Supernova Remnants Selected with X-Rays— Contribution to the Galactic Cosmic-Ray Acceleration

Masaru Ueno
Department of Physics, Graduate School of Science, Kyoto University
Kitashirakawa Oiwake-cho, Sakyo-ku, Kyoto, 606-8502, Japanmasaru@cr.scphys.kyoto-u.ac.jp

This thesis was submitted to the Department of Physics,
Graduate School of Science, Kyoto University
on January 5, 2005
in partial fulfillment of the requirements
for the degree of Doctor of Philosophy in physics.

Abstract
SNRs have been the primary candidate for the galactic cosmic-ray acceleration, becauseof the large energy output. Since the discovery of synchrotron X-ray emission from SN 1006,
nowadays synchrotron X-ray and Tev-gamma rays have been detected from several SNRs.
Thus, existence of high energy particles in SNRs are being established. However, among
more than 200 galactic SNRs, existence of high energy particles (∼ 1 TeV) are shown in only
about 10 SNRs. And as a result, we do not know in which SNRs high energy particles areaccelerated. SNRs have been mainly discovered with radio surveys. However, SNRs emitting
synchrotron X-rays, such as SN 1006 or G347.3−0.5 are tend to be faint in radio. Therefore,
we noticed SNRs discovered X-rays or X-ray selected SNRs.
First, we analyzed XM M -N ewton data of SNRs discovered with the ROSAT All Sky
Survey. With ASCA observations, hard X-ray emission in G156.2+5.7 has already been
known. Spatially resolving the contribution of point sources with XM M -N ewton, we confirmed the hard X-ray emission is truly diffuse. Due to the high background, however, we
could not determine the emission mechanism of the hard X-rays. We performed spatiallyresolved spectroscopy on G272.2−3.2 and G299.2−2.9, resulting in no detection of nonthermal X-ray emission.
Next, we searched SNR candidates with the ASCA Galactic Plane Survey. The ASCA
Galactic Plane Survey is the first imaging survey of the Galactic plane in the energy band>3 keV. We made exposure corrected image and analyses the survey data in detail. As aresult, 15 extended X-ray sources were found in the hard band (2.0–7.0 keV). Although several sources are likely to be star-forming regions, pulsar-wind nebulae or clusters of galaxies,
most of the new sources are likely to be SNRs.
Follow-up observations with ASCA, XM M -N ewton, and Chandra have been performed on 5 of the new extended sources. As a result, G28.6−0.1 and G32.45+0.1 areestablished as SNRs emitting synchrotron X-rays. Although we confirmed non-thermal Xray emission from G11.0+0.0, G25.5−0.1, and G26.6−0.1, observations with higher spatial
resolution is required to find out the emission mechanism and their origins.
SNRs are known to get fainter in radio as they evolve, which is known as the Σ-D(radio surface brightness vs. diameter) relation. When we compared the SNRs emitting
synchrotron X-rays with other SNRs, we found the former are faint in radio even when they
are young. Since SNR brightnesses in radio have dependency on the ambient density, wesuggest SNRs emitting synchrotron X-rays are located in low density media.

Contents

1 Introduction

12 Reviews

3

2.1 Cosmic Ray . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

3

2.1.1

Spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

3

2.1.2

Chemical Composition . . . . . . . . . . . . . . . . . . . . . . . . . .

4

2.1.3

Energy Budget . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

5

2.1.4

Fermi Acceleration . . . . . . . . . . . . . . . . . . . . . . . . . . . .

6

2.2 Supernovae and Supernova Remnants . . . . . . . . . . . . . . . . . . . . . .

6

2.2.1

Supernovae and their Classification . . . . . . . . . . . . . . . . . . .

6

2.2.2

Evolution of Supernova Remnant . . . . . . . . . . . . . . . . . . . .

8

2.2.3

Shock Wave . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

9

2.2.4

Plasma Physics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

11

2.2.5

Non-Equilibrium Ionization . . . . . . . . . . . . . . . . . . . . . . .

11

2.2.6

X-ray emission from Supernova Remnants . . . . . . . . . . . . . . .

12

2.3 Other Diffuse X-ray Sources . . . . . . . . . . . . . . . . . . . . . . . . . . .

15

2.3.1

Pulsar-wind nebula . . . . . . . . . . . . . . . . . . . . . . . . . . . .

15

2.3.2

Star-forming Regions . . . . . . . . . . . . . . . . . . . . . . . . . . .

15

2.3.3

Cluster of Galaxies . . . . . . . . . . . . . . . . . . . . . . . . . . . .

16

i

ii

CONTENTS

3 Instruments
3.1 ASCA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

17
17

3.1.1

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

18

3.1.2

GIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

18

3.1.3

SIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

21

3.2 XM M -N ewton . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

23

3.2.1

X-ray Telescopes . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

24

3.2.2

EPIC Cameras . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

25

3.2.3

EPIC background . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

28

3.3 Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

28

3.3.1

HRMA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

30

3.3.2

ACIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

31

4 SNRs Discovered with ROSAT
4.1 G156.2+5.7 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

35
36

4.1.1

Previous Observations . . . . . . . . . . . . . . . . . . . . . . . . . .

36

4.1.2

Mapping Observations with ASCA . . . . . . . . . . . . . . . . . . .

37

4.1.3

Observations with XM M -N ewton . . . . . . . . . . . . . . . . . . .

42

4.1.4

Discussion on G156.2+5.7 . . . . . . . . . . . . . . . . . . . . . . . .

48

4.2 G272.2−3.2 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

49

4.2.1

Previous Observations . . . . . . . . . . . . . . . . . . . . . . . . . .

49

4.2.2

Observation, Data Reduction and Image . . . . . . . . . . . . . . . .

50

4.2.3

Spectral Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

51

4.2.4

Discussion on G272.2−3.2 . . . . . . . . . . . . . . . . . . . . . . . .

53

4.3 G299.2−2.9 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

54

4.3.1

Previous Observations . . . . . . . . . . . . . . . . . . . . . . . . . .

54

iii

CONTENTS
4.3.2

XM M -N ewton Observations . . . . . . . . . . . . . . . . . . . . . .

55

4.3.3

Discussion on G299.2−2.9 . . . . . . . . . . . . . . . . . . . . . . . .

56

5 ASCA Search for SNRs

61

5.1 ASCA Galactic Plane Survey . . . . . . . . . . . . . . . . . . . . . . . . . .

61

5.2 New Extended Hard X-ray Sources . . . . . . . . . . . . . . . . . . . . . . .

62

5.2.1

Source Finding . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

62

5.2.2

Spatial and Spectral Characteristics of Newly Found Sources . . . . .

67

5.3 Comments on the individual sources . . . . . . . . . . . . . . . . . . . . . . .

75

5.4 Summary of the Extended Source Survey . . . . . . . . . . . . . . . . . . . .

78

6 Follow-up Observations
6.1 G11.0+0.0, G25.5+0.0, and G26.6−0.1 . . . . . . . . . . . . . . . . . . . . .

81
81

6.1.1

Observations and Images . . . . . . . . . . . . . . . . . . . . . . . . .

81

6.1.2

Spectral Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

82

6.1.3

Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

83

6.2 G28.6−0.1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

88

6.2.1

Follow-up Observations with ASCA . . . . . . . . . . . . . . . . . . .

88

6.2.2

Chandra Observations of G28.6−0.1 . . . . . . . . . . . . . . . . . .

89

6.2.3

Discussion on G28.6−0.1 and surroundings . . . . . . . . . . . . . . .

96

6.3 G32.45+0.1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

99

6.3.1

XM M -N ewton Observation . . . . . . . . . . . . . . . . . . . . . . .

6.3.2

Discussion on G32.45+0.1 . . . . . . . . . . . . . . . . . . . . . . . . 101

7 Discussion

99

105

7.1 Properties of SNRs emitting synchrotron X-rays . . . . . . . . . . . . . . . . 105
7.2 Properties of the SNR candidates . . . . . . . . . . . . . . . . . . . . . . . . 111

iv

CONTENTS

A Background Light Curves of XM M -N ewton

115

B Serendipitous Discovery of a Cluster of Galaxy with XM M -N ewton

119

C Diffusive Shock Acceleration

123

List of Figures
2.1 Energy spectrum of cosmic-ray protons. . . . . . . . . . . . . . . . . . . . . .

4

2.2 Chemical composition of cosmic rays. . . . . . . . . . . . . . . . . . . . . . .

5

2.3 The basic classification scheme for supernovae based on spectral features at
early times. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

7

2.4 A schematic view around the shock front. . . . . . . . . . . . . . . . . . . . .

10

3.1 The schematic view of the ASCA satellite. . . . . . . . . . . . . . . . . . . .

17

3.2 The schematic view of the Wolter type-I configuration. . . . . . . . . . . . .

18

3.3 Cross section view of the GIS sensor. . . . . . . . . . . . . . . . . . . . . . .

20

3.4 Schematic diagram of the principle of the GIS detector. . . . . . . . . . . . .

20

3.5 GIS spectra of day earth, blank sky, and night earth from Ishisaki (1997). . .

21

3.6 Sketch of the XM M -N ewton payload. . . . . . . . . . . . . . . . . . . . . .

23

3.7 The light path of the two XM M -N ewton telescopes. . . . . . . . . . . . . .

24

3.8 The on-axis effective areas of the XM M -N ewton telescopes. . . . . . . . . .

25

3.9 (left): The vignetting factor as a function of off-axis angle at a few selectedenergies, of the X-ray telescope in front of the pn camera. (right): Vignetting
effect as a function of azimuth angle of the X-ray telescope in front of the
MOS1 camera. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

26

3.10 A sketch of the FOV of the two types of EPIC cameras; MOS (left) and pn(right). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

27

v

vi

LIST OF FIGURES
3.11 Background spectra for the PN (grey) and MOS1 (black) cameras duringobservations with the filter wheel in the CLOSED position. . . . . . . . . .

29

3.12 Images of Al-K (left) and Si-K (right) fluorescent emissions from the MOS
camera. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

29

3.13 Schematic view of the Chandra satellite. . . . . . . . . . . . . . . . . . . . .

30

3.14 (left):On-axis effective areas of HRMA as a function of X-ray energy. (right):Offaxis effective areas of HRMA in several X-ray energies. . . . . . . . . . . . . 31
3.15 (left):On-axis effective areas of HRMA as a function of X-ray energy. (right):Offaxis effective areas of HRMA in several X-ray energies. . . . . . . . . . . . . 32
3.16 A schematic drawing of the ACIS-array configuration.

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

3.17 (left): The pre-launch energy resolution of ACIS. (right): The energy resolution of chip S3 and I3 as a function of row number, after the proton damage
on orbit. The data were taken at the CCD temperature of −120◦ C. . . . . .

33

33

4.1 ROSAT gray-scale image of G156.7+5.7 with intensity map of the radio polarized component. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

37

4.2 ASCA GIS images of G156.2+5.7 in the energy bands of 0.7–1.5 keV (a),1.5–2.0 keV (b), and 2.0–7.0 keV (c). . . . . . . . . . . . . . . . . . . . . . .

39

4.3 Background-subtracted ASCA spectra of the SE (a), Center (b), NW (c), andNW-shell (d) regions. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

41

4.4 XM M -N ewton/MOS1+2 images of the north-west region of G156.2+5.7 in
the 0.5–2.0 keV (a) and 2.0–7.0 keV (b) bands. . . . . . . . . . . . . . . . . .

44

4.5 XM M -N ewton/MOS1+2 images of the center region of G156.2+5.7 in the
0.5–2.0 keV (a) and 2.0–7.0 keV (b) bands. . . . . . . . . . . . . . . . . . . .

45

4.6 EPIC spectra of the point sources detected in the G156.2+5.7 observations. .

47

4.7 Cumulative source counts N (S) in the 2–10 keV within the 120 -radius regionsof the G156.2+5.7 observations. . . . . . . . . . . . . . . . . . . . . . . . .

48

4.8 MOS1+2 spectra of the source and background regions of the NW observation. 49
4.9 MOS1+2 spectra of the NW region. The best-fit NEI model is shown withthe solid-line. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

49

vii

LIST OF FIGURES
4.10 Comparison of spectral regions of G156.2+5.7 north-west. . . . . . . . . . . .

50

4.11 XM M -N ewton MOS1+2 images of G272.2-3.2 in the energy bands of 0.7–
1.7 keV (a), 1.7–2.0 keV (b), 2.3–2.6 keV (c), and 3.0–7.0 keV (d). . . . . . .

52

4.12 Background subtracted spectra of 3 regions of G272.2-3.2. (a): the Outerregion. (b): the Inner region. (c): the region A. . . . . . . . . . . . . . . . .

53

4.13 XM M -N ewton MOS1+2 images of G299.2−2.9 in the energy bands of 0.5–
1.7 keV (a), 1.7–2.0 keV (b), 2.3–2.6 keV (c), and 2.0–7.0 keV (d). . . . . . .

57

4.14 Background subtracted spectra of 4 regions of G299.2-2.9. . . . . . . . . . .

58

5.1 Flow chart of how to create an exposure-corrected mosaic image. The numbersin the parentheses correspond to the procedures in the text. . . . . . . . . .

65

5.2 ASCA GIS images of newly found extended X-ray sources in the 2.0–7.0 keVband. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

68

5.3 Radial profiles of the newly detected SNR candidates.

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

71

5.4 ASCA GIS spectra of the newly found SNR candidates.

. . . . . . . . . . .

73

6.1 GIS images with Galactic coordinate around (a)G11.0+0.0, (b)G25.5+0.0,
and (c)G26.6−0.1 in the 0.7–7.0 keV band. . . . . . . . . . . . . . . . . . . .

83

6.2 Background-subtracted spectra of (a) G11.0+0.0, (b) G25.5+0.0, and (c)G26.6−0.1. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

84

6.3 ASCA GIS images in the 0.7–2.0 keV (a) and 2.0–7.0 keV (b) bands. . . . .

90

6.4 Background-subtracted spectrum of the combined GIS2 and 3 data of the 3
observations. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

90

6.5 Combined ACIS-I image of observation 1 and observation 2 in the 1.5–8.0 keVband. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

93

6.6 Close-up view of CXOU J184357−035441 in the energy band 1.0–6.0 keV. . .

94

6.7 Chandra spectra of G28.6−0.1. . . . . . . . . . . . . . . . . . . . . . . . . .

95

6.8 Background-subtracted spectrum of CXOU J184357−035441. Solid line: Bestfit NEI model. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

96

viii

LIST OF FIGURES
6.9 XM M -N ewton/MOS1+2 images of G32.45+0.1 in the soft (0.5–2.0 keV; (a))
and the hard (2.0–7.0 keV; (b)) bands. . . . . . . . . . . . . . . . . . . . . . 100
6.10 MOS1+2 hard-band (2.0–7.0 keV) image with the spectral regions. . . . . . 102
6.11 Background-subtracted spectra of G32.45+0.1 observed with MOS1 and MOS2.102
7.1 The surface brightness vs. diameter (Σ–D) relation for shell SNRs with known
distances. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 106
7.2 The ratio between Σs expected by the Σ–D relation and observed Σs plottedagainst ambient densities. . . . . . . . . . . . . . . . . . . . . . . . . . . . . 109
7.3 Plot of diameters and fluxes for the new extended X-ray sources. . . . . . . . 112
7.4 ASCA 0.7–7.0 keV image around G28.6−0.1 overlaid with the contours of the
CO intensity map. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 113
A.1 Background light curves of XM M -N ewton observations (1). . . . . . . . . . 116
A.2 Background light curves of XM M -N ewton observations (2). . . . . . . . . . 117
B.1 MOS1+2 image of XMMJ J0456372+522412 in the 1.0–7.0 keV band. . . . . 121
B.2 Radial profile of XMMJ J0456372+522412 in the 1.0–7.0 keV bands. . . . . . 121
B.3 EPIC spectra of XMMJ J0456372+522412. . . . . . . . . . . . . . . . . . . . 122
B.4 2MASS Ks-band image around XMMJ J0456372+522412 overlaid with the
MOS1+2 contours in the 1.0–7.0 keV band. . . . . . . . . . . . . . . . . . . 122

List of Tables
3.1 Design parameters and performance of the ASCA XRT . . . . . . . . . . . .

19

3.2 Design parameters and performance of the GIS . . . . . . . . . . . . . . . .

19

3.3 The selection criteria for GIS data. . . . . . . . . . . . . . . . . . . . . . . .

22

3.4 Design parameters and performance of the SIS . . . . . . . . . . . . . . . . .

22

3.5 Design parameters and performance of the EPIC MOS and pn . . . . . . . .

27

4.1 SNRs discovered with RASS and follow-up observations . . . . . . . . . . . .

36

4.2 ASCA observations performed on G156.2+5.7. . . . . . . . . . . . . . . . . .

38

4.3 Best-Fit Model Parameters of the ASCA G156.2+5.7 spectra. . . . . . . . .

42

4.4 XM M -N ewton observations performed on G156.2+5.7. . . . . . . . . . . . .

43

4.5 Bright Point Sources Detected in the G156.2+5.7 Observations. . . . . . . .

46

4.6 Comparison of the hard X-ray fluxes in the G156.2+5.7 regiona . . . . . . . .

46

4.7 Best-fit Spectral parameters for the NW region with a NEI model. . . . . . .

48

4.8 XM M -N ewton observations performed on G272.2−3.2. . . . . . . . . . . . .

51

4.9 Best-Fit Model Parameters for G272.2−3.2. . . . . . . . . . . . . . . . . . .

54

4.10 XM M -N ewton observations performed on G299.2−2.9. . . . . . . . . . . . .

55

4.11 Best-Fit Model Parameters for G299.2−2.9. . . . . . . . . . . . . . . . . . .

56

5.1 Parameters of newly found extended hard X-ray sources. . . . . . . . . . . .

66

ix

x

LIST OF TABLES
6.1 Follow-up observations on the SNR candidates . . . . . . . . . . . . . . . . .

81

6.2 ASCA observations performed on G11.0+0.0, G25.5+0.0, and G26.6−0.1 . .

82

6.3 Best-fit Spectral Parameters of G11.0+0.0, G25.5+0.0, and G26.6−0.1 . . .

85

6.4 Physical Parameters of the Diffuse Sources . . . . . . . . . . . . . . . . . . .

85

6.5 Parameters for thermal plasma scenario. . . . . . . . . . . . . . . . . . . . .

87

6.6 ASCA observations performed on G28.6−0.1. . . . . . . . . . . . . . . . . .

89

6.7 Best-fit parameters of G28.6−0.1 for a power-law and a thin thermal plasma
model with the ASCA observations∗ . . . . . . . . . . . . . . . . . . . . . . .

91

6.8 Best-Fit Model Parameters for G28.6−0.1 with Chandra. . . . . . . . . . . .

96

6.9 Best-fit parameters for CXOU J184357−035441 by the NEI model.

. . . . .

97

6.10 Best-Fit Parameters by the SRCUT Model . . . . . . . . . . . . . . . . . . .

98

6.11 The positions and the extents of the point-like sources in the shell (crosses inthe ellipses in Figure 6.9). . . . . . . . . . . . . . . . . . . . . . . . . . . . . 100
6.12 Best-fit parameters of G32.45−0.1 for a power-law and a thin thermal plasma
model with the XM M -N ewton observations. . . . . . . . . . . . . . . . . . 101
6.13 Best-Fit Parameters by the SRCUT Model . . . . . . . . . . . . . . . . . . . 103
7.1 Shell SNRs with Known Distances . . . . . . . . . . . . . . . . . . . . . . . . 107
7.2 Shell SNRs Emitting Synchrotron X-rays . . . . . . . . . . . . . . . . . . . . 108
B.1 Best-fit parameters of XMMJ J0456372+522412 with a APEC model. . . . . 120

Chapter 1
Introduction
Supernovae are energetic explosions of stars at the ends of their lives. The total energy
released in one event is up to 1053 ergs, most part of which is emitted as neutrinos, and1/100 is emitted as kinetic energy of the surrounding matter. Along the explosions, elements produced in the progenitor stars and also those produced during the explosions are
distributed outwards, and change the composition of the interstellar matter. Therefore, supernovae have significant influence on chemical evolution of galaxies. The kinetic energyof supernovae may also trigger next births of stars, compressing ambient molecular clouds.
Supernova are important events in understanding evolutions of the interstellar matter.
We here concentrate on the aspect of supernovae as an origin of cosmic rays. The
origin of cosmic rays has been a mystery since the discovery by Hess in 1912. The spectrumof cosmic rays showering onto the earth has been intensively studied, and nowadays it is
known that the spectrum continues up to 1022 eV. However, the origin of cosmic rays evenwith energy of 1015 eV is yet unknown. Since cosmic rays themselves are charged particles,
their directions are easily changed by the magnetic field filling the galaxy. For example, aproton of energy 1015 eV in a magnetic field of ∼ 1 µG gyrates with a radius of 0.3 pc.
Therefore, it is hard to know where cosmic rays come from. One of the break through wasmade by an observation with X-rays, which travel straight. Koyama et al. (1995) foundsynchrotron X-ray emissions from a supernova remnant, SN1006, which discovery showedexistence of high energy electrons in the supernova remnant. Recently, TeV gamma rays
are also detected from several supernova remnants, and existence of high energy particles insupernova remnants is getting firmly evidenced.
Because of the enormous energy input, supernova remnants have been a primary candidate for the acceleration site of the galactic cosmic rays. Now the discovery of synchrotron
1

2

CHAPTER 1. INTRODUCTION

X-rays and TeV gamma rays evidenced existence of high energy particles in supernova remnants. However, a question of what fraction of the galactic cosmic rays can be explainedwith supernova remnants is not yet solved. In order to answer this question we have toknow roughly two things; How much energy is injected in each supernova remnants andhow many supernova remnants are located in our Galaxy. We know only small number ofsupernova remnants showing evidence of acceleration, which is making the above systematic
study difficult.
In order to work on these problems, we use X-ray observations which have many advantages in studying the acceleration in supernova remnants. First of all, the X-ray band isjust the wavelength of synchrotron emission from high energy electrons of energy ∼10 13 eV
in a typical magnetic field of ∼ 3 µG in the Galactic plane. X-rays, especially in the highenergy band, have strong penetrating power through the interstellar matter, so that evenclusters of galaxies behind the galactic plane can be observed. Nowadays, the instrumentsof X-ray observations are highly developed and as a result they are equipped with large
effective areas, high energy resolution, and high imaging capability. These advantages arereflected to the situation that acceleration in a new supernova remnant has been evidenced
mostly by detection of a synchrotron X-ray emission.
In this thesis, we concentrate on a study of supernova remnants selected with X-rays.
The rate of synchrotron X-ray detection from SNRs which are discovered with X-rays is muchhigher than that from SNRs discovered with other wavelengths. First of all, in Chapter 4, weanalyze supernova remnants discovered with the ROSAT All Sky Survey. In Chapter 5, wereport on search for supernova remnants using the ASCA Galactic Plane Survey. We have
performed follow-up observations on a part of the newly discovered candidates with ASCA,
Chandra, and XM M -N ewton, which will be explained in Chapter 6. We will discuss overallcharacteristics of the X-ray selected SNRs and their contribution to the galactic cosmic-ray
acceleration in Chapter 7.

Chapter 2
Reviews
2.1

Cosmic Ray

Cosmic rays are discovered by Victor Hess in 1912. He measured the ionization level of theatmosphere up to 5 km altitude on a balloon and found that some radiation must comefrom space to explain the ionization state of the atmosphere. After that, many experiments
have been done and the spectrum and the composition of cosmic rays have been measured
extensively. However, the origin of the cosmic rays are yet unknown. In this section, wesummarize the properties of the cosmic rays, which are the key information to constrain
their origin.

2.1.1

Spectrum

Figure 2.1 shows the energy spectrum of cosmic rays (the vertical axis is differential flux
multiplied by E 2.5 ). The spectrum extends up to ∼ 1020 eV at least, which is much higher
than the energy obtained with any particle accelerators on earth. The spectrum is almost
power-law like and featureless with few breaks. The break at ∼10 15.5 eV is called “knee”.
(There is also another break at ∼1020 eV called “ankle”.) The spectrum up to 1020 eV canbe expressed quite well as:
N (E) ∼ 10−2.7 (<∼ 1015.5 eV)

N (E) ∼ 10−3.0 (>∼ 1015.5 eV).

(2.1)
(2.2)

The “knee” is one of the few features informing the origin of the cosmic ray. The energy densities of cosmic rays in the galactic plane is 1 eV cm−3 , which is large compared with that of
the starlight, galactic magnetic field, or the cosmic micro-wave background (' 0.3 eV cm −3 ).
3

4

CHAPTER 2. REVIEWS

Figure 2.1: Energy spectrum of cosmic-ray protons (Sokolsky, 1989). The horizontal axis isenergy of each proton, and the vertical is (differential flux) × (energy)2.5 .

2.1.2

Chemical Composition

Figure 2.2 shows the chemical composition of cosmic rays in comparison with the solar
abundance. In a way similar to the solar abundance, the species with even atomic number
are more abundant than those with odd number, though the difference is smaller than thesolar values. Heavy elements are more abundant than the solar. It is notable cosmic rays
are significantly overabundant in Lithium, Beryllium, and Boron. Since these elements
are not produced either in the primordial nucleosynthesis or in the stellar evolution, it isthought that they themselves are not accelerated but they are produced by spallation, or
inelastic interaction of cosmic rays with the interstellar matter. Using the cross sections of
the spallations, the column density along the pass of cosmic rays from the birth place to theearth is estimated to be ∼5–10 g cm−2 . For a density of 1 proton cm−3 , the travel distanceis estimated to to be ∼1000 kpc, which is much larger than the diameter of our Galaxy.
Therefore, cosmic rays should be confined by the galactic magnetic field.

52.1. COSMIC RAY

Figure 2.2: Chemical composition of cosmic rays (Sokolsky, 1989). The horizontal axis
is atomic number, and the vertical axis is relative abundance. For comparison, the solar
abundance, is designated with the solid line.

2.1.3

Energy Budget

If we assume the lifetime of cosmic rays is detemined mainly by escape from the Galaxy,
the energy for maintaining the cosmic rays in our Galaxy (Lcr ) can be estimated from thelife-time tcr , the energy density cr , and the volume filled by cosmic rays VG as
Lcr =

VG 
.
tcr

(2.3)

The life-time of cosmic rays is estimated from the observation of radio isotopes such as 10 B,
and is known to be tcr ∼ 6×106 yr. As mentioned above, the energy density is  ∼ 1 eV cm−3 .
The volume filled by cosmic rays is approximated as
VG ' π × rg2 × dg ∼ π × (15 kpc)2 × 200 pc ∼ 4 × 1066 cm3 ,(2.4)

where rg and dg are the radius and the thickness of the disk of the galaxy, respectively. As
a result, the energy is estimated to be Lcr ∼ 5 × 1040 ergs s−1 .
Among galactic sources, supernova remnants are the most strong candidate for thecosmic ray acceleration on this point of view. Supernovae are thought to occur one event per

6

CHAPTER 2. REVIEWS∼30 yr, emitting kinetic energy of 1051 erg per each event. Therefore, if 10% of this kinetic
energy is injected into the acceleration, the above required energy for the galactic cosmic ray
acceleration can be fulfilled by SNRs.

2.1.4

Fermi Acceleration

Particles moving between two media which have different speeds can gain energy. By moving
back and forth many times, even high energy can be attainable. When the speed difference
is formed by an infinite plane shock of compression ratio 4, the energy spectrum of theaccelerated particles is expected to be a power-law of index 2 (Bell 1978). Since this value isproper to explain the spectrum of the cosmic ray, the Fermi Acceleration mechanism has been
widely accepted as a standard acceleration mechanism (see e.g., recent review by Malkov &Drury 2001). Acceleration time-scale is summarized in Appendix C.

2.2

Supernovae and Supernova Remnants

2.2.1

Supernovae and their Classification

The term supernova was coined by Baade & Zwicky (1934) to distinguish from the moreordinary novae those anomalously bright stellar outbursts which at maximum often outshine their entire parent galaxy. This phenomenon is now known to be an explosion of a starat the end of its evolutionary lifetime, and the explosion energy is as large as 1051 ergs.
Historically astronomers have recognized two types of supernova which are distinguishedby their optical spectra. Figure 2.3 shows the basic classification scheme. The class of asupernova has traditionally depended on the spectrum near maximum light. Those without
H line in their spectra are called Type I, and those with H line are called Type II.
Among the SN II, a further differentiation has been proposed based on the shape of thelight curve. Those with a pronounced plateau have been termed Type II plateau (SN II-P),
and those with a nearly linear decline in magnitude with time from peak are termed Type II
linear (SN II-L). However, the recently exploded SN 1987A event falls neither of case. The
light curve of SN 1987A was very different due to the blue compact nature of the progenitor,
although it showed some similarity to Type II plateau on the exponential tail (see Dopita(1988); Arnett, Bahcall, Kirshner, & Woosley (1989), and references therein). In addition,
SN 1987A has given a new spectral behavior. Traditional SN II-L and P events have showna nearly pure continuum near maximum light, a week or two after the explosion. SN 1987A,

2.2. SUPERNOVAE AND SUPERNOVA REMNANTS

7

Figure 2.3: The basic classification scheme for supernovae based on spectral features at earlytimes (Harkness & Wheeler, 1990).

in contrast, showed strong P Cygni features of the Balmer lines from the first day.
Type I supernovae should also divided into subclasses. The existence of two populationsof Type I supernova precursors was evidenced by their occurrence throughout spiral andelliptical galaxies and their correlation with the rate of star formation. Firstly, it was foundthat a part of Type I SNe show no Si ii absorption at λ ∼ 6150 Å near maximum light incontrast to the “normal” Type I (Type Ia) events. Those peculiar Type I (Type Ib) events
have common characteristics different from SN Ia. At maximum light, SN Ib tend to be
weaker and redder than SN Ia, and they show stronger non-thermal radio emission than
SN Ia. They occur in S galaxies and regions of star formation. The most clear difference
between Type Ia and Ib show up 25 days or more after the explosion: lines of [Fe ii] and[Fe iii] dominate the spectra of Type Ia events, while lines of [O i] dominate those of TypeIb events (Gaskell et al., 1986). None of these distinctions is purely phenomenological, andthey all bear on the specific physics of the situation. Thus the term SN I will hence beconsidered to be essentially meaningless unless it is further qualified as SN Ia and SN Ib.
Nowadays, the SN Ia is thought be an explosion of white dwarfs in binary systems withmass transfer. Material is added to the white dwarf until the core reaches a critical density
2 × 109 g cm−3 , resulting in an uncontrolled fusion of carbon and oxygen, thus detonating
the white dwarf. The progenitors of SN Ib is thought to be massive stars that have lost
all their hydrogen envelope. Wolf-Rayet stars, for example, are possible progenitors. SN II
is an explosion of a massive star more than 10 M . The iron core produced after nuclear

8

CHAPTER 2. REVIEWS

fusion process along the evolution of the star grows so massive that the core cannot sustain
the gravitation of itself. Then the core collapse and release the gravitational energy, whichturns into a supernova.

2.2.2

Evolution of Supernova Remnant

A supernova remnant (SNR) is the aftermath of the explosion, which expands as a shellwithin the interstellar medium (ISM) for thousands of years. The evolution of an SNR canbe divided into several phases.
Free Expansion Phase
The initial phase of evolution is characterized by the free expansion of ejected shell at avelocity v0 = (5 − 10) × 103 km s−1 , with essentially no deceleration. The deceleration of
the shock front starts when the mass of swept-up matter becomes comparable to that of theejected matter, which is the end of the first phase. The radius and age of the remnant atthat point are
M0
Rs = (3M0 /4πµmH n0 )1/3 = 1.9
1M!1/3 

µ1.36

−1/3 

n0
1cm−3

−1/3

pc,

(2.5)

and
t ≈ Rs v0−1 ≈ 190 yr,(2.6)

in a medium with mean ambient hydrogen density n0 = 1 cm−3 , and with M0 = 1M ,v0 = 104 km s−1 , and µ = 1.36 (mean atomic weight of cosmic material per H atom). In thisstage, almost all the initial explosion energy is manifested by kinetic energy of the ejected
matter.
Sedov Phase (Adiabatic Phase)
The next phase of the SNR evolution is called an adiabatic phase. In this phase, the shock
front is started to decelerate by encountering with the ambient interstellar matter. Shklovskii(1962) proved that the SN explosion within the interstellar matter can be likened to a point
explosion in a gas with constant heat capacity. He also proved that the self-similar Sedov(1959) solution was applicable. At this stage, the evolution of the radius Rs and the velocityof the shock wave Vs , are described as,
Rs

E0
= 5.0
51
10 ergs!1/5 

n0
1 cm−3

−1/5t103 yr

!2/5

pc,

(2.7)

9

2.2. SUPERNOVAE AND SUPERNOVA REMNANTS

Vs ≡

2 Rs
dRs
=
,
dt
5 t

(2.8)

where E0 , n0 , and t are the ambient density, initial explosion energy, and age, respectively.
Numerical calculations by Chevalier (1974) showed that in this stage about 70% of the initial
explosion energy has been transformed into thermal energy of the swept-up interstellar gas.
Radiative Cooling Phase
The remnant continues to expand adiabatically until the radiative cooling is important,when the gas temperature behind the shock front reaches the value that corresponds to the
maximum on the radiative loss curve, kT ∼ (4 − 5) × 105 K. According to Falle (1981), theage tcool , radius Rcool , and expansion velocity Vcool of an SNR entering into this radiative
cooling stage are,tcool = 2.7 × 10

2.2.3

4

E0
51
10 ergs!0.24 !0.295 

Rcool

E0
= 20
51
10 ergs

Vcool

E0
= 277
51
10 ergsn0
1 cm−3n0
1 cm−3!0.055 

−0.52yr,(2.9)

−0.409pc,(2.10)

0.111km s−1 .

(2.11)

n0
1 cm−3

Shock Wave

When the shell of an SNR speed exceeds the sound speed, a shock wave is formed, and thegas is heated. We here define up-stream and down-stream as Figure 2.4 on the coordinate
in which the shock is at rest. There are physical properties between gases in the upstream
and downstream of the shock wave. Conservation of mass, momentum, and energy acrossthe shock front give Rankine-Hugoniot shock jump conditions, which can be expressed as:
ρ1 v1 = ρ 2 v2

(2.12)

ρ1 v12 + p1 = ρ2 v22 + p2
1 2
1 2
v1 + h1 =v + h22
2 2(2.13)(2.14)

where ρ, v, P ,h are the density, velocity, pressure, and specific enthalpy of both upstream(subscript 1) and downstream (subscript 2), respectively. If the fluid can be assumed to beideal gas, the third equation can be written as,
1 2
γ p1
1
γ p2
v1 +
= v22 +
2γ − 1 ρ12
γ − 1 ρ2

(2.15)

10

CHAPTER 2. REVIEWS

upstream

downstream

v1, ρ1, P1

v2, ρ2, P2
shock front

Figure 2.4: A schematic view around the shock front.

where γ is the specific-heat ratio. If we define sound speed (a1 ) and the Mach number M1
in the downstream as,
a1 =

s

M1 =

v1
,
a1

γp1
ρ1

(2.16)(2.17)

then the equations can be written asv1(γ + 1)M12
ρ1
=
=
ρ2
v2
(γ − 1)M12 + 2
p2
2γM12 − (γ − 1)
=
.
p1
γ+1

(2.18)(2.19)

With the equation of state,
T2
p2 ρ1
=
T1
p1 ρ2
[2γM12 − (γ − 1)](γ − 1)M12 + 2]=
(γ + 1)2 M12(2.20)(2.21)

When the shock is strong enough, or M1  1, the density is compressed up to (γ +1)/(γ − 1) = 4 (γ = 35 ). In this case, the temperature in the downstream can also be writtenas32(γ − 1) 2
µv1 = µv12(2.22)kT2 =2(γ + 1)16
where µ is the mean atomic weight and k is the Boltzmann constant. This is the post-shock
temperature and is determined only by the up-stream velocity, or the shock speed. For
example, when the shock velocity is 1000 km s−1 , kT2 = 1 × 107 K (= 1 keV). Since a plasma
at this temperature emits strong X-rays, X-ray observations are strong tool to investigate
the physics of supernova remnants.

2.2. SUPERNOVAE AND SUPERNOVA REMNANTS

2.2.4

11

Plasma Physics

Non-Equipartition
It has been debated whether electron and ion temperatures are quickly equilibrated byplasma processes behind collisionless high Mach number supernova shocks.
Electrons and ions (mainly protons) exchange energy through Coulomb interactions,which provide the minimum level of heating expected in the case of nonequipartition. The
temperature varies as
dTe
Tp − T e(2.23)= 0.13 3/2dt
Te(in cgs units; Spitzer 1978, p. 22).

2.2.5

Non-Equilibrium Ionization

The characteristic timescale for a plasma to reach collisional ionization equilibrium wasderived by Masai (1994, 1984). The collisional ionization rate equation for an element of
atomic number Z is written asdfz= Sz−1 fz−1 + (Sz + αz )fz + αz+1 fz+1
(2.24)d(ne t)where fz is the ionic fraction of the element that is ionized (z − 1) times, andZX

fz = 1,(2.25)

z=0

and where Sz and αz represent the rate coefficients for ionization and recombination froman ion of charge z to z + 1 and z − 1, respectively.
The ionization timescale (τ ) to reach ionization equilibrium is estimated with the zero-thorder approximation to be,ZX(Sz + αz )−1(2.26)≈ [(Sz + αz )min|Sz −αz | ]−1(2.27)

τ = ne t ≈z=0≈ 1012 s cm−3 .

(2.28)

It is notable that this time scale is nearly independent of both the element Z and the electron
temperature Te . Typically, as the electron density is 1 cm−3 in the interstellar medium,shock heated plasmas in young SNRs up to 104 yr old are expected to be in non-equilibrium
ionization. In fact, underionized plasmas have been observed from many SNRs, such as
Cassiopeia A and SN1006.

12

CHAPTER 2. REVIEWS

Heat Conduction
If the mean free path of electrons λe ≈ 1018 (Te /107 K)2 (ne /1 cm−3 )−1 is much shorter than
the scale length of the temperature gradient, lT ≡ T /|∇T |, the heat flux Q is given by,
Q = −κ|∇T |.

(2.29)

Here, κ is the thermal conductivity in a fully ionized hydrogen plasma,κ = 1.31ne kλe' 2.6 × 10

11

kTe
me

!1/2

kT1
1.0 keV(2.30)
!5/2

ln Λ
32.2

!−1

ergs s−1 K−1(2.31)

where for T > 4.2 × 105 K the Coulomb logarithm is ln Λ = 29.7 + ln n−1/2 (Te /106 K).

2.2.6

X-ray emission from Supernova Remnants

In this section, we summarize X-ray emission mechanisms seen in supernova remnants. Distinguishing between these mechanisms is an important work in determining the physicalproperties (such as temperature, density, and energies of particles) of SNRs.
Bremsstrahlung (Free-Free) Emission
Radiation due to the acceleration of a charged particle in the Coulomb field of anothercharged particle is called bremsstrahlung or free-free emission.
The total emission from a medium with ion density ni , electron density ne and a fixed
electron speed v is given bydW
32π 2 e6= √ 3 2 ne ni Z 2 gf f (v, ν)dνdV dt
3 3c m v(2.32)

where gf f (v, ν), e, m, c, and Z are a Gaunt factor, charge unit, electron mass, light speed,and charge number of ion, respectively. Bremsstrahlung from a thermal plasma can beformalized by averaging the above single-speed expression over a thermal distribution of
speeds. The emission power of the thermal bremsstrahlung is given as
dW= 6.8 × 10−38 Z 2 ne ni Te−1/2 e−hν/kTe ḡf f (ergs sec−1 cm−3 Hz−1 )dνdV dt(2.33)

in CGS units, where ḡf f is a frequency average of the velocity averaged Gaunt factor and kis the Boltzmann constant. ḡf f is in the range 1.1 to 1.5.

13

2.2. SUPERNOVAE AND SUPERNOVA REMNANTS

Typical gas in an SNR is thermalized more than 106 K and H and He are totally ionized.
As a result, bremsstrahlung is a major process of X-ray emission from an SNR.
Line Emission (Bound-Bound) Emission
The transitions between two discrete quantum levels are often accompanied by photon emissions, which are called line emission or bound-bound emission. Ionized plasma can be collisionaly excited by electrons and emits line emissions. Since the excited levels come very
quickly, for the first approximation, we can assume that the rate of photon emission is thesame as the excitation. Therefore, the line intensity is proportional to ne ni , which determines
the probability of collisions between electrons and ions. We can derive relative abundancesby comparing line intensities from different elements. Moreover, since emissivity of the continuum emission (bremsstrahlung) is proportional to ne nH , by comparing the intensity of theline and continuum emission, the absolute abundance with respect to H can be derived.
A level of transition changes with a degree of ionization. Therefore, the energy of lines
show the ionization level of that element . However, it is difficult to produce low ionizedplasmas in a ground experiment energies and strengths of emission lines from low-ionizedplasmas have been determined by interactive studies between astronomical observations andatomic theories.
Synchrotron Emission
Particles accelerated by a magnetic field will radiate. For extreme relativistic particles,
this radiation is known as synchrotron radiation. The synchrotron radiation power per unit
frequency from a single charged-particle of energy E (= γmc2 ) in a magnetic field B is√ 3 
3 q B sin αωP (ω) =,
(2.34)
F2
2π mc
ωc
where α is the pitch angle of the particle trajectory and F (x) is a dimensionless function,
which can be derived by integrating a modified Bessel function (Rybicki & Lightman 1979).
The characteristic frequency ωc is
3γ 2 qB sin α.
(2.35)
2mc
Since F (x) peaks at x ∼0.29, P (ω) has emission peak at ∼ 0.29ωc . Therefore, for a relativistic
electron of energy Ee , the peak energy of emitted photon Ep is given as,ωc =BEp ∼ 5.3
10 µG

!

Ee
100 TeV2

keV.

(2.36)

14

CHAPTER 2. REVIEWS

where Ep is averaged over α. By integrating equation (2.34), the total emission power isgiven as
2q 4 B 2 γ 2 sin2 α.
(2.37)
Ptot =
3m2 c3For synchrotron radiation emitted by electrons in a power-law distribution with index p
over a sufficiently broad energy range, the spectrum is also power-law and the spectral index
isp−1(2.38)
s=2(Rybicki & Lightman 1979).
Inverse Compton
The scattering of a high-energy (nearly equal to or larger than the rest mass energy of anelectron) photon by a low-energy free electron is called the Compton scattering, in which thephoton partly loses its energy. Conversely, when an electron is relativistic and has sufficiently
higher energy than a photon, the photon can gain energy by interacting with the electron.
This process is called inverse Compton scattering. The energy transfer by a single collision
averaged over the scattering angle is given as
E ' γ 2 E0 ,(2.39)

where E0 , E, and γ are the photon energies before and after the scattering, the Lorentz
factor of the electron, respectively. The total power transfered from an electron to photons(radiation field) by the inverse Compton scattering is given as
4
PIC = σT cγ 2 β 2 Uph ,3(2.40)

where σT , Uph , and β are cross section of Thomson scattering, energy density of the radiation
field, and the particle speed in a unit of c (light speed). The difference between synchrotron
emission and inverse Compton is the target against which an electron interact; the targetof synchrotron emission is magnetic field and that of inverse Compton is radiation (photon)field. Ratio between the emission powers of synchrotron emission and inverse Compton is
given only by
Psync
UB
=
.
(2.41)PIC
Uph
When the energy of the electron is high enough that the energy of photon in the electron rest
frame (γE0 ) is comparable or higher than the electron rest mass (me c2 ), the cross section ofThomson scattering have to be modified using the Klein-Nishina cross section.

2.3. OTHER DIFFUSE X-RAY SOURCES

2.3

Other Diffuse X-ray Sources

2.3.1

Pulsar-wind nebula

15

Rotation-powered pulsars rotate very fast (∼10 Hz) and have strong magnetic field ∼ 10 12 G.
The pulsar accelerate particles in thier magnetosphere, although the detailed acceleration
region is still controversial between a polar cap region (e.g., Sturrock 1971) and an outer gap
region (e.g., Cheng, Ho, & Ruderman 1986). A relativistic magnetized wind of accelerated
particles injected by the pulsar is thermalized at a shock where the wind comes into pressure
equilibrium with the surroundings, and creates a pulsar-wind nebula. Since a pulsar-windnebula contains relativistic electrons, synchrotron emission is observed in wide spectral band:
radio to γ-rays.
Since rotation-powered pulsars use their spinning energy into the radiation and theacceleration of particles, the spinning period P becomes gradually longer. Since the spinning
power is given as:
1
Erot = Iω 2 = 2πIP −2 ,(2.42)
2
where I is the moment of inertia of the neutron star, and ω is the angular frequency, theenergy loss rate (or spin-down luminosity) is given as:dErot ≡ Ėrot = Iω ω̇ = 4πI Ṗ P −3 .
dt

(2.43)

Pulsars which have large spin-down luminosity are detected in the X-ray and gammaray bands. For example, the famous pulsar-wind nebulae, the Crab Nebula, with the X-ray
luminosity LX = 1 × 1037 ergs s−1 , has a rotation-powered pulsar of spin-down luminosityĖrot = 5 × 1038 ergs s−1 on its center. The relation between X-ray luminosity and spin-downluminosity of pulsars has been sought since Seward & Wang (1988) found a correlation of
the form LX ∝ Ė 1.39 .

2.3.2

Star-forming Regions

Star-forming regions can be roughly divided into two categories with their sizes: high-mass
and low-mass star forming regions.
High-mass star formation is thought to occur in giant molecular clouds, which containa mass of 105 –3 × 106 M . High-mass star forming regions emit strong X-ray radiation withluminosities ∼ 1033 –1034 erg s−1 and with temperatures of several keV, sometimes accompanying diffuse emission (e.g., Townsley et al. 2003). Since the size of those star-forming

16

CHAPTER 2. REVIEWS

regions are 0.1–50 pc, X-ray emissions are very similar to those from SNRs. Moreover, they
may be enclosing even SNRs originating from rapidly evolved highly massive stars. Highmass star forming regions are also known as H ii regions; massive stars ionize ambient dense
media with strong ultra-violet radiation. Since H ii regions radiate strong radio continuum,
nowadays more than 1000 H ii regions are known (e.g., Paladini et al. 2003). Therefore, an
X-ray emitting high-mass star forming regions which is not identified as an H ii regions isquite unlikely.
In smaller molecular clouds, which are generally called “dark clouds” having a mass
about 1/1000–1/100 of giant molecular clouds, only low-mass stars are formed. These regions also emit X-ray emission with luminosity ∼ 1032 erg s−1 . Since those regions do notcontain high-mass stars, they are not accompanied by H ii regions. Therefore, if they are
deeply embedded in the galactic plane and deeply absorbed, they may escape radio surveys.
However, the size of the star-forming region is relatively small (1–2 pc), they would be easily
distinguished from SNRs.

2.3.3

Cluster of Galaxies

Clusters of galaxies are gravitationally bound systems containing typically hundreds of galaxies in a region about 106 pc in size. Clusters of Galaxies are bright X-ray sources with luminosities of 1043−45 ergs s−1 , and the emission mechanism is thermal bremsstrahlung from hot(∼ 108 K), low density (∼ 10−3 atoms cm−3 ) gas. X-ray observations of clusters of galaxies
have provided size of the total mass and about 80–90% of the mass in clusters of galaxies
are known to be dark matter. Since the gas is gravitationally bounded, the structure seenin the X-ray band is center-filled morphology.
When we see clusters of galaxies located behind the Galactic plane, we can only observe
their hard X-ray emission due to the deep interstellar absorption. Therefore, we expect tofind only high-temperature (≥ 2 keV) clusters. Moreover, the positive correlation between
X-ray luminosities and temperatures are known for clusters, we can expect to observe only
clusters of large luminosities (≥ 1044 ergs s−1 ). A deeply absorbed cluster of galaxies locatedbehind the Galactic plane is reported by Nevalainen et al. (2001). The temperature andluminosity of the cluster are ∼ 6 keV and ∼ 4 × 1044 erg s−1 , respectively.

Chapter 3
Instruments
3.1

ASCA

ASCA is a Japanese fourth X-ray satellite launched on February 20th, 1993. ASCA has four
identical X-ray telescopes with nested thin foil mirrors (XRT; Serlemitsos et al. 1995). The
focal-plane instruments consist of two kinds of detectors: two CCD cameras, and two gas
counters. The former are called Solid-state Imaging Spectrometers (SISs: Burke et al. 1994),
which are good in energy resolution. The latter are called two Gas Imaging Spectrometers(GISs: Ohashi et al. 1996; Makishima et al. 1996), which realized large field of view. In thefollowing sections, the above XRT, GIS, and SIS are described respectively.

Figure 3.1: The schematic view of the ASCA satellite.

17

18

CHAPTER 3. INSTRUMENTS
parabolahyperbola

focal plane

Figure 3.2: The schematic view of the Wolter type-I configuration.

3.1.1

XRT (X-Ray Telescope)

The XRT mirrors of ASCA utilize total X-ray reflection under the Wolter type I configuration, which employs paraboloid and hyperboloid surfaces as the primary and the secondary
mirrors (Figure 3.2).
Usually, the effective area of an X-ray mirror is increased by nesting multiple reflective
surfaces, because the critical angle for the total X-ray reflection is very small. Moreover, thecritical angle gets smaller for X-ray photons of higher energy. The ASCA XRT has realized
a drastically large number of nestings by utilizing 120 thin aluminum foils (127 µm thick).
Thus, the XRT had a large effective area even at 5∼10 keV; however, the angular resolution
was so poor as 2∼3 arcmin due to the shape errors of the mirror surfaces. In fact, since sucha thin foil is very difficult to shape into a paraboloid or hyperboloid, a conical surface wasinstead used as an approximation, which degraded the imaging quality. Although surface of
the aluminum foils was coated with lacquer to increase the surface smoothness, the residual
surface waviness added to the image degradation. Design parameters and performance of
the XRT are summarized in table 3.1.

3.1.2

GIS

GIS (Gas Imaging Spectrometer) is a gas scintillation proportional counter with imaging
capability (Ohashi et al. 1996; Makishima et al. 1996). The energy resolution was ∼ 8%
at 6 keV, which was almost the highest value ever achieved with a gas counter. The GIS
had a very low detector background, since charged particles and X-rays could be strongly
discriminated by the “rise time discrimination”. The design parameters and performance of
the GIS are summarized in Table 3.2.

3.1. ASCA

19

Table 3.1: Design parameters and performance of the ASCA XRT
Mirror substrate
Mirror surface
Mirror length
Inner / outer diameter
Focal length
Incident angle
Total weight of four XRTs
Geometrical area
Field of view

Aluminum foil (127 µm)
Acrylic Lacquer (10 µm) + Au (50 µm)
100 mm
120 mm / 345 mm
3500 mm0.24◦ − 0.7◦
∼ 40 kg
558 cm2 / telescope∼ 240 (FVWM at 1 keV)
∼ 160 (FVWM at 7 keV)

Table 3.2: Design parameters and performance of the GIS
Bandpass
Energy Resolution
Effective Area
Entrance Window
Absorption material
Positional Resolution
Minimum Time Resolution

0.7–15 keV
8% at 5.9 keV (FWHM)
50 mm (500 ) diameter
10 µm berylliumXe (96%) + He (4%), 10 mm depth, 1.2 atm at 0◦ C
0.5 mm (0.0 5; FWHM)∼ 61 µs (in PH Mode)

Structure of GIS and the X-ray detection
Two identical detectors (named GIS-2 and GIS-3) are mounted on the focal plane of XRT.
The structure of the GIS is shown in Figure 3.3, where the left side corresponds to the top
while the right to the bottom along the z-axis of the satellite coordinates in Figure 3.1. EachGIS consists of a gas cell, a position-sensitive phototube, and high voltage suppliers. The
X-rays from the XRT are detected by the 25 mm-depth gas cell, which is filled with 96%
of Xe and 4% of He to 1.20 atm pressure (at 0◦ C). The gas is sealed by a beryllium foilof a 10 µm thickness and a 52 mm diameter, supported by a stainless grid and fine mesh.
The gas cell is divided into two regions by the intermediate mesh: the drift region and thefollowing scintillation region. Each region is biased by 1 kV cm−1 and 4 kV cm−1 electricfields.
The principle of the X-ray detection by the GIS is shown in Figure 3.4. First, an X-ray

20

CHAPTER 3. INSTRUMENTS

HV

X-ray

GND
52mm

support grids

Be entrance window

photo electrons
cloud(10um)

Gas cellDrift Region

Xe + He
96%

4%

- 6kV
10mm
- 5kV

intermediate mesh

Fluorescence
UV light

Scintillation
Region15mm

ground mesh

2.5mm
3.0mm

QuartzImaging PhotoMultiplier Tube

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

GND

1400V

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

Pre-AmplifierPHRTXP / YP

Figure 3.3: Cross section view of the
GIS sensor.

Figure 3.4: Schematic diagram of the principleof the GIS detector.

photon is photo-absorbed in the drift region and creates free electrons. These electrons drift
toward the intermediate mesh guided by the electric field, then accelerated by the stronger
field in the scintillation region. These accelerated electrons repeatedly excite xenon atoms,
which produce a large number of scintillation UV photons. These UV photons are finallycollected by the imaging photo-multiplier tube attached beneath the gas cell. The energy
information is obtained from the final dynode signal. As the multiplication does not use
avalanche like gas proportional counters, higher energy resolution (∼8% at 6 keV) could beachieved.

Background Rejection and Data Screening
When X-rays are properly absorbed in the drift region, the signal should exhibit a rise timeof about 3 µs which corresponds to the drift time of elections in the scintillation region. If
a charged particle enters the gas-cell of the GIS detector, it creates a long track in the drift
region and the resultant signal exhibits a long rise time. With these discriminations, 99%
of non-X-ray backgrounds are eliminated and the residual non-X-ray backgrounds in the 17 0

3.1. ASCA

21

Ne
0.85

O
0.52

Mg
1.26

day Earth(4,620 ks)

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

Mg Al

1.26 1.49

Si
1.74

Ar
2.96

Sn-L
3.4

Xe-L
4.8

Cu
8.04

Figure 3.5: GIS spectra of day earth, blank sky, and night earth from Ishisaki (1997).

radius from the detector center is achieved to be 5 × 10−4 cts s−1 cm−2 keV−1 . In Figure 3.5,
the GIS spectra of day earth, blank sky and night earth are displayed. The night earth data
are thought to represent the non-X-ray events.
The characteristics of the RNXB have been investigated by using data from night-earth
observations (Ishisaki 1997):
1. The total counting rate of the RNXB is gradually increasing by about 8% per year.
2. The intensity of the RNXB depends on the position of the satellite in orbit and canbe reproduced as a function of the geomagnetic cut-off rigidity (COR) and H02 counts(GIS monitoring counts: G2 H0+G2 H2+G3 H0+G3 H2).
3. There are some geographical points in orbit where the intensity of the RNXB becomesanomalously high. At those points, RMB counts are founded to be high. RBM is aradiation belt monitor equipped to GIS2. Therefore, these are defined as “flare- events”
and interpreted as a phenomenon associated with a local geomagnetic anomaly.
To reject the “flare-events”, the screening procedure, called “flare-cut”, using H02 counts andRBM counts (RBM CONT) as parameters was introduced (Ishisaki 1997). The screening
conditions of the flare-cut are summarized in Table 3.3.

3.1.3

SIS

The SIS (Burke et al. 1994) is the first satellite-bourne X-ray experiment that utilizes CCDs
in the photon counting mode. Design parameters and performance of the SIS are summarizedin Table 3.4. The main characteristics is its high energy resolution realized simultaneously

22

CHAPTER 3. INSTRUMENTS

Table 3.3: The selection criteria for GIS data.
SAA
Excluded
GIS RBMF
ExcludedCOR> 4 (GV)ANG DIST > 0&&< 0.01
ELV
>10◦
SFR
Excluded
HFR
Excluded
SAA
COR
ANG DIST
ELV
SFR

:
:
:
:
:

HFR

:

South Atlantic Anomaly passage
CutOff Rigidity
ANGular DISTribution
ELeVation from the rim of the earth
Soft Flare Region(RMB CONT < 100 && G2 H0+G2 H2+G3 H0+G3 H2<=45)Hard Flare Region(G2 H0+G2 H2+G3 H0+G3 H2<=0.45×COR 2 − 13×COR+125)

with the imaging capability. One of the short points is the narrow field of view. Because of
the degradation by cosmic ray bombardment in orbit, only a part of the CCDs is operated
in each observation after several years after launch. Therefore, the SIS is not suited forour survey and study of extended SNRs. The knowledge about the X-ray CCD in spaceobtained with the SIS experiment is taken over by the CCD instruments on Chandra and
XM M -N ewton.
Table 3.4: Design parameters and performance of the SIS
Illumination Method
Number of Pixels (1 chip)
Pixel Size
Area (1 chip)
Field of View (1 chip)Thickness of Depletion Layer
Bandpass
Quantum Efficiency
Energy Resolution

Front Illumination
420 ×420
27 µm
11×11 mm square
11×11 mm square∼ 40µm
0.4–12 keV∼ 80% at 5.9 keV
2% at 5.9 keV (FWHM)

3.2. XM M -N EW T ON

3.2

23

XM M -N ewton

XM M -N ewton is a ESA’s X-ray observatory launched on December 10th, 1999. It carries
two distinct types of telescopes: three X-ray telescopes with different X-ray detectors intheir foci, and a 30-cm optical/UV telescope with a microchannel-plate pre-amplified CCD
detector in its focal plane.
XM M -N ewton provides the following three types of scientific instruments: three Xray CCD cameras called European Photon Imaging Camera (EPIC), two high-resolution
X-ray spectroscopy instruments called the Reflection Grating Spectrometer (RGS), and the
Optical Monitor (OM). Three EPIC cameras consist of two different detectors; two MOS
cameras and one pn. They provide X-ray imaging, moderate resolution spectroscopy, and
X-ray photometry, and suit for our study of supernova remnants. Since we mainly used the
EPIC cameras, characteristics of them and the associated X-ray telescopes are described
below.

Figure 3.6: Sketch of the XM M -N ewton payload. The three X-ray mirror modules arevisible at the lower left. At the right end of the assembly, the focal X-ray instruments areshown: the EPIC-MOS cameras with their radiators, the radiator of the EPIC-pn camera
and those of the RGS detectors. The OM telescope is obscured by the lower mirror module.

24

3.2.1

CHAPTER 3. INSTRUMENTS

X-ray Telescopes

Design Structure
XM M -N ewton’s three X-ray telescopes are co-aligned with a relative astrometry between
the three EPIC cameras calibrated to better than 1–200 across the full FOV. One telescope has
a light path simply illuminating the focal plane detector, the pn camera. On the other hand,
two others have grating assemblies in their light paths and part of the incoming radiation
is diffracted onto their secondary foci where RGSs are located. 44% of the incoming lightis directed onto the MOS cameras at the prime foci, while 40% of the radiation is dispersed
onto a linear CCD strips of RGSs. The remaining light is absorbed by the support structuresof the grating assemblies.

Figure 3.7: The light path of the two XM M -N ewton telescopes with grating assemblies(not to scale). Note that the actual fraction of the non-intercepted light that passes to theprimary MOS focus is 44%, while 40% of the incident light is intercepted by grating plates.

Mirror Characteristics
The PSF depends on the source off-axis angle (e.g., the PSF at large off-axis angles iselongated due to the off-axis aberration), and slightly depends on the source azimuth within
the FOV. Values for the half energy width (HEW) of the PSFs on-axes are 13.800 , 13.000 ,and 15.200 for the MOS1, MOS2, and pn telescopes, respectively. These values are about ten
times smaller than those of the ASCA telescopes.

3.2. XM M -N EW T ON

25
EPIC: PN
EPIC: MOS (2 modules)
EPIC: MOS (single)

1000

2

Effective Area, Ae [cm ]stRGS-total: -1 ordernd
RGS-total: -2 orderst
RGS1: -1 ordernd
RGS1: -2 order

100

10

1

Energy [keV]

10

Figure 3.8: The on-axis effective areas of the XM M -N ewton telescopes, combined with thequantum efficiency of the focal X-ray instruments, EPIC and RGS.

Figure 3.8 shows both the on-axis and off-axis effective areas of the XM M -N ewton
telescopes. The XMM-Newton mirrors are most efficient in the energy range from 0.1 to10 keV, with a maximum at about 5 keV and a pronounced edge near 2 keV (Au M edge).
The effective area of each MOS camera is smaller than that of the pn camera, because only
part of the incoming radiation falls onto those detectors as described above. The effective
areas decreases with increasing off-axis angle due to the vignetting effect. The off-axis
vignetting in the two telescopes with the grating assemblies (and associated with the MOS
cameras) has a dependence also with the azimuth angle. A source at an off-axis position
perpendicular to the dispersion direction will be vignetted by a different amount to one ata position parallel to the dispersion direction. This dependency is shown in Figure 3.9.
X-ray stray light in the EPIC is produced by rays which are singly reflected by the
mirror hyperbola and which reach the sensitive area of the camera. Thus, and X-ray baffle
was implemented to shadow those singly reflected rays.

3.2.2

EPIC Cameras

EPIC cameras are X-ray imagers, consists of two MOS (Metal Oxide Semi-conductor) CCD
arrays (Turner et al. 2001) and one CCD camera called pn (Strüder et al. 2001). All EPIC
cameras are operated in a photon counting mode with a fixed frame-readout frequency,
producing event lists. This allows for simultaneous imaging and non-dispersive spectroscopy

26

CHAPTER 3. INSTRUMENTS

11500 eV
4500 eV
6400 eV
8000 eV

Vignetting factor

0.8

0.6

0.4

0.2

0

0

3

6
9
Off−axis angle (arcmin)

12

15

Figure 3.9: (left): The vignetting factor as a function of off-axis angle at a few selectedenergies, of the X-ray telescope in front of the pn camera. (right): Vignetting effect as afunction of azimuth angle of the X-ray telescope in front of the MOS1 camera. The curves
are given for an off-axis angle of 10 arcmin. Due to the presence of grating assemblies in the
beams of the X-ray telescopes, the vignetting functions measured in the MOS cameras are
modulated azimuthally.

due to the intrinsic energy resolution of the pixels.

Design Structure of EPIC MOS and pn
Each of the MOS cameras consists of 7 individual front-illuminated chip. The central CCD isat the focal point on the optical axis of the telescope while the outer six are stepped towards
the mirror by 4.5 mm to follow approximately the focal plane curvature, and improve thefocus for off-axis sources. Each chip is equipped with a frame store region. Two MOS
cameras are rotated by 90◦ with respect to each other, in order to compensate for the deadspaces between the chips.
The pn-CCD, developed by MPE, consists of 12 back-illuminated CCD chips. The main
characteristics of the pn-CCD is thick depletion layer of 300 µm. This realizes high quantum
efficiency for the high energy X-ray photons. Each pixel column of pn chips has its own
readout node; 64 channels per chip and a total of 768 channels for the entire camera. This
allows much faster readout of the pn camera than the MOS cameras.
In Table 3.5, we show several design parameters and the performance of the EPIC

3.2. XM M -N EW T ON

27

Figure 3.10: A sketch of the FOV of the two types of EPIC cameras; MOS (left) and pn(right). The shaded circles depict 300 diameter areas.

Table 3.5: Design parameters and performance of the EPIC MOS and pn
Illumination Method
Number of Pixels (1 chip)
Pixel Size
Area (1 chip)
Field of View
Thickness of Depletion Layer
Time Resolution
Bandpass

EPIC-MOS
Front illuminated600 × 600
40 µm (1.00 1) square10.0 9 × 10.0 9∼ 300 diameter
35–40 µm
2.6 s (Full Frame)
0.15–12 keV

EPIC-pn
Back illuminated378 × 64
150 µm (4.00 1) square
13.0 6 × 4.0 4∼ 300 diameter
300 µm73.4 ms (Full Frame)
0.15–15 keVcameras with “full frame” mode, which we used in this thesis, while the EPIC cameras allow
several other modes of data acquisition.

28

3.2.3

CHAPTER 3. INSTRUMENTS

EPIC background

XM M -N ewton takes an highly elliptical orbit with an apogee of about 115,000 km and aperigee of about 6000 km and thus mostly fly outside of the Earth’s magnetosphere. Thisis thought to be responsible for the violent variability of the background of EPIC cameras
as large as two orders of magnitudes. This variation is caused by protons scattered throughthe mirror system.
The background of the EPIC cameras consists of mainly two components: the cosmic
X-ray background and the instrumental background. The latter component is mainly due
to the interactions of charged paticles with the structure surrounding the detectors and thedetectors themselves. This component can be further divided into 2 components: a transientcomponent and quiescent component. Since spatial and temporal characteristics are different
between these components, it is important to study them separately.
The transient component is thought to be induced by soft protons with energies lessthan a few 100 keV. Since the soft protons are scattered through the mirror system, they are
vignetted. The spectra of soft proton flares are variable and no clear correlation has beenfound betweeen intensity and spectral shape. Therefore, observation intervals affected bythis component have to be screened off.
The quiescent component is induced by high energy cosmic rays with energies larger
than some 100 MeV which penetrate the satellite. As the spectrum in Figure 3.11 shows,
the quiescent background consists of a continuum component and lines. While the continumm component is induced by high energy particles directly in the detector, the lines
correspond to the florescent lines induced by cosmic rays interacting the surrounding structures. Whereas the continuum component is spatially flat, the intensities of the lines spatiallychange depending on the detector structure. Figure 3.12 show images of MOS camera inenergy bands centered around Al-K and Si-K emission. The outer 6 of 7 CCDs detect more
Al K radiation due to their closer proximity to the Al camera housing. Si K emission, however, is concentrated along the edges of some CCDs. This is attributed to detection of Si K
X-rays escaping from the back substrate of a neighbouring CCD placed slightly forward andoverlapping the subject chip.

3.3

Chandra

The Chandra X-Ray Observatory, which was successfully put into its orbit in July 1999 withSpace Shuttle Columbia, is an astronomical X-ray satellite managed by NASA and Chandra

3.3. CHAN DRA

29

Cu
Al
Zn
Ni
Cu

PN
Si

MOS1

Figure 3.11: Background spectra for the PN (grey) and MOS1 (black) cameras during observations with the filter wheel in the CLOSED position. Names of elements are shown near
the corresponding prominent K-fluorescence lines.

Figure 3.12: Images of Al-K (left) and Si-K (right) fluorescent emissions from the MOS
camera Lumb, Warwick, Page, & De Luca (2002).

30

CHAPTER 3. INSTRUMENTS

X-ray Center.
Chandra’s main feature is its high spatial resolution of ∼0.500 , which can be ascribed to
the X-ray telescope called High Resolution Mirror Assembly (HRMA). For the focal plane
instrument, we can choose one of two kinds of X-ray detectors: the Advanced CCD Imaging
spectrometer (ACIS) and the High Resolution Camera (HRC). HRC is a multiple-channel
plate type detector optimized for imaging and is without ability of spectroscopy. Since weused only ACIS in this thesis, HRMA and ACIS are described further below.

Figure 3.13: Schematic view of the Chandra satellite.

3.3.1

HRMA

Just as the telescopes of ASCA and XM M -N ewton, HRMA also utlizes total X-ray reflection under the Wolter type-I configuration. HRMA consists of 4 pairs of concentricthin-walled paraboloid and hyperboloid mirrors, with the focal length of ∼ 10 m and thelargest mirror having a diameter of ∼ 1.2 m. Those eight mirrors were fabricated fromZerodur glass, polished, and coated with iridium on a binding layer of chromium. In orderto keep high accuracy of the mirror surface, the glass is very thick and heavy with the total
mass of ∼ 1500 kg.
Figure 3.14 (left) shows the HRMA effective area as a funciton of X-ray energy in log-log
scale, as well as the HRMA/ACIS and HRMA/HRC effective areas. The effective area also
decreases as source position departs from on axis (Figure 3.14 (right))

3.3. CHAN DRA

31

Figure 3.14: (left):On-axis effective areas of HRMA as a function of X-ray energy. (right):Offaxis effective areas of HRMA in several X-ray energies.

The point-spread function is so sharp that most of parallel X-ray photons are focused
into a circle of ∼ 100 radius on a detector. The encircled energy function of HRMA is shownin Figure 3.15. On axis, 50% photons are included in a radius of ∼ 0.00 35 for 4.51 keV X-rays.
The focus blurrs as the distance from the on axis point becomes large (Figure 3.15 (right)).

3.3.2

ACIS

The ACIS array, developed by the Pennsylvania State University, consists of ten CCDs, which
are divided into two parts; ACIS-I with 2×2 CCDs and ACIS-S with 1×6 CCDs. The former
is mainly used for imaging-spectroscopy, whereas the latter are for grating spectroscopy.
ACIS-S1 and ACIS-S3 are back-illuminated CCDs, whereas all the other CCDs are frontilluminated CCDs.
Figure 3.17 (left) shows the energy resolution of ACIS estimated by ground calibrations.
The FI chips have a good resolution near the theoretical limit (∼ 120 eV at 5.9 keV), whereasthat of the BI chips is slightly poor. After the launch, however, the ACIS-I chips weresignificantly by low energy protons encountered during radiation belt passages. The protons
made many charge traps at buried channels, which cause the increase of charge transfer
inefficiency (CTI). Since the gate structure of the BI chips are at opposite side to HRMA,
the energy resolution of the BI chips remains at their pre-launch values. In the present

32

CHAPTER 3. INSTRUMENTS

Figure 3.15: (left):On-axis effective areas of HRMA as a function of X-ray energy. (right):Offaxis effective areas of HRMA in several X-ray energies.

operation, ACIS is not left at the focal position during the radiation belt passages, hence nofurther degradation of CTI is expected. Furthermore, in order to reduce the effect CTI, theCCD temperature is lowered from −90◦ C to −120◦ C.

3.3. CHAN DRA

33

not constant with Z~22 pixels ~11"22 pixels
= 11"

I0

I1

w203c4r
0

w193c2
1x

I2

I3

w158c4r

w215c2r

2

3

S0

S1

S2

w168c4r
4

w140c4r
5

w182c4r
6}

S3+
w134c4r
7

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

S4

S5

w457c48

w201c3r
9}330 pixels = 163"

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

Image Region

BottomPixel (0,0).

column

three

two

one

Top

BI chip indicator

node zero

+Z

.

row

TargetOffsetCoordinates+Z
Pointing+∆Z

Coordinates-Z+Y

Sim Motion

Frame StoreCCD Key

NodeDefinitions

Row/Column
Definition

Coordinate
Orientations

Figure 3.16: A schematic drawing of the ACIS-array configuration.

Figure 3.17: (left): The pre-launch energy resolution of ACIS. (right): The energy resolutionof chip S3 and I3 as a function of row number, after the proton damage on orbit. The data
were taken at the CCD temperature of −120◦ C.

Chapter 4
X-ray Observations of SNRs
Discovered with ROSATThe ROSAT All-Sky Survey (RASS) was mainly performed from 1990 to 1991 during thefirst phase of the satellite lifetime with the Position Sensitive Proportional Counter (PSPC;
Pfeffermann et al. 1987). From the RASS data, SNRs listed in Table 4.1 have been discovered.
From the fact that they escaped radio surveys it is inferred that they are relatively radio
faint. Since the sensitive energy band of the PSPC was restricted to low energy band(< 2 keV) and the energy resolution was low (the PSPC is proportional counters), follow-up
observations with X-ray satellites with wider energy band and higher energy resolution are
required to reveal physical properties of the SNRs. RX J1713.7−3946 and G266.2−1.2 werelater found to emit synchrotron X-ray with ASCA (Koyama et al., 1997; Slane et al., 1999),
and the results were confirmed with higher spatial resolution of Chandra or XM M -N ewton(Uchiyama, Aharonian, & Takahashi, 2003; Lazendic et al., 2004a; Iyudin et al., 2004).
Among the ∼ 250 known SNRs, non-thermal X-ray emission has been detected from
9 SNRs in total. Up to three of them are SNRs discovered with RASS, and the number islarge in comparison with the total number of SNRs discovered in the RASS.
In order to study the characteristics of these SNRs, we analyzed data of both proposed
observations and archival data. Since G266.2−1.2 and G347.3−0.5 are extensively studied
and many papers have been published, we cite the literature. Since G13.3−1.3, G28.8+1.5,
and G32.1−0.9 are large and faint, follow-up observations have not been done much. Afterthe report of the ASCA observations by Yamauchi et al. (1999) and Tomida (1999), otherparts of G156.2−0.5 were observed with ASCA and follow-up observations with XM M N ewton were also performed. We report here results of these observations. G272.2−3.2
35

36

CHAPTER 4. SNRS DISCOVERED WITH ROSAT

and G299.2−2.9 were observed with XM M -N ewton. Using the high spatial resolution andthe good statistics, the physical conditions of the plasmas, existence of non-thermal X-ray
emission, and the reason for the faintness in radio band were studied.
Table 4.1: SNRs discovered with RASS and follow-up observations
Name
G13.3−1.3
G28.8+1.5
G32.1−0.9
G156.2+5.7
G266.2−1.2
G272.2−3.2
G299.2−2.9
G347.3−0.5

Coordinate(R.A., Dec.)J2000

Diameter(arcmin)Radio Flux
(Jy)

ASCA

XMM

Chandra(18h 19m 20s , −18◦ 000 )(18h 39m 00s , −02◦ 550 )(18h 53m 10s , −01◦ 080 )(04h 58m 40s , +51◦ 500 )(08h 52m 00s , −46◦ 200 )(09h 06m 50s , −52◦ 070 )(12h 14m 50s , −65◦ 280 )(17h 13m 21s , −39◦ 460 )

50
100?
40?
110
120
10
3
60

?
?
?
5
50?
0.4
0.5?
??

no obs.
no obs.
[1][2],[3]
[4]
[6]
[7]
[8,9]

–
no obs.
no obs.
this work[5]
this work
this work[10]

no obs.
no obs.
no obs.
no obs.
–no obs.
no obs.
[11,12]

Ref. [1]Folgheraiter, Warwick, Watson, & Koyama (1997) [2]Yamauchi et al. (1999) [3]Tomida (1999)[4]Slane et al. (2001) [5]Iyudin et al. (2004) [6]Harrus et al. (2001) [7]Bai & Wang (2000) [8]Koyama etal. (1997) [9]Slane et al. (1999) [10]Cassam-Chenaı̈ et al. (2004) [11]Uchiyama, Aharonian, & Takahashi(2003) [12]Lazendic et al. (2004a)

4.1

G156.2+5.7

4.1.1

Previous Observations

G156.2+5.7 (RX 04591+5147) was first discovered in the RASS data as a large shell-like SNRof 108’-diameter (Fig. 4.1; Pfeffermann, Aschenbach, & Predehl 1991). The totally circularshape shown by ROSAT suggests that G156.2+5.7 is a remnant of SN exploded in rathersmooth ISM. The radio emission is highly polarized with a shell of 1.9◦ in diameter (Reich,
Fuerst, & Arnal 1992), which indicates the radio emission is non-thermal. With Ginga
observations, hard X-ray emission extended up to 10 keV was detected from the region ofthis SNR (Yamauchi et al. 1993). However, due to the poor spatial resolution of Ginga,
whether this emission comes from the SNR or an overlapping source was not determined.
From the pointing observation of the north and the center regions with ASCA, the hard
X-ray emission was confirmed and located in the north and the center regions (Yamauchi
et al. 1999; Tomida 1999). However, the emission mechanism of the hard X-ray could not
be determined: synchrotron emission, bremsstrahlung, or inverse compton. They also found4.1. G156.2+5.7

37

Figure 4.1: ROSAT gray-scale image of G156.7+5.7 with intensity map of the radio polarizedcomponent (Reich, Fuerst, & Arnal, 1992). The solid lines filling the radio contours designate
the direction of polarization. Broken-line and solid-line circles indicate the fields of view of
the ASCA/GIS observations and the XM M -N ewton observations, respectively.

that the silicon and sulfur abundances in the center region are 3–5 and 3–10 times higher thanthe cosmic value. From the size (1080 diameter), X-ray flux ((1.2–1.6)×10−9 ergs s−1 cm−2 ),and temperature (0.4 keV), they estimated the age, distance, ambient density, and swept-up
mass to be ∼ 1.5 × 104 yr, ∼ 1.3 kpc, ∼ 0.2 cm−3 , and ∼110 M , respectively, assuming
this SNR is in the Sedov phase. Since this age is larger than that of SN1006, the archetypeof SNRs emitting synchrotron X-ray, G156.2+5.7 is important in studying the history ofparticle acceleration in SNRs: how long particle acceleration continues in one SNR.

4.1.2

Mapping Observations with ASCA

Observations and Images
ASCA observed G156.2+5.7 five times in total (Table 4.2). Since the FOV of GIS is larger
than SIS, we concentrate on the GIS results. We screen the GIS data with the standard
criterion mentioned in §3.1.2. The resultant exposure times are shown in Table 4.2.

38

CHAPTER 4. SNRS DISCOVERED WITH ROSAT

Table 4.2: ASCA observations performed on G156.2+5.7.
Obs. Id

Dateyyyy/mm/dd

Pointed Direction(R.A., Dec.)J2000

Expo.
(ksec)

50010000
50010010
50011000
57058000
57059000

1993/09/03
1993/09/25
1993/09/24
1999/03/04
1999/03/05(04h 59m 16.s 2, 51◦ 480 1000 )(04h 59m 17.s 4, 51◦ 480 1200 )(04h 59m 16.s 7, 52◦ 310 1400 )(05h 02m 23.s 7, 51◦ 250 2200 )(04h 55m 33.s 7, 52◦ 040 3800 )

10
15
19
28
32

We generated exposure-corrected GIS images of G156.2+5.7. First, the GIS2 and the
GIS3 images of the five observations were overlaid, and a mosaicked NXB image normalized
by exposure times was subtracted from it. The exposure map was created from the blanksky observations (see Sakano 1999). Next, dividing the NXB-subtracted image with theexposure map, we constructed a exposure-corrected composite image.
The resultant images in the 0.7–1.5 keV, 1.5–2.0 keV, and 2.0–7.0 keV ranges are shownin Figure 4.2 (a 3-color version of these images is shown in Appendix). The 0.7–1.5 keV band
image is rather consistent with the ROSAT image: a shell like structure with enhancement inseveral regions (north, center, and south-east). A fan-like structure with point-like emission
on its axis is seen in the center of both images. As the spectral analysis below shows, the
X-rays in the 1.5–2.0 keV energy band are dominated by those of Si-K line and ROSAT
was insensitive in this energy band. The image of this band show that Si is concentrated in
the center region and suggests the ejecta exists in this region. The 2.0–7.0 keV band image
shows that the hard X-ray emission is enhanced in the center and the north east regions,just as suggested by the Ginga observation (Yamauchi et al. 1993) and the previous analysisof the ASCA data (Yamauchi et al. 1999; Tomida 1999).

Spectral Analysis
In order to locate hard X-ray emission and to reveal spatial distribution of abundances, we
extracted ASCA spectra from several regions shown in Figure 4.2. Although the data of thecenter observations have already been analyzed by Yamauchi et al. (1999) and Tomida (1999),
we performed spectral analysis on the central circular region of 120 radius in order to easily
compare the spectrum with that of the XM M -N ewton observation. A point-like source(dashed-line circle) is excluded from the SW region with 30 radius. For the background,
we used the blank-sky background provided by HEASARC, which are made from blank4.1. G156.2+5.7

39

Figure 4.2: ASCA GIS images of G156.2+5.7 in the energy bands of 0.7–1.5 keV (a), 1.5–
2.0 keV (b), and 2.0–7.0 keV (c). After the subtraction of the NXB, correction for the
exposure was performed for each image. Scales are logarithmic and the numbers next to thescale bars show surface brightness (×10−6 counts s−1 cm−2 arcmin−2 ). The spectral regions
are designated with solid lines.

40

CHAPTER 4. SNRS DISCOVERED WITH ROSAT

sky observations by extracting point sources with flux larger than 1 × 10−13 ergs cm−2 s−1(Ishisaki, 1997). Assuming the CXB is uniform on the sky, we extracted the background
spectra from the same position of the detectors as the source regions.
The background-subtracted spectra are shown in Figure 4.3. Since the spectra of GIS2
and 3 are consistent with each other, both data are combined. As for the center region,
the data of observations 50010000 and 50010010 are also combined after a check for theconsistency. While the spectra of the SE and the NE regions are soft, those of the centerand the NW-shell regions are hard showing X-ray emission extending up to 7 keV or higher.
Si-K lines are prominent in the center and NW regions.
We first fitted these spectra with a NEI model (vnei model in XSPEC v11.3) modified by
the ISM absorption. In order to correct for the systematic energy shift of GIS (Makishima
et al. 1996), we applied 3% gain shift for all the spectra. Since the effective area of GIS
is small in the soft band (< 1 keV), the spectra have weak constraint on the absorption
column densities. We fixed them to the value 3.5 × 1021 cm−2 determined by Yamauchi etal. (1999) with ASCA/SIS, which is consistent with the total galactic H i column density inthis direction (Dickey & Lockman 1990). The NEI model reproduce the spectrum of the SE
region with the solar abundance (Anders & Grevesse 1989), whereas the spectrum of the NW
region left residual at the Si and S lines. Therefore, we freed these abundances. As a result,the spectra of the SE and NW regions could be fitted with the NEI model with the best-fit
parameters and models shown in Table 4.3 and Figure 4.3, respectively. The spectra of theCenter and the NW-shell regions left residuals in the hard band (>3 keV) against the fitted
NEI model. Accordingly, we allocated a power-law model for the hard X-rays. We thenobtained acceptable fit for the spectrum of the NW-shell, whereas wavy residual remained in
the energy band of ∼ 1.5–2.5 keV in the spectrum of the Center region suggesting that theabundances of Mg, Si, and S are different from the solar values. Therefore, we allowed theabundances of these elements vary and as a result, we obtained an acceptable fit. The bestfit parameters and the best-fit model with a NEI+power-law model are shown in Table 4.3
and Figure 4.3, respectively.
The abundances of Si and S are higher than the solar values not only in the Centerregion, but also in the NW region. The overabundance of Si in the NW region is consistent
with the narrow band image (1.5–2.0 keV) in Figure 4.2. We confirmed hard X-ray emission
from the Center region, and newly detected hard X-ray emission from the north-west shell
region of G156.2+5.7.

41

4.1. G156.2+5.7

(a)
SE

(b)
Center(c)
NW(d)
NW shellFigure 4.3: Background-subtracted ASCA spectra of the SE (a), center (b), NW (c), andNW-shell (d) regions. Best-fit models with parameters in Table 4.3 are shown with solid-lines.
As for the spectra of the center (b) and the NW-shell (d) regions, the NEI and power-law
components are shown with broken and dotted lines, respectively.

42

CHAPTER 4. SNRS DISCOVERED WITH ROSAT

Table 4.3: Best-Fit Model Parameters of the ASCA G156.2+5.7 spectra.
Parameter
SE
CenterNW
NW-shell
NEI model:kT [keV] . . . . . . . .
0.53+0.10
0.85+0.27
0.9+0.6
0.60+0.46
−0.07
−0.21
−0.3
−0.15
Mg . . . . . . . . . . . . . 1.0 (fixed)
0.6+0.1
1.0(fixed)1.0(fixed)−0.1+1.2
Si . . . . . . . . . . . . . . 1.0 (fixed)
3.2−0.6
2.7+1.4
1.0 (fixed)
−0.7
+25
+19
S . . . . . . . . . . . . . . . 1.0 (fixed)
10−5
13−8
1.0 (fixed)
+0.4
+0.2
+0.6
−3
log(ne t [cm s])
10.5−0.3
10.5−0.3
10.4−0.2
10.7+1.0
−0.4
+5
+2.9
+3.3
+2.1
a
norm . . . . . . . . . .
21−4
6.3−1.4
4.9−1.0
3.0−1.4
Power-law model:
Γ. . . . . . . . . . . . . . .
–0.8+0.8–
1.7+0.6
−0.6
−0.7
b
−12
flux . . . . . . . . . .
–1.3 × 10–
6.3 × 10−13
NH [1021 cm−2 ] . . 3.5 (fixed) 3.5 (fixed) 3.5 (fixed) 3.5 (fixed)
χ2 /d.o.f. . . . . . . . . .
51.0/64
33.4/50
37.5/32
31.1/34
ne np V /4πD2 (×1011 cm−5 ).
bUnabsorbed flux (ergs cm−2 s−1 ) in the 2.0–10.0 keV.
a

4.1.3

Observations with XM M -N ewton

In order to find out whether the hard X-ray emission detected with ASCA is truly diffuse or
only a collection of faint point sources and to determine the emission mechanism of the hard
X-rays, we performed XM M -N ewton observations on the the center and the north-west
regions.

Data Reduction, and Image
Observed dates and positions are shown in Table 4.4. In all these observations, the EPIC/PN
and MOS cameras were operated in full frame mode with medium filter selected. Since thebackground showed flares during all the observations, we accumulated the background light
curves of the full-field data between 10–12 keV energy (Appendix A), and screened out
the time intervals when the count rate is larger than specific thresholds for each detector.
We set the thresholds at 0.15 count s−1 for the MOS1 and 2 cameras and 0.7 count s−1
for the PN camera. Resultant exposures are shown in Table 4.4. Since the backgroundlight curves of observation 0148850101 show violent variability and only little exposures areleft, we discarded this observation here. We hereafter call observations 0148850201 and
0205970201 as the center and the north-west observations, respectively, after their positions

43

4.1. G156.2+5.7
on G156.2+5.7.
Table 4.4: XM M -N ewton observations performed on G156.2+5.7.
Obs. Id

Dateyyyy/mm/dd

Pointed Direction(R.A., Dec.)J2000

Total Exposure
M1/M2/PN (ks)

Screened Exp.
M1/M2/PN (ks)0148850201
0148850101
0205970201

2002/10/01
2003/02/18
2004/09/20(04h 59m 00s , 51◦ 450 0000 )(04h 55m 55s , 52◦ 260 0000 )(04h 55m 50s , 52◦ 260 0000 )24/24/20
39/39/34
48/48/42

10/10/5.8
7.5/7.3/3.5
29/31/20

Figure 4.4 and 4.5 shows the MOS1+2 images of the north-west and the center observations in the 0.5–2.0 and 2.0–7.0 keV bands. In the soft band images, we see structuresconsistent with the ROSAT and ASCA images: in the north-west image the shell-like structure is conspicuous, whereas the triangle-like structure is seen in the center image. In thehard band, we see neither a clear diffuse structure nor clustered point sources. It is suggested
that the extended hard X-ray emission found with the ASCA observations is not a collectionof point sources but truly diffuse, although the brightness is too small in comparison with
the EPIC background to be clearly recognized in those images.
Resolved Sources
In order to estimate the point-source contribution into the extended hard X-ray emission
detected with ASCA, we first searched for point sources from the combined MOS1+2 images
in the 2–7 keV band using the emldetect software. Since the PN images have large gaps andtherefore the sky coverage is incomplete, we used only the MOS images here. The images
were binned into pixel size of 3.00 1 × 3.00 1. We accepted source detection with likelihood values
above 10 (about 4σ). As a result, 52 and 18 sources are detected from the north-west andthe center observations, respectively (crosses in Figure 4.4 and 4.5 ).
For the bright sources with total counts larger than 100, we performed spectral analysis. The position of these sources are shown in Table 4.5, and designated with numbers
in Figure 4.4 and 4.5. The source and background regions were selected to be circular regions of 2000 radius and surrounding annular regions of 4000 outer radii, respectively, except
for source 4. If other sources are detected in those spectral regions (both the source andbackground), they were excluded with circular regions of 2000 radius. Source 4 is extended,
and is confirmed to be a cluster of galaxies after the spatial and spectral analysis. The detailis discussed in Appendix B. The background-subtracted spectra of sources 1–3 and 5 areshown in Figure 4.6. Since the spectra of sources 1, 2, and 5 are featureless, we fitted them

44

CHAPTER 4. SNRS DISCOVERED WITH ROSAT

Figure 4.4: XM M -N ewton/MOS1+2 images of the north-west region of G156.2+5.7 in the
0.5–2.0 keV (a) and 2.0–7.0 keV (b) bands. The point sources detected with the emldetect
software in the 2.0–7.0 keV image are shown with plus signs. The spectral regions for the
diffuse emission are shown with the solid-line ellipses. The dotted-line circle shows the
120 -radius region in which the diffuse hard X-ray emission detected with ASCA and the
population of the hard point sources are studied.

with power-law models, resulting in acceptable fits with the best-fit parameters shown in
Table 4.5. The photon indices and the absorption column densities are consistent with thatthese sources are background AGNs (Type 1). The spectrum of source 3 shows a concave
structure and suggests Type 2 AGN. When fitted with a power-law model, the absorption
column density became large and a marginal excess remained in the soft band (< 1 keV)a Type 2 AGN. However, since classification of each source is beyond the scope here, weonly derive the flux in the 2.0–10.0 keV range here. We found that the hard X-ray fluxes(2–10 keV) of these relatively bright sources are well below that of the the diffuse X-ray
detected with ASCA.
For sources with less photon counts (< 100), we estimated their fluxes on assumptionof absorbed power-law spectra of photon indices Γ =1.4 and absorption column densities
NH = 3 × 1021 cm−2 . The photon index is selected so as to be consistent with the analysisof the CXB by Moretti et al. (2003). The column density in this direction is so small thatthe flux in the 2.0–7.0 keV is insensitive to the value, but is chosen to be the same as thetotal galactic H i column density (Dickey & Lockman 1990). We estimated the counts-to-flux
conversion factor for the MOS observations to be 3.8 cts s−1 per 10−10 erg cm−2 s−1 using45

4.1. G156.2+5.7

Figure 4.5: XM M -N ewton/MOS1+2 images of the center region of G156.2+5.7 in the 0.5–
2.0 keV (a) and 2.0–7.0 keV (b) bands. The point sources detected with the emldetect
software in the 2.0–7.0 keV image are shown with plus signs. The dotted-line circle shows
the 120 -radius region in which the diffuse hard X-ray emission detected with ASCA and the
population of the hard point sources are studied.

version 3.6a of the pimms software. Using the resultant fluxes, we compared the cumulativeflux distribution (log N –log S) of the 120 -radius regions shown in Figure 4.4 with that of the
CXB. log N –log S is a distribution of the number of all sources brighter than S weighted by
the sky coverage Ωtot :X
1
N (S) =.
(4.1)
Si >S Ωtot (Si )By combining results of multiple surveys, Moretti et al. (2003) derived the log N –log S
distribution of the CXB in the 2–10 keV band to be:(2 × 10−15 )1.57
N (> S) = 5300 1.57
S
+ (4.5 × 10−15 )1.13 S 0.44"#

deg−2 .

(4.2)

where S is measured in erg cm−2 s−1 . Figure 4.7 shows the log N –log S distributions of theG156.2+5.7 regions (circle and square for the center and the NW observations) and the CXB(solid line). We see the log N -log S in the north-west and the center regions show excess
against that of the CXB about a factor of 1.5–2.5. Therefore, the hard excess seen with the
ASCA observation is partly due to the high population of point sources.
We computed flux sum of the point sources down to the detection limits. The detection
limits for the north-west and the center observations are 2.0×10−14 and 3×10−14 erg cm−2 s−1 ,46

CHAPTER 4. SNRS DISCOVERED WITH ROSAT

Table 4.5: Bright Point Sources Detected in the G156.2+5.7 Observations.
No.

position(R.A., Dec.)J2000

Γ

1 ....
2 ....
3 ....
4 ....
5 ....

(04h 55m 20.s 1, +52◦ 280 5000 )(04h 55m 44.s 4, +52◦ 360 1900 )(04h 56m 08.s 4, +52◦ 290 0400 )(04h 56m 37.s 2, +52◦ 240 1200 )(04h 58m 56.s 1, +51◦ 430 3400 )1.4+0.3−0.3
1.69+0.14
−0.13
+1.1
1.7−0.6—b
1.65+0.27−0.22Power-law Parametersa
NHfluxχ2 /d.o.f.

4.6+2.8−1.3
3.1+0.7
−0.5
+60
60−30
—b
3.0+2.0−1.3

8.4/1050.4/50
19.1/15—b
7.7/13a

7.8 × 10−14
3.7 × 10−13
8.4 × 10−14(1.5 × 10−13 )b
1.2 × 10−13

Best-fit parameters of a power-law model [photon index Γ, absorption column density N H(×1021 cm−2 ), and absorbed flux (ergs cm−2 s−1 ) in 2.0–10.0 keV].
bThis source was fitted with a thermal model. See Appendix B.

respectively. The resultant values are compared with the flux sum of the CXB in the same
flux range, and the flux of the extended emission detected with ASCA in Table 4.6. The flux
sum of the point sources is excess against CXB for 3.5×10−12 , and 4.8×10−12 erg cm−2 s−1 deg−2
in the center and the north-west observations, respectively. In comparison with the fluxes of
the hard X-ray emission detected with ASCA, these values are significantly smaller. Therefore, the hard X-ray emission can not be explained with the population excess of point
sources.
Table 4.6: Comparison of the hard X-ray fluxes in the G156.2+5.7 regiona .
Regions
Center. . . . . . .
NW . . . . . . . . .
aASCAresolved sources

CXB

1.0 × 10−11

6.3 × 10−12 b

2.8 × 10−12 c
4.0 × 10−12 d

1.4 × 10−11

8.8 × 10−12 dFluxes in the 2.0–10.0 keV band normalized for the sky region(× erg cm−2 s−1 deg−2 ) are shown.
b
For sources 3.0 × 10−14 < S < 1.2 × 10−13 erg cm−2 s−1 .
c
For sources 3.0 × 10−14 < S < 1.0 × 10−13 erg cm−2 s−1 .
dFor sources 2.0 × 10−14 < S < 1.0 × 10−13 erg cm−2 s−1 .

Spectrum of the Diffuse Emission
For the next we analyze the X-ray spectrum of the diffuse emission directly. However,since the full field of view of the center observation is filled with the X-ray emission from
G156.2+5.7, it is highly difficult to estimate the background, especially the soft proton
component of the non-X-ray background, we avoid of analyzing the diffuse component of

4.1. G156.2+5.7

47

Figure 4.6: EPIC spectra of the point sources detected in the G156.2+5.7 observations. The
combined MOS1 and 2 data are designated with black, whereas the PN data are designatedwith gray. The solid lines represent the best-fit power-law models (Table 4.5).

the center observation here. On the other hand, since the north-west observation covers theouter region of the soft X-ray emission, we can estimate the background from that region. We
extracted the source and background spectra from the elliptical region shown in Figure 4.4,
which were chosen so that both regions are on the same distance from the nominal point.
The 2000 regions around point sources detected in the previous section were excluded fromboth regions. The extended source, source 4, was excluded with a circular region of 100 00
radius.
The MOS1+2 spectra of the source and background regions are shown in Figure 4.8. Although we expect hard X-ray emission in the source region, no excess against the backgroundregion can be seen. We will discuss the hard X-ray emission, in the following. Figure ??
shows the background-subtracted spectra. We fitted this spectrum with a non-equilibrium
ionization plasma (vnei model in XSPEC 11.3; Borkowski et al. 2001a). As a result, themodel reproduced the spectrum quite well, with the best-fit parameters shown in Table 4.7.

48

CHAPTER 4. SNRS DISCOVERED WITH ROSAT

Figure 4.7: Cumulative source counts N (S) in the 2–10 keV within the 120 -radius regionsof the G156.2+5.7 observations. The data-point of the center and the NW observations aredesignated with the circles and crosses, respectively. The solid line refers to the CXB counts
by Moretti et al. (2003).

All the abundances are consistent with the solar values.
Table 4.7: Best-fit Spectral parameters for the NW region with a NEI model.
kT (keV)
0.29+0.03
−0.03
a
b

O
0.5+0.7−0.2

Ne
0.8+0.9−0.3

Mg
0.6+0.6−0.2

Si
1.0+1.1−0.5

Fe
0.6+0.4−0.2

log(ne t)
11.3+0.3
−0.2

norma
1.9+1.1−0.7

NHb
5.3+0.7−0.5

χ2 /d.o.f.
134/97

ne np V /4πD2 (×1012 cm−5 ).
Absorption column density (×1021 cm−2 ).

According to the ASCA result, hard X-ray emission is expected in the source region,
but we found no excess against the background region. In Figure 4.10, we compare thespectral regions of the ASCA analysis and XM M -N ewton analysis. The background regionis outside of the thermal X-ray emission, wheras on the radio shell. Therefore, the hard
X-ray emission may be extending out to the radio shell beyond the thermal X-ray shell.

4.1.4

Discussion on G156.2+5.7

We newly detected hard X-ray emission from the NW region with ASCA. With spatiallyresolved observations with XM M -N ewton, the hard X-ray emission is found to be not due

4.2. G272.2−3.2

Figure 4.8: MOS1+2 spectra of the source(black) and background (gray) regions in the
NW observation.

49

Figure 4.9: MOS1+2 spectra of the NW region. The best-fit NEI model is shown withthe solid-line.

to collection of point sources but truly diffuse.
The photon index of the hard X-ray emission (Γ ∼ 1.5) is small in comparison with
SN 1006, the established SNR emitting synchrotron X-ray (Γ ∼ 2.2; Bamba et al. 2003b).
If we assume the X-ray emission is inverse compton of the cosmic micro-wave background,the fluxes and the photon (energy) indices of the radio and X-ray emission can be explained
simultaneously with the magnetic field strength of 0.3 µG.

4.2

G272.2−3.2

4.2.1

Previous Observations

G272.2−3.2 was discovered by Greiner, Egger, & Aschenbach (1994) in the RASS. It presents
a centrally filled X-ray morphology and a thermally dominated X-ray spectrum. The shapeis nearly circular with a diameter of about 15 arcmin. In the optical band, Winkler, Hanson,& Phillips (1993) measured the [SII ]/Hα ratio and detected emission from [NII ] 658.3 nm
and [OII ] 732.5 nm. Both of these characteristics are typical of SNRs and confirmed that the
nebulosity is a SNR. From the ROSAT and ASCA observations, Harrus et al. (2001) foundthat the X-ray emission is described by an NEI model with a temperature of ∼0.7 keV, anionization time-scale of the plasma of 3200 cm−3 yr, and a relatively high column density(NH ∼ 1022 cm−2 ). From the thermal emission and the morphology, they interpreted this

50

CHAPTER 4. SNRS DISCOVERED WITH ROSAT]

Figure 4.10: Comparison of spectral regions of G156.2+6.7 north-west. The ASCA NWshell region (dotted-line ellipse), the source and background regions for the XM M -N ewton
observations (solid-line ellipses) are shown on the MOS1+2 image in the 0.5–2.0 keV band.
The intensity contours of the polarized component of the radio emission (2695 MHz; Reich,
Fuerst, & Arnal 1992) is overlaid.

remnant as a “thermal composite” SNR. In order to explain the centrally peaked X-ray
morphology, they introduced cloud evaporation and thermal conduction models. Mainly
because of the uncertainty in the distance toward the SNR, they could not conclude whichmodel is the case.

4.2.2

Observation, Data Reduction and Image

G272.2−3.2 was observed with XM M -N ewton on December 10, 2001. Both the PN, MOS1
and MOS1 cameras were operated in the standard full-frame mode with the thin filters.
Since the coverage of the PN camera is not complete due to the large gap between CCD
chips and statistics of the combined MOS1 and 2 data was enough for our study, we used
the PN data only for consistency check. We report only the result of the MOS1+2 data.
Since the background showed flares during the observations we accumulated the backgroundlight curves of the full-field data between 10–12 keV energy (Appendix A), and screened out

51

4.2. G272.2−3.2

the time intervals when the count rate is larger than 0.15 count s−1 for the MOS1 and 2
cameras, and 0.7 count s−1 for the PN camera. The resultant exposure times are shown in
Table 4.8.
Table 4.8: XM M -N ewton observations performed on G272.2−3.2.
Obs. Id

Dateyyyy/mm/dd
0112930101 2001/12/10

Pointed Direction(R.A., Dec.)J2000h(09 06m 46.0s , −52◦ 070 1200 )Total Exposure
M1/M2/PN (ks)
37/37/28

Screened Exp.
M1/M2/PN (ks)
17/17/11

The MOS1+2 images in several energy bands are shown in Figure 4.11. In order to see
spatial distribution of Si and S line emission, narrow-band images ((b) and (c)) were also
created. The X-ray emission in the soft-band (0.7–2.0 keV) is rather smooth, supporting
the ROSAT result. The emission in the 2.3–2.6 keV band is enhanced in the inner partsuggesting that S is concentrated toward the center. No conspicuous emission is seen in thehard-band (3.0–7.0 keV).

4.2.3

Spectral Analysis

X-ray spectra were extracted from the three regions shown in Figure 4.11: the Inner region,
the Outer region, and the Region A. The “A” region, where the X-ray emission is enhanced in
the west part, corresponds to the region “A” in Harrus et al. (2001). The background for theregion A was extracted from the elliptical region (dotted-line ellipse) shown in Figure 4.11.
Since the vignetting effect is substantial and the center part of the FOV is occupied by
G272.2−3.2 itself, the backgrounds for the Inner and Outer regions were estimated separately
for the non-X-ray background and the CXB. The non-X-ray background was evaluated using
the CLOSED data, whereas the CXB was derived from literature. According to Ishisaki(1997), the spectrum of the cosmic X-ray background above 1.2 keV is consistent with asingle power-law with photon index 1.49±0.06 and the brightness of 10−8 ergs s−1 cm−2 str−1
in the 2–10 keV.
The background-subtracted spectra are shown in Figure 4.12. Since all the spectra show
line emissions and suggested emissions from thin thermal plasmas, we fitted those with a
NEI model (vnei model in XSPEC v11.2). As for the spectra of the Outer and Inner regions,power-law models with photon index 1.49 and Gaussians at energies 1.487 and 1.740 keVwere added to represent the CXB, and the cosmic-ray induced instrumental lines (Al and Silines), respectively. While the normalizations of the power-law models (the CXB) were fixed

52

CHAPTER 4. SNRS DISCOVERED WITH ROSAT

Figure 4.11: XM M -N ewton MOS1+2 images of G272.2-3.2 in the energy bands of 0.7–
1.7 keV (a), 1.7–2.0 keV (b), 2.3–2.6 keV (c), and 3.0–7.0 keV (d). Spectral region are designated with circles and ellipses. Scales are linear. Coordinates are R.A. and Dec. (J2000).

53

4.2. G272.2−3.2

(a)
(Al)

(Al)

(Si)

CXB

(b)
(Si)

CXB

Ar

Ar

(c)

Figure 4.12: Background subtracted spectra of 3 regions of G272.2-3.2. (a): the Outerregion. (b): the Inner region. (c): the region A.

to the value determined from the size of the sky region, the fluxes of the instrumental lines
were fitted.

4.2.4

Discussion on G272.2−3.2

We found that the abundances in the inner part are higher than the solar values, while thosein the outer part are lower than the solar. This suggests that the inner part is dominatedby the ejecta, and that the outer part is dominated by the swept-up ISM. Therefore, it isinferred that this SNR is in the Sedov phase. It is notable that the abundance of the outerregion which corresponds to that of the ambient ISM around G272.−3.2 is substantiallylower than the solar value. Low abundances are reported on several SNRs in our Galaxy, forexample the Cygnus Loop (Miyata et al. 1994).
The temperature of the ejecta is higher than that of the forward shock, which is also the

54

CHAPTER 4. SNRS DISCOVERED WITH ROSAT

Table 4.9: Best-Fit Model Parameters for G272.2−3.2.
Parameter
Outer Region Inner Region
Region A
NEI model:kT [keV] . . . . . . . .
0.51+0.05
0.83+0.04
0.44+0.12
−0.05
−0.08
−0.09
+0.10
+0.5
O ..............
0.20−0.09
1.5−0.7
0.088(< 0.45)
+0.12
+0.3
Ne . . . . . . . . . . . . .
0.34−0.07
1.0−0.2
0.17+0.25
−0.11
+0.06
+0.1
Mg . . . . . . . . . . . . .
0.28−0.03
0.9−0.3
0.20+0.14
−0.08
+0.05
+0.4
+1.4
Si . . . . . . . . . . . . . .
0.50−0.05
2.3−0.4
0.4−0.1+1.0+1.3
S. . . . . . . . . . . .. . .
1.7−0.4
5.5−1.3
1.0 (fixed)
+0.06
+0.3
Fe . . . . . . . . . . . . . .
0.22−0.04
1.9−0.5
0.13+0.14
−0.07
+0.14
+0.05
−3
log(ne t [cm s])
10.75−0.12
10.63−0.08
10.8+0.5
−0.5
+3.8
+0.08
norma . . . . . . . . . .
0.50+0.13
8.3
0.09
−0.13
−1.5
−0.06
+2.9
b
Ar-line flux . . . . . .
2.0(< 5.3)
8.2−3.1...
Additional Line:center . . . . . . . . . .
1.225+0.015
1.235+0.010...
−0.011−0.005+0.4+0.6
fluxc . . . . . . . . . . .
1.3−0.3
2.0−0.4
...
+0.7
+0.3
21
−2
NH [10 cm ] . .
9.4−0.5
10.7−0.3
7.9+2.1−2.4
χ2 /d.o.f. . . . . . . . . .
183/182308/200
44.6/40
ne np V /4πD2 (×1012 cm−5 ).
bFlux of Ar-line (×10−6 cm−2 s−1 ). The center energy is fixed to 3.14 keV.
c×10−4 cm−2 s−1 .
acase for G349.7+0.2, an SNR expanding into a dense medium (Lazendic et al. 2004b). Since
all the abundances of region A are consistent with those of the Outer region (the forward
shock), the X-ray emission of Region A is enhanced not by ejecta but because the density of
the ISM is locally high.

4.3

G299.2−2.9

4.3.1

Previous Observations

G299.2−2.9 was also discovered in the RASS data (Busser, Egger, & Aschenbach 1996). Thisobject has a nearly circular shape with a diameter of ∼ 15 arcmin. The emission from thisregion was identified as two distinct sources in the Einstein Slew Survey: 1ES 1212−651 and1ES 1212−652 (Slane, Vancura, & Hughes 1996). ESO/SRC red survey plates (ESO Schmidt

55

4.3. G299.2−2.9

plate R 2978) show faint filaments around the southern rim of G299.2−2.9, with emission in
the western and central regions as well. A ring emission is also evident in the IRAS data.
The infrared emission is presumably associated with shock heating of the surrounding dust.
Two distinct radio sources coincident with the extended regions of G299.2−2.9 are listed
in the PMN 4.85 GHz radio survey catalog, with a combined flux of ∼ 180 mJy. Slane,
Vancura, & Hughes (1996) proposed two scenarios for this SNR: a nearby young remnant
and a distant Sedov-phase remnant.

4.3.2

XM M -N ewton Observations

G299.2−2.9 was observed with XM M -N ewton twice (Table 4.10). However, one of theseobservations (Obs. Id:0112890201) suffered from high background radiation due to the soft
protons for all the exposure time. Therefore, we concentrate on the other one (Obs. ID:
0112890101). The MOS1 and MOS2 cameras were operated in the standard full-frame modewith the medium filters. High background intervals were screened out in the same manneras the analysis of G156.2+5.7 using the light curve of 10–12 keV events (see Appendix C?).
Here, the threshold count rates were set to be 0.12, 0.12 and 0.5 count s−1 for MOS1, 2 andPN, respectively. The resultant exposures are shown in Table 4.10.
Table 4.10: XM M -N ewton observations performed on G299.2−2.9.
Obs. Id

Dateyyyy/mm/dd
0112890101 2002/02/24
0112890201 2002/08/19

Pointed Direction(R.A., Dec.)J2000h(12 15m 20.s 6, −65◦ 300 2300 )(12h 15m 10.s 4, −65◦ 300 0500 )Total Exposure
M1/M2/PN (ks)
13/13/0.41
15/15/9.5

Screened Exp.
M1/M2/PN (ks)8.8/7.8/0.41
0/0/0

The MOS1+2 images in several energy bands are shown in Figure 4.13. In the soft-band
image (0.5–2.0 keV), we see a ring-like structure, which reminds us of a emission from theforward shock. The shape of the ring is not circular but rather elliptical with major and minor
axes of 60 and 4.0 5, which suggests asymmetrical explosion or non-uniform ambient ISM. The
inner region of the ring-like structure is filled with relatively faint emission, which might bea reverse-shock component or a projected component of the forward-shock emission. Thereis a enhanced emission on the north-east part of the ring structure, which we call “BrightSpot”, hereafter. There is also an emission outside of the ring in the north-east. In the Si
and S bands, differences of the brightnesses between the Inner and Outer regions are lessthan in the soft band, which infer us the abundances of Si and S are larger in the Inner

56

CHAPTER 4. SNRS DISCOVERED WITH ROSATRegion. In the hard band, no emission is significant over the background.
X-ray spectra were extracted from the 4 regions shown in Figure 4.13: Inner Region,
Outer Region, Bright Spot, and NE Edge. From the Outer Region, Bright spot and a hardsource designated with a dashed-line circle was excluded. Background was extracted froma nearby source-free region on approximately the same distance from the nominal point as
the source regions (solid-line ellipse in Figure 4.13). The background-subtracted spectra areshown in figure 4.14.
Table 4.11: Best-Fit Model Parameters for G299.2−2.9.
Parameter
Inner Region Outer Region Bright Spot NE Edge
NEI model:kT [keV] . . . . . . . .
0.46+0.01
0.44+0.03
0.44+0.07
0.42+0.08
−0.01
−0.01
−0.08
−0.08
+0.02
+0.03
+0.07
N, O . . . . . . . . . . .
0.14−0.01
0.14−0.01
0.15−0.08
0.14(fixed)
+0.04
+0.03
+0.17
Ne . . . . . . . . . . . . .
0.40−0.03
0.34−0.03
0.42−0.11
0.34(fixed)
+0.08
+0.03
+0.19
Mg . . . . . . . . . . . . .
0.33−0.03
0.30−0.03
0.49−0.12
0.30(fixed)
Si . . . . . . . . . . . . . .
1.9+0.2
0.61+0.05
0.6+0.3
1.4+0.6
−0.2
−0.07
−0.2
−0.5
+0.8
+0.7
+2.5
S. . . . . . . . . . . . . ..
4.8−0.8
2.1−0.3
2.3−1.7
1.0 (fixed)
+0.06
+0.03
+0.08
Fe . . . . . . . . . . . . . .
0.56−0.04
0.26−0.02
0.32−0.09
0.26(fixed)+0.36
log(ne t [cm−3 s])
11.22+0.03
10.98+0.04
11.01−0.08
10.8+0.8
−0.05
−0.07
−0.2
+0.09
+0.7
+0.2
+0.30
a
norm . . . . . . . . . .
1.11−0.03
5.4−0.8
0.7−0.3
0.33−0.12
Additional Line:center . . . . . . . . . .
1.23+0.01
1.231+0.017
1.23 (fixed)
...
−0.02
−0.006
+2.2
+0.2
+0.10
b
flux . . . . . . . . . . .
7.9−2.0
1.4−0.4
0.15−0.10
...
+0.3
+0.1
+0.8
21
−2
NH [10 cm ] . .
3.2−0.2
3.1−0.2
2.9−0.8
2.9+0.7−0.7
2
χ /d.o.f. . . . . . . . . .
187/148240/191
125/9774.1/75ab

4.3.3

ne np dV (×1012 cm−5 ).
Photon flux of the line (×10−4 cm−2 s−1 ).

R

Discussion on G299.2−2.9

The X-ray emission from the Outer region can be ascribed to the shock-heated swept-up
interstellar matter. The plasma is in non-equilibrium ionization state (ne t ∼ 1×1011 cm−3 s).
We found this SNR is in the Sedov phase of evolution. The absorption is almost consistent
with the ASCA value ((4.0 − 4.7) × 1021 cm−2 ; Bai & Wang 2000).
The abundances of elements lighter than Mg in the Inner and Outer regions are not4.3. G299.2−2.9

57

Figure 4.13: XM M -N ewton MOS1+2 images of G299.2−2.9 in the energy bands of 0.5–
1.7 keV (a), 1.7–2.0 keV (b), 2.3–2.6 keV (c), and 2.0–7.0 keV (d). Scales are linear. Coordinates are R.A. and Dec. (J2000).

58

CHAPTER 4. SNRS DISCOVERED WITH ROSAT

(a)

(b)

(c)

(d)

Figure 4.14: Background subtracted spectra of 4 regions of G299.2-2.9.

4.3. G299.2−2.959different within the 90% error regions. It suggests that the production of light elements areless in the supernova explosion. Therefore, it is inferred that the type of the SN was type Ia.

Chapter 5
Search for New SNRs with the ASCA
Galactic Plane Survey
5.1

ASCA Galactic Plane Survey

The ASCA Galactic Plane Survey was performed from 1996 to 1999 (Yamauchi et al. 2002).
This survey is a collaborative project of the ASCA team, aimed at the systematic study of
the Galactic X-ray sources. It was planned to cover the area of the Galactic inner disk of|l| ≤ 45◦ and |b| ≤ 0.◦ 4 with successive pointing observations with ∼ 10 ks exposure time for
each. Comprehensive study of the survey is reported by Sugizaki (1999) and Sugizaki et al.
(2001).
Soft X-rays are absorbed easily by the interstellar medium (ISM). For example, if we
assume averaged hydrogen density of the ISM is 1 cm−3 , the mean free-path of a 1 keV
X-ray photon is only ∼2 kpc. Therefore hard X-ray observations are essential for finding
new X-ray sources embedded in the galactic plane. Moreover, since the Galactic Plane isbright diffuse X-ray source (the Galactic Ridge emission; e.g., Koyama et al. 1986; Kaneda
et al. 1997), the imaging capability is required to reduce background and therefore essential
for high sensitivity. From these reasons, the ASCA survey, which was the first survey withimaging capability in the energy band > 3keV, is suited for finding faint X-ray sources in thegalactic plane. As a matter of fact, X-ray emissions from about 20 SNRs are newly detectedin this ASCA survey (Yamauchi et al. 2002).
61

62

5.2

CHAPTER 5. ASCA SEARCH FOR SNRSNew Extended Hard X-ray Sources

Although SNRs are the primary candidates for the acceleration sites of the galactic cosmic
rays, what fraction of the flux showering onto the earth can be explained by SNRs is yet
unknown. In order to resolve this issue, we have to proceed with mainly two studies: acceleration efficiency in each SNR and the number of SNRs in our Galaxy. For the former study,
the number of SNRs exhibiting evidence of acceleration is yet too small. Therefore, more
SNRs which show evidence of cosmic ray acceleration are required for the systematic study.
The most sensitive information showing particle acceleration in SNRs is non-thermal X-rayssuch as synchrotron or bremsstrahlung. Since these X-ray emissions tend to extend up-tohigh energy than thermal X-rays, a survey in the hard X-ray band is suited. The study of thenumber of SNRs in our Galaxy, the latter study, should be performed with an X-ray survey.
SNRs have been extensively searched for with radio observations. However, the Galactic
plane is a strong radio emitter and the extended emission makes high background, and as aresult, the sensitivity becomes lower in the regions near the galactic plane. As shown by the
RASS, an X-ray survey should be a efficient way for discovering new SNRs.
Thus we searched for unidentified extended X-ray sources in the hard band images of
the ASCA Galactic Plane Survey. Sugizaki et al. (2001) also searched for discrete sources,
however, their survey was mainly sensitive for sources with extension smaller than the PSF
of ASCA (FWHM ∼ 10 ) and they have not yet determined each source is extended or not.
Therefore, many extended X-ray sources might remain to be found in the data set of the
ASCA survey.
There are mainly 4 possibilities for extended X-ray objects located in the directionof the Galactic plane: SNRs, pulsar-wind nebulae, star-forming regions, and clusters ofgalaxies. Among the numbers of X-ray detected sources at this moment, that of SNRs is thelargest. Therefore, SNRs are naively the most probable objects for unidentified extended
X-ray sources. Although study of clusters of galaxies located behind the Galactic plane,
especially toward the inner Galaxy, has been almost impossible due to the large absorption,
the expected number can be estimated from the high latitude observations.

5.2.1

Source Finding

We used 173 pointing observations of the ASCA Galactic Plane Survey. Since about 34% of
the surveyed regions are suffered from stray lights from nearby bright sources, we excluded
these regions. The regions analyzed are the same as Sugizaki et al. (2001) (Figure 1) except

5.2. NEW EXTENDED HARD X-RAY SOURCES

63

that the Galactic ridge observations (from l = 57.◦ 5 to l = 62.◦ 5) are not included in thisthesis. We screened these data with the criterion shown in §3.1.2. To gain the statisticalaccuracy, we summed the data of GIS-2 and GIS-3, hereafter. Exposure corrected imagesof the AGPS data in the 0.7–2.0 and 2.0–7.0 keV bands were produced with the following
procedures (see Figure 5.1):(1) First, in order to make each image containing only cosmic X-ray events, non-X-ray
background (NXB) is subtracted from the raw images. We used the night earth observations from 1993 to 1994 provided by HEASARC for estimating the NXB. The night
earth observations are screened with the same criterion as the survey data (§3.1.2)
except that ELV< −5◦ . Then we normalized the image of the night earth observation
for the exposure time, and subtracted it from each raw image.
(2) An exposure map for each observation is made from blank-sky observations with brightsources excluded, which Dr. Ishisaki kindly supplied. If we properly subtract NXB fromthe blank-sky observations, the data contain only the cosmic X-ray background. Since
the cosmic X-ray background (CXB) can be thought to be flat, each pixel of the image
shows efficiency map of the detector. We normalized the NXB subtracted blank-sky
image for the exposure time and regard it as an exposure map. According to Ishisaki(1997), the brightnesses of the the CXB in the 0.7–2.0 keV and 2.0–7.0 keV bands
are 8.5×10−7 and 6.0×10−7 counts cm−2 s−1 arcmin−2 , respectively. (These values
are calculated from the energy flux in Ishisaki (1997) using the PIMMS software, on
assumption that the spectrum of the CXB is an absorbed power-law with photon index
Γ = 1.5 and NH = 1 × 1021 cm−2 .)(3) Next, we mosaicked both the NXB subtracted images and the exposure maps by castingthem onto the sky coordinate and by superposing them. Since the data of the survey
is very large (long), the surveyed area is divided into 1.◦ 5 long along the Galactic
plane and analyzed. In order to avoid dividing extended structures with borders, wemade margins of 0.◦ 25 long on each end of the images and as a result, the size of eachmosaicked image is ∼ 2.◦ 0 × 0.◦ 8.
(4) The mosaicked images are divided by the exposure maps. Here, the resultant images are
exposure corrected images, with a number in each pixel designating surface brightness(counts cm−2 s−1 arcmin−2 ) at that position.
We searched for extended sources from the exposure-corrected images in the 2.0–7.0 keVband. In order to avoid mistaking bright sources for diffuse sources, we first excluded sources
with flux larger than 4 × 10−12 ergs cm−2 s−1 (0.7−10 keV) in Sugizaki et al. (2001). Next,
in order to distinguish significant excess in these images, we made distributions of brightness

64

CHAPTER 5. ASCA SEARCH FOR SNRS

in each pixel and made images showing only pixels with brightness 1σ larger than the mode.
Then, extended structures are searched for from those images. When a picked-up extended
source has already been identified as a SNR or some other kind of objects by Sugizaki et al.
(2001), it is excluded here. Images and names of extended X-ray sources discovered in thisway are shown in Figure 5.2 and Table 5.1, respectively. Each source is designated with thegalactic coordinate, such as G22.0+0.0.

5.2. NEW EXTENDED HARD X-RAY SOURCES

65

Figure 5.1: Flow chart of how to create an exposure-corrected mosaic image. The numbersin the parentheses correspond to the procedures in the text.

66

Source Name

χ2ν /ν d
36.8/46
52.1/46
40.5/56
47.3/46
67.7/46
50.4/46—46.5/46
34.4/46
52.2/46
31.2/46
51.9/4650.2/46
16.0/2134.3/46

Best-fit parameters of a power-law model [photon index Γ, absorption column density N H (×1022 cm−2 ), and absorbed flux
(×10−12 ergs cm−2 s−1 ) in 2.0–10.0 keV] are shown.
b
The fitting results of the radial profiles with a PSF + constant. Since the fitting was rejected for all the sources, only χ 2 is shown.
cBest-fit parameters of the radial-profile with a combined model of a PSF + a Gaussian + constant [Flux of the central point source(×10−5 counts cm−2 s−1 ), σ and normalization (counts cm−2 −1 arcmin−2 ) of the Gaussian, the brightness of the constant component(×10−6 counts cm−2 s−1 arcmin−2 ), and reduced χ2 ] are shown.
d 2
χ /degree of freedom (ν).
a

CHAPTER 5. ASCA SEARCH FOR SNRS

G11.0+0.0
G22.0+0.0
G23.5+0.1
G25.0+0.0
G25.5+0.0
G26.0+0.0
G26.6−0.1
G28.6−0.1
G32.45+0.1
G35.5+0.0
G36.0+0.0
G37.0−0.1
G316.8−0.1
G318.6+0.0
G337.1+0.0

size
(0 )
15
14
16×12
11
12
16
14
11
10
14×10
17
7
7
18×12
15

Table 5.1: Parameters of newly found extended hard X-ray sources.
PL Parametersa
PSFb
PSF + Gaussc
ΓNH
Flux χ2ν /ν d
χ2ν /ν dflux
σ (0 )
norm
const.
+0.4
+0.6
+9
+1.1
+0.5
+0.08
1.7−0.4 0.8−0.5 4.0 31.2/33
173/4822−8 6.2−0.9 2.2−0.4 2.58−0.10
+0.7
+1.8
+0.09
0<1.3
2.2 17.3/18
153/48
5<10 6.4−1.5
1.7+0.6
2.25−0.11
1.0−0.3
−0.4
+1.1
+1.3
+0.12
6+4
2.5
12.3/9103/483<10 6.6−1.0
2.5+0.8
2.00−0.11
2.5−0.8
−3
−0.6
+1.0
+1.6
+1.4
+0.06
3.1−1.1
2.2
22.9/887.6/48
2<5 4.6−1.0
1.3+0.6
2.44−0.07
2.9−0.7
−0.4
+0.8
+1.3
+1.5
+0.08
1.9−0.9
1.8 12.5/17
211/480<3 6.4−0.9
1.9+0.3
2.45−0.11
1.9−0.6
−0.3
+1.3
+2.4
+0.1
5+6
2.7
4.0/10187/480.9<5 9.7−1.5
1.6+0.2
2.2−0.3
2.2−0.9
−3
−0.2
+0.6
+0.9
0.1−0.1
2.8 20.1/12
—
—
—
—
—
1.2−0.3
+1.0
+2.1
+0.9
+0.7
+0.07
<6
136/481
4.4−0.7 2.3−0.6 2.22−0.07
2.1−0.8 2.4−1.0 2.2 10.0/11
+1.8
+0.8
+0.07
7+9
2.0
2.9/8
128/482<8 3.9−0.6
3.0+0.9
2.26−0.07
2.6−1.2
−4
−0.8
+0.2
+2.0
+0.07
0<0.2
2.3 19.8/14
123/4812+5
6.5−1.5
1.2+0.4
1.60−0.09
1.2−0.2
−5
−0.3
+0.7
+2.4
+3.4
+0.08
2.7−1.5
2.6 25.2/14
97.0/48
7+4
8.3−1.7
0.9+0.2
1.54−0.14
1.3−0.4
−3
−0.2
+1.1
+0.07
7+7
2.2
8.2/5
90.3/48
26+9
5+4
1.2+1.7
1.45−0.10
2.0−0.9
−4
−12
−2
−0.6
+1.0
+2.3
+0.7
+0.05
3.1−1.6
1.3
1.6/5
76.2/48
7<19 1.6−0.1
6+4
1.38−0.04
2.4−0.8
−3
+1.6
+1.6
+0.05
10+13
2.5 19.3/12
71.9/23
1<5 4.7−1.0
1.2+0.5
1.35−0.07
2.2−1.1
−6
−0.4
+0.9
+1.8
+0.12
7+6
3.9 30.8/21
183/4821+5
7.9−1.2
1.9+0.3
2.25−0.16
1.8−0.6−3−7−0.3

5.2. NEW EXTENDED HARD X-RAY SOURCES

5.2.2

67

Spatial and Spectral Characteristics of Newly Found Sources

We tested whether these newly found sources are significantly extended or not by making
radial profiles. The NXB was subtracted from the radial profiles and the vignetting andexposure were corrected, utilizing the exposure map made in the above.
The center of the radial profile of each source was determined as follows: (1) If there isa (point) source listed by Sugizaki et al. (2001) inside the candidate, the center of the profile
was determined to its position. (2) If multiple sources by Sugizaki et al. (2001) are locatedinside, the center was determined to the position of the brightest one. (3) If no source is
detected inside, the center of the profile was manually determined to be approximately thecenter of the emission. Before binning the images, regions of nearby sources listed in Sugizaki
et al. (2001) are excluded with circular regions of 30 or 60 radii, depending on the flux inthe 0.7–10 keV range (60 for flux > 5 × 10−12 ergs cm−2 s−1 , and 30 for else; thick solid-line
circles in Figure 5.2).
The resultant radial profiles are shown in Figure 5.3. We see the profiles are flat at
the large radii and therefore can confirm that the correction for the vignetting and exposure
is properly performed. In order to determine whether these profiles can be explained witha point source or not, we first fitted these profiles with a point-spread function (PSF) +constant. The PSF was evaluated with the observations of Cyg X-1 (Takahashi et al. 1995).
We made radial profiles of Cyg X-1 observed on several positions of GIS and fitted those with
3 gaussians. As a result, we found that best-fit 3 gaussians for pointing “P3” with GIS 3 canreproduce any PSF in the FOV better than 20% for radius between 00 and 20 . Thus we used
the best-fit 3 gaussians as the PSF. If the fitting was rejected with a 90% confidence level,
we fitted the profile adding a Gaussian for the next. The results are tabulated in Table 5.1.
We found no radial profiles of the candidates can be explained with a point source (a PSF +constant), and thus we can conclude that all the 15 sources listed here are extended sources.
X-ray spectra of the newly found sources are extracted with the source and backgroundregions shown in Figure 5.2 (thin solid lines and thin broken lines, respectively). The point
sources located inside the source regions are included into the spectra, whereas those locatedinside the background regions are excluded with circular regions of 2’ radii.
The background-subtracted spectra are shown in Figure 5.4. In order to estimate thehardness and absorption of each source, we first fitted each spectrum simply with a powerlaw model. The resultant best-fit parameters are shown in Table 5.1. For several sources,
the power-law model was rejected with a 90% confidence level. Fittings with other models
are discussed in §5.3.

68

CHAPTER 5. ASCA SEARCH FOR SNRS

10150
9125

+11°40′G11.0+0.0

7

7 75
50

6
+11°20′+22°20′525+0°20′

5 04+0°20′+22°00′-25
4-50

G22.0+0.0

8

8100

6

10
9

+11°00′3-0°20′ +10°40′-75+0°00′+21°40′-0°20′

75 100 125 150 175 200 225 250 275
10 IMSMO(Gauss) / Mosaic image: exposure NOT corrected
10
BLOW
G23.5+0.1
9225
9
8200

7

6150

6

+0°20′+0°20′

4+25°00′50
25+25°20′5+23°40′

100
4
75
3

G25.0+0.0

8+24°00′

7
175

5125

+0°00′+23°20′-0°20′+0°00′+23°00′3-0°40′

25 50 75 100 125 150 175 200 225
BLOW IMSMO(Gauss) / Mosaic image: exposure NOT corrected+24°40′ -0°20′+0°00′

Figure 5.2: ASCA GIS images of newly found extended X-ray sources in the 2.0–7.0 keVband. See text about the production of these images. Scales are logarithmic and the numbers
next to the scale bars show surface brightness (×10−6 counts cm−2 s−1 arcmin−2 ). The
orientation is equatorial (J2000), whereas the coordinates shown are galactic. The sources
detected in Sugizaki et al. (2001) are designated with white crosses. The circular regions
excluded from the radial profiles are shown with thick solid-line circles (small, and large ones
are with radii of 30 and 60 , respectively). As for G23.5+0.0, only the northern region of thedashed-lines is used for the radial profile. The source and background regions for the spectra
are designated with thin solid lines and dotted lines (a circle or an ellipse), respectively.

69

5.2. NEW EXTENDED HARD X-RAY SOURCES
109325

G25.5+0.0+26°00′

10

8300

8 75

7275

7 50

6250
225
5
200

+26°20′6 25+25°40′+0°20′

4175
150

G26.0+0.0

9100+0°20′

4-50-75+25°20′125
30
5-25 +26°00′-0°20′+25°00′+0°00′-1003+25°40′-0°20′+0°00′ +25°20′-25 0 25 50 75 100 125 150 175
75 100 125 150 175 200 225 250 275 300
10 IMSMO(Gauss) / Mosaic image: exposure NOTBLOW
10 IMSMO(Gauss) / Mosaic image: exposure NOT corrected
BLOW
corrected
225
350
G26.6-0.1
G28.6-0.1
99200 +27°00′
325
8+29°00′
8
175
300
7
7
+0°20′
150
275
6
6125
250
+0°20′
100
5
75
4 50
25

+26°40′5225+28°40′200
4175-0°40′+26°00′+0°00′

150+0°00′-0°40′+28°20′-0°20′+28°00′3
0 25 50 75 100 125 150 175 200-50 -25 0 25 50 75 100 125 150
10225IMSMO(Gauss) / Mosaic image: exposure NOTBLOW
8 IMSMO(Gauss) / Mosaic image: exposure NOT corrected
BLOW
corrected+33°00′G32.45+0.1
G35.5+0.0
225
9
7
200
+36°00′
200
8
6
175
175
7
5
150
150
6
125 +32°40′
4125 +35°40′
5100
100
75
3 75
+0°20′
4+0°20′
50
50+32°20′25+35°20′+0°00′
25-0°20′+32°00′-0°20′+35°00′+0°00′
3
2
50 75 100 125 150 175 200 225 250
25 50 75 100 125 150 175 200 225 250
BLOW IMSMO(Gauss) / Mosaic image: exposure
corrected
IMSMO(Gauss) / Mosaic image: exposure NOT corrected
Figure NOT
5.2:BLOW(Continued)+26°20′-0°20′

70

CHAPTER 5. ASCA SEARCH FOR SNRS

8350

G36.0+0.0

7325

200
10
9175
8
7150

6300
5275 +36°20′
250
4
225
200
3
175

+0°20′+36°00′

6125
5
100
4
75
3 50

150
-0°40′

225+35°40′-0°20′+0°00′

2

G37.0-0.1+37°20′+0°20′+37°00′25-0°40′ +36°40′-0°20′+0°00′

2125-50 -25 0 25 50 75 100 125 150
25 50 75 100 125 150 175 200 225
10300IMSMO(Gauss) / Mosaic image: exposure NOTBLOW
BLOW
corrected
IMSMO(Gauss) / Mosaic image: exposure NOT corrected
175
G316.8-0.1
G318.6+0.0
9
6
275
8
150
7
5125
250
6
+0°20′+0°20′
225
100
4
5-41°00′200 -43°00′
4
175

75
G316.3+0.03150+0°00′

125
2100
75

-43°20′-43°40′-0°20′

3 50+0°00′

2502-25-44°00′-41°20′-41°40′-0°40′-0°20′ -42°00′-100 -75 -50 -25 0 25 50 75 100
125 150 175 200 225 250 275 300 325
BLOW IMSMO(Gauss) / Mosaic image: exposure NOTBLOW
corrected
IMSMO(Gauss) / Mosaic image: exposure NOT corrected
G337.1+0.0
250
225
10
9200
8
175
7
6150
5125
100
4
75
3 50

-22°20′+0°20′-22°40′-23°00′-0°20′-23°20′+0°00′

0 25 50 75 100 125 150 175 200
BLOW IMSMO(Gauss)/ Mosaic
image: exposure NOT corrected
Figure 5.2:
(Continued)

71

5.2. NEW EXTENDED HARD X-RAY SOURCES

G11.0+0.0

G22.0+0.0

G23.5+0.1

G25.0+0.0

G25.5+0.0

G26.0+0.0

G28.6-0.1

G32.45+0.1

Figure 5.3: Radial profiles of the newly detected SNR candidates. The vertical axis showsbrightness (×10−6 counts cm−2 s−1 arcmin−2 ). The best-fit models (Table 5.1) are shownwith solid lines.

72

CHAPTER 5. ASCA SEARCH FOR SNRS

G35.5+0.0

G36.0+0.0

G37.0-0.1

G316.8-0.1

G318.6+0.0

G337.1+0.0

Figure 5.3: (Continued)

73

5.2. NEW EXTENDED HARD X-RAY SOURCES

G11.0+0.0

G22.0+0.0

G23.5+0.1

G25.0+0.0

G25.5+0.0

G26.0+0.0

G26.6-0.1

G28.6-0.1

Figure 5.4: ASCA GIS spectra of the newly found SNR candidates. The best-fit power-law
models are shown with solid-lines. The best-fit parameters are shown in Table 5.1

74

CHAPTER 5. ASCA SEARCH FOR SNRS

G32.45+0.1

G35.5+0.0

G36.0+0.0

G37.0-0.1

G316.8-0.1

G318.6+0.0

G337.1+0.0

Figure 5.4: (Continued)

5.3. COMMENTS ON THE INDIVIDUAL SOURCES

5.3

75

Comments on the individual sources

Characteristics of the individual sources are discussed in this section. Since sources listed
in Table 6.1 except for G32.45−0.1 (G11.0+0.0, G25.5+0.0, G26.6−0.1, G28.6−0.1) are
observed further with ASCA, ASCA results of these sources are discussed in the following
sections.
G22.0+0.0
The spatial structures of G22.0+0.0 in the soft and hard X-ray bands are quite different. It
was, however, impossible to make spectra spatially resolving the soft and hard components,
because of the poor statistics and low spatial resolution of ASCA. This is likely to be
the reason for the extremely low NH obtained with the spectral fitting (Table 5.1). An
observation with higher spatial resolution with enough statistics is required to determine thespectra of these components separately.
G23.5+0.1
The large absorption (NH > 3.2 × 1022 cm2 ) of the X-ray spectrum indicates that G23.5+0.1
is located at distance of at least ∼ 10 kpc. For the discussion below, we estimate thelower limit of the X-ray luminosity. With the X-ray spectrum given, the smallest fluxin the 0.5–10 keV band (6.1 × 10−12 ergs cm−2 s−1 within 90% error) corresponds to a
combination of the smallest photon index and absorption within the error ranges (Γ = 1.6
and NH = 3.2 × 1022 cm2 , respectively). Then, on assumption of the 10 kpc distance, thelower limit of the X-ray luminosity is 3 × 1034 erg s−1 (0.5–10 keV).
A radio pulsar, PSR 1830−08 (Clifton & Lyne 1986), is located at (l, b) = (23.◦ 39,0.◦ 06), inside G23.5+0.1. This pulsar is relatively active with spin-down luminosity of 6 ×1035 ergs s−1 and young (age ∼ 1 × 105 yr), and shows glitches (Clifton & Lyne 1986; Shemar& Lyne 1996). The ASCA point-source, AX J183345−0828, is located at 1.5 arcmin distancefrom this pulsar, and may be the counterpart (Sugizaki et al. 2001). A part or whole of the
diffuse X-ray emission might be a pulsar-wind nebula created by this pulsar. If all the Xray emission is powered by this pulsar, the effeciency with which the pulsar’s spin-down
is converted to the nebula X-ray is LX /Ė = 0.5. Though there is large uncertainty in
the X-ray luminosity, this value is significantly larger than those of other pulsars (Becker &Truemper (1997)). Therefore, it is unlikely all of the X-ray emission of G23.5+0.1 is poweredby PSR 1830−08.
A H ii region with the radio flux of 2.6 Jy at 4.85 GHz is located at (l, b)=(23.◦ 538,-0.◦ 041) (Kuchar & Clark 1997). The X-ray emission may be originate from the associated

76

CHAPTER 5. ASCA SEARCH FOR SNRS

star-forming region, or a supernova of a massive star which evolved in the star-forming
region.
G25.0+0.0
Since the X-ray emission around (l, b)=(25.◦ 2, 0.◦ 0) is very complicated with many point-likesources and extended structures, it is difficult to determine the spatial extent of each discrete
source. This situation may be the reason for the complicated X-ray spectrum of G25.0+0.0;
the spectrum leaves wavy residuals against the fitted power-law model and the model isrejected (χ2ν /ν = 22.9/8). Spatially resolved observation is required for a disscussion on theorigin of this extended X-ray emission.
G26.0+0.0
The highly absorbed (NH ≥ 2.1 × 1022 cm−2 ) X-ray spectrum suggests that G26.0+0.0 is
at large distance (≥ 5 kpc). Then, the physical diameter and the X-ray luminosity in the
2.0–10.0 keV band are ≥22 pc and ≥ 8.0 × 1033 ergs −1 , respectively. Since no radio sourcein the NVSS data (Condon et al. 1998) spatially correlates well, G26.0+0.0 is not likely to
be a H ii region. When the spectrum is fitted with a thin thermal plasma model (APEC;
Smith et al. 2001), the best-fit temperature is 3.8 (1.9–16.3) keV. Since this temperatureis relatively high for a middle age SNR with ∼20 pc diameter, the X-ray emission may be
non-thermal, such as synchrotron X-rays.
G32.45+0.1
We found a shell-like structure of 60 diameter is located at the same position in the radio
continuum image (1.4 GHz; Condon et al. 1998). Therefore, G32.45+0.1 is very likely tobe an SNR. The X-ray spectrum of G32.45+0.1 is fitted well with deeply absorbed (NH ≥3 × 1022 cm−2 ) power-law model. Although the error range is large, the photon indexof Γ ∼ 2.5 reminds us of the non-thermal X-ray emissions from SN1006 and G347.3−0.5(Koyama et al. 1995, 1997; Slane et al. 1999), the archetypal SNRs emitting synchrotron
X-rays. Spatially resolved observation with XM M -N ewton will be reported in §6.3.
G35.5+0.0
The point-like source at (l, b)=(35.◦ 52, 0.◦ 01), AX J185608+0218 (Sugizaki et al. 2001), is
accompanied by extended emission also in the soft X-ray band (0.7–2.0 keV), though theextension (∼6 diameter) is smaller than in the hard band. An extended hard X-ray emission
with a point source on its center is typical of pulsar-wind nebulae. However, in pulsar-windnebulae, X-ray emission ordinarily becomes softer as the distance from the center becomeslarger, because of synchrotron loss. The G35.5+0.0 case conflicts with this. Spectral analysis

5.3. COMMENTS ON THE INDIVIDUAL SOURCES

77

spatially resolving hard extended emission and AX J185608+0218 is required to reveal theorigin, which could not be performed with the survey data due to poor statistics.
G36.0+0.0
The X-ray spectrum clearly shows excess around 6.0–7.0 keV against the fitted power-law
model. When an emission line is added to the model, we obtained a much improved fit
with χ2ν /ν = 8.6/12. The center energy of the emission line was 6.43+0.15−0.08 keV, indicatinga K-line from low-ionized Fe. In this fitting the best-fit photon index and the absorption+1.5
22
column density were Γ = 1.9+1.1
cm−2 . Several SNRs have been
−0.5 and NH = 3.3(−1.4 ) × 10
reported to show emission lines from low-ionized Fe (Tycho’s SNR, Kepler’s SNR, RCW86,
and G344.7−0.1; Hwang, Hughes, & Petre 1998; Kinugasa & Tsunemi 1999; Vink, Kaastra,& Bleeker 1997; Yamauchi et al. 2005).
The large absorption (NH ≥ 1.4 × 1022 cm−2 ) indicates G36.0+0.0 is located at distanceof further than 4 kpc, and therefore the diameter is larger than ∼ 20 pc. This physical size
suggests G36.0+0.0 is an SNR similar to RCW 86 (28 pc diameter), rather than the young
SNRs, Tycho and Kepler (5.4 and 2.5 pc diameters, respectively). If this analogy is the case,the hard X-ray emission of G36.0+0.0 may be synchrotron X-ray emission accompanied bya Fe-K line (Rho et al. 2002).
G37.0−0.1
The best-fit photon index obtained with the power-law model is consistent with this emission
being synchrotron X-ray emission from a SNR. The large absorption column indicates, thissource is at distance of ∼ 10 kpc or further.
G316.8−0.1
G316.8−0.1 is also know as a H ii region with several kinds of masers (e.g., OH and H 2 O)
and molecular lines (e.g, CO) (Shaver et al. 1981 and references therein). After extensive
disussion on kinematic distance and other observations, Shaver et al. (1981) conclude thedistance is 2.9 kpc.
The radio continuum emission (1415 MHz; Shaver et al. 1981) and the X-ray emission have good spatial correlation with sizes of ∼ 70 . If the X-ray emission is due tothe H ii region, the spectrum is likely to be thermal. We also tried spectral fitting witha thin thermal plasma model (APEC; Smith et al. 2001), resulting in a statistically acceptable fit. The best-fit temperature and absorption column density are 3.1+5.1−1.2 keV and
+1.7
22
−2
2.8(−1.3 ) × 10 cm . The observed flux and the absorption-corrected luminosity at 2.9 kpc(0.7–10 keV) is 1.3 × 10−12 ergs cm−2 s−1 and 3.0 × 1033 erg s−1 , respectively. The tem-

78

CHAPTER 5. ASCA SEARCH FOR SNRS

perature and the luminosity is similar to those observed in other H ii regions (e.g., W3 andNGC6334; Hofner & Churchwell 1997; Sekimoto et al. 2000). Therefore, the X-ray emissionof G316.8−0.1 is likely to arise from the multiple young stars embedded in the H ii region.
Since the absorption column density is slightly large for the 2.9 kpc distance, intrinsic absorption by the associated molecular cloud would be significant.
G318.6+0.0
In the X-ray spectrum (Figure 5.4), we see an excess around 6.0 keV against the fitted
power-law model. When an emission line is added to the model, χ2ν /ν is much improved tobe 11.9/11 with the best-fit center energy of the line 5.9+0.2−0.2 keV. Since no atomic line of any
abundant element corresponds to this energy, this line is likely to be a red-shifted Fe K-linefrom a cluster of galaxies with z ∼ 0.14. In this case, the X-ray luminosity in the 2.0–10 keV
is calculate to be ∼ 2 × 1044 ergs s−1 (H0 = 50 km Mpc−1 ). According to the relation
between temperature and X-ray luminosity of clusters of galaxies (Arnaud & Evrard 1999),
a cluster with this luminosity is expected to have ∼3–4 keV temperature, which is consistent
with the hardness of the spectrum. The large absorption column (NH > 4.1 × 1022 cm−2 ) is
also consistent with that G316.8+0.0 is located behind the Galactic plane.
G337.1+0.0
In the X-ray spectrum (Figure 5.4), a line-like excess against the best-fit power-law model
is seen around 6.7 keV, which suggests the emission is due to a thin thermal plasma. When
fitted with a thin thermal model (APEC; Smith et al. 2001), χ2ν was much improved (χ2ν /ν =20.5/21), with the best-fit temperature and absorption column of kT = 6.8+6.0−3.3 keV and+3.3
NH = 7.0(−2.6 ) × 1022 cm−2 . If we assume the average density of the ISM on the line of sight
is nH = 1 cm−3 , the distance is estimated to be at least 11 kpc. Then, the physical diameteris at least 48 pc.
An radio SNR candidate, G337.2+0.1, listed by Whiteoak & Green (1996) is locatednearby at (l, b) = (337.◦ 18, 0.◦ 06). Since the size of G337.2+0.1 is 30 × 20 , and is slightly
smaller than the X-ray source G337.1+0.0, a part of G337.1+0.0 might have been enhanced
and detected in the radio band.

5.4

Summary of the Extended Source Survey(1) We searched extended X-ray sources in the hard band image of the ASCA Galactic
Plane Survey.

5.4. SUMMARY OF THE EXTENDED SOURCE SURVEY

79(2) As a result, we found 15 unidentified extended sources.
(3) Many sources (11 sources) are deeply absorbed N > 1022 cm−2 , and suggests they are
distant (∼ 10 kpc).
(4) G316.8−0.1 is identified with a H ii region.
(5) G318.6+0.0 may be a cluster of galaxies located behind the Galacitc plane, which issuggested by a line emission around 5.5 keV.
(6) Follow-up observations with higher spatial resolution and rich statistics are required
for all the candidates to establish their origin.

Chapter 6
Follow-up Observations of the New
ASCA Sources
Since the data set of the ASCA Galactic Plane Survey is limited in statistics and spatial
resolution, we performed follow-up observations with ASCA, Chandra, and XM M -N ewton
as shown in Table 6.1.
Table 6.1: Follow-up observations on the SNR candidates
ASCAa
(ks)
G11.0+0.0
49
G25.5+0.0
50
G26.6−0.1
45
G28.6−0.1
81
G32.45+0.1
10Target

a
b

CXO
(ks)
(20)b
100
-

XM M
(ks)
27

results
?
?
?
SNR
SNR

Including the exposures of the ASCA Galactic Plane Survey.
Observed with Chandra, but the data is not archived.

6.1

Follow-ups of G11.0+0.0, G25.5+0.0, and G26.6−0.1

6.1.1

Observations and Images

In order to study spatial and spectral characteristics, we performed ASCA follow-up observations on G11.0+0.0, G25.5+0.0, and G26.6−0.1 for longer exposures than the galactic81

82

CHAPTER 6. FOLLOW-UP OBSERVATIONS

Table 6.2: ASCA observations performed on G11.0+0.0, G25.5+0.0, and G26.6−0.1
Obs. Id
54005000
57027000
55002090
57029000
55004010
57026000

Dateyyyy/mm/dd
1996/08/04
1999/09/28
1997/10/14
1999/09/26
1997/04/20
1999/10/03

Pointed Direction
Exposure(R.A., Dec.)J2000(ksec)(18h 09m 50.s 4, −19◦ 240 4000 )
10.3
h
m s
◦
0
00
(18 09 36.0, −19 27 08 )
39.0
h
m s
◦
0
00
(18 37 48.0, −06 36 42 )12.4
(18h 37m 45.s 6, −06◦ 360 4300 )
37.3
h
m s
◦
0
00
(18 39 43.2, −05 42 58 )
9.8
h
m s
◦
0
00
(18 40 14.4, −05 40 44 )35.9

Observation
Type
SurveyFollow-up
SurveyFollow-up
SurveyFollow-up

Observation
Number12
3
456

plane survey. The observation dates and the on-axis positions are summarized in Table 6.2.
We screened the GIS data using the standard procedures with a little different criteria
shown in §3.1.2; the thresholds for the cutoff rigidity and the elevation angle for selected
events were set to be COR > 6 (GV) and ELV > 10 (◦ ), respectively, and screening for SFR
and HFR were not performed. The resultant exposures after these screenings are shown in
Table 6.2.
The GIS images in the 0.7–7.0 keV band around the SNR candidates are shown in
Figure 6.1. The diffuse structures discovered in the ASCA Galactic Plane survey are clearlyseen in those images. Weak point-like sources 1, 7, and 12 are detected in G11.0+0.0,
G25.5+0.0, and G26.6−0.1, respectively. We also checked the ROSAT and Einstein pointsource catalogs and found only one source, 1RXS J184049.1−054336 in G26.6−0.1 (at thesame position as source 12) within 150 radius regions around the diffuse sources.

6.1.2

Spectral Analysis

For the spectral analyses, we combined the survey and the follow-up observation data: Observation 1 and 2 for G11.0+0.0, Observation 3 and 4 for G25.5+0.0, and Observation 5 and
6 for G26.6−0.1 (see Table 6.2). The spectra were made using the events in the circularregions of 150 diameter for G11.0+0.0, and 120 diameter for both G25.5+0.0 and G26.6−0.1,
respectively (solid-line circles in Figure 6.1). X-ray events of point source 1 for G11.0+0.0,
7 for G25.5+0.0, and 12 for G26.6−0.1 were excluded from the spectra by neglecting circular regions of 30 around these sources. In order to properly subtract the Galactic ridge
X-rays (Koyama, Makishima, Tanaka, & Tsunemi 1986; Kaneda et al. 1997), the background83

6.1. G11.0+0.0, G25.5+0.0, AND G26.6−0.1(a)1+11°20′(b)(c)+0°20′

4

13
11

5

+25°40′73+26°40′10+11°00′2+0°20′

912+25°20′-0°20′+10°40′+0°00′-0°20′

6
8+0°00′+25°00′+26°20′ -0°20′+0°00′

Figure 6.1: GIS images with Galactic coordinate around (a)G11.0+0.0, (b)G25.5+0.0, and(c)G26.6−0.1 in the 0.7–7.0 keV band, where exposure time and vignetting effect are corrected. The images are smoothed with a Gaussian profiles of σ = 0.0 5. The contour levels
are linearly spaced with 6.4 × 10−4 counts arcmin−2 s−1 interval from the lowest level of
1.9 × 10−3 counts arcmin−2 s−1 for G11.0+0.0, 1.5 × 10−3 counts arcmin−2 s−1 for G25.5+0.0,
and 2.4 × 10−3 counts arcmin−2 s−1 for G26.6−0.1. For bright point sources, the higher contour levels are truncated. The source and background regions for the spectral analysis (seetext) are shown with the solid and dashed-line circles. The crosses indicate the position ofpoint-like sources.

spectra were extracted from the source-free regions near the targets within the same FOVs,
and at the same distance from the Galactic plane as the source region. For the backgroundof G11.0+0.0, we selected the annular region between 180 and 240 diameter, but excluded
circular regions of 120 diameter around sources 3 and 4.
The background-subtracted spectra are shown in Figure 6.2. We see relatively flat spectra and no line feature. Those spectra were fitted with two models of a thin thermal plasmas
in nonequilibrium ionization (NEI; Borkowski et al. 2001b) and a power-law model. Bothof these models are statistically acceptable for all the spectra, with the best-fit parameters
shown in Table 6.3. The best-fit power-law models are shown in Figure 6.2 (solid-line).

6.1.3

Discussion

On assumption that the average density in the Galaxy toward the inner disk is nH = 1 cm−3 ,we estimated the distances using the best-fit absorption column densities. The results arelisted in Table 3 together with the physical diameters and X-ray luminosities. We used thebest-fit parameters with the power-law models (Table 6.3), whereas using those with the

84

CHAPTER 6. FOLLOW-UP OBSERVATIONS

(a)

(b)

(c)

Figure 6.2: Background-subtracted spectra (crosses) of (a) G11.0+0.0, (b) G25.5+0.0, and(c) G26.6−0.1. The spectra of GIS 2 and 3 are combined. The best-fit power-law models(Table 6.3) are designated with solid-lines.

NEI model makes no essential difference.
Counterparts in other wavelengths
In order to discuss the origins of the extended X-ray emissions of G11.0+0.0, G25.5+0.0,
and G26.6−0.1, we first looked for counterparts in other wavelengths using the SIMBAD
database.
In the region of G11.0+0.0, a radio young pulsar, PSR 1809−1917, was found in thenorth (Morris et al. 2002), whose position corresponds to source 1. This pulsar is relatively
powerful with spin-down luminosity of 2×1036 ergs s−1 . If all the X-ray emission of G11.0+0.0
is powered by this pulsar, the efficiency with which the pulsar’s spin-down is converted into
nebular X-ray is LX /Ė = 1.9 × 10−3 . This value is typical for young pulsars (Becker &Truemper 1997). Therefore, significant fraction of X-ray emission from G25.5+0.0 may be85

6.1. G11.0+0.0, G25.5+0.0, AND G26.6−0.1

Table 6.3: Best-fit Spectral Parameters of G11.0+0.0, G25.5+0.0, and G26.6−0.1
Parameter
G11.0+0.0
G25.5+0.0
G26.6−0.1
Power-law Model:
Photon Index . . . . . . . . . . . . .
NH (1022 cm−2 ) . . . . . . . . . . . .
Fluxa (ergs cm−2 s−1 ) . . . . .
χ2 /degrees of freedom . . . . .
srcut modelb :Break frequency (×1019 Hz)
Flux densitya (mJy). . . . . . .
NH (1022 cm−2 ) . . . . . . . . . . .
χ2 /degrees of freedom . . . . .
NEI model:kT (keV) . . . . . . . . . . . . . . . . .
Abundance. . . . . . . . . . . . . . . .
ne t (109 s cm−3 ) . . . . . . . . . .
NH (1022 cm−2 ) . . . . . . . . . . . .
Fluxc (ergs cm−2 s−1 ) . . . . .
χ2 /d.o.f. . . . . . . . . . . . . . . . . .

1.6 (1.4–1.9)
0.8 (0.5–1.1)
3.8 × 10−12
44.2/421.8 (1.6–2.2)
2.4 (1.8–3.2)
2.0 × 10−12
55.4/48

1.3 (1.2–1.5)
0.4 (0.2–0.6)
3.5 × 10−1273.2/55

3.2 (> 0.4)
9.0 (8.4–9.6)
0.8 (0.6–0.9)
45.3/420.4 (0.07–36)
9.6 (8.8–11)
2.3 (2.0–2.7)
55.3/48

1.3 × 103 (> 12)
5.6 (5.2–5.9)
0.6 (0.4–0.7)
75.6/55

11.5 (7.3–21.0)
7.1 (4.5–14.0)
0.7 (0.1–1.8)
0.4 (0.1–0.7)
0.9 (< 4.2)
3.2 × 102 (> 58)
0.6 (0.4–1.1)
2.3 (1.8–3.0)−12
3.8 × 102.0 × 10−12
42.7/40
63.5/53

14.5 (8.7–34.8)
0.5 (0.1–1.1)
5.1 (3.0–9.1)
0.9 (0.5–1.4)
3.4 × 10−12
31.1/45aAbsorbed flux in the 0.7–10.0 keV band. b The radio spectral index is fixed to 0.5. c At the
1 GHz band.

a

Table 6.4: Physical Parameters of the Diffuse Sources
Parameter
G11.0+0.0 G25.5+0.0 G26.6−0.1
Distancea (kpc) . . . . . . .
2.6
7.8
1.3
Diameter (pc) . . . . . . . .
11
27
4.5
b
−1
33
34
Luminosity (ergs s ) 3.7 × 10
2.3 × 10
8.1 × 1032

Calculated from the best-fit absorption column in Table 6.3 assuming themean density in the Galactic plane is 1 cm−3 .
bAbsorption-corrected luminosity in the 0.7–10.0 keV.

86

CHAPTER 6. FOLLOW-UP OBSERVATIONS

powered by this pulsar. However, the X-ray emission is spatially asymmetric with respect
to the pulsar. We need X-ray observation with higher spatial resolution to conclude whatfraction of the G25.5+0.0 can be explained by the pulsar.
H ii regions were found within all the three sources; G10.9+0.1 and G11.0+0.0 in
G11.0+0.0, G25.4+0.0 in G25.5+0.0, and G26.6−0.1 in G26.6−0.1 (Kuchar & Clark 1997;
Note several names of the H ii regions are the same as the X-ray sources). However, since
their spatial sizes are significantly smaller than those of the X-ray emissions, these H ii regions
are not likely to be the direct power source of the X-ray, although a supernova originatingthese H ii regions may be the power source.
Thermal Emission Scenario
The high temperatures derived with the NEI models indicate G11.0+0.0, G25.5+0.0, andG26.6−0.1 are very young SNRs. On assumption of a uniform density plasma sphere, weestimate the emission measures (EM), dynamical times (t =radius/sound velocity), electron
densities (ne ), total masses (M ), and thermal energies (E), which are listed in Table 6.5.
We find that total masses of G11.0+0.0 and G26.6−0.1 are only a few times solar,
consistent with young SNRs. However, the thermal energies are significantly smaller thanthe canonical explosion energy of a supernova (1 × 1051 ergs). Therefore, only a small
fraction of explosion energy must be converted to the thermal energy; hence, they may bein the free expansion phase. In this case, the X-ray emitting plasma is attributable to thesupernova ejecta with large metallicity, but the observations show no such large abundance.
Thus, thermal X-ray emissions from young SNRs are not likely the case for G11.0+0.0, andG26.6−0.1.
For G25.5+0.0, the physical parameters are consistent with the adiabatic phase SNR;
however, the plasma density (< 0.1 H cm−3 ) is much smaller than that of the interstellar
gas in the Scutum arm. We thus regard that a thermal SNR for this diffuse source is also
unlikely.
SN 1006-like caseWe discuss the possibility that they are SN 1006-like SNRs. Photon indices of G11.0+0.0
and G25.5+0.0 are ∼ 2, which are smaller than those of typical SN 1006-like SNRs, SN 1006
itself, G347.3−0.5, and G266.2−1.2. This may indicate that the synchrotron cut-off energyof electrons is higher in G11.0+0.0 and G25.5+0.0 than in those typical SN 1006-like SNRs.
In order to establish the SN 1006-like SNRs, the presence of synchrotron radio emis-

6.1. G11.0+0.0, G25.5+0.0, AND G26.6−0.1

a87

Table 6.5: Parameters for thermal plasma scenario.
Parameter
G11.0+0.0 G25.5+0.0 G26.6−0.1
Emission measure EMa (cm−3 ) 2.3 × 1056 1.5 × 1057 4.9 × 1055
Dynamical time tb (s) . . . . . . . . 1.3 × 1011 3.9 × 1011 4.7 × 1010
Electron density ne (cm−3 ) . . .
0.11
0.070
0.18
Total mass M (M ) . . . . . . . . . .
1.8
18
2.3
c
50
50
Thermal energy E (ergs) . . . . . 1.2 × 10
7.3 × 10
1.9 × 1049

Emission measure ne V 2 , where ne is the electron density and V is the plasma
volume.
b
The ratio of the radius to the sound velocity.
c
E = (3/2)(ni + ne )V kT , where ni is the proton density.

sion is essential. However, no cataloged radio source is found in the NVSS archival data
(1.4 GHz; Condon et al. 1998). In order to estimate the radio flux for the case of SN 1006like SNRs, we tried spectral fitting with a SRCUT model (Reynolds 1998; Reynolds &Keohane 1999). The spectral index at 1 GHz was frozen to α = 0.5, the expected value by
the first-order Fermi acceleration. The fittings were statistically acceptable, and the bestfit parameters are listed in Table 6.3. We also tried the same fitting with α = 0.6 andfound no essential difference in the best-fit parameters. The best-fit flux density for eachsource suggests that the expected surface brightness at 1 GHz is 9.8×10−24 , 1.7×10−23 , and9.8×10−24 W m−2 Hz−1 sr−1 for G11.0+0.0, G25.5+0.0, and G26.6−0.1, respectively. On
the other hand, the minimum surface brightness of the cataloged radio SNRs in our survey
field is 2 × 10−21 W m−2 Hz−1 sr−1 (for G3.8+0.3; Case & Bhattacharya 1998), which issimilar to SN 1006 (3×10−21 W m−2 Hz−1 sr−1 ; Long et al. 1988; Winkler & Long 1997) andG347.3−0.5 (4×10−21 W m−2 Hz−1 sr−1 ; Ellison et al. 2001; assuming α = 0.5), but is two or
three orders of magnitude larger than our new sources. Even the minimum surface brightnessof the all cataloged radio SNRs is larger than our sample, 6.2 × 10−23 W m−2 Hz−1 sr−1 (forG156.2+5.7; Case & Bhattacharya 1998). Therefore, no radio counterpart for our new X-ray
sources would be simply due to limited detection threshold of the current radio observations.
Therefore, no radio counterpart of our new X-ray sources would be simply due to limited
detection threshold of the current radio observations.
Crab-like case
The power-law spectra and the luminosities (1034 –1035 ergs s−1 ) suggest that the new sources
are also Crab-like (plerionic) objects. We also discuss no detection of radio source in thiscase. The luminosity ratio between the radio and X-ray band is largely different among

88

CHAPTER 6. FOLLOW-UP OBSERVATIONS

the plerionic SNRs: for example, the radio luminosity of 3C58 is about 10 times smaller,
while that of X-ray is about 1000 times smaller than that of the Crab Nebula (Hands etal., 2001). We estimated the expected surface brightness using the flux density of 3C58
and the Crab Nebula (Green, 2004). For the 3C58-like case, the surface brightness of
G11.0+0.0, G25.5+0.0, and G26.6−0.1 at 1 GHz is expected to be 2.5×10 −20 , 5.1×10−20 ,and 1.9×10−20 W m−2 Hz−1 sr−1 , respectively. Since these values are larger than the detection limit in the surveyed region (see previous paragraph), we exclude the 3C58-like case.
On the other hand, for the Crab-like case, their respective surface brightness becomes tobe 7.3×10−23 , 1.5×10−22 , and 5.6×10−23 W m−2 Hz−1 sr−1 . Therefore, a possibility of the
Crab-like case can not be excluded.

6.2

ASCA and Chandra Follow-up of G28.6−0.1

6.2.1

Follow-up Observations with ASCA

G28.6−0.1 was observed with ASCA twice other than the ASCA Galactic plane survey(Table 6.6). One of which was performed for the study of the Galactic Ridge Emission(Kaneda et al. 1997), and the other one was for a follow-up observation of G28.6−0.1 itself.
In total, 81 ks observations were performed on G28.6−0.1. The exposure-corrected smoothed
image is shown in Figure 6.3. Spatially extended X-rays with an elliptical shape is evident
especially in the hard band (2.0–7.0 keV), and the extended structure seen in the survey
data is confirmed.
The X-ray spectrum was extracted using the photons in an elliptical area of 130 × 90
(the inner ellipse in Figure 6.3), whereas the background spectrum was obtained from anelliptical ring around the source (between the outer and inner ellipses in Figure 6.3). The
background-subtracted spectrum is shown in Figure 6.4. The spectrum is hard and without
conspicuous structure of lines. In order to find out the origin of this X-ray emission, we fit
the spectrum with two models of a power-law and a thin thermal plasma in non-equilibrium
ionization (NEI; Masai 1984; Borkowski et al. 2001b). As a result, both of these models
were statistically acceptable with the best-fit parameters and the best-fit models shown in
Table 6.7 and Figure 6.4, respectively.
The best-fit NH value (2.6 − 4.0) × 1022 cm−2 is consistent with that obtained with thesurvey data. This value is similar to the nearby SNRs, Kes 73 [(1.6 − 2.3) × 1022 cm−2 ;Gotthelf & Vasisht (1997)] and Kes 75 [(2.9 ± 0.4) × 1022 cm−2 ; Blanton & Helfand (1996)]
at (27.◦ 4, 0.◦ 0) and (29.◦ 7, −0.◦ 3), respectively. Using the H i absorption data, the distances of6.2. G28.6−0.1

89

Table 6.6: ASCA observations performed on G28.6−0.1.
Obs. Id
Date
Pointed Direction
Exposure
yyyy/mm/dd(R.A., Dec.)J2000(ksec)hm s◦
000
90006000
1993/10/22 (18 44 40.7, −03 36 14 )
19
a
h
m s
◦
0
00
55004050
1997/04/22 (18 42 38.2, −04 09 13 )1057028000
1999/03/23 (18h 43m 10.s 4, −04◦ 030 0600 )42aThis observation is a part of the ASCA Galactic Plane Survey.

Kes 73 and Kes 75 are estimated to be 6–7.5 kpc (Sanbonmatsu & Helfand 1992) and 9–21 kpc(Caswell et al. 1975; Becker & Helfand 1984), respectively. Koyama et al. (1990) found manytransient sources at the tangent point of the Scutum arm (about 8.5 kpc) with absorptionsof around 1023 cm−2 , which are larger than that of G28.6−0.1 Accordingly, we infer thatthe distance of G28.6−0.1 is in between 6–8.5 kpc. Hereafter, we assume the distance to be
7 kpc. The X-ray luminosity (0.1–10 keV) and the source size (the mean diameter) are thenestimated to be 3 × 1034 ergs s−1 and 20 pc, respectively. These parameters are similar to
those of shell-like young SNRs.
The high temperature of ∼5 keV corresponds to the shock speed of ∼2000 km s−1 andis also consistent with that G28.6−0.1 is a young SNR. The high speed shock indicates thatthe swept-up mass is not very large and if we assume the explosion energy is 1 × 1051 ergs,
the mass is estimated to be 4M . Since this value is comparable to the ejected mass, theabundance is expected to be very high with prominent emission lines, which can be seen
in most of the young SNRs, such as Cas A, Tycho’s SNR, and Kepler’s SNR (Holt et al.
1994; Hwang, Hughes, & Petre 1998; Kinugasa & Tsunemi 1999). However, the abundanceof G28.6−0.1 with the NEI model is only ∼1 solar. Therefore, it is difficult to attribute thehard and featureless X-ray spectrum of G28.6−0.1 to thermal emission. Thus, we expect
the emission from this remnant to be non-thermal. The luminosity is almost consistent withthat of the non-thermal component of SN 1006 (∼ 2 × 1034 ergs s−1 ; Dyer, Reynolds, &Borkowski 2004) and the diameter is similar to SN 1006.

6.2.2

Chandra Observations of G28.6−0.1

The ASCA observations suggested that G28.6−0.1 is an SNR emitting non-thermal X-rays.
However, two other possibilities remained; this source may be a collection of point sourcesor the origin of the X-ray emission may be thermal plasma in non-equilibrium ionization.
In order to exclude these possibilities by resolving point sources and a putative thermal

90

CHAPTER 6. FOLLOW-UP OBSERVATIONS

Figure 6.3: ASCA GIS images in the 0.7–2.0 keV (a) and 2.0–7.0 keV (b) bands. The VLA
20 cm radio contours by Helfand et al. (1989) are superposed on the 2.0–7.0 keV image (b).
The ASCA images are corrected for the exposure and smoothed with a Gaussian kernel of
σ = 7500 . The source and the background regions for the spectrum are shown with broken-line
ellipses.

(a)

(b)

Figure 6.4: Background-subtracted spectrum of the combined GIS2 and 3 data of the 3
observations. The solid lines represent the best-fit power-law model (a) and the best-fit NEI
model (b).

91

6.2. G28.6−0.1

Table 6.7: Best-fit parameters of G28.6−0.1 for a power-law and a thin thermal plasma
model with the ASCA observations∗ .

∗
†

Model

Γ / kT.../(keV)

Abundance

power-law . . . . . . . . . .
thin thermal plasma

2.1+0.3−0.4
5.4+3.4−1.8...
1.5+1.9−1.2†

log ne t(s cm−3 )NH
22(10 cm−2 )χ2 /d.o.f.

...
9.6+0.3−0.3

2.6+0.8−0.6
4.0+1.8−1.2

45.7/47
32.5/46

Errors and upper-limits are at 90% confidence for one relevant parameter.
Assuming the solar abundance ratio (Anders & Grevesse 1989).

component, we performed X-ray observations with the high spatial resolution of Chandra.
Observations and Data Reduction
The Chandra deep observations on the G28.6−0.1 region were performed on 2000 February
24–26 (here observation 1) and on 2001 May 20–21 (here observation 2). The targeted
positions are R.A. = 18h 43m 57.s 8, decl. = −04◦ 040 45.00 9 (epoch 2000) and R.A. = 18h 43m 32.s 1,
decl. = −03◦ 540 44.00 8 (epoch 2000) for observation 1 and 2, respectively. The observation
regions were 170 × 170 fields of the ACIS-I arrays.
G28.6−0.1 lies near the northwest edge and the east part of the ACIS-I array, for observation 1 and 2, respectively. Data acquisition from ACIS-I was made in Timed-Exposure
Faint mode with the chip readout times of 3.24 s. Data reduction and analysis were madeusing the Chandra Interactive Analysis of Observations (CIAO) software version 2.2. Using
Level 2 processed events provided by the pipeline processing at the Chandra X-ray Center,
we selected the ASCA grades 0, 2, 3, 4 and 6, as X-ray events. The other events due tocharged particles and/or hot and flickering pixels were removed. We also removed high background data in the observation 1 where total count rates were larger than 5.5 counts s −1 .
The effective exposures were then about 94 ks and 99 ks for observation 1 and 2, respectively.
Over-all Morphology
In the raw images obtained by these observations, we see numerous point sources. We at firstpicked up point sources from the raw images in the energy bands of 0.5−3.0 keV, 3.0−8.0 keV,
and 0.5−8.0 keV using the program wavdetect (Freeman et al. 2002) with the significancecriterion at 4.0σ, which corresponds to about 6 – 33 counts s−1 , depending on the size of thepoint-spread function (PSF), the energy band, and the background level. Then 274 sources
are found in the two observations, of which 225, 116, and 271 are detected in the energy

92

CHAPTER 6. FOLLOW-UP OBSERVATIONS

bands of 0.5−3.0 keV, 3.0−8.0 keV, and 0.5−8.0 keV, respectively. Surface density of thepoint sources detected in the 3.0−8.0 keV band is not significantly higher than that expectedfrom the extragalactic sources seen through the Galactic plane. Therefore, most of the hard
X-ray sources are considered to be extragalactic (Ebisawa et al. 2001). On the other hand,
soft X-ray point sources show significant excess over the extragalactic population. They have
low temperature (< 1 keV) thermal spectra and some of them exhibit flare-like temporalvariations (Ebisawa et al. 2002). Therefore, most of these soft X-ray sources are presumably
nearby active stars.
The combined image (observation 1 + 2), after the exposure correction and smoothingwith a Gaussian kernel of σ = 600 is shown in Figure 6.5, where the point sources of > 5.0σ
detection in the 0.5−8.0 keV band are given by crosses. In Figure 6.5, we can see an
extended complex at the north (shown by the solid-line ellipse). Since the point sources arescattered uniformly over the field of view, the extended X-ray structure is unlikely to be a
local enhancement of many point sources, but will be a really diffuse source. The center
position of the diffuse source is at around R.A. = 18h 43m 50s , decl. = −03◦ 520 0000 (epoch2000) and the emission is extended in the elliptical shape of 110 × 7.0 5.
We found that G28.6−0.1 consists of numerous clumps. We name the brightest sphericalclump in the south-eastern part of G28.6−0.1 as CXOU J184357−035441 from its central
position (Although this source is designated as “CXO J184357−035441” in Ueno et al. (2003),“CXOU J184357−035441” fits the Chandra source naming convention). A closed-up viewof this clump in the 1.0–6.0 keV band is shown in Figure 6.6, together with the 20 cm (1.5
GHz) contour of a radio source F (Helfand et al. 1989). A notable feature is its peculiar
shape: an elliptical head with a jet-like tail 3000 long to the southeast.
The spatial correlation between the diffuse X-ray and radio is demonstrated by the
20 cm VLA contours in Figure 6.5, with the designation of the radio source (A–I) given by
Helfand et al. (1989). The X-ray emission consists of many clumps filling an ellipse, which ispartly outlined by the radio sources C and F, the SNR candidates proposed by Helfand et al.
(1989). The south tail of the radio source G is also correlated with the X-ray enhancement.
Though the source H is said to be extragalactic because of its infrared emission (Helfand etal. 1989), it is also inside G28.6−0.1 and may be associated. Note, however, no significant
X-ray emission is associated with the source E, which is a partial radio shell and an SNR
candidate suggested by Helfand et al. (1989).

6.2. G28.6−0.1

93

Figure 6.5: Combined ACIS-I image of observation 1 and observation 2 in the 1.5–8.0 keVband overlaid with the VLA 20 cm radio contours by Helfand et al. (1989). The X-ray image
is corrected for the exposure, smoothed with a Gaussian kernel of σ = 600 , and plotted from
5.7 × 10−6 to 1.1 × 10−4 counts cm−2 s−1 arcmin−2 in logarithmic scale. Plotted are FOVs ofobservations 1 and 2 (large squares), G28.6−0.1 source region(ellipse) and background region(circles), and point sources detected by the wavdetect software with significance greater than
5.0σ (plus signs). Radio sources are labeled following Helfand et al. (1989).

94

CHAPTER 6. FOLLOW-UP OBSERVATIONS

Figure 6.6: Close-up view of CXOU J184357−035441 in the energy band 1.0–6.0 keV. The
image is smoothed with a Gaussian kernel of σ = 1.s 5. Shown are the 20 cm radio continuum source F (Helfand et al. 1989; solid contours, the elliptical region for extracting thespectrum of CXOU J184357−035441 (dashed ellipse), and the four point sources detected in
CXOU J184357−035441 (white plus signs).

X-ray Spectra of G28.6−0.1
The X-ray spectra of G28.6−0.1 were separately made for observations 1 and 2, from the elliptical region of 110 × 7.0 5 shown by the ellipse in Figure 6.5. An elliptical region surrounding
CXOU J184357−035441 (dotted line in Figure 2) was excluded to extract the energy spectrum. Although observation 2 covers full region, observation 1 covers only about the southern
1. The point-source events, which are in the PSF (90% encircled radii) circles around the3point sources detected with the wavdetect software, were removed from the source events.
The background regions were selected at the same Galactic latitude as the source region, sothat the GRXE are properly eliminated (e.g., Kaneda et al. 1997). The background regions
for observations 1 and 2 are shown by circles in Figure 1. The point-source contribution was
excluded from the background spectra, in a way similar to the source regions.
The background-subtracted spectra for observations 1 and 2 have no emission line, as
shown in Figure 6.7. We fitted the spectra with a power-law model. The fits are acceptablewith χ2 /degree of freedom (d.o.f.) = 18.9/27 and 48.1/52 for observations 1 and observation
2, respectively. The best-fit parameters are given in Table 6.8. Although the spectrumof observation 1 represents only the southern part of G28.6−0.1, the spectral parameters,

95

6.2. G28.6−0.1

except the flux, are consistent with those of observation 2, the full spectrum of G28.6−0.1.
We therefore carried out a simultaneous fit to the observations 1 and 2 spectra with apower-law model leaving each normalization of the two observations independent. The fit isacceptable with the best-fit parameters given in Table 6.8.
(a)

(b)

Figure 6.7: Spectra of G28.6−0.1 derived from observation 1 (squares) and observation 2(circles) with the best-fit models in the simultaneous fitting for observation 1 (dotted lines)
and observation 2 (solid lines). Fitted models are a power-law model (a) and a NEI model(b), respectively.

To examine the possibility of a thermal plasma origin, we also proceeded with a simultaneous fit with a model of a thin thermal plasma in non-equilibrium ionization (NEI;
Borkowski et al. 2001b). This NEI model is acceptable with the best-fit parameters in Table 6.8. The metal abundances, however, are lower (∼ 0.2 solar) and the temperature ishigher (∼ 5.4 keV) than any of the diffuse thermal sources (e.g., SNRs and star forming
regions) in the Galaxy.

X-ray Spectrum of CXOU J184357−035441
The X-ray spectrum of CXOU J184357−035441 was made from the full elliptical region of
5000 ×3000 shown in Figure 6.6. We used the same background spectrum as that for G28.6−0.1
in observations 2. As is shown in Figure 6.8, the background-subtracted spectrum is muchsofter than that of G28.6−0.1 and shows clear emission lines at 1.85 and 2.41 keV, which
are equal to or slightly lower than those from He-like Si and S. We fitted the spectrum withan NEI model. This model is acceptable with the best-fit parameters given in Table 6.9 andthe best-fit model in Figure 6.8.

96

CHAPTER 6. FOLLOW-UP OBSERVATIONS

Table 6.8: Best-Fit Model Parameters for G28.6−0.1 with Chandra.

Parameter
Γ/kT [keV]. . . . . . . . . . .
Abundance . . . . . . . . . . .
log(ne t [cm−3 s]) . . . .
NH [1022 cm−2 ] . . . . . .
Absorbed fluxa . . . . . . .
Unabsorbed fluxc . . . .
χ2 /d.o.f. . . . . . . . . . . . .

Observation 1

power-lawObservation 2

2.4+1.1−0.9

2.1+0.4
−0.4

···
···
4.7+2.8
−1.8
2.3
3.6
18.9/27······
3.7+0.7
−0.6
3.1
4.3
48.1/52

Observation 1+2

NEI
Observation 1+2

2.1+0.4
−0.3
···
···
3.8+0.7
−0.6
3.1b
4.3b
68.1/81

5.4+3.5−1.6
0.17 +0.17−0.14
11.1+0.4−0.7
3.5+0.6
−0.5
3.0b
4.1b
65.3/79

Observed flux (10−14 ergs cm−2 s−1 arcmin−2 ) per unit area (arcmin2 ) in 2.0–10.0 keV.
b
The fluxes of observation 1+2 are calculated from the normalizations for the spectrum of observation 2,
which covers whole G28.6−0.1.
cAbsorption-corrected flux per unit area (10−14 ergs cm−2 s−1 arcmin−2 ) in 2.0–10.0 keV.
a

Figure 6.8: Background-subtracted spectrum of CXOU J184357−035441. Solid line: Best-fit
NEI model.

6.2.3

Discussion on G28.6−0.1 and surroundings

G28.6−0.1
We found that the newly identified SNR, G28.6−0.1, has many X-ray clumps, which areglobally filling an elliptical region. The X-ray spectrum is well-fitted with either a power-law
or a thin thermal plasma model, similar to the ASCA spectrum. Chandra gives more severe
constraint on the thin thermal parameters than ASCA, due mainly to the removal of the97

6.2. G28.6−0.1

Table 6.9: Best-fit parameters for CXOU J184357−035441 by the NEI model.
Parameter
ValuekT . . . . . . . . . . . . . . . . .
0.71+0.22−0.16 keV
Abundance . . . . . . . . .
1.1+1.4−0.7
log[ne t (cm−3 s)] . . . .
10.4+1.5−0.3+1.6
NH . . . . . . . . . . . . . . . . .
7.4−1.2 × 1022 cm−2
Absorbed fluxa . . . . . . 1.1 × 10−13 erg cm−2 s−1
Unabsorbed fluxb . . . 4.5 × 10−11 erg cm−2 s−1
χ2 /d.o.f. . . . . . . . . . . . .
22.2/39ab

Observed flux in the energy band 0.7–10.0 keV.
Absorption-corrected flux in the energy band 0.7–10 keV.

contamination from the thermal source CXOU J184357−035441. The thermal scenario of
G28.6−0.1 requires uncomfortably high temperature (> 3.8 keV) and low metal abundance(< 0.34 solar; Table 6.8). This strengthens the conclusion made with ASCA that the X-raysof G28.6−0.1 is non-thermal origin with a power-law spectrum.
The absorption (NH ) of (3.2 − 4.5) × 1022 cm−2 determined with Chandra is slightly
larger than, but roughly similar to the ASCA result. We therefore adopt the source distanceof 7 kpc, following the discussion given in §6.2.1. The X-ray luminosity (2.0–10.0 keV) andthe source size are then estimated to be 1.5 × 1034 ergs s−1 and 18 pc (mean diameter),respectively.
Since the spectrum is a power-law with the photon index of 2.1 (the energy index of 1.1),
the X-rays would be due to synchrotron emission from high energy electrons in a power-law
distribution. The corresponding spectral index of the electrons is 3.2, which is larger thanthe value of 2 which is expected by the first-order Fermi acceleration. This steepening in the
X-ray band is also seen in all of the well established SNRs as a site of high energy electrons(SN 1006, G347.3−0.5, and G266.2−1.2, e.g., Koyama et al. 1995, 1997; Slane et al. 2001),
and their photon indices are very similar (between 2.2 − 2.6).
In order to study the wide-band spectrum of the synchrotron emission, we fitted the
X-ray spectrum with an SRCUT model (Reynolds 1998), on an assumption that the energy
index (α) in the radio band (1 GHz) to be 0.5 or 0.6. The derived cutoff energy and theflux density at 1.5 GHz are shown in Table 6.10. Here the absorption is consistent withthe fit with the power-law model. Both of the predicted radio fluxes at 1.5 GHz (= 20 cm)
by the SRCUT model are smaller than the summed flux of the radio clumps C, F, G, and98

CHAPTER 6. FOLLOW-UP OBSERVATIONS

Table 6.10: Best-Fit Parameters by the SRCUT Model
Parameter
α = 0.5a
α = 0.6aCut-off frequency [1017 Hz]
7.2 (2.8 − 27)
11 (2.2 − 92)b−2Flux density [Jy] . . . . . . . . 1.8 (1.1 − 4.6) × 10
0.11 (0.07 − 0.45)
+0.3
22
−2
NH [10 cm ] . . . . . . . . . . . .
3.6−0.2
3.7+0.2−0.3
2
χ /d.o.f. . . . . . . . . . . . . . . . . . . .
75.1/8275.0/82ab

The fixed value of α.
Flux density at 1.5 GHz.

H (∼1 Jy; Helfand et al. 1989). This apparent inconsistency may be due to our simple
assumption that magnetic field strength (B) and density of high-energy electrons (ne ) are
uniform in G28.6−0.1. A more realistic assumption is that the radio clumps have higher B
than average, hence producing the high contrast in the radio image. On the other hand, the
X-ray flux contrast in the radio clumps is small, because the number ne of higher energy
electrons (which emit X-rays) is smaller than average, because the synchrotron energy loss islarge at high B. This scenario is consistent with the observed X-ray and radio morphology.
In this case, the SRCUT model predicts that in the radio clumps the cutoff energy is lower
and the X-ray spectrum is steeper than in the other regions. The limited photon flux,
however, does not allow us to verify this scenario. Discrepancy between synchrotron X-ray
and radio morphology is also seen in G347.3−0.5, an established source as an SNR emitting
synchrotron X-rays (Slane et al. 1999).

CXOU J184357−035441
The X-ray energy spectrum of CXOU J184357−035441, thin thermal spectrum with 0.7 keVtemperature and solar abundance, is typical for young and intermediate-age SNRs. Since theprojected position of this source is in the SNR, G28.6−0.1, and is associated with the nonthermal radio source F (Helfand et al., 1989), a naive scenario is that CXOU J184357−035441
is a thermal counterpart of the SNR. It may be possible that we have a large shell whose soft
X-ray emission is invisible because heavily absorbed. While a reverse shock is propagating
toward the center, the core region makes harder spectrum. Hence, CXOU J184357−035441
might be one of the clumps made by the reverse shock. Adopting the same distance asG28.6−0.1 (7 kpc), we estimate that the X-ray luminosity (0.7–10.0 keV) is 2.7×10 35 erg s−1
and the physical size is 1.7 × 1.0 pc of “head” and 1.0 pc-long “jet”. Assuming a uniform
density of a prorate sphere with 1.7 × 1.0 × 1.0 pc diameters, the density of the “head” is6.3. G32.45+0.1

99∼ 15 cm−3 . From the ionization parameter (ne t) of (1.4 − 83) × 1010 cm−3 s, this plasma
would be heated-up very recently (30 − 1700 years ago), possibly by the reverse shock.
Since CXOU J184357−035441 exhibits 2 times larger absorption than G28.6−0.1, thissource is likely located at the distance of 14 kpc. Derived plasma luminosity, size and density
are still consistent with being a young SNR behind AX J1843.8−0352.
Since the absorption of CXOU J184357−035441 is nearly the same as the total Galactic
absorption through the Galactic plane (Ebisawa et al., 2001), a possibility of an extragalactic
source can not be excluded. The jet-like morphology is also unusual for an SNR plasma, butrather resembles the hot plasmas associated with the jet sources, like SS433 or AGN. It is
unclear whether CXOU J184357−035441 is a thermal component associated with G28.6−0.1,
another SNR located in the far side of our Galaxy, or an extragalactic source.

6.3

XM M -N ewton Follow-up of G32.45+0.1

6.3.1

XM M -N ewton Observation

Data reduction and image
G32.45+0.1 was observed with XM M -N ewton in September 25, 2003. Although both the
MOS and PN cameras were operated, the major part of G32.45+0.0 was not covered with
the PN camera due to the CCD gaps and bad columns, and we verified that cleaned PN
data do not improve precision of results obtained with the MOS cameras only. Therefore,
we ignore the PN data. For data reduction and analysis, we have used the version 5.4.1of the Standard Analysis System (SAS) software. Since the background showed flares, we
accumulated the background light-curve of the full-field events with 10–12 keV energy, andscreened out the time intervals when the count rate is larger than 0.3 count s−1 . After thisfiltering of the high background intervals, the effective exposure times of MOS1 and MOS2
are 20.5 and 21.2 ks, respectively.
Figure 6.9 shows the soft (0.5–2.0 kev) and the hard (2.0–7.0 keV) band X-ray images.
Only in the hard band image (b), we see a shell-like diffuse structure, whose radius is about
4 arcmin.
In order to estimate the accurate flux of the diffuse X-rays, we picked up point-sources
using the ewavelet software of the SAS, with the detection threshold of 5σ. Detected point
sources are designated with crosses in Figure 6.10. Three sources are located in the diffuse
shell, however, the extents of them determined by the ewavelet software are larger than the

100

CHAPTER 6. FOLLOW-UP OBSERVATIONS

Figure 6.9: XM M -N ewton/MOS1+2 images of G32.45+0.1 in the soft (0.5–2.0 keV; (a))
and the hard (2.0–7.0 keV; (b)) bands superimposed the NVSS 20 cm radio contours (Condon
et al. 1998). The images are smoothed with a Gaussian kernel of σ = 9.00 6. The coordinates
are galactic.

Table 6.11: The positions and the extents of the point-like sources in the shell (crosses inthe ellipses in Figure 6.9).
(R.A., Dec.)J2000
regionms◦
000(18 50 19 , −00 26 20 ) southeast(18h 50m 06s , −00◦ 250 5800 ) southwest(18h 50m 05s , −00◦ 260 0400 ) southwest
hextent (FWHM in arcsec)
60 ± 2560 ± 25
26 ± 6.1

point spread function of XM M -N ewton (FWHM ∼ 500 ). We, therefore, conclude that these
are not discrete point sources, but local enhancements of the diffuse emission. Thus, thediffuse structure discovered with the ASCA survey data was confirmed with the high spatial
resolution of XM M -N ewton.
Spectral Analysis
The spectrum was extracted from a region consists of three ellipses shown in Figure 6.10
(solid-line). The background was extracted from the neighborhood ellipses (Figure 6.10;

101

6.3. G32.45+0.1

dotted-line). This region was selected so as that the intrinsic brightnesses of the background
X-ray emission (the Galactic Ridge Emission; Kaneda et al. 1997) and the vignetting effectare the same for both the source and background regions; the background region is on
the same distances from both the galactic plane and from the nominal point as the source
region. The background-subtracted spectra of MOS1 and MOS2 are shown in Figure 6.11.
Since these spectra are consistent with each other and have no emission line, we fitted themsimultaneously with a power-law model modified by the interstellar absorption. The crosssections of the absorption were adopted from Morrison & McCammon (1983), and the solar
abundance (Anders & Grevesse 1989) was assumed. The fitting was statistically acceptablewith the best-fit parameters given in Table 6.12.
To examine the possibility of a thermal plasma origin, we also tried fitting with a thinthermal plasma model in non-equilibrium ionization (an NEI model; Borkowski et al. 2001b),
and as a result, obtained a statistically acceptable fit with the best-fit parameters shown in
Table 6.12.
Table 6.12: Best-fit parameters of G32.45−0.1 for a power-law and a thin thermal plasma
model with the XM M -N ewton observations.
†Γ / kT
Abundancelog(ne t)NHχ2 /d.o.f.
Model.../(keV)[cm−3 s] (1022 cm−2 )power-law . . . . . . . . . .
thin thermal plasma†

2.2+0.8−0.4
5.2+6.8
−2.2

...
−2

1.0 × 10 (< 0.35)

...
13.6

5.2+2.3−1.3+1.5
4.6−1.2

52.4/46
52.4/44

Assuming the solar abundance ratio (Anders & Grevesse 1989).

6.3.2

Discussion on G32.45+0.1

We found a clear shell-like structure with a radius of ∼ 40 in the hard X-ray band (2.0–
7.0 keV) with an XM M -N ewton observation. The X-ray spectrum is well fitted with either
a power-law or a thin thermal plasma model. Here, the thermal model requires temperatureof 3.0–12 keV, which is similar to that of typical young SNRs such as Cas A and Tycho’s SNR(kT ∼ 2 − 4 keV; e.g., ). However, the thermal model also requires an low metal abundance(< 0.35 solar), which conflicts with the case of young SNRs. Typical X-ray emissions fromyoung SNRs are dominated by those from ejecta, and the abundances are high (∼ 2 solar).
Therefore, the origin of the X-ray emission is not likely to be thermal, but non-thermal withthe power-law spectrum.
The best-fit parameter of NH = 5.2(3.9 − 7.5) × 1022 cm−2 gives rough estimation on

102

Figure 6.10: MOS1+2 hard-band (2.0–
7.0 keV) image with the spectral regions.
The source and background regions areshown with solid-line and dotted-line
ellipses, respectively.
The point-likesources detected with the ewavelet software are designated with white crosses.
The coordinates are R.A. and DEC.
(J2000).

CHAPTER 6. FOLLOW-UP OBSERVATIONS

Figure 6.11: Background-subtracted spectra of
G32.45+0.1 observed with MOS1 (circle) andMOS (square). The solid lines show the bestfit power-law models.

the source distance to be 17 kpc on an assumption that the averaged density in the galactic
plane is 1 cm−3 . Then the X-ray luminosity (0.5–10 keV) and the radius are estimated to be∼ 9.5 × 1034 ergs s−1 and ∼ 20 pc, respectively. These values are not largely different fromthose of the non-thermal component of SN 1006 (∼ 2.0 × 1034 ergs s−1 and ∼ 9.5 pc; Dyer,
Reynolds, & Borkowski 2004).
We found a radio shell at the position of G32.45+0.1 from the NRAO/VLA Sky Survey(NVSS) data at 1.4 GHz (Condon et al. 1998) with the flux sum ∼ 0.17 Jy. Although thesize of the shells seen in the X-ray and radio bands are quite similar, the enhanced positions
are different between them. Bad correlation between X-ray and radio synchrotron emissions
in SNRs was already pointed out in the discussion of G28.6−0.1, which can be explained by
non-uniformity of the magnetic field strength. In order to discuss the wide-band spectrum
assuming the synchrotron emission, we fitted the X-ray spectrum with a SRCUT model(Reynolds 1998). The best-fit parameters are shown in Table 6.13. The radio flux obtained6.3. G32.45+0.1

103

Table 6.13: Best-Fit Parameters by the SRCUT Model
Parameter
α = 0.5a
α = 0.6aCut-off frequency [1017 Hz]
3.4 (0.79–91)
5.4 (<35)b−2Flux density [Jy] . . . . . . . . 1.6 (1.4–1.8)×10
8.6 (7.4–9.4)×10−2
NH [1022 cm−2 ] . . . . . . . . . . . .
5.2 (4.5–6.2)
5.2 (4.4–6.1)
χ2 /d.o.f. . . . . . . . . . . . . . . . . . .
52.4/46
52.4/46ab

The fixed value of α.
Flux density at 1.4 GHz.

with the SRCUT model is lower than the observed flux in both the α = 0.5 and α = 0.6
cases. This is consistent our argument of magnetic field non-uniformity; the radio emission
is detected from regions of stronger magnetic field, whereas the X-ray emission is detectedfrom those of weaker magnetic field.

Chapter 7
Discussion
7.1

Properties of SNRs emitting synchrotron X-rays

In this thesis, we newly discovered SNRs emitting synchrotron X-rays: G28.6−0.1, G32.45+0.1.
Moreover we found many candidates for such SNRs with the ASCA Galactic Plane Survey.
It is notable that these SNRs are discovered with X-ray observations. Nowadays, up to 231
galactic SNRs are known (Green 2004), which are mostly discovered with radio observations.
However, synchrotron X-rays are detected from only ∼10 SNRs. On the other hand, SNRs
discovered with X-rays often show synchrotron X-rays (or non-thermal X-rays of other processes). These facts infer us that most SNRs emitting synchrotron X-rays are faint in theradio band.
In order to discuss this supposition, we utilize the Σ–D relation adding newly found
SNRs. The Σ–D relation is a phenomenological relation between radio surface brightness anddiameters of SNRs (e.g., Clark & Caswell 1976; Case & Bhattacharya 1998) introduced for
estimating a distance of each SNR. The mean surface brightness at a specific radio frequency,Σν , is a distance-independent parameter and, to a first approximation, is an intrinsic propertyof the SNR (Shklovskii 1960). Therefore, if the Σ–D relation is reliable, we can determine
physical size (D) of a SNR from Σ, and as a result the distance is determined. Although it
has been controversial that the Σ–D relation can be used for estimating distance of a SNR(e.g., Green 2004), there is a clear trend for the surface brightness of SNRs to decrease withincreasing diameter. Moreover, diameters show approximate ages of SNRs at least. In order
to see evolution stages of the newly found synchrotron X-ray sources and to compare theirsurface radio brightness with the other radio selected SNRs, we plotted them to the Σ–D
relation.
105

106

CHAPTER 7. DISCUSSION

Figure 7.1: The surface brightness vs. diameter (Σ–D) relation for shell SNRs with known
distances (see Table 7.1 and 7.2). The circles and crosses designate SNRs with and without
synchrotron X-rays listed in Table 7.1. The squares designate SNRs we added (Table 7.2).
The solid line shows the best-fit power-law function (see text)

As radio selected SNRs, we consider those lied in Table 2 of Green (2004). The utilized
values of distances and sizes, and calculated diameters and radio surface brightnesses at
1 GHz are shown in Table 7.1. Those distances are measured mainly 4 methods independently of X-ray observations: e.g., optical expansion and proper motion studies, 21 cm H iabsorption spectra, and association with H i or CO features in the surrounding ISM (seeGreen 2004). In order to restrict the discussion to shell SNRs, Crab-like “SNRs” were excluded here, which are cataloged with the type flag “F” in Green (2004). We also excluded
W50 (G39.7−2.0) and RCW 89 (G20.4−1.2), since whether these objects are SNRs or not
is controversial.
Figure 7.1 shows the resultant Σ–D relation. SNRs listed in Table 7.1 are designatedwith the circles and crosses depending on whether synchrotron X-rays have been detected
or not, respectively. SNRs emitting synchrotron X-rays are Kepler’s SNR (Cassam-Chenaı̈
et al. 2004), Cas A (Vink & Laming 2003), Tycho’s SNR (Hwang et al. 2002), RCW 86(Bamba et al. 2000; Borkowski et al. 2001b), and SN 1006 (Koyama et al. 1995). Although7.1. PROPERTIES OF SNRS EMITTING SYNCHROTRON X-RAYS

107

Table 7.1: Shell SNRs with Known Distances
Remnant
G4.5+6.8
G6.4−0.1
G11.2−0.3
G18.8+0.3
G21.5−0.9
G27.4+0.0
G33.6+0.1
G34.7−0.4
G43.3−0.2
G49.2−0.7
G53.6−2.2
G55.0+0.3
G74.0−8.5
G84.2−0.8
G89.0+4.7
G93.3+6.9
G93.7−0.2
G109.1−1.0
G111.7−2.1
G114.3+0.3
G116.5+1.1
G119.5+10.2
G120.1+1.4
G132.7+1.3
G166.0+4.3
G166.2+2.5
G189.1+3.0
G205.5+0.5
G260.4−3.4
G263.9−3.3
G292.0+1.8
G292.2−0.5
G296.8−0.3
G315.4−2.3
G327.4+0.4
G327.6+14.6
G332.4−0.4
G337.0−0.1
G348.5+0.1
G348.7+0.4
G349.7+0.2

Other name
Kepler’s SNR
W28···
Kes 67···
Kes 73
Kes 79
W44
W49B
W51C
3C400.2···
Cygnus Loop···
HB21
DA 530
CTB 104A
CTB 109
Cas A
···
···
CTA 1
Tycho’s SNR
HB3
VRO 42.05.01
OA 184
IC443
Monoceros
Pup A
Vela (XYZ)
MSH 11-54···
1156−62
RCW 86
Kes 27SN 1006
RCW 103
CTB 33
CTB 37A
CTB 37B···

Distancea
(kpc)
2.9
1.9
4.4
14
4.6
6.8
7.8
2.8
10
6.0
2.8
14
0.4
4.5
0.8
2.2
1.5
3.0
3.4
0.7
1.6
1.4
2.3
2.2
4.5
8.0
1.5
1.6
2.2
0.3
6.0
8.4
9.6
2.3
4.8
2.2
3.1
11
8.0
8.0
15

Sizea
(0 )
3
48
4
14
4
4
10
31
3.5
30
30
17
190
18
100
23
80
28
5
70
70
90
8
80
44
80
45
220
55
255
10
17
17
42
21
30
10
1.5
15
17
2.2

a Cited from Green (2004).

Db
2.5
27
5.1
57
5.4
7.9
23
25
10
52
24
69
22
24
24
15
35
24
4.9
14
32
37
5.4
51
58
180
20
100
35
22
17
42
47
28
29
19
9.0
4.8
35
40
9.5

Σc
3.2 × 10−19
2.0 × 10−20
2.1 × 10−19
2.5 × 10−20
5.6 × 10−20
5.6 × 10−20
3.3 × 10−20
3.6 × 10−20
4.7 × 10−19
2.7 × 10−20
1.3 × 10−21
2.6 × 10−22
8.6 × 10−22
5.1 × 10−21
3.1 × 10−21
2.6 × 10−21
1.5 × 10−21
3.8 × 10−21
1.6 × 10−17
1.8 × 10−22
3.5 × 10−22
6.7 × 10−22
1.3 × 10−19
1.1 × 10−21
5.4 × 10−22
2.7 × 10−22
1.2 × 10−20
5.0 × 10−22
6.5 × 10−21
4.1 × 10−21
2.4 × 10−20
3.6 × 10−21
4.7 × 10−21
4.2 × 10−21
1.0 × 10−20
3.2 × 10−21
4.2 × 10−20
1.0 × 10−19
4.8 × 10−20
1.4 × 10−20
6.2 × 10−19

nd0
···
0.1
···
···
···
0.8
0.2
0.08
0.5
0.14
0.01
···
0.5
···
0.04
0.01
···
0.16
4.5
···
···
0.04
0.4
0.015
0.01
···
1.0
0.003
1
···
0.51
···
···
variese
0.1
0.1
2
···
···
···
1.0

Reference···
Rho & Borkowski (2002)
···
···
···
Gotthelf & Vasisht (1997)
Sun et al. (2004)
Harrus et al. (1997)
Kawasaki et al. (2002)
Koo et al. (2002)
Yoshita et al. (2001)···
Miyata et al. (1994)···
Leahy & Aschenbach (1996)
Landecker et al. (1999)···
Sasaki et al. (2004)
Willingale et al. (2003)······
Slane et al. (1997)
Seward et al. (1983)
Leahy et al. (1985)
Burrows & Guo (1994)···
Kawasaki et al. (2002)
Leahy et al. (1986)
Petre et al. (1982)···
Gonzalez & Safi-Harb (2003)
···
···
···
Enoguchi et al. (2002)
Bamba et al. (2003b)
Nugent et al. (1984)
···
···
···
Lazendic et al. (2004b)b Physical diameter (pc).
c Radio surface brightness at 1 GHz (W m−2 Hz−1 sr−1 ).
d Ambient medium density (cm−3 ) cited from the references. We adopted those calculated from the emission measures ofthe forward-shock plasmas. We assumed the forward-shock plasma is 4 times compressed than the ambient medium.
0 is large (e.g., Vink, Kaastra, & Bleeker 1997).

e Spatial variation of n

108

CHAPTER 7. DISCUSSION

Table 7.2: Shell SNRs Emitting Synchrotron X-rays
Remnant
G28.6−0.1
G32.45+0.1
G266.2−1.2
G347.3−0.5

Other name
A J1843.8−0352···
RX J0852.0−4622
RX J1713.7−3946

Distance
(kpc)
7
17
<1
1.0

Size
(0 )
11
8
120
60

a Physical diameter (pc).

Da
22
40
< 30
17

Σb
1 × 10−21
4 × 10−22
5 × 10−22
2 × 10−22 c

Reference
this thesis
this thesis
Duncan & Green (2000)
Lazendic et al. (2004a)b Radio surface brightness at 1 GHz (W m−2 Hz−1 sr−1 ).
c The surface brightness of G347.3−0.5 as a whole was estimated from Fig. 6 in Lazendic et al. (2004a), regarding thelowest contour level (0.5 mJy beam−1 ) is the average surface brightness of G347.3−0.5.

the scatter is large, we can see a trend that Σ becomes smaller as D becomes larger. When
we fitted the Σ–D relation in Figure 7.1 with a power-law function, the best-fit function was,
Σ1 GHz = 3.6 × 10−18 × D −1.9 W m−2 Hz−1 sr−1 ,(7.1)

which is shown with the solid line in Figure 7.1. In this fitting, SNRs listed in Table 7.1(crosses and circles) were included.
We next added SNRs emitting synchrotron X-rays, both of those already known andthose we found, which are shown in Table 7.2. These SNRs are designated with the squares
in Figure 7.1. The distances of G28.6−0.1 and G32.45−0.1 are estimated from the X-ray
absorption column density (§6.2.3 and 6.3.2). Just as expected, SNRs emitting synchrotron
X-rays are mainly located in a region of small Σ. However, three sources (Cas A, and Tycho’s
and Kepler’s SNRs) are located in a region of large Σ.
On the ground that diameters approximately show the ages of SNRs, we can summarize
this situation as follows: when SNRs are young, most of them emit synchrotron X-rays. As
they evolve, some SNRs faint in radio continue to emit synchrotron X-rays, whereas theothers stop the emission. Therefore, it is suggested that the faintness in the radio band isconnected with the lifetime of synchrotron X-ray emission in each SNR.
Then what determines the radio surface brightness other than the diameter. The powerof synchrotron radiation is proportional to the product of the magnetic field strength andthe electron density, and furthermore, the magnetic field strength is known to correlate withthe ambient density. Therefore, we can first suspect ambient density is responsible for the
scatter of Σ. In order to test whether the scatter of Σ is due to ambient densities, we derived
the ratios between Σs expected by the Σ–D relation (equation 7.1) and the observed Σs,and plotted them against ambient densities. Ambient densities can be derived from X-ray

7.1. PROPERTIES OF SNRS EMITTING SYNCHROTRON X-RAYS

109

Cas A

Monoceros

Cygnus Loop

Figure 7.2: The ratio between Σs expected by the Σ–D relation and observed Σs plottedagainst ambient densities.

observations; ambient densities are cited from literature as shown in Table 7.1. We here
picked up only those derived from X-ray luminosities.
Figure 7.2 shows the resultant plot. Although the scatter is large, we can see a positive
correlation between ambient densities and the ratios of observed Σs against the Σ–D relation.
The correlation index is calculated to be 0.41, and supports the positive correlation. This
result confirms that one of the causes for the scatter of Σ is ambient densities of SNRs.
Therefore, we can conclude that SNRs which emit synchrotron X-rays for longer time arelikely to be located in low density regions.
If we assume the electrons are accelerated by the standard diffusive shock acceleration(e.g., Drury 1983; see Appendix C), we have a relation between the cut-off frequency ofsynchrotron emission νcutof f and the shock velocity V such as:νcutof f ≈ 5 × 10

17



V
2000 km s−1

2

ξ −1 Hz,(7.2)

where ξ is the gyrofactor. This relation shows that whether electron synchrotron emission
is extending up to the X-ray band or not is determined only by the velocity of the shock(ξ is assumed to be ∼ 1 and rather constant even for various SNRs). Since SNRs expand
sweeping up the ambient matter, in order to maintain the shock velocity high for longer
time, the ambient density must be low. For example, in order to maintain the shock velocity
V larger than 2000 km s−1 until the shell radius is 20 pc, on assumption the SNR is in the

110

CHAPTER 7. DISCUSSION

Sedov phase, the ambient density n0 must be:
E0
n0 < 0.13
51
10 erg!

D
20 pc!−3/5 V
2000 km−2/5

cm−3 ,(7.3)

where E0 , D are the initial explosion energy of the supernova, and the diameter, respectively.
Thus, our hypothesis that SNRs emitting synchrotron X-rays longer time are located in low
density region is also supported theoretically.
Our result predicts that small SNRs or radio faint SNRs emit synchrotron X-rays. Thismay be the case for SNR 0509−67.5, a shell-type SNR in the Large Magellanic Cloud with
7 pc-diameter, in which Warren & Hughes (2004) suggested existence of synchrotron X-rays
after a detailed analysis of a Chandra observation. However, it is clear that not all the SNRs
with small diameters or not all the SNRs faint in radio emit synchrotron X-rays. We suspect
there are other causes for the scatter of Σ, for example, the explosion energies of supernovae.
Since the number of accelerated electrons is expected to be less in supernova remnants of low
energy, brightnesses in radio wwill be small in such SNRs, which is analytically discussed
by Berezhko & Völk (2004). However, variation of explosion energies of supernova remnants
are not well understood yet.
Our result supports that the shock speed is an important parameter to determine
whether we can expect synchrotron X-ray emission from an SNR. Since the thermal energyof the shocked plasma is a conversion of the kinetic energy of the shock, the temperature of
the plasma would indicate the shock speed. However, the temperatures of ions and electrons
may not have been equilibrated in such thin plasmas in SNRs. For example, though the
speed of the north-west shock of SN 1006 (∼2900 km s−1 ; Winkler, Gupta, & Long 2003)
suggests the temperature of the plasma behind the shock is ∼ 10 keV, the electron temperature of the plasma measured to be only ∼ 0.9 keV (Long et al. 2003). The measurement of
the ion temperatures will be possible in the near future; high energy resolution spectroscopy
with such instruments as XRS onboard Astro-E2 will enable us to measure thermal Doppler
broadening of the emission lines from ions. Hα observations have also been determining the
speed of SNR shocks (e.g., Chevalier et al. 1980). However, not only SNRs do not always
emit Hα lines but also the regions accompanying synchrotron X-ray emissions are often weakin Hα emissions (e.g., Rho et al. 2002). Therefore, future X-ray observations with high energy resolution would be important in studying the relation between shockk speed and X-ray
synchrotron emission.
Protons and ions are expected to be accelerated simultaneously with electrons. Since
synchrotron energy loss is negligible for them, the maximum energy is only determined by
the acceleration time scale. Under the diffusive shock acceleration theory, the acceleration

111

7.2. PROPERTIES OF THE SNR CANDIDATES

time scale is proportional to V −2 (see Appendix C). Since shock speeds of SNRs in lower
density medium is maintained large for longer time, those SNRs may be responsible for the
maximum energy of galactic cosmic rays. According to the discussion above, those SNRs
are likely to be faint in radio, and may be found with X-ray observations.

7.2

Properties of the SNR candidates

Although the number of SNRs emitting synchrotron X-ray emission have been small, ourdiscovery of new SNRs and candidates indicate that many such SNRs are located in thegalactic plane, and that those are likely to be discovered with X-ray observations.
Assuming the brightness of SN 1006, we estimate the deepness of our search with thedata set of the ASCA Galactic Plane Survey. In order to compare brightnesses and sizesof the new extended X-ray sources with those of SN 1006, we plotted them in Figure 7.3.
The solid-line corresponds to the same brightness as SN 1006, and corresponds to the casethat SN 1006 were located more distant. From this plot, we see that we are detecting
sources dimmer than 1/2 brightness of SN 1006. If we exclude the galactic center region,
the maximum absorption column density within the galaxy is ∼ 5 × 1022 cm−2 (Ebisawa
et al. (2001)). With this absorption column density, 50% of X-ray photons are absorbed in
the 2.0–10.0 keV band. However, we can detect sources with dimmer then 1/2 brightnessof SN 1006, the sensitivity is enough to detect a distant “SN 1006”. On the other hand, adistant “SN 1006” becomes small, the spatial resolution is required to detect the extension.
Since the resolution of ASCA corresponds to 60 , “SN 1006” at a distance of 10 kpc is thedetection limit.
SNRs located at low longitude are likely to originate from Type II supernovae. Therefore, their distribution is expected to be similar to that of the H ii regions, and further thatof CO molecular cloud. In order to estimate coverage of the survey, we simply assume afinite disk model for the number density of SNRs η: 2  η exp − 1 |z|η(r, z) = 

0

0

2

Zs

(r <= Rd )(7.4)

(r > Rd )According to Sanders et al. (1984), the scale hight of CO (Zs) is ∼ 60 pc. We assume
Rd = 5 kpc, at which radius the galactocentric arms are located (Sanders et al. 1984).
We found 37% of the putative SNRs are located in the region covered by the ASCA
survey (|l| < 45◦ , |b| < 0.4◦ , distance <10 kpc). However, 34% of the survey region

112

CHAPTER 7. DISCUSSION

Spatial Resolution

x2(4 kpc)

SN1006x 1/2(8 kpc)Figure 7.3: Plot of diameters and fluxes for the new extended X-ray sources. SN 1006 is
designated with the star, new synchrotron X-ray SNRs, G28.6−0.1 and G32.45+0.1 are withthe circles, and other sources are with the triangles. The solid-line shows the points where
the surface brightness is the same as SN 1006. If SN 1006 were located at the distance inthe parentheses, the diameter and the brightness have the values at the point.

is suffered from stray light and could no be used for extended-source search. Therefore,
1/0.37/(1 − 0.34) ∼4 times the number of SNRs detected with the ASCA survey is expectedin the galactic plane. In the ASCA survey region, at least 3 SNRs emitting synchrotron
X-rays (G347.3−0.5, G28.6−0.1, and G32.45+0.1) are located. Among our 15 candidates,
G26.6−0.1, G316.8−0.1, and G318.6+0.0 are not likely to be SNRs. Therefore, the numberof SNRs emitting synchrotron X-rays in the survey is at most 12 + 1 (G347.3−0.5)= 13.
Therefore, we expect 12∼52 SNRs emitting synchrotron X-rays in the galactic plane.
The discussion in §7.1 indicate that SNRs emitting synchrotron X-rays are located in
low-density regions. This seems to be conflicting with that our new SNRs and candidates arelocated on the low Galactic latitude, where the most of the mass in the galaxy is concentrated.
One plausible situation is that the progenitor of these SNRs are exploded in cavities made bystellar wind. One example is recently demonstrated for G347.5−0.5 with a CO observation
by Fukui et al. (2003). They argue G347.3−0.5 exploded in a cavity and is expanding athigh velocity ∼5000 km s−1 . RCW 86 is also suggested to be an SNR exploded in a cavity
by Vink, Kaastra, & Bleeker (1997). G28.6−0.1 may be another example for the cavity
explosion. We found a strong anti-correlation between the ASCA X-ray image aand the CO
map as shown in Figure 7.4. This anti-correlation is not likely to be by chance, because thedistance determined with X-ray is consistent with the kinematic distance with the CO-line;7.2. PROPERTIES OF THE SNR CANDIDATES

113

Figure 7.4: ASCA 0.7–7.0 keV image around G28.6−0.1 overlaid with the contours of the
CO intensity map (courtesy of Y. Moriguchi, the NANTEN group). The CO intensity isderived by integrating the CO spectra from 80 to 90 km s−1 . The contours are plotted from30 K km s−1 with intervals of 5 K km s−1 .

the former is 7 kpc (see §6.2.1) and the latter is 7.5 kpc.

Appendix A
Background Light Curves of
XM M -N ewton
All the background light curves of the XM M -N ewton observations appearing in this thesis
are accumulated here. These lightcurves are created over the whole detector for signle
events in the energy range 10–15 keV and with FLAG values as defined by #XMMEA EP
or #XMMEA EM depending on the PN or MOS cameras.

115

116

APPENDIX A. BACKGROUND LIGHT CURVES OF XM M -N EW T ON

0148850201 (G156.2+5.7 center)0148850101 (G156.2+5.7 north-west)

0205970201 (G156.2+5.7 north-west)

PN

PN

PN

MOS1

MOS1

MOS1

MOS2

MOS2

MOS2

Figure A.1: Background light curves of XM M -N ewton observations (1).

117

0112930101 (G272.2-3.2)

0112890101 (G299.2-2.9)0136030101 (G32.45+0.1)

PN

PN

PN

MOS1

MOS1

MOS1

MOS2

MOS2

MOS1

Figure A.2: Background light curves of XM M -N ewton observations (2).

Appendix B
Serendipitous Discovery of a Clusterof Galaxy with XM M -N ewton
In the XM M -N ewton observation of G156.2+5.7, we serendipitously discovered a new cluster of galaxies. In this section, we report on the discovery.
This source was detected in the MOS1+2 image of the 2.0–7.0 keV band using theemldetect software. Figure B.1 shows the 0.5–7.0 keV image around the source. We cansee a diffuse emission centering on at R.A.=04h 56m 37.s 2, Dec.=+52◦ 240 1200 (J2000). We will
designate this source as XMMJ J0456372+522412 (or XMMJ J0456372) after its position.
In order to demonstrate the extension of XMMJ J0456372, we made a radial profilearound XMMJ J0456372, excluding the nearby point sources detected with the emldetect
software (dotted-line circles in Figure B.1). Figure B.2 shows the result. For comparison,
we plotted a radial profile of a point source in the same FOV (source 2 in §4.1.3). Since
off-axis angles of XMMJ J0456372 and source 2 are 70 and 100 , respectively, the point spreadfunction at XMMJ J0456372 is sharper than the radial profile of source 2. Thus we can see
clearly that the X-ray emission of XMMJ J0456372 is extended about ∼ 45 00 .
Source and background regions for the X-ray spectra were chosen to be a circular regionof 5000 and a surrounding annular region with inner and outer radii of 5000 and 10000 , respectively (solid-line circes in Figure B.1). The background-subtracted PN and MOS spectra areshown in Figure B.3. The MOS1 and 2 spectra are combined after checking for the consistency. The spectra show a line-like structure around 5.5 keV, and a wavy structure around
1–2 keV, suggesting the X-ray emission is thermal. Since no abundant element emit X-ray
line at 5.5 keV, the line is likely to be a redshifted Fe-K line. Accordingly, the extended119

120APPENDIX B. SERENDIPITOUS DISCOVERY OF A CLUSTER OF GALAXY WITH XM M -N EW T
X-ray emission is suspected to be thin thermal plasma emission from a cluster of galaxies.
We tried fitting the spectra with a thermal plasma model (APEC in XSPEC v11.3; Smith etal. 2001) allowing the redshift vary. As a result, we obtained an acceptable fit, with best-fit
parameters and models shown in Table B.1 and Figure B.3, respectively. The absorption column density is consitent with the total galactic H i column density in this direction (Dickey& Lockman 1990).
If we use a hubble constant H0 = 50 h50 km s−1 Mpc−1 , the spatial extent of 4500 radius
at a redshift of 0.21 corresponds to a physical diameter of 0.7 h−1
50 Mpc.
−1
The luminosity is calculated to be 1.0×1044 h−3
50 erg s . According to the LX –T relation (Arnaud & Evrard 1999), the temperature of 3.2 keV expects the luminosity to be
1.9 × 1022 erg s−1 , which agrees well with the observed value. This supports further thatXMMJ J0456372 is a cluster of galaxies.

When we searched for counterparts in other wavelength using the NED data system, wefound a 2MASS extended source, 2MASX J04563758+5224148, within the extended emissionof XMMJ J0456372 (see Figure B.4). This infrared object must be a member galaxy of thecluter, XMMJ J0456372+522412, especially a cD galaxy resting in the center of the cluster.
Table B.1: Best-fit parameters of XMMJ J0456372+522412 with a APEC model.
kT
NH
Fluxa
(keV)
redshift(cm−2 )
(erg cm−2 s) χ2 /d.o.f.
3.2 (3.0–3.7) 0.21 (0.15–0.23) 3.9 (3.2–5.0)×1021
a

Absorption-corrected flux in the 0.4–10.0 keV.

3.5×10−13

65.0/67

121

420

18
+52°26′source161412Brightness

410

10
9
400
8
7
390
6

10

380
5
4
370

+52°24′

8

3
360

6

350

4

2

340
330
1

background+52°22′hms

04 56 50hms

04 56 40hms

04 56 30240 250 260 270 280 290 300 310 320 330
BLOW mos12_pat12_flag0_E10007000_bin64.img0

FigureB.1:
MOS1+2
imageof
XMMJ J0456372+522412 in the 1.0–7.0 keVband. The PN and MOS1+2 spectra shownwith gray and black, respectively. The spectral
regions (source and background) are designatedwith solid-line circles. The regions aroundnearby point sources are excluded with thedotted-line circles from the spectral regions.

2
0

0

10

20

30 40 50 60
Radius (arcsec)

70

80

FigureB.2:
Radialprofileof
XMMJ J0456372+522412 in the 1.0–7.0 keVbands (crosses). The vertical acis showsbrightness (×10−5 cnts cm− 2 s−1 arcmin−2 ).
The contribution of nearby point sources
are excluded with 2000 - radius regions. For
reference, the radial profile of the point
source, “source 2” (§4.1.3), normalized at
the center is shown as the histogram.

122APPENDIX B. SERENDIPITOUS DISCOVERY OF A CLUSTER OF GALAXY WITH XM M -N EW T

Figure
B.3:
EPIC
spectraof
XMMJ J0456372+522412. The best-fit APEC
model is shown with the solid-line.

Figure B.4: 2MASS Ks-band imagearound XMMJ J0456372+522412 overlaid
with the MOS1+2 contours in the 1.0–
7.0 keV band. Both the image and thecontours are in logarithmic scale.

Appendix C
Diffusive Shock Acceleration
Since particles are accelerated by moving back and forth between the shocks, the acceleration
time scale is determined by a time which a particles takes for one cycle ∆t and a momentum
gain per cycle ∆p.
The time scale for a particle to complete one cycle is a sum of the mean residence times
in the upstream and the downstream regions of each particle and given as
∆t =

κ2
4 κ1
+
,
v U1 U2




(C.1)

where U1 (U2 ) and κ1 (κ2 ) are the fluid velocity and the diffusion coefficient in the upstream(downstream) region, respectively (Drury 1983). On the other hand, the mean momentum
gain per cycle is given as
4 U1 − U 2
∆p =
p,
(C.2)
3
v
Then, the acceleration time scale is3
p∆t=tacc =∆p
U1 − U 2



κ1
κ2
+
.
U1 U2


(C.3)

For the strong shock, the compression ratio can be assummed to be 4, and thus U2 = 41 U1 .
Then the acceleration time scale can be written simply astacc =

4(κ1 + 4κ2 )
.
U12

(C.4)

Therefore, the diffusion coeffecients κ1 , κ2 are essential for the acceleration time scale (justas the name of the theory shows). The diffusion coefficient can be written as κ = 13 λD c,where λD is the mean free path of particle scattering. Since the particles are diffused by the
123

124

APPENDIX C. DIFFUSIVE SHOCK ACCELERATION

interaction with the magnetic field, the mean free path λD is thought to be longer than the
Lamour radius rL = E/ZeB, and usually written as
λD = ξrL(C.5)

with the so-called gyrofactor ξ, which is always larger than 1. Using ξ, we can write thediffusion coefficient as
crL
ξ.
(C.6)
κ=
3
Then, if we take κ1 = κ2 for simplicity, equations (C.4) and (C.5) yield
tacc =

20 crL
ξ.
3 V2

(C.7)

where V is the shock velocity (= U1 ).
For electrons, the synchrotron energy loss is significant and the maximum energy is likely
to be determined by the balance between the acceleration time scale and the synchrotron
times scale. From equation (2.37), the synchrotron loss time scale istsynch = 1.25 × 10

3



Ee
100 TeV

−1

B
10 µG

!−2

yr.

(C.8)

The maximum energy of accelerated electrons Ee is determined under the condition of tacc =tsynch , and given as,
B
Ee ≈ 67
10 µG!−1/2 V
ξ −1/2 yr.
2000 km s−1


(C.9)

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Acknowledgement
I am deeply grateful to Prof. K. Koyama for his continuous guidance and support throughout
the five years of my graduate school period. I also thank my collaborators, Prof. S. Yamauchi,
Dr. K. Ebisawa, Dr. A. Bamba, Mr. H. Nakajima, and Mr. H. Yamaguchi for their helpful
and constructive discussions.
I am grateful to all the members of the Cosmic-ray laboratory in Kyoto University for
their continuous discussions, encouragements, and technical advice. I thank Prof. T. Kifune,
Prof. M. Mori and all the members of the Cangaroo Team for the constructive discussions.
I would like to thank all the members of the ASCA team. The longer I get involved
the X-ray astronomy, the more I recognize how the background of ASCA was low and how
it is important.

133