a (cm−3 )The electron density ne b (cm−3 )Total mass M c (M )Thermal energy E d (ergs)
The dynamical time scale tdyn e (s)The ionization time scale tioni f (s)Value
1.4 × 1057
0.9
1.3
2.4 × 1049
1.3 × 1011
2.1 × 1011

Note—The plasma is assumed to be a uniform density ellipsoidwith the 3-axis radii of 2.7 pc, 2.7 pc and 6 pc (see text).
a
EM = ne nH V , where ne and nH are the electron and hydrogendensity.
bderived from EM assuming that ne and nH are equal.
c
M = ne mp V , where mp is the proton mass and V is the plasma
volume.
d
E = 3ne kT V .
ederived from the radius of the major axis of the plasma ellipsoid
divided by the sound velocity of the ∼3 keV plasma.
fderived from the ionization parameter (see table 5.1) divided by
the electron density.

scale (tioni ) is estimated to be ∼ 7 × 103 years. Since the source size of 2.2 ×4.8
is comparable to the half power diameter (∼ 2 ), the real size of G 0.61+0.01 mustbe smaller. Therefore the quoted value of tdyn ∼ 4×103 years should be an upper
limit. That of tioni ∼ 7×103 years is also an upper limit, because ne is inversely
proportional to the square root of the plasma volume. Thus the age of G 0.61+0.01
is probably younger than 7×103 years.
Another possibility is that G 0.61+0.01 is a part of a larger SNR. Since G 0.61+0.01
is found at the edge of the XIS ﬁeld, other parts of a candidate SNR may be out of
the XIS ﬁeld. In this scenario, G 0.61+0.01 may be a part of the expanding radio
shell discovered by Oka et al. (1998). The kinetic energy of the radio shell is a
few times 1052 erg s−1 , which is within the range of a single or multiple supernova
explosions. Thus follow-up X-ray observations including this expanding radio shellis highly required.

5.2.2

XRN Candidate — M 0.79−0.09

We have made a narrow band image at 6.40 keV (the 6.33–6.46 keV band) in ﬁgure
5.6. We see two bright spots in the north. One is the “Sgr B2 cloud” which has
been already found as a strong 6.40 keV source (Koyama et al. 1996), and the otherCHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES67

Figure 5.6: The 6.40 keV line map (the 6.33–6.46 keV band map) showing brightspots at the Sgr B2 cloud and M 0.74-0.09. The ﬂuxes are normalized by the 6.40 keV
ﬂat-ﬁeld image. The sources and background regions are shown by the solid circles
and dotted ellipse, respectively.

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES68
is a newly discovered source. We again referred to the same data of Chandra and
XMM as in the case of G 0.61+0.01 (shown in ﬁgure 5.7). In the Chandra data, thisexcess is found near the edge of the ACIS FOV. The XMM data show a clear excess
near this source. The presence of the 6.40 keV line supports the presence of cool anddense gas clouds. We therefore designate this new source as Suzaku J1747.7−2821.2(M 0.74−0.09) from its peak position. We also rename Sgr B2 cloud to M 0.66−0.02
to avoid the confusion between the whole area and the individual molecular cloud.
We made X-ray spectra of M 0.66−0.02 and M 0.74−0.09 from the circles given inﬁgure 5.6, for comparison. The background spectrum is made from the dotted ellipse
and subtracted in the same procedure as in the case of G 0.61+0.01.
The background-subtracted spectra are shown is ﬁgures 5.8 and 5.9. We simultaneously ﬁt the FIs and BI spectra with a model of absorbed power-law plus two
Gaussians near 6.40 and 7.06 keV, which are for the Kα and Kβ lines of Fe I. The
best-ﬁt parameters are shown in table 5.3. This model nicely ﬁts the data except foran excess near the 6.67 keV line in the M 0.66−0.02 spectra. In fact, the 6.67 keVline map (ﬁgure 5.3) shows a weak enhancement at the position of M 0.66−0.02. One
possibility is that the 6.67 keV enhancement is a part of the new SNR candidate
G 0.61+0.01, because it is located in the close vicinity of M 0.66−0.02. The otherpossibility is that the 6.67 keV enhancement is due to the young steller objects(YSOs) embedded in the center of M 0.66−0.02. In fact, this region is relatively
crowded with the Chandra point sources (13 point sources), and at least some ofthem are YSOs with a hint of the 6.67 keV line emission (Takagi et al. 2002). The
total ﬂux (in the 2–10 keV band) of the point sources is ∼10−13 ergs cm−2 s−1 , whichis ∼6% of the M 0.66−0.02 ﬂux (see table 5.3).
The M 0.66−0.02 (Sgr B2 cloud) has been studied extensively with ASCA andChandra. Koyama et al. (1996) and Murakami et al. (2001) concluded that the6.40 keV emission is due to ﬂuorescence by strong X-rays coming from Sgr A∗ ,hence named the X-ray reﬂection nebula (XRN). In this paper, we found a clear Kβline at 7.06 keV with consistent ﬂux ratio to the Kα line (6.40 keV) in the ﬂuorescent
X-ray origin and a deep Fe edge at 7.1 keV. These discoveries provide additional
supports for the XRN scenario of Sgr B2.
For the M 0.74−0.09 region, Miyazaki & Tsuboi (2000) reported ﬂux peaks of the
CS (J=1-0) line emission at (l, b) = (0.◦761, −0.◦ 117) and (0.◦ 764, −0.◦ 064), which isclear evidence for the presence of a molecular cloud. The XIS spectrum of this region
exhibits a strong 6.40 keV line with an equivalent width of 1.6 keV, a 7.06 keV line
and an Fe edge structure at 7.1 keV (see table 5.3). All these features are consistentCHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES69

Figure 5.7: The exposure corrected 6.40 keV line maps of (a) Chandra in 2000 and(b) XMM-Newton in 2004. The source regions are shown by the solid circles.

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES70

Figure 5.8: Top: the X-ray spectrum of the sum of the 3 FI CCDs (XIS0, 2 and 3) forM 0.66−0.02 (Sgr B2 cloud) with an absorbed power-law model and two Gaussianlines. Bottom: same as the top but of the BI CCD (XIS1).

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES71

Figure 5.9: Same as ﬁgure 5.8, but for a new source M 0.74-0.09.

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES72
with being from Kα, Kβ and K-edge from Fe I. The ﬂux of the 7.06 keV line isabout 10% of that of the 6.40 keV line, which is also consistent with the ﬂuorescent
X-ray origin (Kaastra & Mewe 1993). We note here that the background region is
the same as in the case of G 0.61+0.01. With the same argument in section 5.2.1,
possible point source contribution can be ignored.
Unlike M 0.66−0.02, no hint of YSOs is found so far. No bright point source isfound in the Chandra image. Therefore, the X-rays cannot be the scattering andﬂuorescence by embedded YSOs. If the X-rays from M 0.66−0.02 and M 0.74−0.09
are due to the Thomson scattering and ﬂuorescence of the same irradiating externalsource like Sgr A∗ , then the NH ratio between M 0.66−0.02 and M 0.74−0.09 should
be similar to their 6.40 keV line ﬂux ratio. The observed NH ratio is 0.42, while thatof the 6.40 keV line ﬂux is 0.36, in good agreement of the ﬂuorescence scenario bya single irradiation source. Therefore the XRN scenario by past activity of Sgr A∗ ,which was successfully applied to M 0.66−0.02, may also be applied to M 0.74−0.09.
A diﬀerent scenario than the XRN is that the 6.40 keV line emission is producedby the collision of electrons. As we reviewed in section 2.3, the cross section of ironK-shell ionization is maximum at the electron energy of a few times 10 keV. Thus
the most probable source is from low energy electrons (LEE) as proposed for theorigin of the Galactic Ridge iron K-shell emission (Valinia et al. 2000). Since a few
10 keV electrons are absorbed in less than 1022 H cm−2 of depth, the produced Xray spectrum should have no signiﬁcant absorption edge. Our observation, however,shows a clear absorption of (4.0–9.6) ×1023 H cm−2 , in far excess of the Galacticinterstellar absorption (Sakano et al. 2002). Thus the LEE origin is unlikely, unless
we assume a special geometry such that the 6.40 keV source is deep in or behindthe dense cloud.

CHAPTER 5. RESULT I — SUZAKU DISCOVERIES OF NEW DIFFUSE SOURCES73

Table 5.3: The result of spectral ﬁttings of M 0.66−0.02 (Sgr B2 cloud) andM0.74−0.09 with a power-law and two Gaussian models.

Parameter
Absorbed power-law model:
Column density NH (1023 cm−2 )Photon index Γ
Gaussian 1 (Fe I K α):Line energy (eV)Intensity (10−5 photons cm−2 s−1 )Equivalent Width (keV)
Gaussian 2 (Fe I K β ):Line energy (eV)aIntensity (10−5 photons cm−2 s−1 )Equivalent Width (keV)
Observed ﬂux† (10−12 ergs cm−2 s−1 )Luminosity‡ (1034 ergs s−1 )χ2 /dof

M 0.66−0.02 (Sgr B2 cloud)M 0.74−0.09

9.6+2.5−0.8
3.2+0.9−0.6

4.0+1.4−1.1
1.4+0.4−0.7

6399+5−5
16.5+0.8−0.3
1.13

6406+6−6
5.9+1.4
−1.0
1.55

7057
1.4+0.5−0.5
0.13
1.5+0.1−0.9
9.7+0.1−5.1
154.7/897064
0.6+0.3−0.3
0.18
1.3+0.2−0.8
2.6+0.4−0.9
54.4/66

Note—The uncertainties indicate the 90% conﬁdence limit.
a
The energy gap between Kα and Kβ is ﬁxed at the theoretical value (+658 eV) (Kaastra
and Mewe 1993).
†
Observed ﬂux in the 4.0–10.0 keV band.
‡
Absorption corrected luminosity in the 4.0–10.0 keV band.

Chapter 6
Result II — Time Variabilities of
Neutral Iron Line
6.1

Spectral Analysis and Error Estimation

The determination of the iron line ﬂuxes is inﬂuenced by a continuum model and abackground. In order to estimate the ﬂux errors caused by models and backgrounds,
we make a comparison analysis for the whole Sgr B2 region. Figure 6.1 shows
relevant diﬀuse sources in the Sgr B2 complex overlaid on the XMM-Newton 6–
7 keV image. Since the FOVs and pointing positions of the Suzaku, Chandra,
XMM-Newton and ASCA observations are slightly diﬀerent from each other, weselected a region in which all the observed areas overlapped (referred to hereafteras the Sgr B2 region). The Sgr B2 region (the overlapping region) is indicated byby the solid squares in ﬁgure 6.1. The subsequent data analysis and discussion are
based on the sources in the Sgr B2 region. Hereafter, we refer to the Sgr B2 cloudas M 0.66−0.02. G 0.61+0.01 is a new source with a strong 6.67 keV line, hence
possibly a new young SNR (Koyama et al. 2007b). G 0.570−0.0018 is an unusual
diﬀuse source discovered by Senda et al. 2002. M 0.74−0.09 is also a new 6.40 keV
source, but it is not discussed in this paper, because it is outside of the Sgr B2
region.
We extracted X-ray spectra from the Sgr B2 region and subtracted the oﬀ-planeblank-sky regions. Thus, the CXB and NXB were subtracted but the GCDX were
not subtracted. As the oﬀ-plane blank-sky, we selected the north ecliptic pole data
for XIS, distributed blank-sky databases for MOS and ACIS, and Lockman holedata for PN and SIS.
The blank-sky-subtracted spectra are shown in ﬁgure 6.2. These spectra exhibit
continuum emissions and pronounced peaks due to Fe I Kα (6.40 keV), Fe XXV Kα74

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE75

Figure 6.1: The XMM-Newton image in 6–7 keV band. The source region for the
Sgr B2 region is described in a solid line. Dashed line circles are individual sources
discovered in this region.
(6.67 keV), and the composite of Fe XXVI Kα (6.96 keV) and Fe I Kβ(7.06 keV). The
continuum emission is composed of: 1) the continuum of the GCDX (Koyama et al.
2007d), 2) the continuum (thermal bremsstrahlung) from G 0.61+0.01, a new SNR
in this region (Koyama et al. 2007b), and 3) a non-thermal power-law continuumwith a deep iron K-edge related to the 6.40 keV line (Koyama et al. 1996).
It is practically impossible to resolve these three continuums; we therefore represented these three continuum by a single power-law continuum with absorption(the wabs*power-law model in the xspec code). We then added four K-shell lines
by Gaussian functions. These are the Fe XXV Kα (∼6.67 keV) and Fe XXVI Kα(∼6.96 keV) lines due to the GCDX, and Fe I Kα (6.40 keV) and Fe I Kβ(7.06 keV)lines due mainly to the Sgr B2 region. Thus, the model spectrum in the xspec codeis given as;
wabs*power-law + 4 Gaussians.

6.1.1

SuzakuThe spectra of the Sgr B2 region and background subtractions of the blank sky for
XIS0, 1, 2 and 3 were performed separately, however all the FI-CCD spectra (XIS0,

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE76

6
7Energy [ keV ]

8

0.1
0.01

Counts sec−1 keV−1χ5−4−2 0 2 4

0.1Counts sec−1 keV−1

0.01−4−2 0 2 4

χ1(b) Suzaku XIS−BI 2005

1(a) Suzaku XIS−FI 2005

1
0.1
0.01

Counts sec−1 keV−1χ67
Energy [ keV ]

8

5

8

0.1

1
6
7
Energy [ keV ]0.01

Counts sec−1 keV−1χ5−4−2 0 2

0.1

Counts sec−1 keV−1

0.01−2 0 2

χ

6
7
Energy [ keV ](f) XMM−Newton PN 2001

1(e) XMM−Newton MOS 2001

8

5

6
7
Energy [ keV ]

8

χ
5

6
7
Energy [ keV ]

8

0.1
0.01
−4−2 0 2−2 0 20.01

0.1

Counts sec−1 keV−11(h) Chandra ACIS 2000

1(g) Chandra ACIS 2001
Counts sec−1 keV−1

8−2 0 2 4

0.1
0.01

Counts sec−1 keV−1
χ−2 0 2 45χ67Energy [ keV ](d) XMM−Newton PN 2004

1(c) XMM−Newton MOS 2004

5

5

6
7
Energy [ keV ]

8

χ

0.1
0.01
−2 0 2

Counts sec−1 keV−11(i) ASCA SIS 1994

5

6
7
Energy [ keV ]8

Figure 6.2: The blank-sky-subtracted spectra of the Sgr B2 region, with the Suzaku
XIS FI sensors (a) and BI sensor (b) in 2005, XMM-Newton MOS (c) and PN (d) in
2004, (e) and (f) in 2001, Chandra ACIS-I (g) in 2001 and (h) in 2000, and ASCASIS (i) in 1994. The solid lines indicate the best-ﬁt models (see text).

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE77
1 and 3) were combined, because their response functions are almost identical witheach other. We ﬁtted the FI-CCD and BI-CCD spectra simultaneously in the energy
band of 5–8 keV with the model given by equation (1). All the parameters were free
under the following constraints: the line widths were ﬁxed to be 0 eV except that of
the Fe XXV Kα line (Koyama et al. 2007d), and the energy interval between Fe I Kα
(6400 eV) and Kβ (7058 eV) was ﬁxed at the theoretical value of 658 eV (Kaastra& Mewe 1993). The best-ﬁt spectra for the FI and BI are shown in ﬁgures 2(a) and2(b) respectively, while the best-ﬁt parameters are listed in table 6.1.
Although the K-shell lines from highly ionized iron, the Fe XXVI Kα and Fe XXV
Kα lines are due to the largely extended GCDX, the best-ﬁt center energy of the
Fe XXVI Kα line (6665 eV) and the ﬂux ratio of Fe XXVI Kα to Fe XXV Kα (0.30)
are slightly smaller than those (6680 eV and 0.34, respectively) found in the GC
region (Koyama et al. 2007d). This may be due to the contamination of a strong
Fe XXV Kα line of a new SNR G 0.61+0.01 (Koyama et al. 2007b).
When ﬁtting the XMM-Newton, Chandra and ASCA spectra, we ﬁxed the line
centroids and the Kα ﬂux ratio of Fe XXVI and Fe XXV to the best-ﬁt Suzaku values
given in table 6.1, because the statistics and energy response of these satellites werelimited compared to Suzaku.
Then we deﬁned four types of powerlaw continuum models as follows,
Model I Absorption column density NH and powerlaw index Γ are left free.
Model II NH is a free parameter but Γ is ﬁxed at the Suzaku value.
Model III NH is ﬁxed at the Suzaku value and Γ is free.
Model IV NH and Γ are ﬁxed at the Suzaku values.
With these four continuum models, we operated line ﬂux estimations on theother spectra.

6.1.2

XMM-Newton

At ﬁrst, we ﬁtted the MOS and the PN spectra independently with all parameters
free (predeﬁned as Model I). Since the lines and edge structure cannot be ﬁtted wellfor the PN spectra, we therefore added some systematic gain oﬀsets.
Then we used Model II–IV and ﬁt the spectra. The best-ﬁt parameters are shownin table 6.2 and 6.3. The line ﬂuxes are consistent within 10% among the models,
which is comparable with statistical errors.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE78

Table 6.1: The Suzaku result of spectral ﬁtting for the Sgr B2 region
Suzaku 2005 XIS—Continuum—(1) NH (1023 H cm−2 )(2) Γ(3) Fpow (10−4 photons cm−2 s−1 )—Neutral iron lines—
Gaussian 1(4) E640 (eV)(5) σ640 (eV)(6) F640 (10−4 photons cm−2 s−1 )Gaussian 2(7) E706 (eV)(8) σ706 (eV)(9) F706 (10−4 photons cm−2 s−1 )—Ionized iron lines—
Gaussian 3(10) E667 (eV)(11) σ667 (eV)(12) F667 (10−4 photons cm−2 s−1 )Gaussian 4(13) E697 (eV)(14) σ697 (eV)(15) F697 (10−4 photons cm−2 s−1 )(16) δGain (eV)
χ2 /dof3.0+0.3−0.5
2.6+0.2−0.2
5.2+0.4−0.46400+4−1
0 (ﬁx)
1.14+0.04
−0.03
7058b
0 (ﬁx)
0.11+0.05−0.076665+4−327+7−9
1.07+0.04−0.05
6966+14−130 (ﬁx)
0.32+0.05
−0.04
0(ﬁx)
574.43/514

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption
column density. (2): Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band. (4),(7),(10),(13): Line centroid.
(5),(8),(11),(14): Line width in the Gaussian sigma. (6),(9),(12),(15):Observed line ﬂux . (16): Additional gain oﬀset.
b
The energy gap between Kα (6400 eV) and Kβ (7058 eV) is ﬁxed at
the theoretical value (+658 eV) (Kaastra & Mewe 1993).

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE79

Table 6.2: The XMM-Newton 2004 result of spectral ﬁtting for the Sgr B2 region
NH
Γ
Fpow
(1)
(2)
(3)
XMM-Newton 2004 MOS
3.6+0.7
4.8+0.2
I
5.3+1.4
−1.3
−0.7
−0.2
+0.7
a
II
3.6−0.7
2.6
4.9+0.2−0.2+0.2
2.6+0.3
4.8
III
3.0a
−0.3
−0.2
2.6a
4.8+0.2
IV
3.0a
−0.2
XMM-Newton 2004 PN
I
2.1+0.4
1.6+0.2
7.0+0.2
−0.4
−0.1
−0.2
+0.7
a
II
3.7−0.1
2.6
6.8+0.2−0.2+0.2
1.9+0.2
7.1
III
3.0a
−0.2
−0.2
2.6a
6.6+0.2
IV
3.0a
−0.2
Model

F640
(4)

F706
(5)

F667
(6)

F697
(7)

δGain
(8)

χ2 /dof(9)

1.03+0.08−0.08
1.04+0.08−0.08
1.07+0.07−0.07
1.07+0.07−0.07

0.13+0.07−0.07
0.10+0.07−0.07
0.12+0.07−0.07
0.12+0.07−0.07

0.91+0.04−0.08
0.92+0.08−0.04
0.95+0.07−0.07
0.95+0.07
−0.07

0.27c
0.27c
0.28c
0.28c

0(ﬁx)
0(ﬁx)
0(ﬁx)
0(ﬁx)

88.85/93
95.10/94
97.19/94
97.20/95

1.06+0.05−0.07
1.02+0.08−0.05
1.01+0.07−0.05
1.07+0.06−0.06

0.15+0.05−0.15
0.06+0.05−0.06
0.10+0.05−0.10
0.08+0.04−0.05

0.89+0.06−0.06
0.88+0.05−0.09
0.85+0.05−0.06
0.92+0.05
−0.05

0.27c
0.26c
0.27c
0.27c

21+3
−3
23+4
−3
21+4
−3
21+6
−3

185.62/169
209.71/170
188.47/170
223.78/171

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(5),(6),(7): Observed line ﬂux of 6.40, 7.06, 6.67, 6.97 keV line (10−4 photons cm−2 s−1 ). (8): Additional
gain oﬀset (eV). (9): χ2 value per Degree of Freedom.
aFixed at the value obtained with Suzaku ﬁt.
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα determined by the Suzaku ﬁt.

Table 6.3: The XMM-Newton 2001 result of spectral ﬁtting for the Sgr B2 region
NH
Γ
Fpow
(1)
(2)
(3)
XMM-Newton 2001 MOS
0.4+0.6
8.8+0.4
I
0.5+1.3
−0.5
−0.5
−0.4
+0.3
a
II
4.2+1.0
2.6
8.2
−0.9
−0.3
1.3+0.4
8.9+0.4
III
3.0a
−0.4
−0.4
2.6a
7.7+0.3
IV
3.0a
−0.3
XMM-Newton 2001 PN
I
1.5+1.0
0.9+0.4
9.4+0.3
−0.9
−0.4
−0.3
+0.3
a
II
4.7+0.3
2.6
9.2
−0.4
−0.3
+0.3
1.4+0.2
9.6
III
3.0a
−0.2
−0.3
a
a
2.6
8.6+0.3
IV
3.0
−0.3
Model

F640
(4)

F706
(5)

F667
(6)

F697
(7)

δGain
(8)

χ2 /dof(9)

1.52+0.15−0.16
1.53+0.12−0.26
1.43+0.15−0.16
1.55+0.14−0.16

0.16+0.16−0.14
0.31+0.12−0.22
0.16+0.14−0.16
0.32+0.14−0.14

1.22+0.15−0.16
1.16+0.17−0.19
1.10+0.15−0.15
1.28+0.14
−0.14

0.36c
0.34c
0.32c
0.38c

0(ﬁx)
0(ﬁx)
0(ﬁx)
0(ﬁx)

88.51/101
119.45/102
98.09/102
124.30/103

1.27+0.11−0.12
1.19+0.11−0.11
1.21+0.11−0.11
1.29+0.11−0.10

0.17+0.09−0.09
0.19+0.09−0.08
0.14+0.10−0.08
0.28+0.09−0.08

0.81+0.11−0.11
0.72+0.10−0.10
0.73+0.10−0.10
0.88+0.09
−0.10

0.24c
0.27c
0.27c
0.27c

35+4
−11
35+3
−11
32+4
−9
36+4
−10

188.08/164
230.89/165
194.13/165
248.33/166

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(5),(6),(7): Observed line ﬂux of 6.40, 7.06, 6.67, 6.97 keV line (10−4 photons cm−2 s−1 ). (8): Additional
gain oﬀset (eV). (9): χ2 value per Degree of Freedom.
aFixed at the value obtained with Suzaku ﬁt.
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα determined by the Suzaku ﬁt.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE80

Table 6.4: The XMM-Newton 2004 results of spectral ﬁtting for the Sgr B2 region
with multiplied blank-sky backgrounds
NH
Γ
Fpow
(1)
(2)
(3)
†
XMM-Newton 2004 MOS
I
...
...
...
II
...
...
...
III
...
...
...
IV
...
...
...
XMM-Newton 2004 PN
I
4.4+0.2
3.1+0.2
5.5+0.2
−0.9
−0.2
−0.2
+0.2
a
II
3.2+0.3
2.6
5.6
−0.2
−0.2
III
3.0a
2.6+0.3
5.5+0.2
−0.2
−0.2
2.6a
5.5+0.2
IV
3.0a
−0.2
Model

F640
(4)

F706
(5)

F667
(6)

F697
(7)

δGain
(8)

χ2 /dof
(9)

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

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

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

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

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

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

1.05+0.06−0.06
1.07+0.06−0.06
1.08+0.06−0.06
1.08+0.06−0.06

0.06+0.05−0.05
0.06+0.06−0.05
0.07+0.05−0.05
0.07+0.04−0.06

0.90+0.05−0.06
0.91+0.06−0.06
0.92+0.06−0.06
0.92+0.06
−0.06

0.27c
0.27c
0.27c
0.27c

20+4
−3
24+4
−4
23+1
−4
23+1
−4

238.40/169
240.94/170
240.73/170
241.10/171

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photos cm−2 s−1 ).
(4),(5),(6),(7): Observed line ﬂux of 6.40, 7.06, 6.67, 6.97 keV line (10−4 photons cm−2 s−1 ). (8): Additional
gain oﬀset (eV). (9): χ2 value per Degree of Freedom.
†
No correction is needed.
aFixed at the value obtained with Suzaku ﬁt.
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα determined by the Suzaku ﬁt.

We must note that the intensities of the XMM-Newton NXB are time dependent(Carter & Read 2007). The residual of the NXB lines may aﬀect the estimatesof the iron line ﬂuxes. We checked each count rate of MOS and PN over 10 keV,
where the NXB is dominant, and multiplied the background spectra to match thecount rates with those of the source. The background correction factors for thedata in 2004 are relatively small (1.0 for MOS, and 1.4 for PN), while those forthe data in 2001 are 2.1 for MOS, and 2.6 for PN. We re-ﬁtted the spectra withthese new backgrounds. The results are shown in table 6.4 and 6.5. The derived
line ﬂuxes are almost consistent within the statistical errors, while the inconsistencyof the continuum ﬂuxes between MOS and PN is improved. We also mention thatin this method we overestimate the CXB component. The CXB surface brightness
with a typical Galactic absorption of 1 × 1023 cm−2 is about 1 × 10−15 ergs cm−2 s−1
arcmin−2 . The CXB ﬂux in the Sgr B2 region, the size of which is 106 arcmin2 , isexpected to be 1 × 10−13 ergs cm−2 s−1 arcmin−2 , which is just a few percent of thecontiuum ﬂux.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE81

Table 6.5: The XMM-Newton 2001 results of spectral ﬁtting for the Sgr B2 region
with multiplied blank-sky backgrounds
NH
Γ
Fpow
(1)
(2)
(3)
XMM-Newton 2001 MOS
6.0+3.7
3.9+0.3I
7.2+5.5
−3.6
−2.4
−0.3
+0.3
a
2.6
4.4
II
1.8+1.3
−1.3
−0.3
III
3.0a
3.7+0.8
4.1+0.3
−0.7
−0.3
2.6a
4.6+0.3
IV
3.0a
−0.3
XMM-Newton 2001 PN
I
6.2+1.8
4.1+0.7
5.0+0.3
−1.4
−0.9
−0.3
+0.3
a
II
3.5+1.1
2.6
5.1
−0.9
−0.3
III
3.0a
2.8+0.1
4.9+0.3
−0.3
−0.3
2.6a
5.0+0.3
IV
3.0a
−0.3
Model

F640
(4)

F706
(5)

F667
(6)

F697
(7)

δGain
(8)

χ2 /dof(9)

1.49+0.16−0.16
1.51+0.16−0.16
1.52+0.16−0.15
1.46+0.15−0.15

0.21+0.15−0.15
0.14+0.14−0.14
0.17+0.15−0.14
0.10+0.14−0.10

1.23+0.16−0.16
1.23+0.15−0.16
1.25+0.15−0.15
1.17+0.15
−0.15

0.36c
0.34c
0.32c
0.35c

0(ﬁx)
0(ﬁx)
0(ﬁx)
0(ﬁx)

84.18/101
92.55/102
88.01/102
94.70/103

1.28+0.12−0.11
1.32+0.11−0.12
1.35+0.03−0.06
1.34+0.11−0.11

0.24+0.09−0.10
0.21+0.10−0.09
0.24+0.05−0.03
0.23+0.10−0.09

0.84+0.11−0.11
0.87+0.11−0.05
0.92+0.04−0.03
0.90+0.09
−0.10

0.25c
0.27c
0.27c
0.27c

33+4
−10
32+4
−9
35+1
−12
34+4
−11

260.40/164
267.51/165
267.11/165
267.94/166

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(5),(6),(7): Observed line ﬂux of 6.40, 7.06, 6.67, 6.97 keV line (10−4 photons cm−2 s−1 ). (8): Additional
gain oﬀset (eV). (9): χ2 value per Degree of Freedom.
aFixed at the value obtained with Suzaku ﬁt.
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα determined by the Suzaku ﬁt.

6.1.3

ChandraWith the same method as the previous section, we ﬁtted the Chandra spectra with
Model I. The spectrum in 2000 needs a systematic oﬀset of ∼ −20 eV to match thespectral shape. The spectrum in 2001 required too strange photon index (Γ ∼ −2.2)
and made large residuals over 7 keV when we ﬁxed a photon index at 2.6. We checked
the count rates in the 10.0–12.5 keV band between the source and the blank-sky
spectra, then we found that no correction is needed for the data in 2000 while wemust multiply the background by a factor of 2.4 to match the count rate for thedata in 2001. The best-ﬁt values are derived from the spectral ﬁt with this newcorrected background. Obtained values are shown in table 6.6. The ﬂux of each line
is determined within 10% error among the four continuum models.

6.1.4

ASCA

The ASCA spectrum needs large systematic gain oﬀset about 80 eV to adjust ironline centroids. This large oﬀset is caused by the chip-to-chip variation of the gain1 .
1

See http://heasarc.gsfc.nasa.gov/docs/asca/4ccd.html

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE82

Table 6.6: The Chandra 2000 and 2001 results of spectral ﬁtting for the Sgr B2
regionNHΓ(1)(2)
Chandra 2001 ACIS
1.0+1.7I
2.4+3.4
−2.4
−1.5
a
2.6
II
4.6+1.4
−1.4
III
3.0a
1.3+0.9
−0.7
a
2.6a
IV
3.0
Chandra 2000 ACISI
1.0+0.5
0.7+0.4
−0.4
−0.5
a
II
3.8+0.4
2.6
−0.3
III
3.0a
1.7+0.3
−0.2
a
2.6a
IV
3.0
Model

Fpow
(3)

F640
(4)

F706
(5)

F667
(6)

F697
(7)

δGain
(8)

χ2 /dof(9)

7.2+0.5−0.5
6.6+0.5−0.5
7.1+0.5−0.5
6.0+0.4−0.4

1.37+0.26−0.26
1.36+0.27−0.25
1.33+0.27−0.24
1.48+0.24−0.24

0.10+0.32−0.10
0.12+0.28−0.12
0.12+0.32−0.12
0.20+0.28−0.20

1.14+0.27−0.27
1.16+0.27−0.27
1.13+0.14−0.41
1.30+0.24
−0.24

0.34c
0.34c
0.34c
0.39c

0(ﬁx)
0(ﬁx)
0(ﬁx)
0(ﬁx)

37.51/50
39.02/51
37.67/51
42.33/52

6.7+0.2−0.2
6.0+0.2−0.2
6.5+0.2−0.2
5.7+0.2−0.2

1.71+0.08−0.08
1.72+0.07−0.10
1.67+0.08−0.08
1.76+0.08−0.08

0.17+0.09−0.04
0.26+0.12−0.08
0.21+0.08−0.09
0.31+0.09−0.09

1.00+0.08−0.07
1.02+0.09−0.08
0.97+0.08−0.07
1.08+0.07
−0.08

0.30c
0.30c
0.29c
0.32c

−20+4
−3
−17+6
−3
−18+8
−3
−19+8
−1

195.97/169
218.89/170
205.41/170
224.15/171

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(5),(6),(7): Observed line ﬂux of 6.40, 7.06, 6.67, 6.97 keV line (10−4 photons cm−2 s−1 ). (8): Additional
gain oﬀset (eV). (9): χ2 value per Degree of Freedom.
aFixed at the value obtained with Suzaku ﬁt.
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα determined by the Suzaku ﬁt.

We ﬁt the SIS spectrum with the four continuum models used in the previous sections. The best-ﬁt parameters are shown in table 6.7. Model dependent ﬂuctuationon the line ﬂuxes are far small compared to the statistical errors.

6.1.5

Comparison Analysis

Model Dependence
The 6.67 keV ﬂuxes obtained by ﬁttings with Model I–IV in a unit of surface brightness by dividing the size of the source region (106 arcmin2 ) are shown in ﬁgure 6.3.
While the values with Model III are systematically low and those with Model IV
are high, the ﬂuctuations are entirely within the statistical errors. Then we use thevalue with Model I as the result of each observation.
Systematic Errors
The 6.67 keV line emission is due to the GCDX and a strong 6.67 keV source,
G0.61+0.01, which contributes about 30% of the total 6.67 keV line emission. Theseemissions can be considered to be constant in time. However, a constant model on
the 6.67 keV ﬂux is rejected by the χ2 statistics (χ2 /dof=56.32/8).

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE83

Table 6.7: The ASCA 1994 result of spectral ﬁtting for the Sgr B2 region
NH
Γ
(1)
(2)
ASCA 1994 SIS
1.1+0.7I
0.7+1.0
−0.7
−0.5
a
II
3.1+0.7
2.6−0.3
2.1+0.3
III
3.0a
−0.3
a
2.6a
IV
3.0
Model

Fpow
(3)

F640
(4)

F706
(5)

F667
(6)

F697
(7)

δGain
(8)

χ2 /dof(9)

5.8+0.2−0.2
5.5+0.2−0.2
5.8+0.2−0.2
5.4+0.2−0.2

1.63+0.18−0.17
1.65+0.16−0.18
1.62+0.14−0.19
1.65+0.17−0.10

0.22+0.15−0.13
0.30+0.17−0.15
0.24+0.13−0.15
0.30+0.17−0.14

1.20+0.18−0.18
1.14+0.16−0.18
1.07+0.18−0.13
1.15+0.15
−0.16

0.35c
0.33c
0.32c
0.33c

71+21
−13
81+18
−20
81+19
−19
81+18
−5

81.81/76
92.69/77
89.00/77
92.87/78

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(5),(6),(7): Observed line ﬂux of 6.40, 7.06, 6.67, 6.97 keV line (10−4 photons cm−2 s−1 ). (8): Additional
gain oﬀset (eV). (9): χ2 value per Degree of Freedom.
aFixed at the value obtained with Suzaku ﬁt.
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα determined by the Suzaku ﬁt.

1.5

MOS

PN

PN

1

MOS

0.5

6.67 keV surface brightness

2

Sgr B2 region [10−6 photons cm−2 s−1 arcmins−2]

1995

2000
Time [year]

2005

Figure 6.3: The 6.67 keV line surface brightness obtained with each observation in
the Sgr B2 region. In this plot, each error bar indicates the 90% conﬁdence. The
values are marked with red circles for Model I, green x marks for Model II, blue
squares for Model III and aqua triangles for Model IV.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE84

1.5

MOS

PN

PN

1

MOS

0.5

6.67 keV surface brightness

2

Sgr B2 region [10−6 photons cm−2 s−1 arcmins−2]

1995

2000
Time [year]

2005

Figure 6.4: The 6.67 keV line surface brightness with Model I but each error bar
in this plot indicates one sigma statistical error plus the systematic error of thedetector. The data are marked with a triangle mark for Suzaku, a circle for XMMNewton, an asterisk for Chandra and a square for ASCA. The average ﬂux is shownin a solid horizontal line.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE85

1.5
1
0.5

corrected 6.40 keV surface brightness

2

Sgr B2 region [10−6 photons cm−2 s−1 arcmins−2]

1995

2000
Time [year]

2005

Figure 6.5: The light curve of the corrected 6.40 keV line surface brightness in the
Sgr B2 region. The XMM-Newton values are the average of MOS and PN. In thisplot, each error bar indicates the 90% conﬁdence limit.
Then we considered systematic responce errors of the eﬃciencies. The absoluteresponce uncertainty of EPIC is 10% and the MOS ﬂux is 10–15% higher than
the PN (Carter & Read 2007). The systematic errors of the other detectors are
10% for Suzaku (Serlemitsos et al. 2007), 5% for Chandra (Schwartz et al. 2000),
and 13% for ASCA SIS2 . The 6.67 keV ﬂux considering these systematic error ineach observation period is ploted in ﬁgure 6.4. The χ2 /dof value with a constantmodel on this plot is 4.848/8, which is acceptable in 70% conﬁdence. Therefore
we conclude that the obtained line ﬂuxes are ﬂuctuated by the systematic errors ofmirror responces. Then we calculate the ﬂux correction factor to normalize the 6.67
keV ﬂux of each observation to the average value, and correct each 6.40 keV ﬂux
with the factor.
Figure 6.5 shows the corrected 6.40 keV line surface brightness in each observation period. In this ﬁgure, the XMM-Newton values are the averaged ﬂuxes of MOS
and PN. The 6.40 keV line ﬂux is signiﬁcantly variable with χ2 /dof = 28.9/6, shows2

See http://heasarc.gsfc.nasa.gov/listserv/ascanews/msg00143.html

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE86
a decline in 2001, and ﬁnally drops into 60% from 2000 to 2005.

6.2

Variability Veriﬁcations on Individual Sources

6.2.1

Surface Brightness Mapping

In order to examine the 6.40 keV line ﬂux in morphology, we make the surface
brightness map of iron lines in each observation. Since the energy resolution of
Chandra or ASCA are not good enough to separate the 6.40 keV line from the
6.67 keV line, we make images in the 6–7 keV band. The surface brightness maps
of 6–7 keV band are made by the raw images by dividing the exposure maps witheﬀective areas. To eliminate the component of the powerlaw continuum, we makethe surface brightness maps in the 5–6 keV band in the same method. We subtract
65% of the 5–6 keV maps from the 6–7 keV maps as a continuum emission assuming
the photon index is 2.6. Finally we rebin all the maps to make each pixel size about
50-arcsec square, about a size of the PSF of ASCA.
Figure 6.6(a)–(d) show the surface brightness maps of iron lines. The Suzaku and
the XMM-Newton images are relatively dim in the whole area compared to thoseof Chandra and ASCA, which is consistent with the result of the spectral ﬁtting.
In the Chandra image, brighter point is found in the west where G 0.570−0.018 is
located (Senda et al. 2002).
We select the brightest point M 0.66−0.02 (Sgr B2 cloud, shown as a solid circle
in ﬁgure 6.6) and G 0.570−0.018 (shown as a dashed circle in ﬁgure 6.6). Then we
examine the variations of the 6.40 keV ﬂux from these regions. Hereafter we exclude
Chandra data in 2001 because they do not have enough statistics for the analysis.

6.2.2

M 0.66−0.02

We extract the spectra of M 0.66−0.02 (Sgr B2 cloud) from the solid circle with aradius of 3.2 arcmin in ﬁgure 6.6 and the blank-sky background spectra in the samedetector position. In the ASCA image, the peak point of M 0.66−0.02 is seems to
be diﬀerent from the others. The diﬀerence is ∼ 1 arcmin, which is similar to the
SIS astrometry uncertainty. It is not clear whether this shift is due to a relatively
large error of the ASCA astrometry (no ﬁne-tuning of the ASCA astrometry was
performed) or due to a real shift in the variable 6.40 keV line. Therefore, in the
ASCA analysis, we selected two circles for M 0.66−0.02, at the same as the positionin the original coordinate of Suzaku, XMM-Newton, and Chandra (P) and at thecenter on the peak of the SIS image (Q).

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE87

Figure 6.6: The surface brightness maps of iron lines obtained with (a) Suzaku XIS
2005, (b) XMM-Newton MOS and PN 2004, (c) Chandra ACIS-I 2000, and (d)ASCA SIS 1994. Each pixel size is 50 × 50

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE88

χ
5

6
7
Energy [ keV ]

0.1
0.01
−2 0 2 4

0.01

0.1

Counts sec−1 keV−1(b) Suzaku XIS−BI 2005

−2 0 2 4χCounts sec−1 keV−1(a) Suzaku XIS−FI 2005

8

5

χ
6
7
Energy [ keV ]

8

0.1
0.01
−4−2 0 2

Counts sec−1 keV−1

0.1
0.01
5

χ

0.1
5

8

χ
8

0.1
0.01
−2 0 2

Counts sec−1 keV−1

0.1
0.01

Counts sec−1 keV−1χ67
Energy [ keV ](h) ASCA SIS 1994

−2 0 2

6
7
Energy [ keV ]

8

0.01

8

(g) Chandra ACIS 2000

5

6
7
Energy [ keV ]−2−10 1 2

Counts sec−1 keV−1

0.1
0.01

6
7
Energy [ keV ]5(f) XMM−Newton PN 2001

−2 0 2χCounts sec−1 keV−1(e) XMM−Newton MOS 2001

58(d) XMM−Newton PN 2004

−2 0 2χCounts sec−1 keV−1(c) XMM−Newton MOS 2004

6
7
Energy [ keV ]567Energy [ keV ]8

Figure 6.7: The spectra for M0.66−0.02 obtained with Suzaku XIS FI sensors (a)and BI sensor (b) in 2005, XMM-Newton MOS (c) and PN (d) in 2004, (e) and (f)in 2001, Chandra ACIS-I (g) in 2000, and ASCA SIS (h) in 1994. The solid lines
indicate the best-ﬁt models with Model I.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE89

Table 6.8: The result of spectral ﬁtting for M 0.66−0.02 with Model I

Observatory
SuzakuDetector
XIS(1) NH
5.0+0.5−0.3(2) Γ
3.0+0.3−0.3(3) Fpow
2.0+0.1−0.1—Neutral iron lines—(4) E640
6399+4−3(5) F640
0.66+0.02−0.02(6) E706
7057b(7) F706
0.06+0.04−0.04—Ionized iron lines—(8) E667
6670+5−5(9) σ667
29+12−16(10) F667
0.36+0.02−0.02(11) E697
6967+21−21(12) F697
0.12+0.03−0.02(13) δGain0(ﬁx)χ2 /dof348.41/354

XMM-Newton
MOS
PN
+1.1
5.4+2.1
3.7
−2.0
−1.1
+0.2
3.2+0.5
1.7
−1.0
−0.3
+0.1
1.6−0.1
2.4+0.1−0.1

XMM-Newton
MOS
PN
+2.8
0.2+2.4
2.5
−0.2
−2.5
+1.1
0.1+1.0
1.2
−0.5
−1.1
+0.2
3.0−0.2
2.4+0.2−0.2

ChandraACIS
4.3+1.1−1.6
1.2+0.5−0.9
2.7+0.1
−0.1

ASCA
SIS†
3.0+2.9−1.4
1.7+1.9−1.4
2.3+0.2−0.2

6399a
0.58+0.04−0.04

0.57+0.04−0.04

0.92+0.11−0.11

0.82+0.10−0.09

1.00+0.05−0.06

0.99+0.13−0.14

0.08+0.04−0.04

0.02+0.03−0.02

0.05+0.09−0.05

0.11+0.08−0.07

0.06+0.06−0.04

0.12+0.13
−0.11

6670a
29a
0.29+0.04
0.25+0.03
0.42+0.09
0.26+0.08
0.25+0.04
0.42+0.13
−0.04
−0.03
−0.09
−0.07
−0.04
−0.14
6967a
0.10c
0.08c
0.14c
0.09c
0.08c
0.14c
0(ﬁx)
23+6
0(ﬁx)
23+13
−19+8
81+21
−4
−17
−3
−11
36.92/40 71.27/68 35.29/38 16.64/27 142.13/144 36.14/37

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(6),(8),(11): Line centroid (eV). (5),(7),(10),(12): Observed line ﬂux (10−4 photons cm−2 s−1 ). (9): Linewidth in gaussian sigma (eV). The other line widths are ﬁxed at zero. (13): Additional gain oﬀset (eV).
†Values are obtained from region ’P’.
aFixed at the value obtained with Suzaku ﬁt.
b
The energy gap between Kα (6400 eV) and Kβ (7058 eV) is ﬁxed at the theoretical value (+658 eV)(Kaastra & Mewe 1993).
c
The intensity of Fe XXVI Kα is ﬁxed at 34% of that of Fe XXV Kα obtained with Suzaku.

With the same method in section 6.1, we ﬁt the Suzaku spectra and determine theline centroids and ﬂux ratio of the ionized lines. Then we ﬁx the line centroids andthe ionized line ratio to the Suzaku values, we ﬁt the other spectra with Model I. The
best-ﬁt spectra and parameters are shown in ﬁgure 6.7 and table 6.8, respectively.
We also operated the spectral ﬁttings with Model II–IV. The ﬁtting results of the
6.67 keV ﬂuxes in a unit of surface brightness are shown in ﬁgure 6.8 (more details
are shown in tables A.1–A.3 in the appendix). The model dependence of the line
ﬂuxes is comparable to the statistical error. For the spectra of XMM-Newton, we
checked the dependence of the NXB with the same method in section 6.1.2, resulting
little inﬂuence on the line ﬂuxes (shown in table A.4,A.5). Therefore we estimatedline ﬂuxes with Model I.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE90

1.5

MOS

PN

PN

1

MOS

0.5

6.67 keV surface brightness

2

M0.66−0.02 [10−6 photons cm−2 s−1 arcmins−2]

1995

2000
Time [year]

2005

Figure 6.8: The 6.67 keV line surface brightness obtained with each observation in
M 0.66−0.02. In this plot, each error bar indicates the 90% conﬁdence. The values
are marked with red circles for Model I, green x marks for Model II, blue squares
for Model III and aqua triangles for Model IV.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE91

3.5
3
2.5
2
1.5

corrected 6.40 keV surface brightness

4

M0.66−0.02 [10−6 photons cm−2 s−1 arcmins−2]

1995

2000
Time [year]

2005

Figure 6.9: The light curve of the 6.40 keV line surface brightness from M0.66−0.02(Sgr B2 cloud). In this plot, each error bar indicates the 90% conﬁdence limit.
The derived SIS ﬂuxes of the 6.40 keV line are 0.99+0.13−0.14 for the spectrum of−4region ’P’, and 0.93+0.11
photons cm−2 s−1 for that of region ’Q’. These ﬂuxes−0.15 10are consistent within the statistical error. Then we used the result of region ’P’ as

the SIS value.
Since the variation of the 6.67 keV ﬂuxes is quite similar to that obtained withthe whole Sgr B2 region shown in ﬁgure 6.4, it is also caused by the systematic errorof each mirror responce. In order to minimize the statistical errors of the correction
factors, we used the values derived in section 6.1.5 to correct the 6.40 keV ﬂuxes.
The surface brightness of the 6.40 keV line is ploted in ﬁgure 6.9. The MOS and
the PN values were averaged for each XMM observation. The time trend is quite
similar to that of the whole region (ﬁgure 6.5) and the constant ﬂux hypothesis isrejected with χ2 /dof=30.7/6. The ﬂux in 2000 is 1.5 times higher than that in 2004
and 2005, and also shows a decline in 2001.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE92

χ
5

6
7
Energy [ keV ]

0.1
0.01
−2 0 20.01

0.1

Counts sec−1 keV−1(b) Suzaku XIS−BI 2005

−2 0 2χCounts sec−1 keV−1(a) Suzaku XIS−FI 2005

8

5

8

(d) XMM−Newton PN 2004

χ
5

6
7
Energy [ keV ]

0.1
0.01
−2 0 2

Counts sec−1 keV−1

0.1
0.01
−2 0 2χCounts sec−1 keV−11(c) XMM−Newton MOS 2004

6
7
Energy [ keV ]

8

5

χ
6
7
Energy [ keV ]

0.1
0.01
−2 0 2

Counts sec−1 keV−1

0.1
0.01
5

8

5

8

χ
8

0.1
0.01
−2−1 0 1

Counts sec−1 keV−1

0.1
0.01

6
7
Energy [ keV ]6
7Energy [ keV ](h) ASCA SIS 1994

−2 0 2χCounts sec−1 keV−1(g) Chandra ACIS 2000

58(f) XMM−Newton PN 2001

−2 0 2χCounts sec−1 keV−1(e) XMM−Newton MOS 2001

6
7
Energy [ keV ]567Energy [ keV ]8

Figure 6.10: The spectra for G0.570−0.018 obtained with Suzaku XIS FI sensors(a) and BI sensor (b) in 2005, XMM-Newton MOS (c) and PN (d) in 2004, (e) and(f) in 2001, Chandra ACIS-I (g) in 2000, and ASCA SIS (h) in 1994. The solid lines
indicate the best-ﬁt models.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE93

6.2.3

G0.570−0.018

G 0.570−0.018 is a diﬀuse source extending to about 10 arcmins (Senda et al. 2002).
We extracted the 6.40 keV ﬂux from G 0.570−0.018 (the dashed circle with a radiusof 2.5 arcmin in ﬁgure 6.6) with the same way as M 0.66−0.02. Since the spectra
do not have enough statistics to derive the Fe I Kβ (7.06 keV) line, the intensity ofKβ is ﬁxed at 12.5% of that of Kα (Kaastra & Mewe 1993). With the same method
as M 0.66−0.02, we ﬁrstly ﬁt the Suzaku spectra with Model I. Then we obtained
6415 eV as a line centroid of the 6.40 keV line, which is slightly higher than thevalues obtained in the other region, and 35 eV as a line width of the 6.67 keV line,
which is also a little wider value. We re-ﬁtted them with the line centroids ﬁxed
at the values of the Sgr B2 region (table 6.1) and we found some residuals between
6.4–6.6 keV. We added one Gaussian function in this energy band, and obtained an
acceptable ﬁt. The centroid of the additional line is 6505 eV, which is consistent
with the Fe XX line. This line energy is consistent with the result of Senda et al.
(2002), which determined G 0.570−0.018 as a SNR candidate.
Since the ﬂux of the 6.51 keV line is also constant in time, the ﬂux ratio among
the ionized iron lines was ﬁxed at the following analysis. Then we ﬁtted the otherspectra with this model (Model I + one Gaussian line). Table 6.9 shows the bestﬁt parameters for G 0.570−0.018. The continuum models and the XMM-Newton
backgrounds have little impact on determinating the line ﬂuxes (the results areshown tables B.1–B.3, B.4 and B.5 in the appendix). Therefore we used the values
obtained with Model I as the results.
We corrected the 6.40 keV ﬂuxes using the correction factors as we did for
M 0.66−0.02. Figure 6.11 shows the time trend of the 6.40 keV line from G 0.570−0.018,
which indicates a clear variability in the 6.40 keV line emission. A constant ﬂux
hypothesis is marginally rejected with χ2 /dof = 13.2/6.
Senda et al. (2002) mentioned that G 0.570−0.018 is a simple SNR candidate,but the 6.40 keV variability suggests that G 0.570−0.018 includes an XRN (seeDiscussion). In fact, Miyazaki & Tsuboi (2000) reported a ﬂux peak of the CS(J=1-0) line emission at (l, b) = (0.◦ 601, − 0.◦ 025), which indicates the presence of amolecular cloud. The light curve of G 0.570−0.018 is similar to that of M 0.66−0.02.
Hence the distances of the two clouds from Sgr A∗ are the same. G 0.570−0.018
may be located in front of M 0.66−0.02 with some oﬀ-sets of the line of sight.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE94

Table 6.9: The result of spectral ﬁtting for G 0.570−0.018 with Model I plus one
Gaussian function

Observatory
SuzakuXMM-Newton
XMM-Newton
Chandra ASCA
Detector
XIS
MOS
PN
MOS
PN
ACIS
SIS
+1.6
+1.7
+1.4
+3.6
+2.7
+3.4
(1) NH
1.5+1.3
6.2
0.4
1.3
1.5
2.0
0.2
−1.0
−2.5
−0.4
−1.3
−1.5
−2.0
−0.2
+1.9
+0.8
+1.2
+1.6
+1.4
+1.5
(2) Γ
2.6+0.6
4.1
1.8
0.5
2.0
2.0
1.2
−0.5
−1.6
−0.6
−0.9
−1.2
−0.6
−0.8
+0.1
+0.1
+0.1
+0.1
+0.1
+0.1
(3) Fpow
1.1−0.1
1.1−0.1
1.3−0.1
1.9−0.1
1.2−0.1
1.2−0.1
0.9+0.1−0.1—Neutral iron lines—
6400a(4) E640(5) F640
0.16+0.03
0.11+0.04
0.14+0.03
0.18+0.06
0.19+0.06
0.26+0.04
0.12+0.05
−0.05
−0.04
−0.03
−0.06
−0.06
−0.04
−0.05
(6) E706
7058a(7) F706
0.02b
0.01b
0.02b
0.02b
0.02b
0.03b
0.01b
—Ionized iron lines—
6665a(8) E667(9) σ667
27a(10) F667
0.29+0.02
0.20+0.04
0.22+0.03
0.18+0.06
0.23+0.06
0.25+0.03
0.20+0.04
−0.02
−0.04
−0.03
−0.06
−0.06
−0.03
−0.02
(11) E697
6966a(12) F697
0.09c
0.07c
0.07c
0.05c
0.07c
0.08c
0.06c
(13) E651
6505+25
6505d
−31
+0.03
(14) F651
0.07−0.02
0.05d
0.06d
0.04d
0.05d
0.06d
0.05d
(15) δGain
0(ﬁx)
0(ﬁx)
35(ﬁx)0(ﬁx)
33(ﬁx)-18(ﬁx)71(ﬁx)
2
192.33/181 56.06/51 70.76/78 40.89/25 21.01/15 82.35/93 6.81/18
χ /dof

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(6),(8),(11),(13): Line centroid (eV). (5),(7),(10),(12),(14): Observed line ﬂux (10−4 photons cm−2 s−1 ).
(9): Line width in gaussian sigma (eV). The other line widths are ﬁxed at zero. (15): Additional gain oﬀset(eV).
aFixed at the value obtained in section 6.1.1.
b
The intensity of Kβ is ﬁxed at 12.5% of that of Kα (Kaastra & Mewe 1993).
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα.
dFixed at the value and ﬂux ratio against Fe XXV Kα line obtained with Suzaku.

CHAPTER 6. RESULT II — TIME VARIABILITIES OF NEUTRAL IRON LINE95

1
0.5
0

corrected 6.40 keV surface brightness

1.5

G0.570−0.018 [10−6 photons cm−2 s−1 arcmins−2]

1995

2000
Time [year]

2005

Figure 6.11: The light curve of the 6.40 keV line surface brightness from
G 0.570−0.018. In this plot, each error bar indicates the 90% conﬁdence limit.

Chapter 7
Summary and Discussion
7.1

Ionized Iron Line Sources

7.1.1

Galactic Center Diﬀuse X-ray Emission

With the Suzaku deep observation, we conﬁrm that the Sgr B2 region is embedded
in a thin hot plasma. The line ratio of 6.97/6.67 keV is 0.39 all over the Sgr B region
except for G 0.61+0.01, which is quite consistent with that of the Galactic Center(Koyama et al. 2007d). Thus we conclude that this largely diﬀuse emission is the
GCDX extending to this region. The intensity is about a half of the Sgr A region
on the Galactic plane, decreasing to the Galactic latitude. This result can constrain
the distribution of the GCDX against the Galactic longitude more strictly than thatderived by Ginga, 1◦ –2◦ .

7.1.2

Supernova Remnants

In addition to the distribution of the GCDX, we discover a clear excess in the
6.67 keV narrow band image. The source named as G 0.61+0.01 is a SNR candidate
which is characterized by a strong 6.67 keV line emission. The high temperatureand overabundance of iron are similar to those of Sgr A East (Koyama et al. 2007c).
On the other hand, there is no low-temperature component which exists in Sgr A
East. The lack of the low-temperature component may be explained by the tail-likestructure below 6.67 keV line of G 0.61+0.01.
We found the iron line of low-ionization at 6.5 keV in G 0.570−0.018, which isconsistent with the result of Senda et al. (2002). However, since the enhancement
at the position of G 0.570−0.018 in the iron line map (shown in ﬁgure 6.6) fadedaway in the time span from 2000 to 2004, the image structure is mostly due tothe 6.40 keV emission. The 6.5 keV line is located bewteen the strong lines of96

CHAPTER 7. SUMMARY AND DISCUSSION

97

6.40 and 6.67 keV, we cannot distinguish the source with the narrow band imaging
technique of Suzaku. Thus this low ionization line may be originated to a nearby
source G 0.61+0.01. We, however, cannot exclude the origin is G 0.570−0.018 itself
as a SNR source proseped in (Senda et al. 2002).

7.2

Neutral Iron Line Sources

7.2.1

The Origin

We measure the intensities of the 6.40 and the 7.06 keV line independently andconﬁrm the ﬂux ratio (Kα/Kβ) is consistent with 0.1. The ratio indicates the ironatom is nearly neutral (Kaastra & Mewe 1993). Then we discuss whether X-ray
photons or energetic particles make the inner-shell ionization and illuminate theneutral molecular clouds.
The equivalent width of Fe I Kα is 1–2 keV by the X-ray irradiation, while theequivalent width is 300 eV by the collisional excitaion by electrons assuming a molecular cloud with the solar abundance, which is reviewed in section 2.3. The observed
equivalent width prefers the X-ray origin. However, since the metal abundance of
GC is typically high, we cannot exclude the electron origin if the metalicity of the
Sgr B2 region is 4 times a solar. Discussed in section 5.2.2, the deep absorption
edge of Fe K also prefers the X-ray origin, but the electron origin cannot be excluded when electron bomberding takes plase at the rear sides of the clouds. These
spectral features strongly support the X-ray origin but cannot exclude the electron
origin.
The time variability of the 6.40 keV line emission is the most decisive evidence for
the X-ray origin. We found the variabilities not only in M 0.66−0.02, Sgr B2 cloud,
but also in G 0.570−0.018. The sizes of M 0.66−0.02 and G 0.570−0.018 are ∼ 10
light years, comparable to the variation time scale. This fact requires X-ray lights
or relativistic particles as the external factor, however, high energy electrons havefew cross section to neutral iron. On the other hand, protons, suggested by a TeV
gamma-ray emission in the Galactic molecular clouds (Aharonian et al. 2006), canbe also considered as the origin. Reviewed in section 2.3, the spectrum by proton
impact is characterised by the broad bump features and the weak continuum withrespect to the lines. We ﬁnd no clear evidences for bump structures. In addition, we
examine the valiability of the continuum emission. Since the inverse bremsstrahung
by protons is very weak, no variability can be seen while the continuum is variablewith the X-ray reﬂection by Thomson scattering. Figure 7.1 shows the variation

CHAPTER 7. SUMMARY AND DISCUSSION

98

2.5
2
1.5

corrected continuum flux

3

M0.66−0.02 [10−4 photons cm−2 s−1]

1995

2000
Time [year]

2005

Figure 7.1: The time trend of the continuum ﬂuxes in M 0.66−0.02. XMM values in2001 and 2004 are the average of MOS and PN. In this plot, each error bar indicatesthe 90% conﬁdence limit.
of the continuum (5–8 keV) ﬂux in M 0.66−0.02. We can see a slight variation,
whereas we do not discuss the signiﬁcance bacause this continuum includes thethermal GCDX component and the CXB and suﬀers large systematic error causedby the background substraction particularly in the XMM-Newton data.
With these observational results, we have concluded that the most likely originof the neutral line emissions in the Sgr B2 complex is X-ray ﬂuorescence.

7.2.2

Relation to the Past Activity of Sgr A∗

Since the similar spectral features in M 0.66−0.02 and M 0.74−0.09 and the similarvariation in M 0.66−0.02 and G 0.570−0.018, we consider these three sources havethe same origin. The variaion time scale and intensity of the 6.40 keV line requiresa bright source with a luminosity of more than 1037 erg s−1 located at the distanceof a few persec from the 6.40 keV sources. The point sources in the Galactic centerregion have been catalogued with many satellites, but no tranjent source brighter

CHAPTER 7. SUMMARY AND DISCUSSION

99

than 1037 erg s−1 in a time scale of ten years has been ever found. Thus we haveconcluded that the most probable source is our Galactic massive black hole Sgr A∗ .
With this scenerio, the light curves of the 6.40 keV line ﬂux in the molecular
clouds can be regarded as the light curve of Sgr A∗ in the past. The required
luminosity to account for the 6.40 keV ﬂux in 2000 is 2 × 1039 ergs s−1 (Murakami
et al. 2001), and decreases to about a half in 2005. Sgr A∗ shows frequent shortterm (∼ 1 hour) and low-level (∼ 10–100 times of the quiescent level) ﬂares (Baganoﬀ
et al. 2001), however, no variation with long time scale (10 years) is ever found in
the present Sgr A∗ . Similar valiabilities are also found in the molecular cloud in the
Sgr A region, M0.11−0.11, and indicate an outburst of Sgr A∗ with a luminosity ofat least 1038 ergs s−1 , which would have occurred about 60 years ago and lasted 2–3
years (Muno et al. 2007). The variabilities of the Sgr B2 clouds also suggest that Sgr
A∗ was in an active phase about 300 years ago, 106 times brighter than the present,
and experienced a giant ﬂare. The X-rays illuminate the molecular clouds after 300
years travel. We now see a decay of the ﬂare as the ﬂuorescence iron emission ofmolecular clouds.

Acknowledgement
First of all, I am deeply grateful to my supervisor, Prof. K. Koyama, Assoc. Prof. T. Tsuru,
and Res. Assoc. H. Matsumoto for their continual guidance, advice and support
throughout my graduate career. I also thank my collaborators, Prof. S. Yamauchi,
Dr. H. Murakami, Mr. Y. Hyodo, Mr. Y. Takikawa and Mr. M. Nobukawa for thier
helpful comments and constructive discussions. I thank all the members of theCosmic-ray laboratory in Kyoto University for the daily discussion and insightfulcomments. I express my gratitude to all the members of the Suzaku team.
Finally I thank all my family and friends for thier encouragement.

100

CHAPTER 7. SUMMARY AND DISCUSSION

101

References
Aharonian, F., et al. 2004, A&A, 425, L13
Aharonian, F., et al. 2006, Nature, 439, 695
Atoyan, A., & Dermer, C. D. 2004, ApJL, 617, L123
Baganoﬀ, F. K., et al. 2001, Nature, 413, 45
Baganoﬀ, F. K., et al. 2003, ApJ, 591, 891
Balick, B., & Brown, R. L. 1974, ApJ, 194, 265
Bautz, M. W., Kissel, S. E., Prigozhin, G. Y., LaMarr, B., Burke, B. E., & Gregory, J. A.
2004, Proc. SPIE, 5501, 111
Borkowski, K. J., Lyerly, W. J., & Reynolds, S. P. 2001, ApJ, 548, 820
Burke, B .E., et al. 1994, IEEE, 41, 375
Carter, J. A., & Read, A. M. 2007, A&A, 464, 1155
Dogiel, V. A., Ichimura, A., Inoue, H., & Masai, K. 1998, PASJ, 50, 567
Eckart, A., Schödel, R., Meyer, L., Trippe, S., Ott, T., & Genzel, R. 2006, A&A, 455, 1
Eisenhauer, F., et al. 2005, ApJ, 628, 246
Garcia, J. D., Fortner, R. J., & Kavanagh, T. M. 1973, Reviews of Modern Physics, 45,
111
Gaume, R. A., Claussen, M. J., de Pree, C. G., Goss, W. M., & Mehringer, D. M. 1995,
ApJ, 449, 663
Gendreau, K. C. 1995, Ph.D. Thesis,
Genzel, R., Schödel, R., Ott, T., Eckart, A., Alexander, T., Lacombe, F., Rouan, D., &Aschenbach, B. 2003, Nature, 425, 934
Ghez, A. M., Salim, S., Hornstein, S. D., Tanner, A., Lu, J. R., Morris, M., Becklin, E. E.,& Duchêne, G. 2005, ApJ, 620, 744
Goldston, J. E., Quataert, E., & Igumenshchev, I. V. 2005, ApJ, 621, 785
Goldwurm, A., Brion, E., Goldoni, P., Ferrando, P., Daigne, F., Decourchelle, A.,
Warwick, R. S., & Predehl, P. 2003, ApJ, 584, 751
Gordon, M. A., Berkermann, U., Mezger, P. G., Zylka, R., Haslam, C. G. T., Kreysa, E.,
Sievers, A., & Lemke, R. 1993, A&A, 280, 208
Green, D. A. 2004, Bulletin of the Astronomical Society of India, 32, 335
Hoﬀmann, D. H. H., Brendel, C., Genz, H., Löw, W., Müller, S., & Richter, A. 1979,
Zeitschrift fur Physik , 293, 187
Inoue, H. 1985, Space Science Reviews, 40, 317
Kaastra, J. S., & Mewe, R. 1993, A&AS, 97, 443
Kosack, K., et al. 2004, ApJL, 608, L97
Koyama, K., Awaki, H., Kunieda, H., Takano, S., & Tawara, Y. 1989, Nature, 339, 603
Koyama, K., Maeda, Y., Sonobe, T., Takeshima, T., Tanaka, Y., & Yamauchi, S. 1996,
PASJ, 48, 249

CHAPTER 7. SUMMARY AND DISCUSSION

102

Koyama, K. et al. 2007a, PASJ, 59, S23,
Koyama, K. et al. 2007b, PASJ, 59, S221
Koyama, K. et al. 2007c, PASJ, 59, S237
Koyama, K. et al. 2007d, PASJ, 59, S245
Koyama, K. et al. 2008, PASJ, accepted (#3188) (astroph/0711.2853)
Liu, S., Petrosian, V., & Melia, F. 2004, ApJL, 611, L101
Maeda, Y., et al. 2002, ApJ, 570, 671
Makishima, K., et al. 1996, PASJ, 48, 171
Markoﬀ, S. 2005, ApJL, 618, L103
Mayer-Hasselwander, H. A., et al. 1998, A&A, 335, 161
Mehringer, D. M., Palmer, P., & Goss, W. M. 1995, ApJS, 97, 497
Mehringer, D. M., & Menten, K. M. 1997, ApJ, 474, 346
Mitsuda, K. et al. 2007, PASJ, 59, S1
Miyazaki, A., & Tsuboi, M. 2000, ApJ, 536, 357
Mori, H., et al.. 2008, PASJ, accepted (#3131)
Morris, M., & Serabyn, E. 1996, ARA&A, 34, 645
Muno, M. P., et al. 2003, ApJ, 589, 225
Muno, M. P., Baganoﬀ, F. K., Brandt, W. N., Park, S., & Morris, M. R. 2007, ApJL, 656,
L69
Murakami, H., Koyama, K., & Maeda, Y. 2001, ApJ, 558, 687
Murakami, H. 2002, Ph.D. Thesis,
Nobukawa, M. et al. 2008, PASJ, accepted (#3110) (astroph/0712.0877)
Oka, T., Hasegawa, T., Sato, F., Tsuboi, M., & Miyazaki, A. 1998, ApJS, 118, 455
Ohashi, T., et al. 1996, PASJ, 48, 157
Park, S., et al. 2005, ApJ, 631, 964
Reid, M. J., Schneps, M. H., Moran, J. M., Gwinn, C. R., Genzel, R., Downes, D., &Roennaeng, B. 1988, ApJ, 330, 809
Rho, J., & Petre, R. 1998, ApJL, 503, L167
Sakano, M., Koyama, K., Murakami, H., Maeda, Y., & Yamauchi, S. 2002, ApJS, 138, 19
Sakano, M., Warwick, R.S., Decourchelle, A., & Predehl, P. 2004, MNRAS, 350, 129
Sato, F., Hasegawa, T., Whiteoak, J. B., & Miyawaki, R. 2000, ApJ, 535, 857
Schwartz, D. A., et al. 2000, Proc. SPIE, 4012, 28
Senda, A., Murakami, H., & Koyama, K. 2002, ApJ, 565, 1017
Serlemitsos, P. J., et al. 1995, PASJ, 47, 105Serlemitsos, P. J., et al. 2007, PASJ, 59, S9Strüder, L., et al. 2001, A&A, 365, L18

CHAPTER 7. SUMMARY AND DISCUSSION

103

Takagi, S., Murakami, H., & Koyama, K. 2002, ApJ, 573, 275
Tanaka, Y., Inoue, H., & Holt, S. S. 1994, PASJ, 46, L37
Tatischeﬀ, V. 2003, in Final Stage of Stellar Evolution, ed. C. Motch & Hameury (EAS
publication Series vol.7), 79
Tawa, N., et al. 2008, PASJ, accepted (#3176)
Tsuboi, M., Handa, T., & Ukita, N. 1999, ApJS, 120, 1
Tsuchiya, K., et al. 2004, ApJL, 606, L115
Turner, M. J. L., et al. 2001, A&A, 365, L27
Valinia, A., Tatischeﬀ, V., Arnaud, K., Ebisawa, K., & Ramaty, R. 2000, ApJ, 543, 733
Weisskopf, M. C., Brinkman, B., Canizares, C., Garmire, G., Murray, S., & Van
Speybroeck, L. P. 2002, PASP, 114, 1
Weisskopf, M. C., et al. 2003, Experimental Astronomy, 16, 1
Yamauchi, S., Kawada, M., Koyama, K., Kunieda, H., & Tawara, Y. 1990, ApJ, 365, 532
Yuan, F., Quataert, E., & Narayan, R. 2004, ApJ, 606, 894
Yusef-Zadeh, F., et al. 2006, ApJ, 644, 198
Yusef-Zadeh, F., Muno, M., Wardle, M., & Lis, D. C. 2007, ApJ, 656, 847

Appendix A
Tables for Spectral Results of
M 0.66−0.02
A.1

Continuum Models

Using NH and Γ values obtained by Suzaku ﬁt, we ﬁt spectra with three types ofcontinuum models (Model II–IV descibed in section 6.1.1).

A.2

Background Analysis of the MOS and the PN
detector

Using the background correction factors derived in section 6.1.2, we multiplied theblank-sky background spectra. We ﬁt source spectra applied with the new backgrounds. The results are shown in table A.4 and A.5.

104

APPENDIX A. TABLES FOR SPECTRAL RESULTS OF M 0.66−0.02

105

Table A.1: The result of spectral ﬁtting for M 0.66−0.02 with Model II

Observatory
SuzakuDetector
XIS(1) NH
5.0+0.5−0.3(2) Γ
3.0+0.3−0.3(3) Fpow
2.0+0.1−0.1—Neutral iron lines—
6399+4(4) E640−3(5) F640
0.66+0.02−0.02(6) E706
7057b(7) F706
0.06+0.04−0.04—Ionized iron lines—
6670+5(8) E667−5(9) σ667
29+12−16(10) F667
0.36+0.02−0.02(11) E697
6967+21−21(12) F697
0.12+0.03−0.02(13) δGain0(ﬁx)χ2 /dof348.41/354

XMM-Newton
MOS
PN
+0.3
4.9+1.1
5.4
−1.0
−0.2
3.0a
1.7+0.1
2.3+0.1−0.1−0.1

XMM-Newton
MOS
PN
+1.9
5.0+1.5
5.8
−1.5
−1.7

ChandraACIS
7.0+1.1
−0.9

ASCA
SIS†
5.0+0.7−1.0

2.6+0.2−0.2

2.3+0.2−0.2

2.5+0.1−0.1

2.2+0.2−0.2

6399a
0.59+0.02−0.05

0.57+0.03−0.04

0.89+0.10−0.10

0.80+0.10−0.10

1.00+0.06−0.05

0.99+0.12−0.13

0.07+0.04−0.04

0.05+0.02−0.03

0.10+0.09−0.09

0.11+0.08−0.06

0.12+0.04−0.06

0.17+0.08
−0.12

6670a
29a
0.30+0.02
0.25+0.03
0.40+0.09
0.24+0.08
0.25+0.04
0.40+0.12
−0.04
−0.03
−0.09
−0.08
−0.03
−0.11
a
6967
0.10c
0.08c
0.13c
0.08c
0.08c
0.13c
0(ﬁx)
23+13
0(ﬁx)
36+1
−18+9
81+20
−4
−28
−3
−19
37.03/41 90.79/69 50.69/39 22.66/28 152.47/145 37.29/38

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(6),(8),(11): Line centroid (eV). (5),(7),(10),(12): Observed line ﬂux (10−4 photons cm−2 s−1 ). (9): Linewidth in gaussian sigma (eV). The other line widths are ﬁxed at zero. (13): Additional gain oﬀset (eV).
†Values are obtained from region ’P’.
aFixed at the value obtained with Suzaku ﬁt.
b
The energy gap between Kα (6400 eV) and Kβ (7058 eV) is ﬁxed at the theoretical value (+658 eV)(Kaastra & Mewe 1993).
c
The intensity of Fe XXVI Kα is ﬁxed at 34% of that of Fe XXV Kα obtained with Suzaku.

APPENDIX A. TABLES FOR SPECTRAL RESULTS OF M 0.66−0.02

106

Table A.2: The result of spectral ﬁtting for M 0.66−0.02 with Model III

Observatory
SuzakuDetector
XIS(1) NH
5.0+0.5−0.3(2) Γ
3.0+0.3−0.3(3) Fpow
2.0+0.1−0.1—Neutral iron lines—
6399+4(4) E640−3(5) F640
0.66+0.02−0.02(6) E706
7057b(7) F706
0.06+0.04−0.04—Ionized iron lines—
6670+5(8) E667−5(9) σ667
29+12−16(10) F667
0.36+0.02−0.02(11) E697
6967+21−21(12) F697
0.12+0.03−0.02(13) δGain0(ﬁx)χ2 /dof348.41/354

XMM-Newton
MOS
PN
5.0a
3.1+0.6
2.1+0.3
−0.5
−0.2
+0.1
1.6−0.1
2.4+0.1−0.1

XMM-Newton
MOS
PN

Chandra
ACIS

ASCA
SIS†

2.0+0.8−0.7
3.0+0.2−0.2

2.0+0.7−0.8
2.5+0.2−0.2

1.5+0.5−0.5
2.7+0.1−0.1

3.0+0.7−1.0
2.2+0.2−0.2

6399a
0.59+0.04−0.05

0.56+0.03−0.05

0.85+0.11−0.11

0.79+0.10−0.10

0.99+0.05−0.05

0.99+0.11−0.13

0.08+0.04−0.04

0.03+0.02−0.03

0.06+0.09−0.06

0.10+0.07−0.07

0.07+0.05−0.04

0.17+0.08
−0.12

6670a
29a
0.30+0.04
0.23+0.03
0.35+0.10
0.23+0.08
0.24+0.04
0.40+0.12
−0.04
−0.03
−0.10
−0.08
−0.04
−0.11
a
6967
0.10c
0.08c
0.13c
0.08c
0.08c
0.13c
0(ﬁx)
23+6
0(ﬁx)
33+4
−19+10
81+35
−4
−27
−3
−20
36.98/41 74.91/69 46.19/39 18.78/28 142.83/145 37.29/38

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(6),(8),(11): Line centroid (eV). (5),(7),(10),(12): Observed line ﬂux (10−4 photons cm−2 s−1 ). (9): Linewidth in gaussian sigma (eV). The other line widths are ﬁxed at zero. (13): Additional gain oﬀset (eV).
†Values are obtained from region ’P’.
aFixed at the value obtained with Suzaku ﬁt.
b
The energy gap between Kα (6400 eV) and Kβ (7058 eV) is ﬁxed at the theoretical value (+658 eV)(Kaastra & Mewe 1993).
c
The intensity of Fe XXVI Kα is ﬁxed at 34% of that of Fe XXV Kα obtained with Suzaku.

APPENDIX A. TABLES FOR SPECTRAL RESULTS OF M 0.66−0.02

107

Table A.3: The result of spectral ﬁtting for M 0.66−0.02 with Model IV

Observatory
SuzakuDetector
XIS(1) NH
5.0+0.5−0.3(2) Γ
3.0+0.3−0.3(3) Fpow
2.0+0.1−0.1—Neutral iron lines—
6399+4(4) E640−3(5) F640
0.66+0.02−0.02(6) E706
7057b(7) F706
0.06+0.04−0.04—Ionized iron lines—
6670+5(8) E667−5(9) σ667
29+12−16(10) F667
0.36+0.02−0.02(11) E697
6967+21−21(12) F697
0.12+0.03−0.02(13) δGain0(ﬁx)χ2 /dof348.41/354

XMM-Newton
MOS
PN
5.0a
3.0a
1.7+0.1
2.2+0.1−0.1−0.1

XMM-Newton
MOS
PN

Chandra
ACIS

ASCA
SIS†

2.6+0.2−0.2

2.2+0.2−0.2

2.2+0.1−0.1

2.2+0.2−0.2

6399a
0.58+0.05−0.04

0.58+0.04−0.05

0.89+0.10−0.10

0.81+0.10−0.09

1.05+0.05−0.06

0.99+0.11−0.13

0.07+0.04−0.04

0.05+0.03−0.04

0.10+0.09−0.09

0.13+0.07−0.07

0.14+0.05−0.04

0.17+0.08
−0.12

6670a
29a
0.30+0.04
0.26+0.04
0.40+0.09
0.26+0.07
0.31+0.04
0.40+0.12
−0.04
−0.04
−0.09
−0.07
−0.04
−0.11
a
6967
0.10c
0.09c
0.13c
0.09c
0.10c
0.13c
0(ﬁx)
28+6
0(ﬁx)
36+6
−18+8
81+22
−7
−28
−11
−20
37.04/42 95.88/70 50.70/40 23.15/29 165.18/146 37.29/39

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(6),(8),(11): Line centroid (eV). (5),(7),(10),(12): Observed line ﬂux (10−4 photons cm−2 s−1 ). (9): Linewidth in gaussian sigma (eV). The other line widths are ﬁxed at zero. (13): Additional gain oﬀset (eV).
†Values are obtained from region ’P’.
aFixed at the value obtained with Suzaku ﬁt.
b
The energy gap between Kα (6400 eV) and Kβ (7058 eV) is ﬁxed at the theoretical value (+658 eV)(Kaastra & Mewe 1993).
c
The intensity of Fe XXVI Kα is ﬁxed at 34% of that of Fe XXV Kα obtained with Suzaku.

APPENDIX A. TABLES FOR SPECTRAL RESULTS OF M 0.66−0.02

108

Table A.4: The XMM-Newton 2004 results of spectral ﬁtting for M 0.66−0.02 withmultiplied blank-sky backgrounds
NH
Γ
Fpow
(1)
(2)
(3)
†
XMM-Newton 2004 MOS
I
...
...
...
II
...
...
...
III
...
...
...
IV
...
...
...
XMM-Newton 2004 PN
I
5.9+0.3
3.0+0.3
2.0+0.1
−1.5
−0.4
−0.1
+0.1
a
3.0
1.9
II
5.6+0.4−0.8−0.1
2.7+0.3
2.0+0.1
III
5.0a
−0.4
−0.1
IV
5.0a
3.0a
1.9+0.1
−0.1
Model

F640
(4)

F706
(5)

F667
(6)

F697
(7)

δGain
(8)

χ2 /dof
(9)

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

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

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

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

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

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

0.57+0.04−0.04
0.57+0.04−0.04
0.58+0.04−0.04
0.58+0.04−0.04

0.03+0.03−0.02
0.04+0.02−0.03
0.04+0.02−0.03
0.04+0.03−0.02

0.25+0.03−0.03
0.25+0.03−0.03
0.26+0.03−0.03
0.27+0.03
−0.03

0.08c
0.08c
0.09c
0.09c

23+6
−4
23+6
−4
17+5
−4
23+7
−4

85.48/68
85.23/69
85.47/69
86.93/70

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(5),(6),(7): Observed line ﬂux of 6.40, 7.06, 6.67, 6.97 keV line (10−4 photons cm−2 s−1 ). (8): Additional
gain oﬀset (eV). (9): χ2 value per Degree of Freedom.
†
No correction is needed.
aFixed at the value obtained with Suzaku ﬁt.
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα determined by the Suzaku ﬁt.

APPENDIX A. TABLES FOR SPECTRAL RESULTS OF M 0.66−0.02

109

Table A.5: The XMM-Newton 2001 results of spectral ﬁtting for M 0.66−0.02 withmultiplied blank-sky backgrounds
NH
Γ
(1)
(2)
XMM-Newton 2001 MOS
I
4.4+8.1
3.9+3.1
−4.4
−2.5
+2.1
3.0a
II
2.9−2.2
4.2+0.8
III
5.0a
−1.2
IV
5.0a
3.0a
XMM-Newton 2001 PN+11.0
6.1+3.9I
10.6−3.9
−3.4
+3.8
II
5.0−3.6
3.0a
3.8+1.9
III
5.0a
−1.5
IV
5.0a
3.0a
Model

Fpow
(3)

F640
(4)

F706
(5)

F667
(6)

F697
(7)

δGain
(8)

χ2 /dof(9)

1.4+0.2−0.2
1.5+0.2−0.2
1.4+0.2−0.2
1.6+0.2−0.2

0.90+0.11−0.11
0.91+0.11−0.11
0.90+0.11−0.11
0.88+0.11−0.10

0.07+0.09−0.07
0.06+0.09−0.06
0.07+0.09−0.07
0.04+0.08−0.04

0.41+0.09−0.09
0.41+0.09−0.09
0.41+0.09−0.09
0.37+0.09
−0.09

0.14c
0.14c
0.14c
0.12c

0(ﬁx)
0(ﬁx)
0(ﬁx)
0(ﬁx)

35.64/38
35.89/39
35.66/39
38.11/40

1.1+0.2−0.2
1.2+0.2−0.2
1.1+0.2−0.2
1.2+0.2−0.2

0.83+0.10−0.10
0.83+0.09−0.10
0.85+0.10−0.10
0.83+0.10−0.09

0.12+0.08−0.07
0.12+0.08−0.08
0.12+0.09−0.07
0.12+0.08−0.08

0.27+0.08−0.08
0.29+0.08−0.07
0.30+0.08−0.08
0.29+0.07
−0.07

0.09c
0.10c
0.10c
0.10c

33+1
−32
21+13
−16
33+6
−27
21+17
−15

17.68/27
19.94/28
19.28/28
19.94/29

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(5),(6),(7): Observed line ﬂux of 6.40, 7.06, 6.67, 6.97 keV line (10−4 photons cm−2 s−1 ). (8): Additional
gain oﬀset (eV). (9): χ2 value per Degree of Freedom.
aFixed at the value obtained with Suzaku ﬁt.
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα determined by the Suzaku ﬁt.

Appendix B
Tables for Spectral Results of
G 0.570−0.018
B.1

Continuum Models

Suzaku spectral analysis mentioned in 6.2.3 requires one additional line component
for G 0.570−0.018. We ﬁt spectra with Model II–IV including this line component.

B.2

Background Analysis of the MOS and the PN
detector

Using the background correction factors derived in section 6.1.2, we multiplied theblank-sky background spectra. We ﬁt source spectra applied with the new backgrounds. The results are shown in table B.4 and B.5.

110

APPENDIX B. TABLES FOR SPECTRAL RESULTS OF G 0.570−0.018

111

Table B.1: The result of spectral ﬁtting for G 0.570−0.018 with Model II plus one
Gaussian function

Observatory
SuzakuXMM-Newton
XMM-Newton
ChandraASCA
Detector
XIS
MOS
PN
MOS
PN
ACIS
SIS
+1.3
+0.8
+0.6
+0.9
+1.1
+1.2
(1) NH
1.5−1.0
3.7−0.8
1.8−0.6
5.0−1.0
2.6−1.3
2.9−1.2
2.9+1.9
−1.9
+0.6
a
(2) Γ
2.6−0.5
2.6(3) Fpow
1.1+0.2
1.1+0.1
1.3+0.1
1.8+0.1
1.2+0.1
1.1+0.1
0.9+0.1
−0.2
−0.1
−0.1
−0.1
−0.1
−0.1
−0.1
—Neutral iron lines—(4) E640
6400a(5) F640
0.16+0.02
0.11+0.04
0.13+0.03
0.17+0.06
0.19+0.06
0.26+0.04
0.11+0.054
−0.03
−0.04
−0.03
−0.06
−0.06
−0.04
−0.05
a
(6) E706
7058(7) F706
0.02b
0.01b
0.02b
0.02b
0.02b
0.03b
0.01b
—Ionized iron lines—(8) E667
6665a(9) σ667
27a(10) F667
0.29+0.02
0.21+0.04
0.22+0.03
0.17+0.06
0.23+0.06
0.25+0.03
0.20+0.05
−0.02
−0.04
−0.03
−0.06
−0.06
−0.03
−0.05
(11) E697
6966a(12) F697
0.09c
0.07c
0.07c
0.05c
0.07c
0.08c
0.06c
+25
d
(13) E651
6505−316505+0.03(14) F651
0.07−0.02
0.05d
0.05d
0.04d
0.06d
0.06d
0.05d
(15) δGain
0(ﬁx)
0(ﬁx)
35(ﬁx)0(ﬁx)
33(ﬁx)-18(ﬁx)71(ﬁx)
192.33/181 58.38/52 73.07/79 48.12/26 21.35/16 82.71/94 9.21/19
χ2 /dofNote—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(6),(8),(11),(13): Line centroid (eV). (5),(7),(10),(12),(14): Observed line ﬂux (10−4 photons cm−2 s−1 ).
(9): Line width in gaussian sigma (eV). The other line widths are ﬁxed at zero. (15): Additional gain oﬀset(eV).
aFixed at the value obtained with Suzaku ﬁt.
b
The intensity of Kβ is ﬁxed at 12.5% of that of Kα (Kaastra & Mewe 1993).
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα obtained with Suzaku.

APPENDIX B. TABLES FOR SPECTRAL RESULTS OF G 0.570−0.018

112

Table B.2: The result of spectral ﬁtting for G 0.570−0.018 with Model III plus one
Gaussian function

Observatory
SuzakuXMM-Newton
XMM-Newton
Chandra ASCA
Detector
XIS
MOS
PN
MOS
PN
ACIS
SIS
+1.3
a
(1) NH
1.5−1.0
1.5
+0.6
(2) Γ
2.6−0.5
2.0+0.6
1.8+0.8
0.6+0.7
2.0+0.8
1.8+0.6
1.7+0.8
−0.7
−0.6
−0.6
−0.8
−0.5
−0.7
+0.1
+0.1
+0.1
+0.1
+0.1
+0.1
(3) Fpow
1.1+0.2
1.1
1.3
1.9
1.2
1.0
0.9
−0.2
−0.1
−0.1
−0.1
−0.1
−0.1
−0.1
—Neutral iron lines—(4) E640
6400a(5) F640
0.16+0.02
0.11+0.04
0.13+0.03
0.20+0.04
0.19+0.06
0.26+0.04
0.11+0.05
−0.03
−0.03
−0.03
−0.08
−0.06
−0.04
−0.05
a
(6) E706
7058(7) F706
0.02b
0.01b
0.02b
0.02b
0.02b
0.03b
0.01b
—Ionized iron lines—(8) E667
6665a(9) σ667
27a(10) F667
0.29+0.02
0.22+0.04
0.22+0.03
0.18+0.06
0.23+0.06
0.25+0.03
0.20+0.05
−0.02
−0.04
−0.03
−0.06
−0.06
−0.03
−0.05
(11) E697
6966a(12) F697
0.09c
0.07c
0.07c
0.05c
0.07c
0.08c
0.06c
+25
d
(13) E651
6505−316505+0.03(14) F651
0.07−0.02
0.06d
0.05d
0.04d
0.06d
0.06d
0.05d
(15) δGain
0(ﬁx)
0(ﬁx)
35(ﬁx)0(ﬁx)
33(ﬁx)-18(ﬁx)71(ﬁx)
192.33/181 62.86/52 71.81/79 41.20/26 21.01/16 82.39/94 7.24/19
χ2 /dofNote—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(6),(8),(11),(13): Line centroid (eV). (5),(7),(10),(12),(14): Observed line ﬂux (10−4 photons cm−2 s−1 ).
(9): Line width in gaussian sigma (eV). The other line widths are ﬁxed at zero. (15): Additional gain oﬀset(eV).
aFixed at the value obtained with Suzaku ﬁt.
b
The intensity of Kβ is ﬁxed at 12.5% of that of Kα (Kaastra & Mewe 1993).
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα obtained with Suzaku.

APPENDIX B. TABLES FOR SPECTRAL RESULTS OF G 0.570−0.018

113

Table B.3: The result of spectral ﬁtting for G 0.570−0.018 with Model IV plus one
Gaussian function

Observatory
SuzakuDetector
XIS(1) NH
1.5+1.3−1.0(2) Γ
2.6+0.6−0.5(3) Fpow
1.1+0.2−0.2—Neutral iron lines—(4) E640
6400a(5) F640
0.16+0.02−0.03(6) E706
7058a(7) F706
0.02b—Ionized iron lines—(8) E667
6665a(9) σ667
27a(10) F667
0.29+0.02−0.02(11) E697
6966a(12) F697
0.09c(13) E651
6505+25−31(14) F651
0.07+0.03−0.02(15) δGain0(ﬁx)192.33/181
χ2 /dof

XMM-Newton
MOS
PN
a
1.5
2.6a
1.0+0.1
1.3+0.1−0.1−0.1

XMM-Newton
MOS
PN

Chandra
ACIS

ASCA
SIS

1.5+0.1−0.1

1.2+0.1−0.1

1.2+0.1−0.1

0.8+0.1−0.1

0.12+0.04−0.03

0.14+0.03−0.03

0.22+0.06−0.06

0.20+0.06−0.06

0.27+0.03−0.03

0.12+0.05
−0.05

0.01b

0.02b

0.03b

0.02b

0.03b

0.01b

0.23+0.04−0.03

0.22+0.03−0.03

0.24+0.05−0.05

0.24+0.05−0.05

0.26+0.03−0.03

0.22+0.04
−0.04

0.07c
0.07c
0.07c
0.08c
0.08c
0.07c
d
6505
0.06d
0.06d
0.06d
0.06d
0.06d
0.05d
0(ﬁx)
35(ﬁx)0(ﬁx)
33(ﬁx)-18(ﬁx)71(ﬁx)
64.69/53 73.28/80 51.44/27 22.09/17 86.00/95 10.54/20

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(6),(8),(11),(13): Line centroid (eV). (5),(7),(10),(12),(14): Observed line ﬂux (10−4 photons cm−2 s−1 ).
(9): Line width in gaussian sigma (eV). The other line widths are ﬁxed at zero. (15): Additional gain oﬀset(eV).
aFixed at the value obtained with Suzaku ﬁt.
b
The intensity of Kβ is ﬁxed at 12.5% of that of Kα (Kaastra & Mewe 1993).
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα obtained with Suzaku.

APPENDIX B. TABLES FOR SPECTRAL RESULTS OF G 0.570−0.018

114

Table B.4: The XMM-Newton 2004 results of spectral ﬁtting for G 0.570−0.018 withmultiplied blank-sky backgrounds
NH
Γ
Fpow
(1)
(2)
(3)
†
XMM-Newton 2004 MOS
I
...
...
...
II
...
...
...
III
...
...
...
IV
...
...
...
XMM-Newton 2004 PN
I
0.4+1.7
1.8+0.8
1.3+0.1
−0.4
−0.6
−0.1
+0.1
a
2.6
1.3
II
1.8+0.6−0.6−0.1
2.2+0.5
1.3+0.1
III
1.5a
−0.4
−0.1
IV
1.5a
2.6a
1.3+0.1
−0.1
Model

F640
(4)

F706
(5)

F667
(6)

F697
(7)

F651
(8)

δGain
(9)

χ2 /dof
(10)

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

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

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

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

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

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

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

0.14+0.03−0.03
0.13+0.03−0.03
0.13+0.03−0.03
0.14+0.03
−0.03

0.02b
0.02b
0.02b
0.02b

0.22+0.03−0.03
0.22+0.03−0.03
0.21+0.03−0.03
0.22+0.03
−0.03

0.07c
0.07c
0.07c
0.07c

0.06d
0.05d
0.05d
0.06d

35(ﬁx)35(ﬁx)35(ﬁx)35(ﬁx)

70.76/78
73.07/79
71.81/79
73.28/80

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(5),(6),(7),(8): Observed line ﬂux of 6.40, 7.06, 6.67, 6.97, 6.51 keV line (10−4 photons cm−2 s−1 ). (9):Additional gain oﬀset (eV). (10): χ2 value per Degree of Freedom.
†
No correction is needed.
aFixed at the value obtained with Suzaku ﬁt.
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα determined by the Suzaku ﬁt.

APPENDIX B. TABLES FOR SPECTRAL RESULTS OF G 0.570−0.018

115

Table B.5: The XMM-Newton 2001 results of spectral ﬁtting for G 0.570−0.018 withmultiplied blank-sky backgrounds
NH
Γ
Fpow
(1)
(2)
(3)
†
XMM-Newton 2004 MOS
I
1.3+1.4
0.5+1.2
1.9+0.1
−1.3
−0.9
−0.1
+0.1
a
2.6
1.8
II
5.0+0.9
−1.0
−0.1
+0.7
0.6−0.6 1.9+0.1
III
1.5a
−0.1
IV
1.5a
2.6a
1.5+0.1−0.1
XMM-Newton 2004 PN
I
1.5+3.6
2.0+1.6
1.2+0.1
−1.5
−1.2
−0.1
+0.1
a
2.6
1.2
II
2.6+1.1−1.3−0.1
2.0+0.8
1.2+0.1
III
1.5a
−0.8
−0.1
IV
1.5a
2.6a
1.2+0.1
−0.1
Model

F640
(4)

F706
(5)

F667
(6)

F697
(7)

F651
(8)

δGain
(9)

χ2 /dof(10)

0.18+0.06−0.06
0.17+0.06−0.06
0.20+0.04−0.08
0.22+0.06
−0.06

0.02b
0.02b
0.02b
0.03b

0.18+0.06−0.06
0.17+0.06−0.06
0.18+0.06−0.06
0.24+0.05
−0.05

0.05c
0.05c
0.05c
0.07c

0.04d
0.04d
0.04d
0.06d

0(ﬁx)
0(ﬁx)
0(ﬁx)
0(ﬁx)

40.89/25
48.12/26
41.20/26
51.44/27

0.19+0.06−0.06
0.19+0.06−0.06
0.19+0.06−0.06
0.20+0.06
−0.06

0.02b
0.02b
0.02b
0.02b

0.23+0.06−0.06
0.23+0.06−0.06
0.23+0.06−0.06
0.24+0.05
−0.05

0.07c
0.07c
0.07c
0.08c

0.06d
0.06d
0.06d
0.06d

33(ﬁx)
33(ﬁx)
33(ﬁx)
33(ﬁx)

21.01/15
21.35/16
21.01/16
22.09/17

Note—Parentheses indicate the 90% conﬁdence limit. (1): Absorption column density (1023 H cm−2 ). (2):Photon index. (3): Observed ﬂux of a powerlaw continuum in the 5–8 keV band (10−4 photons cm−2 s−1 ).
(4),(5),(6),(7),(8): Observed line ﬂux of 6.40, 7.06, 6.67, 6.97, 6.51 keV line (10−4 photons cm−2 s−1 ). (9):Additional gain oﬀset (eV). (10): χ2 value per Degree of Freedom.
†
No correction is needed.
aFixed at the value obtained with Suzaku ﬁt.
c
The intensity of Fe XXVI Kα is ﬁxed at 30% of that of Fe XXV Kα determined by the Suzaku ﬁt.