X-ray Study of Neutral Iron Line Emission
in the Galactic Ridge:
Contribution of Low-Energy Cosmic Rays

Kumiko Nobukawa (Kawabata)
Department of Physics, Graduate School of Science, Kyoto University
Kitashirakawa Oiwake-cho, Sakyo-ku, Kyoto, 606-8502, Japan
kumiko@cr.scphys.kyoto-u.ac.jp
February 10, 2016

Abstract
The Galactic diﬀuse X-ray emission (GDXE) is unresolved X-rays prevailing along the Galactic
plane (e.g. Worrall et al. 1982, Astrophysical Journal 255, 111–121). One of the most remarkablefeatures of the GDXE is strong K-shell lines of neutral, helium-like and hydrogen-like iron at 6.4 keV,
6.7 keV and 7.0 keV, respectively. A popular idea for the origin of the GDXE is a superposition ofmany faint point sources, such as cataclysmic variables (CVs) and coronally active binaries (ABs)because more than 80% of the GDXE has been resolved into point-like sources at a small area near the
Galactic center (Revnivtsev et al. 2009, Nature, 458, 7242, 1142–1144). However, the spectra of the
GDXE are diﬀerent from position to position (Uchiyama et al. 2013, Publications of the Astronomical
Society of Japan, 65, 19). If many point sources contribute to the GDXE, the spectra should be spatially
uniform. Previous debate on the GDXE origin has been based on the results of limited spectral andspatial resolution. In this dissertation, using the Suzaku satellite, we have measured the detailed
spatial distributions (the scale heights) and spectral features (the equivalent widths) of the 6.4 keV,
6.7 keV, and 7.0 keV lines separately. The scale height of the 6.7 keV line is roughly consistent
with those of CVs and ABs, while that of the 6.4 keV line is far smaller than those of the CVs andABs, and similar to molecular clouds. Furthermore clear enhancements of the 6.4 keV line relative
to the 6.7 keV line are found at 1.◦ 5 < l < 3.◦ 5 (GC east) and l = 20◦ regions. These enhancements
are associated with molecular gas. The X-ray spectra of the enhanced regions show large equivalent
widths of > 300 eV (Nobukawa et al. 2015, The Astrophysical Journal Letters, 807, L10). Even in the
GDXE of non-excess areas, an assembly of CVs and ABs does not reproduce the spectra; at least ahalf of the 6.4 keV line emission along the Galactic plane cannot be explained by point sources. Thus
additional contribution of other sources with a strong 6.4 keV line is required. Since the 6.4 keV linefrom molecular gas irradiated by X-rays, cosmic-ray protons or electrons has a large equivalent width(> 300 eV), the 6.4 keV line should be dominated by the diﬀuse emission from molecular gas. The
scattered photons from the interstellar gas originated from bright X-ray sources are less than 10%of the total 6.4 keV line emission. Therefore the most possible origin is low-energy cosmic rays
(LECRs). The energy densities of the LECRs would be ∼ 10 eV cm−3 and ∼ 0.01 eV cm−3 for protons
and electrons, respectively. Our study demonstrates that the 6.4 keV line can be a unique probe toinvestigate the LECRs.

3

Acknowledgement
I would like to express my special appreciation and thanks to my advisor, Professor Emeritus Dr.
Katsuji Koyama for providing comprehensive guidelines and for encouraging my research throughout my Ph.D. years. His advice has consistently been based on far-sighted vision. I consider it as agreat opportunity to do my Ph.D research under his guidance and to learn from his research expertise.
I would like to thank my supervisor, Professor Dr. Takeshi Tsuru, who has guided me through my
Ph.D. years and provided me with the valuable opportunity to develop the SXI onboard ASTRO-H.
The experience of developing the SXI occupies an important place in my career. I have learned to
combine observational astrophysics research and detector development.
I am vey grateful to the members of my Ph.D. committee for their time and valuable feedback
on a preliminary version of this dissertation. I would particularly like to acknowledge Professor Dr.
Toru Tanimori and Dr. Hidetoshi Kubo, who have also provided valuable comments in the laboratory
colloquiums.
I gratefully acknowledge my research collaborators. Professor Dr. Shigeo Yamauchi at Nara
Women’s University analyzed the Suzaku data with me and provided a lot of information on previous studies. Dr. Kazufumi Torii and Professor Dr. Yasuo Fukui at Nagoya University provided
the NANTEN data and discussed with me which molecular clouds corresponds to the enhancementsof the 6.4 keV line and what is the origin of low-energy cosmic rays. I am also very grateful to Dr.
Takaaki Tanaka. His comments and calculation based on his expertise in cosmic-ray accelerationin SNRs have been valuable and they strengthen the discussion in my Ph.D. research. He has also
corrected English errors in this dissertation.
I am very thankful to members of the Cosmic-Ray laboratory for all the great times that we haveshared.
I am deeply thankful to my parents for their love, support and encouragements. I am very muchindebted to my husband Dr. Masayoshi Nobukawa, who has supported me in my personal andacademic life. Without him, I would never have been complete my Ph.D. nor this dissertation.
This work is supported by Research Fellowships of the Japan Society for the Promotion of Sciencefor Young Scientists.
Kumiko K. Nobukawa
Kyoto
February 2016

5

Contents
1 Review
1
1.1 Interstellar medium and spiral arms in the Milky Way galaxy . . . . . . . . . . . . . . . 1
1.1.1 Spiral arm structures . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1
1.1.2 Interstellar medium (ISM) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1
1.2 Galactic diﬀuse X-ray emission with iron K-shell lines . . . . . . . . . . . . . . . . . . . 5
1.3 Debates over the origin of the iron K-shell lines in the GDXE . . . . . . . . . . . . . . . 7
1.4 6.4 keV clumps in the Galactic center region . . . . . . . . . . . . . . . . . . . . . . . . . 11
2 Introduction
15
2.1 Problems in the previous study on the iron lines in the GDXE . . . . . . . . . . . . . . . 15
2.2 Aim of this dissertation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17
3 Instrument
3.1 Suzaku satellite . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2 X-ray Telescope (XRT) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.1 Mechanism and design . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.2 Performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.3 Stray light . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3 X-ray Imaging Spectrometer (XIS) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.1 X-ray CCDs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.2 Energy response . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.3 Performance and calibration . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.4 Non X-ray background . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

19
19
19
19
21
23
27
27
29
30
32

4 Observations and Data Reductions

35

5 Results I: Global structure of the GDXE
5.1 Analytical procedure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.2 Latitudinal distribution of the GDXE . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3 Equivalent widths of the iron K-shell lines . . . . . . . . . . . . . . . . . . . . . . . . . .
5.4 Intensity profile along the Galactic longitude . . . . . . . . . . . . . . . . . . . . . . . .

41
41
42
46
46

6 Results II: Local enhancements in the 6.4 keV line
49
6.1 Estimation of stray light . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 49
6.1.1 Stray light in the GC east region . . . . . . . . . . . . . . . . . . . . . . . . . . . . 49i

CONTENTS

ii

6.1.2 Stray light in the l = −20◦ region . . . . . . . . . . . . . . . . . . . . . . . . . . . 51
6.2 Excess in the GC east . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 53
6.3 Excess in the l = −20◦ region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 56
7 Discussion
7.1 Origin of the 6.4 keV enhancements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.1.1 The GC east region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.1.2 The l = −20◦ region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.2 Scale heights and equivalent widths of the GDXE . . . . . . . . . . . . . . . . . . . . . .
7.2.1 Comparison with previous study . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.2.2 Comparison with point sources . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.2.3 GRXE . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.2.4 GCXE . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.2.5 GBXE . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.3 Where low-energy cosmic rays come from? . . . . . . . . . . . . . . . . . . . . . . . . . .

59
59
61
63
64
64
65
67
69
69
69

8 Prospect for ASTRO-H

73

9 Conclusion

75

Appendix: X-ray image and selected region85

List of Figures
1.1
1.2
1.3
1.4

Sketch of face-on view of the Milky Way . . . . . . . . . . . . . . . . . . . . . . . . . . .
CO velocity-integrated map and longitude-velocity map in the Milky Way . . . . . . . .
Spectrum taken from the Galactic plane observed by Tenma . . . . . . . . . . . . . . . .
Intensity distribution of the 6.7 keV line along the Galactic plane taken from a Ginga
survey . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
1.5 Spectrum of the GCXE and GRXE observed by ASCA . . . . . . . . . . . . . . . . . . . .
1.6 X-ray spectra of the Galactic center and ridge observed by Suzaku. . . . . . . . . . . . .
1.7 Profiles of the surface brightness of the 6.7 keV line along the Galactic plane measured
by RXTE/PCA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
1.8 Spectra extracted by Chandra near the Galactic center (l = 0.◦ 1, b = −1.◦ 4) . . . . . . . .
1.9 Fraction of resolved X-ray emission around the 6.7 keV iron emission line as a functionof the limiting source flux/luminosity . . . . . . . . . . . . . . . . . . . . . . . . . . . .
1.10 Longitude distributions of the 6.4 keV and 6.7 keV lines in the Galactic center region .
1.11 Comparison between an X-ray image in the 6.4 keV and a CS J = 1–0 map in the
Galactic center region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
1.12 Surface brightness maps of the iron lines obtained with Suzaku, XMM-Newton, Chandra, and ASCA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
1.13 Equivalent widths of K-shell line of various neutral atoms in the GC . . . . . . . . . . .

2
4
6
6
7
7
9
10
11
12
12
13
14

2.1 6.7 keV- and 7.0 keV-line intensity distributions along the Galactic longitude as a function of the distance from Sgr A∗ . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17
2.2 6.7 keV line intensity distribution along the Galactic longitude and Galactic latitude
as a function of the distance from Sgr A∗ . . . . . . . . . . . . . . . . . . . . . . . . . . . 17
3.1 Schematic picture of the Suzaku satellite and side view of the instruments and telescopes 20
3.2 Picture of an XRT module . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 21
3.3 Schematic view of an X-ray telescope composed of nested reflectors . . . . . . . . . . . 21
3.4 PSF and EEF of the XRT module . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 22
3.5 Vignetting of the four XRT modules using the observation data of the Crab Nebula . . 22
3.6 Reflection paths causing stray light . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 23
3.7 Wide-range vignetting function of the XRT . . . . . . . . . . . . . . . . . . . . . . . . . 24
3.8 Mosaic images of oﬀset observations of Crab . . . . . . . . . . . . . . . . . . . . . . . . . 25
3.9 Stray patterns for a stray source . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 26
3.10 Schematic view of the XIS CCD . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 28
iii

LIST OF FIGURES

iv

3.11 Picture and cross section of the XIS sensor . . . . . . . . . . . . . . . . . . . . . . . . . .
3.12 Picture of the XIS CCD installed in the base . . . . . . . . . . . . . . . . . . . . . . . . .
3.13 Eﬀective area of one XRT+XIS system . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.14 Energy-resolution history of w Kα line for calibrated data of the XIS . . . . . . . . . . .
3.15 Gain history of manganese Kα line for uncalibrated data of the XIS . . . . . . . . . . . .
3.16 Gain history of manganese Kα line for calibrated data of the XIS . . . . . . . . . . . . .
3.17 Energy-resolution history of manganese Kα line for calibrated data of the XIS . . . . .
3.18 Background counting rate as a function of energy for the XIS and other instruments . .
3.19 Spectra of the NXB for the XIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.20 Coordinate and cut-oﬀ rigidity dependency of the NXB. . . . . . . . . . . . . . . . . . .

28
29
29
30
31
31
32
33
34
34

5.1 Example of a spectrum fitted by the phenomenological model and the CXB . . . . . . .
5.2 Latitudinal distributions of the 6.4 keV line and 6.7 keV line flux in the regions of|l| < 0.◦ 5, l = 358.◦ 5, l = 356.◦ 0–356.◦ 4 and 10◦ < |l| < 30◦ . . . . . . . . . . . . . . . . . . .
5.3 Latitudinal distributions of the 7.0 keV line and the 5–8 keV flux in the regions of|l| < 0.◦ 5, l = 358.◦ 5, l = 356.◦ 0–356.◦ 4, and 10◦ < |l| < 30◦ . . . . . . . . . . . . . . . . . . .
5.4 Equivalent width diagram of the 6.4 keV and 6.7 keV lines in the GCXE, GBXE, and
GRXE . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.5 Intensity profiles of X-rays along the Galactic plane . . . . . . . . . . . . . . . . . . . . .

42

6.1 Stray light pattern of the GC east region with the position of the stray light source
GX 3+1. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.2 Stray-light spectrum of GX 3+1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.3 Stray light pattern of the l = −20◦ region with the position of the stray light source
GX 340+0. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.4 Stray-light spectrum of GX 340+0 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.5 Intensity profiles of X-rays and 12 CO molecular clouds in −4.◦ 5 < l < 4.◦ 5 . . . . . . . . .
6.6 Integrated intensity map of 12 CO in −4.◦ 5 < l < 4.◦ 5 . . . . . . . . . . . . . . . . . . . . .
6.7 X-ray spectra extracted from the east and west sides and X-ray spectrum of the excess
emission in the east . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.8 Integrated intensity map of 12 CO around the −20◦ region . . . . . . . . . . . . . . . . .
6.9 Channel map of 12 CO around the l = −20◦ region . . . . . . . . . . . . . . . . . . . . . .
6.10 X-ray spectra extracted from l = 20◦ and the other regions on the Galactic plane andthe X-ray spectrum of the excess emission at l = 20◦ . . . . . . . . . . . . . . . . . . . . .
7.1
7.2
7.3
7.4

43
44
47
48

50
51
52
52
53
53
55
56
57
58

Equivalent width of the 6.4 keV line as a function of particle index or photon index . . 60
Cross section for producing the 6.4 keV line by the impact of fast electrons and protons 61
Calculated photon index of radiated X-ray as a function of the particle index . . . . . . 62
Latitudinal distribution of integrated intensity of 12 CO . . . . . . . . . . . . . . . . . . 66

8.1 Simulation of an SXS/ASTRO-H observation in the GC east region . . . . . . . . . . . . 74

List of Tables
4.1 Observation log. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 36
5.1 Scale height of the GDXE . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 45
6.1 Best-fit parameters of the X-ray spectra from the east and west sides . . . . . . . . . . . 55
6.2 Best-fit parameters of the X-ray spectra from the l = −20◦ regions and the other region
on the Galactic plane . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 58
7.1 Equivalent widths of CVs and ABs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 67
7.2 Scale heights of CVs, ABs, and molecular gas (12 CO) . . . . . . . . . . . . . . . . . . . . 68
8.1 Simulation result of an SXS/ASTRO-H observation in the GC east region . . . . . . . . 74v

Chapter 1

Review
1.1 Interstellar medium and spiral arms in the Milky Way galaxy
The Milky Way, where we live, is the closest among all galaxies in the universe, and hence a basic
target to study and understand an incredible number of galaxies which exist in the universe. Many
researchers have observed the Galaxy in the optical band since the days of Galileo Galilei. However,
optical surveys cannot see far into the Galactic plane because of large extinction due to interstellar
dust and gas. Thanks to development in technology, we can today observe the Galaxy by various
wavelengths other than the optical band: radio, far and near infrared, X-ray, and gamma-ray.
The Milky Way is classed into a bar spiral galaxy. In the nucleus of the Milky Way, a supermassive
black hole with a mass of ∼ 4 × 106 M⊙ 1 , Sagittarius A* (Sgr A*), has been found (Genzel et al., 2000;
Ghez et al., 2000, 2008). Sgr A* sits at a dynamical center of the Galactic gravity (e.g. Ghez et al.,2008). The main part of the Galaxy is the disk, where stars and gas are concentrated in. The Sun,
which is located at a distance between 7.5 kpc and 8.5 kpc from the Galactic center, is also placed in
the disk.

1.1.1 Spiral arm structures
Figure 1.1 shows an artists concept of a face-on view of the Milky Way. Spiral structures in agalaxy, which are called spiral arms, are regions where star formation actively occurs. The Milky Way
galaxy is found to have four principal spiral arms from star, interstellar-gas, and dust distributions:
the Scutum-Centaurus arm, the Sagittarius arm, the Norma arm, and the Perseus arm. In addition,the (Near and Far) 3 kpc arm, which could be a ring-like structure, also has been reported (e.g.
Van Woerden and Rougoor, 1957; Dame and Thaddeus, 2008).

1.1.2 Interstellar medium (ISM)
The space between stars in the disk is filled with the interstellar medium (ISM), which is mainlycomposed of hydrogen (neutral atoms, molecules, or ions).
1 The symbol of M represents the solar mass of 1.99 × 1033 g.
⊙

1

CHAPTER 1. REVIEW

2

Figure 1.1. Sketch of face-on view of the Milky Way taken from Churchwell et al. (2009).

Neutral hydrogen atoms (HI gas)
Cold (∼100 K) neutral hydrogen atoms can be observed by a radio emission line at 21 cm. Sinceelectrons and protons have intrinsic magnetic moments, interaction between those magnetic dipole
causes the hyperfine structure. The energy state of the electron spinning parallel has slightly higher
than that of the electron spinning antiparallel. When an electron transition occurs from the higher
to the lower energy orbit, radiation at a wavelength of 21 cm is emitted from the atom. This is anextremely rare phenomenon; its lifetime is 1.1 × 107 years. However, there are an enormous amountof hydrogen atoms in the interstellar space, and thus a considerable number of transition occur. The
HI gas is distributed through the Milky Way galaxy, and traces the spiral-arm structure of the MilkyWay. If we assume that most of gas follows an circular orbit around the Galactic center and know
the rotation curve for the Galaxy (e.g. Fich et al., 1989), we obtain the relation between the radial
velocity along the line of sight and distance from the Sun. Combining this relation with the HI line
spectra, which reflect motion of the gas (Doppler eﬀect), we can yield spatial distribution of HI gas
in the Milky Way (e.g. Oort et al., 1958; Nakanishi and Sofue, 2003). The vertical thickness of the
HI gas increases with the Galactocentric distance; the scale height at the Galactocentric distance of∼ 1 kpc is ∼ 100 pc while that of ∼ 10 kpc is ∼ 300 pc (Nakanishi and Sofue, 2003).

1.1. INTERSTELLAR MEDIUM AND SPIRAL ARMS IN THE MILKY WAY GALAXY

3

Molecular clouds
When the density of hydrogen becomes higher, hydrogen exists as molecular gas (H2 ). A hydrogen molecule is a homonuclear linear molecule with no permanent dipole moment, and hencecannot radiate emission in low temperature of ∼10 K. The line emission from carbon-monoxide (CO)is widely used as a main tracer of molecular gas. The transitions of CO is more frequent than thatof H2 and the emission lines are detected even in quite tenuous molecular gas. Furthermore, a proportional relationship between the intensity of the CO line and amount of H2 is observationallyconfirmed. A factor to convert the 12 CO intensity (K km s−1 ) to the H2 column density NH (cm−2 )is called X-factor. The X-factor in the solar neighborhood of ∼ 1.8 × 1020 is reported (e.g. Dameet al., 2001). Torii et al. (2010) obtained the X-factor of ∼ 0.7 × 1020 in the Galactic center (GC).
Figure 1.2 shows a CO velocity-integrated map and longitude-velocity map in the Milky Way (Dameet al., 2001). The GC region represents the strongest concentration of molecular gas and a large velocity dispersion in the CO maps. This corresponds to the central molecular zone (CMZ) of about
200 pc radius, in which the molecular gas is characterized by large densities, high temperatures, andapparently strong magnetic fields (e.g. Morris and Serabyn, 1996).
Ionized hydrogen (HII gas)
The ionization energy of an electron in hydrogen atoms is 13.6 eV. Hydrogen atoms are ionizedby photoionization by ultraviolet, collisional ionization by high energy particles, and thermal ionization. Ionized hydrogen (HII) regions in the ISM are ionized by ultraviolet from young massive (O- or
B-type) stars which were formed inside the giant molecular clouds. HII regions trace star formation
at the present time, and have the temperature of 104 K typically. Also young white dwarfs radiate
ultraviolet. Old low-mass (1–8M⊙ ) stars evolve to red giants and finally to white dwarfs. Youngwhite dwarfs are hot enough to ionize the stellar envelope material that had been ejected during thered giant phase. This phenomenon is called as planetary nebulae.
There exists much higher temperature (≳ 106 K) gas emitting X-rays in the Milky Way. We will
describe the hot gas (plasma) in the next section.

Figure 1.2. (a) Velocity-integrated CO map and (b) longitude-velocity map of Milky Way. Both figures are taken from Dame et al. (2001).

(b)

(a)

4

CHAPTER 1. REVIEW

1.2. GALACTIC DIFFUSE X-RAY EMISSION WITH IRON K-SHELL LINES

5

1.2 X-ray view of the Milky Way galaxy:
the Galactic diﬀuse X-ray emission (GDXE) with iron K-shell lines
X-ray observations by the HEAO 1 satellite discovered unresolved X-ray emission in the 2–10 keVband with the half thickness of ∼ 240 pc in the inner disk (50◦ < l < 120◦ and −130◦ < l < −70◦ )2 of
the Milky Way (Worrall et al., 1982). This emission has been called the Galactic ridge X-ray emission(GRXE). Worrall et al. (1982) estimated the luminosity of the GRXE within a 7.8 kpc radius of the
Galactic disk and obtained ∼ 1.4 × 1038 erg s−1 in the 2–10 keV band. They argued that the origin of
the X-ray emission is unresolved point sources. The EXOSAT found that the GRXE extends from the
GC to |l| = 40◦ with the scale height of ∼ 2◦ (Warwick et al., 1985).
The Japanese X-ray astronomy satellite Tenma detected an intense iron line in the spectra takenfrom the Galactic ridge (Koyama et al., 1986b). An example of the spectra is shown in figure 1.3. Theline energy of 6.7 keV is consistent with the fluorescence line from helium-like iron, and therefore
the line indicates that the GRXE is composed of an optical thin thermal plasma.
Performing a survey observation along the Galactc plane, the Ginga satellite discovered that theiron line intensity has a large peak in the GC region (Koyama et al., 1989; Yamauchi et al., 1990;
Yamauchi and Koyama, 1993). The intense line-emitting region is distributed within the angularsize of 1.◦ 8 × 1.◦ 0 around Sgr A* (Yamauchi et al., 1990). The intense X-ray emission from the GC has
been referred to as the Galactic center X-ray emission (GCXE).
The ASCA satellite separated the previously observed iron line at 6.7 keV in the GCXE into threeiron K-shell lines: the 6.4 keV line from neutral iron, the 6.7 keV line from helium-like iron, and the
7.0 keV line from hydrogen-like iron. The emission mechanism of the 6.4 keV line diﬀers from thatof the 6.7 keV and 7.0 keV lines. The 6.7 keV and 7.0 keV lines are due to collisional excitation in
optically-thin thermal plasma. On the other hand, the 6.4 keV line originates from photoionization
by external X-rays or inner-shell ionizaion by high-energy particles. The external X-rays and particles
should have higher energy than the neutral-iron K edge at 7.1 keV. Figure 1.5a shows a spectrum
extracted from the GC obtained by ASCA. In the GRXE, the iron K-shell line was resolved into thethree iron lines for the first time by Suzaku (Ebisawa et al., 2008).
ASCA revealed that the GRXE cannot be explained by one-temperature plasma model. The GRXE
spectrum exhibits K-shell lines from helium-like magnesium, silicon, sulfur, and iron as is shown infigure 1.5b, and is fairly fitted by two-temperature plasma with the temperatures of ∼ 0.8 keV and∼ 7 keV (Kaneda et al., 1997); the former temperature plasma is required to explain the magnesium, silicon, and sulfur lines while the latter temperature plasma is for the iron lines. Similar
two-temperature models are required also in the GCXE (Muno et al., 2004; Uchiyama et al., 2013).
Figure 1.6a and b show X-ray spectra of the GCXE and GRXE fitted with a two-temperature model
as well as the 6.4 keV line and continuum radiated from cold material. We will hereafter call theplasma with the lower temperature as the low-temperature plasma (LP) and the one with the higher
temperature as the high-temperature plasma (HP).
The distribution of the GRXE and GCXE has often been expressed by an exponential model.
Yamauchi and Koyama (1993) measured the latitudinal distribution of the GRXE (|l| > 10◦ ) and foundthat the e-folding scale height of the 6.7 keV line emission is ∼100 pc. Kaneda et al. (1997) obtained
2 Studies of the Milky Way commonly use the Galactic coordinate system: the Galactic longitude (l) and the Galactic

latitude (b). See also figure 1.1, which shows the Galactic longitude.

CHAPTER 1. REVIEW

6Counts/sec/keV

1

10−1

10 −2

1

10
Energy (keV)

6.7 keV line intensity (a.u.)

Figure 1.3. Spectrum taken from the Galactic plane observed by Tenma (Koyama et al., 1986b).

60

50

40

30

20
10
0
–10 –20
Galactic Longitude (degree)

–30

–40

–50

–60

Figure 1.4. Intensity distribution of the 6.7 keV line along the Galactic plane taken from a Ginga survey(Koyama et al., 1989).

a similar result by observing the Scutum arm (l ∼ 28◦ ). Revnivtsev et al. (2006a) measured the scaleheight of the GRXE of 130 ± 20 pc in the 3–20 keV band at l = 20◦ . Also Revnivtsev and Molkov(2012) performed a scan survey across the Galactic plane with the stripe from b ∼ 0◦ to b ∼ 30◦
at l ∼ 18.◦ 5 and obtained the scale height of ∼ 107 pc in the 4.3–10.5 keV. Uchiyama et al. (2013)
measured intensity profiles of the 6.4 keV, 6.7 keV, and 7.0 keV lines and the continuum band in 5–8
keV, and thus separated the GCXE from the GRXE with scale lengths of 0.◦ 6 (80 pc) and 50◦ (8 kpc)
and with scale heights of ∼ 0.◦ 2–0.◦ 4 (30–60 pc) and ∼ 1◦ –5◦ (150–750 pc), respectively3 .
Yamauchi and Koyama (1993) discovered a clear excess near the GC, which is broader than the
GRXE component. Its scale height is as large as ∼ 5◦ . A similar result was obtained by Valinia andMarshall (1998), who observed the |l| < 30◦ region and found the presence of two components: a thinand thick disk components with the scale heights of ≲ 70 pc (0.◦ 5) and ∼ 500 pc (4◦ ), respectively.
The thick disk component was found also by Revnivtsev (2003). They observed the GRXE near the
GC in the energy band of 3–20 keV and found that the intensity distribution at |b| > 2◦ could be
3 On the assumption of the distance of 8 kpc.

1.3. DEBATES OVER THE ORIGIN OF THE IRON K-SHELL LINES IN THE GDXE(b)Ca XIX Kα

Ar XVII Kα
Ar XVIII Kα

Si XV KαSi XVI Kα

Si XIII Kα
Si XIV Kα

0.01
10

−3

Counts/sec/keV/1chip

0.1

Fe (neutral) Kα
Fe XXV Kα
Fe XXVI Kα

(a)

7

1

2

5
Channel Energy (keV)

10

Figure 1.5. Spectrum of (a) the GCXE (Koyama et al., 1996) and (b) the GRXE (Kaneda et al., 1997) observed
by ASCA.

(a)(b)

Figure 1.6. X-ray spectra of (a) the GCXE and (b) the GRXE observed by Suzaku (Uchiyama et al., 2013). Boththe spectra are fitted with a two-temperature model.

well described by an exponential model with scale height of ∼ 3◦ . This bulge-like emission is called
as the Galactic bulge X-ray emission (GBXE). Hereafter, we will call the GCXE, GBXE, and GRXE
collectively as the Galactic diﬀuse X-ray emission (GDXE).

1.3 Debates over the origin of the iron K-shell lines in the GDXE
The origin of the GDXE, particularly that of the iron lines, has been caused a lot of discussion.
The most popular idea is a superposition of many point sources as is the case in the cosmic X-ray
background (CXB), which is explained by a sum of active galactic nuclei below several hundredskeV (e.g. Ueda et al., 2003). Yamauchi et al. (1996) performed Poisson-fluctuation analysis of thesurface brightness of the GRXE and found that point sources which have the luminosity larger than
∼ 2 × 1033 erg s−1 in the 2–10 keV band are excluded as the origin of the GRXE.
Cataclysmic variables (CVs) and coronally active binaries (ABs) have the luminosity of ≲ 1033 erg s−1

8

CHAPTER 1. REVIEW

in the 2–10 keV band, and their spectra exhibit iron K-shell lines as seen in the GDXE. The number
densities of CVs and ABs are considered to be large in the Milky Way. Therefore, CVs and ABs havebeen major candidates for the GDXE. CVs are close binaries that consist of a white dwarf and a companion star, and show violent variability in the time-scales of seconds to years. The two objects are
so close to each other that the gravity of white dwarf tears oﬀ gas from a companion star and the gas
falls into the white dwarf. X-rays are radiated during the accretion.
ABs are late-type close binaries like RS Canum Venaticorum (RS CVn) systems and BY Draconis(BY Dra-), Algol-, and W Ursae Majoris (W UMa-) type binaries. Most of ABs fall into the first two
types. RS CVn systems typically consist of a G- or K-type giant or subgiant with a late-type mainsequence or subgiant companion, while BY Dra-type binaries typically consist of two late-type mainsequence stars (Dempsey et al., 1993, 1997). Algol binaries contain late-type component, typicallyspectral type K, filling its Rosh robe and losing mass to a A- or B-type companion (Dempsey et al.,1993). W UMa-type binaries are composed of F-K stars which are in contact via a common convective
envelope (McGale et al., 1996).
Sazonov et al. (2006) constructed an X-ray luminosity function of nearby X-ray sources in theluminosity range of 1030 –1034 erg s−1 (2–10 keV) by using RXTE. They also extended this luminosity
function down to 1027.5 erg s−1 by using the Rosat All-Sky Survey in soft X-rays; they converted
the observed soft X-ray (0.5–2 keV) luminosity to the 2–10 keV energy band based on the 0.1–10 keV
spectra of typical sources. Although they ascribed 50% errors in converting, they found that the local
cumulative X-ray (2–10 keV) emissivities (per unit stellar mass) of ABs and CVs are (2.0 ± 0.8) × 1027
and (1.1 ± 0.3) × 1027 erg s−1 M⊙−1 , respectively.
Revnivtsev et al. (2006a) studied the surface brightness of the GBXE (|l| < 10◦ and 1◦ < b < 8◦ )and GRXE (|b| < 0.◦ 5) and found that the distribution of the 3–20 keV band closely follows a Galactic
near-infrared (NIR) surface brightness measured by COBE/DIRBE (3.5 µm). The NIR data trace the
Galactic stellar mass distribution. The authors also calculated a ratio of X-ray luminosity to stellar
mass in the GBXE and obtained the ratio of (3.5 ± 0.5) × 1027 erg s−1 M⊙−1 . This value is almost consistent with the sum of the emissivities of ABs and CVs of ∼ 3.1×1027 erg s−1 M⊙−1 (Sazonov et al., 2006).
Therefore, they claimed that the bulk of the GDXE is composed of point sources. Also Revnivtsev
et al. (2006b) measured the distribution of the 6.7 keV line along the Galactic plane. Here, they did
not resolved the iron line at 6.7 keV into the 6.4 keV, 6.7 keV, and 7.0 keV lines due to performance
limitations of the detector. They found that the line distribution spatially corresponds to the NIR
surface brightness (4.9 µm) as is shown in figure 1.7. They also demonstrated the proportionality
between the 6.7-keV surface brightness and the NIR in the |l| > 10◦ and |l| < 3◦ regions.
XMM-Newton and Chandra have the good spatial resolution of 5′′ and 0.6′′ (FWHM; Aschenbach
et al., 2000; Weisskopf et al., 2002), respectively, and therefore they have an advantage to resolve point
sources. Observing by XMM-Newton, Muno et al. (2004) resolved only 10% of the GCXE into discrete
point sources under their detection limit for point sources of ∼ 3 × 10−15 erg s−1 cm−2 (∼ 1031 erg s−1
in luminosity) in the 2–10 keV band. A similar result was obtained in Chandra observations of the
GRXE performed by Ebisawa et al. (2001, 2005) under the same detection limit as Muno et al. (2004).
Revnivtsev et al. (2009), on the other hand, claimed that they resolved ∼80% of the GDXE inthe 6–7 keV band into discrete point sources by performing deep observation (1 Msec) at (l, b) =(0.◦ 1, −1.◦ 4) by Chandra (Chandra bulge field: CBF). Figure 1.8 shows the spectrum of total emission
observed by Chandra as well as the summed spectra of the detected sources and unresolved emission.

1.3. DEBATES OVER THE ORIGIN OF THE IRON K-SHELL LINES IN THE GDXE

9Figure 1.7. Profiles of the surface brightness of the 6.7 keV line along the Galactic plane at |l| < 2◦ measured
by RXTE/PCA (Revnivtsev et al., 2006b). The dashed line shows the profile of the surface brightness of 4.9 µm
measured by COBE/DIRBE.

The detected sources exhibit a prominent 6.7 keV line although the authors did not separate theiron line into the three lines due to finite energy resolution of the instrument. They also madethe histogram which shows the fraction of resolved X-ray emission around the 6.7 keV line as afunction of the limiting source flux/luminosity (figure 1.9). Most of the detected point sources areprobably ABs (with luminosity of < 1032 erg s−1 ) and CVs (with luminosity of 1031 –1032 erg s−1 ).
They predicted the attributable fraction of the ABs and CVs of 2 : 1. This seems to be a crucial result.
However, there are several problems with the point-source scenario as discussed later (section 2.1).
The other scenario for the origin of the GDXE is a truly diﬀuse emission. Since the GDXE involves
the emission lines from highly ionized iron (the 6.7 keV and 7.0 keV lines), diﬀuse thermal plasma
can contribute to the GDXE. However, there are some problems in the diﬀuse-plasma scenario. The
first one is an energy source. Koyama et al. (1986a) estimated the contribution of uncatalogued
supernova remnants (SNRs) and found that the supernova (SN) rate of one SN per 10 years can
explain the observational result. This rate is much higher than the canonical value of one SN per
50–100 years. The second problem is that hot plasma with temperature as high as 5–10 keV cannot
be bound to the small scale height of ∼ 100 pc by the gravitational potential of the Milky Way.
The discovery of non-thermal emission (Yamasaki et al., 1997; Valinia and Marshall, 1998) raised
another scenario in which the interaction of low-energy cosmic-ray (LECR) electrons with the ISM
contributes to the GDXE. Valinia et al. (2000) argued that the hard X-ray emission above 2 keV comesfrom two components: a thermal component with the temperature of 2–3 keV and a non-thermal
component produced by LECR electron. In this scenario, the 6.7 keV line is due to the former while

CHAPTER 1. REVIEW

10Intensity

the 6.4 keV line originates from the latter. Masai et al. (2002) claimed that the GRXE can be explainedby non-equilibrium emission due to quasi-thermal electrons with the temperature of a few tens ofkeV. These electrons are accelerated in situ, and can be bound gravitationally in the Galactic disk.
Tanuma et al. (1999) proposed that magnetic reconnection plays a role in producing the hot plasma
in the GDXE. Magnetic reconnection, which can occur by SNs, cosmic rays, Parker instability, or
turbulence, heats the interstellar gas by releasing interstellar magnetic energy. Sunyaev et al. (1993)
predicted that a part of the GCXE may be due to Thomson scattering of X-rays from nearby sources
by dense gas. Molaro et al. (2014) applied this idea to the GRXE and suggested that 10–30% of the
GRXE arise from the scattering of X-rays from bright X-ray binaries by the ISM.

10‒4Resolved fraction

1
0.8
0.6
0.4
0.2
0

4

Energy (keV)

68

Figure 1.8. (upper) Spectra obtained by the Chandra deep observation of the region near the Galactic center(the CBF: l = 0.◦ 1, b = −1.◦ 4). Black, blue and red data points show the spectrum of the total emission, the collective spectrum of all detected sources, and the spectrum of the remaining unresolved emission, respectively.
(bottom) Fraction of the X-ray emission resolved by Chandra into point sources as a function of X-ray photon
energy. Both the figures are taken from Revnivtsev et al. (2009).

1.4. 6.4 KEV CLUMPS IN THE GALACTIC CENTER REGION

11

Flux 2‒10 keV (erg s ‒1 cm ‒2)
10‒16

10‒17

10‒15

10‒14

10‒13

1032

1033

100

80

60

ABs

40

CVs

20

0

extragalactic sources
1029

1030

1031

2‒10 keV source luminosity at 8 kpc (erg s ‒1)

Figure 1.9. Fraction of resolved X-ray emission around the 6.7 keV iron emission line as a function of thelimiting source flux/luminosity (Revnivtsev et al., 2009). The histogram shows the fraction of the total fluxin the 6.5–7.1 keV energy band in the CBF provided by discrete sources with fluxes above a given detection
threshold in the 2–10 keV energy band. The wide yellow curve shows the corresponding dependence expected
for a combination of Galactic sources with luminosities below 1032 erg s−1 (all located at the Galactic center
distance of 8 kpc) and extragalactic sources with fluxes below 2 × 10−14 erg s−1 cm−2 in this energy band.
The blue, green, and red curves show the expected contributions of extragalactic sources, CVs, and ABs,
respectively.

1.4 6.4 keV clumps in the Galactic center region
Figure 1.10 shows the longitude distribution of the 6.4 keV and 6.7 keV lines in the GC obtainedby Suzaku. Both the profiles are clearly diﬀerent from each other; the profile of the 6.7 keV line
distributes smoothly while that of the 6.4 keV line are clumpy. Therefore, the 6.4 keV clumps diﬀers
from the 6.7 keV line in the origin. Koyama et al. (1996) discovered that these clumpy structures
correspond to giant molecular clouds in the Sgr B2 and the Radio Arc regions. Figure 1.11 shows the
X-ray image in the 6.4 keV band and velosity-integrated CS (J = 1 − 0) map in the GC. CS emission
traces highly dense gas, and we see the regions which is bright in 6.4 keV corresponding to thosewith dense molecular clouds. The 6.4 keV clumps should be originate from large amount of neutral
iron atoms in the clouds.
Koyama et al. (1996) proposed that the 6.4 keV emission from the Sgr B2 cloud is due to photoionization of neutral iron by external X-rays. This idea is called an X-ray reflection nebulae (XRN)
scenario, which is essentially the same as the prediction by Sunyaev et al. (1993). Murakami et al.
(2000) found the large equivalent width (1–2 keV) of the 6.4 keV line and the deep absorption edge at
the iron K-edge (7.1 keV) corresponding to NH ∼ 1024 cm−2 . They also discovered that the peak position of the 6.4 keV line emission is shifted from the core of the molecular cloud toward the Galactic

CHAPTER 1. REVIEW

12

2

10−5

1
0
Galactic Longitude (degree)

−1

−2

10−6
10−7
10

−9

10

(b) 6.7 keV line

10−8

10−8

10−7

10−6

Line Intensity (ph cm−2 s−1 arcmin−2 )(a) 6.4 keV line

−9Line Intensity (ph cm−2 s−1 arcmin−2 )10−5

center by about 1′ –2′ . These features are well explained by the XRN model. Murakami et al. (2001a)discovered another 6.4 keV clump in the Sgr C region. In order to explain the 6.4 keV line intensity,a bright source with the luminosity of ≥ 1038 –1039 is required (Koyama et al., 1996; Murakami et al.,2001a,b). A candidate source is Sgr A*, the super-massive black hole at a distance of several hundred
light-years from the molecular clouds (Ghez et al., 2008). Although the present luminosity of Sgr A*
is only 1033 –1034 erg s−1 (Baganoﬀ et al., 2003), it is possible that Sgr A* was bright several hundredyears ago with an X-ray luminosity a million times higher than the present.

2

10
Galactic Longitude (degree)−1−2

Figure 1.10. Longitude distributions of (a) the 6.4 keV line and (b) the 6.7 keV line in the Galactic centerregion. The data points are taken from Uchiyama et al. (2011).

Galactic latitude (degree)

0.2
0
–0.2
0.2
0
–0.2
1

0.5

0
Galactic longitude (degree)

–0.5

–1

Figure 1.11. (top) X-ray image in the 6.4 keV line in the Galactic center region obtained by Suzaku (Nobukawa,2011). (bottom) Velosity-integrated CS (J = 1–0) emission in the same region as the top panel. The velocity
integrated range is from −200 km s−1 to 200 km s−1 (Tsuboi et al., 1999).

Inui et al. (2009) discovered time variability of the 6.4 keV flux in the Sgr B2 region from thearchival data for 10 years obtained with Suzaku, XMM-Newton, Chandra, and ASCA; the 6.4 keV

1.4. 6.4 KEV CLUMPS IN THE GALACTIC CENTER REGION

13

line intensity decreased to 60% in the time span 1994–2005 (figure 1.12) in contrast to the constantflux of the 6.7 keV line. Also other authors detected similar time variability of the 6.4 keV line flux
in other regions (Ponti et al., 2010; Nobukawa et al., 2011). Although these results favored the XRN
scenario, another scenario in which the 6.4 keV line is produced by inner-shell ionization by LECR
electrons could not be rejected.

Figure 1.12. Surface brightness maps of the iron lines obtained with (a) Suzaku/XIS in 2005, (b) XMMNewton/MOS and PN in 2004, (c) Chandra/ACIS-wein 2000, and (d) ASCA/SIS in 1994. Pixel size is 50′′ ×50′′
in each case (Inui et al., 2009).

The equivalent widths consolidated the XRN scenario. Nobukawa et al. (2010) discovered K-shelllines of neutral argon, calcium, chrome, and manganese atoms in the GC region. The equivalent
widths of these K-shell lines indicate that the metal abundances should be 1.6 solar in the XRN
scenario and 4 solar in the electron-bombardment scenario. However, the metal abundance in the
GC was found to be 1–2 solar (Muno et al., 2004; Nobukawa et al., 2011; Uchiyama et al., 2013).
Therefore, the origin of the 6.4 keV clumps are believed to the XRN.
In addition to the 6.4 keV clumps, a more uniformly distributed 6.4 keV line emission extending
from the GC to the Galactic plane is found (see figure 1.10; Ebisawa et al., 2001; Yamauchi et al., 2009;
Uchiyama et al., 2011). This component should have an origin other than the XRNe reflecting X-rays
from Sgr A∗ . Its origin is under debate as a part of the GDXE. It is uncertain even whether the 6.4 keVline emission shares the origin with the other iron lines and the continuum. At present, a common
candidate for the origin of the 6.4 keV line emission is CVs because they exhibit a prominent 6.4 keVline (Yuasa et al., 2012). However, the previous studies supporting point-source scenario have several
problems to be considered and addressed as discussed in the next chapter.

14

CHAPTER 1. REVIEW

Figure 1.13. Equivalent widths of K-shell line of various neutral atoms in the GC. Black and red lines are thecalculated values for the X-ray (with photon index of 1.9) and electron (with particle index of 3) scenarios,
respectively. The data points marked with the open circles are the observed values. Errors were estimatedat the 90% confidence level. The black and red dashed lines are to guide eyes, which are X-ray and electron
scenarios in 1.6 solar and 4.0 solar abundances, respectively (Nobukawa et al., 2010).

Chapter 2

Introduction
2.1 Problems in the previous study on the iron lines in the GDXE
The true origin of the GDXE would be a mixture of point sources and diﬀuse components. The
most popular scenario is that a superposition of point sources dominates the GDXE. In fact, it seems
to be crucial that Revnivtsev et al. (2009) resolved 80% of the iron line flux into discrete point sourcesin the CBF. However, the previous studies on the GDXE, especially those supporting the point-source
scenario, leave several severe problems. We will consider them in the following.
1. Is the result of 80% resolution in the CBF obtained by Revnivtsev et al. (2009) universal in
the GDXE?
Previous study revealed that the GDXE is composed of the three components: the GCXE,
GBXE, and GRXE. Whether they share the origin is unknown. However, there are some hints
which suggests that the three components are diﬀerent in character.
Uchiyama et al. (2013) investigated the distribution of the 6.7 keV and 7.0 keV line intensity
along the Galactic plane as shown in figure 2.1. While the intensity of the 6.7 keV (helium-likeiron) line in the GCXE is 10 times as large as the GRXE, that of the 7.0 keV (hydrogen-like iron)line in the GCXE is 30 times as large as the GRXE. This result suggests that the X-ray spectrumof the GCXE is diﬀerent from that of the GRXE. Even in the point-source scenario, the spectra
and the mixing ratio of CVs and ABs should be diﬀerent between the GCXE and GRXE.
Also a stellar mass distribution model suggests a possibility that the mixing ratio of point
sources in the CBF is diﬀerent from the GC region and the Galactic plane. Colored lines infigure 2.2 show the stellar mass distribution model originally compiled by Muno et al. (2006),
which is composed of the nuclear stellar cluster, nuclear stellar disk, Galactic disk, and Galactic
bulge. Although the CBF is located near the GC, its latitudinal coordinate of b = 1.◦ 4 is a bit
farther out from the Galactic plane. The arrow in figure 2.2b shows the position of the CBF,
where the Galactic bulge and Galactic disk components contribute half and half. On the other
hand, the central (|l∗ | < 1◦ and |b∗ | < 0.◦ 2 ) region is dominated by the nuclear stellar cluster
and nuclear stellar disk components, and the Galactic plane (|l| > 1◦ ) is contributed mainly by
the Galactic disk component. Furthermore, figure 2.2 also shows a large excess over the stellar
mass distribution model in the GC region (|l∗ | < 1.◦ 5 and |b∗ | < 1◦ ) assuming that all the other15

16

CHAPTER 2. INTRODUCTIONGDXE are due to unresolved point sources. This indicates the possibility that other components
than point sources are dominated the GCXE.
There is still a room for consideration about the argument that Revnivtsev et al. (2009)
resolved 80% of the iron line flux in the CBF by Chandra and hence point sources should
account for the iron lines of GDXE in other regions. We should verify whether their argument
is valid throughout the GDXE: the GCXE, GBXE, and GRXE.
2. The GDXE origin has been debated based on the observational results with limited spatialand spectral resolution.
The GDXE has been described often by an exponential model with a scale height and ascale length (e.g. Uchiyama et al., 2013). Especially the scale height would provide important
information to solve the origin of the GDXE.
In the previous study, a stellar mass distribution model constructed from NIR surface brightness observed by COBE/DIRBE has been used and compared with the spatial distribution of
the GDXE (Revnivtsev et al., 2006a; Muno et al., 2006; Uchiyama et al., 2011). However, theangular resolution of the NIR maps was 0.◦ 7, and thus it is not simple way to compare them to
the GCXE and GRXE maps with the scale height of ∼ 0.◦ 5–0.◦ 7. In fact, as shown in figure 2.2b,
the scale height of the GCXE would be smaller than that of point sources (Uchiyama et al.,2011). Furthermore, the NIR maps could be aﬀected by bright stars, and hence it is uncertainthat they can thoroughly follow low mass stars such as ABs and CVs.
As described in section 1.2, the iron three lines, namely 6.4 keV, 6.7 keV, and 7.0 keV lines,
have diﬀerent radiation mechanism. Therefore, the origin of each line could be diﬀerent fromeach other. However, the previous studies have often not separated the three lines due to limitations of detector performance such as energy resolution and eﬀective area (e.g. Revnivtsev
et al., 2006b, 2009). The previous debate over the origin of the GCXE, especially the origin of
the bright 6.4 keV line emission, shows indispensability to separate each line as described in
section 1.4.
Also equivalent widths1 of the three iron lines in the GDXE are controversial issue. Uchiyama
et al. (2013) found the equivalent widths of the 6.4 keV, 6.7 keV, and 7.0 keV lines in the
GRXE are ∼ 110 eV, ∼ 490 eV, ∼ 110 eV, respectively (see also Ebisawa et al., 2008; Yamauchi
et al., 2009). Averaged equivalent widths of CVs with the luminosity of 1030 –1033 erg s−1 are∼ 130 eV for the 6.4 keV line, ∼ 160 eV for the 6.7 keV line, and ∼ 90 eV for the 7.0 keV line
(Hellier et al., 1998; Ezuka and Ishida, 1999; Hellier and Mukai, 2004). Suzaku archive of ABs
shows that their mean equivalent widths of the 6.4 keV, 6.7 keV, and 7.0 keV lines are ∼ 60 eV,∼ 560 eV, and ∼ 80 eV, respectively (Uchiyama et al., 2013) .
As noted by Uchiyama et al. (2013) and Warwick et al. (2014), any mixture of well-known
X-ray CVs and ABs cannot simultaneously explain the equivalent widths of the three lines in
the GRXE, and faint sources with lager equivalent width than those of the well-known CVs andABs are required in order to explain by the point-source scenario. Yuasa et al. (2012) found thatthe spectrum of the GDXE near the GC is explained by CVs with a specific mass. However, in
1 An equivalent width is the area of the line on a plot of counts versus X-ray energy. Unit is energy (eV or keV).

2.2. AIM OF THIS DISSERTATION

17

order to explain the observed equivalent width, the iron abundance is required ∼ 3 times larger
than the typical value in the well-known CVs.

(a) 6.7 keV(b) 7.0 keV
10 −5

10 −5
−2

photons cm s arcmin

GRXE

GCXE

10 −6
10 −7

GRXE

−2

−1

10 −7

−2

−1

photons cm s arcmin

−2

GCXE
10 −6

10 −8
10 −9
0.01

0.1

1

10

100

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

10 −8
10 −9
0.01

0.1

1

10

100

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

Figure 2.1. (a) 6.7 keV-line intensity distribution along the Galactic plane as a function of l∗ (distance from SgrA∗ ). (b) 7.0 keV-line intensity along the Galactic longitude. Both the figures are taken from Uchiyama et al.
(2013).

Galactic disk
Galactic bulge
nuclear stellar
cluster

nuclear stellar
disk

Galactic disk
Galactic bulge
nuclear stellar cluster
nuclear stellar disk

Figure 2.2. (a) 6.7 keV line intensity distribution along the Galactic plane as a function of l∗ (distance from SgrA∗ ). The black solid line is the sum of the four components. (b) Same as the left panel (a), but those along the
Galactic longitude at l∗ = −0.◦ 114 as a function of b∗ (distance from the Galactic plane). The arrow shows the
position of b = −1.◦ 4, which is almost the same region as the CBF (Revnivtsev et al., 2009). Both the figures aretaken from Uchiyama et al. (2011).

2.2 Aim of this dissertation
The aim of this dissertation is to reveal the origin of the 6.4 keV line emission along the Galactic
plane by addressing the problems laid above: (1) We confirm that the GDXE consists of the GCXE,
GBXE, GRXE. (2) We separately measure their distributions (the scale heights) and spectral features(the equivalent widths) for the 6.4 keV line, and for comparison, those of the 6.7 keV and 7.0 keVlines in the GCXE, GBXE, and GRXE. (3) We investigate the origin of the 6.4 keV line emission basedon the observational result.

18

CHAPTER 2. INTRODUCTION

In this dissertation, we use the Suzaku satellite. The X-ray charge coupled device camera, the
XIS, onboard Suzaku provides lower background and large eﬀective area than ever before, especiallyfor a diﬀuse emission with energy above 5 keV, and thus Suzaku have demonstrated its capability
to resolve the 6.4 keV, 6.7 keV, 7.0 keV lines in the GDXE. We analyze archival data from a huge
number of Suzaku observations along the inner Galactic disk (|l| ≲ 80◦ ).
In chapter 3, we overview the Suzaku satellite (section 3.1), the X-ray telescope (section 3.2), andthe CCD camera (section 3.3), which are the instruments we used for the observations. The detailed
observation log and data reduction are described in chapter 4. In chapter 5, we make the global
intensity profiles of the 6.4 keV, 6.7 keV, and 7.0 keV lines and the hard continuum and measure theirscale heights to spatially resolve the GCXE, GBXE, and GRXE. We also separately obtain equivalent
widths of the three lines for the GCXE, GBXE, and GRXE. In order to constrain the origin of the 6.4
keV line, we extracted X-ray spectra from the regions where the 6.4 keV line intensity is especiallystrong (chapter 6). In the involved regions, some XIS fields of view (FOVs) are contaminated by "straylight", which occurs by anomaly reflection of X-ray which comes from a very bright point source
outside an FOV. We evaluate its eﬀect before spectral analysis (section 6.1). We draw a conclusion forthe origin of the 6.4 keV line emission from the Galactic plane in chapter 7. At the last, we briefly
give prospects for ASTRO-H, the next Japanese X-ray astronomy satellite (chapter 8).
Throughout this dissertation, the distance to the GC is assumed to be 8 kpc. We call the regionsin positive and negative Galactic longitude as "east" and "west", respectively. Errors quoted in thispaper are at 68% confidence levels unless otherwise specified.

Chapter 3

Instrument:
Suzaku X-ray Satellite
3.1 Suzaku satellite
Suzaku is the fifth Japanese X-ray astronomy satellite launched by Japan Aerospace Exploration
Agency (JAXA) with the M-V Launch Vehicle from JAXA’s Uchinoura Space Center on July 10, 2005.
A schematic view of Suzaku is shown in figure 3.1(left). This satellite was orbiting in a near-circularpath with an altitude of 550 km, an inclination of 31.9 degrees, and an orbital period of about 96
minutes. Suzaku had been operated for about 10 years exceeding its planned observation lifetimeof about two years. However, its communication with the ground had been established only in fragments since June 1, 2015. The completion of its scientific mission was announced from JAXA on
August 26, 2015.
Suzaku carried the X-ray spectrometer (XRS, Kelley et al., 2007), the X-ray Imaging Spectrometer(XIS, Koyama et al., 2007), and the Hard X-ray Detector (HXD, Kokubun et al., 2007; Takahashi et al.,2007). The XRS was intended to provide for the first time both a good energy resolution (∼ 7 eV
at 6 keV; FWHM) and a large eﬀective area. Unfortunately, because of a liquid helium coolant loss
on August 8, 2005, the XRS was no longer able to provide the planned science. We will not discuss
the XRS further in this dissertation. The XIS consists of four X-ray charge-coupled device (CCD)
cameras, each placed on the focal plane of each module of the X-Ray Telescope (XRT, Serlemitsos
et al., 2007). A side view of the instrument and telescopes are shown in figure 3.1(left).
In this dissertation, we used the data of the XIS in the energy range of 0.2–12 keV, and hence wewill describe the detail of the XRT and the XIS in the following sections.

3.2 X-ray Telescope (XRT)
3.2.1 Mechanism and design
X-ray optics is completely diﬀerent from that for visible light. While lenses used for visible lightare produced with transparent materials that have a refractive index much larger than unity, refractive indices of matter to X-rays are less than unity. Thus the most common method to focus X-rays isreflection. X-ray reflection occurs, only when the grazing incidence angle is shallow enough (≤ 1◦ ).
19

CHAPTER 3. INSTRUMENT

20

XRT-S
Sun shade
XRT-I(4 units)

Star tracker

Side panel(8 panels in total)

Gyro-scopes

XRS
HXD
Base panel

XIS(4 units)

Figure 3.1. (left) Schematic picture of the Suzaku satellite. (right) Side view of the instruments and telescopeson Suzaku. Both are taken from Mitsuda et al. (2007)

The XRTs contain in total ∼ 170 pairs of concentric thin-foil reflectors, whose thickness is only
∼ 165 µm (figure 3.2). The reflectors are thin aluminum foils coated with gold as an X-ray-reflecting
material. The gold layer was sputtered onto the glass tube with a highly smooth surface, and then
was transferred to the foil reflectors (Serlemitsos and Soong, 1996). The incident grazing angle is≲ 0.6◦ so that even the outermost mirror shell can reflect the iron K-shell line. The diameter of the
XRT is ∼ 400 mm. The reflectors are radially nested as tight as possible, which results in an apertureeﬃciency as high as ∼ 60%. Thus the XRTs achieve at the same time a large eﬀective area and lightweight. In fact, although the total weight is only ≲ 80 kg, the total eﬀective area of the four sets of
XRTs at 7 keV is comparable to that of XMM-Newton mirrors, whose total weight is ∼ 1300 kg.
The XRT adopted an approximation of the Wolter-I type geometry. As shown in figure 3.3, X-rays
are reflected twice on the surface of the primary and secondary reflectors and are gathered on focal
plane. The shapes of the primary and secondary reflectors are a paraboloid and hyperboloid, respectively, in the original Wolter-I type optics. In Suzaku, the reflectors have a double-cone geometry
to approximate the Wolter-I system (Serlemitsos et al., 2007). This approximation ultimately limits
the achievable angular resolution. Reflectors are fabricated in quadrant cones, and the quadrants arealigned so that the optical axes are within 2′ from the mechanical axis. The focal length is 4.75 m for
the XIS.
X-rays from the FOV are correctly focused when they are reflected by the two surfaces. However,
the tight nesting of the reflectors raises the possibility of reflection other than the normal reflection
path; X-rays from the outside of the FOV are reflected only once by the surface or by the backsideof the reflectors. When such unintended X-ray reflections cause "stray light" in the FOV. We will

3.2. X-RAY TELESCOPE (XRT)

21

Figure 3.2. Picture of an XRT module (Serlemitsos et al., 2007).

(a)

(b)

paraboloid“primary”

hyperboloid“secondary”
focus

Figure 3.3. Schematic view of an X-ray telescope composed of nested reflectors: (a) front view; (b) the crosssection including the optical axis (Mori et al., 2005).

later describe characteristics of stray lights (section 3.2.3). In order to reduce the stray light, a precollimator is equipped on the top of the XRT.
The top of the pre-collimator stage is covered with a thermal shield in order to isolate the XRTthermally from the space and to keep the temperature within the XRT’s designed function rangeof 20.0 ± 7.5◦ C. The thermal shield also function to block optical light from the sky and the Earth’s
surface illuminated by the Sun.

3.2.2 Performance
The XRTs are designed to achieve the maximum possible eﬀective area, particularly at the energyof the iron K-shell lines even under a severe weight limit imposed by the launch vehicle for Suzaku.
Thanks to the tight nesting of the reflectors, the XRTs have large eﬀective areas (450 cm2 at 1.5 keV,
250 cm2 at 7 keV per module). On the other hand, the XRTs trade oﬀ the spacial resolution. In
order to evaluate the spacial resolution, a half power diameter (HPD) is commonly used. The HPD
is defined as the diameter within which 50% of X-rays from the focused point source are enclosed.
The HPDs of the XRTs are 1.′ 8–2.′ 3, which are about one order of magnitude larger than the X-ray

22

CHAPTER 3. INSTRUMENT

telescopes onboard the other contemporary satellites such as Chandra and XMM-Newton.
An intensity profile of a point source as a function of radius from the point source is called as apoint spread function (PSF). An enclosed energy function (EEF) is an integration of a PSF and shows
the ratio of the flux in the enclosed circle to the total flux of a point source. The PSF and EEF of aunit of the XRTs are shown in figure 3.4.

Figure 3.4. Point-spread function (PSF; left), and encircled energy functions (EEF; right) of the XRT module
for XIS 0. The plots were obtained by an observation of a moderately bright point source, SS Cyg. The EEF isnormalized to unity at the edge of the CCD chip (Serlemitsos et al., 2007).

Figure 3.5. Vignetting of the four XRT-I modules using the data of the Crab Nebula observed during August
22–27, 2005 in the two energy bands of 3–6 keV and 8–10 keV. The model curves were calculated with a raytracing simulator with spectral parameters of NH = 0.33×1022 cm−2 , a photon index of 2.09, and normalizationof 9.845 photons cm−2 s−1 keV−1 at 1 keV (Serlemitsos et al., 2007).

The eﬀective area for oﬀ-axis sources is smaller than that for on-axis ones. This eﬀect is calledas "vignetting", and is prominent in high energy. A critical angle is smaller for higher energy X-rays
than for lower ones, and therefore less of oﬀ-axis high-energy photons reach the focal plane (seefigure 3.3). Figure 3.5 shows the oﬀ-axis angle dependence of the eﬀective area for the soft (3–6 keV)and the hard (8–10 keV) bands.

3.2. X-RAY TELESCOPE (XRT)

3.2.3

23

Stray light

There are two major components in stray light appearing in the FOV. Reflection paths causing
stray light are shown in figure 3.6 as well as the normal double reflection. When an incident X-ray
is reflected only once by a secondary reflector, which is so-called secondary reflection, the brightest
stray light appears (figure 3.6b). The other component is referred to as backside reflection, whichis the result of reflection on the backside of a primary, followed by the normal double reflection(figure 3.6c). As represented in figure 3.6, if the oblique angle of a primary reflector measured fromthe optical axis is defined as τ, the secondary reflection can occur when the oﬀ-axis angle of anincident X-ray θ is in the range of τ < θ < 2τ, and the backside reflection appears in the case of2τ < θ < 3τ. The oblique angles of all the primary reflectors are in the range of 0.◦ 18 < τ < 0.◦ 60.
Therefore, the secondary and backside reflections can possibly cause stray images when the oﬀ-axis
angle of the source is in the range of 0.◦ 18 < θ < 1.◦ 26 and 0.◦ 36 < θ < 1.◦ 89, respectively, which meansthat the two major components appear on the opposite sides with respect to the detector center (Moriet al., 2005).

Figure 3.6. Reflection paths causing stray light as well as the normal reflection. The oblique angle of the
primary reflector measured from the optical axis of the XRT is represented by τ. (a) Normal double reflectionof the X-ray. (b) Secondary reflection, which can occur when the incident angle of the X-ray measured fromthe optical axis θ is in the range of τ < θ < 2τ. (c) Backside reflection, which is followed by the normal double
reflection. This pattern occurs in the case of 2τ < θ < 3τ (Mori et al., 2005).

Ray-tracing simulation demonstrates a capability of the pre-collimator. Figure 3.7 shows a widerange vignetting function of the XRT at the three energies of 1.49 keV, 4.51 keV, and 8.04 keV foroﬀ-axis angles up to 80′ . The pre-collimator cannot completely reduce the secondary reflection at
the oﬀ-axis angle of < 30′ , and also a leak of secondary reflection occurs at oﬀ-axis angle of ∼ 60′ –70′
through the gap between the pre-collimator and the primary reflectors. However, the intensity ofthe secondary reflection is reduced by a few orders of magnitude on the average in the oﬀ-axis angle
range of 20′ –70′ .

24

CHAPTER 3. INSTRUMENT

Figure 3.7. Wide-range vignetting function of the XRT at the three energies of 1.49 keV, 4.51 keV, and 8.04 keV
for oﬀ-axis angles up to 80′ . All of the curves were simulated by a ray-tracing simulator (Mori et al., 2005).

Studies of the stray light were carried out with Crab nebula observations with oﬀ-axis angles of
20′ , 50′ , 65′ , and 120′ . Examples of mosaic images taken in the observations are shown in figures 3.8.
The stray images diﬀer from oﬀset to oﬀset and from azimuth to azimuth. They also strongly dependon the incident X-ray energy. The Crab observations revealed that there is discrepancy of the straylight intensity between the ray tracing simulation and observed data. The spectral shapes above
1.5 keV are well reproduced by the simulation. However, the absolute fluxes diﬀer from the raytracing simulation with a factor of ∼ 2. This is because the estimated gap between the pre-collimatorand the primary reflectors in the simulation does not reflect the real one. Therefore stray light flux
cannot be estimated by the simulation (Takata et al., 2015).
Figure 3.9 shows a schematic diagram of the stray patterns for a stray light source at an azimuth
angle of α and an oﬀset angle of R. The square represents the FOV of the XIS. The FOV near side toa stray light source is contaminated by stray light of the secondary-reflection component, but at thesame time free from the backside-reflection component. On the other hand, the far side of the FOV
suﬀers from the backside component, but free from the secondary component. There is another straylight free region corresponding to the quadrant boundary with a cone angle of 12.◦ 8 at an azimuth
angle of 45◦ with a 90◦ pitch (the white radial lines in figure 3.9).

3.2. X-RAY TELESCOPE (XRT)

25

(a)

(b)

Figure 3.8. Mosaic images of oﬀset observations of an extremely bright point source, Crab, in the sky coordinate (right ascension and declination). (a) Observation with oﬀ-axis of 20′ and 50′ . (b) Same as (a), but withaﬀ-axis of 50′ and 65′ (Suzaku Technical Description).

26

CHAPTER 3. INSTRUMENT

Figure 3.9. Schematic diagram of stray light patterns for a stray source at an azimuth angle of α and oﬀset
angle of R. The square shows the FOV of the XIS. This figure is taken from Suzaku Technical Description.

3.3. X-RAY IMAGING SPECTROMETER (XIS)

3.3

X-ray Imaging Spectrometer (XIS)

3.3.1

X-ray CCDs

27

The XIS consists of MOS-type CCDs, in which each pixel is composed of metal (poly-silicon),oxide (silicon dioxide) as an insulator and p-type semiconductor (silicon). When an X-ray photon
enters into the depletion layer in a CCD, photo-electrons are created and then electron-hole pairs are
generated. The number of the generated electrons is proportional to the incident photon energy. The
average energy required to produce an electron-hole pair is 3.65 eV for silicon. The electrons andholes are separated by the electronic field in the depletion layer before they recombine. The electrons
are collected into the potential wells under the electrodes as signal charges. After an exposure period,
the voltages of the electrodes are clocked to transfer the signal charges to the readout nodes.
The XIS adopted frame transfer CCDs, which consist of two distinct area: an imaging area anda frame-store region as shown in figure 3.10. Whereas the imaging area is exposed, the frame-store
region is shielded from X-rays. Each CCD chip is composed of four segments called segment A, B, C,
and D. Each segment has its own readout node. After a certain exposure period, the integrated signal
charges are transferred in a short time from the imaging area to the frame-store region. The charges
in the frame-store region are transferred vertically to the horizontal shift register by each row. At theend, the signal charges in the horizontal shift register are transferred to the readout node by eachpixel. This sequence is repeated until all the pixels in the frame-store region are read out. When the
charges are transferred to the frame-store region, new exposure starts in the imaging region. The
exposures and readouts are performed continuously. The coordinates describing the actual pixel
location on the chip are named ACTX and ACTY coordinates (see figure 3.10).
Figure 3.11 shows a photograph and cross section of the XIS sensor. Each XIS sensor has a single
CCD chip. Figure 3.12 shows a photograph of the CCD installed in the base. The imaging areaof the CCD has 1024 × 1024 pixels, and the pixel size is 24µm×24µm. Combined with the XRT,
the imaging area covers the 18′ × 18′ region in the sky. Each XIS sensor has three 55 Fe calibration
sources. One is for the ground calibration and initial orbit calibration, and is attached to the door.
For normal observations, this calibration source is out of the FOV because the door was opened. The
other calibration sources illuminate the two far-end corners from the read-out node of the imaging
area.
The XIS employs two types of CCDs: front-illuminated (FI) and back-illuminated (BI). Three (XIS
0, 2, and 3) of the four sensors are FI type whereas the other (XIS 1) is BI type. The front side of the
CCDs have a structure of electrode and insulator made of poly-silicon and silicon dioxide layers,
whose thicknesses are ∼ 0.28 µm and ∼ 0.44 µm, respectively. The FI CCDs are less sensitive than the
BI CCD to X-rays in the < 2 keV band because the front-side structure absorbs soft incident X-rays.
On the other hand, the depletion layer of the FI CCDs is thicker (∼ 65 µm) than that of the BI CCD(∼ 42 µm), and hence the FI CCDs have high quantum eﬃciency to hard X-rays than the BI CCD(figure 3.13). The whole region of XIS 2 and the segment-A of XIS 0 have been out of function since2006 November and 2009 June, respectively, possibly due to micro-meteorite impact on the CCDs.

CHAPTER 3. INSTRUMENT

28

Figure 3.10. Schematic view of the XIS CCD (Koyama et al., 2007).

Hood
105 mmBaffleElectric valve

Door calibration
sourceBonnet
125 mmDoor

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

Base
70 mm

Wall calibration
sourceFlexible print

Figure 3.11. Picture (left) and cross section (right) of the XIS sensor (Koyama et al., 2007).

3.3. X-RAY IMAGING SPECTROMETER (XIS)

29

Figure 3.12. Picture of the XIS CCD installed in the base (Koyama et al., 2007).

FI
BI
Effective Area (cm2)

100

10

1
0.1

1
Energy (keV)

10

Figure 3.13. Eﬀective area of one XRT+XIS system. The FI CCDs (XIS 0, 2, and 3) have larger eﬀective area
than the BI CCD (XIS 1) in the high energy band of ≳ 5 keV.

3.3.2 Energy response
The pulse height distribution of the signals for a monochromatic incident X-ray energy is calleda response function of the detector. The schematic distribution of the XIS response is shown infigure 3.14. The XIS response consists of at least six components: (1) a main peak, (2) a sub peak, (3) atriangle component, (4) a silicon escape, (5) a silicon line, and (6) a constant component (Matsumoto
et al., 2006). The main-peak component is made in the case where X-ray photons are absorbed in
the depletion layer of the CCD and all electrons are collected as signal. The sub-peak, triangle andconstant components come from events in which a part of charges are lost. The silicon escape line
appears at the pulse height channel lower than the main peak by 1.74 keV (which is the energy of

30

CHAPTER 3. INSTRUMENT

the silicon Kα line). The silicon escape event happens when a fluorescent silicon X-ray photon is
generated by a photoelectric absorption and the photon escapes from the pixel that absorbed theincident X-ray photon. The silicon Kα line at 1.74 keV occurs when the fluorescent silicon photon isabsorbed in another pixel.

Figure 3.14. Schematic description of the XIS response, which consists of six components: (1) main peak, (2)sub peak, (3) triangle component, (4) silicon escape, (5) silicon Kα line, and (6) constant component (Matsumoto et al., 2006).

3.3.3 Performance and calibration
CCDs in orbit suﬀer radiation damage in a space environment; for example, cosmic rays generatecharge traps that capture signal carriers and increase the charge transfer ineﬃciency (CTI). The resulting pulse heights of the charge generated in a pixel with a given energy depend on the distance
between the pixel and the readout node. Since the charge loss is a stochastic process, the amount oflost charge fluctuates, and energy resolution degrades.
In order to reduce the CTI, a spaced-row charge injection (SCI) technique has been introduced
into the XIS (Prigozhin et al., 2000; Bautz et al., 2004; LaMarr et al., 2004). The principle of the SCI
technique is to inject artificial charges into given pixels1 from serial registers attached to the top ofeach column (Bautz et al., 2004; Nakajima et al., 2008; Uchiyama et al., 2009; Ozawa et al., 2009). The
injected charges fill traps so that subsequent charges are transferred with less charge loss.
Figure 3.15 shows the long-term trend of the measured peak energy and width of the manganeseKα line (5.9 keV) from the calibration sources. The peak energy decreases whereas the width increases gradually. In the middle of 2006, the SCI technique was put into routine operation, which
resulted in a drastic improvement of the gain and energy resolution. The amount of charges wereequivalent to that generated by a 6 keV X-ray photon (6 keV-equivalent) for the FI devices whereas
the smaller amount of 2 keV-equivalent was injected for the BI device in oder to decrease the expected SCI-related increase in noise in the soft X-ray band, where the BI device has an advantage
over the FI devices. The long-term trend in XIS1 had shown increment of the CTI at a faster rate,
1 In the case of the XIS, the artificial charges are injected periodically in space: every 54 rows.

3.3. X-RAY IMAGING SPECTROMETER (XIS)31

and therefore, the injected charge amount was changed from 2 keV to 6 keV equivalent also for the
BI sensor in 2011.

Figure 3.15. Gain history of manganese Kα line for uncalibrated data of the XIS (Suzaku Technical Description).

Mn K Line Center (keV)

Mn K Line Center (keV)

5.92

XIS 0 (FI)

5920
5900
5880
5920

XIS 1 (BI)

5900
5880

XIS 0

5.9
5.88
5.92

XIS 1

5.9
5.88
5.92
5.9

XIS 3

5.88
0

200
400
Time since 2005-7-10 (day)

600

2008

2010
Time (year)

2012

2014

Figure 3.16. Gain history of manganese Kα line for calibrated data of the XIS before turning on the charge
injection technique (left) and after (right). The black and red data in the right panel show the segment A and
D of each sensor, respectively (Suzaku Technical Description).

Although the SCI technique extremely improves the gain and energy resolution, the degree of
CTI depends on the distance from the preceding charge-injected row, and hence the SCI technique
can cause positional variations in gain on the CCD chip. This results in degradation of the energy
resolution and the determination accuracy of the incident energy. Therefore, all the XIS data in the
SCI mode are calibrated with a new method to correct the non-uniformity of the pulse heights. The
details of the CTI correction are described in Uchiyama et al. (2009).
Figure 3.16 shows the time history of the center energy of the manganese Kα line in the calibrated data. The center energy has been constantly 5.895 keV. The uncertainty of the energy-scale
determination is less than ∼ 20 eV (0.3%). The energy resolution of the XIS is defined as the widthCHAPTER 3. INSTRUMENT

32

250

XIS 0 (FI)200

MnKa Line Width (eV)

MnKa Line Width (eV)

250

150
100
250

XIS 1 (BI)

200
150

XIS 0

200
150
250

XIS 1

200
150
250

XIS 3

200

100
0

200
400
Time since 2005-7-10 (day)

600

150
2008

2010
Time (year)

2012

2014

Figure 3.17. Energy-resolution history of manganese Kα line for calibrated data of the XIS before turning on
the charge injection technique (left) and after (right; Suzaku Technical Description). The black and red datain the left panel show the segment A and D of each sensor, respectively.

of the main peak in the response function. The time history of the line widths of manganese Kα isshown in figure 3.17. The energy resolution of the BI, which had degraded more rapidly than that of
the FI in the SCI mode, was improved by the change of the injected charge amount.

3.3.4 Non X-ray background
One of the advantages of the XIS compared with Chandra and XMM-Newton is its low and stable
non X-ray background (NXB; figure 3.18), which is thanks to its lower altitude of the orbit. It is
known that the NXB of the XIS is very stable on time scales of months and thus the NXB spectrum
can be constructed using data obtained when the spacecraft is pointed toward the Earth at night. The
NXB of the XIS is caused by ionization losses of charged cosmic-ray particles, fluorescence X-rays
from materials used in the spacecraft, and the 55 Fe calibration sources. Figure 3.19 shows spectraof the NXB extracted from the whole region of the CCD, except for the regions illuminated by thecalibration sources. The spectra show fluorescence lines from aluminum, silicon, gold, manganese,and nickel in the XIS and the XRT. The 5.9 keV and 6.5 keV lines are due to scattered X-rays from thecalibration sources (manganese Kα and Kβ). The XIS 0 shows relatively strong manganese K-shelllines, which is due to stray X-rays from the 55 Fe calibration source although the reason is unknown
why the radiation was detected outside the calibration source regions. The background counting ratein the 0.4–12 keV band is 0.1–0.2 counts s−1 for the FI CCDs and 0.3–0.6 counts s−1 for the BI CCD.
The intensity of the NXB is not uniform over a CCD chip. Figure 3.20a represents position dependence of the NXB. When the pixel position is farther from the readout node (namely, larger ACTY),
the NXB becomes higher. This is because some fraction of the NXB is generated in the frame-store
region and pixels in larger ACTY take longer time to be read out: the integrated intensity of the NXB
during readout becomes higher in larger ACTY pixels. The NXB strongly depends on the geomagnetic cut-oﬀ rigidity. Figure 3.20b shows the NXB count rate as a function of the cut-oﬀ rigidity.
Tawa et al. (2008) successfully modeled the NXB spectra as a function of the cut-oﬀ rigidity and theposition on the CCDs. The model reproduce the NXB spectra within the accuracy of 5%.
The good energy resolution, large eﬀective area, and the low and stable NXB background of the3.3. X-RAY IMAGING SPECTROMETER (XIS)33

Figure 3.18. Background counting rate as a function of energy for the XIS and other instruments (Mitsuda
et al., 2007). The background rate was normalized with the eﬀective area and the field of view, which is a goodmeasure of the sensitivity determined by the background for spatially extended sources.

XIS oﬀer a good spectral performance especially for diﬀuse sources with low surface brightness suchas the GDXE in the 5–10 keV band, including iron K-shell lines.

CHAPTER 3. INSTRUMENT

XIS 0
XIS 1 (CI=2 keV)
XIS 1 (CI=6 keV)
XIS 3

10−7

0

2

4

6
8
Energy (keV)

10

Au Lβ

Au Lα

Ni Kα

Ni Kβ

10−9

Mn Kα
Mn Kβ

10−8

Al Kα
Si Kα
Au Mα

Counts s−1 keV −1 pixel−1

34

12

14

Figure 3.19. Spectra of the NXB for XIS0 (black), XIS1 (CI=2 keV: red, CI=6 keV: green), and XIS3 (blue),
which are constructed from night Earth observations (Suzaku Technical Description).

(a)

(b)

Figure 3.20. Coordinate and cut-oﬀ rigidity dependency of the NXB. (a) Coordinate dependence of the intensity of the NXB for XIS 0. Black lines indicate the continuum component (2.5–5.5 keV), while the red lines
indicate the nickel Kα line component (7.2–7.8 keV). (b) Cut-oﬀ rigidity dependency of the NXB (averageintensity for 5–10 keV). Both are taken from Suzaku Technical Description.

Chapter 4

Observations and Data Reductions
In this dissertation, we used the Suzaku observation data of the Galactic plane, including theseries of the Key and Large projects on and near the Galactic center region as well as the archival
data. We chose the data which were obtained in the normal clocking mode without window option.
We reprocessed the cleaned data by using xispi in the analysis software package HEAsoft, and theSuzaku calibration database (CALDB) version 2015 March. The observation log is listed in table 4.1.
We used in total 153 pointings with the total exposure time of 8.8 Ms. This is 2.5 times larger thanthe data set analyzed by Uchiyama et al. (2013).
Using xissimarfgen (Ishisaki et al., 2007) and xisrmfgen, the eﬀective area of the XRT and the
response of the XIS were calculated. The ancillary response files (arf) and redistribution file (rmf)
for the XIS sensors were produced for each observation. The NXB with the cut-oﬀ rigidity values wasestimated with xisnxbgen (Tawa et al., 2008) released by the Suzaku team. We subtracted the NXB
from the raw spectral data with the same cut-oﬀ rigidity.

35

CHAPTER 4. OBSERVATIONS AND DATA REDUCTIONS

36

Table 4.1. Observation log.
Obs. ID
100026020
100026030
100027020
100028010
100028020
100028030
100037010
401026010
401052010
401054010
401056010
401092010
401101010
402094010
403005010
404020010
404056010
404081010
405025010
405027010
405028010
406069010
406078010
407018010
407020010
407021010
407091010
407092010
407094010
408021010
500007010
500008010
500009010
500009020
501008010
501009010
501042010
501043010
501044010
501045010
501046010
501052010
501053010
501060010
501099010
501105010
502001010

Pointing directionl (◦ )b (◦ )
347.630.71
345.81 −0.54
359.76 −0.05
332.40 −0.15
332.00 −0.15
332.70 −0.15
359.76 −0.05
25.27 −0.10
340.000.11
341.320.58
333.490.30
9.94 −0.27
12.87
0.01
9.95 −0.27
288.24 −0.47
33.17
0.84
338.000.08
29.77 −0.20
311.640.22
340.44 −0.18
54.62
0.11
10.00 −0.23
340.17 −0.12
333.61 −0.20
333.720.22
316.980.19
333.89
0.41
8.14
0.19
322.06 −0.42
352.17 −0.27
8.44 −0.05
8.04 −0.04
28.46 −0.21
28.46 −0.21
359.84 −0.19
359.930.18
331.57 −0.53
330.40 −0.38
17.87 −0.70
18.44 −0.84
359.830.34
358.50 −0.00
358.17 −0.00
1.50
0.00
287.89 −0.37
348.80 −0.54
11.95 −0.09

Obs. start(UT)
2005-09-25 19:11:40
2005-09-28 07:09:13
2005-09-24 14:17:17
2005-09-19 12:00:02
2005-09-18 22:47:36
2005-09-20 19:40:17
2005-09-29 04:35:41
2007-03-05 12:49:14
2006-09-09 09:12:56
2006-10-05 21:10:30
2006-09-20 20:25:12
2006-09-09 22:13:43
2007-03-01 21:35:58
2007-10-14 05:35:49
2008-11-30 23:02:01
2010-03-12 00:22:33
2010-03-12 23:40:40
2009-04-15 19:37:17
2011-01-28 13:24:43
2011-02-11 03:42:23
2010-04-27 10:28:40
2012-03-24 10:47:46
2012-02-23 22:39:06
2012-08-21 23:55:27
2012-08-19 12:30:06
2012-08-04 06:47:27
2012-08-18 19:16:15
2012-09-21 08:05:31
2012-08-15 15:34:02
2013-09-05 00:04:43
2006-04-06 14:41:18
2006-04-07 11:49:16
2005-10-28 02:40:08
2006-10-15 02:15:12
2006-09-26 14:18:16
2006-09-29 21:26:07
2006-09-15 16:00:48
2006-09-16 11:02:03
2006-10-17 19:37:16
2006-10-19 04:03:16
2007-03-10 15:03:10
2006-10-10 06:45:09
2006-10-10 21:18:59
2007-03-17 05:07:04
2006-06-05 13:21:42
2007-02-23 08:36:38
2007-10-02 14:17:21

Exp.a
(ks)
34.9
37.5
42.8
41.4
19.3
21.9
43.7
42.2
22.5
21.2
39.1
48.9
63.8
52.2
100.4
34.9
50.6
104.3
24.6
20.9
23.9
70.6
149.8
40.5
44.3
41.1
29.3
32.0
30.0
37.3
37.5
40.7
93.3
98.9
129.6
51.2
40.2
43.6
50.3
52.2
25.2
21.0
23.4
68.1
76.9
20.7
53.8

XIS 2bWin.c

CId

Target name

OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
OK
-

XIS3
XIS3
-

oﬀ
oﬀ
oﬀ
oﬀ
oﬀ
oﬀ
oﬀ
on
oﬀ
oﬀ
oﬀ
oﬀ
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
oﬀ
oﬀ
oﬀ
on
on
on
oﬀ
oﬀ
oﬀ
oﬀ
on
on
on
on
oﬀ
oﬀ
on

RXJ1713−3946−BKGD1
RXJ1713−3946−BKGD2
Sgr_A_west
HESS J1616−508
HESS J1616−508_BGD1
HESS J1616−508_BGD2
Sgr_A_west
HESS J1837−069
IGRJ16465−4507
IGRJ16493−4348
IGRJ16195−4945
SGR 1806−20
HESSJ1813−178
SGR 1806−20
1E 1048.1−5937
V603 AQL
IGRJ16393−4643
PSR J1846−0258
EMS0918
EMS01095
EMS1308
SGR 1806−20
IGRJ16479−4514
AX J1622.1−5005
AX J1620.7−4942
AX J1445.7−5931
J1620−4927
J1803−2149
J1522−5734
CXO J172641.7−354052
HESS J1804−216
HESS J1804−216 BGDGALACTIC RIDGEGALACTIC RIDGEGC SOUTHGC SOUTH BGDHESS J1614-518
HESS J1614-518 BG
HESS J1825-137
HESS J1825-137 BG
GALACTIC CENTERGALACTIC CENTERGALACTIC CENTER
GC SGR D EAST
CAR-D1HESS J1718-385
G12.0-0.1

a Eﬀective exposure time of the cleaned data.
b Availability of the XIS 2 sensor.
c The sensor which is not used for analysis because its observation is performed in the window mode.
d The SCI technique was operated (on) or not (oﬀ) in the observation.

37
Table 4.1. Observation log.
Obs. ID
502002010
502003010
502004010
502005010
502006010
502007010
502008010
502009010
502049010
502053010
502059010
502068010
502077010
503007010
503008010
503014010
503015010
503016010
503017010
503018010
503018020
503018030
503019010
503020010
503022010
503023010
503026010
503027010
503028010
503029010
503030010
503072010
503073010
503074010
503078010
503079010
503081010
503087010
503099010
503103010
503108010
504001010
504002010
504003010
504034010
504036010
504043010

Pointing directionl (◦ )b (◦ )
0.17 −0.67
359.83 −0.67
0.17 −1.00
359.83 −1.00
0.17
0.33
0.17
0.66
359.830.66
1.83 −0.00
344.26 −0.22
15.82 −0.85
360.00 −2.00
296.03 −0.32
54.06
0.21
0.33
0.17
0.00 −0.38
357.90 −0.05
357.65 −0.05
357.40 −0.05
357.15 −0.05
356.90 −0.05
356.90 −0.05
356.90 −0.05
356.65 −0.05
356.40 −0.05
356.000.70
356.330.70
5.89 −0.38
5.72 −0.06
17.61 −0.84
17.73 −0.44
17.47 −0.58
359.580.17
331.30 −0.76
331.47 −0.64
11.03
0.07
10.84
0.04
0.03 −1.66
12.82 −0.02
359.781.13
359.991.20
348.92 −0.45
358.53 −0.26
358.47 −0.58
358.55 −0.87
314.40 −0.14
357.71 −0.12
40.17 −0.88

Obs. start(UT)
2007-10-09 16:40:54
2007-10-10 03:41:13
2007-10-10 15:21:17
2007-10-11 01:01:17
2007-10-11 11:34:01
2007-10-11 23:09:15
2007-10-12 09:52:59
2007-10-12 21:52:24
2008-03-25 11:00:23
2007-10-07 02:16:29
2007-09-29 01:40:51
2007-08-09 06:05:33
2007-10-30 15:56:49
2008-09-02 10:15:27
2008-09-03 22:53:29
2008-09-18 04:46:49
2008-09-19 07:33:05
2008-09-22 06:47:49
2008-09-23 08:08:10
2008-09-24 09:27:54
2008-10-03 18:05:13
2009-02-19 07:32:01
2009-02-19 16:37:49
2009-02-21 01:15:55
2009-03-18 23:11:24
2009-03-26 06:37:01
2008-04-06 07:34:41
2008-04-07 00:21:13
2008-10-15 21:49:50
2008-10-17 11:01:12
2008-10-19 04:41:55
2009-03-06 02:39:12
2008-09-20 18:18:35
2008-09-21 13:31:03
2008-03-31 14:05:55
2008-04-01 16:33:52
2009-03-09 15:41:50
2009-03-04 19:45:39
2009-03-10 19:39:08
2009-03-11 10:56:59
2008-08-28 05:36:23
2010-02-26 09:15:00
2010-02-27 16:14:41
2010-02-25 04:33:17
2010-01-13 13:49:39
2009-08-29 12:05:20
2009-04-15 03:40:44

Exp.a
(ks)
23.2
21.5
19.9
20.6
22.6
22.0
23.8
22.9
215.7
71.5
136.8
77.2
101.6
52.2
53.7
59.7
61.4
57.4
56.5
31.9
13.3
12.9
56.8
66.2
41.3
31.2
31.7
31.0
57.2
57.2
55.5
140.6
53.7
52.6
51.5
44.2
59.2
56.2
29.7
18.3
23.5
51.2
53.1
50.9
104.1
136.5
24.5

XIS 2bWin.c

CId

Target name

-

-

on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on

GC14
GC15
GC16
GC17
GC18
GC19
GC20
GC21
HESS J1702-420
M17 EAST BKG
GALACTIC BULGE
G296.1-0.5N
G54.1+0.3
GC_LARGEPROJECT1
GC_LARGEPROJECT2
GC_LARGEPROJECT8
GC_LARGEPROJECT9
GC_LARGEPROJECT10
GC_LARGEPROJECT11
GC_LARGEPROJECT12
GC_LARGEPROJECT12
GC_LARGEPROJECT12
GC_LARGEPROJECT13
GC_LARGEPROJECT14
LOOP 1 L=356.00
LOOP 2 L=356.33
HESS J1800−240B
HESS J1800−240C
HESS J1825−137
HESS_J1825-137_2
HESS_J1825-137_3
EXTENDED CHIMNEYHESSJ1614−A
HESSJ1614−B
HESS_J1809-193_N
HESS_J1809-193_S
GC(0,−1.6)
G12.8−0.0
GCL1
SLX 1737−282
HESS J1718−385
HESS J1745−303A
HESS J1745−303B
HESS J1745−303C
HESS J1427−608
TORNADO
ASO0431a Eﬀective exposure time of the cleaned data.
b Availability of the XIS 2 sensor.
c The sensor which is not used for analysis because its observation is performed in the window mode.
d The SCI technique was operated (on) or not (oﬀ) in the observation.

CHAPTER 4. OBSERVATIONS AND DATA REDUCTIONS

38

Table 4.1. Observation log.
Obs. ID
504044010
504045010
504046010
504049010
504050010
504052010
504077010
504078010
504079010
504088010
504089010
504090010
504091010
504092010
504093010
504094010
504095010
504099010
505025010
505026010
505027010
505029010
505030010
505049010
505050010
505051010
505052010
505074010
505076010
505079010
505080010
505082010
505083010
505084010
505085010
505086010
505088010
505089010
505090010
505091010
506019010
506021010
506051010
507027010
507044010
507068010
507069010

Pointing directionl (◦ )b (◦ )
65.88 −0.33
285.07 −0.49
65.26
0.41
356.290.99
0.10 −1.42
26.45
0.13
11.61 −0.25
11.33 −0.06
10.72
0.33
360.00 −0.83
359.95 −1.20
358.51 −1.18
358.50 −1.60
358.56 −2.15
358.50 −2.80
358.50 −3.80
358.50 −5.00
25.50
0.01
22.00
0.00
23.49
0.04
37.00 −0.09
319.620.29
319.320.29
339.01 −0.93
339.79 −1.14
339.43 −0.80
340.77 −1.01
304.580.13
347.85 −0.23
359.96 −2.79
359.95 −3.80
356.41 −0.39
356.39 −0.81
356.41 −1.49
356.39 −2.31
356.39 −3.51
26.31 −0.00
26.40 −0.31
26.71 −0.15
27.13 −0.28
36.00
0.05
23.30
0.31
18.78
0.40
318.36 −0.44
19.57
0.01
337.21 −0.73
2.00 −0.04

Obs. start(UT)
2009-05-23 15:19:43
2009-07-08 08:12:56
2009-04-07 18:02:07
2009-09-08 03:36:55
2010-03-06 03:55:37
2009-04-13 15:32:05
2009-09-09 04:55:29
2009-09-10 11:36:47
2009-09-11 17:59:44
2009-10-14 11:30:56
2009-10-09 04:05:59
2009-10-13 04:17:20
2009-09-14 19:37:36
2009-09-16 07:21:35
2009-09-17 13:54:31
2009-09-19 03:40:31
2009-10-15 15:32:28
2009-04-06 02:57:46
2010-04-16 14:27:26
2010-10-20 13:34:39
2010-04-17 17:35:13
2010-08-05 03:45:12
2010-08-05 22:21:12
2010-09-10 18:33:35
2010-09-12 04:52:53
2010-09-23 06:09:00
2010-09-24 09:32:21
2010-09-03 21:54:15
2011-02-16 01:17:29
2011-03-12 06:36:20
2010-04-07 17:15:10
2011-03-15 13:54:28
2010-10-10 14:04:38
2011-03-06 05:36:44
2010-10-13 06:04:47
2010-10-14 13:31:50
2011-03-25 07:00:01
2011-03-26 10:41:12
2011-03-27 15:21:16
2011-03-28 19:01:23
2011-09-18 22:04:29
2011-04-08 06:06:16
2012-03-08 22:01:58
2012-09-15 02:32:40
2012-10-15 13:05:48
2012-09-02 13:16:23
2013-03-15 09:48:19

Exp.a
(ks)
38.7
22.8
20.5
17.9
100.4
41.1
51.9
52.5
51.0
47.2
55.3
41.3
51.3
50.9
53.2
93.1
48.3
52.7
50.5
49.0
51.0
51.4
51.6
51.9
52.7
50.2
49.6
99.6
32.6
50.2
56.1
54.8
56.9
50.3
55.0
53.1
49.7
50.0
49.6
51.3
40.9
40.3
52.0
31.0
171.8
304.2
110.3

XIS 2bWin.c

CId

Target name

-

-

on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on
on

ASO0449
ASO0218
ASO0447
ASO0376
NONE
ASO0418HESS J1809−193A
HESS J1809−193B
HESS J1809−193E
GALACTIC_BULGE1GALACTIC_BULGE2GALACTIC_BULGE7GALACTIC_BULGE8GALACTIC_BULGE9GALACTIC_BULGE10GALACTIC_BULGE11
GALACTIC_BULGE12
G25.5+0.0
G22.0+0.0
G23.5+0.1
G37.0−0.1
FVW319.8+0.3
FVW319.8+0.3_BGD
WESTERLUND1_1
WESTERLUND1_2
WESTERLUND1_3
WESTERLUND1_4
SNR G304.6+0.1
AX J1714.1−3912
GALACTIC_BULGE4GALACTIC_BULGE5GALACTIC_BULGE13
GALACTIC_BULGE14
GALACTIC_BULGE15
GALACTIC_BULGE16
GALACTIC_BULGE17
HESS_J1741-055_1
HESS_J1741-055_2
HESS_J1741-055_3
HESS_J1741-055_4
HESSJ1857−A
HESSJ1832
G18.8+0.3
HESS J1457−593
NONE
G337.2−0.7
GRXE_E_2

a Eﬀective exposure time of the cleaned data.
b Availability of the XIS 2 sensor.
c The sensor which is not used for analysis because its observation is performed in the window mode.
d The SCI technique was operated (on) or not (oﬀ) in the observation.

39
Table 4.1. Observation log.
Obs. ID
507070010
507071010
507072010
507073010
507074010
507075010
508021010
508022010
508075010
508076010
508077010
509038010

Pointing directionl (◦ )b (◦ )
2.25 −0.04
2.50 −0.04
2.75 −0.04
3.00 −0.04
3.15
0.16
3.25
0.41
35.57 −0.59
35.61 −0.40
1.75 −0.04
3.25 −0.04
3.50 −0.04
39.19 −0.30

Obs. start(UT)
2013-03-17 18:39:56
2013-03-20 02:41:04
2013-03-22 07:20:36
2013-03-24 08:46:03
2013-04-03 21:33:46
2013-03-11 09:06:19
2013-10-27 20:00:53
2013-10-28 23:25:55
2014-03-10 01:33:32
2014-02-28 12:46:16
2014-03-02 17:00:51
2014-04-26 06:42:08

Exp.a
(ks)
111.8
112.3
110.7
108.9
104.1
109.6
51.7
52.6
109.3
109.8
109.4
82.8

XIS 2bWin.c

CId

Target name

-

-

on
on
on
on
on
on
on
on
on
on
on
on

GRXE_E_3
GRXE_E_4
GRXE_E_5
GRXE_E_6
GRXE_E_9
GRXE_E_10
HESSJ1858−A
HESSJ1858−B
GRXE_E_1
GRXE_E_7
GRXE_E_8
3C 396a Eﬀective exposure time of the cleaned data.
b Availability of the XIS 2 sensor.
c The sensor which is not used for analysis because its observation is performed in the window mode.
d The SCI technique was operated (on) or not (oﬀ) in the observation.

Chapter 5

Result I: Global structure of the GDXE–the 6.4 keV and 6.7 keV lines and thecontinuum emission
5.1 Analytical procedure
We make an X-ray image in the energy band of 2–8 keV for each FOV. We exclude point-likesources with flux higher than 1 × 10−13 erg s−1 cm−2 and bright diﬀuse sources from the observation
data. The FOV images and excluded regions are summarized in Appendix. In some FOVs, a fan-likeshaped emission is clearly seen. Since the structures are diﬀerent among XIS 0, 1, 2 (if it is valid),
and 3, this emission must be artificial and caused by stray light from a bright X-ray point source
outside the FOV. Such regions require special care. Most of the stray light sources which aﬀect ourused data are low mass X-ray binaries: GX 3+1, GX 5−1, GX 340+0, and Cir X−1. Low mass X-ray
binaries generally have continua (e.g. blackbody, disk blackbody, or power law) often accompanied
by a broad emission line; do not have narrow emission lines. This means that the stray light does not
aﬀect the observed intensities of the 6.4 keV, 6.7 keV and 7.0 keV lines while it aﬀects the continuum
flux.
Next, we extract a spectrum from each pointing. The spectra in the 4–10 keV band are fitted witha phenomenological model which consists of four Gaussian lines and an absorbed bremsstrahlungas well as the CXB model (Kushino et al., 2002). The line energies of the Gaussians are fixed to6.40 keV (neutral iron Kα), 6.68 keV (helium-like iron Kα), 6.97 keV (hydrogen-like iron Lyα), and7.06 keV (neutral iron Kβ) (Smith et al., 2001; Kaastra and Mewe, 1993). The intensity of the neutral
iron Kβ line is fixed to 0.125 times that of the neutral iron Kα line (Kaastra and Mewe, 1993), while
those of the other lines are free parameters. The interstellar absorption column densities are fixed to
6 × 1022 cm−2 for the GCXE region (|l| < 5◦ and |b| < 0.◦ 5 ), 1 × 1022 cm−2 for the GBXE region (|b| > 1◦ ),and 3 × 1022 cm−2 for the GRXE region (|l| > 5◦ and |b| < 1◦ ) (e.g. Uchiyama et al., 2013; Ebisawa et al.,2001). The normalization and temperature of the bremsstrahlung component are allowed to vary.
We then obtain the intensities of the 6.4 keV, 6.7 keV, and 7.0 keV lines and the continuum in the
5–8 keV band. Here we exclude the contribution of the CXB. We show an example of the spectral
fitting in figure 5.1.
41

CHAPTER 5. RESULTS I: GLOBAL STRUCTURE OF THE GDXEχCounts s−1 keV−1 arcmin−2

42He-like Fe Kα
H-like Fe Kα
neutral Fe Kβ

neutral Fe Kα−410

bremsstrahlung

10−5

CXB

10−6
4
2
0
−2
−4
4

5

6
Energy (keV)

7

8

9

Figure 5.1. Example of the spectrum (ObsID=503020010) fitted by the phenomenological model consisting ofan absorbed bremsstrahlung and Gaussians as well as the CXB.

5.2 Latitudinal distribution of the iron lines and continuum of the GDXE
Using the observed intensities of the 6.4 keV, 6.7 keV, and 7.0 keV lines and the continuum of the
GDXE, we investigate the spatial distribution of them. The intensity of each component decreases
with the Galactic latitude approximately in an exponential manner with an e-folding scale height,
which is one of the most important keys to resolving its origin. Figure 5.2 and 5.3 shows latitudinal
distributions of the 6.4 keV, 6.7 keV, and 7.0 keV lines, and the 5–8 keV flux in the regions of (a)|l| < 0.◦ 5, (b) l = 358.◦ 5, (c) l = 356.◦ 0–356.◦ 4, (d) 10◦ < l < 30◦ , and (e) 330◦ < l < 350◦ .
Each distribution of the upper panels in figure 5.2 and figure 5.3 cannot be represented by asingle exponential function. Thus, the data plots in and near the GC region, especially for the |l| < 0.◦ 5
region, require two exponential functions. We simultaneously fit the data of (a), (b) and (c) with themodel described by

AGCXE × exp(−|b∗ |/bGCXE ) + AGBXE × exp(−|b∗ |/bGBXE ).

(5.1)

The normalizations, AGCXE and AGBXE , would vary depending on the Galactic longitude. We
scaled the AGCXE assuming the e-folding scale length of 0.◦ 63 (Uchiyama et al., 2013). The normalization of AGBXE are free parameters among the three regions. The e-folding scales of bGCXE andbGBXE in (b) and (c) are linked to those in (a).
The data of (d) and (e) are fitted with an exponential model described by

AGRXE × exp(−|b∗ |/bGRXE ) + A2 × exp(−|b∗ |/bGRXE ).

(5.2)

5.2. LATITUDINAL DISTRIBUTION OF THE GDXE

43
l = 0.°0, 6.7 keV line

0

1

2
3
Absolute b* (degree)

4

10−6
10−7
10−8

Intensity (photons s−1 cm−2 arcmin−2)

10−7
10−8
10−9

Intensity (photons s−1 cm−2 arcmin−2)

10−6

l = 0.°0, 6.4 keV line

5

0

1

3
2
Absolute b* (degree)

4

5

4

5

4

5

l = 358.°5, 6.7 keV line

0

1

3
2
Absolute b* (degree)

4

10−6
10−7
10−8

Intensity (photons s−1 cm−2 arcmin−2)

10−7
10−8
10−9

Intensity (photons s−1 cm−2 arcmin−2)

10−6

l = 358.°5, 6.4 keV line

5

0

1

3
2
Absolute b* (degree)l = 356.°0−356.°4, 6.7 keV line

0

1

2
3
Absolute b* (degree)

4

10−6
10−7
10−8

Intensity (photons s−1 cm−2 arcmin−2)

10−7
10−8
10−9

Intensity (photons s−1 cm−2 arcmin−2)

10−6

l = 356.°0−356.°4, 6.4 keV line

5

0

3
2
Absolute b* (degree)
10° < |l| < 30°, 6.7 keV line

0

0.5
1
Absolute b* (degree)

1.5

10−6
10−7
10−8

Intensity (photons s−1 cm−2 arcmin−2)

10−7
10−8
10−9

Intensity (photons s−1 cm−2 arcmin−2)

10−6

10° < |l| < 30°, 6.4 keV line

1

0

0.5
1
Absolute b* (degree)

1.5

Figure 5.2. Latitudinal distributions of the 6.4 keV line and the 6.7 keV line flux in the regions of |l| < 0.◦ 5,
l = 358.◦ 5, l = 356.◦ 0–356.◦ 4 and 10◦ < |l| < 30◦ . The solid lines show the best-fit model. The red and black
data and lines in the bottom panel show the data and best-fit result in the east (10◦ < l < 30◦ ) and the west
(330◦ < l < 350◦ ), respectively.

CHAPTER 5. RESULTS I: GLOBAL STRUCTURE OF THE GDXE

44

l = 0.°0, 5−8 keV

0

1

2
3
Absolute b* (degree)

4

10−5
10−6
10−7

Intensity (photons s−1 cm−2 arcmin−2)

10−7
10−8
10−9

Intensity (photons s−1 cm−2 arcmin−2)

10−6

l = 0.°0, 7.0 keV line

5

0

1

l = 358.°5, 7.0 keV line32
Absolute b* (degree)

4

5

4

5

4

5

0

1

2
3
Absolute b* (degree)

4

10−5
10−6
10−7

Intensity (photons s−1 cm−2 arcmin−2)

10−7
10−8
10−9

Intensity (photons s−1 cm−2 arcmin−2)

10−6

l = 358.°5, 5−8 keV

0

5

1

l = 356.°0−356.°4, 7.0 keV line32
Absolute b* (degree)

0

1

2
3
Absolute b* (degree)

4

10−5
10−6
10−7

Intensity (photons s−1 cm−2 arcmin−2)

10−7
10−8
10−9

Intensity (photons s−1 cm−2 arcmin−2)

10−6

l = 356.°0−356.°4, 5−8 keV

50

10° < |l| < 30°, 7.0 keV line132
Absolute b* (degree)

0

0.5
1
Absolute b* (degree)

1.5

10−5
10−6
10−7

Intensity (photons s−1 cm−2 arcmin−2)

10−7
10−8
10−9

Intensity (photons s−1 cm−2 arcmin−2)

10−6

10° < |l| < 30°, 5−8 keV

0

0.5
1
Absolute b* (degree)

1.5

Figure 5.3. Same as figure 5.2, but those of the 7.0 keV line and the continuum in the 5–8 keV band.

5.2. LATITUDINAL DISTRIBUTION OF THE GDXE

45Since the longitudinal distribution of the GDXE on the plane is asymmetric with respect to the
GC especially for |l| > 10◦ (e.g. Revnivtsev et al., 2006a), the normalization AGRXE are varied between
the east and west longitude regions. The scale heights of bGRXE in both the regions are linked toeach other. We exclude the regions contaminated by stray light in fitting the distributions of thecontinuum. The best-fit parameters are summarized in table 5.1.
If we assume the distance to the emission region to be 8 kpc, the scale heights of the 6.7 keV lineof the GCXE, GBXE, and GRXE are 36 ± 3 pc, 314 ± 95 pc, and 142 ± 17 pc, respectively, and thoseof the 7.0 keV and the continuum coincide within the error with them. On the other hand, the scaleheights of the 6.4 keV line in the GCXE, GBXE, and GRXE are 31 ± 3 pc, 161 ± 50 pc, and 70 ± 17 pc,
respectively, which are systematically smaller than those of the 6.7 keV line.

Table 5.1. Scale height of the GDXE*.
Region(a) l = 0◦(b) l = 358.◦ 5(c) l = 356.◦ 2(d) l = 10◦ –30◦(e) l = 330◦ –350◦

Component

6.4 keV
6.7 keV
7.0 keV
5–8 keV
6.4 keV
6.7 keV
7.0 keV
5–8 keV
6.4 keV
6.7 keV
7.0 keV
5–8 keV
6.4 keV
6.7 keV
7.0 keV
5–8 keV
6.4 keV
6.7 keV
7.0 keV
5–8 keVGCXE
AGCXE †bGCXE
(degree)4.1±0.2 0.22±0.02
11.9±0.6 0.26±0.02
4.9±0.2 0.24±0.02
77.0±3.6 0.25±0.02
0.45‡
0.22§
‡
1.31
0.26§
0.54‡
0.24§
‡
8.47
0.25§
0.02‡
0.22§
‡
0.05
0.26§
0.02‡
0.24§
‡
0.31
0.25§
-

GBXE
IGBXE †
0.31±0.15
1.14±0.34
0.40±0.1211.9±1.8
0.35±0.10
1.15±0.27
0.39±0.10
10.6±1.4
0.28±0.07
1.04±0.21
0.19±0.06
7.2±0.9-bGBXE(degree)1.15±0.36
2.25±0.68
2.13±0.66
1.96±0.25
1.15§
2.25§
2.13§
1.96§
1.15§
2.25§
2.13§
1.96§
-

GRXE
AGRXE †
0.23±0.03
0.76±0.02
0.09±0.02
5.8±0.4
0.28±0.04
0.54±0.03
0.09∥
4.9±0.5

bGRXE(degree)0.50±0.12
1.02±0.12
0.71±0.29
1.04±0.20
0.5¶
1.0¶
0.71¶
1.0¶

* The error ranges in this table are calculated at 68% confidence levels.
† The units are 10−7 photons s−1 cm−2 arcmin−2 .
‡ These normalizations are fixed to I◦◦
GCXE(l=0◦ ) × 0.11 at l = 358. 5 and IGCXE(l=0◦ ) × 0.004 at l = 356. 2 based on thescale length of the GCXE obtained by Uchiyama et al. (2013).
§ These scale heights are linked to those of l = 0◦ .
∥ The normalization is linked to that of l = 10◦ –30◦ .
¶ These scale heights are linked to those of l = 10◦ –30◦ .

46

5.3

CHAPTER 5. RESULTS I: GLOBAL STRUCTURE OF THE GDXE

Equivalent widths of the iron K-shell lines

We investigate a typical value of the equivalent width of the GCXE (|l| < 1.◦ 5 and |b| ≤ 0.◦ 5), GBXE(|l| < 4.◦ 0 and |b| ≥ 1.◦ 0), and GRXE (10◦ < |l| < 30◦ and |b| < 1.◦ 0). The mean equivalent widths of the
6.4 keV line in the GCXE, GBXE and GRXE are calculated to be 145 ± 3 eV, 61 ± 11 eV, and 97 ± 12 eV,
respectively, while those of the 6.7 keV lines are 527 ± 4 eV, 443 ± 14 eV, 428 ± 15 eV, and those of the
7.0 keV lines are 221±3 eV, 160±14 eV, and 117±19 eV, respectively. Figure 5.4 shows the equivalent
width diagrams of the 6.4 keV and 6.7 keV lines in the GCXE, GBXE, and GRXE.

5.4

Intensity profile along the Galactic longitude

In order to investigate the longitudinal distributions of the 6.4 keV line along the Galactic plane,
especially in the GRXE region (outside the range in figure 1.10), we make the intensity profile of
the 6.4 keV line, and for a comparison, those of the 6.7 keV and 7.0 keV lines and the continuum
in the 5–8 keV band in the |l| > 1.◦ 5 region as shown in figure 5.5. The data points are extracted
from the |b| < 0.◦ 5 and are not corrected for the b-dependence. The continuum emissions in 1.◦ 5 < l <3.◦ 5, l ∼ −20◦ , and l ∼ −38◦ are contaminated by stray light from GX 3+1, GX 340+0, and Cir X−1,
respectively. In particular, the 1.◦ 5 < l < 3.◦ 5 and l ∼ −38◦ regions are severely contaminated by brightsecond-reflection components. We excluded the regions contaminated by the second reflection fromthe 1.◦ 5 < l < 3.◦ 5 and the l ∼ −38◦ regions; the data points in the two regions are still contaminatedby backside reflection. In the l ∼ −20◦ region, we leave the stray-light component in the data points
because the stray light in the region is not so bright.
We discover enhancements of the 6.4 keV line intensity distribution compared with the 6.7 keVline. The flux ratio of the 6.4 keV and 6.7 keV lines in figure 5.5e clearly shows local excesses in the1.◦ 5 < l < 3.◦ 5 and −20◦ < l < −30◦ regions. These enhancements should have a diﬀerent origin fromthat of the 6.7 keV and and the continuum. In the following chapter, we will discuss the origin of
these enhancements.

5.4. INTENSITY PROFILE ALONG THE GALACTIC LONGITUDE

47

0

200

EW6.4keV

400

600

GCXE

0

200

400

600
EW6.7keV

800

1000

1200

800

1000

1200

800

1000

1200

0

200

EW6.4keV

400

600

GBXE

0

200

400

600
EW6.7keV

0200

EW6.4keV

400

600

GRXE

0

200

400

600
EW6.7keV

Figure 5.4. Equivalent width diagrams of the 6.4 keV and 6.7 keV lines in the GCXE (top), the GBXE (middle),
and the GRXE (bottom).

CHAPTER 5. RESULTS I: GLOBAL STRUCTURE OF THE GDXE

48

10 −5
10 −6

(a) 6.4 keV

10 −7

−2 −1

−2

X−ray flux (ph cm s arcmin )

10 −8
10 −5
10 −6

(b) 6.7 keV

10 −7
10 −8
10 −5
10 −6

(c) 7.0 keV

10 −7
10 −8
10 −4
10 −5

(d) 5−8 keV

10 −6
10 −7

Ratio

1.5

(e) 6.4 keV / 6.7 keV

1
0.5
0
1

2

20
10
Galacitc longitude (degree)

100

Figure 5.5. Intensity profiles of X-rays along the Galactic plane. The red and black circles indicate the datain the east and west regions, respectively. (a) The 6.4 keV line. (b) The 6.7 keV line. (c) The 7.0 keV line. (d)The continuum in the 5–8 keV band. The open circles are the data contaminated by stray light (see text). (e)Intensity ratio of the 6.4 keV and 6.7 keV lines. The dashed line shows the averaged value on the Galactic
plane (|l| > 1.◦ 5). The two regions indicated by the arrows, namely the 1.◦ 5 < l < 3.◦ 5 and the −20◦ < l < −30◦
regions, show excesses in the ratio.

Chapter 6

Results II: Discovery of local
enhancements in the 6.4 keV line
Figure 5.5 shows local enhancements in the two regions in the longitudinal distribution of the
6.4 keV line; (a) 1.◦ 5 < l < 3.◦ 5 and (b) −20◦ < l < −30◦ regions. In this chapter, we perform a moredetailed analysis of the two regions. We analyze all the data in the 1.◦ 5 < l < 3.◦ 5 region for theregion (a), and call them the GC east region. For analysis of the region (b), we use the data of thethree FOVs (ObsID=406078010, 401052010, and 404056010) which have the strongest intensity of
the 6.4 keV line, in order to increase the signal-to-noise ratio, and call them the l = −20◦ region. We
caution that most of the FOVs in the (a) and (b) regions, unfortunately, are aﬀected by stray light. It
is important to estimate the stray light in extracting and analysing the spectra of the enhancements.
First we describe the estimation of the stray light, and then we investigate the spectral features of theenhancements.

6.1 Estimation of stray light
6.1.1 Stray light in the GC east region
A bright low mass X-ray binary, GX 3+1, is located at (l, b) = (2.◦ 294, 0.◦ 794), which is ∼ 50′ –90′
apart from each pointing in the GC east region. Some regions of the FOVs are contaminated by thestray light from GX 3+1. Stray light pattern in the FOVs are shown in figure 6.1 with the position
of GX 3+1 and the ObsIDs. Most of the pointings are contaminated by stray light of secondary andbackside reflection. Since two FOVs (ObsID = 508077010 and 508076010) are 89′ and 77′ apartfrom GX 3+1, respectively, these observations are free from secondary-only stray (Suzaku Technical Description); the region in the FOVs near side to the stray light source is free from stray light.
Some FOVs (ObsID = 508075010, 502009010, 507075010, and 507074010) are located at quadrant
boundary, and thus a part of the FOVs are stray-light free region.
49

CHAPTER 6. RESULTS II: LOCAL ENHANCEMENTS IN THE 6.4 KEV LINE50

1.0
Secondary-only stray
Backside stray
Stray light free region
including quadrant boundaryID
Obs 75010
0
507

0.5
401

07
507

0

0
20
09
50
0
80 10
75
01
0
50
10
60
01
0

69
01

50

70
50

70

01

0
50
70

71
70
50

70
7
50

73
70
50

01

0
20
1

0
01

0
01
50

80

76

01
77
80
50

-0.5

0

0.0
0

Galactic latitude (degree)

GX 3+13

2

1

Galactic longitude (degree)
Figure 6.1. Stray light pattern for each FOV in the observation of the GC east region. Red and green regions
show those contaminated by the stray light of secondary reflection and backside reflection, respectively. Blue
regions show the stray-light free region including the quadrant boundary region. The position of the straylight source GX 3+1 is represented by a circle. The observed images are shown in Appendix.

Since the low mass X-ray binary GX 3+1 does not emit narrow lines (Piraino et al., 2012), thestray light does not aﬀect the line intensity of the GDXE, but aﬀects the continuum. We exclude
the region contaminated by secondary reflection because the flux of the secondary-reflection component is far beyond the intensities of the GDXE. On the other hand, the flux of the backside-reflection
component is about 10% of that of the GDXE. In order to estimate the backside stray, we integrate
X-ray spectra from the region contaminated by the backside reflection and the stray-light free region, and subtract the latter from the former. Here, we performed the correction for the Galactic
coordinate dependence of the GDXE emission on the basis of chapter 5. Figure 6.2 shows the straylight spectrum. For comparison, we simulate the stray light spectrum by xissim using the X-ray
spectral model of GX 3+1 adapted in Piraino et al. (2012), which consists of a blackbody with temperature of 1.7 keV and a multicolor disk blackbody with temperature of 0.93 keV and has the fluxof ∼ 4.5 × 10−9 erg s−1 cm2 in the 2–10 keV band. The simulated spectrum is presented by the red
curve in figure 6.2. The observed spectrum is well fitted with the simulated spectrum with the null
hypothesis probability of 3% (figure 6.2). The absolute observed flux is required to be multiplied by−9 erg s−1 cm2 . This would be due to calibration uncertainty as
1.3 and obtained to be (5.8+1.2−1.1 ) × 10
described in section 3.2.3. Thus we will adopt the obtained flux by the subtraction.

χCounts s−1 keV−1 arcmin−2

6.1. ESTIMATION OF STRAY LIGHT

51

10−5
10−6
10−7
10−8
4
2
0
−2
−4
5

10
Energy (keV)

Figure 6.2. Stray-light spectrum of GX 3+1. The data were obtained by subtracting the stray-light free spectrum from the spectrum contaminated by backside reflection. The red curve shows the simulated spectrumof stray light from GX 3+1 that was estimated by xissim with the model and parameters reported by Piraino
et al. (2012) (see text). The bottom panel shows the diﬀerence between the observed and simulated spectra.

6.1.2 Stray light in the l = −20◦ region
A bright X-ray binary, GX 340+0, which is also a low mass X-ray binary, is 30′ –60′ apart fromtwo FOVs in the −20◦ region, and they are contaminated by stray light from the binary. In theoﬀ-set angle of 30′ –60′ , the secondary reflection is considerably low (see figure 3.7). Therefore weuse the entire FOV, of course excluding point sources. Church et al. (2006) reported that the X-ray
spectrum of GX 340+0 is represented by a cut-oﬀ power law and multiclolor disk blackbody as well
as a broad iron line. In order to estimate the stray light spectrum in a similar way as in the caseof the GC east region, we use quadrant boundary region in a FOV (ObsID=406078010). A nearby
FOV (ObsID=405027010) is also contaminated by stray light from GX 340, but in the same time the
FOV has a quadrant boundary, and thus we used also the region to increase statistics. We perform
simulation by xissim to determine the position of the quadrant boundary regions. The results areshown in figure 6.3.
We integrate spectra extracted from the quadrant boundary regions (the blue regions in figure 5.3) in ObsID=406078010 and 405027010. We then extract the spectrum from the region contaminated by stray light (the green region in figure 5.3) in ObsID=406078010, and subtract the former from the latter. The stray light spectrum is shown in the left-side panel in figure 6.4. In a similarway, the stray light spectrum of ObsID=401052010 is also obtained as is shown in the right-side
panel. The red curves in both the panel are simulated spectra by using xissim, assuming the spectral model of GX 340+0, which consists of a multicolor disk blackbody with temperature of 2.44 keV,a cut-oﬀ power law with photon index of 1.7 and cut-oﬀ energy of 5.8 keV and a broad line at 6.9 keVwith the line width of 0.5 keV, and has the flux of ∼ 9.0 × 10−9 erg s−1 cm−2 in the 2–10 keV band
based on Church et al. (2006). The observed spectra are well fitted with the simulated spectrum with
the null hypothesis probability of 0.2% and 30% for ObsID=406078010 and 401052010, respectively.
The absolute observed fluxes are required to to be multiplied by 2.1 for both the FOVs and obtained

CHAPTER 6. RESULTS II: LOCAL ENHANCEMENTS IN THE 6.4 KEV LINE52

Galactic latitude (degree)

to be (1.9 ± 0.1) × 10−8 erg s−1 cm−2 . We will adopt the obtained flux by the subtraction as the flux ofthe stray light.

0.2

Regions contaminated by stray light
Quadrant boundary
401052010

0.0

ObsID
405027010

GX 340+0406078010
−0.2

341

340
Galactic longitude (degree)

10−4Counts s−1 keV−1 arcmin−2

Counts s−1 keV−1 arcmin−2

Figure 6.3. Stray light pattern for each FOV in the observation of the l = −20◦ region estimated with xissim.
Green and blue regions show the region contaminated by the stray light and the quadrant boundary region,
respectively. The position of the stray light source GX 340+0 is represented by a circle. See Appendix for theobserved images.

10−5

10−6
2
1
0.5
0

10−5

10−6
2
1.5

ratio

ratio

1.5

10−4

1
0.5

5

10
Energy (keV)

0

5

10
Energy (keV)

Figure 6.4. Stray-light spectrum of GX 340+0. (left) Spectra diﬀerential between the observed spectra extracted from the region contaminated by stray light in the observation of ObsID=406078010 and the quadrant
boundary regions in the two observations (ObsID=405027010 and 406078010). The red curve shows the simulated spectrum of stray light from GX 340+0 that was estimated by xissim with the model and parameters
reported by Church et al. (2006) (see text). The bottom panel shows the diﬀerence between the observed
and simulated spectra. (right) Same as the left panel, but in the case of the stray light in the observation of
ObsID=401052010.

6.2. EXCESS IN THE GC EAST

6.2

53

Excess in the GC east

104
10−7
103
10−8
102

Continuum
6.7 keV line
6.4 keV line

10−9
4

3

2

−2
−1
1
Longitude (degree)−3

CO Intensity (K km/s)

10−6

12

X-ray Flux (photons s−1 cm−2 arcmin−2 )

105

−4

Figure 6.5. Intensity profiles of X-rays and 12 CO molecular clouds in −4.◦ 5 < l < 4.◦ 5. The triangles, squares,
and circles show the fluxes of the continuum in the 4–10 keV band (where the 6–7 keV band was excluded),the 6.7 keV line, and the 6.4 keV line, respectively. Errors are quoted at 68% confidence levels. The dashedlines show a model of stellar distribution, which is symmetric with respect to Sgr A∗ (Nishiyama et al., 2013;
Launhardt et al., 2002; Revnivtsev et al., 2006a). The red lines are the 12 CO intensity profile (the unit is
the right-side axis). The blue lines are the sum of the symmetric distribution model and the 12 CO intensity
multiplied by α, where α is a factor to convert the 12 CO intensity to the 6.4 keV flux (see text). The figure istaken from Nobukawa et al. (2015a, figure 1).

Clump 2CMZ

Figure 6.6. Integrated intensity map of 12 CO in −4.◦ 5 < l < 4.◦ 5. Integrated range is −300 km s−1 to +300 km s−1(Torii et al., 2010), where the local components are excluded (see text). White and black squares show the FOVs
of the XIS.

54

CHAPTER 6. RESULTS II: LOCAL ENHANCEMENTS IN THE 6.4 KEV LINE

We show a longitudinal profile of the intensity of the 6.4 keV and 6.7 keV lines and the continuum
in the 4–10 keV band (where the 6–7 keV band was excluded) in the GC east region in figure 6.5.
A popular scenario for the origin of the GDXE is a collection of faint stars (Revnivtsev et al., 2009,2006a; Yuasa et al., 2012; Warwick et al., 2014). The stellar mass distribution has been constructedfrom infrared observations, and has an east-west symmetry (Nishiyama et al., 2013; Launhardt et al.,
2002; Revnivtsev et al., 2006a). The profile is given by the dashed lines in figure 6.5. All the X-ray flux
distributions, except that of the 6.4 keV line in the east, are consistent with the stellar distribution
curve. A remarkable structure is an excess of the 6.4 keV line flux in the east compared to the west,
and hence the excess should have a diﬀerent origin from the symmetrical components which would
be due to point sources.
We plot in figure 6.5 the 12 CO line intensity taken by NANTEN integrated over the velocity
range from −300 km s−1 to +300 km s−1 (Torii et al., 2010). Here, the local components uniformly
distributed around −20 km s−1 to +30 km s−1 and the Near 3 kpc arm having a velocity gradientfrom −80 km s−1 at l = −5◦ to −30 km s−1 at l = 5◦ (Dame and Thaddeus, 2008) are excluded. The
12 CO profile is similar to the excess distribution of the 6.4 keV line in the east, which suggests thatits origin is due to molecular gas. The excess at l = 3◦ coincides with the intersection point of the
Galactic plane and the giant molecular cloud, Bania’s Clump 2, which extends towards the north(Bania, 1977; Tsuru et al., 2014).
Figure 6.7a shows integrated X-ray spectra from the east and west sides. The CXB is subtracted
from both the spectra according to Kushino et al. (2002). The stray light from GX 3+1, which isestimated in section 6.1.1, is also subtracted from the east spectrum. Each spectrum is fitted with anabsorbed power law plus four Gaussian lines at 6.4, 6.7, 7.0, and 7.06 keV. The interstellar absorption
column density is fixed to 6 × 1022 cm−2 . The intensity of 7.06 keV (the neutral iron Kβ) line is fixed
to 0.125 times that of 6.4 keV (the neutral iron Kα) line. The best-fit parameters are summarized intable 6.1.
We subtract the west spectrum from the east one to make the X-ray spectrum of the east excess
as is shown in figure 6.7b. The excess spectrum is fitted with a power law plus a Gaussian line at
6.4 keV. The best-fit line and continuum fluxes are (2.5 ± 0.6) × 10−8 photons s−1 cm−2 arcmin−2 and(1.9±0.3)×10−8 (E/6.4 keV)−3±1 photons s−1 keV−1 cm−2 arcmin−2 , respectively, as shown in table 6.1.
The 6.7 keV line flux has the symmetry between the east and west, but at the same time, there arestatistic errors by ±7%. It means that the 6.7 keV line flux has a possible asymmetry between theeast and west by ±7%. Thus we take this error into account as the uncertainty of the symmetry. Thenthe equivalent width is estimated to be 1.3 ± 0.4+4.2−0.2 keV, where the second and third terms are thestatistical and systematical errors, respectively.

6.2. EXCESS IN THE GC EAST

55

2×10
(a)

5×10

−4

(b)
Counts s−1 keV −1 arcmin −2

Counts s−1 keV −1 arcmin −2

1×10

−5

East
West

−5

−5

1.5×10

10

5×10

−5

−6

0

0
5

6

7

8

9

10

Energy (keV)

5

6

7

8

9

10

Energy (keV)

Figure 6.7. (a) X-ray spectra extracted from the east (red) and west (black) sides. The red and black lines show
the best-fit model consisting of an absorbed power law plus four Gaussian lines at 6.4 keV, 6.7 keV, 7.0 keV,and 7.06 keV. (b) X-ray spectrum of the excess emission in the east. The red line shows the best-fit model
consisting of a power law plus a Gaussian line at 6.4 keV. Errors are quoted at 68% confidence levels for theboth panels. The both figures are taken from Nobukawa et al. (2015a, figure 2).

Table 6.1. Best-fit parameters of the X-ray spectra from the east and west sides (Nobukawa et al., 2015a,
table 2).*

Unit
east
west
excessEmission lines
6.4 keV 6.7 keV 7.0 keV(10−8 )†
(10−8 )†
(10−8 )†
5.1±0.4 10.3±0.5 2.5±0.4
2.5±0.4 10.3±0.5 2.8±0.4
2.5±0.6––

Continuum
photon index normalization(10−8 )‡
2.3±0.1
22.9±0.1
2.2±0.1
21.7±0.1
3±1
1.9±0.3* The error ranges in this table are calculated at 90% confidence levels.
† Absorption-corrected line flux in the units of photons s−1 cm−2 arcmin−2 .
‡ Normalization at 6.4 keV in the units of photons s−1 keV−1 cm−2 arcmin−2 .

CHAPTER 6. RESULTS II: LOCAL ENHANCEMENTS IN THE 6.4 KEV LINE

56

6.3

Excess in the l = −20◦ region

2

400

1

300

401052010

404056010
200

0406078010
−1
−2
−18

cloud 1−20

cloud 2−22

K km/sGalactic latitude (degree)

As in the GC east region, the excess of the intensity of the 6.4 keV line in the l = −20◦ region would
be due to a diﬀerent origin from the 6.7 keV and 7.0 keV lines and the continuum. Interestingly, theNear 3 kpc arm tangent is located at l ∼ −20◦ (see figure 1.1), and thus it is expected that there
is enhancement in molecular gas distribution. Integrated intensity of 12 CO obtained by NANTEN
around l ∼ −20◦ is shown in figure 6.8 (Kazufumi Torii, private communication). The FOVs of ObsID= 406078010, 401052010, and 404056010 are represented by squares, and we find enhancementsin the molecular gas distribution in the three FOV regions. Actually, the FOV regions of ObsID =406078010 and 404056010 have the highest integrated intensity of 12 CO among all of the Suzaku
FOVs in −40◦ < l < −5◦ region that we analyzed. This fact suggests that the 6.4 keV excess in thel = −20◦ region is also due to the molecular gas. We call the enhancement in the molecular gas
distribution around the FOVs of ObsID = 406078010 and 40105201 as cloud 1 and that of the FOVof 404056010 as cloud 2. Figure 6.9 gives a velocity channel distribution of 12 CO in the l = −20◦
region. Cloud 1 consists of molecular gas in the velocity ranges from −140 km s−1 to −120 km s−1
and from −60 km s−1 to −20 km s−1 . The former component would come from the Near 3 kpc arm
while the latter component would belong to molecular clouds located in ∼ 4 kpc distance (Kazufumi
Torii, private communication). Cloud 2, given its corresponding velocity range, would belong to the
Near 3 kpc arm.

100

−24

0

Galactic longitude (degree)
Figure 6.8. Integrated intensity map of 12 CO around the l = −20◦ region. All velocities are integrated. Blacksquares show the FOVs of the XIS.

6.3. EXCESS IN THE L = −20◦ REGION

57

0

100

-160 - -140 km/s

-140 - -120 km/s

-120 - -100 km/s

-100 - -80 km/s

-80 - -60 km/s

-60 - -40 km/s

-40 - -20 km/s

-20 - 0 km/s

0 - 20 km/s

342
340
338
Galactic Longitude [degree]

342

200

K km/s1
0

Galactic Latitude [degree]

-1
-2
1
0
-1
-2
1
0
-1
-2
342

340

338

340

338

Figure 6.9. Channel map of 12 CO around the l = −20◦ region. Black squares show the FOVs of the XIS.

We extract the spectrum associated with the 6.4 keV excess in the same manner as the case of
the GC east. We construct an integrated spectrum from the −20◦ region and an integrated GRXE
spectrum using the data in the 10◦ < |l| < 30◦ and |b| < 0.5◦ region. The CXB is subtracted from boththe spectra according to Kushino et al. (2002). The stray light from GX 340+0, which was estimatedin the previous subsection as the model spectrum, is subtracted from the l = −20◦ -region spectrum.
The Galactic-coordinate dependence is corrected by adjusting the 6.7 keV line intensity as is shownin figure 6.10a. We, then, extract the spectrum of the excess emission by subtraction of the GRXE
spectrum.
The excess spectrum is fitted with a power law plus a Gaussian line at 6.4 keV (figure 6.10b).
The best-fit line and continuum fluxes are (3.9 ± 1.1) × 10−8 photons s−1 cm−2 arcmin−2 and (5.7 ±1.9) × 10−8 (E/6.4 keV)−1.2 photons s−1 keV−1 cm−2 arcmin−2 , respectively. Here, there is a statisticaluncertainty of ±21% in the adjustment of the 6.7 keV line intensity: there are the statistical errors by±21% in the ratio of 6.7 keV line intensity between the integrated spectrum from the −20◦ region andthe integrated GRXE spectrum. This makes a systematical uncertainty in the subtraction of spectra.
Then the equivalent width is estimated to be 0.67 ± 0.22+1.6−0.1 keV, where the second and third terms
are the statistical and systematical uncertainties, respectively.

CHAPTER 6. RESULTS II: LOCAL ENHANCEMENTS IN THE 6.4 KEV LINE

58

(a)

(b)
−4

−4

10

Counts s−1 keV −1 arcmin −2

Counts s−1 keV −1 arcmin −2

10

−5

10

−6

10

−5

10

−6

5

6

7

8

9

10

Energy (keV)

10

5

6

7

8

9

10

Energy (keV)

Figure 6.10. (a) X-ray spectrum extracted from l = 20◦ (red) and the integrated spectrum extracted from theother regions on the Galactic plane (10◦ < |l| < 30◦ ; black). The red and black lines show the best-fit model
consisting of an absorbed powerlaw plus four Gaussian lines at 6.4 keV, 6.7 keV, 7.0 keV, and 7.06 keV. (b)X-ray spectrum of the excess emission in the l = −20◦ region. The red line shows the best-fit model consistingof a power law plus a Gaussian line at 6.4 keV. Errors are quoted at 68% confidence levels for both the panels.

Table 6.2. Best-fit parameters of the X-ray spectra from the l = −20◦ regions and the other region on the
Galactic plane (10◦ < |l| < 30◦ ) before correcting the Galactic-coordinate dependence by adjusting the 6.7 keVline intensity.*Unitl = −20◦◦
10 < |l| < 30◦
excessEmission lines
6.4 keV 6.7 keV 7.0 keV(10−8 )† (10−8 )† (10−8 )†6.3±1.0 8.9±1.0< 1.4
1.6±0.2 5.8±0.3 0.9±0.2
3.9±1.1––

Continuum
photon index normalization(10−8 )‡
1.6±0.127±1
2.5±0.1
12.6±0.1< 1.2
5.7±1.9* The error ranges in this table are calculated at 90% confidence levels.
† Absorption-corrected line flux in the units of photons s−1 cm−2 arcmin−2 .
‡ Normalization at 6.4 keV in the units of photons s−1 keV−1 cm−2 arcmin−2 .

Chapter 7

Discussion
In chapter 5 and 6, we have separately measured the scale heights and equivalent widths of the
6.4 keV, 6.7 keV, and 7.0 keV lines in the GCXE, GBXE, GRXE. Furthermore, we have found that there
are enhancements of the 6.4 keV line relative to the 6.7 keV distribution in the GC east and l = −20◦
regions and that the enhancements are associated with molecular gas. In this chapter, we will first
discuss the origin of the 6.4 keV enhancements based on the result of spectral analysis. Then, we
will discuss the origin of the GDXE on the basis of the observed scale heights and equivalent widths.

7.1 Origin of the 6.4 keV enhancements
Recently, a common candidate of the origin of 6.4 keV line emission in the GDXE has been considered to be CVs (Yuasa et al., 2012). However, the excess emissions discovered in this dissertation
have only the 6.4 keV line while the CVs usually exhibit the 6.4 keV, 6.7 keV, and 6.9 keV lines(Hellier et al., 1998; Ezuka and Ishida, 1999; Hellier and Mukai, 2004). The spatial distribution of
the 6.4 keV line intensity strongly indicates that the excess emissions are emitted from neutral ironatoms in molecular gas. Since temperature of the molecular gas in interstellar space is at most 100 K,
X-rays cannot be generated. Some particles with higher energy than the ionization energy of aniron K-shell electron, such as cosmic-ray protons, electrons, or X-rays, should be bombarded on themolecular gas.
The equivalent width of a 6.4 keV line depends on the irradiating particles and their spectral
index as shown in figure 7.1. In the case of cosmic-ray bombardment, the 6.4 keV line is produced via
inner-shell ionization by protons (Dogiel et al., 2011) or by electrons (Valinia et al., 2000; Yusef-Zadeh
et al., 2002), while the continuum is due to inverse bremsstrahlung (for protons) or bremsstrahlung(for electrons). Figure 7.2 shows the cross sections for producing the 6.4 keV line and bremsstrahlung
at 6.4 keV by the impact of electrons and protons. The energy range of the protons that produce the
6.4 keV line is mainly the MeV band while that of the electrons is mainly the keV band. The photon
index of the observed X-ray spectrum is not the same as the particle index. We calculate the photon
index in the case of a certain particle index assuming optically-thin line emission based on Baring
et al. (2000). The result is shown in figure 7.3.
As for the X-ray irradiation, the 6.4 keV line is produced via photoionization by X-rays withenergy higher than 7.1 keV (K-edge) while the continuum is due to Thomson scattering. The photon
index of the observed X-ray spectrum is the same as that of the irradiating X-rays in the energy band
59

CHAPTER 7. DISCUSSION

Equivalent width of 6.4 keV line (keV)

60
10 2

ProtonElectron

X-ray

10

the GC eastthe GC east1

l = −20°0.1

0l = −20°
1

2

3

particle index

4

0

1

2

3

4

photon index

Figure 7.1. (left) Equivalent width of the 6.4 keV line as a function of particle index (Dogiel et al., 2011) with
the iron abundance of 1–1.5 solar (see text). The error regions of the particle index and equivalent width for
the GC east and the l = 20◦ regions are shown by the hatched region. The red and blue belts are the calculated
values for protons and electrons, respectively. The thick red and blue regions are the acceptable ranges forthe proton and electron scenarios, respectively. (right) Same as the left, but of photon index (Tsujimoto et al.,2007). The green belt and the thick green region are the calculated values and the acceptable range for the
X-ray scenario, respectively.

of 2–10 keV, which we treat in this dissertation.
In both the scenarios, the equivalent width depends on the iron abundance (Tsujimoto et al., 2007;
Dogiel et al., 2011). The metal abundances were determined in the GC region by X-ray observationsof the HP to be ∼ 1.9 solar for sulfur, argon, calcium but ∼1.2 solar for iron (Nobukawa et al., 2010;
Uchiyama et al., 2013). Muno et al. (2004) reported a similar result. The iron abundance of 1–1.5
solar was also measured in the XRNe in Sgr B with the absorption of the iron K-edge (Nobukawa
et al., 2011). Giveon et al. (2002) obtained 2 solar for heavy elements of neon, sulfur, and argon(not of iron), and Cunha et al. (2007) and Martin et al. (2015) measured the iron abundance to be
1–1.5 solar. Therefore, we adopt the iron abundance of 1–1.5 solar as is shown in figure 7.1 withthe colored belts. The equivalent width in the X-ray scenario also depends on the angle θ between
the line of sight and the incident photon direction (Tsujimoto et al., 2007). Figure 7.1 shows theresult of the reflection angle θ = 90◦ , which gives the maximum equivalent width. The shapes of theequivalent width curves roughly reflect the ratio of cross sections for producing the 6.4 keV line andbremsstrahlung (figure 7.2). The hatched regions in figure 7.1 indicates the equivalent width rangesof the statistical and systematic errors in the GC east and l = −20◦ regions.
All the scenarios are accepted for the l = −20◦ region. On the other hand, in the GC east region,
the allowed parameter region completely excludes the electron origin. Then, we discuss the others,
the X-ray and proton scenarios for the GC east in the next subsection.

61

-1

σ (b H-atom or b H-atom keV )7.1. ORIGIN OF THE 6.4 KEV ENHANCEMENTS-1

dσB(E =6.4 keV)
dEX X
10

Electrons

ProtonsdσB(E =6.4 keV)
dEX X

σ6.4 keV-2-1

σ6.4 keV

a)
10

b)

-3

1

10

10

2

10

3

10

4

10

5

6

10 1

Ee (keV)

10

10

2

10

3

10

4

Ep (MeV)Figure 7.2. Cross section for producing the 6.4 keV line by the impact of fast electrons and protons. Solid lines
show the cross section (in units of barn per ambient H-atom) for producing the 6.4 keV line by the impact offast protons or electrons with the abundance of 1 solar. Dashed lines show the diﬀerential cross section (inbarn per H-atom per keV) for producing 6.4 keV X-rays by bremsstrahlung (electrons; left panel) and inverse
bremsstrahlung (protons; right panel) in a medium composed of hydrogen and helium with the density ratioof 1 : 0.1. Both the panels are taken from Tatischeﬀ et al. (2012).

7.1.1 The GC east region
X-ray irradiation
In the X-ray irradiation scenario, the 6.4 keV line flux depends on the irradiating source flux,
the hydrogen column density (NH ) of the molecular gas (or cloud), and the distance (D) between theirradiating source and the target. The hydrogen column density NH =(2–6)×1022 cm−2 is obtainedby multiplying the 12 CO line intensity in figure 6.5 by the conversion factor (called X-factor) of
0.7 × 1020 cm−2 (K · km s−1 )−1 (Torii et al., 2010), which is measured near the GC.
A possible irradiating source is the supermassive black hole Sgr A∗ , as for the XRN model invokedfor the GC region (Sunyaev et al., 1993; Koyama et al., 1996; Ponti et al., 2010). In this case, the source
luminosity should be LX ∼ 1040 (D/450 pc)2 erg s−1 about 1500×(D/450 pc) years ago, where D is the
distance of the 6.4 keV line emitting region from Sgr A∗ . For the XRN in Sgr B, the luminosity of Sgr
A∗ was ∼ 1039 erg s−1 about 300 years ago (Murakami et al., 2000; Revnivtsev et al., 2004). Brighter
flares of ∼ 1041−43 erg s−1 in the far past (105 –107 years ago) have been proposed to explain the Fermi
bubbles and the recombining plasma in the south region of the GC (Su et al., 2010; Nakashima et al.,2013). Thus, many big flares in the interval period between these two epochs (∼700–1500 years ago)
may be conceivable. However, this scenario has two diﬃculties.
The first diﬃculty comes from the longitudinal distribution of the excess flux of the 6.4 keV line.
The observed excess flux is proportional to the 12 CO intensity, which means that the X-ray intensityfrom Sgr A∗ is almost constant, and does not decrease as the square of the Galactic longitude, or the
distance from Sgr A∗ ; the past Sgr A∗ flares should have smoothly increasing flux with the square of
the look-back time. A monotonous smooth decrease of flux should continue during 700–1500 years
ago, which is artificial, although not completely rejected.
The second diﬃculty comes from the flux ratio of the 6.4 keV line to the 12 CO intensity (here, α)CHAPTER 7. DISCUSSION

62
5Photon index

4

3

2

1
0

1

2
Particle index

3

4

Figure 7.3. Calculated photon index of X-ray radiated by cosmic-ray particles as a function of the particle
index based on Baring et al. (2000).

in the molecular cloud Clump 2 (see figure 6.6). Clump 2 extends toward the north from the Galactic
plane and has an elliptical shape with major and minor axes of ∼ 1.◦ 1 and ∼ 0.◦ 5, respectively. We estimate the excess 6.4 keV line flux from the on-plane (b ∼ −0.◦ 04; Obs. ID=507073010) and oﬀ-plane(b ∼ 0.◦ 15 − −0.◦ 40; Obs.ID=507074010 and 507075010) parts of Clump 2 by subtracting the symmetrical component from the observed flux, and obtained α of (5.9 ± 0.7) × 10−11 and (2.1 ± 0.6) × 10−11
for the on- and oﬀ-plane parts, respectively (unit is photons s−1 cm−2 arcmin−2 (K · km s−1 )−1 ). Since
the flux from Sgr A∗ would be equal in the small separation angle between on- and oﬀ-plane parts,
α should be also equal, which is in conflict with the observed values. From these two diﬃculties, we
regard that the X-ray irradiation by Sgr A∗ is unlikely.
The other possibility is that irradiation of many X-ray binaries in the east is responsible for the
6.4 keV line. This idea is essentially the same as Molaro et al. (2014). Assuming that a mean spectrumof the relevant X-ray binaries is a power law with the photon index of Γ = 3, and is surrounded by coldmaterial with NH = 2–6 × 1022 cm−2 , a fraction of 0.03–0.1% of the total 2–10 keV flux is converted
to 6.4 keV X-rays. We estimate the excess 6.4 keV luminosity in total to be 3 × 1033 erg s−1 in therelevant area of 2.◦ 0 × 0.◦ 2. Then, a total luminosity of 3 × 1036 –1037 erg s−1 is required to the sources
inside the area. This flux exceeds the total luminosity of the GRXE in this region (∼ 3 × 1035 erg s−1) by one or two orders of magnitude. No source brighter than 1035 erg s−1 is found in this area.
Therefore, this scenario requires a source with more than one or two orders of magnitude brighter
flares within the relevant regions and duration of the past ∼ 1000 years. In addition, the reflection
angle θ should be distributed randomly. Then the equivalent width integrated over the reflection
angle θ is ∼ 1.3 times lower than the green belt in figure 7.1 (reflection angle θ = 90◦ ), and hence
is out of the allowed range. Thus this scenario is also unlikely. We regard the proton-bombardment
scenario as the most probable origin of the excess in the GC east region.

7.1. ORIGIN OF THE 6.4 KEV ENHANCEMENTS

63

Cosmic-ray Bombardment
In the proton model, the excess flux of the 6.4 keV line I6.4 keV is given by
1
I6.4 keV =
N
4π H

(

∫
σ6.4 keV υ A

Ep
1 MeV

)−p
dEp ,

(7.1)

where σ6.4 keV , υ, A, Ep and NH are the cross section to produce the 6.4 keV line by protons, thevelocity, the number density at 1 MeV and the energy of protons, and the line-of-sight hydrogencolumn density, respectively. The spectral index p of 2.5 is adopted from the best-fit photon index
of 3 (see figre 7.3). Using the X-factor of 0.7 × 1020 cm−2 (K · km s−1 )−1 (Torii et al., 2010), the 6.4
keV line flux is expressed as the 12 CO line intensity multiplied by the conversion factor α, which ismeasured to be 4.2 × 10−11 photons s−1 cm−2 arcmin−2 (K · km s−1 )−1 in figure 6.5. The cross section
σ6.4 keV has a peak at 10 MeV and rapidly decreases below 1 MeV and above 50 MeV (Paul andSacher, 1989). We set the integration range to be 1–50 MeV, and then obtained the normalization
A = 1.4 × 10−5 protons cm−3 and the energy density of 20 eV cm−3 . When the integration range is
expanded to be 0.1–1000 MeV, the energy density becomes 80 eV cm−3 . This is about one or twoorders of magnitude higher than the canonical value ∼ 1 eV cm−3 that is determined by observing
high-energy cosmic rays (Neronov et al., 2012).
Tatischeﬀ et al. (2012) calculated an energy conversion rate from protons to the 6.4 keV line to
be 10−6 or less. Since the total luminosity of the 6.4 keV line emission is 3 × 1033 erg s−1 , the proton
power of > 3 × 1039 erg s−1 is estimated. This is not far from the energy ∼ 2 × 1039 erg s−1 that protons
input to the GC (Dogiel et al., 2013).

7.1.2 The l = −20◦ region
X-ray irradiation
Figure 7.1 does not exclude any scenario in the l = −20◦ region. We first discuss the X-ray origin
in the following. The hydrogen column density NH =(5–6)×1022 cm−2 is obtained by multiplying the
12 CO line intensity in figure 6.8 by the conversion factor (X-factor) of 2 × 1020 cm−2 (K · km s−1 )−1(Dame et al., 2001).
A possible irradiating source is a nearby bright X-ray binary (Molaro et al., 2014). The source
luminosity should be LX ∼ 4 × 1038 (D/100 pc)2 erg s−1 , where D is the distance of the 6.4 keV line
emitting region from the irradiating source. This luminosity is comparable to or higher than the
Eddington luminosity of a neutron star LEddington = 2 × 1038 (D/100 pc)2 erg s−1 . There are two
possible sources which located at the projected distance of 70–150 pc: the LMXB GX 340+0 and theblackhole candidate 4U 1630−472. However, both the sources are located at 10–11 kpc from thesun (Grimm et al., 2002; Seifina et al., 2014)1 while the molecular clouds which associated with the
6.4 keV excess located at 4–8 kpc. There is no possible irradiating source in the vicinity. In the GC
region, the supermassive black hole, Sgr A* was active and irradiated the surrounding molecular
clouds several hundred years ago (e.g. Koyama et al., 1996). However, there is no such a special
source in this area.
1 Since no optical counter part has not been identified for 4U 1630−472, its distance from the Sun has an uncertainty.

Seifina et al. (2014) estimate that the distance is large (⩾ 10 kpc) because the source is heavily absorbed in soft X-rays.

CHAPTER 7. DISCUSSION

64

In the same manner as the case of the GC east, we also discuss contribution of many faint point
sources. This scenario was rejected for the GC east. The hydrogen column density NH = 5–6 ×1022 cm−2 and the surface intensity of the excess 6.4 keV line is ∼ 4×10−8 photons s cm−2 arcmin−2 arecomparable to those of the GC east. Thus, the luminosity of the many point sources LX > 1036 erg s−1
is required assuming the same emitting region as the GC east. However, the total luminosity of the
GDXE in this region is ∼ 2 × 1035 erg s−1 , and hence this scenario is rejected. Therefore, the X-ray
scenario would be unlikely.
Cosmic-ray bombardment
Also in the l = −20◦ region, the most plausible scenario is the cosmic-ray bombardment. Since
whether the irradiating particles are protons or electrons cannot be constrained, we estimate theenergy density for both the case. Using equation (7.1), we obtained the energy density of 20 eV cm−3in the energy range of 0.1–1000 MeV and 0.03 eV cm−3 in the energy range of 0.1–1000 keV for
protons and electrons, respectively. Although the LECR electron origin cannot be ruled out, theproton origin is acceptable as is adopted to the case of the GC east region.
The total observed luminosity of the excess 6.4 keV line emission is at least 1 × 1033 erg s−1 . Here,
we assume the excess mainly comes from molecular clouds located at the distance of 5 kpc. Since
the eﬃciency for the production of the 6.4 keV line is 10−6 or less in both of the proton and electron
cases (Tatischeﬀ et al., 2012), the kinematic power of > 1 × 1039 erg s−1 is estimated.

7.2

Scale heights and equivalent widths of the GDXE

7.2.1

Comparison with previous study

Scale height
The Tenma satellite, which discovered the intense iron line, found that the scale heights of thecontinuum (2–10 keV) and the iron lines are 100–300 pc assuming the emission comes from a disk
with a radius of 10 kpc (Koyama et al., 1986b). Yamauchi and Koyama (1993) determined the scaleheight of the 6.7 keV line to be 100 pc by the Ginga survey of the Galactic ridge. Valinia and Marshall (1998) found that the continuum emission in the 5–8 keV band from the |l| < 15◦ region has
two exponential components and the scale heights of the thin and broad ones are 70 pc and 500 pc,
respectively. Revnivtsev and Molkov (2012) made the surface brightness profile of the GRXE continuum in the 4.3–10.5 keV band across the Galactic plane from b = 0◦ to b = −30◦ at l = 18.◦ 5 by using
RXTE, and obtained the scale height of 107 ± 5 pc. Uchiyama et al. (2013) measured the latitudinal
distributions of the 6.4 keV and 6.7 keV lines separately as well as the 5–8 keV continuum in the
GCXE and GRXE by using Suzaku. They reported that the e-folding scale heights of the 6.4 keVin the GCXE and GRXE are ∼30 pc and ∼150 pc, respectively, while those of the 6.7 keV and thecontinuum in the GCXE and the GRXE are ∼40 pc and ∼400 pc, respectively.
The measured scale heights of the 6.4 keV, 6.7 keV, and 7.0 keV lines and the hard X-ray continuum of the GCXE in this work are almost consistent with the results in Uchiyama et al. (2013), butthere is discrepancy for the GRXE. Our result of the scale height of the 6.7 keV and 7.0 keV lines andthe hard continuum in the GRXE of ∼140 pc is almost consistent with the results obtained by Tenma,

7.2. SCALE HEIGHTS AND EQUIVALENT WIDTHS OF THE GDXE

65

Ginga, and RXTE. The scale heights of the GRXE measured by Uchiyama et al. (2013) are systematically large but are closer to those in the GBXE in our work. Actually, a majority of the data that
Uchiyama et al. (2013) used for the GRXE are extracted from the 2◦ < |l| < 4◦ region, which mainly
consists of the GCXE and GBXE based on our result.
Revnivtsev et al. (2009) performed a Chandra deep observation (∼1 Msec) at (l, b) = (0.◦ 1, −1.◦ 4)(CBF). It is true that the CBF is near the GC. However, figure 5.2 and figure 5.3 show that contributionof the GCXE is ≤ 10% of the GBXE at b∗ = 1.◦ 4 in the 0.◦ 0 < l < 385.◦ 5 region, and thus the CBF isdominated by the GBXE.
Equivalent width
In section 4.3, we obtained the equivalent widths of the iron lines as is listed in table 7.1. Uchiyama
et al. (2013) show the mean equivalent widths of the 6.4 keV and 6.7 keV lines in the GRXE ∼ 110 eVand ∼ 490 eV, respectively. Given the mixture of the GCXE and GBXE in the data used in Uchiyama
et al. (2013) as mentioned before, our result would be consistent with them.

7.2.2

Comparison with point sources

A common idea for the origin of the GDXE is a superposition of many point sources such as CVs
and ABs (e.g. Revnivtsev et al., 2009). Previous studies focused on an X-ray luminosity function or
spectra of point sources, and thus verified whether the integrated surface brightness and spectrum
match the observed GDXE. In chapter 5, we separately measured the scale heights and equivalent
widths of the 6.4 keV, 6.7 keV, and 7.0 keV lines and discovered that the values are diﬀerent among
the GCXE, GBXE, and GRXE. We below examine the point-source scenario by comparing the scaleheights and equivalent widths of CVs and ABs with those of the GCXE, GBXE, and GRXE.
Scale heights of CVs in the solar neighbourhood (distance ≲ 1 kpc) are measured by several
authors (Patterson, 1984; Ak et al., 2008; Revnivtsev et al., 2008). Most ABs composed of G or K stars
and a small fraction (∼15%) are F stars (Strassmeier et al., 1993). We adopt scale heights of G and K
stars as that of ABs. We show the scale heights of CVs and ABs in table 7.2.
Molecular gas distribution may provide a hint of the origin of the GDXE. We measure latitudinal
distributions of the molecular gas using NANTEN data in the GC region (|l| < 1◦ ) and the Galactic
plane (2◦ < |l| < 40◦ ). Latitude distributions of the two regions are shown in figure 7.4. By fitting theplots in the range of |b| < 2◦ with an exponential function, we obtained the scale heights of the 12 CO
as shown in table 7.2.
The typical equivalent widths of CVs are listed in table 7.1. CVs are divided into two categories
by their magnetic field strengths, namely magnetic CVs and non-magnetic CVs. The mass accretionin magnetic CVs and non-magnetic CVs are diﬀerent; typical equivalent widths of iron K-shell lines
from the two categories are also diﬀerent. In magnetic CVs, namely polars and intermediate polars,
magnetic filed of a white dwarf is so strong that it controls the accretion flow from a companion
star. The accretion matter forms an accretion column, which radiate mainly in X-rays. Magnetic CVs
typically have luminosities of ∼ 1032 –1034 erg s−1 . The equivalent widths of the three iron K-shelllines of magnetic CVs are well studied by many authors (e.g. Hellier et al., 1998; Ezuka and Ishida,
1999; Hellier and Mukai, 2004). Although there is dispersion of ∼ 30 eV between authors, typicalequivalent widths are ∼ 130 eV, ∼ 160 eV, and ∼ 90 eV for the 6.4 keV, 6.7 keV, and 7.0 keV lines,CHAPTER 7. DISCUSSION

66
104

103

(b) 2° < | l | < 40°

Integrated intensity (K km s -1)(a) | l | < 1°
10

3

102

102
101
10

1

100

100
10-1

-4

-2
0
2
Galactic latitude (degree)

4

10-1

-4

-2
0
2
Galactic latitude (degree)4

Figure 7.4. (a) Latitudinal distribution of integrated in intensity of 12 CO in the CMZ (|l| < 1.◦ 0). Integrated
range is the same as figure 6.5 and 6.6. The dashed line shows the best-fit model consisting of an exponentialfunction, which is symmetric with respect to the Galactic plane (b = −0.◦ 046). (b) same as (a), but in the
Galactic plane (2◦ < |l| < 40◦ ). All velocities are integrated. The data are provided by Kazufumi Torii (private
communication).

respectively. On the other hand, 90% of CVs fall into non-magnetic CVs. Most common type of
non-magnetic CVs are dwarf novae with luminosities of ∼ 1030 –1032 erg s−1 (e.g. Patterson, 1984).
Dwarf novae are considered to change their luminosities via instabilities in accretion disk (Osaki,1974). Dwarf novae show the 6.7 keV line whereas they do not so strongly emit the 6.4 keV line.
ABs, which have luminosities of 1027 –1031 erg s−1 , emit a strong 6.7 keV line, but they show little
emission for the other two lines. Although X-ray spectra from thermal plasma in non-magnetic
CVs and ABs are well studied, there is few study that gives direct information about the equivalent
widths for the three iron K-shell lines. We use the averaged equivalent widths obtained by Suzaku
observations of 9 dwarf novae and 9 ABs (Nobukawa et al. in prep.).
Sazonov et al. (2006) estimate contribution of ABs and CVs to the GDXE of (2 − 3) : 1 in theluminosity range of 1028 –1031 erg s−1 , based on an X-ray luminosity function of nearby X-ray sources.
In a similar way, Warwick (2014) also claimed that the bulk of the GRXE signal (80% in the restricted
6–10 keV band) can be attributed to the X-ray emission of these two populations, with a rough 3 : 1
split between the ABs and CVs contributions. Since ABs and CVs represent old stellar populations,
their X-ray luminosity function normalized to the stellar mass is not expected to vary significantly
across the Galaxy (Sazonov et al., 2006). We thus assume that ABs and CVs contribute to the ironK-shell iron with a 3 : 1 split. We also assume that magnetic CVs and non-magnetic CVs accountfor 90% and 10% in the total CVs, respectively, because the number density of non-magnetic CVs is
considered to be 10 times larger than the magnetic CVs but the luminosity of non-magnetic CVs is
100 times less than the magnetic CVs. Then superposition of these point sources gives the equivalent
widths of 50 eV, 300 eV, and 70 eV for the 6.4 keV, 6.7 keV, and 7.0 keV lines, respectively. We
compare these values (hereafter 3ABs+CVs) with equivalent widths in the GRXE, GCXE and GBXE.

7.2. SCALE HEIGHTS AND EQUIVALENT WIDTHS OF THE GDXE

67

Table 7.1. Equivalent widths of CVs and ABs.

Objectmagnetic CVs*
non-magnetic CVs†
ABs†
3ABs+CVs
GCXE
GBXE
GRXE

Equivalent width (eV)
6.4 keV 6.7 keV 7.0 keV
∼ 130
∼ 160
∼ 90
∼ 70
∼ 500
∼ 130
∼ 20
∼ 340
∼ 60
∼ 50
∼ 300
∼ 70
145 ± 3 527 ± 4
221 ± 361 ± 11 443 ± 14 160 ± 14
97 ± 12 428 ± 15 117 ± 19

* Hellier et al. (1998); Ezuka and Ishida (1999); Hellier and Mukai (2004).
† Nobukawa et al. in prep.

7.2.3 GRXE
Diﬀerence between 6.4 keV and 6.7 keV lines
We found that the scale height of the 6.4 keV line in the GRXE is systematically smaller thanthose of the 6.7 keV and continuum X-rays. This indicates that the 6.4 keV and 6.7 keV lines in the
GRXE have diﬀerent origins. The 6.7 keV line in the GRXE of ∼ 140 pc is consistent with the point
sources (ABs and CVs). The equivalent width of the 6.7 keV line (428 ± 15 eV) cannot be explainedby superposition of 3ABs+CVs (∼300 eV) as estimated in section 6.2.2 (table 7.1). Dwarf novae withthe equivalent width ∼ 500 eV might contribute to the GRXE more significantly.
On the other hand, any mixing ratio of the point sources cannot explain the equivalent widths of
the 6.4 keV line. In the case of the combination of 3AB+CVs, the 6.4 keV line equivalent width is
only 50 eV, which is 50% of the 6.4 keV line intensity. Another origin is required to account for atleast the half of the 6.4 keV line emission in the GRXE. The scale height of the 6.4 keV line is similar
to that of the molecular gas on the Galactic plane. As already mentioned, if the molecular gas isirradiated by X-rays or cosmic-ray particles, prominent 6.4 keV line emission is generated. In fact,
as investigated in chapter 6, the excesses of the 6.4 keV line emission in the GC east and l = −20◦
regions are associated with the molecular gas.
Origin of the 6.4 keV line
As with discussion about the origin of the 6.4 keV excesses, we discuss the X-ray-irradiating
scenario in the global scale. This scenario is in principle the same as Molaro et al. (2014), who claimed
that scattered photons from the ISM originating LMXBs are able to contribute, on average, 10%–30%of the GRXE. In the 10◦ < |l| < 30◦ region, averaged 12 CO integrated intensity is ∼ 150 K (km/s)−1 ,which is translates into NH ∼ 3 × 1022 cm−2 using the X-factor (Dame et al., 2001). X-ray observationsin this region also give NH of 3–4×1022 cm−2 (e.g. Yasumi et al., 2014; Takata et al., 2015; Nobukawa
et al., 2015b).
Here we caution that an X-ray source illuminate only its vicinity and the source cannot illuminate
all the cold matter along the (∼ 10 kpc) line of sight because the X-ray intensity decreases with

CHAPTER 7. DISCUSSION

68

Table 7.2. Scale heights of CVs, ABs, and molecular gas (12 CO).

ObjectCV*AB†
molecular gas
GCXE
GBXE
GRXE

(|l| < 1◦ )(2◦ < |l| < 40◦ )6.4 keV
6.7 keV
6.4 keV
6.7 keV
6.4 keV
6.7 keV

Scale height (pc)130–160
150–300
28 ± 192 ± 131 ± 336 ± 3161 ± 50314 ± 95
70 ± 17142 ± 17

* Obtained from the vicinity of the solar system (< 1 kpc)
(Patterson, 1984; Ak et al., 2008; Revnivtsev et al., 2008).
† Obtained from the vicinity of the solar system (< 1 kpc)
(Gilmore & Zeilik 2000).

the squared distance between the source and the target. The scale height of the cold matter, i.e.
molecular gas, is ∼ 100 pc, and the line-of-sight length is assumed to be the same as the Galactic
diameter, ∼ 10 kpc. If the X-ray source with the same spectral shape as GX 3+1 (Piraino et al.,2012) is surrounded by molecular gas with the above size, 0.005% of the total 2–10 keV flux of theirradiating source is converted to the 6.4 keV line emission. We estimate the 6.4 keV luminosity intotal to be 3 × 1035 erg s−1 in the relevant area of 2 × (20◦ × 1.◦ 0) using the parameters in table 5.1.
Then, a total luminosity of 6 × 1039 erg s−1 is required to the sources inside the area. Since the totalluminosity of LMXBs in the same region is ∼ 6 × 1038 erg s−1 (Grimm et al., 2002; Molaro et al., 2014),
the 6.4 keV line due to the X-ray irradiation contributes only 10% of the 6.4 keV line emission in the
GRXE.
The most possible origin is still cosmic-ray bombardment as with the GC east and the l = −20◦
regions. We already estimated that at least ∼ 50% of the 6.4 keV line emission should have another origin than the point sources. If this is of the cosmic-ray bombardment, the energy densities would be ∼ 10 eV cm−3 and ∼ 0.01 eV cm−3 for the protons and electrons cases, respectively.
Combined with the energy density obtained in the GC east region (80 eV cm−3 ) and the l = −20◦
region (20 eV cm−3 ), the energy density of LECR protons on the Galactic plane would commonly behigher than the canonical value ∼ 1 eV cm−3 that is determined by observing high-energy cosmic
rays (Neronov et al., 2012). The total luminosity of the 6.4 keV line emission requires the energy
injection to cosmic rays of > 3 × 1041 erg s−1 .
Among the Galactic plane, the regions where molecular gas is rich would show the 6.4 keV line
enhancements. The GC east region contains the giant molecular cloud Clump 2. The l = −20◦ region
is located at the line of sight tangent to the Near 3 kpc arm; the molecular gas column density
increases at a tangent to a spiral arm. There exists other tangent regions in the Suzaku observation
range: the tangents to the Far 3 kpc arm at l ∼ 25◦ and the Scutum-Centaurus arm at l ∼ −50◦ .
Also these regions exhibit the indication of the 6.4 keV line enhancement (figure 5.5e) although not7.3. WHERE LOW-ENERGY COSMIC RAYS COME FROM?

69

statistically significant. The 6.4 keV line enhancements should be also due to the higher density of
LECRs than other regions on the Galactic plane.

7.2.4

GCXE

The scale heights of the 6.4 keV and 6.7 keV lines in the GCXE of ∼ 30 pc and ∼ 40 pc aremuch smaller than the CVs and ABs although there is an uncertainty whether the scale heights ofpoint sources obtained in solar neighbourhood can be adopted even in the peculiar environment of
the GC. Uchiyama et al. (2011) measured the scale height of the 6.7 keV in the GC, and found thatit shows 3.8 times larger excess than the stellar distribution model assuming all of the GBXE and
GRXE are due to point sources. A similar result was obtained by comparing a stellar density map
constructed from an infrared observation with the latitudinal profiles of the 6.7 keV line (Nishiyama
et al., 2013; Yasui et al., 2015). The smaller scale heights of the 6.4 keV and 6.7 keV lines than thestellar distribution require contribution of other origins, which would be diﬀuse emission in the
GCXE. Interestingly, the scale heights of the 6.4 keV and 6.7 keV lines in the GCXE are consistent
with that of the molecular gas in the GC. It is possible that the major fraction of the 6.4 keV linein the GC come from the molecular gas origin and the 6.7 keV line may also originate from sources
correlated with the molecular clouds.
The equivalent widths of the three lines in the GCXE are larger than those in the GBXE and
GRXE. Any mixing ratio of ABs and CVs cannot explain either iron line. The assumption that the ironabundance in the GC is 2–3 times larger than the solar one may be explained the observed equivalent
widths. However, observational results are against the assumption. As discussed in section 7.1, areasonable iron abundance in the GC should be 1–1.5 solar, and thus the large equivalent widths in
the GCXE also require other origins than the superposition of point sources.

7.2.5

GBXE

The scale heights of the 6.4 keV and 6.7 keV lines in the GBXE of ∼ 160 pc and ∼ 320 pc are
consistent with those of the CVs and ABs although those values are obtained from data in the solar
neighbourhood (section 7.2.2). The equivalent width of the 6.4 keV and 6.7 keV lines in the GBXE
is marginally consistent with those of 3ABs+CVs while the equivalent widths of the 7.0 keV line in
the GBXE are 2 times higher. The equivalent widths of the three lines are almost explained whendwarf novae more contribute to the GBXE than the assumption. Combination of ABs and CVs would
significantly contribute to the GBXE. Our result is consistent with Revnivtsev et al. (2009), whoreported that more than 80% of the X-ray emission in the GBXE is resolved into point sources.

7.3

Where low-energy cosmic rays come from?

We demonstrated that significant fraction of the 6.4 keV line emission in the GRXE is generated
by the LECR protons (∼MeV) or electrons (∼keV). LECRs are aﬀected by solar modulation, and hence
their intensity and spectral shape in the interstellar space cannot be measured inside the solar system. In fact, while characteristics of GeV and TeV cosmic rays have gradually been revealed thanks
to GeV and TeV gamma-ray observations, the LECRs have been a totally unknown component. Only

70

CHAPTER 7. DISCUSSION

Voyager-I, which has got away from the heliosphere and reached 122 astronomical units in 2012,
measured LECRs in the interstellar space (Stone et al., 2013).
The energy density of the LECR protons measured by Voyager-I is ∼1 eV cm−3 although the LECR
intensity immediately outside the heliosphere may be lower than the Galactic intensity because ofmodulation in the local interstellar medium (Stone et al., 2013). Our measurement using the 6.4 keVline has been performed on the Galactic intensity for the first time. The energy density of LECR
protons measured in the GC east, l = −20◦ region and the entire GRXE ranges from ∼ 10–80 eV cm−3in the 0.1–1000 MeV band, which is one order of magnitude higher than the Voyager-I result. Above
the GeV band, the cosmic-ray spectrum has the particle index of 2.7 and the energy density above
the GeV band is 1 eV cm−3 . In the case that the cosmic-ray spectrum keeps its shape below 1 GeV,
the energy density in the 0.1–1000 MeV band is estimated to be ∼ 600 eV cm−3 . Our result is oneorder of magnitude lower than the estimation.
The origin of LECRs have rarely discussed so far due to poor information. One clue to the origin may come from the diﬀusion length of LECRs: only a few tens of parsecs (Dogiel et al., 2011).
Therefore LECRs should be produced in situ. SNRs and pulsar wind nebulae are thought to generate
Galactic cosmic rays. However, in either the GC east, l = −20◦ region or the entire GRXE, no candidate source is found in the vicinity of the observed regions. Another acceleration mechanism isrequired.
One possible scenario is stochastic acceleration. However, the acceleration is too slow in general. Amano et al. (2011) argued that stochastic acceleration powered by Alfvénic turbulence caneﬃciently generate cosmic rays and that large velocity dispersion of ∼ 100 km s−1 in the CMZ indicates the presence of turbulence. Since also Clump 2 at l = 3◦ exhibits large velocity dispersionof ∼ 100 km s−1 (Bania, 1977; Torii et al., 2010), the stochastic acceleration can produce the LECR
protons in the GC east region. Nevertheless, the stochastic acceleration would not be necessarilyeﬃcient in the other regions.
Another possibility is a reconnection of magnetic field (Lesch and Reich, 1992; Yokoyama et al.,2000). The magnetic field is amplified by the rotation of the Galactic disk, and expands from the disk
to outside by the Parker instability. The expanding magnetic loop collides with ambient field lines
and induces magnetic reconnection. Since the velocity dispersion of Clump 2 suggests the existenceof magnetic loop (Fukui et al., 2006), the LECR protons in the GC east can be accelerated by thereconnection. Also High-velocity clouds or supernovae can trigger the magnetic reconnection in the
Galactic disk (Kerp et al., 1994; Tanuma et al., 1999). The LECRs in the l = −20◦ region and the entire
ridge region can be energized by the magnetic reconnection.
The LECRs may be secondary particles that are produced by interactions of higher-energy cosmic
rays with the ISM. In the proton scenario, the energy density of ∼ 10–80 eV cm−3 in the MeV bandis far higher than 1 eV cm−3 that is measured above the GeV band, and thus the secondary protons
cannot explain the measured LECRs. In the electron case, on the other hand, the energy density inthe keV band would be ∼ 0.01–0.03 eV cm−3 , which is much below that of the GeV band. Therefore
the LECR electrons could be the second electrons.
The LECRs could be possibly produced by undiscovered SNRs on the Galactic plane or an unknown acceleration mechanism. In any origin, since no observational result that demonstrates excess emission in the GeV gamma-ray has been reported in ether the GC east or l = −20◦ regions, thecosmic-ray particles are not necessary to be accelerated up to the GeV band.

7.3. WHERE LOW-ENERGY COSMIC RAYS COME FROM?

71

In order to investigate the origin of LECRs, we need more information about their spectra andspatial distribution as well as about the ambient ISM. ASTRO-H (Takahashi et al., 2014), the next
Japanese X-ray astronomy satellite, will provide new clues as described in the next chapter.
LECRs are an eﬃcient heating source of the ISM, and hence play a crucial role in the energetics
and chemical evolution of the Milky Way. Our results demonstrate that the 6.4 keV line can be aunique probe to investigate LECRs. We found that their energy density could be one or two orders ofmagnitude higher than the canonical value of 1 eV cm−3 on the Galactic plane including the excess
regions (in the proton origin). This means that the LECRs, which have been little focused on so far,should be taken into account in exploring characteristics of hydrodynamic process in our galaxy.

Chapter 8

Prospect for ASTRO-H
We have investigated the origin of the GDXE, especially of the 6.4 keV line emission, through thisdissertation, and found that interaction of cosmic rays (protons) with molecular gas contributes to the
6.4 keV line emission in the GRXE. However, its attributable fraction has a large uncertainty because
the equivalent width of the diﬀuse emission and the particle index of cosmic rays are unknown on
the Galactic plane; even in the GC east and l = −20◦ regions, there remain uncertainty in the spectral
shape. ASTRO-H, which is planned to be launched in February 2016, will oﬀer new insights into theorigin of 6.4 keV line in the GDXE and contribution of LECRs.
ASTRO-H will provide high resolution spectroscopy (energy resolution < 7 eV at 6 keV; FWHM)
with a micro-calorimeter array, the soft X-ray spectrometer (SXS). The SXS will discriminate theorigin of the 6.4 keV emission by the line width. While electrons and X-rays generate a very narrow
line with the line width of 1 eV, a proton-induced fluorescence line would be broad to ∼ 10 eV due
to multiple simultaneous ionization (Garcia et al., 1973).
Figure 8.1 shows simulation of an SXS observation in the GC east region. We assume the spectral
parameters of the GRXE obtained by Uchiyama et al. (2013) in principle. The ratio of intensities ofneutral iron Kα1 and Kα2 lines is assumed to be 2 : 1 and the sum intensity of two lines is assumed
to be the best-fit 6.4 keV line intensity in the integrated spectrum from the GC east region. We assume that a half of the 6.4 keV emission comes from point sources, where the line does not broad.
Figure 8.1a shows the case that the other half originates from X-ray reflection. Only the upper limitof the line width of < 2 eV is obtained as shown in table 8.1. On the other hand, figure 8.1b shows
the case that the other half of the 6.4 keV line emission is due to proton bombardment. The obtained
line width is 5±2 eV. We can discriminate the two cases. As in the Suzaku observation of the GC
east region, it would be meaningful to observe not only the east side but also the west side for comparison. We will definitely conclude the proton origin in the GC east region, and probably in thel = −20◦ region in a similar way.
The new satellite will also provide a broadband observation combined with a soft X-ray imaging
spectrometers covering 0.4–12 keV and hard X-ray imaging spectrometers covering 5–80 keV, which
will achieve more accurate determination of the spectral index of the continuum associated withcosmic rays induced the 6.4 keV line. ASTRO-H would open a new chapter of the study of LECRs as
well as the GDXE study.

73

CHAPTER 8. PROSPECT FOR ASTRO-H

74

Counts s−1 keV−1

10−1
10−2
10−3

10−1
Counts s−1 keV−1(a) X-ray origin(b) X-ray origin + proton origin

10−2
10−3

6.3

6.4

6.5

6.6

6.7

6.8

Energy (keV)
Figure 8.1. Simulation of an SXS/ASTRO-H observation in the GC east region with the exposure time of
300 ksec. (a) In the case that all the 6.4 keV line originates from X-ray reflection. (b) In the case that a half of
the 6.4 keV line emission originates from X-ray reflection and the rest are due to proton bombardment.

Table 8.1. Simulation result of an SXS/ASTRO-H in the GC east region with the exposure time of 300 ksec
(figure 8.1)*.

Origin
X-ray
X-rays+Protons

Line width of neutral iron Kα1 and Kα2< 2 eV5±2 eV* Errors are quoted at 68% confidence level.

Chapter 9

Conclusion
We compiled the Suzaku observation data of the Galactic plane in order to measure spatial distributions and equivalent widths of the 6.4 keV line, and for the comparison, the 6.7 keV and 7.0 keVlines and 5–8 keV continuum emission in the GDXE. A summary of this dissertation is as follows:• The latitudinal distribution of each line of the GDXE demonstrates that the GDXE consists ofthe three components: the GCXE, GBXE, and GRXE.
• We discover the enhancements of the 6.4 keV line emission relative to the 6.7 keV line in thetwo regions: (1) the GC east (1.◦ 5 < l < 3.◦ 5) and (2) l = −20◦ . Both the enhancements areassociated with molecular clouds. The most plausible irradiating particles are LECR protons in
the GC east. Also the l = −20◦ region favors the LECR scenario although whether the irradiating
particles are protons or electrons cannot be constrained.
• The scale height and equivalent width of the 6.4 keV line emission in the GRXE indicate thatat least a half of the 6.4 keV line emission originates from molecular clouds irradiated by Xrays or cosmic-ray particles. Since the 6.4 keV line due to the X-ray irradiation from the LMXB
contributes only 10% of the total 6.4 keV line emission in the GRXE, the most possible origin
is cosmic-ray bombardment.
• The energy density of LECR protons in the GC east would be ∼ 80 eV cm−3 while those in thel = −20◦ region would be ∼ 20 eV cm−3 and ∼ 0.03 eV cm−3 for the proton and electron cases,
respectively. In the case that a half of the 6.4 keV line emission in the GRXE originates fromthe cosmic-ray bombardment, the energy densities of LECRs along the Galactic plane would be∼10 eV cm−3 and ∼0.01 eV cm−3 for protons and electrons, respectively.
• Since the diﬀuse length of LECRs is only a few tens of parsecs, they should be produced in situ.
There is neither SNRs nor pulsar wind nebulae in the vicinity of the GC east, l = −20◦ , andthe entire GRXE region. The cosmic rays in the two enhancement regions can be energized bystochastic acceleration or magnetic reconnection. The magnetic reconnection may produce the
LECR protons in the entire GRXE region. In the electron scenario, they could be the second
particles that are produced via interaction of high-energy cosmic rays and the ISM.
• The scale heights and equivalent widths of the iron K-shell lines in the GBXE indicate that the
majority of the iron lines is due to point sources although non-magnetic CVs should contribute75

CHAPTER 9. CONCLUSION

76

much more than previous study discussed.
• The scale heights and equivalent widths of the iron K-shell lines in the GCXE shows that the
majority of the iron lines cannot be explained by point sources.

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Appendix:
X-ray image and selected region
We show the X-ray images in the energy band of 2–8 keV with the FOVs of the XIS (squares) andexcluded regions (those with a red diagonal).

85

APPENDIX86

RXJ1713-3946-BKGD2

RXJ1713-3946-BKGD1

Sgr_A_west

1.000-0.3000.200-0.400

0.100

0.700

0.600

Galactic latitude

0.800

Galactic latitude

Galactic latitude

0.900-0.500-0.6000.000-0.100-0.700-0.200-0.800-0.3000.500

347.900 347.800 347.700 347.600 347.500 347.400

346.100 346.000 345.900 345.800 345.700 345.600 345.500

Galactic longitude

Galactic longitude
0.00010.00020.0004 0.0007

0.0001HESS J1616-508

Galactic latitude-0.2000.00020.0004 0.0007

HESS J1616-508_BGD2
0.100

0.000

0.000

359.500

0.00010.0004 0.0007

0.100-0.100359.700

Galactic longitude

HESS J1616-508_BGD1

0.100

Galactic latitude

0.00020.000 359.900

0.000

Galactic latitude

0.400-0.100-0.200-0.100-0.200-0.300-0.300-0.300-0.400-0.400-0.400332.700 332.600 332.500 332.400 332.300 332.200 332.100

332.300 332.200 332.100 332.000 331.900 331.800 331.700

333.000 332.900 332.800 332.700 332.600 332.500 332.400

Galactic longitude

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.0001

Sgr_A_west

0.00020.0004 0.0007

0.00010.00020.0004 0.0007

IGRJ16465-4507

HESS J1837-069
0.400

0.200
0.100

0.300

0.000

-0.1000.000

Galactic latitude

Galactic latitude

Galactic latitude

0.100-0.100-0.200

0.200

0.100

0.000

-0.200-0.300-0.100-0.300-0.4000.000 359.900

359.700

359.500

0.00010.000225.500 25.400 25.300 25.200 25.100 25.000

340.300 340.200 340.100 340.000 339.900 339.800 339.700

Galactic longitude

Galactic longitude

Galactic longitude
0.00010.0004 0.0007

0.00020.0004 0.0007

IGRJ16195-4945

IGRJ16493-4348

0.500

0.400

0.300

0.200

Galactic longitude
0.00020.0004 0.0007

0.0001HESSJ1813-178

0.0002Galactic latitude

0.100

0.000

-0.100-0.20013.100 13.000 12.900 12.800 12.700 12.600

0.0004 0.0007-0.200-0.100-0.300-0.200-0.3000.00020.0004 0.0007-0.500-0.600-0.500-0.7009.800

288.500 288.400 288.300 288.200 288.100 288.000

9.700

Galactic longitude

Galactic longitude
0.0001-0.400-0.4009.900

0.0004 0.0007

1E 1048.1-5937

0.000

10.200 10.100 10.000

Galactic longitude

9.900

Galactic longitude

Galactic latitude

0.200

Galactic latitude

10.200 10.100 10.000

SGR 1806-20

0.300

0.0002-0.500Galactic longitude
0.00010.0001-0.3000.000
333.800 333.700 333.600 333.500 333.400 333.300 333.200

341.600 341.500 341.400 341.300 341.200 341.100

9.700-0.200-0.400

0.100

0.300

9.800

-0.1000.400

Galactic latitude

Galactic latitude

Galactic latitude

0.600

0.0004 0.0007

0.000

0.500

0.700

0.0002

SGR 1806-20

0.600
0.800

0.00010.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

87

IGRJ16393-4643

V603 AQL

PSR J1846-0258

1.100
0.300

0.000

0.200

-0.100

0.900

0.800

Galactic latitude

Galactic latitude

Galactic latitude

1.000

0.100

0.000

-0.200-0.3000.700-0.100-0.4000.600-0.20033.400 33.300 33.200 33.100 33.000 32.900-0.500338.300 338.200 338.100 338.000 337.900 337.800 337.700

Galactic longitude

0.00010.00020.0004 0.0007

0.00010.00020.00010.0004 0.0007

EMS01095

EMS0918

0.300

0.100

Galactic latitude

Galactic latitude

Galactic latitude

0.200-0.100-0.200-0.3000.000

0.0004 0.0007

0.400

0.000

0.300

0.0002

EMS1308

0.100
0.400

30.000 29.900 29.800 29.700 29.600 29.500

Galactic longitude

Galactic longitude

0.200

0.100

0.000

-0.100-0.400311.900 311.800 311.700 311.600 311.500 311.400

Galactic longitude

54.900 54.800 54.700 54.600 54.500 54.400

340.700 340.600 340.500 340.400 340.300 340.200

Galactic longitude

Galactic longitude
0.00010.00020.0004 0.0007
0.00010.00020.00010.0004 0.0007

SGR 1806-20

0.00020.0004 0.0007

AX J1622.1-5005

IGRJ16479-4514
0.100

Galactic latitude

Galactic latitude-0.100-0.200-0.300-0.400

0.100

0.000

0.000

-0.100Galactic latitude

0.000-0.100-0.200-0.300-0.500
9.900

9.800

-0.500333.900 333.800 333.700 333.600 333.500 333.400 333.300

340.400 340.300 340.200 340.100 340.000 339.900

Galactic longitude-0.300-0.400-0.40010.300 10.200 10.100 10.000-0.200Galactic longitude
0.00010.00020.0004 0.0007

Galactic longitude
0.0001

AX J1620.7-4942

0.00020.0004 0.0007

0.00010.00020.0004 0.0007

J1620-4927

AX J1445.7-5931
0.700

0.500
0.400

0.600

0.300

0.200

0.300

Galactic latitude

Galactic latitude

Galactic latitude

0.400

0.200

0.100

0.100

0.500

0.400

0.300

0.000
0.200

0.000
-0.100334.000 333.900 333.800 333.700 333.600 333.500

317.200 317.100 317.000 316.900 316.800 316.700
334.200 334.100 334.000 333.900 333.800 333.700 333.600

Galactic longitude

Galactic longitude

Galactic longitude
0.00010.00020.00010.0004 0.0007

0.00020.0004 0.0007

0.0001

J1522-5734

J1803-2149

0.00020.0004 0.0007

CXO J172641.7-354052
0.000

0.400

-0.2000.300-0.300

0.200

0.100

Galactic latitude

Galactic latitude

Galactic latitude-0.100-0.400-0.500-0.200-0.300-0.400-0.6000.000-0.500-0.700-0.100
8.400

8.300

8.200

8.100

8.000

7.900

322.300 322.200 322.100 322.000 321.900 321.800

Galactic longitude352.400 352.300 352.200 352.100 352.000 351.900

Galactic longitude
0.00010.00020.0004 0.0007

Galactic longitude
0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

APPENDIX88

HESS J1804-216 BGDHESS J1804-216
0.200

0.200

0.100

0.100

GALACTIC RIDGE

0.000-0.100Galactic latitude

Galactic latitude

Galactic latitude

0.000

0.000

-0.100-0.200-0.200-0.300-0.300-0.100-0.200-0.300-0.400-0.500
8.700

8.600

8.500

8.400

8.300

8.200

8.300

8.200

Galactic longitude

8.100

8.000

7.900

7.800

28.700 28.600 28.500 28.400 28.300 28.200

Galactic longitude

Galactic longitude
0.00010.00020.0004 0.0007

0.00010.00020.00010.0004 0.0007

0.00020.0004 0.0007

GC SOUTH BGDGC SOUTHGALACTIC RIDGE
0.100

0.400
0.000
0.300

-0.200-0.300-0.200-0.300

0.200

0.100

0.000

-0.400-0.500-0.100Galactic latitude-0.100Galactic latitude

Galactic latitude

0.000-0.400-0.10028.700 28.600 28.500 28.400 28.300 28.200

0.000 359.900

Galactic longitude359.700

0.200

Galactic longitude
0.00010.00020.0004 0.0007

0.0001HESS J1614-518

0.000 359.900

359.700

Galactic longitude
0.00020.0004 0.0007

0.0001HESS J1614-518 BG

0.00020.0004 0.0007

HESS J1825-137-0.100-0.300-0.500-0.200-0.500-0.600Galactic latitude

Galactic latitude

Galactic latitude-0.400-0.300-0.400-0.600-0.700-0.800-0.500-0.700-0.900-0.600-0.800-1.000330.700 330.600 330.500 330.400 330.300 330.200 330.100

331.800 331.700 331.600 331.500 331.400 331.300

0.00010.00020.0004 0.0007

0.00010.00020.00010.0004 0.0007

0.00020.0004 0.0007

GALACTIC CENTERGALACTIC CENTERHESS J1825-137 BG

18.100 18.000 17.900 17.800 17.700 17.600

Galactic longitude

Galactic longitude

Galactic longitude

0.300
0.600

-0.6000.200

Galactic latitude

Galactic latitude-0.800-0.900Galactic latitude

0.500-0.700

0.400

0.300

0.100

0.000

-0.1000.200-1.000-0.2000.100-1.10018.700 18.600 18.500 18.400 18.300 18.200

0.000 359.900

Galactic longitude

0.00010.00020.0004 0.0007

0.00020.0004 0.0007

0.300

0.200

0.200

0.100

0.100

0.000

-0.100-0.200-0.2000.000-0.1001.800

0.00020.0004 0.0007-0.300-0.400-0.500-0.600
1.700

1.600

1.500

1.400

1.300

1.200

Galactic longitude

Galactic longitude

0.0004 0.0007-0.100-0.200358.400 358.300 358.200 358.100 358.000 357.900

0.0002

CAR-D1

Galactic latitude

0.300

0.00010.0001GC SGR D EAST

Galactic latitude

Galactic latitude

358.800 358.700 358.600 358.500 358.400 358.300 358.200

Galactic longitude
0.0001GALACTIC CENTER-0.300359.700

Galactic longitude

0.00010.00020.0004 0.0007

288.100 288.000 287.900 287.800 287.700 287.600

Galactic longitude

0.00010.00020.0004 0.0007

89

HESS J1718-385

GC14G12.0-0.1
0.200-0.400-0.3000.100-0.500-0.500-0.6000.000

Galactic latitude

Galactic latitude

Galactic latitude-0.400-0.100-0.200-0.600-0.700-0.800-0.700-0.300-0.900-0.80012.200 12.100 12.000 11.900 11.800 11.700

349.100 349.000 348.900 348.800 348.700 348.600 348.500

0.00010.00020.400

Galactic longitude

Galactic longitude

0.00010.0004 0.0007

GC15

0.200

0.000 359.900

Galactic longitude
0.00020.0004 0.0007

0.00010.00020.0004 0.0007

0.00020.0004 0.0007

0.00020.0004 0.0007

GC17

GC16
-0.700-0.700-0.800-0.800-0.900-0.900-0.400-0.600-0.700Galactic latitude

Galactic latitude

Galactic latitude-0.500-1.000-1.100-1.000-1.100-0.800-1.200-1.200-0.9000.000 359.900-1.300359.700

0.400

0.200

-1.3000.000 359.900

Galactic longitude

Galactic longitude
0.00010.00020.0004 0.0007

359.700

Galactic longitude
0.0001GC18

0.000 359.900

0.00020.0004 0.0007

0.0001

GC20

GC19

0.600
0.900

0.900

0.800

0.800

0.400

0.300

Galactic latitude

Galactic latitude

Galactic latitude

0.500

0.700

0.600

0.700

0.600

0.200
0.500

0.500
0.100

0.400

0.400
0.400

0.200

0.000 359.900

0.400

Galactic longitude

0.200

0.000 359.900

0.000 359.900

Galactic longitude
0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

0.0001HESS J1702-420

GC21359.700

Galactic longitude

M17 EAST BKG

0.300-0.600
0.200

0.000

0.100

-0.1000.000-0.100Galactic latitude

Galactic latitude

Galactic latitude-0.700-0.200-0.300-0.800-0.900-1.000-0.400-0.200-1.100-0.300-0.500
2.100

2.000

1.900

1.800

1.700

1.600

16.100 16.000 15.900 15.800 15.700 15.600

344.500 344.400 344.300 344.200 344.100 344.000

Galactic longitude

Galactic longitude

Galactic longitude
0.00010.00020.0004 0.0007

0.0001GALACTIC BULGE

0.00020.00010.0004 0.0007

0.0004 0.0007

G54.1+0.3

G296.1-0.5N-1.7000.00020.500-0.100-1.8000.400-1.900-2.000Galactic latitude

Galactic latitude

Galactic latitude-0.200-0.300-0.400-2.100

0.300

0.200

0.100

-0.500-2.2000.000
296.200 296.100 296.000 295.900 295.800-2.3000.200

0.000 359.900

359.700

Galactic longitude

54.300 54.200 54.100 54.000 53.900 53.800

Galactic longitude

Galactic longitude
0.00010.00010.00020.0004 0.0007

0.00020.0004 0.0007
0.00010.00020.0004 0.0007

APPENDIX90

GC_LARGEPROJECT1

GC_LARGEPROJECT8

GC_LARGEPROJECT2-0.1000.200

0.400-0.200

0.100

0.200

0.100

Galactic latitude

Galactic latitude

Galactic latitude

0.300-0.300-0.4000.000-0.100-0.500-0.2000.000-0.600-0.300-0.100
0.600

0.500

0.400

0.300

0.200

0.200

0.100

0.00010.00020.0004 0.0007

358.200 358.100 358.000 357.900 357.800 357.700 357.600

Galactic longitude
0.00010.00020.0004 0.0007

0.0001

0.200

0.100

0.100

-0.100Galactic latitude

0.200

0.100

Galactic latitude

0.200

0.000

0.00020.000-0.1000.000-0.100-0.200-0.200-0.200-0.300-0.300-0.300357.900 357.800 357.700 357.600 357.500 357.400

357.700 357.600 357.500 357.400 357.300 357.200 357.100

357.400 357.300 357.200 357.100 357.000 356.900

Galactic longitude

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.0001

GC_LARGEPROJECT120.00020.0004 0.0007

0.0001

0.200

0.100

0.100

Galactic latitude

0.200

0.100

Galactic latitude

0.200-0.1000.00020.0004 0.0007

GC_LARGEPROJECT12

GC_LARGEPROJECT120.000

0.0004 0.0007

GC_LARGEPROJECT11

GC_LARGEPROJECT10

GC_LARGEPROJECT9

Galactic latitude359.700

Galactic longitude

Galactic longitude

Galactic latitude

0.000 359.900

0.000-0.1000.000-0.100-0.200-0.200-0.200-0.300-0.300-0.300357.200 357.100 357.000 356.900 356.800 356.700 356.600

357.200 357.100 357.000 356.900 356.800 356.700 356.600

357.200 357.100 357.000 356.900 356.800 356.700 356.600

Galactic longitude

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

0.0001

GC_LARGEPROJECT14

GC_LARGEPROJECT13

0.00020.0004 0.0007

LOOP 1 L=356.00
1.000

0.200

0.200

0.100

0.100

0.000

-0.100Galactic latitude

Galactic latitude

Galactic latitude

0.900

0.000

-0.100-0.200-0.200-0.300-0.300

0.800

0.700

0.600

0.500

356.900 356.800 356.700 356.600 356.500 356.400

356.700 356.600 356.500 356.400 356.300 356.200 356.100

356.300 356.200 356.100 356.000 355.900 355.800 355.700

Galactic longitude

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.0001

LOOP 2 L=356.33

0.0004 0.0007

0.0001-0.100

0.900

0.600

0.100

Galactic latitude

Galactic latitude

0.700-0.300-0.400-0.500

0.500

6.000

5.900

5.800

5.700

6.000

5.600

0.00020.0004 0.0007

5.900

5.800

5.700

5.600

5.500

Galactic longitude

Galactic longitude
0.0001-0.100-0.3006.100

Galactic longitude

0.000-0.200-0.600356.600 356.500 356.400 356.300 356.200 356.100

0.0004 0.0007

0.200-0.2000.800

0.0002HESS J1800-240C

HESS J1800-240B

1.000

Galactic latitude

0.00020.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

91

HESS J1825-137

HESS_J1825-137_2

HESS_J1825-137_3-0.300-0.600-0.200-0.700-0.300-0.800-0.900Galactic latitude

Galactic latitude

Galactic latitude-0.400-0.400-0.500-0.500-0.600-0.700-1.000-0.600-1.100-0.700-0.80017.900 17.800 17.700 17.600 17.500 17.400

17.700 17.600 17.500 17.400 17.300 17.200

18.000 17.900 17.800 17.700 17.600 17.500

Galactic longitude

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

0.00010.0004 0.0007

HESSJ1614-B

HESSJ1614-A

EXTENDED CHIMNEY

0.0002-0.500-0.4000.400-0.600-0.500

0.200

0.100

Galactic latitude

Galactic latitude

Galactic latitude

0.300-0.700-0.800-0.600-0.700-0.900-0.8000.000-1.000-0.900-0.100359.800 359.700 359.600 359.500 359.400 359.300

331.600 331.500 331.400 331.300 331.200 331.100 331.000

Galactic longitude

331.700 331.600 331.500 331.400 331.300 331.200

Galactic longitude
0.00010.00020.0004 0.0007

Galactic longitude
0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

0.00020.0004 0.0007

GC(0 -1.6)

HESS_J1809-193_S

HESS_J1809-193_N
0.300-1.400

0.300
0.200

-1.500

0.100

0.000

Galactic latitude

Galactic latitude

Galactic latitude

0.200
0.100

0.000

-0.100-1.600-1.700-1.800-0.100-0.200-1.900-0.20011.300 11.200 11.100 11.000 10.900 10.800

11.100 11.000 10.900 10.800 10.700 10.600

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.200

0.00010.00020.0004 0.0007

0.0001

GCL1

G12.8-0.0

0.000 359.900

Galactic longitude

SLX 1737-282
1.500

1.400
0.200
1.400
1.300

0.000

-0.100Galactic latitude

Galactic latitude

Galactic latitude

0.100
1.200

1.100

1.300

1.200

1.100

1.000
-0.2001.000
0.900-0.30013.100 13.000 12.900 12.800 12.700 12.600

0.000 359.900

Galactic longitude

0.00010.00020.0004 0.0007

0.900

359.500

0.0002-0.100-0.400-0.5000.0001-0.700-0.500358.800 358.700 358.600 358.500 358.400 358.300

Galactic longitude

Galactic longitude
0.00020.0004 0.0007-0.500-0.600-0.700-0.800349.200 349.100 349.000 348.900 348.800 348.700

0.00010.0004 0.0007-0.400-0.300-0.4000.0002-0.300-0.200-0.600359.700

HESS J1745-303B

Galactic latitude-0.300Galactic latitude

0.000

0.000 359.900

0.0004 0.0007

HESS J1745-303A-0.2000.200

Galactic longitude
0.0001HESS J1718-385

Galactic latitude359.700

Galactic longitude

358.700 358.600 358.500 358.400 358.300 358.200

Galactic longitude
0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

APPENDIX

92

HESS J1745-303C

TORNADOHESS J1427-608-0.6000.100-0.7000.000-0.800-0.9000.000

Galactic latitude

Galactic latitude

Galactic latitude

0.100-0.100-0.200-1.000-0.300-1.100-0.400-0.100-0.200-0.300-0.400358.000 357.900 357.800 357.700 357.600 357.500 357.400

314.600 314.500 314.400 314.300 314.200
358.800 358.700 358.600 358.500 358.400 358.300

Galactic longitude

Galactic longitude

Galactic longitude
0.00010.00020.00010.0004 0.0007

0.00020.0004 0.0007

0.00010.00020.0004 0.0007

ASO0218

ASO0449

ASO0431-0.600-0.200-0.100-0.700-0.300-0.800-0.900-1.000Galactic latitude

Galactic latitude

Galactic latitude-0.200-0.300-0.400-0.400-0.500-0.600-0.500-1.100-0.700-0.60040.400 40.300 40.200 40.100 40.000 39.900

0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

0.0001

ASO03760.00020.0004 0.0007

NONE

1.300

0.700

-1.2001.200

Galactic latitude

0.500

0.400

0.300

-1.3001.100

Galactic latitude

0.600

Galactic latitude

Galactic longitude

Galactic longitude

ASO0447

1.000

0.900

-1.500-1.700
0.400

0.700

65.500 65.400 65.300 65.200 65.100 65.000-1.400-1.600

0.800

0.200

356.500 356.400 356.300 356.200 356.100 356.000

Galactic longitude
0.00010.00020.0004 0.0007

0.00010.00020.0001

0.200

0.100

0.000

-0.10026.700 26.600 26.500 26.400 26.300 26.200

0.000

0.200

-0.1000.100-0.200-0.300-0.100-0.200-0.500-0.30011.900 11.800 11.700 11.600 11.500 11.400

11.600 11.500 11.400 11.300 11.200 11.100

Galactic longitude
0.00010.00020.0004 0.0007

Galactic longitude

0.00010.00020.0004 0.0007

0.0001GALACTIC_BULGE1HESS J1809-193E

0.0004 0.0007

0.000-0.400Galactic longitude

0.0002HESS J1809-193B

Galactic latitude

Galactic latitude

0.300

0.000 359.900

0.0004 0.0007

HESS J1809-193A

0.400

0.200

Galactic longitude

Galactic longitude

ASO0418

Galactic latitude285.300 285.200 285.100 285.000 284.900 284.800

66.100 66.000 65.900 65.800 65.700 65.600

Galactic longitude

0.00020.0004 0.0007

GALACTIC_BULGE2-0.9000.600-0.600-1.000
0.500

0.400

0.300

Galactic latitude

Galactic latitude

Galactic latitude-0.700-0.800-0.900-1.100-1.200-1.3000.200-1.000-1.4000.100-1.10011.000 10.900 10.800 10.700 10.600 10.500

0.200

Galactic longitude

0.000 359.900

359.700

0.00010.00020.0004 0.0007-1.5000.200

0.000 359.900

359.700

Galactic longitude

Galactic longitude
0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

93

GALACTIC_BULGE7GALACTIC_BULGE8GALACTIC_BULGE9-1.300-0.900-1.900-1.400-1.000-1.200-1.300-1.500Galactic latitude

Galactic latitude

Galactic latitude-2.000-1.100-1.600-1.700-2.100-2.200-2.300-1.800-1.400-2.400358.800 358.700 358.600 358.500 358.400 358.300

358.800 358.700 358.600 358.500 358.400 358.300 358.200

358.800 358.700 358.600 358.500 358.400 358.300

Galactic longitude

Galactic longitude

Galactic longitude
0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

0.0001GALACTIC_BULGE11-2.600-3.600-4.800-2.700-3.700-4.900-2.900Galactic latitude-3.500-2.800-3.800-3.9000.0004 0.0007-5.000-5.100-5.200-4.000-3.0000.0002GALACTIC_BULGE12-2.500Galactic latitude

Galactic latitude

GALACTIC_BULGE10

358.800 358.700 358.600 358.500 358.400 358.300 358.200

358.800 358.700 358.600 358.500 358.400 358.300 358.200

358.800 358.700 358.600 358.500 358.400 358.300 358.200

Galactic longitude

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.00010.00020.00010.0004 0.0007

G22.0+0.0

G25.5+0.0
0.300

0.300

0.200

0.200

0.100

0.100

0.00020.0004 0.0007

G23.5+0.1
0.300

0.000

-0.100Galactic latitude

Galactic latitude

Galactic latitude

0.200

0.000

-0.100

0.100

0.000

-0.100-0.200-0.200-0.20025.800 25.700 25.600 25.500 25.400 25.300

23.700 23.600 23.500 23.400 23.300 23.200

22.200 22.100 22.000 21.900 21.800 21.700

Galactic longitude

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

0.00010.0004 0.0007

FVW319.8+0.3_BGD

FVW319.8+0.3

G37.0-0.1

0.00020.200-0.100-0.200

0.500

0.500

0.400

0.400

Galactic latitude

0.000

Galactic latitude

Galactic latitude

0.100

0.300

0.200

37.300 37.200 37.100 37.000 36.900 36.800

319.600 319.500 319.400 319.300 319.200 319.100

319.900 319.800 319.700 319.600 319.500 319.400

Galactic longitude

0.200

0.100

0.100

-0.3000.300

Galactic longitude

Galactic longitude
0.00010.00020.0004 0.0007

0.0001WESTERLUND1_1

0.00020.00010.0004 0.0007

0.00020.0004 0.0007

WESTERLUND1_3

WESTERLUND1_2-0.500-0.900-0.700-0.600Galactic latitude

Galactic latitude-0.900-1.000Galactic latitude-1.000-0.800-1.100-1.200-0.700-0.800-0.900-1.300-1.100-1.000-1.400-1.200-1.100339.300 339.200 339.100 339.000 338.900 338.800 338.700

340.100 340.000 339.900 339.800 339.700 339.600 339.500

339.700 339.600 339.500 339.400 339.300 339.200

Galactic longitude

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

APPENDIX94

WESTERLUND1_4

AX J1714.1-3912
SNR G304.6+0.1

Galactic latitude-0.900-1.000-0.1000.200

Galactic latitude

Galactic latitude

0.000

0.300

-0.800

0.100

0.000

-1.100-0.200-0.300-0.400-1.200304.800 304.700 304.600 304.500 304.400

Galactic longitude-0.500-1.300348.100 348.000 347.900 347.800 347.700 347.600

341.000 340.900 340.800 340.700 340.600 340.500

Galactic longitude

Galactic longitude
0.00010.00010.00020.0004 0.0007
0.0001GALACTIC_BULGE4-3.500-0.100-2.600-3.600-0.200-2.700-3.700-0.300-2.900Galactic latitude-2.500-2.800-3.800-3.900-4.000-3.0000.200

0.000 359.900

0.200-0.400-0.5000.000 359.900

359.700

356.700 356.600 356.500 356.400 356.300 356.200

Galactic longitude

Galactic longitude
0.00010.00020.0004 0.0007-0.600-4.100359.700

0.0002GALACTIC_BULGE13

GALACTIC_BULGE5

Galactic latitude

Galactic latitude

0.00020.0004 0.0007

Galactic longitude
0.00010.0004 0.0007

GALACTIC_BULGE14

0.00020.0004 0.0007

0.00010.00020.0004 0.0007

GALACTIC_BULGE16

GALACTIC_BULGE15-1.200-2.100-1.300-0.800-0.900-1.000-1.400Galactic latitude-0.700Galactic latitude

Galactic latitude-0.600-1.500-1.600-2.300-2.400-2.500-1.700-2.600-1.100356.600 356.500 356.400 356.300 356.200 356.100

356.700 356.600 356.500 356.400 356.300 356.200

356.600 356.500 356.400 356.300 356.200 356.100

Galactic longitude

Galactic longitude

Galactic longitude
0.00010.00020.0004 0.0007

0.0001GALACTIC_BULGE17

0.00020.00010.0004 0.0007

HESS_J1741-055_1-0.100-3.4000.100-0.200-3.600-3.700Galactic latitude

0.200-3.5000.000-0.100-0.300-0.400-0.60026.600 26.500 26.400 26.300 26.200 26.100

356.600 356.500 356.400 356.300 356.200 356.100

0.00010.00020.00010.0004 0.0007

HESS_J1741-055_3

26.600 26.500 26.400 26.300 26.200 26.100

Galactic longitude

Galactic longitude

Galactic longitude

0.0004 0.0007-0.500-0.200-3.8000.0002

HESS_J1741-055_2-3.300Galactic latitude

Galactic latitude-2.2000.00020.0004 0.0007

0.0001

HESS_J1741-055_4

0.00020.0004 0.0007

HESSJ1857-A

0.000
0.100

0.300
-0.100-0.200Galactic latitude

0.200

Galactic latitude

Galactic latitude

0.000-0.100-0.200-0.300

0.100

0.000

-0.400-0.300-0.100-0.500-0.400-0.20036.200 36.100 36.000 35.900 35.800 35.700

27.400 27.300 27.200 27.100 27.000 26.900

27.000 26.900 26.800 26.700 26.600 26.500

Galactic longitude

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

95

HESS J1457-593

G18.8+0.3

HESSJ1832
0.600-0.200
0.500

0.600

0.400

0.500

0.300

0.200

Galactic latitude

Galactic latitude

Galactic latitude-0.300

0.400

0.300

-0.400-0.500-0.600
0.200

0.100

-0.7000.100

23.500 23.400 23.300 23.200 23.100 23.000

318.600 318.500 318.400 318.300 318.200 318.100

19.000 18.900 18.800 18.700 18.600 18.500

Galactic longitude

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007

G337.2-0.7

NONE
0.300

GRXE_E_2-0.500
0.200

0.200

0.000

-0.1000.100

Galactic latitude

Galactic latitude

Galactic latitude-0.6000.100-0.700-0.8000.000-0.100-0.200-0.900-0.200-0.300-1.00019.800 19.700 19.600 19.500 19.400 19.300

2.300 2.200 2.100 2.000 1.900 1.800 1.700

Galactic longitude
0.00010.00020.0004 0.0007

0.0001

GRXE_E_3

0.00020.0004 0.0007

GRXE_E_4
0.200

0.100

0.100

-0.100Galactic latitude

0.200

0.100

0.000

0.000

-0.100

2.500

2.400

2.300

2.200

2.100

2.800

2.000

2.700

0.0002

2.600

2.500

2.400

0.0004 0.0007

2.300

0.00010.0002

3.100

3.000

0.00010.0002

1.600

1.500

0.00010.0002-0.100

3.000

2.200

GRXE_E_6

2.900

2.800

2.700

Galactic longitude
0.00010.00020.0004 0.0007

0.0004 0.0007

GRXE_E_10

GRXE_E_9

0.200

2.500

0.000

Galactic longitude

Galactic longitude
0.00012.600

0.0004 0.0007-0.300-0.300-0.3000.0002-0.200-0.200-0.2000.0001

GRXE_E_5

0.200

Galactic latitude

Galactic latitude

Galactic longitude337.500 337.400 337.300 337.200 337.100 337.000 336.900

Galactic longitude

0.700

0.400
0.600

0.000

-0.100-0.200Galactic latitude

0.300

Galactic latitude

Galactic latitude

0.100

0.200

0.100

0.500

0.400

0.300

0.000
0.200

-0.3003.300-0.100
3.200

3.100

3.000

2.900

2.800

2.700

3.400

Galactic longitude

3.300

3.200

3.100

3.000

2.900

0.00010.00023.500

0.00010.0002

3.400

3.300

3.200

Galactic longitude

Galactic longitude
0.0004 0.0007

HESSJ1858-A

GRXE_E_1

HESSJ1858-B-0.300-0.100-0.400-0.200-0.500-0.3000.0004 0.0007

0.0004 0.0007

0.200-0.600-0.700Galactic latitude

Galactic latitude

Galactic latitude

0.100-0.400-0.5000.000-0.100-0.200-0.600-0.800-0.30035.900 35.800 35.700 35.600 35.500 35.400

35.800 35.700 35.600 35.500 35.400 35.300

2.000

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

1.900

1.800

1.700

Galactic longitude
0.00010.00020.0004 0.0007

0.0004 0.0007

APPENDIX96

GRXE_E_7

GRXE_E_8

3C 396
0.000

0.200

0.200

0.100

0.100

0.000

-0.100Galactic latitude

Galactic latitude

Galactic latitude-0.1000.000-0.100-0.200-0.200-0.300-0.300-0.200-0.300-0.400-0.500

3.500

3.400

3.300

3.200

3.100

3.800

3.000

3.700

3.600

3.500

3.400

3.300

3.200

0.00010.00020.0004 0.0007

39.400 39.300 39.200 39.100 39.000 38.900

Galactic longitude

Galactic longitude

Galactic longitude

0.00010.00020.0004 0.0007

0.00010.00020.0004 0.0007