X-ray Study of Stellar Winds with Suzaku
Yoshiaki Hyodo
Department of Physics, Graduate School of Science, Kyoto University
Kitashirakawa Oiwake-cho, Sakyo-ku, Kyoto, 606-8502, Japanhyodo@cr.scphys.kyoto-u.ac.jp

This thesis was submitted to the Department of Physics,
Graduate School of Science, Kyoto University
on January 5 2009in partial fulfillment of the requirements
for the degree of Doctor of Philosophy in physics.

Abstract
Stellar winds from massive stars play an important role for dynamical and chemical
evolution of galaxies. To evaluate their effects in a galaxy scale, both the total number andthe properties of massive stars should be understood. Most massive stars have soft X-ray
emission, while some have additional hard emission. Although it is indicated that the binarity
is important, the condition for hard X-ray production has not yet been elucidated. Also,
some massive star clusters exhibit diffuse (pc scale) X-ray emission, where stellar winds may
be depositing energy to the interstellar space. However, previous studies were not able toclarify the origin of the diffuse emission. It is even unclear whether the emission is physicallyrelated to the star cluster or not.
The X-ray Imaging Spectrometer (XIS) on board Suzaku has a good energy resolution,a low and stable background, and a large effective area, which are suited for narrow-bandsurveys and spectroscopy of both point-like and extended X-ray emission. Using the XIS,
we performed a very deep survey in the direction of the central region of our Galaxy andpointing observations of three well-known star forming regions in our Galaxy (M 17, theCarina Nebula, and the Arches cluster).
A bright point-like source (CXOGC J174645.3–281546) was found in the 6.7 keV intensity map of the Galactic center region. The temperature and luminosity are similar
to those of wind-collision binary. The spectral energy distribution in the infrared is very
similar to those of eponymous Quintuplet cluster members, indicating that the source is a
WC star+OB star binary.
We also present the results of X-ray spectroscopy of nine massive stars in M 17, theCarina nebula, and the Arches cluster. Many sources show hard X-ray emission. We obtained
important clues for hard X-ray production mechanism in massive star binaries. WR 25(WN6+O4f) in the Carina region shows hard X-ray emission with a flux variation correlatedwith the binary orbital phase, while the close binary HD 93205 (O3+O8) lacks hot component
and is stable. These two facts indicate that the binary separation is a key parameter for
hard X-ray emission. We speculate that the other two massive stars OI 352 in M 17, andHD 93250 in the Carina Nebula with hard X-ray emission are also colliding-wind binaries.
A variety of diffuse emission is apparently associated with some massive star clusters.
We argue that the diffuse source in M 17 and in the west to η Carinae are prototypicalstellar wind bubbles; the absorption columns typical to these regions indicate that the diffuse<
emission is physically associated with the star clusters and the O/Fe abundance ratio ( ∼1 solar) is too low for supernova remnants caused by core-collapse type supernovae. On theother hand, the diffuse emission in the south side of η Carinae is unlikely to have physicalrelation with the Carina Nebula, because the absorption column is very low (N H < 6 ×31020 cm−2 ). A possible nature is a foreground type Ia supernova remnant. The absorption
column of diffuse emission in the Arches cluster is consistent with that of the Galactic center,indicating a physical relation to the Arches cluster. Although massive stars in the Arches
cluster are the sources of strong stellar wind, the diffuse emission is unlikely to be hot plasma
created by the stellar wind. It shows a strong line from neutral iron. The origin is likely
due to X-ray and/or electron fluorescent, but the flux of the X-ray photons or electrons are
extremely insufficient to produce the diffuse X-rays.

Contents
1 Introduction

15

2 Review

17

2.1

Evolution of Massive Stars . . . . . . . . . . . . . . . . . . . . . . . . . . . .

17

2.2

Wolf-Rayet Stars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

21

2.3

Diffuse X-ray Emission from Massive Stars . . . . . . . . . . . . . . . . . . .

24

2.4

Point-like X-ray Emission from Massive Stars . . . . . . . . . . . . . . . . .

27

3 Instruments

31

3.1

Overview of Suzaku . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

32

3.2

XRT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

34

3.2.1

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

34

3.2.2

Onborad Performance . . . . . . . . . . . . . . . . . . . . . . . . . .

35

XIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

38

3.3.1

Overview of the XIS . . . . . . . . . . . . . . . . . . . . . . . . . . .

38

3.3.2

Onboard Performance . . . . . . . . . . . . . . . . . . . . . . . . . .

40

HXD . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

44

3.4.1

Overview of the HXD . . . . . . . . . . . . . . . . . . . . . . . . . . .

44

3.4.2

Onboard Performance . . . . . . . . . . . . . . . . . . . . . . . . . .

46

3.3

3.4

4 Observations & Results

49
5

6

CONTENTS
4.1

4.2

4.3

4.4

A New Wolf-Rayet Star in the Galactic Center . . . . . . . . . . . . . . . . .

51

4.1.1

Objectives . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

51

4.1.2

Observations & Reduction . . . . . . . . . . . . . . . . . . . . . . . .

52

4.1.3

Band-limited Images & Source list

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

54

4.1.4

A Discrete source . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

55

Carina Nebula . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

59

4.2.1

Objectives . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

59

4.2.2

Observations & Reduction . . . . . . . . . . . . . . . . . . . . . . . .

61

4.2.3

Band-limited Images & Source list

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

61

4.2.4

Discrete Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

62

4.2.5

Diffuse Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

70

M 17 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

73

4.3.1

Objectives . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

73

4.3.2

Observations & Reduction . . . . . . . . . . . . . . . . . . . . . . . .

76

4.3.3

Band-limited Images & Source list

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

76

4.3.4

Discrete sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

80

4.3.5

Diffuse sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

87

4.3.6

HXD-PIN Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . .

90

Arches Cluster. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

92

4.4.1

Objectives . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

92

4.4.2

Observations & Reduction . . . . . . . . . . . . . . . . . . . . . . . .

94

4.4.3

Band-limited Images & Source list

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

94

4.4.4

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

95

5 Discussion
5.1

101

A New Wolf-Rayet Star in the Galactic Center . . . . . . . . . . . . . . . . . 102
5.1.1

Chandra and XMM-Newton Counterparts . . . . . . . . . . . . . . . 102

CONTENTS

7

5.1.2

Longer Wavelength Data . . . . . . . . . . . . . . . . . . . . . . . . . 105

5.1.3

Location of Extinction Matter . . . . . . . . . . . . . . . . . . . . . . 106

5.1.4

Nature of the Source . . . . . . . . . . . . . . . . . . . . . . . . . . . 107

5.2

5.3

5.4

Carina Nebula . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110
5.2.1

Discrete Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110

5.2.2

Diffuse Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 112

5.2.3

The Plasma Properties . . . . . . . . . . . . . . . . . . . . . . . . . . 112

M 17 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 113
5.3.1

Discrete Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 113

5.3.2

Diffuse Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 117

5.3.3

HXD signals . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 122

Arches Cluster
5.4.1

Discrete Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 124

5.4.2

Origin of Three Spectral Components . . . . . . . . . . . . . . . . . . 130

5.4.3

Diffuse Medium . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 131

5.4.4

Cause of Line and Power-law Emission . . . . . . . . . . . . . . . . . 131

6 Conclusions
6.1

6.2

6.3

. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 124

135

A New Wolf-Rayet Star in the Galactic Center . . . . . . . . . . . . . . . . . 135
6.1.1

Summary of the Results . . . . . . . . . . . . . . . . . . . . . . . . . 135

6.1.2

Conclusion — Contribution to the Galactic Center Diffuse Emission . 136

Discrete Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 137
6.2.1

Summary of the Results . . . . . . . . . . . . . . . . . . . . . . . . . 137

6.2.2

Conclusion — Mechanism for the Hard X-rays from Massive Stars . . 138

Diffuse Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 140
6.3.1

Summary of the Results . . . . . . . . . . . . . . . . . . . . . . . . . 140

8

CONTENTS
6.3.2

Conclusion — Various origins of Diffuse Emission in Star-forming
Regions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 141

List of Figures
2.1

Diagram of CNO cycle. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

18

2.2

Energy production rate of the pp chain and CNO cycle. . . . . . . . . . . . .

19

2.3

HST image of M1-67 around WR 124 . . . . . . . . . . . . . . . . . . . . . .

21

2.4

WR158 spectrum in the NIR bands. . . . . . . . . . . . . . . . . . . . . . . .

22

2.5

WR star surface number density as a function of Galactocentrinc distance . .

23

2.6

Velocity structure of line-driven stellar winds. . . . . . . . . . . . . . . . . .

24

2.7

Temperature and density structure of a stellar wind bubble. . . . . . . . . .

25

2.8

Color composite image of a stellar wind bubble around WR 136 . . . . . . .

26

2.9

Bolometric luminosity and X-ray luminosity in the 0.2–4.0 keV band. . . . .

27

2.10 Structure of stellar wind. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

28

2.11 X-ray luminosity history of WR 140.

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

29

2.12 Temperature map of θ 1 Ori C. . . . . . . . . . . . . . . . . . . . . . . . . . .

30

3.1

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

32

3.2

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

33

3.3

Picture of the XRT-I1 module. . . . . . . . . . . . . . . . . . . . . . . . . . .

34

3.4

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

35

3.5

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

36

3.6

Image of SS Cyg at the optical axis taken with XRT-I2. . . . . . . . . . . . .

36

3.7

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

37

9

10

LIST OF FIGURES
3.8

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

38

3.9

Difference of front illumination and back illumination . . . . . . . . . . . . .

39

3.10 Schematic view of the XIS CCD. . . . . . . . . . . . . . . . . . . . . . . . .

40

3.11 QE of the XIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

41

3.12 Long-term trend of energy resolution. . . . . . . . . . . . . . . . . . . . . . .

41

3.13 Long-term trend of on-axis contaminant column density . . . . . . . . . . . .

42

3.14 XIS spectra obtained during night earth observations.

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

43

3.15 History of the observing condition of the XIS. . . . . . . . . . . . . . . . . .

43

3.16 Picture of the HXD . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

44

3.17 Effective areas of PIN and GSO . . . . . . . . . . . . . . . . . . . . . . . . .

45

3.18 Schematic view of a well-counter unit . . . . . . . . . . . . . . . . . . . . . .

45

3.19 Angular response of PIN . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

46

3.20 NXB spectra of the PIN sorted with COR. . . . . . . . . . . . . . . . . . . .

47

4.1

XIS image of the Galactic center region in the 6.55–6.8 keV band . . . . . .

53

4.2

XIS images of Sgr B North region. . . . . . . . . . . . . . . . . . . . . . . . .

55

4.3

XIS light curve of CXOGC J174645.3–281546. . . . . . . . . . . . . . . . . .

56

4.4

XIS spectra of CXOGC J174645.3–281546. . . . . . . . . . . . . . . . . . . .

57

4.5

2MASS image of the Carina Nebula . . . . . . . . . . . . . . . . . . . . . . .

59

4.6

Carina Nebula image taken with Einstein . . . . . . . . . . . . . . . . . . . .

60

4.7

XIS images of Carina Nebula. . . . . . . . . . . . . . . . . . . . . . . . . . .

62

4.8

Three color image of the Carina Nebula in the soft band. . . . . . . . . . . .

63

4.9

XIS light curves of WR 25 . . . . . . . . . . . . . . . . . . . . . . . . . . . .

64

4.10 XIS light curves of HD 93205 . . . . . . . . . . . . . . . . . . . . . . . . . . .

64

4.11 XIS light curves of HD 93250 . . . . . . . . . . . . . . . . . . . . . . . . . . .

65

4.12 XIS spectra of WR25. Only FI spectra are shown for simplicity. The lowerand higher temperature components are shown in red and blue, respectively.

65

11

LIST OF FIGURES
4.13 XIS spectra of HD 93205. The summed FI and BI spectra are shown in blackand red, respectively. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

66

4.14 XIS spectra of HD 93250. Only FI spectra are shown for simplicity. The lowerand higher temperature components are shown in red and blue, respectively.

66

4.15 Merged XIS spectrum of HD 93205 . . . . . . . . . . . . . . . . . . . . . . .

67

4.16 Merged XIS spectrum of HD 93250 . . . . . . . . . . . . . . . . . . . . . . .

68

4.17 Merged XIS spectrum of WR 25 . . . . . . . . . . . . . . . . . . . . . . . . .

68

4.18 ACIS spectrum of integrated point sources in the Diffuse 3 region. . . . . . .

70

4.19 XIS of diffuse emission in the Carina Nebula region.

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

72

4.20 2MASS image of M 17 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

73

4.21 Chandra/ACIS image and spectrum of M 17 . . . . . . . . . . . . . . . . . .

74

4.22 ROSAT image and spectrum of M 17 . . . . . . . . . . . . . . . . . . . . . .

75

4.23 Mosaic XIS image of the three Suzaku observations in the M 17 region. . . .

77

4.24 XIS images of the M 17 region. . . . . . . . . . . . . . . . . . . . . . . . . . .

78

4.25 XIS light curves of the OB association. . . . . . . . . . . . . . . . . . . . . .

80

4.26 XIS light curves of M 17 North. . . . . . . . . . . . . . . . . . . . . . . . . .

80

4.27 XIS light curves of OI 352 . . . . . . . . . . . . . . . . . . . . . . . . . . . .

81

4.28 XIS light curves of Suzaku J182132–1610.4. . . . . . . . . . . . . . . . . . . .

81

4.29 XIS light curves of Suzaku J182132–1613.4. . . . . . . . . . . . . . . . . . . .

81

4.30 XIS light curves of Suzaku J182140–1604.1. . . . . . . . . . . . . . . . . . . .

82

4.31 XIS spectra of point sources in M 17 region. . . . . . . . . . . . . . . . . . .

84

4.32 Background-subtracted spectra of three point sources in the M17 East field. .

85

4.33 Background-subtracted spectra of diffuse emission in the M 17 West region. .

88

4.34 PIN spectra of the M 17 region observations. . . . . . . . . . . . . . . . . . .

90

4.35 Background-subtracted PIN spectra of the M 17 regions . . . . . . . . . . . .

90

4.36 HST image of the Arches cluster . . . . . . . . . . . . . . . . . . . . . . . . .

92

4.37 XIS images of the Arches cluster. . . . . . . . . . . . . . . . . . . . . . . . .

95

12

LIST OF FIGURES
4.38 Band-limited XIS images of the Galactic center region. . . . . . . . . . . . .

96

4.39 XIS spectra of the Arches cluster. . . . . . . . . . . . . . . . . . . . . . . . .

97

4.40 XIS spectrum of the Arches cluster. . . . . . . . . . . . . . . . . . . . . . . .

98

5.1

EPIC spectra of CXOGC J174645.3–281546. . . . . . . . . . . . . . . . . . 103

5.2

Long-term X-ray flux trend of CXOGC J174645.3–281546. . . . . . . . . . . 104

5.3

NIR pseudo color image around CXOGC J174645.3–281546 . . . . . . . . . . 105

5.4

NIR SED of CXOGC J174645.3–281546. . . . . . . . . . . . . . . . . . . . . 106

5.5

Cartoon of the CXOGC J174645.3–281546 system. . . . . . . . . . . . . . . . 108

5.6

Long-term X-ray flux variation of WR 25. . . . . . . . . . . . . . . . . . . . . 110

5.7

SED of three point sources in the M17 east field. . . . . . . . . . . . . . . . . 113

5.8

Contaminating emission to the M 17 diffuse emission. . . . . . . . . . . . . . 118

5.9

Simultaneous fit of XIS and PIN spectra of M17 West . . . . . . . . . . . . . 122

5.10 Close-up ACIS view of the Arches cluster. . . . . . . . . . . . . . . . . . . . 125
5.11 ACIS spectra and best-fit models of point sources. . . . . . . . . . . . . . . . 127
6.1

Abundace of diffuse emission in M 17 and the Carina Nebula . . . . . . . . . 140

List of Tables
4.1

Observation log of the Sgr B North region. . . . . . . . . . . . . . . . . . . .

54

4.2

Point source list in the Sgr B North region . . . . . . . . . . . . . . . . . . .

54

4.3

Best-fit spectral parameters. . . . . . . . . . . . . . . . . . . . . . . . . . . .

58

4.4

Observation log of the Carina Nebula.

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

61

4.5

Point source list in the Carina Nebula region . . . . . . . . . . . . . . . . . .

61

4.6

Best-fit parameters of WR 25. . . . . . . . . . . . . . . . . . . . . . . . . . .

66

4.7

Best-fit parameters of HD 93205. . . . . . . . . . . . . . . . . . . . . . . . . .

67

4.8

Best-fit parameters of HD 93250. . . . . . . . . . . . . . . . . . . . . . . . . .

67

4.9

Best-fit APEC parameters for HD 93205, HD 93250, WR 25 . . . . . . . . . .

69

4.10 Best-fit parameters of the APEC model for the diffuse emission. . . . . . . .

71

4.11 Observation log around the M17 region.

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

76

4.12 Point source list of the M 17 region . . . . . . . . . . . . . . . . . . . . . . .

79

4.13 Best-fit APEC parameters for the discrete spectra.

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

83

4.14 Best-fit parameters of three point sources in the M 17 EAST region . . . . .

86

4.15 Best-fit parameters of the APEC model for the diffuse emission in the M 17
West field . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

89

4.16 Best-fit parameters of PIN spectra of the M 17 regions. . . . . . . . . . . .

91

4.17 Observation log of the Arches cluster. . . . . . . . . . . . . . . . . . . . . . .

94

4.18 Source list in the Sgr A East region . . . . . . . . . . . . . . . . . . . . . . .

94

4.19 Best-fit parameters for the Arches cluster . . . . . . . . . . . . . . . . . . . .

99

13

14

LIST OF TABLES
5.1

Observation log of the Sgr North region. . . . . . . . . . . . . . . . . . . . . 102

5.2

Best-fit spectral parameters. . . . . . . . . . . . . . . . . . . . . . . . . . . . 103

5.3Counterparts of three unidentified point sources in the M17 East field . . . . 116

5.4

Observation Log . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 124

5.5

ACIS photometry of point sources. . . . . . . . . . . . . . . . . . . . . . . . 126

5.6

ACIS spectroscopy of point-like sources. . . . . . . . . . . . . . . . . . . . . 129

6.1

Plasma parameters of M 17 and western Carina . . . . . . . . . . . . . . . . 141

Chapter 1
Introduction
The outer atmosphere of massive stars obtains momentum from UV photons radiatedin their photosphere, and eventually the velocity exceeds the star’s surface escape velocity.
This outflow is called stellar wind and shows a terminal velocity (v∞ ) of ∼ 103 km s−1 and a
mass-loss rate of 10−7 –10−5 M yr−1 . The integrated mass, momentum, and energy deposited
to the surrounding environment via stellar winds through the lifetime can be comparable tothose by a supernova (SN) at its death (Leitherer et al., 1992). Stellar winds also contaminate
the environment chemically by releasing heavy metals produced via thermonuclear reactions.
Stellar winds make a ring nebula around individual stars by sweeping out their surrounding matter (Chu, 1991). The most dramatic case is seen in starburst galaxies (e.g., M 82), in
which a galaxy scale bubble is produced by superwind, a combination of stellar winds andsupernovae (McCarthy et al., 1987). In our Galaxy, stellar winds is one of the prospective
candidates for large scale structures, such as the bipolar dust lobe in the Galactic center(Bland-Hawthorn & Cohen, 2003), the Galactic ridge X-ray emission (Kaneda et al., 1997;
Ebisawa et al., 2001), and the Galactic center diffuse X-ray emission (Koyama et al., 1989;
Yamauchi et al., 1990). Of course, SN explosions are vital in these phenomena. However, thepre-explosion effects are equally important, which is poorly understood.
The interaction between stellar winds and ambient matter, hydrodynamic instability
within stellar winds, collision of two stellar winds, or magnetically confined winds cause ashock to produce hot (106 − 107 K) plasma that emits X-rays. X-ray photons can penetrate
through the natal molecular cloud from which massive stars are born, and also the interstellar
matter in the Galactic plane where most of massive stars exist (Wood & Churchwell, 1989).
Emission lines from highly ionized atoms in the X-ray band enable a reliable derivation ofelemental abundances and plasma temperature because they are based on simple physics.
Therefore, X-ray is an important yet unexplored tool to investigate the physics and chemistryof stellar winds quantitatively.
15

16

CHAPTER 1. INTRODUCTION

This thesis focuses on the X-ray emission of stellar winds from massive stars and massive
star clusters using the Suzaku satellite. Other X-ray satellites, XMM-Newton and Chandra
with better spatial resolutions, are supplementary used to investigate the point source properties. Infrared data by ground-based and space-based observatories in infrared bands arealso used to constrain the nature of X-ray sources.
We select four fields that represents a wide variety of phenomena related to the X-ray
emission from massive stars. The first field is the central region of our Galaxy. We perform
a very deep survey of the Galactic center region to find new massive stars. We also selectthree Galactic fields, M 17, the Carina Nebula, and the Arches cluster, because these regions
have unambiguous diffuse X-ray emission but with unclear physical origin.
The plan of this thesis is as follows. In chapter 2, we review the theoretical basis ofstellar winds and survey recent observational results to clarify the problems to address in thiswork. The basic properties of the Suzaku observatory are shown in chapter 3. We present
the timing and spectral features of massive stars and star clusters in our Galaxy obtainedwith Suzaku in chapter 4. In chapter 5, we discuss the nature and the emission mechanisms
for both the discrete sources and diffuse emission. Finally in chapter 6, we summarize theresults and conclude the thesis.

Chapter 2
Review
In the first subsection, we review the structure and evolution of massive stars. In
section 2.2, the spatial distribution of WR stars is described. Theoretical studies on the
X-ray emission from massive stars are summarized in the following subsections.

Contents

2.1

2.1

Evolution of Massive Stars . . . . . . . . . . . . . . . . . . . . . .

17

2.2

Wolf-Rayet Stars . . . . . . . . . . . . . . . . . . . . . . . . . . . .

21

2.3

Diffuse X-ray Emission from Massive Stars . . . . . . . . . . . .

24

2.4

Point-like X-ray Emission from Massive Stars . . . . . . . . . .

27

Evolution of Massive Stars

The radiation energy of stars at the main-sequence phase is produced by the nuclear< 1.5 M is the pp
fusion of H. The major energy-generating process of stars with mass of ∼chain,
11

→ 2 H + e + + νe
 H+ H2(2.1)
H + 1H→ 3 He + γ 3 He + 3 He → 4 He + 2 1 H.

Considering that the first and second steps occur twice for the third one, the net
process is 4 1 H → 4 He + 2γ + 2νe + 2e+ . The rate controlling reaction is the firststep, which involves a weak
 The energy-generating rate is expressed as interaction.
−1/3
5
2 −2/3erg s−1 g−1 , where ρ is the density in g cm−3 ,pp = 5.3 × 10 ρ X T7
exp −15.7 T717

18

CHAPTER 2. REVIEW

X is the H mass fraction, and T7 is the temperature in 107 K.

Figure 2.1: Diagram of CNO cycle (Bahcall, 1969). Protons and neutrons are shown in red andblack dots, respectively.

Hydrogen burning can take place in another way, utilizing C, N, and O as catalysts.
As shown in figure 2.1, the net reaction is 4 1 H → 4 He + 3γ + 2νe + 2e+ . Since reaction
involving C, N, and O have higher potential barriers than those involving only H and He,
the CNO cycle requires a higher temperature than ∼ 1.8 × 107 K. The slowest step is the
proton capture of 14 N (14 N + 1 H → 15
 O + γ). The energy-generating rate is expressed as
−2/3
−1/3
26
CNO = 4.9 × 10 ρ X ZCNO T7
exp −70.5 T7
erg s−1 g−1 , where ZCNO is the massfraction of C, N, and O.
Figure 2.2 shows the energy production rate of both processes at the center of a star.
With the central temperature of 1.6 × 107 K for the Sun, the energy generation is dominated(> 90%) by the pp chain. Since the reaction rate of the CNO cycle is much higher thanthat of pp chain (CNO ∝ T 19 and pp ∝ T 4 at T = 1.6 × 107 K) for higher temperatures, it
dominates the hydrogen burning for massive stars.
The difference of the dominant hydrogen burning process causes the difference in stellar
structure. In the inner part of low-mass stars, the energy is transferred by radiation since
the temperature gradient of the pp chain is small, while that in massive stars is transported
by convection. The energy transfer in the outer layer of high mass stars is dominated by

19

10

4

10

3

10 2
1

1
10−1

pp

n
ai
h
c

ycle

10

CN O cEnergy production rate (erg s−1 g −1 )2.1. EVOLUTION OF MASSIVE STARS

10−2
−3

10

10−4
10 6

10 7

10 8

Temperature (K)Figure 2.2: Energy production rate of the pp chain (red) and CNO cycle (blue) as a function oftemperature. We adopted the values of the center of the Sun; ρ = 156, X = 0.33, Z CNO = 0.005(Bahcall et al., 1995).

radiation, whereas that of low mass stars is by convection because of the high opacity due tothe presence of H− . A differential rotation can amplify the magnetic field in the outer layerof low mass stars, but not of high mass stars. Plasma production by magnetic reconnections(Shibata & Yokoyama, 1999) cannot take place in high mass stars.
When hydrogen is exhausted at the central core, the core starts to contract until He
burning is ignited. Since the pressure gradient also rises, the star begins to expand and theeffective temperature drops. Stars in this stage are called red giant.
If the initial mass exceeds ∼ 30M , the outer hydrogen layer is blown away by stellar
winds and the helium layer is exposed without becoming a red giant. This class is called
Wolf-Rayet stars, which is described in the next subsection. The post-main-sequence stellar
evolution differs according to the initial mass ranges.
 O–Of–BSG–LBV–WR–SN

M2 < M
O–BSG–YSG–RSG–WR–SN
M 1 < M < M2
 O–(BSG)–RSG–YSG–(Ceheid)–RSG–SN
M < M1 ,(2.2)

where BSG, YSG, and RSG are blue, yellow, and red supergiants, respectively. LBV repre-

20

CHAPTER 2. REVIEW

sents luminous blue variable. Various values are proposed for the lower- and upper- limits ofinitial mass depending on the calculation method, the metal abundance, and the initial spin
period (Chiosi & Maeder, 1986; Maeder & Meynet, 1987; Meynet & Maeder, 2005). The
typical values of M1 and M2 are ∼ 25–40 and ∼ 40–60, respectively.

2.2. WOLF-RAYET STARS

2.2

21

Wolf-Rayet Stars

Figure 2.3: HST image of M1-67 around WR 124 (Grosdidier et al., 1998).

Wolf & Rayet (1867) discovered three peculiar early-type stars (now designated as
HD 191765 = WR 134, HD 192103 = WR 135, HD 192641 = WR 137) having exceptionallybroad emission lines in the Cygnus region. The widths of emission lines from highly excited He, C, and N are equivalent to Doppler velocities of ∼ 103 km s−1 . This class of stars
are named Wolf-Rayet (WR) stars after their discovery. Figure 2.3 shows the Hα image takenwith the Hubble space telescope (Grosdidier et al., 1998). A typical spectrum of a WR star
(WR 158) is shown in figure 2.4 (Rayner et al., 2003). Although the original definition of
WR stars includes central stars in planetary nebulae (population II), we hereafter discuss
only early-type stars (population I). Conti & Walborn (1976) suggested that WR stars are
evolved massive stars that have lost the H shell through stellar winds with a mass-loss rateof 10−5 –10−4 M yr−1 , and the ”peeled onion skin” (Paczynski, 1973) is seen.
With the initial mass exceeding 20 M (Meynet & Maeder, 2005), a WR star will causea core-collapsed supernova (Type Ib or Ic). Actually, Soderberg et al. (2008) detected themoment that a WR star caused a supernova explosion for the first time with an intense
X-ray outburst. Furthermore, WR stars are one of the strong candidates for the progenitorsof gamma-ray bursts, the most energetic events in the universe. The pre-explosion activity
also plays an important role by fueling mass, momentum, and energy into the circumstellar
environment, which is comparable to that of a supernova (Leitherer et al., 1992).

22

CHAPTER 2. REVIEW

Figure 2.4: WR158 spectrum in the NIR bands. (Rayner et al., 2003)

WR stars are divided into three sub-groups according to line intensities (Beals , 1938;
Barlow & Hummer, 1982).
 WN star: N and He emission lines are intense.

WC star: He, C, and O emission lines are intense.

(2.3)



 WO star: O emission lines are intense.

Some exhibit the property of both the WN and WC types. In the latest catalogue compiling
298 WR stars identified in our Galaxy (van der Hucht, 2006), the spectral sub-type distribution is: 171 WN stars, 10 WN/WC stars, 113 WC stars, and 4 WO stars. In the localgroup besides our Galaxy, there are 132 stars in the Large Magellanic Cloud (Breysacher
et al., 1999; Moffat, 2008), 12 stars in the Small Magellanic Cloud (Massey & Duffy, 2001;
Massey et al., 2003), 29 stars in M 31 (Moffat & Shara, 1987; Armandroff & Massey, 1991),
and 141 stars in M 33 (Massey & Johnson, 1998; Abbott et al., 2004).
Since the mass-loss rate of WR stars is larger than that of main-sequence O stars bymore than ten-fold, it is vital to estimate the total number in the Galaxy. Although thenumber is rapidly increasing, it is considered that much more WR stars exist in the Galaxy(van der Hucht, 2001). Figure 2.5 shows the WR star surface number density as a functionof Galactocentric distance. If we assume that the sample is complete between 7–12 kpc fromthe Galactic center, the difference between the extrapolated model and the observed number

2.2. WOLF-RAYET STARS

23

Figure 2.5: WR star surface number density as a function of Galactocentrinc distance (van der
Hucht, 2001). The best-fit power-law is shown in the dashed line using the five points between 7
and 12 kpc.

implies that there are still more than 6000 WR stars undiscovered in the Galaxy. The innerpart of the Galaxy is a promising region that harbors most of the unseen WR stars.

24

2.3

CHAPTER 2. REVIEW

Diffuse X-ray Emission from Massive Stars

In early-type stars, stellar winds are driven by resonance lines in the UV bands. Whenan ion in the stellar atmosphere receives a photon corresponding to a resonance line, the ion
is accelerated outward and the line is blue shifted. The ion can continue to be accelerated
without shielding by inner ions. Since the opacity is large at resonance level, this mechanism effectively accelerate stellar winds (line-driven wind). The fundamental equation for
spherically-symmetric winds is expressed as,v
2

dv 1 dp GM++ 2 (1 − ΓR ) = 0,
dr ρ dr
r

(2.4)

∞
4πr
in which ΓR (= GMκν Fν dν) represents the radiative pressure effect. The basic theory ofc 0stellar winds was constructed by Lucy & Solomon (1970) and Castor et al. (1975). Figure 2.6
shows the velocity structure of line-driven winds (Abbott, 1980). The velocity is shown in
the unit of Mach number. The critical point is located in the supersonic region.

R

Figure 2.6: Velocity structure of line-driven stellar winds. The thick (dashed) line shows criticalsolution curve where a subsonic flow becomes a supersonic flow. The vertical line at r/r ∗ ∼ 6 shows
the critical point (Abbott, 1980).

2.3. DIFFUSE X-RAY EMISSION FROM MASSIVE STARS

25

Weaver et al. (1977) presented a self-similar solution of stellar winds interacting with
the ISM. Radiative cooling and thermal conduction between the hot interior and ambient
matters were taken into account. They showed a single O7 star (Ṁ = 10−6 M ) with anage of 106 yr forms a hot (∼ 106−7 K) bubble by the shock, which can be observed as soft
extended (∼ 10 pc) X-ray emission (figure 2.7).

Figure 2.7: Temperature (dashed line) and density structure (solid line) of a stellar wind bubble
around a single O7 star (Weaver et al., 1977).

Although Weaver et al. (1977) predicted X-ray emission with LX ∼ 2 × 1033 erg s−1 fromall O stars, there are only two single massive stars that showed X-ray emission from bubbles(S 308; Bochkarev 1988, NGC 6888; Wrigge et al. 1994, Wrigge 1999, Chu et al. 2003). Bothof the central stars are Wolf-Rayet stars, hence no single early-type main-sequence star is
known to show diffuse X-ray emission from the bubble. Figure 2.8 shows the composite
image of NGC 6888.

26

CHAPTER 2. REVIEW

Figure 2.8: Color composite image of a stellar wind bubble around WR 136. The intensities ofHα line, O III line, and X-ray are shown in red, green, and blue respectively. The central star isout of the FOV. (Chandra Photo Album: http://chandra.harvard.edu/photo/2003/ngc6888/)27

2.4. POINT-LIKE X-RAY EMISSION FROM MASSIVE STARS

2.4

Point-like X-ray Emission from Massive Stars

The first detections of X-ray emission from early-type stars were made with the Einsteinobservations of the Cyg OB2 region and the Carina Nebula (Harnden et al., 1979; Seward
et al., 1979). It is hard to explain the emission with coronal activity, since X-ray emission
from corona should suffer significant photoelectric absorption by stellar wind, which is absent
in the observed spectra. X-ray emission should arise farther out in the wind. The X-ray
luminosity in the 0.2–4.0 keV band shows strong linear correlation LX /Lbol ∼ 10−7 (Harnden
et al., 1979; Pallavicini et al., 1981). The correlation is independent of the spectral type
between A5–O3, indicating a universal mechanism for X-ray emission (figure 2.9).

Log L X (erg s−1)

34
32
Algol

30
28
26
34

38
36
Log L bol (erg s−1)

40

Figure 2.9: Bolometric luminosity and X-ray luminosity in the 0.2–4.0 keV band of early-typestars. Crosses, circles, and squares are for O stars, B stars, and A stars, respectively. The dashedstraight line corresponds to the relationship LX = 10−7 Lbol . The only deviated star is the peculiar
binary β Per (Algol). Data are taken from Pallavicini et al. (1981).

Lucy & White (1980) and Lucy (1982) proposed a theoretical model for the structure of
unstable line-driven winds in early-type stars. In their model, supersonic motion relative tothe mean outflow is developed due to hydrodynamic instability. Feldmeier (1995) presented
time-dependent numerical simulations for the O supergiant star ζ Pup (Teff = 4.2 × 104 K,
R∗ = 19 R , M = 42 M , v∞ = 2 × 103 km s−1 , and Ṁ = 3 × 10−6 M yr−1 ). They assumed

28

CHAPTER 2. REVIEW

Figure 2.10: Structure of stellar wind where sound-wave instability produces many small-scale
shock waves (Feldmeier, 1995). Top, middle, and bottom panels show the gas density, wind velocity,and temperature respectively. Dotted lines show the initial conditions.

1-dimensinal symmetry and gave perturbation with 1% amplitude and a period of 5000 s.
Figure 2.10 shows a snapshot of the wind structure 10 days after the perturbation is given.
A sequence of very narrow and dense shells with a temperature of ∼ 106 K are formed byhydrodynamical instability.
The ASCA satellite detected hard X-rays from massive stars for the first time (Yamauchi
et al., 1996; Tsuboi et al., 1997). The temperatures exceed ∼ 3 keV and some sources show
flux-variability, which cannot be explained by the wind-instability model described above.
Head-on collision between two fast winds can yield a postshock temperature simply because
the relative velocity is the sum of two winds’ velocities.

2.4. POINT-LIKE X-RAY EMISSION FROM MASSIVE STARS

29

Figure 2.11: Luminosity history of WR 140 (Pollock et al., 2002). Data from EXOSAT (1–10 keV),ROSAT (0.3–2.0 keV), and ASCA (1-10 keV) are shown with squares, pluses, and triangles, respectively. The sold line shows the best-fit X-ray orbit of e = 0.55. The lightcurves of soft and hardbands were seperately fitted.

One example of the wind-wind collision is a binary composed of two massive stars.
Stevens et al. (1992) presented a 2-dimensional hydrodynamical simulation of colliding windsof binary of two massive stars. If we assume that the stellar winds have reached the terminal
velocity even at the periastron, the electron density (ne ) and the volume (V ) of the plasma
scale as ne ∝ D−2 and V ∝ D3 , where D is the binary separation. Recalling that the twowinds collide at their terminal velocities, the plasma temperature does not depend on D.
We thus find the luminosity scale as LX ∝ ne 2 V ∝ D−1 . From a Kepler’s law of motion,
D ∝ 1/(1 + e cos φ), hence LX ∝ 1 + e cos φ. The prototype is WR 140 comprised of a WC 7
and a O4–5 V star. The X-ray emission is hard (Koyama et al., 1994) and shows a year-scale
flux-variability. Pollock et al. (2002) reported the X-ray light curve spanning ∼ 15 years(figure 2.11). Although the luminosity is well correlated with the orbital phase, the best-fit
eccentricity of 0.55 is significantly lower than that derived by optical observations (e = 0.88;
Marchenko et al. 2003).

30

CHAPTER 2. REVIEW

Figure 2.12: Temperature map of θ 1 Ori C obtained by two-dimensional hydrodynamical simulation (Gagné et al., 2005).

Another mechanism is magnetically channeled wind shock (MCWS). Stellar winds inlarge-scale dipole magnetic field are channeled along the magnetic field line and meet at theequator. This picture was first proposed for magnetic Ap/Bp stars (Shore & Brown, 1990;
Babel & Montmerle, 1997), but can also be applied to O stars. The O5.5 V star θ 1 Orionis C
is the only O star with measured magnetic field strength. This object is the primary ionizing
source of M 42, and shows hard X-ray emission modulated with the spin period of 15.4 days
(Stahl et al., 1993; Gagné et al., 1997). Donati et al. (2002) conducted spectropolarimetricobservations of θ 1 Orionis C. Their model indicates that the dipole magnetic field has anintensity of ∼ 1 kG and a tilt of 42◦ ± 6◦ with respect to the rotation axis. Gagné et al.
(2005) showed analyzed the phase-resolved Chandra/HETG spectra. They also performed
2-dimensional MHD simulation (figure 2.12) and showed that the observed spectra were well
reproduced by MCWS produced at a distance two times the stellar radius.

Chapter 3
Instruments
In this chapter, we review the Suzaku satellite and the onboard instruments. The first
subsection describes the overview of the spacecraft. In the following subsections, we show
the principles of the telescope and the detectors. Then, the actual performance in the orbit
and the calibration accuracy are shown.

Contents
3.1

Overview of Suzaku . . . . . . . . . . . . . . . . . . . . . . . . . .

32

3.2

XRT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

34

3.2.1

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

34

3.2.2

Onborad Performance . . . . . . . . . . . . . . . . . . . . . . . . .

35

XIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

38

3.3.1

Overview of the XIS . . . . . . . . . . . . . . . . . . . . . . . . . .

38

3.3.2

Onboard Performance . . . . . . . . . . . . . . . . . . . . . . . . .

40

HXD . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

44

3.4.1

Overview of the HXD . . . . . . . . . . . . . . . . . . . . . . . . .

44

3.4.2

Onboard Performance . . . . . . . . . . . . . . . . . . . . . . . . .

46

3.3

3.4

31

32

CHAPTER 3. INSTRUMENTS

3.1

Overview of Suzaku

(a)

(b)

Sun shade
XRT
EOB

31
Equator570 kmStar
tracker
Gyroscopes(8 panels)

XRS
Base
panel

Orbital path

Side panel

HXD
XIS

(4 units)

Figure 3.1: (a) Schematic view of the Suzaku satellite in orbit (Mitsuda et al., 2007). (b)Schematic cross-section of Suzaku with the internal structure after the EOB extension (Mitsuda etal., 2007).

Suzaku (Mitsuda et al., 2007) is the fifth Japanese X-ray astronomy satellite developed
under a Japan–US collaboration. It was launched by a Japanese M-V 6 rocket from the
Uchinoura Space Center (USC) on July 10, 2005, and was successfully put into a nearcircular orbit at a ∼ 570 km altitude with an inclination angle of ∼ 31◦ . The orbital periodis ∼ 96 minutes. Figure 3.1(a) shows a schematic view of Suzaku in orbit. The spacecraft
length is 6.5 m along the telescope axis after extending the extensible optical bench (EOB)
and the total weight is ∼ 1700 kg.
The solar panel, providing the total electric power of ∼ 660 W, is fixed on the satellite.
This constrains the pointing direction to be 65◦ − 110◦ from the Sun for the battery charge.
With this constraint, almost all targets are occulted by the earth for a certain fraction of
every 96 minutes. Only objects near the orbital pole can be observed nearly continuously.
The averaged observing efficiency is ∼ 45%
Suzaku carries three co-aligned scientific detector systems: the Hard X-ray Detector(HXD: Takahashi et al. 2007; Kokubun et al. 2007), the X-ray Imaging Spectrometer (XIS:Koyama et al. 2007a), and the X-ray Spectrometer (XRS: Kelley et al. 2007). Figure 3.2
shows the configuration of the detectors on the base panel. The XRS and four XIS cam-

33

3.1. OVERVIEW OF SUZAKUA B C D

C

C

B

D
XIS 0

D

B

XIS 1

ActX

ActY

A

ActX

ActY

ActY

ActX

ActX

HXD

A B C D

A
ActY

XIS 3

XIS 2
Calibration source

XRS

Figure 3.2: Configuration of the detectors in the Suzaku satellite. (the Suzaku technical description). The definition of the XIS coordinates is also shown.

eras are located at the focal plane of the corresponding X-ray Telescope modules (XRT:
Serlemitsos et al. 2007). We will not discuss the XRS further, which became unoperational
due to the evaporation of the liquid helium coolant a few weeks after the launch. The
configuration of the detectors are shown in figure 3.2.

34

CHAPTER 3. INSTRUMENTS

3.2

XRT

3.2.1

Overview of the XRT

Figure 3.3: Picture of the XRT-I1 module (Serlemitsos et al., 2007).

Five XRT modules are installed on the top plate of the EOB. One is designated as
XRT-S, dedicated for the XRS. The other four modules are adapted for each XIS camera
and referred to as XRT-I0, I1, I2, and I3 (figure 3.3).
X-rays are reflected by smooth surfaces of high-electron-density material at very shallow
angles of ≤ 1◦ (grazing incidence). Since the reflection is non-dispersive, circular X-ray mirrors for grazing incidence can focus X-rays over a broad energy band. In this configuration,
only a fraction of the geometric area of the telescope can be used to collect X-ray photons
for imaging, so multiple layers are combined to increase the effective area.
In Wolter type-I reflective optics, X-rays are reflected twice by a parabolic surface and a
hyperbolic surface sequentially (figure 3.4). The precise Wolter design yields a spatial resolution of ∼ 1 arcsec, such as the X-ray telescope adopted for the Chandra X-ray Observatory.

35

3.2. XRT

Mirrors

X-rays
Focal point

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

However, this sacrifices the effective area, as the small surface roughness requires a heavy
substrate and limits the number of nesting layers.
Suzaku takes the opposite approach. It increases the effective area with 175 thin foil
layers in each XRT-I module at the sacrifice of good spatial resolution. It is difficult toconstruct a precise Wolter geometry with thin foils with a thickness of ∼ 180 µm adopted
for XRT, so both the primary and secondary reflectors of XRT are approximated by simple
cones, which limits the spatial resolution to about ∼ 18 arcsec. Since the electron density
strongly depends on the atomic number, the reflector surface is coated with Au with athickness of ≥ 1000 Å.
Both the primary and secondary reflectors are made of four sections (”quadrants”).
Since the spaces between two quadrants are dead areas, point source images appear to be across or a butterfly (figure 3.6).

3.2.2

Onborad Performance

The left panel of figure 3.5 shows the energy-dependent on-axis effective area including
the detector efficiency. At larger off-axis angles, incident X-rays are partially obscured byneighboring reflectors. This vignetting effect is prominent at high energy because highenergy photons have smaller critical angles and are efficiently reflected at innermost mirrors,
where the cone angles is small and the reflectors are tightly packed, thus highly subject to36

CHAPTER 3. INSTRUMENTS

(a)

(b)

Figure 3.5: (a) Total effective area of the XRT combined with the XIS (Serlemitsos et al., 2007).
Those of other X-ray missions are also shown for comparison. (b) Off-axis angle dependence of theeffective normalized at the optical axes in the soft and hard bands shown with black and red points
respectively (Serlemitsos et al., 2007). Data are obtained with the Crab nebula observations. The
solid lines show the simulated model.

shadowing. Figure 3.5 (b) shows the vignetting curve in the soft (3–6 keV) and hard (8–10 keV) bands. Since the effective area is defined by the number of source photons within
the CCD field of view (FOV) of a 17.0 8×17.0 8 square, the model curves show sharp drops at∼ 80 .
The angular resolution was verified with the SS Cyg data in quiescence (figure 3.6).
Figure 3.7 shows the point spread function (PSF) and the encircled energy function, which
are typical indicators of spatial resolution. Since the primary factor of the XRT spatial
resolution is the misalignment of the nested reflectors, the PSF is nearly energy-independent.

37

3.2. XRT

0.1

1

10

ActX (arcmin)

8
4
0
−4
−8
−8 −4 0 4 8
ActY (arcmin)1(a)Brightness

0.1
0.01
−3

10

−4

10

−5

10

0

2
4
6
8
Radius from focus (arcmin)

Encircled photon fraction

Figure 3.6: Image of SS Cyg at the optical axis taken with XRT-I2 (Serlemitsos et al., 2007).
The color scale shows the X-ray counts pixel−1

1

(b)

0.5

0
0

4
2
6
8
Radius from focus (arcmin)

Figure 3.7: Point spread function (PSF, left panel) and the encircle energy fraction (EEF, rightpanel) of the XRT (for XIS2) at the focal plane (Serlemitsos et al., 2007).

38

CHAPTER 3. INSTRUMENTS

3.3

XIS

3.3.1

Overview of the XISElectric
valve

Hood
105 mmBaffleWall calibration
sourceOptical blocking
filter (OBF)
Cover shield
Thermo-electric
CCD
cooler (TEC)
Heat sink

Base
70 mm

Door calibration
sourceBonnet
125 mmDoor

Flexible print

Figure 3.8: Left panel: Picture of one of the XIS sensors. Right panel: Cross section of the XIS
sensor (Koyama et al., 2007a).

The XIS is comprised of four X-ray charge coupled device (CCD) camera systems (figure 3.8). X-ray CCD converts an incident X-ray photon to a cloud of electrons, the amountof which is proportional to the energy of the absorbed X-ray photon. The cloud madein the depletion layer in each pixel is transferred to the gate of the output transistor by
three-phase electric potential clocking. In MOS-type CCDs (figure 3.9), the gate is made ofpoli-Si (Metal), the insulator is made of SiO2 (Oxide), and the depletion layer is made of Si(Semiconductor). In the frame transfer mode, the charges are transferred from the imaging
area (IA) to the frame store (FS) region after an exposure. Then, the charges in the FS
region is transferred vertically along the ActY direction and horizontally along the ActX
direction sequentially.
One side of a CCD is coated with the gate structure mainly made of Si and SiO2 (front

39

3.3. XIS

X-ray

X-ray

Gate
Insulator

Depletion layer
Front illumination

Depletion layer
Insulator
Gate
Back illumination

Figure 3.9: Difference of front illumination and back illumination in MOS-type CCD.

side). Front-illuminated (FI) CCDs detect X-rays that penetrate the gate structure with athickness of ∼ 0.7 µm, whereas the surface dead layer of back-illuminated (BI) CCDs have athickness of ∼ 10 nm (figure 3.9). Due to the absence of the gate structure, BI CCDs achievea high quantum efficiency (QE) even below ∼ 1 keV. On the other hand, FI CCDs employ
thicker depletion layer than BI CCDs (∼ 76 µm for FI CCDs and ∼ 42 µm for BI CCD),
and have higher QE for high energy photons.
In both the IA and FS regions, the CCD chips have a format of 1024 × 1024 pixels andfour read-out nodes. Thus, a CCD chip is composed of four segments of 256 × 1024 pixels(figure 3.10). The pixel size of the IA is 24 µm×24 µm. Combined with XRT, each pixel
corresponds to 1.0400 ×1.0400 of the sky, and the field of view is 180 ×180 . Three of the chips(XIS 0, XIS 2, and XIS 3) are FI CCDs, and the remaining one (XIS 1) is a BI CCD. The QEof the FI CCDs is better than that of BI CCD above ∼ 4 keV (figure 3.11). The detector isoperated at a temperature of ∼ −90◦ C controlled by Pertier coolers.
Each XIS has an optical blocking filter (OBF) to block optical and UV photons. The
OBF is made of ∼ 1400 Å polyimide (C22 H10 N2 O4 ) coated with Al of a ∼ 1200Å thickness.
For on-orbit calibration of the energy scale, each device is equipped with two radioactive sources (55 Fe). 55 Fe decays to leave 55 Mn via electron capture mainly (88.5%) from

40

CHAPTER 3. INSTRUMENTS

D
C
B
A

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

the K-shell. Subsequently, fluorescent X-ray is emitted, accompanied by an L→K electrontransition (Kα: 5.895 keV) or an M→K electron transition (Kβ: 6.513 keV). These sources
illuminate two far corners from the read-out nodes. A small fraction of the X-rays from thecalibration sources is scattered onto the entire CCD.

3.3.2

Onboard Performance

The charge transfer inefficiency (CTI) has been increasing since the launch due to charge
traps caused by radiation damage (Nakajima et al., 2008; Ozawa et al., 2009). Consequently,
the energy resolution has been degrading, which was initially of ∼ 130 eV at 5.9 keV in full
width at half maximum. (figure 3.12). To mitigate the degradation, spaced-row charge
injection (SCI) technique is applied for almost all observations since October 2006. In the
SCI technique, electrons are injected from the furthest side from the read-out nodes to fill
the traps as ”sacrificial charges”. This prevents charges generated by X-rays from being

41

Quantum efficiency

3.3. XIS

1
FI (XIS 0)
BI (XIS 1)

0.5

0

0

5
10
Energy (keV)

180

220 XIS 0
180

Apr.2007

Sep.2005

Typical error

140
220 XIS 1
180
140
0

200

400

Time since 2005-7-10 (day)600

MnI Kα FWHM (eV)MnI Kα FWHM (eV)Figure 3.11: Quntum efficiency of the XIS as a function of incident photon energy. The black andred curves are for XIS 0 and XIS 1 respectively (Koyama et al., 2007a). The effect of contaminanton the OBF is not included.

XIS 1

170
160

Oct.
2006

Mar.
2008

150
XIS 0
140
400

600
800
Time since 2005-7-10 (day)1000

Figure 3.12: Long-term trend of energy resolution (Ozawa et al., 2009; Uchiyama et al., 2009).

captured by the traps and actively reduces the CTI. A new calibration method dedicated for
SCI observations has been established (Uchiyama et al., 2009). With this method, the energy
resolution still retains ∼ 155 eV and ∼ 175 eV at 5.9 keV for FI and BI CCDs respectively
as of March 2008 (figure 3.12).
The low energy QE has degraded by unexpected contamination (due to out-gassing

42

6

systematic error

Mass surface density (µg cm −2 )C column density (10 18 cm−2 )CHAPTER 3. INSTRUMENTS

100
4
50

2

0
Aug Jan
2006

XIS3
XIS2
XIS1
XIS0

0
Jul

Jan
2007

Jul

Jan
2008

Jul

Figure 3.13: Long-term trend of on-axis contaminant represented by carbon column density (The
Suzaku technical description). An empirical model is shown in solid lines.

from the satellite) accumulating on the OBF (Koyama et al., 2007a). The contaminant ismainly composed of carbon and oxygen. The time evolution of the contamination has beenmeasured using regular observations of 1E 0102–072 (figure 3.13). Although the chemicalcomposition of the contaminant is still uncertain, it is assumed to be C/O=6 in number.
The effect of contamination is included in ancillary response files made by xissimarfgen(Ishisaki et al., 2007).
Figure 3.14 shows the non–X-ray background (NXB) accumulated during night-earth
observations. The data obtained during the South Atlantic anomaly passages and in thecalibration source areas are excluded. The flux of NXB depends on the geo-magnetic cut-off
rigidity (COR, figure 3.14) and the position of CCDs. The COR distribution was recently
updated and is called COR2. We hereafter use COR2. Tawa et al. (2008) modeled the NXB
spectra as a function of COR2. The NXB can be constructed for any observations with anaccuracy of ∼ 5%.
On November 9, 2006, XIS 2 suddenly showed an anomaly. About 2/3 of the image wasunusable due to a large amount of charge leakage. This event was possibly caused by a micrometeoroid impact, as seen in XMM-Newton and Swift. Since this anomaly, XIS 2 has been
dysfunctional. Figure 3.15 shows the time of XIS 2 anomaly and start of SCI observations.
The observations presented in this thesis are also shown sequentially.

43

Au Lα
−7

10

Mn Kα
Al Kα
Si Kα
Au Mα

Ni Kα

Au Lβ
−8

10

XIS0
XIS1
XIS2
XIS3

0.2

Counts s 1

Counts s −1 keV −1 pixel −1

3.3. XIS

0.1

0.05
−9

10

0

XIS0 XIS2
XIS1 XIS3

Mn Kβ Ni Kβ5510Cut−off rigidity (GV)

10
Energy (keV)

Figure 3.14: XIS spectra obtained during night earth observations.

Time

Number of
XIS cameras

4

SCI

Observation

Off

Carina Nebula 1stArches cluster 1st
M17
Carina Nebula 2nd
Arches cluster 2nd
CXOGC J1745.6−281546

2006 Sep.
2006 Nov.

3

On

Carina Nebula 3rdArches cluster 3rd
M17 East, M17 BG
Carina Nebula 4th

Figure 3.15: History of the observing condition of the XIS.

44

CHAPTER 3. INSTRUMENTS

3.4

HXD

3.4.1

Overview of the HXD

Figure 3.16: Picture of the HXD (Kokubun et al., 2007).

The HXD (figure 3.16) is a non-imaging collimated detector. It has a compound eye
configuration to increase the effective area. The sensor consists of 16 well-counter units and
20 surrounding bismuth germanate crystal (BGO) anti-coincidence counters (figure 3.16).
Each well-counter unit is comprised of four (2 × 2) Si PIN diodes, which are locatedin front of a GSO scintillator. The PIN is sensitive in 10–60 keV, and gradually becomestransparent to higher energy photons. The GSO detects photons that pass through the PIN
and is sensitive in 40–600 keV (figure 3.17). A coarse collimator and an active shield is madeof BGO scintillators. A BGO scintillator is placed also below the GSO scintillators. The
aspect ratio of the coarse collimators restricts the FOV of the GSO to be 4.◦ 5 in FWHM. The
PIN has extra ”fine collimators” made of bronze and achieves a narrower FOV of 340 × 340 .

45

Effective area (cm 2)

3.4. HXD

PIN

GSO

100
10
10

100
Energy (keV)

Figure 3.17: Effective areas of PIN and GSO (Takahashi et al., 2007)

BGO

GSO

BGO

PMT

PIN diodeassembly

Fine collimator

Figure 3.18: Schematic view of a well-counter unit (Takahashi et al., 2007).

All scitillator signals in one unit is read out from one photomultiplier tube (PMT). Since
the BGO and GSO have different rise/decay times (∼ 700 ns for BGO and ∼ 120 ns for GSO
at a working temperature of −20◦ C), a single PMT can discriminate signals in the two typesof scintillators.

CHAPTER 3. INSTRUMENTS

0.5

6

0.4

4

Y-axis (arcmin)

Counts s−1 (15−40 keV)

46

0.3
0.2
0.1
0

−40 −20 0 20 40
X-axis (arcmin)

2
0
−2
−4
−6
−2

0

2 4 6 8 10
X-axis (arcmin)

Figure 3.19: Angular response of PIN (Kokubun et al., 2007).

3.4.2

Onboard Performance

The angular response of PIN was measured using multiple pointing observations of the
Crab Nebula. The observations were performed with offset angles of 0, ±3.0 5, ±70 , ±100 ,±200 in both the X and Y directions of the satellite. Background-subtracted spectra wereconstructed for all of the 64 PIN diodes individually. The resultant count rates in 15–40 keV
are shown in figure 3.19. The fluxes are plotted against X- and Y- axes and the peak point
was searched by fitting with a triangular function. Figure 3.19 shows the optical axes of 64
PIN diodes. Since the weighted mean is slightly shifted by ∼ 40 from the XIS optical axis,
the effective area of PIN decreases by ∼ 10% for the sources at the XIS-nominal position.
The NXB of the HXD depends on COR like that of the XIS. Figure 3.20 shows the
NXB spectra with three COR regions, together with their ratios. The NXB increases atlower COR, and this trend is more significant at higher (≥ 25 keV) energy band. The
reproducibility of the latest NXB model (bgd d) is 1.3% in 1σ.

47

ratio

Counts s−1 keV−1

3.4. HXD

10−1
10−2
10−3

2.0
1.0

10

20
50
Energy (keV)

Figure 3.20: NXB spectra of the PIN sorted with COR (Kokubun et al., 2007). Events are
extracted from the Earth occultation data.

Chapter 4
Observations & Results
In this chapter, we describe the Suzaku observations on four fields containing massive
stars and massive star clusters. In section 3.1, we identify a new source serendipitously
discovered during the Galactic center survey conducted with Suzaku. In sections 3.2 and
3.3, we present the analysis of two famous star forming regions, the Carina Nebula andM 17 (the Omega Nebula). Both regions contain discrete massive stars and diffuse
emission. Finally in section 3.4, we show the analysis of the Arches cluster. The cluster
also has point sources and diffuse emission, which are not spatially resolved with Suzaku’s
telescope. The goal of this chapter is to present light curves of all discrete sources andspectra of all discrete sources and diffuse sources.

Contents
4.1

4.2

4.3

A New Wolf-Rayet Star in the Galactic Center . . . . . . . . .

51

4.1.1

Objectives . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

51

4.1.2

Observations & Reduction . . . . . . . . . . . . . . . . . . . . . . .

52

4.1.3

Band-limited Images & Source list . . . . . . . . . . . . . . . . . .

54

4.1.4

A Discrete source . . . . . . . . . . . . . . . . . . . . . . . . . . . .

55

Carina Nebula . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

59

4.2.1

Objectives . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

59

4.2.2

Observations & Reduction . . . . . . . . . . . . . . . . . . . . . . .

61

4.2.3

Band-limited Images & Source list . . . . . . . . . . . . . . . . . .

61

4.2.4

Discrete Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . .

62

4.2.5

Diffuse Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

70

M 17 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

73

49

50

CHAPTER 4. OBSERVATIONS & RESULTS

4.4

4.3.1

Objectives . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

73

4.3.2

Observations & Reduction . . . . . . . . . . . . . . . . . . . . . . .

76

4.3.3

Band-limited Images & Source list . . . . . . . . . . . . . . . . . .

76

4.3.4

Discrete sources. . . . . . . . . . . . . . . . . . . . . . . . . . . .

80

4.3.5

Diffuse sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

87

4.3.6

HXD-PIN Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . .

90

Arches Cluster . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

92

4.4.1

Objectives . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

92

4.4.2

Observations & Reduction . . . . . . . . . . . . . . . . . . . . . . .

94

4.4.3

Band-limited Images & Source list . . . . . . . . . . . . . . . . . .

94

4.4.4

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

95

4.1. A NEW WOLF-RAYET STAR IN THE GALACTIC CENTER

4.1

A New Wolf-Rayet Star in the Galactic Center

4.1.1

Objectives

51

In the Galactic center region, WR stars as well as main-sequence massive stars areconcentrated in three clusters near the Galactic center; the Arches (Nagata et al., 1995;
Cotera et al., 1996; Serabyn et al., 1998; Figer et al., 1999b), the Quintuplet (Kobayashi
et al., 1983; Okuda et al., 1990; Nagata et al., 1990; Figer et al., 1999a), and the Centralcluster (Becklin & Neugebauer, 1968; Krabbe et al., 1995; Ghez et al., 2005). However, it isquite natural to expect that a much larger number of early-type stars remain unidentified.
Portegies Zwart et al. (2001) claimed that the number of young massive star clusters may
exceed 50 in the Galactic center region. The inconsistency with the observed value indicatesthat these clusters are too obscured to be visible in the optical and infrared bands, or thatthe cluster members are dissipated before the earliest member reaches its end, or that thepredicted number is a gross overestimate. Kim et al. (1999) showed that the strong tidal
disruption in the Galactic center region makes massive star clusters become unbound on avery short time scale comparable to the lifetime of an O star. Isolated early-type stars may
be distributed throughout the region, unlike the other parts of the Galaxy where they arefound in clusters.
Some early attempts have been made to discover unidentified early-type stars in the
Galactic center region with the combination of X-ray and infrared (IR) observations. Muno
et al. (2006a) made trailblazing observations of radio, IR, and X-ray to reveal the populationof young massive stars in the Galactic center region. They detected strong Brγ and He I lines
from two sources, and classified them as either Of or candidate LBV stars. Mauerhan et al.
(2007) also conducted K -band spectroscopy of six bright hard X-ray sources with very red
IR colors. Two of them, with IR colors of J –K = 4–5 mag, show broad hydrogen and helium
emission lines characteristic of evolved O-type stars. They are classified as a Wolf-Rayet starof a spectral type of WN6 and an O-type Ia supergiant.
Likewise, Mikles et al. (2006) performed infrared spectroscopy of CXO J174536.1–
285638, which is one of new Chandra sources discovered by Muno et al. (2003). They
detected P Cygni profiles of He II lines of a 170 km s−1 wind. The spectral features in the Xray and IR bands are most consistent with a colliding wind binary system such as η Carinae.
This source is of particular interest in terms of a strong Fe XXV emission line, which is also
prominent in our source, presented below. We see a high prospect of these methods to findsimilar sources en masse.
Near-IR spectroscopy may confront a challenge for even redder sources. In addition toa ∼ 30 mag visual extinction ubiquitous toward the Galactic center region (Catchpole et

52

CHAPTER 4. OBSERVATIONS & RESULTS

al., 1990; Schultheis et al., 1999), evolved early-type stars are behind an additional local
extinction by their own mass loss. Photospheric emission is reprocessed into longer wavelength radiation in the optically thick circumstellar matter. The IR spectra turn out to befeatureless by strong veiling (Figer et al., 1999a; Crowther et al., 2006). For such sources,
the spectroscopy of the hard X-ray emission directly from the vicinity of the star is the only
practical tool. These stars contribute to the 6.7 keV line emission, and the diagnosis of theline emission helps to identify their nature.
Suzaku has made a total of ∼ 2 Ms mapping observations toward the central ∼ 1◦ × 3◦
of our Galaxy. With the large effective area, this deep survey provides the best opportunity
to perform a systematic search for WR stars. Figure 4.1 shows the XIS mosaic image ofthe Galactic center in the 6.55-6.8 keV. In addition to the diffuse emission prevailing thewhole region, we see some local excess, which we identify as the SNR G0.61+0.01 (Koyama
et al., 2007b), the pulsar wind nebula G0.9 + 0.1 (Mereghetti et al., 1998; Sidoli et al., 2000;
Porquet et al., 2003), the Great Annihilator 1E 1740.7–2942 (Bouchet et al., 1991; Sakano
et al., 1999; Gallo & Fender, 2002), the low-mass X-ray binary 1A 1742–294 (Lewin et al.,
1976; Pavlinsky et al., 1994; Sidoli et al., 1999), and the Arches cluster. We found a bright,
hard, and point-like source at (l, b) ∼ (0.7, 0.1), the nature of which is not addressed before.
To identify the nature of this source, we perform temporal and spectral analysis.

4.1.2

Observations & Reduction

Galactic latitude (degree)

CXOGC J1745.6-281546
G0.61+0.01

Arches

0.4
0.2
0.0
− 0.2
− 0.4
1A 1742−2982

1.5

1.0

0.5

0.0
359.5
Galactic longitude (degree)

1E 1740.7−2942

359.0

358.5

Figure 4.1: XIS mosaic image of the Galactic center region in the 6.55–6.8 keV band. The exposure correction is applied after the
NXB is subtracted.

4.1. A NEW WOLF-RAYET STAR IN THE GALACTIC CENTER

G0.9+0.1

53

54

CHAPTER 4. OBSERVATIONS & RESULTS

Table 4.1: Observation log of the Sgr B North region.

∗

Start Date

ObsID

2006-09-21
2006-09-24

501040010
501040020

R. A.
Decl.
(J2000.0)hm
17 46 46s –28◦ 220 5100
17h 46m 46s –28◦ 220 5100

texp ∗(ks)
70
50

Exposure time after good time interval filtering described in the text.

The source was detected in the vicinity of the Sgr B molecular complex. The observations were conducted twice in 2006 September 21–23 and 24–25 (table 5.1). We removed
the events taken during the South Atlantic Anomaly passages and at earth elevation angles
below 3 degrees. After the filtering, the combined net integration time is ∼ 120 ks.

4.1.3

Band-limited Images & Source list
Table 4.2: Point source list in the Sgr B North region

Source NameCXOGC J174645.3–281546

Position (J2000.)
RA
Dec
h
m
s
17 46 45 –28◦ 150 5600

Figure 4.2 shows the XIS images in the (a) 0.5–2.0 keV and (b) 2.0–7.0 keV bands. The
two extended sources seen in the hard band image are Sgr B2 molecular cloud (Koyama etal., 1996; Murakami et al., 2000, 2001) and G0.61–0.01 (Koyama et al., 2007b). The brightest
source near the northern edge of the FOV is the main topic of this section.
We extracted the source and background photons from the polygonal and dashed elliptical regions, respectively (figure 4.2b). We simulated a point source image at the source
position using xissim (Ishisaki et al., 2007), and extracted source photons at various enclosed
photon fractions. Consequently, we found that a 60% enclosed photon polygon maximizes
the signal-to-noise ratio. The background region was selected from an ellipse adjacent toand at a similar off-axis angle with the source region.

55

4.1. A NEW WOLF-RAYET STAR IN THE GALACTIC CENTER−28:10:00

Background(b) Sgr B2(a)Declination

15:00

Source

20:00
25:00
30:00

G0.61−0.01

35:00
30

17:47:00

30

Right ascension

46:00

30

17:47:00

30

46:00

Right ascension

Figure 4.2: XIS images of the study field in the (a) 0.7–2.0 and (b) 2.0–7.0 keV bands. All four
XIS images in the two observations are combined. The images were adaptively binned so thatevery bin has a signal-to-noise ratio larger than 8 using the weighted Voronoi tessellation algorithm(Cappellari et al., 2003; Diehl et al., 2006). Then, the NXB was subtracted and the vignetting wascorrected. The FOV is shown with solid squares.

4.1.4

A Discrete source

Timing Analysis
Figure 4.3 shows the XIS light curve in the 2.0–7.0 keV band. The light curve with abin width of 4000 s shows no variation. The constant model is not rejected with χ 2 /degreeof freedom (dof) = 67.6 / 57.
Spectral Analysis
Figure 4.4 shows the background-subtracted spectrum. Since the response function ofeach FI camera is essentially same and the energy gain is well calibrated, events of XIS 0,
XIS 2, and XIS 3 are summed throughout this thesis. The spectrum is characterized byseveral emission lines as well as a hard continuum extending above 10 keV with a strong
cut-off at ∼ 3 keV. The lines are identified as Kα emission from highly ionized (He-likeand H-like) atoms of Ar, Ca, and Fe. The spectrum was first fitted by a bremsstrahlung

56

CHAPTER 4. OBSERVATIONS & RESULTSCounts s−1 (2.0−7.0 keV)

0.1

0.05

0

0

105

2×105
Time (s)3×105

Figure 4.3: XIS light curve of CXOGC J174645.3–281546in the 2.0–7.0 keV band. The light curveis binned for each 4000 s.

model with Gaussian lines convolved by interstellar extinction (Morrison & McCammon,
1983) of solar abundances (Anders & Grevesse, 1989). We generated the ancillary response
files (ARFs) of the XRT using xissimarfgen (Ishisaki et al., 2007) and response matrix
files (RMFs) of the XIS using xisrmfgen. The strongest line is attributed to Fe XXIV Kα at
6.66 ± 0.01 keV with an energy flux of (2.5 ± 0.3) × 10−13 erg s−1 cm−2 , which corresponds
to an equivalent width (EW) of 950 ± 100 eV.
The line feature at 6.4 keV, which stems from quasi-neutral Fe and is a characteristicof the diffuse emission in this region, is not detected. By adding a narrow Gaussian line at
6.4 keV, we found a 90% upper limit of 1.5×10−14 erg s−1 cm−2 , which corresponds to an
EW of 50 eV. Since the temperature derived from the continuum shape and those from theintensity ratios of He-like and H-like Kα lines are consistent with ∼ 4 keV, the emission isinferred to originate from thermal plasma in a collisional equilibrium.
We then fitted the spectrum with an optically-thin thermal plasma model (APEC:
Smith et al. 2001) with interstellar extinction. The abundances of Ar, Ca, and Fe relative
to solar were left to vary, while other elements were fixed to 1 solar abundance. A singletemperature model yielded an acceptable fit with a plasma temperature (kB T ) of ∼ 3.8 keV,
1.0×10−12 erg s−1 cm−2 in X-ray flux (FX ) in the 2.0–8.0 keV band and 2.4×1023 cm−2 in hydrogen column extinction (NH ). No systematic deviation was found in the single temperature

57

χCounts s−1 keV−1

4.1. A NEW WOLF-RAYET STAR IN THE GALACTIC CENTER

Ni Fe

0.01

10−3
Ar

Ca

Fe

10−4
2
0
−2
2

5
Energy (keV)

10

Figure 4.4: XIS spectra of CXOGC J174645.3–281546. The BI and summed FI spectra arerespectively shown in black and red. The upper panel shows the data with crosses and the bestfit model with solid lines, while the lower panel shows the residual to the best-fit. Conspicuous
emission lines are shown with gray dotted (He-like ion) and solid (H-like ion) lines.

fit (figure 4.4), indicating that no extra component is required.
To estimate the systematic error caused by the choice of background region, we repeated< 10%
the same procedure with several different regions. Although the X-ray flux suffers ∼systematic uncertainty, other best-fit parameters are consistent with each other.

58

CHAPTER 4. OBSERVATIONS & RESULTS

Table 4.3: Best-fit spectral parameters.∗

Parameter
NH
kB T
ZAr
ZCa
ZFe
FX †
LX ‡
χ2 /d.o.f.
∗

Unit(10 cm−2 )(keV)(solar)(solar)(solar)−12(10
erg s−1 cm−2 )(1034 erg s−1 )
23

Value
+0.3
2.4−0.2+0.5
3.8−0.6+6.3
8.7−4.9+2.3
3.2−2.0+0.1
0.8−0.1+0.03
1.00−0.03
2.8
89.4/106

The uncertainty indicates the 90% confidence ranges of the fit.
Energy flux in the 2.0–8.0 keV band.
‡
Absorption-corrected luminosity in the 2.0–8.0 keV band. A
distance of 8 kpc is assumed.
†

59

4.2. CARINA NEBULA

4.2

Carina Nebula

4.2.1

Objectives

In colliding wind binary systems with high eccentricity, the X-ray luminosity is expected
to vary as a function of orbital phase due to the variation of volume emission measure of
the plasma (Stevens et al., 1992). The spectral shape may also be modified by the variable
absorption column of the cold wind material. The thermal X-ray emission from Cygnus
OB2#8A (De Becker et al., 2006), η Carinae (Hamaguchi et al., 2007b), and WR 140 (Pollock
et al., 2005) clearly depends on the orbital phase, implying that a major part of the plasma
is produced by wind collision. For such study, repeated X-ray observations with a typicalorbital period of massive stars binary of 100–1000 days are necessary.

Declination (J2000.0)−59:35:00
40:00
45:00
50:00

46:00

45:00
10:44:00
43:00
Right ascension (J2000.0)

Figure 4.5: 2MASS image of the Carina Nebula. J, H, and K band intensities are shown in blue,
green, and red, respectively.

60

CHAPTER 4. OBSERVATIONS & RESULTS

The Carina Nebula is one of the brightest H II regions in the sky, which is composed ofseveral star clusters (Trumpler 14–16, Collinder 228 and 232, and Bochum 10 and 11). Thisregion harbours more than 64 O stars (Smith, 2006), including the luminous blue variable η
Carinae, the WR star HD 93162 (=WR 25), the main-sequence O3 star (HD 93205), and themain-sequence O3+O8 binary (HD 93250). With the proximity of ∼ 2.3 kpc (Davidson &Humphreys, 1997), the Carina region is one of the best regions in the Galaxy for systematicstudies of massive stars. Actually, this region is extensively studied by Chandra (Evans etal., 2003, 2004; Sanchawala et al., 2007; Albacete-Colombo et al., 2008) and XMM-Newton(Albacete Colombo et al., 2003; Antokhin et al., 2008). Fortunately, the region is observedby Suzaku and other satellites repeatedly, as η Car is often used for a routine calibration
source. Many early-type stars are within the fields of view. This rich data set is ideal to
study long-term behaviors of these massive stars.

Figure 4.6: Einstein/IPC contour of the Carina Nebula overlaid on the B band image.

This region is also distinctive for the diffuse X-ray emission discovered by the Einsteinobservatory (Seward et al., 1979; Seward & Chlebowski, 1982). The emission extends over
∼ 1 deg2 with a total X-ray luminosity of ∼ 1035 erg s−1 . Hamaguchi et al. (2007) showed that
the X-ray spectra are different between the north and south sides of η Carinae, especiallybelow 1 keV. They argued that Fe and Si abundances in the south region are higher than
< 0.4). The
those in the north region with a factor of 2–4 and N to O abundance is low (∼authors concluded that the diffuse emission originates not from wind-blown bubbles, butfrom one or multiple supernova remnant(s).

61

4.2. CARINA NEBULA

4.2.2

Observations & Reduction
Table 4.4: Observation log of the Carina Nebula.

Start Date

ObsID

2005-08-29
2006-02-03
2007-06-23
2008-06-10

100012010
100045010
402039010
403035010

∗

R. A.
Decl.
(J2000.0)hm
10 45 07s –59◦ 410 0400
10h 44m 59s –59◦ 410 0900
10h 44m 52s –59◦ 430 4100
10h 44m 55s –59◦ 430 5300

texp ∗
(ks)
68
22
63
35

Exposure time after good time interval filtering described in the text.

Suzaku observed η Carinae four times (table 4.4). The former two observations wereconducted at the XIS nominal position and without the SCI, while the latter two observations
were conducted at the HXD nominal position and with the SCI. XIS 2 was not used in thelatter two observations. Data taken at elevation angles below 4◦ from the Earth rim and
10◦ from the sunlit Earth rim were discarded. All data were processes using ver. 2.3 of the
Suzaku pipeline software.

4.2.3

Band-limited Images & Source list
Table 4.5: Point source list in the Carina Nebula region

Source Name
WR 25
HD 93205
HD 93250

Position (J2000.)
RA
Dec
h
m
s
10 44 10 –59◦ 430 1100
10h 44m 34s –59◦ 440 1600
10h 44m 45s –59◦ 330 5400

Spectral Type

Reference

WN6+O4O3 V+O8 V
O3 V

1
2, 3
2

References. — (1) van der Hucht (2001) (2) Levato et al. (1991) (3) Morrell et al. (2001)

Figure 4.7 shows the XIS images of the Carina region in the (a) soft (0.5–2.0 keV) and(b) hard (2.0–5.0 keV) bands. We see extended emission in the northeastern and south partof the field in the soft band, whereas the emission from point sources dominate the hardband image. Since Hamaguchi et al. (2007) argued that the diffuse emission shows spatial
variation particularly in the soft band, we further divided the 0.5–2.0 keV events into threebands (0.5–0.7 keV, 0.7–1.0 keV, and 1.0–2.0 keV) to make a three color image (figure 4.8).
The uncertainty of the XIS sky coordinates is ∼ 2000 (Uchiyama et al., 2008). To achieve

62

CHAPTER 4. OBSERVATIONS & RESULTS

a higher astrometric accuracy, we registered the XIS frame using the three bright point
sources (WR 25, HD 93205, and HD 93250). We do not use η Car, which has an extended
structure in X-rays. We referred the sky position determined with Chandra (Evans et al.,2003), and shifted the XIS frame in each observation by 8–1500 . The accuracy of the Chandra
frame is ∼ 0.00 5.
We analyze the three point sources, WR 25, HD 93205, and HD 93250, which have notbeen discussed in detail in previous work (Hamaguchi et al., 2007). HD 93250 was out of
the FOV in the latter two observations, while WR 25 and HD 93205 were detected in allof the four observations. We also analyze the diffuse emission according to the emission
morphology (figure 4.8).

−59:30:00

HD 93250(a)contaminationfrom Tr 14(b)

Declination

35:00
η Car

40:00

WR 25

45:00
50:00
5’
HD 93205

55:00
46:00

10:45:00

44:00

Right ascension

3.3 pc @ 2.3 kpc

46:00

10:45:00

44:00

Right ascension

Figure 4.7: XIS images of Carina Nebula in the (a) 0.5–2.0 keV and (b) 2.0–5.0 keV bands. The
FOVs of each observation are shown with solid squares.

4.2.4

Discrete Sources

Timing Analysis
We present the temporal and spectral behaviors of the three massive stars (table 4.5).
Source photons are accumulated from elliptical regions of 4000 –9000 in the axis lengths, while
the background photons are from adjacent regions free of sources. The source light curves are63

4.2. CARINA NEBULA

Diffuse 3
BG

Diffuse 2

Diff

use

1

Figure 4.8: Three color image of the Carina Nebula in the soft band. The events in 0.5–0.7 keV,
0.7–1.0 keV, and 1.0–2.0 keV are shown in red, green, and blue, respectively. The FOVs of the four
observations and the photon extraction regions are shown in are shown in green and white solidlines.

shown in figures 4.9–4.11. None of the three sources shows flux variation in each observation.
Spectral Analysis
Figures 4.12–4.14 show the spectra of the three sources. We see emission lines fromhighly ionized O, Ne, Mg, Si, S, and Fe, indicating an optically-thin thermal plasma origin.
We then fitted the spectra with the APEC model attenuated by interstellar absorption.
The summed FI and the BI spectra were fitted simultaneously. We first fitted the spectrawith a single temperature model. If this model is rejected, an additional component with

64

CHAPTER 4. OBSERVATIONS & RESULTSCounts s−1 (2.0−7.0 keV)Counts s−1 (0.5−2.0 keV)1(a) 2005−08−29(c) 2007−06−23(b) 2006−02−03(d) 2008−06−10

0.5

0

0

2×104

4×104 6×104
Time (s)8×104 0(a) 2005−08−29

104

2×104 3×104 4×104
Time (s)

0

5×104
105
Time (s)0(c) 2007−06−23(b) 2006−02−03

104

2×104 3×104 4×104
Time (s)(d) 2008−06−10

0.4
0.3
0.2
0.1
0

0

2×104

4×104 6×104
Time (s)8×104 0104

2×104 3×104 4×104
Time (s)

0

5×104
105
Time (s)0104

2×104 3×104 4×104
Time (s)Figure 4.9: XIS light curves of WR 25. Upper and lower panels show the count rate in the soft
(0.5–2.0 keV) and hard (2.0–7.0 keV) bands respectively.

Counts s−1 (0.5−2.0 keV)

0.2(a) 2005−08−29(b) 2006−02−03(c) 2007−06−23(d) 2008−06−10

0.15
0.1
0.05
0

0

2×104

4×104 6×104
Time (s)8×104 0104

2×104 3×104 4×104
Time (s)

0

5×104
105
Time (s)0104

2×104 3×104 4×104
Time (s)Counts s−1 (2.0−7.0 keV)

0.05(a) 2005−08−29(b) 2006−02−03(c) 2007−06−23(d) 2008−06−10

0.04
0.03
0.02
0.01
0

0

2×104

4×104 6×104
Time (s)8×104 0104

2×104 3×104 4×104
Time (s)

0

5×104
105
Time (s)0104

2×104 3×104 4×104
Time (s)Figure 4.10: XIS light curves of HD 93205. Symbols follow figure 4.9.

a different temperature was added. The elemental abundances other than He was assumed
to be common. The best-fit parameters are shown in tables 4.6–4.8. While the spectra of
WR 25 and HD 93250 required two components of different temperatures, that of HD 93205
was well reproduced with a single-temperature component.

65Counts s−1 (0.5−2.0 keV)

4.2. CARINA NEBULA(b) 2006−02−03(a) 2005−08−29
0.4
0.3
0.2
0.1
0

0

2×104

4×104 6×104
Time (s)8×104 0104

2×104 3×104 4×104
Time (s)Counts s−1 (2.0−7.0 keV)

0.2(a) 2005−08−29(b) 2006−02−03

0.15
0.1
0.05
0

0

2×104

4×104 6×104
Time (s)104

8×104 02×104 3×104 4×104
Time (s)Figure 4.11: Same as figure 4.9, but for HD 93250.

(a) 2005−08−29(c) 2007−06−23(b) 2006−02−03(d) 2008−06−10

0.1
0.01

χ

Counts s −1 keV −1

1

5
0
−5
1

2

5

Energy (keV)

1

2

5

Energy (keV)

1

2

5

Energy (keV)

1

2

5

Energy (keV)

Figure 4.12: XIS spectra of WR25. Only FI spectra are shown for simplicity. The lower andhigher temperature components are shown in red and blue, respectively.

Since spectral shapes of the all sources did not change significantly, we summed the photons from the four observations of WR 25 and HD 93205 and, the two former observations of
HD 93250 in order to increase the photon statistics. The spectra were not well reproduced byone- or two- temperature APEC models. Allowing the elemental abundance of conspicuous
emission lines to vary, we obtained acceptable fitting results. The best-fit parameters andthe spectra are presented in table 4.9 and figures 4.15–4.17.

66

CHAPTER 4. OBSERVATIONS & RESULTSCounts s−1 keV−1

0.1(a) 2005−08−29(b) 2006−02−03(c) 2007−06−23(d) 2008−06−10

0.05
0.02
0.01

ǿ

5×10−3
4
2
0
−2
−4

1
Energy (keV)

2

1
Energy (keV)

2

1
Energy (keV)

2

1
Energy (keV)

2

Figure 4.13: XIS spectra of HD 93205. The summed FI and BI spectra are shown in black andred, respectively.

Counts s−1 keV−1(a)2005−08−29(b)

2006−02−03

0.1

0.01

χ

5
0
−5
0.5

1

2
Energy (keV)

5

0.5

1

2
Energy (keV)

5

Figure 4.14: XIS spectra of HD 93250. Only FI spectra are shown for simplicity. The lower andhigher temperature components are shown in red and blue, respectively.

Table 4.6: Best-fit parameters of WR 25.

Obs. DateNH
21(yyyy-mm-dd) (10 cm−2 )2005-08-29
8.1±0.3
2006-02-03
8.2±0.5
2007-06-23
7.3±0.2
2008-06-10
7.2±0.5

kB T 1(keV)
0.63±0.01
0.61±0.03
0.63±0.01
0.64±0.02

kB T 2(keV)
2.3±0.1
2.3±0.1
2.1±0.1
1.8±0.1

Z(solar)
0.7
0.7
0.6
0.4

FX
−1

(erg s cm−2 )4.8×10−12
4.1×10−12
6.8×10−12
3.8×10−12

67

4.2. CARINA NEBULA

Table 4.7: Best-fit parameters of HD 93205.

Obs. DateNH
21(yyyy-mm-dd) (10 cm−2 )2005-08-29<2.7
2006-02-03<4.7
2007-06-23<2.0
2008-06-10
2.7±1.1

kB T(keV)
0.58±0.05
0.62±0.11
0.58±0.05
0.54±0.06

Z(solar)
0.06
0.08
0.04
1.0

FX
−1

(erg s cm−2 )4.3×10−13
4.5×10−13
4.4×10−13
4.0×10−13

Table 4.8: Best-fit parameters of HD 93250.

Obs. DateNH
kB T 121−2(yyyy-mm-dd) (10 cm )(keV)
2005-08-29
1.8±0.5
0.73±0.03
2006-02-03
4.8±1.4
0.60±0.03

kB T 2(keV)
3.0±0.2+4.5
4.8−1.6

Z(solar)
0.7
0.7

FX
−1

(erg s cm−2 )1.3×10−12
1.3×10−12Counts s−1 keV−1

0.01

χ

0.1

10−3
3
0
−3
0.5

1
Energy (keV)

2

Figure 4.15: XIS spectrum of HD 93205. The summed FI and BI spectra are shown in black and
red.

68

CHAPTER 4. OBSERVATIONS & RESULTSCounts s−1 keV−1

0.01

χ

1

10−4
3
0
−3
0.5

0.1

10−3

1

2
Energy (keV)

5

10Counts s−1 keV−1

0.01

χ

Figure 4.16: XIS spectrum of HD 93250. The summed FI spectrum shown with crosses. The
lower and higher temperature components are shown in red and blue, respectively.

10−4
3
0
−3
0.5

0.1

10−3

1

2
Energy (keV)

5

10

Figure 4.17: XIS spectrum of WR 25. Only FI spectra are shown for simplicity. The lower andhigher temperature components are shown in red and blue, respectively.

69

4.2. CARINA NEBULA

Table 4.9: Best-fit APEC parameters for HD 93205, HD 93250, WR 25

Pars.
NH ∗
kB Thigh ∗†kB Tlow ∗†
ZNe ∗
ZMg ∗
ZSi ∗
ZS ∗
ZAr ∗
ZCa ∗
ZFe ∗
FX ∗‡
LX §
χ2 /d.o.f.
∗

Units(1021 cm−2 ) . . . . . . . .
(keV) . . . . . . . . . . . . . .
(keV) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(10−13 erg s−1 cm−2 )(1032 erg s−1 ) . . . . . . .
.....................

HD 93205+0.3
5.2−0.9...
+0.04
0.25−0.02
0.3
0.49 (0.24–0.87)
1.07 (0.58–1.72)
0.3
0.3
0.3
0.33 (0.19–0.48)
3.9 (3.8–4.0)
3.2
100.2/134

HD 93250
WR 25
2.2±0.4
7.1±0.4
3.1±0.2
2.33±0.06
0.62±0.03
0.62±0.01
0.3
0.51 (0.34–0.71)
0.3
0.51 (0.38–0.65)
0.3
0.53 (0.39–0.71)
0.3
1.17 (0.98–1.38)
1.67 (0.72–3.00) 1.44 (1.00–1.90)
0.3
0.46 (0.01–0.90)
0.36 (0.22–0.50) 0.28 (0.25–0.31)
7.4 (7.1–7.7)
58.3 (57.8–58.8)
4.9
30.0
66.4/76473.3/405

The uncertainties in the parentheses are the 90% confidence range. Fixed values are shown
without ranges.
†
The plasma temperatures for the higher and lower temperature components. Only the higher
temperature value is given for HD 93205, which is fitted by a single temperature model.
‡
The X-ray flux in the 1.0–8.0 keV band.
§
The absorption-corrected X-ray luminosity in the 1.0–8.0 keV band. A distance of 2.3 kpc is
assumed.

70

4.2.5

CHAPTER 4. OBSERVATIONS & RESULTS

Diffuse SourcesχCounts s−1 keV−1

Based on the morphology, we extracted the source photons of diffuse emission from three
regions (”Diffuse 1”, ”Diffuse 2”, and ”Diffuse 3” in figure 4.8). The background photons wereaccumulated from a region devoid of bright point sources and intense diffuse emission (”BG”
in figure 4.8).

0.01

10−3
5
0
−5
0.5

1

2
Energy (keV)

5

Figure 4.18: ACIS spectrum of integrated point sources in the Diffuse 3 region.

Many point sources contaminate the source spectra, especially in the ”Diffuse 3” region.
We estimate the point source emission observed with the XIS based on the position and flux
derived by a Chandra observation. We used the Trumpler 14 observation (ObsID = 4495)
covering the whole ”Diffuse 3” region and detected 30 point sources using the wavdetectalgorithm. Source and background photons were extracted using the ACIS Extract package(Broos et al., 2002). The spectrum was well described with a single temperature (k B T ∼4.3 keV) APEC model (figure 4.18). Using the position and flux of each source, we simulated
the XIS events of point sources with xissim. We extracted the integrated point source
spectrum and subtracted it from the diffuse spectrum.
Out-of-time events from η Carinae are not negligible in the SCI-on observations. It takes∼ 127 ms to inject the SCI charges and ∼ 25 ms to read out the events. The ”Diffuse 3”

71

4.2. CARINA NEBULA

region lies at almost the same ActY coordinate with ∼ 60 length in ActY direction. Therefore,∼ (0.127+0.025)/8×(6/18) ∼ 0.6% of the η Carinae events contaminates the Diffuse 3 spectra
obtained by XIS 0 and XIS 3. This effect was canceled by subtracting a normalized η Carinae
spectrum by XIS.
The resultant spectra were well reproduced by an APEC model modified by an interstellar photo-electric absorption. The best-fit model and parameters are shown in figure 4.19 andtable 4.10. Since the temperatures and the relative metal abundances are consistent between
diffuse 1 and 2, we simultaneously fitted the spectra of two regions to obtain more constrained
abundances. We tied the temperature and abundances of O, Ne, Mg, and Fe. The resultant
+0.04
+0.04
abundances are ZO = 0.05−0.03, ZNe = 0.09−0.03, ZMg = 0.15 ± 0.04, and ZFe = 0.27 ± 0.03.
Table 4.10: Best-fit parameters of the APEC model for the diffuse emission.∗

Par.
NH
kB T
ZO
ZNe
ZMg
ZFe
FX †
SX ‡
LX §
∗

Unit
Diffuse 1−2(10 cm )
< 0.1
+0.01
(keV)
0.55−0.03+0.14(solar)0.25−0.08(solar)0.25 ± 0.10
(solar)0.22 ± 0.10
+0.14
(solar)0.51−0.08(10−13 erg s−1 cm−2 )17 ± 3
−13
−1
−2
−2
(10
erg s cm arcmin ) 0.9 ± 0.1(1033 erg s−1 )1.1
21

Diffuse 2
Diffuse 3
0.6 ± 0.2
3.6 ± 0.3
0.56 ± 0.01 0.18 ± 0.01
< 0.060.09 ± 0.01
0.04 ± 0.04 0.16 ± 0.03
0.17 ± 0.04 0.32 ± 0.11
0.21 ± 0.02
0.19+0.09−0.05
14 ± 232 ± 20.6 ± 0.1
2.1 ± 0.1
1.1
16

The uncertainties indicate the 90% confidence range.
X-ray flux in the 0.5–2.0 keV band.
‡
X-ray surface brightness in the 0.5–2.0 keV band.
§
Absorption-corrected X-ray luminosity in the 0.5–2.0 keV band. A distance of
2.3 kpc is assumed.
†

72Counts s−1 keV−1

CHAPTER 4. OBSERVATIONS & RESULTS

1

Diffuse 1(a)Diffuse 2

(b)

0.1

χ

0.01
5
0
−5
1
Energy (keV)Counts s−1 keV−1

0.5

2

0.5

1
Energy (keV)

2

1
(c)

Diffuse 3

0.1

χ

0.01
5
0
−5
0.5

1
Energy (keV)

2

Figure 4.19: Background-subtracted spectra of diffuse emission in the Carina Nebula region. The
summed FI and BI spectra are shown in black and red.

73

4.3. M 17

4.3

M 17

4.3.1

Objectives−16:02:00

}

r

he
t
r
o

04:00

N

06:00

}

10:00
12:00
14:00
16:00
18:00
OI 352

20:00
22:00

22:00

21:30

18:21:00

ar

08:00

Weste
rn b

Declination (J2000.0)

ar
b
n

20:30

Right ascension (J2000.0)
Figure 4.20: 2MASS J (blue), H (green), and K (red) mosaic image of the M 17 region.

M 17 (Omega Nebula) is a nearby Galactic H II region associated with giant molecular
clouds (figure 4.20). Hoffmeister et al. (2008) identified 53 OB stars in the central OB
association with near infrared spectroscopy and derived the distance to be 2.1 ± 0.2 kpc. A
much larger number of young OB stars are suggested by near infrared photometry (Ladaet al., 1991; Jiang et al., 2002). The earliest systems are two O4 binaries (CEN 1a and< 1 Myr based on the H-R diagram
CEN 1b). The age of the cluster is estimated to be ∼(Hanson et al., 1997; Hoffmeister et al., 2008). Strong winds and radiation from the central
OB association sculpted the ambient matter to form a V-shaped cloud, which was traced bymolecular and atomic hydrogen lines (Chrysostomou et al., 1992; Felli et al., 1984; Brogan& Troland, 2001). Each side of the V is called Northern and Western bars (figure 4.20).

74

CHAPTER 4. OBSERVATIONS & RESULTS

The diffuse X-ray emission was found by Chandra and ROSAT (Townsley et al., 2003;
Dunne et al., 2003). It has asymmetric morphology with respect to the OB association andfills the cavity of the molecular cloud toward the negative Galactic latitude (Povich et al.,2007; Pellegrini et al., 2007).

χCounts s −1 keV −1

(a)
0.6

(b)

0.4
0.2
0
2
−2
0.5

1
Energy (keV)

2

Figure 4.21: (a) Chandra image of M 17 with adaptive smoothing and exposure corrction. Low(0.5–2.0 keV) and high- (2.0–8.0 keV) energy events are shown in red and blue. (b) Chandra/ACIS
spectrum of the diffuse emission in M 17. The top panel shows the data in crosses and the best-fit
model in the solid line, while the bottom panel shows the residual (Townsley et al., 2003).

Townsley et al. (2003) claimed the first unambiguous detections of such emission in the
H II regions M 17 and the Rosette Nebula using the Advanced CCD Imaging Spectrometer(ACIS; Garmire et al. 2003) onboard the Chandra X-ray Observatory (Weisskopf et al.,2002). With a ∼ 40 ks integration time of M 17 (ObsID = 972), diffuse soft X-ray emission was
detected (figure 4.21) apart from 886 point sources above ∼1029.3 erg s−1 (Broos et al., 2007).
Dunne et al. (2003) showed the entire structure of the soft X-ray diffuse emission (figure 4.22)
using the Position-Sensitive Proportional Counter (PSPC; Pfeffermann et al. 1987) onboard
ROSAT (Truemper, 1982). They measured the total X-ray luminosity and compared to
the wind-blown bubble models with and without heat conduction. They concluded that
only the bubble without heat conduction can account for the observed X-ray luminosityof ∼ 2.5 × 1033 erg s−1 . The magnetic field may be responsible for suppressing the heat
conduction and mass evaporation between the hot gas and cold ISM (Dunne et al., 2003).
The results obtained by these high-resolution imaging studies are generally consistent
with the wind-blown bubble models. However, the observational results are not still accurate

75

4.3. M 17(a)(b)

Region A

Counts

100
50
0
0.1

1
Energy (keV)

10

Figure 4.22: (a) ROSAT/PSPC image of M 17 smoothed with a Gaussian kernel of σ = 20 00 . (b)ROSAT/PSPC spectrum of a part of the diffuse emission in M 17. The data and best-fit model are
shown in squares and dashed line (Dunne et al., 2003).

enough to compare to the theoretical works. It is often ambiguous whether the observed
diffuse emission is from wind-blown bubbles or from supernovae, which give rise to diffuse
emission with a similar X-ray spectral hardness and luminosity in a similar spatial scale.
The largest uncertainty stems from the lack of spectral analysis based on line diagnostics ina spatially-resolved manner.
Resolving emission lines is crucial to examine whether the spectrum is thermal and to
determine the temperature and the chemical composition of the plasma. For example, theintensity ratio of Kα lines between O VII and O VIII and that between Ne IX and Ne X aresteep functions of the plasma temperature at 1–10 MK (Tucker & Gould, 1966). In between
the O and Ne Kα complex, Fe L series lines dominate the spectrum. The metallicity ofthese elements is one of the factors to determine the X-ray luminosity expected from a windblown bubble (Chu et al., 1995). The anomaly in the O and Fe abundance ratio can beused to discriminate different types of supernovae (Tsujimoto et al., 1995; Thielemann etal., 1996; Nomoto et al., 1997) if the emission is of a supernova origin. The previous studies
using ROSAT and Chandra were incapable of resolving these lines, limiting their ability todiagnose the plasma emission.
The XIS onboard Suzaku has a good spectral resolution, a low background, and a large
effective area, which are particularly suited for spectroscopy of extended X-ray emission.
The capability to resolve key elements with sufficient statistical significance has been il-

76

CHAPTER 4. OBSERVATIONS & RESULTS

lustrated by several initial studies on extended emission in H II regions (Hamaguchi et al.,2007; Tsujimoto et al., 2007). Hamaguchi et al. (2007) resolved various emission lines fromthe diffuse emission in the Carina Nebula. Based on the low nitrogen-to-oxygen ratio andthe spatial variation of the Fe and Si abundances, they suggested that the diffuse emission
originates not from wind-blown bubbles but from one or multiple old supernova remnant(s).

4.3.2

Observations & Reduction
Table 4.11: Observation log around the M17 region.

Obs. Name

Start Date

ObsID

M 17 WestM 17 East
M 17 BG

2006-03-11
2007-10-04
2007-10-07

501003010
502052010
502053010

Aim point (J2000.0)
RA
Dec
h
m
s
18 20 49 –16◦ 100 5800
18h 21m 33s –16◦ 120 2200
18h 22m 34s –15◦ 340 4900∗

Exposure time after good time interval filtering described in the text.

†

Before the dead time correction.

texp ∗ (ks)
XIS HXD†
110
86
121
94
75
63

Suzaku made two on-source observations of the M 17 region. Since the diffuse emission
is prevailing over the entire FOV, an off-source region ∼ 350 apart from the OB association
was also observed to construct the background spectra. The layout of the three observed
fields are shown in figure 4.23. We discarded data taken during the elevation angles below
4◦ from the Earth rim and 10◦ from the day Earth rim. We used the data processed withthe pipe-line process version 2.3. The observation log is shown in table 4.11.

4.3.3

Band-limited Images & Source list

Figure 4.24 shows the mosaic image of M 17 East and West region in the (a) 0.5–2.0 keV
and (b) 2.0–5.0 keV bands. The two images are strikingly different. In the hard-band, theimage is dominated by the OB association consisting of hundreds of point sources (Broos etal., 2007). We also see excess emission from a group of protostars in M 17 North (Wilson
et al., 1979; Henning et al., 1998; Broos et al., 2007), OI 352 (Ogura & Ishida, 1976; Whiteet al., 1994), and three point sources in the M 17 East field. We designate the three point
sources Suzaku J182132–1610.4, Suzaku J182132–1613.4, and Suzaku J182140–1604.1. In thesoft band, The diffuse emission is clearly seen as was claimed by Dunne et al. (2003) andTownsley et al. (2003).
We tuned the sky coordinates using Suzaku J182132–1610.4 Suzaku J182132–1613.4 and77

4.3. M 17-15:10:00

M17 BGlaGa

3EG J1824−1514

ctic

20:00

n
pla

30:00

e

Declination

40:00

AX J1819.2−1548

50:00
AX J1819.2−1601-16:00:00
GSO
PIN

10:00
20:00

M17

30:00

t
Wes

M17 East

24:00

22:00
Right ascension

18:20:00

Figure 4.23:
Mosaic XIS image of the three Suzaku observations in the M 17 region.
The FOVs of PIN are displayed with solid lines, while the soft-band XIS images are shown in gray scale. Point sources in the INTEGRAL general reference catalog(http://isdc.unige.ch/Data/cat/29/catalog.html) are shown with crosses.

OI 352. Suzaku J182140–1604.1 was not used because the image is smeared near the FOV
edge. The three sources are also detected in the Chandra/ACIS observation (ObsID = 7391).
We shifted the Suzaku coordinates so that the positions match with those of the Chandra
observations. Since no bright point source are detected in the background observation, we
did not perform the astrometry tuning for M 17 BG.

78

CHAPTER 4. OBSERVATIONS & RESULTS−16:00:00

(a)

5
4
3

05:00
Declination

2
1
10:00
15:00

5’

20:00

3.1 pc @ 2.1 kpc−16:00:00

M17 North(b) Suzaku J182140−1604.1

OB assoc.

Declination

05:00
Suzaku J182132−1610.4

10:00
15:00
Suzaku J182132−1613.4

20:00
OI 352

22:00

21:30
18:21:00
Right ascension

20:30

Figure 4.24: XIS images of the on-source observations in the (a) 0.5–2.0 keV and (b) 2.0–5.0 keVbands. Both images are smoothed with a Gaussian kernel of 1200 after the NXB is subtracted. The
effects of exposure, vignetting, and contamination on the OBF are corrected. The coordinates are
reregistered using the point sources referring the Chandra frame. The FOVs and the source regionsof diffuse emission are shown in white and green solid lines.

79

4.3. M 17

Table 4.12: Point source list of the M 17 region

Source ID
OB association
OI 352
M 17 North
Suzaku J182132–1610.4
Suzaku J182132–1613.4
Suzaku J182140–1604.1

Position (J2000.)
RA
Dec
h
m
s
18 20 30
–16◦ 110
18h 20m 36s –16◦ 150 4200
18h 20m 30s–16◦ 030
18h 21m 32s –16◦ 100 2800
18h 21m 32s –16◦ 130 2400
18h 21m 40s –16◦ 040 1000

NatureReference

O8 V
protostars

1, 2
3, 4, 5

References. — (1) Ogura & Ishida (1976) (2) White et al. (1994) (3) Wilson et al. (1979)(4) Henning et al. (1998) (5) Broos et al. (2007)

80

4.3.4

CHAPTER 4. OBSERVATIONS & RESULTS

Discrete sources

Timing Analysis
0.6Counts s−1(a) 0.5−2.0 keV(b) 2.0−7.0 keV

0.4

0.2

0

0

5×104

105 1.5×105 2×105
Time (s)0

5×104

105 1.5×105 2×105
Time (s)Figure 4.25: XIS light curves of the OB association in the (a) 0.5–2.0 keV and (b) 2.0–7.0 keVbands. The events of four chips are summed. The background is not subtracted.
0.1

Counts s−1(a) 0.5−2.0 keV(b) 2.0−7.0 keV

0.05

0

0

5×104

105 1.5×105 2×105
Time (s)0

5×104

105 1.5×105 2×105
Time (s)Figure 4.26: XIS light curves of M 17 North. The XIS 1 events are not included.

We extracted source photons of point sources (Suzaku J182132–1610.4, Suzaku J182132–
1613.4, Suzaku J182140–1604.1, and OI 352) from polygonal regions along the PSF as is
described in §3.1.1. The OB association photons were extracted from a circular region witha radius of 10 around the cluster center at (RA, Dec) = (18h 20m 29.9s , −16◦ 10 0 4500 ) in theequinox J2000.0 (Chini et al., 1980). We collected M 17 North photons from an ellipticalregion with semi-minor and semi-major axes of 1.0 5 and 2.0 5, respectively. The corresponding
background photons were accumulated from nearby regions individually.
The XIS light curves of the discrete sources are shown in figures 4.25–4.30. Among thesix sources, only Suzaku J182132–1610.4 showed significant flux variation. Soon after the81

4.3. M 17
0.04(a) 0.5−2.0 keV(b) 2.0−7.0 keVCounts s−1

0.03
0.02
0.01
0

0

5×104

105 1.5×105 2×105
Time (s)0

5×104

105 1.5×105 2×105
Time (s)Figure 4.27: XIS light curves of OI 352.
0.2Counts s−1(a) 0.5−2.0 keV(b) 2.0−7.0 keV

0.15
0.1
0.05
0

0

5×104

105 1.5×105 2×105
Time (s)0

5×104

105 1.5×105 2×105
Time (s)Figure 4.28: XIS light curves of Suzaku J182132–1610.4 The events of three chips are summed.
0.06

Counts s−1(a) 0.5−2.0 keV(b) 2.0−7.0 keV

0.04

0.02

0

0

5×104

105 1.5×105 2×105
Time (s)0

5×104

105 1.5×105 2×105
Time (s)Figure 4.29: XIS light curves of Suzaku J182132–1613.4. Symbols follow figure 4.28.

observation start, a flare was detected in both the soft and hard bands. In the decay phase,a weaker flare was also detected. The source was still in the decay phase at the end of the82

CHAPTER 4. OBSERVATIONS & RESULTS
0.02

Counts s−1(a) 0.5−2.0 keV(b) 2.0−7.0 keV

0.01

0

0

5×104

105 1.5×105 2×105
Time (s)0

5×104

105 1.5×105 2×105
Time (s)Figure 4.30: XIS light curves of Suzaku J182140–1604.1. Symbols follow figure 4.28.

observation.
Spectral Analysis
The spectra of the discrete sources are shown in figures 4.31 and 4.32. The summed
FI and BI spectra were fitted simultaneously, but only the summed FI data and the best-fit
model are shown for simplicity. The spectra of the three sources in the M 17 West observation
are characterized by a hard continuum and Kα lines from He- and H-like ions of Mg, Si, S, Ar,
Ca, and Fe. On the other hand, the three point sources in the M 17 East observation show
only Fe emission line. The spectra were first fitted with a single thin-thermal plasma model(APEC; Smith et al. 2001) attenuated by interstellar extinction (Morrison & McCammon,
1983). If the model is rejected, we added an additional plasma component of a different
temperature. The elemental abundances with distinctive emission lines were left free. The
best-fit parameters are summarized in table 4.13.
Since the three point sources in M 17 East region show only Fe emission lines, we assumed
that all elemental abundances relative to the solar values are the same. All spectra were well
reproduced with a single temperature plasma model modified by interstellar photo-electric
absorption. The spectra and the best-fit parameters are shown in table 4.14 and figure 4.32.

83

4.3. M 17

Table 4.13: Best-fit APEC parameters for the discrete spectra.

Pars.
NH ∗
kB Thigh ∗†kB Tlow ∗†
ZNe ∗
ZMg ∗
ZSi ∗
ZS ∗
ZAr ∗
ZCa ∗
ZFe ∗
FX ∗‡
LX §
χ2 /d.o.f.
∗

Units(1022 cm−2 ) . . . . . . . .
(keV) . . . . . . . . . . . . . .
(keV) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(solar) . . . . . . . . . . . . . .
(10−13 erg s−1 cm−2 )(1032 erg s−1 ) . . . . . . .
.....................

OI 352
1.7 (1.4–2.0)
3.8 (3.1–4.7)
0.56 (0.47–0.65)
0.3
0.49 (0.24–0.87)
0.30 (0.16–0.58)
1.07 (0.58–1.72)
0.3
0.3
0.33 (0.19–0.48)
3.9 (3.8–4.0)
5.5
100.2/134

M 17 North
OB association
2.4 (2.1–2.7)
1.3 (1.2–1.4)
2.7 (2.4–3.2)
4.0 (3.9–4.1)...
0.59 (0.56–0.62)
0.3
0.51 (0.34–0.71)
0.3
0.51 (0.38–0.65)
0.3
0.53 (0.39–0.71)
0.3
1.17 (0.98–1.38)
1.67 (0.72–3.00) 1.44 (1.00–1.90)
0.3
0.46 (0.01–0.90)
0.36 (0.22–0.50) 0.28 (0.25–0.31)
7.4 (7.1–7.7)
58.3 (57.8–58.8)
8.4
52
66.4/76473.3/405

The uncertainties in the parentheses are the 90% confidence range. Fixed values are shown
without ranges.
†
The plasma temperatures for the higher and lower temperature components. Only the higher
temperature value is given for M 17 North, which is fitted by a single temperature model.
‡
The X-ray flux in the 1.0–8.0 keV band.
§
The absorption-corrected X-ray luminosity in the 1.0–8.0 keV band. A distance of 2.1 kpc is
assumed.

84

χCounts s−1 keV−1χCounts s−1 keV−1χCounts s−1 keV−1

CHAPTER 4. OBSERVATIONS & RESULTS(a)

OI 352(b)

M17 North(c)

OB association

0.01

10−3
5
0
−5
0.01

10−3
5
0
−5
0.1
0.01
10−3
5
0
−5
1

2
5
Energy (keV)

10

Figure 4.31: Background-subtracted spectra of (a) OI 352, (b) M 17 North, and (c) the OB
association. Red and blue lines show two components separately. We show the summed FI datain crosses and the best-fit models in solid lines. Lower panels show the residual between the data
and model.

85

4.3. M 17

Suzaku J182132−1610.4(b)

Suzaku J182132−1613.4(c)Suzaku J182140−1604.1

10−3
5
0
−5

0.01

10−3

10−4
5
0
−5

10−3

10−4
5
0
−5

χ

Counts s−1 keV−1χCounts s−1 keV−1

(a)
0.01

χ

Counts s−1 keV−1

0.1

1

2
5
Energy (keV)

10

Figure 4.32: Background-subtracted spectra of (a) Suzaku J182132–1610.4, (b) Suzaku J182132–
1613.4, and (c) Suzaku J182140–1604.1.

86

Table 4.14: Best-fit parameters of three point sources in the M 17 EAST region

Pars.
NH
kB T
Z
FX
LX ∗
χ2 /d.o.f.

Unit
Suzaku J182132–1610.4
21−2(10 cm )3.3 ± 0.6(keV)
6.1 ± 0.7
(solar)0.3 ± 0.1
−13
−2
(10
erg s cm )9.5 ± 0.3
32
−1
(10 erg s )55
181.6/209

Suzaku J182132–1613.4
4.6+2.0−1.5
9.5+3.5−2.7
0.9+0.5−0.4
2.3 ± 0.1
1.3
72.4/81

Suzaku J182140–1604.1
91+29
−24
12.7+6.7−3.3
2.9+4.1−1.5
3.4 ± 0.2
3.8
45.8/25∗

The absorption-corrected X-ray luminosity in the 1.0–8.0 keV band. A distance of
2.1 kpc is assumed.

CHAPTER 4. OBSERVATIONS & RESULTS

4.3. M 17

4.3.5

87

Diffuse sources

As was claimed by Dunne et al. (2003), the surface brightness is the highest near the cluster core and spreads out westward in a fan-like form, indicating that the emission originatesfrom the cluster core. We therefore extracted source photons from five concentric pie-shaped
regions with radii of 30 –60 , 60 –90 , 90 –120 , 120 –150 , and 150 –180 , centered at the cluster center.
We excluded the regions around the three point sources in M 17 East field. The background
events were accumulated from the M 17 BG field in the same detector coordinate with thesource regions. M 17 West observation was performed ∼ 1.5 yr earlier than the M 17BG
observation. Since the amount of the OBF contaminant changed significantly during thisinterval (figure 3.13), we processed raw spectra in the following way to compensate for thedifference in the quantum efficiency: (1) We subtracted the NXB provided by the XIS team.
Since the NXB count rates anti-correlate with COR2, we compiled night Earth observations
so that the COR2 distribution becomes the same with that of the source observation. (2)The vignetting including the spatial non-uniformity of the OBF contaminant was correctedby multiplying the effective area ratio between the source and the background regions foreach energy bin of the background spectra. On the other hand, for M 17 East field observed
just before the M 17 BG observation, we simply subtracted the background photons fromsource photons.
Although we collected photons avoiding the point source regions, spill-over photons may
contaminate the diffuse emission spectra. Using xissim, we estimate the spill-over fraction
from the point sources to the diffuse source regions and subtracted them from the diffuse
emission spectra. For regions 2–4, we summed the events from M 17 West and East fields.
The resultant spectra were fitted with an APEC model attenuated by interstellar extinction. The best-fit model parameters and spectra are shown in figure 4.33 and table 4.15.

88Counts s−1 keV−1

CHAPTER 4. OBSERVATIONS & RESULTS(a) Region1(b) Region 2(c) Region 3(d) Region 4

0.1

0.01

χ

5
0

Counts s−1 keV−1

−5

0.1

0.01

χ

5
0
−5
0.5

2

1
Energy (keV)

0.5

1Energy (keV)

2Counts s−1 keV−1(e) Region 5
0.1

0.01

χ

5
0
−5
0.5

1
Energy (keV)

2

Figure 4.33: Background-subtracted spectra of diffuse emission in the M 17 East region in the
0.4–2.0 keV band. Lower panels show the residual between the data and the best-fit model. The
summed FI and BI spectra are shown in black and red.

Par.
NH
kB T
ZO
ZNe
ZMg
ZFe
SX †
∗
†

Unit
(10 cm−2 )(keV)(solar)(solar)(solar)(solar)−14−1(10
erg s cm−2 arcmin−2 )21

Region 1
5.3 ± 0.3
0.23 ± 0.01
0.04 ± 0.01
0.10 ± 0.01
0.09 ± 0.01
0.05 ± 0.01
2.34 ± 0.05

Region 2
4.4 ± 0.3
0.23 ± 0.01
0.03 ± 0.01
0.09 ± 0.01
0.09 ± 0.01
0.05 ± 0.01
1.77 ± 0.03

Region 3
4.1 ± 0.2
0.23 ± 0.01
0.03 ± 0.01
0.08 ± 0.01
0.07 ± 0.01
0.05 ± 0.01
1.79 ± 0.03

Region 4
4.0 ± 0.2
0.24 ± 0.01
0.06 ± 0.01
0.14 ± 0.01
0.11 ± 0.02
0.07 ± 0.01
2.00 ± 0.04

Region 5
4.0 ± 0.2
0.25 ± 0.01
0.10 ± 0.01
0.20 ± 0.02
0.17 ± 0.02
0.12 ± 0.01
1.60 ± 0.05

4.3. M 17

Table 4.15: Best-fit parameters of the APEC model for the diffuse emission in the M 17 West field. ∗The uncertainties indicate the 90% confidence range.
X-ray surface brightness in the 0.5–2.0 keV band.

89

90

CHAPTER 4. OBSERVATIONS & RESULTS

4.3.6

HXD-PIN Analysis

Counts s−1 keV−1

0.1

(a)

M17 West
NXB(b)

M17 East(c)

M17 BG

RAW

0.01
CXB

10−3
10

20
Energy (keV)

50 10

20
Energy (keV)

50 10

20
Energy (keV)

50

χCounts s−1 keV−1

Figure 4.34: PIN spectra of the M 17 region observations. The raw spectra, reconstructed NXB
spectra, and estimated CXB spectra are shown in black, red, green respectively.

0.01

(a)

M17 West(b)

M17 Eest(c)

M17 BG

10−3
10−4
2
0
−2
10

1.3% NXB

20Energy (keV)

50 10

20
Energy (keV)

50 10

20
Energy (keV)

50

Figure 4.35: Background-subtracted PIN spectra of the M 17 regions. Both the NXB and CXB
are subtracted in data shown with crosses. Solid black lines show the best-fit power-law model.
The NXB reproducibility is shown in red lines. Lower panels show the residual between the data
and the best-fit model.

We used the PIN data processed with the pipeline processing ver 2.0. In this process,data are discarded during the elevation angles below 5◦ and COR less than 6 GV. Slightly
tighter screening criteria yield shorter net exposures for PIN than XIS (table 4.11). We used
the time-dependent NXB model constructed as a function of PINUD provided by the HXDteam. We also corrected the source exposure for the instrumental dead time, which reduced
exposure time by ∼ 7%.
The raw spectra of each observation are shown in figure 4.34 overlaid with simulated
cosmic X-ray background (CXB) and NXB spectra. CXB spectra were made using the model

91

4.3. M 17

by Boldt (1987) and Gruber et al. (1999) in the same manner as Reeves et al. (2007). After
subtracting the CXB and NXB, we see excess counts in the ∼ 13–30 keV band above thecurrent background uncertainty (figure 4.35). We fitted each spectrum with simple powerlaw, which yielded acceptable results (table 4.16).
Table 4.16: Best-fit parameters of PIN spectra of the M 17 regions

Obs. Name
M 17 WestM 17 East
M 17 BG*Γ+0.6
3.3−0.5+0.7
3.0−0.6+0.7
1.9−0.6FX *5.1 ± 0.5
4.2 ± 0.5
5.5 ± 0.9

LX †
2.7
2.2
2.9

χ2 /d.o.f.
5.7/6
4.4/6
4.4/6

The X-ray flux in the 15–30 keV band in the unit of
10−12 erg s−1 cm−2 . The source is assumed to be at the HXD
nominal position.
†
The X-ray luminosity in the 15–30 keV band in the unit of1033 erg s−1 . A distance of 2.1 kpc is assumed.

92

CHAPTER 4. OBSERVATIONS & RESULTS

4.4

Arches Cluster

4.4.1

Objectives

Figure 4.36: HST image of the Arches cluster (Figer et al., 1999b). Red, green, and blue scales
show the flux in F205W, F160W, and F110W, respectively.

The central ∼ 200 pc of our Galaxy is particularly rich in curious spatial and spectralfeatures. One of the features distinctively seen in the Galactic center region is the diffuse
iron Kα line emission at ∼ 6.4 keV (Koyama et al., 1996). The line originates from neutral or
low-ionized iron (hereafter quoted as Fe I), but its ionizing mechanism is not yet understood.
Two major ideas have been developed. One is fluorescence from photo-ionized iron (Sunyaev
et al., 1993; Koyama et al., 1996; Revnivtsev et al., 2004). The X-ray spectrum from the
Sagittarius B2 cloud, which is characterized by the strong ∼ 6.4 keV emission line and the∼ 7.1 keV iron K edge, is best accountable by this model (Koyama et al., 1996; Murakami
et al., 2000). Another idea is that iron is ionized by accelerated electrons (Yusef-Zadeh etal., 2002b; Predehl et al., 2003; Wang et al., 2006).
The distribution of the diffuse 6.4 keV emission is highly asymmetric around the Galactic
center with a bias in the north-east direction (Koyama et al., 1996, 2007b) and appears

4.4. ARCHES CLUSTER

93

filamentary (Park et al., 2004). Of particular interest is the 6.4 keV emission associated with
the Arches cluster, which is one of the richest and the most densely packed massive starclusters in our Galaxy.
The Arches cluster has a total mass of ∼ 104 M , a compact size with a diameter of ∼ 1500(∼ 0.6 pc at an assumed distance of 8.5 kpc), a stellar mass density of ∼ 3 × 105 M pc−3 ,and an estimated age of 2–4.5 Myr (Nagata et al., 1995; Cotera et al., 1996; Serabyn et al.,
1998; Figer et al., 1999b; Blum et al., 2001; Yang et al., 2002; Stolte et al., 2002; Figer etal., 2002; Figer, 2005). It contains ∼ 5% of all the known Wolf-Rayet stars in our Galaxy(Figer, 2005), which are characterized by a large mass loss rate reaching ∼ 10−4 M yr−1
(Lang et al., 2001, 2005). The level of high mass star-forming activity of the Arches cluster,which can be measured by the number of O stars, is comparable only to NGC 3603 (Moffat,1983), W49A (Conti & Blum, 2002; Alves & Homeier, 2003), and Westerlund 1 (Clark etal., 2005) in our Galaxy and R136 at the center of 30 Doradus (Campbell et al., 1992) inthe Large Magellanic Cloud.
Yusef-Zadeh et al. (2002) first reported the detection of two point sources in addition
to diffuse emission from this region using the Chandra X-ray Observatory (Weisskopf et al.,2002). The extended X-ray emission is elongated beyond the cluster boundary with a size of∼ 6000 ×9000 and its spectrum can be fit with a thermal model with a temperature of ∼ 5.7 keVin addition to a ∼ 6.4 keV line (Yusef-Zadeh et al., 2002). Based on these claims, the hard
diffuse emission is considered an example of the cluster wind plasma, which is a collection ofinteracting stellar winds from massive stars in a cluster (Cantó et al., 2000; Raga et al., 2001;
Silich et al., 2004; Rockefeller et al., 2005). However, the poor spectrum of the extended
emission does not allow a conclusion to be drawn about a thermal origin. Law & Yusef-Zadeh(2004) cautioned that the continuum emission can also be fit by a power-law, and can thus
be explained by non-thermal emission. Recent results based on a longer exposure Chandra
observation by Wang et al. (2006) indicated that the cluster wind is only seen at the cluster
center, and the majority of the diffuse emission is composed of a strong 6.4 keV line and apower-law continuum emission of non-thermal origin.
In order to understand the nature of the X-ray emission in the Arches cluster and the
diffuse 6.4 keV emission ubiquitous in, but exclusive to, the Galactic center region, hard-band
X-ray spectroscopy with high signal-to-noise ratio (S/N) is obviously a key. The bandpass
includes both thermal (Fe XXV Kα and Fe XXVI Kα at ∼ 6.7 and ∼ 7.0 keV) and fluorescent(Fe I Kα and Kβ at ∼ 6.40 and ∼ 7.05 keV) features as well as the iron K absorption edge
feature at ∼ 7.11 keV. The underlying power-law continuum, if it ever exists, is visible at> 8 keV, where the thermal contribution plays a minor role. Moreover, the hardenergies ∼X-rays penetrate through the large attenuation of AV ∼ 30 mag common toward the Galactic
center region (Morris & Serabyn, 1996).

94

CHAPTER 4. OBSERVATIONS & RESULTS

Here, we present the results of a spectroscopic study of the hard emission in the Arches
cluster using the Suzaku Observatory. The XIS onboard Suzaku provides high S/N spectrain the hard-band aided by the excellent spectroscopic performance of the CCDs and thelarge effective area of the X-ray optics. To compliment the moderate spatial resolution of
Suzaku, we refer to the recent Chandra results by Wang et al. (2006) and results obtainedby our own reduction of the same dataset.

4.4.2

Observations & Reduction

Four observations of the Galactic center region were conducted using Suzaku (table 4.17). The aim points and roll angles among four observations are almost identical.
The events were removed when they were taken during the elevation angles below 5◦ , theday earth elevation angles below 10◦ , and the SAA passages. The net exposure after thisfiltering is ∼ 220 ks in total. For the first and second observations, we corrected for the CTI
and fine-tuned the energy gain across the chips assuming that the center energy of Fe andS Kα lines in the diffuse spectrum is spatially uniform. These calibration processes were
performed using the same dataset with the presented analysis. Resultantly, the systematicenergy uncertainty of XIS is as low as ∼ +3/−6 eV at 6 keV. The details of the calibrationof the Galactic center datasets are described in Koyama et al. (2007b). The third and fourth
observation data (ObsID=100048010 and 102013010) were processed with ver. 2.3 pipeline
software.
Table 4.17: Observation log of the Arches cluster.

∗

4.4.3

Start Date

ObsID

2005-09-23
2005-09-30
2006-09-08
2007-09-03

100027010
100037040
100048010
102013010

R. A.
Decl.
(J2000.0)hm
17 46 02s –28◦ 550 3200
17h 46m 02s –28◦ 550 3200
17h 46m 03s –28◦ 550 3600
17h 46m 03s –28◦ 550 3900

texp ∗
(ks)
48
46
69
57

Exposure time after good time interval filtering described in the text.

Band-limited Images & Source list

Figure 4.38 shows the band-limited XIS images in the (a) 6.3–6.5, (b) 6.6–6.8, and(c) 7.5–10.0 keV bands. The former two bands contain Kα emission lines from iron at low
and high ionization stages, respectively. The 7.5–10.0 keV image complements the one below

95

4.4. ARCHES CLUSTER

Arches cluster-28:45:00 HD 316314

Declination

50:00

55:00

-29:00:00

Sgr A East
05:00

AX J1745.6-2901
47:00

30

17:46:00

Right ascension

45:30

47:00

30

17:46:00

45:30

Right ascension

Figure 4.37: XIS images in the (a) 0.7–2.0 keV and (b) 2.0–5.0 keV bands. The making of the
figures and symbols follow that of figure 4.24.
Table 4.18: Source list in the Sgr A East region

Source Name

Position (J2000.)
RA
Dec
h
m
s
Arches cluster 17 45 50 –28◦ 490 2200

8.0 keV by Chandra; the effective area at 8 keV of Suzaku XIS is larger than that of Chandra
ACIS by more than ten-fold. A significant excess is seen at the position of the Arches cluster
in all bands, which is elevated from the diffuse emission at the north east of the Galactic
center. Due to the limited spatial capability of Suzaku XIS, it is difficult to discriminate
point sources from the underlying diffuse emission in the Arches cluster.
The 7.5–10.0 keV image (figure 4.38c) morphologically resembles the 6.4 keV image(figure 4.38a) more than the 6.7 keV image (Fig. 4.38b). In particular, both the 7.5–10.0 keV
and the 6.4 keV images have a local excess centered at the radio arc bubble located at about(l, b)=(0.13◦ , –0.11◦ ). The correlation between the 6.4 keV emission and the radio arc bubble
was found by Rodrı́guez-Fernández et al. (2001). Also, the 6.7 keV image has the dominating
peak at the Saggitarius A East, while the 6.4 keV and 7.5–10.0 keV images have extended
emission across the field with a similar surface brightness.

Declination

96

CHAPTER 4. OBSERVATIONS & RESULTS

46:00
48:00
−28:50:00
52:00
54:00
56:00
58:00
−29:00:00
02:00
04:00

(a) 6.3−6.5 keV(b) 6.6−6.8 keV

46:30

46:30

17:46:00
45:30
Right ascension

17:46:00
45:30
Right ascension

Declination

46:00
48:00 (c) 7.5−10 keV−28:50:00
52:00
54:00
56:00
58:00
−29:00:00
02:00
04:00
46:30

17:46:00
45:30
Right ascension

Figure 4.38: XIS images with a smoothing of σ = 800 in (a) the 6.3–6.5 keV band, (b) the 6.6–
6.8 keV band, and (c) 7.5–10.0 keV band. The NXB is subtraceted and the exposure correction isapplied.

4.4.4

Spectral Analysis

The source photons are collected from a circle centered at the local intensity peak of
the XIS image. The background are from a position devoid of intense diffuse emission andat a similar off-axis angle. We examined several position for the background and found that

97

Counts s−1 keV−1

4.4. ARCHES CLUSTER

0.01(a) 2005 Sep.

(b) 2006 Sep.

(c) 2007 Sep.

−3

10

−4

χ

10

3
0
−3
5
Energy (keV)

105Energy (keV)

105Energy (keV)

10

Figure 4.39: Background-subtracted spectra of the Arches cluster in 3–12 keV. The thermal,
power-law, and Gaussian components are shown in red, green, and blue, respectively.

the spectrum above ∼ 3 keV is intense to the background, On the other hand, the spectrum
below ∼ 3 keV is contaminated by inhomogeneous diffuse emission prevailing the Galactic
center region, and hence is ignored hereafter. A radius of 1.0 6 is used for the background,while a smaller radius of 1.0 4 is used for the source region to increase the signal-to-noise ratio.
Figure 4.39 shows the background-subtracted spectra in each observation. The spectra
shows intense emission lines in the 6.0–8.0 keV band over a hard continuum. The threemost conspicuous lines are attributable to the Kα line from highly ionized iron (Fe XXV) at∼6.7 keV, and Kα and Kβ lines from neutral or low-ionized iron (Fe I) at ∼6.4 and ∼7.1 keV,
respectively. The existence of Fe XXV Kα in addition to Ca XIX Kα (∼3.9 keV) and Ni XXVII
Kα+ Fe XXV Kβ (7.8–7.9 keV) lines indicates a thermal plasma with a temperature of a few
keV, while Fe I K lines and a hard continuum up to ∼12 keV suggest additional components.
The events of the first and second observations were summed, while the third and forth
observations were treated separately. We first fitted the summed first and second spectrumwith photo-electrically absorbed thin-thermal plasma (APEC; Smith et al. 2001) model todetermine the plasma temperature (kB T ) and the amount of interstellar extinction (NH ).
The abundance relative to solar was derived separately for elements with prominent
emission lines. Two Gaussian components were added to the model to account for the linesat ∼ 6.4 and ∼ 7.1 keV. Their line center and flux of the 7.1 keV were fixed to 1.103 and0.113 times those of the 6.4 keV line. The widths of the lines were consistent with zero.
The resultant best-fit model was marginally acceptable with the null hypothesis probabilityof the χ2 value (Pχ2 ) of 0.01. A significant residual above ∼7 keV obviously requires anadditional component. We added a power-law component to reach an acceptable fit. An
additional thermal component also yielded an acceptable fit, but the best-fit temperature of> 20 keV is unphysical. We therefore continued with the power-law component.

98χCounts s−1 keV−1

CHAPTER 4. OBSERVATIONS & RESULTS

0.01
10−3
10−4
3
0
−3
5
Energy (keV)

10

Figure 4.40: XIS spectrum of the Arches cluster. Data from four observations are summed.
Symbols follow figure 4.39

The final fit result (Pχ2 =0.28) is summarized in table 4.19. Here, we fixed the abundance
to be solar for all the elements upon the confirmation that the best-fit value is consistent
with solar. By separating the Fe I Kβ line at ∼7.05 keV, the Fe XXVI Kα line at ∼ 6.97 keVand the Fe K edge at ∼ 7.11 keV are found to be absent. The absence of the Fe XXVI line
as well as the existence of the Fe XXV line contributes for a tight constraint of the plasma
temperature. The NH is common for all the components. The fitting results are shown intable 4.19.
For the third and fourth observations, we fixed NH , kB T , and Γ to the best-fit values
derived above due to the limited photon statistics. The normalization of the thermal andpower-law components and Gaussian line parameters were consistent with the first andsecond observations.
Finally, we summed all the events from four observations, and performed spectral analysis in the same way of the first and second observation data. We obtained more constrained
parameters (table 4.19).

Observation

1st & 2nd
3rd
4th
All merged∗

Absorption
Thermal
NH
kB T
fth ∗
(1023 cm−2 )(keV)(erg s−1 cm−2 )
+0.5
+1.8
−13
1.4−0.4
2.2−0.7
5+8−3 × 10
1.4 (fixed) 2.2 (fixed) 8 ± 2 × 10−13
1.4 (fixed) 2.2 (fixed) 7 ± 3 × 10−13+0.6−131.6 ± 0.4
1.9−0.4
6+7−3 × 10

Power-lawΓfPL ∗(erg s−1 cm−2 )−130.72 ± 0.7
7+21−5 × 100.72 (fixed) 5 ± 2 × 10−130.72 (fixed) 8 ± 4 × 10−13−130.60+0.74
7+14
−0.31
−4 × 10Gaussian
E
N
−5
(eV)
(10 cm−2 s−1 )6407+72.1+0.4−5−0.3
6408+142.6 ± 0.4
−166399+172.7+0.5−23−0.6+96404−82.4 ± 0.3

4.4. ARCHES CLUSTER

Table 4.19: Best-fit parameters for the Arches cluster

Values in the 3.0–10.0 keV band.

99

Chapter 5
Discussion
In this chapter, we discuss the nature and the X-ray production mechanisms of diffuse
and discrete sources. In section 5.1, making use of the infrared data, we discuss the natureof the bright Suzaku X-ray source in the Galactic center to be a WR star. Section 5.2 isdevoted for the interpretaion of the X-ray emission mechanisms of the discrete and diffuse
sources in the Carina Nebula. In section 5.3, we speculate the nature of the newly detected
three discrete sources. We also discuss the X-ray emission mechanisms of previously known
three sources and the diffuse sources. We study the X-ray emission mechanisms of discreteand diffuse sources in the Arches cluster in section 5.4. The cross-regional study ispresented in section 5.5.

Contents
5.1

5.2

5.3

A New Wolf-Rayet Star in the Galactic Center . . . . . . . . . 102
5.1.1

Chandra and XMM-Newton Counterparts . . . . . . . . . . . . . . 102

5.1.2

Longer Wavelength Data . . . . . . . . . . . . . . . . . . . . . . . 105

5.1.3

Location of Extinction Matter . . . . . . . . . . . . . . . . . . . . 106

5.1.4

Nature of the Source . . . . . . . . . . . . . . . . . . . . . . . . . . 107

Carina Nebula . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110
5.2.1

Discrete Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110

5.2.2

Diffuse Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 112

5.2.3

The Plasma Properties . . . . . . . . . . . . . . . . . . . . . . . . . 112

M 17 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 113
5.3.1

Discrete Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . 113

5.3.2

Diffuse Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 117

101

102

CHAPTER 5. DISCUSSION
5.3.3
5.4

HXD signals . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 122

Arches Cluster . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 124
5.4.1

Discrete Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . 124

5.4.2

Origin of Three Spectral Components . . . . . . . . . . . . . . . . 130

5.4.3

Diffuse Medium . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 131

5.4.4

Cause of Line and Power-law Emission . . . . . . . . . . . . . . . . 131

5.1

A New Wolf-Rayet Star in the Galactic Center

5.1.1

Chandra and XMM-Newton Counterparts
Table 5.1: Observation log of the Sgr North region.

Start Date

Observatory

2000-09-23
2000-10-27
2001-07-16
2001-07-16
2001-07-16
2003-03-12
2006-09-21
2006-09-24

XMM-Newton 0112970201
Chandra1036
Chandra2271
Chandra2274
Chandra2285
XMM-Newton 0144220101
Suzaku501040010
Suzaku501040020

∗

ObsID

R. A.
Decl.
(J2000.0)hm
17 47 21s –28◦ 090 0200
17h 47m 22s –28◦ 110 3600
17h 47m 28s –28◦ 160 2900
17h 46m 42s –28◦ 100 2300
17h 46m 18s –28◦ 200 2300
17h 47m 23s –28◦ 090 1500
17h 46m 46s –28◦ 220 5100
17h 46m 46s –28◦ 220 5100

texp ∗
(ks)
11
35
10
10
10
34
70
50

θ†
(0 )
11.7
10.5
9.5
5.7
7.9
10.4
7.4
7.4

Exposure time after removing periods with high background level. The exposure times

for the XMM-Newton observations refer to those obtained with the EPIC-pn camera.
Angular distance of the source from the optical axis.

†

We retrieved the archived Chandra (Weisskopf et al., 2002) and XMM-Newton (Jansen
et al., 2001) data to study the long-term behaviors of the source, and to locate its position
more precisely using telescopes of smaller HPDs than that of Suzaku. Four Chandra observations using the Advanced CCD Imaging Spectrometer (ACIS: Garmire et al. 2003) and
two XMM-Newton observations using the European Photon Imaging Camera (EPIC: Turner
et al. 2001; Strüder et al. 2001), which is comprised of two MOS and a PN, were found tocover the Suzaku source (table 5.1).
Within the Suzaku positional uncertainty of 5000 , we found only one Chandra
source (CXOGC J174645.3–281546) in Muno et al. (2006b) and one XMM-Newton source

103

Counts s−1 keV−1

5.1. A NEW WOLF-RAYET STAR IN THE GALACTIC CENTER

0.01

10−3

χ

Fe

10−4
2
0
−2
2

5
Energy (keV)

10

Figure 5.1: Background-subtracted EPIC spectra. The pn spectrum is shown in red, while themerged MOS spectrum is in black. The symbols follow figure 4.4.
Table 5.2: Best-fit spectral parameters.∗

Parameter
NH
kB T
ZAr
ZCa
ZFe
FX †
LX ‡
χ2 /d.o.f.

Unit
XMM-Newton+0.4−2(10 cm )2.3−0.3+0.7(keV)
3.8−0.7+7.1(solar)2.9−2.9+2.8(solar)0.9−0.9+0.2(solar)0.9−0.2+0.07(10−12 erg s−1 cm−2 )1.05−0.07(1034 erg s−1 )2.9
69.9/74
23∗

The uncertainty indicates the 90% confidence ranges of the fit.
Energy flux in the 2.0–8.0 keV band.
‡
Absorption-corrected luminosity in the 2.0–8.0 keV band.
A distance of 8 kpc is assumed.
†(2XMMp J174645.2–281547) in the Second XMM-Newton Serendipitous Source Pre-release
Catalogue, XMM-Newton Survey Science Centre (2006). Also, hereafter we refer the Suzaku
source as CXOGC J174645.3–281546.
Since no spectroscopic and temporal behaviors are presented for both of the Chandra

104

CHAPTER 5. DISCUSSION

Flux (10 −13 erg s −1 cm −2 )

15

10

٠

5

0
2000

2002
2004
2006
Observation date (year)

2008

Figure 5.2: Long-term trend of X-ray flux in the 2.0–8.0 keV using XMM-Newton (squares),
Chandra (circles), and Suzaku (a diamond). Error bars on the data points are plotted at the 90%
confidence level.

and XMM-Newton sources in the literatures, we analyze the data and present the results here.
One of the two XMM-Newton observations showed a net exposure of ∼ 34 ks after removing
data during high background. Spectra with high statistics were obtained from MOS and PN,
for which we conducted spectral fits in a similar manner as with the Suzaku spectrum. The
best-fit model and parameters are shown in figure 5.1 and table 5.2, respectively. Although
the emission lines are less conspicuous in the XMM-Newton data, the Suzaku and XMMNewton results are consistent with each other.
For the remaining one XMM-Newton and four Chandra observations, the photon statistics were too poor for a detailed spectral analysis, due to short exposures, chip gaps, andlarge off-axis angles of the source position. We therefore applied the best-fit Suzaku model
to derive their flux.
Figure 5.2 shows a long-term flux variation, which spans ∼ 6 years at seven different
epochs. Although no flux variation within each observation is found, a long-term variation
having a factor ∼ 2 is clearly found.

105

5.1. A NEW WOLF-RAYET STAR IN THE GALACTIC CENTER

5.1.2

Longer Wavelength Data

Declination−28:15:00
30
16:00
30

30"
1.2 pc @ 8 kpc

17:46:48
44
Right ascension

40

Figure 5.3: 2MASS image around CXOGC J174645.3–281546. J, H, and K band intensities areshown in blu, green, and red, respectively.

We retrieved other databases to characterize the multi-wavelength features of
CXOGC J174645.3–281546. Because of the extreme extinction, the source is inaccessible
in the optical bands.
In the near-infrared (NIR) bands, we examined the Two Micron All-Sky Survey (2MASS:
Cutri et al. 2003; Skrutskie et al. 2006) and the NIR Galactic center survey (Nishiyama et al.,2006) using the Simultaneous Infrared Imager for Unbiased Survey (SIRIUS: Nagashima etal. 1999; Nagayama et al. 2003) on the Infrared Survey Facility (IRSF) telescope. In the midinfrared (MIR) bands, we used the Midcourse Space Experiment (MSX: Egan et al. 2003)
and the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire program (GLIMPSE:
Benjamin et al. 2003) using the Spitzer Space Telescope (Werner et al., 2004).
Figure 5.3 shows an NIR pseudo-color image, in which an isolated point-like source(2MASS J17464524–2815476 and MSX C6 G000.7036+00.1375) is found in the 2MASS and

106

CHAPTER 5. DISCUSSION

MSX images within the positional uncertainty range of the Chandra source. The sourceis very bright and red with 2MASS magnitudes of (J, H, K s ) = (> 15.4, 10.0, > 7.2) mag.
The J - and K s -band magnitudes are upper limits due to nearby source contamination. In
the SIRIUS data, while the Ks -band is unavailable due to saturation, the magnitudes of (J,H ) = (15.53, 10.05) were derived with photometric errors of ∼ 0.01 mag.

Intensity (10−11 erg s−1 cm−2 µm−1)

In the MIR bands, the source is in the linearity range in the MSX photometry, but is toobright in Spitzer images, from which no meaningful photometry was obtained. Figure 5.4
shows the spectral energy distribution (SED) in the 1.2–30 µm band. We fitted the datawith a single-temperature blackbody and an assumed extinction proportional to the inverse square of the wavelength. We found the best-fit blackbody temperature (TBB ) of
980 ± 20 K, with a visual extinction (AV ) of 31 ± 1 mag, and a bolometric luminosity (Lbol )of (8.3 ± 0.5)×104 L (d/8.0 [kpc])2 .

1000
100
10
1
0.1

SIRIUS
MSX

1

2

5

10

20

Wavelength (µm)
Figure 5.4: SED constructed from SIRIUS (open squares) and MSX (open triangles) photometry.
The solid curve shows the best-fit blackbody model with TBB = 980 K, AV = 31 mag and Lbol =8.3 × 104 L . SEDs of the eponymous Quintuplet cluster members (Okuda et al., 1990; Figer etal., 1999a) are also shown with colors for comparison.

5.1.3

Location of Extinction Matter

Adopting NH /AV = 1.79 × 1021 cm−2 mag−1 (Predehl & Schmitt, 1995), AV = 31 mag isconverted to NH ∼ 5.6 × 1022 cm−2 . This is far smaller than that determined with the X-rayof NH ∼ 2.4 × 1023 cm−2 . Since the visual extinction of 31 mag is the typical value toward5.1. A NEW WOLF-RAYET STAR IN THE GALACTIC CENTER

107

the Galactic center (Catchpole et al., 1990), the excess X-ray absorption is probably due tolocal obscuring matter that radiates the detected IR emission. If such an obscuring matterwith solar abundance spherically surrounds the X-ray source, the 6.4 keV line with an EW of∼ 200 eV should be detected (Inoue, 1985). The lack of 6.4 keV line (EW < 50 eV) requires< 0.25 solar or that the local matter is concentrated ineither that the iron abundance is ∼front of the source along the line of sight.

5.1.4

Nature of the Source

WR Binary Origin
The SED of CXOGC J174645.3–281546 in the IR bands is very similar to those of theeponymous Quintuplet cluster members (figure 5.4). The five stars have cool (T ∼ 1000 K)
spectra and are bright (∼ 7 mag in the K-band). The deep absorption at ∼ 10 µm is due tointerstellar silicate. The lack of either emission lines or intrinsic absorption features (Okuda
et al., 1989; Figer et al., 1999a) allows no spectral classification, though their luminosityof ∼ 105 L corresponds to those of supergiant stars. These intriguing stars were recently
spatially resolved (Tuthill et al., 2006). Two (GCS 3–2 and GCS 4 of Nagata et al. 1990, or
Q2 and Q3 of Moneti et al. 2001) out of five showed beautiful pinwheel nebulae of dust plume,
which are seen in a few WC stars (Tuthill et al., 1999; Monnier et al., 1999). Circumstellar
dust emission with temperatures of 700–1700 K is a common character of WC stars (Williams
et al., 1987).
Interestingly, the two sources that exhibit the pinwheel morphology are the brightest
hard X-ray sources (Law & Yusef-Zadeh, 2004; Wang et al., 2006) among the Quintuplet
members. Since both the pinwheel dust plumes and the hard X-ray emission indicate fast
wind-wind collision (Tuthill et al., 1999; Monnier et al., 1999; Oskinova et al., 2003), thiscoincidence is not probably accidental. The summed spectrum of three X-ray emitting
sources (GCS 3-2, GCS 4, and source D of Nagata et al. 1990) in the Quintuplet cluster isvery hard (kT ∼ 9 keV) and shows a hint of iron K emission line, but the luminosity is smallerby ∼ 2 orders of magnitude (7.6×1032 erg s−1 in the 0.3–8.0 keV band; Wang et al. 2006)
than CXOGC J174645.3–281546. In the colliding wind binary scenario, the X-ray luminosity
scales as LX ∝ D−1 (Stevens et al., 1992), where D is the binary orbital separation. The
large X-ray luminosity of CXOGC J174645.3–281546 indicates the close separation, and themoderate flux variation having a factor of ∼ 2 can be interpreted as the variation of theseparation due to the orbital motion.
Some of evolved massive star+OB star systems have very large X-ray luminosities up to> 2 keV∼ 10 erg s−1 in the 0.5–10.0 keV band with thermal spectra with a temperature of ∼
35

108

CHAPTER 5. DISCUSSION

and an iron abundance of 1–2 (η Carinae: Tsuboi et al. 1997, Hamaguchi et al. 2007b;
WR 140: Koyama et al. 1994, Pollock et al. 2005, A1 S, A1 N, and A2 in the Arches cluster;
Wang et al. 2006). The X-ray spectral property of CXOGC J174645.3–281546 coincides
with those of these massive star binaries, though near- and mid-IR spectroscopy and highspatial resolution observations are essential for further constraints. All of the observational
results are consistent with the idea that CXOGC J174645.3–281546 is a WC star+massive
star binary near the Galactic center, which has local obscuring matter (figure 5.5).

NH ~ 2e23 cm-2

ISM
NH ~ 6e22 cm-2

IR

X-ray

massive star?

50 MK
WC star?

980 K

Figure 5.5: Cartoon of the source system. Winds from the primary WC star and the companion
massive star collide to produce X-rays. The obscuring material of NH ∼ 2 × 1023 cm−2 aroundthe binary system is heated to T ∼ 980 K by photospheric emission to emit IR. Whereas theIR emission suffers only interstellar medium (ISM) extinction, the X-ray emission is subject toadditional intrinsic absorption.

Because the initial mass of a WR star exceeds ∼20 M and the age is less than ∼ 6 Myr(Maeder & Meynet, 1987; Massey et al., 1995; Meynet & Maeder, 2005), we naturally expectthat WR stars are found in large young star clusters, assuming a usual initial mass function.
Our source, however, seems to be isolated (figure 5.3). Muno et al. (2006a) and Mauerhan etal. (2007) also recently discovered such isolated massive stars in the Galactic center region.
One possibility is that these stars were formed initially in a cluster, and the cluster dissipated
due to the strong tidal force of the Galactic center (Portegies Zwart et al., 2002a). Future
IR and X-ray observations with a higher spatial resolution and sensitivity may reveal that

5.1. A NEW WOLF-RAYET STAR IN THE GALACTIC CENTER

109

either this source is really a cluster member, or truly an isolated star.
Other Possible Origins
The most populous class of hard X-ray sources that exhibits iron K-shell emission lines is cataclysmic variables (CVs). The X-ray luminosity of CXOGC J174645.3–
281546 of 3 × 1034 erg s−1 is at the brightest end of CVs. The bolometric luminosity of∼ 105 L (d/8.0 [kpc])2 , however, is far larger than CVs, because most of the optical companions of CVs are late-type main-sequence stars. The spectra of CVs are characterized byhard continuum, with three iron K-shell emission lines at 6.4 keV, 6.7 keV, and 6.97 keV.
The equivalent widths of both the 6.4 keV and 6.7 keV lines are, however, around 100–
200 eV, although in some cases that of the 6.7 keV line shows an exceptionally high value
possibly due to resonance scattering (Terada et al., 2001). The 6.97 keV/6.7 keV flux ratioof CVs is larger than 0.1, indicating that the ionization temperature is higher than ∼ 4 keV(Ezuka & Ishida, 1999; Hellier & Mukai, 2004; Rana et al., 2006). CXOGC J174645.3–281546
has a significantly larger equivalent width of the 6.7 keV line (∼ 1 keV) and lower flux ratioof the 6.97 keV/6.7 keV lines than those of CVs. The upper limit of 50 eV for the 6.4 keVline, on the other hand, can not give any constraint on the possibility of a CV. Altogether,
CXOGC J174645.3–281546 is very unlikely to be a CV.
The X-ray spectra of YSOs are characterized by a thin thermal plasma with a temperature of 1–5 keV and a metal abundance of around 0.3 solar in the quiescent phase (Feigelson
et al., 2002; Imanishi et al., 2003; Ozawa et al., 2005). The CXOGC J174645.3–281546 ironabundance of 0.8 ± 0.1 solar is higher than the typical YSO value of 0.3 solar, but is not
exceptional. Although the ratio of the X-ray and bolometric luminosity (L X /Lbol ) of 10−5
is in the range of low-mass YSOs (Imanishi et al., 2001; Feigelson et al., 2002), the absolute luminosities are too high for a YSO if CXOGC J174645.3–281546 is located at 8 kpc.
The possibility still remains that CXOGC J174645.3–281546 is a foreground YSO in a dense
molecular cloud.

110

CHAPTER 5. DISCUSSION

5.2

Carina Nebula

5.2.1

Discrete Sources

WR 25
WR 25 is a very bright WR star in the optical band (V ∼ 8.1 mag). From the WNtype emission lines and O-type absorption lines (Smith, 1968), this object is classified as a
WN6h+O4f binary (Smith et al., 1996; van der Hucht, 2001). The X-rays from WR 25 weredetected by Einstein for the first time (Seward et al., 1979). The luminosity is unusuallyhigh (∼ 1034 ergs s−1 in the 0.5–10.0 keV band; Raassen et al. 2003) and shows a year-scalevariability (Pollock & Corcoran, 2006). After the report by Pollock & Corcoran (2006),
Gamen et al. (2006, 2008) discovered a periodicity in radial velocity of the N IV emission lineat λ =4058 Å for the first time, which should arises from the binary motion. The orbital
period is 207.7 ± 0.1 days and the eccentricity is 0.56 ± 0.01.

Flux (10−12 erg s−1 cm−2)

108

FX=F0(1+ecosφ)

6
4
2
0

0

0.5
φ

1

Figure 5.6: Long-term flux variation of WR 25 as a function of orbital phase. Best-fit Stevens’model is shown with a solid line.

Figure 5.6 shows the long-term X-ray flux variation of WR 25 as a function of theorbital phase (Gamen et al., 2006). The X-ray flux clearly depends on the orbital phase.
Stevens et al. (1992) argued that the X-ray luminosity is proportional to the inverse of thebinary orbital separation in the wind collision scenario. Since the binary separation D isexpressed as D ∝ (1 + e cos φ)−1 , the X-ray flux from a colliding-wind binary (FCWB ) scales5.2. CARINA NEBULA

111

as FCWB ∝ 1 + e cos φ. Allowing e to vary, we fitted the light curve with the above function(figure 5.6). The best-fit value of e was ∼ 0.3, which is significantly lower than that derived
by Gamen et al. (2008). This indicates that the wind collision alone is not responsible for
the X-ray emission.
To separate the X-ray flux into the intrinsic and colliding-wind component, we
fixed e to 0.56 and added a constant value, i.e., we fitted the flux with a function
FX = F0 + F1 (1 + 0.56 cos φ). The best-fit values are F0 = 2.2 ± 0.1 × 10−12 and F1 =4.1 ± 0.1 × 10−12 erg s−1 cm−2 in the 1.0–8.0 keV band.
HD 93205
Although HD 93205 is a binary of two high mass stars (O3 V + O8 V; Conti & Walborn1976) with a large eccentricity (e = 0.42; Morrell et al. 2001), the X-ray emission is soft andsteady, as was pointed out by a previous study (Albacete-Colombo et al., 2008). Assumingthat Lbol = 1.3 × 106 L , we obtain LX /Lbol ∼ 1 × 10−7 , the typical value of massive stars(Lucy & White, 1980). HD 93205 does not show a signature of wind collision probably
because they are too close (a1 sin i ∼ 1.0 × 107 km and a2 sin i ∼ 2.4 × 107 km; Morrell et al.
2001). If we adopt the inclination angle of ∼ 60◦ calculated by Benvenuto et al. (2002), theseparation is comparable to the radius of an O3 V star, 15 R (Cox, 2000). Since the criticalpoint of stellar wind acceleration is located at the distance of ∼ 6 times the stellar radius(Abbott 1980, figure 2.6), the stellar winds from the two stars probably collide before they areaccelerated enough to produce hard X-rays. Morrell et al. (2001) reported an apsidal motion
at a rate of 0.◦ 032 ± 0.◦ 0031 per cycle. They argued that tidal and rotational distortions is
the main cause of this apsidal motion, which also supports the closeness of the binary.
HD 93250
HD 93250 has some features similar to WR 25. The temperatures of HD 93250 are
kB T low = 0.62 (0.59–0.65) keV and kB Thigh = 3.1 (2.9–3.3) keV, while those of WR 25 are
kB T low = 0.62 (0.61–0.63) keV and kB Thigh = 2.33 (2.27–2.39) keV (table 4.9). The Fe abundance in HD 93250 is 0.36 (0.22–0.50) solar and 0.28 (0.25–0.31) solar in WR 25. The only
clear difference is the X-ray luminosity with a factor of 4–8 (depending on the orbital phaseof WR 25). These similarities in the spectral shapes indicate that some fraction of the X-ray
in HD 93250 arises from wind collision, although there is no evidence for binarity from opticalobservations.

112

5.2.2

CHAPTER 5. DISCUSSION

Diffuse Sources

Comparison with Previous Results
In section 4.2., we selected three regions according to the morphology in the 0.5–2.0 keVband (figure 4.8), and performed spectral analysis. The spectra were fitted with single
temperature plasma model and emission above 2 keV was not found. On the other hand,the spectra in Hamaguchi et al. (2007) have harder components. This difference probably
arises from the choice of photon extraction regions. Hamaguchi et al. (2007) chose larger
regions with areas of ∼ 110 and ∼ 70 arcmin2 , while we accumulated photons from smaller
regions that have high surface brightness with areas of ∼ 10–20 arcmin2 to obtain higher
signal-to-noise ratios. Hamaguchi et al. (2007) attributed the difference of spectral shapein north and south regions to the difference of the metal abundance. However, we foundthat the absorption columns and temperatures also show significant spatial variations. The
temperature of diffuse 1 and 2 regions is ∼ 0.55 keV and the absorption value is small, while
the temperature of diffuse 3 region is lower (∼ 0.18 keV) and the absorption column is typicalof this field (table 4.10).

5.2.3

The Plasma Properties

Since the absorption columns of diffuse 1 and 2 are significantly smaller than othersources in this region, it is likely that the emitting source is located in front of the Carina
Nebula and may not be physically related to it. The iron abundance is ∼ 5 times larger than
Ne or O which implies a type Ia SN origin.
Diffuse 3 is likely to be located at the Carina Nebula distance from its absorption column.
Assuming a distance of 2.3 kpc, we can estimate the plasma mass. From the spectral analysis,
a volume emission measure of ∼ 1058 cm−3 is obtained. If we further assume that the plasma
is uniformly distributed in an ellipsoid of 1 pc×1 pc×2 pc, the electron density and theplasma mass are 5 cm−3 and 1.5 M . This amount of plasma can be fueled either by WR 25
with Ṁ ∼ 3.6 × 10−5 M yr−1 (Crowther et al., 1995) or O stars in Tr 14 and Tr 16 within
105 yr. We see no N emission line at 0.5 keV and obtained upper-limit values of N/O < 0.5
and N/Fe < 0.3.

113

5.3. M 17

M 17

5.3.1

Discrete Sources

Intensity (10−11 erg s−1 cm−2 µm−1)

5.3

100
Suzaku J182132−1613.4

10
Suzaku J182132−16130.4

1
Suzaku J182140−16104.1

0.1
0.01

1

2
5
Wavelength (µm)

10Figure 5.7: NIR to MIR SED of three point sources in the M17 east field. 2MASS and Spitzer of
Suzaku J182132–1610.4, Suzaku J182132–1613.4, and Suzaku J182140–1604.1 are shown in black,
red, and blue squares respectively.

We detected three point sources in the M17 East field. In order to identify their nature,
we retrieved databases of other wavelengths. We examined 2MASS All-sky Point Source
Catalog (Cutri et al., 2003; Skrutskie et al., 2006) and GLIMPSE (Benjamin et al., 2003).
We found counterparts for all three sources in both catalogs (table 5.3). Their SEDs areshown in figure 5.7.
Suzaku J182132–1610.4
The absorption column (NH ) of Suzaku J182132–1610.4 (table 4.14) is consistent withthe typical value of AV = 2 mag in this region (Hoffmeister et al., 2008) using NH /AV =1.79 × 1021 cm−2 mag−1 (Predehl & Schmitt, 1995). After correcting for the distance modulus and the extinction, the magnitude and the color indicate that this is a B8 V star. Since

114

CHAPTER 5. DISCUSSION

intermediate mass stars are dim in the X-ray and Suzaku J182132–1610.4 shows flux variability with a time-scale of ∼ 10 ks similar to those seen in young stellar objects (Imanishi
et al., 2003; Getman et al., 2008), a probable nature is a foreground (∼ 200 pc) young stellar
object. In this case, the absorption is mainly of an intrinsic origin but not of ISM origin.
The absorption column is typical of young stellar objects (Feigelson et al., 2005).
Suzaku J182132–1613.4
The NH value of Suzaku J182132–1613.4 is also typical to this region. The J, H, andK magnitudes of Suzaku J182132–1613.4 are consistent with a B0 V star at a distance of
2.1 kpc. However, the X-ray spectrum is too hard and has a too high LX /Lbol value (∼ 10−6 )compared with the conventional value of ∼ 10−7 (Lucy & White, 1980; Pallavicini et al.,1981). We speculate that this object is also a binary comprised of a B0 V star and anunresolved massive companion.
Suzaku J182140–1604.1
This source, in contrast to the previous two sources, shows a high NH value and a verystrong iron K emission line with an EW of ∼ 2 keV. The probable counterpart shows red
excess (J − H ∼ 1.85 mag). From the NIR brightness, it is unlikely to be a background active
galactic nucleus. A possible nature is an embedded protostar.
OI 352
OI 352 is classified as an O8 V star based on U BV photometry (Ogura & Ishida, 1976).
This source is too bright and too hard for a single O8 V star, which indicates the X-ray
emission is of wind collision origin. The absorption column is larger than those of diffuse
sources and diffuse emission is seen in the NIR image (figure 4.20), implying that this sourceis embedded in a molecular cloud.
M17 North
M17 North is not point like with the PSF of XRT. Broos et al. (2007) detected ∼ 10
sources from this region. Together with the fact that the absorption column (N H = 2.4 ±0.3 × 1022 cm−2 ) is consistent with those of molecular cloud clumps A8–A13 in Wilson etal. (2003), M17 North is likely to be embedded young stellar objects. The temperature+0.5(kB T = 2.7−0.3
keV, table 4.13) is a typical value of class-I protostars (Imanishi et al., 2003).

5.3. M 17

115

OB association
We detected Kα emission lines from highly ionized ions (figure 4.31c). The spectrumof the OB association is comprised of hundreds of point sources, but the emission from an
O4–O5 binary dominates the spectrum. Broos et al. (2007) claimed that both componentsof the binary (sources 543 and 536) have plasma temperatures exceeding 10 keV. However,the strong Fe XXV Kα line at 6.7 keV and the weak Fe XXVI Kα line at 7.0 keV in the XIS
spectrum (figure 5.10c) do not support such high temperatures.

Source ID
Suzaku J182132–1610.4
Suzaku J182132–1613.4
Suzaku J182140–1604.1. . . . . . . . . . . . 2MASS source . . . . . . . . . . . .
J (mag)H (mag)
K (mag)2MASS J18213225–1610285
12.42 ± 0.02 11.58 ± 0.02 11.34 ± 0.02
2MASS J18213201–1613239
9.66 ± 0.02 9.16 ± 0.02 8.73 ± 0.02
2MASS J18214023–1604128
14.56 ± 0.03 12.71 ± 0.02 12.00 ± 0.02

116

Table 5.3: Counterparts of three unidentified point sources in the M17 East field. . . . . . . . . . . . . . . . . . . . . . . Spitzer source . . . . . . . . . . . . . . . . . . . . . . .
3.6 µm (mag) 4.5 µm (mag) 5.8 µm (mag) 8.0 µm (mag)SSTGLMC G015.1779–00.9065
11.06 ± 0.05
11.05 ± 0.08
11.00 ± 0.12
10.68 ± 0.22
SSTGLMC G015.1334–00.9285
8.21 ± 0.04
7.93 ± 0.04
7.65 ± 0.4
7.25 ± 0.03
SSTGLMC G015.2850–00.8857
11.44 ± 0.05
11.43 ± 0.09
10.56 ± 0.11—

CHAPTER 5. DISCUSSION

5.3. M 17

5.3.2

117

Diffuse Sources

Contamination to the Extended Emission
The spectra of the diffuse emission in figure 4.33 are contaminated by other sources ofemission. We evaluate the levels of contamination by unresolved point sources, the Galactic
Ridge X-ray emission (GRXE), the cosmic X-ray background (CXB), and the local hot
bubble (LHB), and argue that their contributions to the background-subtracted spectra are
negligible. About 70% of the XIS field is covered in the Chandra observation which has amuch better spatial resolution and sensitivity for faint point sources (Townsley et al., 2003;
Broos et al., 2007). Among the five sub-regions, the region 2 has a complete coverage by
Chandra. We therefore use this sub-region as a representative to evaluation the levels of
various contaminations.
First, we examine the contribution of unresolved point sources. We extracted 19 point
sources from the ACIS data in the region 2, constructed the composite spectrum, and fitted
it with a thermal plasma model. The spectrum of each source is too poor to fit individually,so we assume that all sources have the same spectral shape that best describes the composite
spectrum. Using their positions, flux, and the assumed spectral shape, we generated their
XIS events using a ray-tracing simulator (xissim; Ishisaki et al. 2007). In figure 5.8, wecompare the integral of the simulated spectra of unresolved point sources (PS1) to theobserved diffuse spectrum (pluses). In the displayed observed spectrum, we subtracted the
NXB spectrum but not the background spectrum in the neighboring region. The point source
contribution accounts for ∼ 8% of the emission in the 0.4–1.8 keV band.
To have an estimate of the maximum contamination to the lines, we repeated the sameprocedure using the model with the maximum allowable abundance values in the best-fit
model for simulating unresolved point source events. The resultant integrated spectrum(PS2) is also shown in figure 5.8. At the Ne IX line at 0.92 keV, the contribution by the
point sources accounts only for <10% of the observed emission.
Second, we derive the GRXE contribution, which is ubiquitous along the Galactic Plane< 45◦ (Kaneda et al., 1997; Sugizaki et al., 2001). We refer
at Galactic longitudes (l) of |l| ∼to a Chandra result (Ebisawa et al., 2005) for the spectral shape in the soft band and to a
Rossi X-ray Timing Explorer result (Revnivtsev et al., 2006) for the surface brightness atthe position of M17 (l ∼ 15◦ ). The GRXE contribution to the observed emission is estimatedto be ∼ 3% and is shown in figure 5.8.
Third, for the CXB contribution, we consulted the Suzaku XIS observation of the North
Ecliptic Pole (Fujimoto et al., 2007) both for the spectral shape and the surface brightness.
The CXB was observed and fitted by a power-law model. We convolved the model with the

118Counts s−1keV −1

CHAPTER 5. DISCUSSION

0.1
PS2

0.01
PS1
LHB
10 −3

CXB
GRXE

0.5

1
Energy (keV)

Figure 5.8: Comparison of the observed diffuse spectrum to the simulated spectra of contaminating sources. The NXB signal is subtracted from the observed spectrum. PS1 and PS2 are forthe contribution of unresolved point sources with different spectral models, GRXE for the Galactic
Ridge X-ray emission, CXB for the cosmic X-ray background, and LHB for the local hot bubble.
None of them play a significant contribution. Moreover, most of them are removed by subtracting
a background spectrum from the neighboring region.

5.3. M 17

119

XIS responses and found that the contribution is ∼ 2%.
Finally, we constrain the contribution by the LHB emission. We estimate its surface
brightness to be ∼ 4×104 counts s−1 arcmin−2 in the PSPC R1 and R2 bands (Snowden
et al., 1998). We assume that the spectrum is a thin-thermal plasma (the Raymond-Smith
model; Raymond & Smith 1977) with a temperature of 0.1 keV. With the derived emission
measure of ∼ 2.8 × 10−3 cm−6 pc, the LHB contributes ∼ 2% of the observed emission.
None of the above components play a significant role in the observed diffuse spectrum.
Moreover, we removed most of them by subtracting a background spectrum from a neighboring region. The background events show no indication of time variability caused by solar
flares (Fujimoto et al., 2007). We conclude that the background-subtracted XIS spectrum
represents the spectrum of the uncontaminated diffuse emission quite well.
The background spectrum accounts for ∼ 35% of the source spectrum of the sub-region”diffuse 1”. Besides the NXB contributing ∼ 7% of the source flux, the sum of unresolved
point sources, GRXE, CXB, and LHB contributes ∼ 15%. Therefore, emission with a fluxof ∼ 13% of the source flux is additionally included in the background spectrum. The contribution from the bright sources outside of the extraction region is negligible. We attributethe remaining emission to the diffuse emission in the background region. If this is the case,the flux estimate of the diffuse emission (table 4.15) is underestimated by ∼ 13%.
Comparison with Previous Studies
We compare our results with the previous works using ROSAT (Dunne et al., 2003) andChandra (Townsley et al., 2003; Broos et al., 2007). We derived that the diffuse spectrum is
explained by an absorbed single temperature thin-thermal plasma model of kB T ∼ 0.25 keV
and NH ∼ 4.5 × 1021 cm−2 . The total luminosity (0.5–2.0 keV) in the combined diffuse1–6
regions is ∼ 5.2×1033 erg s−1 .
The ROSAT study (Dunne et al., 2003) shows that the total luminosity of the diffuse
emission is ∼ 2.5×1033 erg s−1 . The smaller estimate than our result is more noticeable if
we consider that the ROSAT value was derived from a larger area and in a wider energy
range (0.1–2.4 keV). This stems from an underestimate of the extinction. Without a sufficient
spectral resolution to resolve lines, two different models were not disentangled in the ROSATPSPC spectra; one is a low plasma temperature with a large extinction (∼ 0.2 keV and∼ 1022 cm−2 ) and the other is a high plasma temperature with a small extinction (∼ 0.7 keVand ∼ 1020 − 1021 cm−2 ). Dunne et al. (2003) derived the luminosity based on the latter, but
our spectroscopy shows that the former should be in the case.
The best-fit XIS value of NH is consistent with that presented in the Chandra study

120

CHAPTER 5. DISCUSSION(Townsley et al., 2003), in which NH = (4 ± 1)×1021 cm−2 . However, the plasma temperatures are different between the two studies. In Chandra, the primary component of
kB T = 0.6 ± 0.1 keV and the secondary component of ∼ 0.13 keV were claimed. In Suzaku,
however, we confirmed that a single temperature component of ∼ 0.25 keV is adequate fromthe diagnosis of resolved oxygen and neon lines.
Spatial Difference of the Plasma Properties
A high signal-to-noise ratio spectrum by XIS enabled us to conduct spatially-resolved
spectroscopy of the diffuse emission. The plasma temperature and the chemical composition
are uniform. The Chandra study (Townsley et al., 2003) also show no evidence for spatial
variation of plasma temperature.
The observed uniformity indicates that the entire plasma is at a thermal equilibriumin the observed spatial scale of ∼ 10 pc, unless the plasma is patchy at equilibria locally bymagnetic confinement. The global equilibrium is reasonable considering the fact that theplasma sound crossing time (∼ 4 × 104 yr) is much smaller than the time scale of the system(∼ 106 yr), thus the constant pressure is achieved (Weaver et al., 1977). Here, we used theplasma volume and the electron density as ∼ 100 pc3 and ∼ 1 cm−3 , respectively, by assumingthat the plasma distribution has a conical shape with its apex at the OB association andwith a filling factor of 1. Given the uniformity of the plasma temperature and pressure, we
speculate that the density is also spatially uniform.
The observed surface brightness, however, is different among the sub-regions (figure 4.24a, table 4.15). It is ∼ 1.3 times more intense in the region 1 than region 2. This
is not attributable entirely to the different extinction, as NH is larger in the sub-region 1
than in 2. Because the plasma has a uniform temperature and density, we speculate thatthe difference of the surface brightness is likely due to the different line-of-sight depths ordifferent filling factors.
With the derived plasma volume and the density, the total mass of the plasma is ∼ 2 M .
This is comparable to the integrated ejecta mass by stellar winds at a mass loss rate of∼ 10−6 M for ∼ 106 yr and agrees with the estimates by the previous works (Dunne et al.,
2003; Townsley et al., 2003). We also speculate that the swept-up and evaporated ISM does
not make a significant contribution to the plasma mass.
Chemical Composition
The chemical composition of the diffuse emission is revealed for the first time in thisstudy. The metallicity of the diffuse emission is 0.1–0.3 solar (table 4.15), which is signifi-

5.3. M 17

121

cantly lower than those derived for the discrete sources (0.3–1.5 solar). The metal abundances
relative to each other in the diffuse emission are consistent with solar among O, Mg, and Fe(Anders & Grevesse, 1989), but Ne is enhanced in all three sub-regions by a factor of ∼2(table 4.15). This is also evident in the simultaneous spectral fits of the three sub-regions, inwhich we tied the abundance values of these elements. The resultant values are ZO = 0.12
(0.11–0.13), ZNe = 0.22 (0.21–0.25), ZMg = 0.10 (0.08–0.12), and ZFe = 0.12 (0.11–0.13) solar.
Such Ne enhancement from other metals is widely seen in coronally active stars(Brinkman et al., 2001; Kastner et al., 2002; Audard et al., 2003; Imanishi et al., 2003;
Stelzer & Schmitt, 2004; Maggio et al., 2007) for unknown reasons. We consider that the Ne
enhancement is an intrinsic feature of the diffuse plasma, and is not influenced by the contamination of Ne-enhanced point sources spectra (§ 5.3.2). One explanation for the anomaly
is that the poorly-constrained solar Ne abundance is underestimated by a factor of a few(Drake & Testa, 2005; Liefke & Schmitt, 2006). This would account for the observed Ne
enhancement in the M17 diffuse plasma as well.
The lack of a clear spatial variation of the chemical composition comprises a sharp
contrast to the diffuse emission in the Carina Nebula (Hamaguchi et al., 2007), where thedifferent abundance patterns across a similar spatial scale suggest the supernova origin for
the emission. In the diffuse emission in M17, the O and Fe ratio is consistent with thesolar abundance, which is another line of evidence against the supernova interpretation. If a
supernova has occurred in M17, it should have been caused by a star earlier than the earliest(O4–O5) star in the OB association. Such a massive source causes a core-collapse–typesupernova. It would have yielded a measurably larger ratio of O against Fe than the solar
value by a factor of a few (Tsujimoto et al., 1995; Thielemann et al., 1996; Nomoto et al.,1997).

122

5.3.3

CHAPTER 5. DISCUSSION

HXD signals

1
Counts s−1 keV−1

10−4
100
50
0

keV (keV cm−2 s −1 keV −1 )

0.01

Ratio

(a)

0.1

10−3

(b)

0.01

10−3

10−4

1

2

5
10
Energy (keV)

30

Figure 5.9: (a) Simultaneous fit of XIS (black) and PIN (red) spectra of M17 West. The low- andhigh- temperature components are respectively shown in green and blue. The lower panel shows
the residual (data/model). (b) Unfolded νFν spectra.

In section 4.3.6, we detected significant signals with HXD/PIN in the 13–30 keV bands.
We try to simultaneously fit the XIS and PIN spectra of OB association, the brightest object
in the XIS FOV. Since the OB association is not located at the XIS nor HXD nominal
position, the effective area of PIN decreases by ∼ 10%. We calculated corresponding arf
using hxdarfgen, and applied it to the spectral analysis.

5.3. M 17

123

We fitted the XIS and PIN spectra simultaneously with two thermal components commonly modified by photoelectric absorption. Figure 5.9 shows the XIS and PIN spectra. The
PIN data show large excess by a factor of 10–50 against the XIS. Any physical model cannot
reproduce both the XIS and PIN spectra. Neither OI 352 nor M 17 North contributes to the
PIN data. We conclude that the PIN signals are from outside the XIS FOV and inside thePIN FOV. The nearest bright source is AX J1819.2–1601. (figure 4.23, Sugizaki et al. 2001).
At the position of AX J1819.2–1601, the effective area decreases by ∼ 80% (See figure 6 of
Takahashi et al. 2008). Extrapolating the power-law spectrum (table 4.16) to lower energy
band, we estimate the 2.0–8.0 keV band flux to be ∼ 2–4×10−10 erg s−1 cm−2 , which is larger
than that reported in Sugizaki et al. (2001) by more than ten-fold.
Likewise, there is no bright source in the XIS FOVs in the observations of M 17 Eastand M 17 BG. Both of the nearest INTEGRAL sources 3EG J1824–1514 and AX J18192–
1601 (figure 4.23) are out of HXD FOVs, which suggests that there are unknown bright hard
X-ray sources in the FOVs of HXD but out of XIS FOVs.

124

CHAPTER 5. DISCUSSION

5.4

Arches Cluster

5.4.1

Discrete Sources
Table 5.4: Observation Log

Obs ID
945
2276
2284
4500
∗
†

R. A.
(J2000.0)
17h 46m 20s
17h 46m 05s
17h 45m 37s
17h 45m 56s

Decl.
(J2000.0)–28◦ 520 1900
–28◦ 460 2800
–28◦ 560 2700
–28◦ 490 0100

Date
2000-07-07
2001-07-18
2001-07-18
2004-06-09

texp ∗
(ks)
49
12
11
99

θ†
(0 )
7.2
4.2
7.8
1.2

Exposure time.
Angular distance of the Arches cluster from the optical axis.

We retrieved the archived data taken by ACIS on-board Chandra and found that four
observations cover the Arches cluster. ACIS has a field of view of ∼ 170 ×170 with a PSF
radius of ∼ 0.00 5 at the optical axis and is sensitive in the 0.5–8.0 keV band with an energy
resolution of ∼ 150 eV at 6 keV and an effective area of ∼ 50 cm2 at 8 keV. The details of theobservations are summarized in table 5.4. All the datasets were collected with the Timed
Exposure mode, the nominal frame time, and the CCD temperature of –120◦ C.
Previous Chandra reports of the Arches cluster are based on ObsID=945 (Yusef-Zadeh
et al., 2002), ObsID = 945, 2276, and 2284 (Law & Yusef-Zadeh, 2004), or ObsID = 4500
(Wang et al., 2006). We combine all the data in the presented analysis. Instead, we discard
ObsID = 2284 because the observation is too short and the off-axis angle is too large. The
total integration time of the three observations is ∼ 160 ks, which is more than twice longer
than the datasets used in the previous papers. We used the level 2 data processed at the
Chandra X-ray Center. The known offset of the aspect solution1 was compensated before
merging the three datasets.
Figure 5.10 shows a close-up view of the Arches cluster. The four point-like sources werepreviously reported (Yusef-Zadeh et al., 2002; Law & Yusef-Zadeh, 2004; Wang et al., 2006).
The ACIS frame was shifted so that A1S and A1N match their radio continuum counterparts(Lang et al., 2001). The photometric properties of these sources are summarized in table 5.4.
No significant variation in the count rate is seen from all of these sources.
The three brightest sources are hard with almost all photons above 2 keV and are1

See http://cxc.harvard.edu/cal/ASPECT/fix offset/ for detail.

125

10

5.4. ARCHES CLUSTER

0.2 pc

AR5
AR9
AR4

-28:49:20

Decl. (J2000.0)

A2

A1N
AR2

A1S

AR3

AR1
AR6 56

A6

AR7

51

AR10

17:45:50
R. A. (J2000.0)

Figure 5.10: Close-up ACIS view of the Arches cluster using the best-focused ObsID=4500 data.
The 0.5–8.0 keV intensity is plotted in linear scale to emphasize point-like emission. The positionof X-ray, radio (Lang et al., 2001, 2005), and 30 brightest NIR (Figer et al., 2002) sources are
represented by diamonds, crosses and pluses, respectively. The systematic offset of the NIR sources
mentioned in (Lang et al., 2001) is corrected. Names of X-ray and radio sources are also shown.
The upper one third of the field is out of the radio studies (Lang et al., 2001, 2005).

exceptionally bright among all near-by sources. Within 1.0 4 around the cluster (the XIS
source extraction region), the three sources alone account for ∼ 80% of the total net countof all 21 point sources. There are dozens of sources with similar NIR brightness (Figer etal., 2002) yet no X-ray detections in the cluster. These NIR sources are considered to haveless than ∼ 10 counts, which is the lowest count of all the detected sources in the region.
The three sources are thus an order of magnitude brighter than A6 and at least two ordersof magnitude brighter than any other cluster members in the hard X-ray band.

126

CHAPTER 5. DISCUSSION

Table 5.5: ACIS photometry of point sources.

ID∗
A1N
A1S
A2
A6
∗

R. A.†(J2000.0)
17h 45m 50.46s
17h 45m 50.42s
17h 45m 50.26s
17h 45m 49.74s

Decl.†(J2000.0)–28◦ 490 19.200–28◦ 490 21.900–28◦ 490 11.500–28◦ 490 25.900

Csrc ‡

Cbkg ‡

Pvar §

744
1167
1031
98

10.4
10.3
20.6
25.2

1.1×10−2
2.9×10−2
4.4×10−1
3.3×10−1

MEk
(keV)
3.5
3.7
3.8
4.3

NIR#

N10, C5, B21, F7 AR4
N8, C8, B23, F6 AR1
N5, C1, B26, F9.......
F19
AR6

Identification by Law & Yusef-Zadeh (2004) except for A6, which is a new source found in this study.
The astrometric correction is applied to the X-ray data so that A1N and A1S match their VLA counterparts.
‡
Number of counts in the source region (Csrc ) and the contribution of background counts estimated froma surrounding region (Cbkg ) in 0.5–8.0 keV.
§The null hypothesis probability of the χ2 value for the fit of light curves by a constant flux model.
kMedian energy.
#
The NIR counteparts (N: Nagata et al. (1995), C: Cotera et al. (1996), B: Blum et al. (2001), and F:Figer et al. (2002)). The identification by Blum et al. (2001), Figer et al. (2002), and Lang et al. (2001)
are compiled for the complete list.
∗∗
The radio counteparts (Lang et al., 2001). A2 is out of the study field by Lang et al. (2001, 2005).
†

Radio∗∗

127

5.4. ARCHES CLUSTERχCounts s−1 keV−1χCounts s−1 keV−1χCounts s−1 keV−1

0.01(a) A1N

10−3

10−4
5
0
−5
0.01
(b) A1S

10−3

10−4
5
0
−5
0.01
(c) A2

10−3

10−4
5
0
−5
1

2
5
Energy (keV)

10

Figure 5.11: ACIS spectra and best-fit models of point sources: (a) A1N, (b) A1S, and (c) A2.
Conspicuous emission lines are shown with arrows. The lower and higher temperature component
are respectively shown with red and blue lines for A1S and A2.

128

CHAPTER 5. DISCUSSION

The spectra of the three brightest X-ray sources are shown in figure 5.11. All the spectra
are characterized by a strong cut-off at the soft end and emission lines from highly ionizedions including S XV (∼ 2.4 keV), S XVI (∼ 2.6 keV), Ar XVII (∼ 3.1 keV), Ca XIX (∼ 3.9 keV),and Fe XXV (∼ 6.7 keV). From the intensity ratio of these lines, we can have a crude estimateof appropriate models for spectral fits. The brightest two sources (A1S and A2) show the
Fe XXV line as well as the S XV line stronger than S XVI, requiring two thermal components
of different temperatures. The faintest one (A1N) shows S XVI and no prominent Fe XXV,
thus one-temperature model would be adequate.
Indeed, fits by two-temperature plasma for A1S and A2 and one-temperature plasma for
A1N yielded acceptable fits. The best-fit parameters are summarized in table 5.6. Additionalfeatures of Gaussian lines and power-law tail, as was found in the XIS spectrum, are notconfirmed in any of the spectra nor the composite spectrum of the three. The addition ofthese components in the fitting is not justified on a statistical basis, indicating that the point
source spectra lack these components.
The fit by Yusef-Zadeh et al. (2002) was conducted for the composite spectrum of A1N
and A1S, which was unavoidable due to the worse image quality of the ObsID = 945 datawith the Arches cluster at a large off-axis angle. The composite spectrum deteriorates theinformation of S XV and S XVI line intensity of each source, wherein the former is stronger in
A1S and the latter in A1N. As a result, it was impossible to decide whether a power-law or
lower temperature component was necessary in addition to a higher temperature component.
We detected three bright and a faint point-like X-ray sources in the ACIS image. Most,
if not all, of the emission in the three bright sources are from single stars (i.e., they are
not the composite of unresolved multiple bright stars) for the following reasons. First, theradial profiles of the three sources perfectly fit with the PSF. We constructed an X-ray image
deconvolved for the PSF using a maximum-likelihood reconstruction technique (Lucy, 1974),
but no additional peak was found at near-by sources in the existing NIR and radio catalogs.
Second, the positional coincidence of the X-ray and radio/NIR counterparts is superb withno measurable offset, suggesting that there are few blending sources.
The NIR identifications of the X-ray sources are given in table 5.5. Three of them
(A1N, A1S, and A2) are detected by Blum et al. (2001) in their narrow-band images. Fromthe He II (2.19 µm) and Brγ (2.17 µm) line ratio, Blum et al. (2001) discussed that thesesources have a spectral type of either OIf4 or WN7–8. The degeneracy can not be solved byimaging information alone. Figer et al. (2002) presents high-resolution K-band spectra byNIRSPEC on the Keck II telescope. We confirmed N III (2.25 µm), Brγ, and He I (2.11 µm),
He II (2.19 µm) emission as well as the deficiency of C IV (2.07 µm) in all these sources. We
conclude that the three X-ray sources are WN7–8 stars. The remaining source (A6), on theother hand, has a featureless K-band spectrum without any strong H and He emission lines

129

5.4. ARCHES CLUSTER

Table 5.6: ACIS spectroscopy of point-like sources.

ID
A1N
A1S
A2

NH ∗
(1022 cm−2 )
7.3 (6.2–8.7)
15.1 (11.5–20.5)
16 (9.1–21)

kB T high ∗kB T low ∗
FX ∗†
(keV)
(keV)
(10−13 erg s−1 cm−2 )
2.0 (1.6–2.6) . . . . . . . . . . . . 1.3 (1.2–1.4)
2.5 (2.0–3.5) 0.4 (0.2–0.7) 2.3 (2.2–2.4)
2.1 (1.5–4.5) 0.3 (0.2–0.9) 1.7 (0.52–4.4)

ZS ∗
(Z )
A1N 2.7 (1.3–5.2)
A1S 0.7 (0.0–2.5)
A2
1.5 (0.6–3.1)
ID

ZAr ∗
(Z )
2.7 (0.7–5.8)
5.6 (2.7–9.9)
4.2 (0.7–8.5)

ZFe ∗
(Z )
1.6 (0.9–2.7)
1.8 (1.1–3.2)
1.9 (1.1–2.6)

LX ‡(1033 erg s−1 )
2.3
15
6.0

P χ2 k
0.92
0.45
0.63

∗

Uncertainties in the paralentheses give the 90% confidence range.
Those without the uncertainty indicate that they are fixed values.
†
Energy flux in the 2.0–8.0 keV band.
‡
Absorption-corrected luminosity in the 2.0–8.0 keV band. A distance of 8 kpc
is assumed.
§
The null hypothesis probability of the χ2 value in the best-fit model.

(Figer et al., 2002). We consider, therefore, that it is likely an O star. It should be notedthat this is the only source in Lang et al. (2001, 2005) with a negative index at 8.5 GHz,
which is indicative of a colliding wind shock.
The three X-ray brightest sources are also among the twelve NIR brightest sources with> 120 M (Figer et al., 2002). Eleven of them (source 1–9,
an estimated initial mass of ∼11, 12 in Figer et al. 2002) are WN7–8 and the remaining one (source 9) is OIf4, based on
the line ratio (Blum et al., 2001) and our inspection of the spectra presented in Figer et al.
(2002); Najarro et al. (2004). Without an exception, the NIR brightest WN7–8 sources haveradio continuum counterparts (Lang et al. (2001). Note that sources 9 and 11 are out of
the radio study field). This is in sharp contrast with the hard X-ray result, in which only
three of them are exceptionally bright and the remaining eight are fainter by more than twoorders of magnitude. The NIR and radio emission may have the same origin, but the result
strongly indicates that an additional factor (or several additional factors) is required for the
hard X-ray production.
One plausible idea is the binarity (Law & Yusef-Zadeh, 2004). Among WR stars, thosein a binary system with another massive star tend to have higher X-ray luminosity and harder

130

CHAPTER 5. DISCUSSION

X-ray spectra (e.g., Pollock 1987). There has been no hint to suggest the binary nature of
the X-ray sources in the Arches cluster, except for a possible K-band brightness change ofA1N between Cotera et al. (1996) and Figer et al. (2002) after correcting a systematic offset
between the two measurements. Long-term and homogeneous datasets in photometry or
spectroscopy are necessary to test this idea.
It is interesting to note the extreme brightness of the brightest source A1S. Although
the uncertainty is large, the estimated luminosity exceeds ∼ 1034.2 erg s−1 in the hard (2–8 keV) band alone. The luminosity in the soft (0.5–2.0 keV) band, which we can not quote
due to the lack of photons in this range, would increase the value by a several factor. In
fact, the luminosity of A1N is comparable to those of the brightest known stellar ( = nondegenerated) X-ray sources in our Galaxy, including η Carinae (a luminous blue variablewith ∼ 1034.6 erg s−1 in 0.5–10 keV; Corcoran et al. 2001), V444 Cygni (WN5+O6 with∼ 1033.4 erg s−1 in 0.2–10 keV; Maeda et al. 1999), WR20a (WN6+WN6 with ∼ 1033.8 erg s−1
in 0.9–8.0 keV; Rauw et al. 2005), WR25 (WN6+O4f with ∼ 1034.5 erg s−1 in 0.5–8 keV; Evanset al. 2004), HD50896 (WN5 with ∼ 1032.9 erg s−1 in 0.5–10 keV; Skinner et al. 1998), WR147(WN8+O8–9 with ∼ 1032.6 erg s−1 in 0.5–10 keV; Skinner et al. 1999), WR140 (WC7+O4–5
with ∼ 5 × 1033 erg s−1 in 1–10 keV; Koyama et al. 1994), and γ 2 Velorum (WC8+O7.5 with∼ 1032.6 erg s−1 in 0.5–10 keV; Stevens et al. 1996). Stars with higher luminosities are only
seen in the R136 in 30 Doradus represented by CX5 (WN4.5 with ∼ 1035.1 erg s−1 in 2–8 keV;
Portegies Zwart et al. 2002b; Townsley et al. 2006). The combination of the extreme massloss rate (∼ 2.2×10−4 ; Lang et al. 2005), which is also one of the Galactic largest, and thepossible binarity makes this source one of the brightest Galactic stellar source in the hard
X-ray band.

5.4.2

Origin of Three Spectral Components

We identified three spectral components (thermal, power-law, and two Gaussian lines)
in the XIS spectrum. Wang et al. (2006) conducted spatially-resolved Chandra spectroscopyof the Arches cluster in the 2–8 keV band and revealed a more detailed spatial structure
upon the results by Yusef-Zadeh et al. (2002); Law & Yusef-Zadeh (2004). The Arches Xrays consist of three spatial components: (1) point-like sources with a thermal spectrum, (2)diffuse emission with a thermal spectrum in a ∼ 3000 scale, and (3) diffuse emission with a∼ 6.4 keV line over a power-law continuum in a more extended scale (∼ 6000 ×9000 by YusefZadeh et al. 2002). The former two thermal components are concentrated at the cluster
center, in which three brightest point sources dominates the total emission. The last nonthermal component extends toward the southeast direction of the cluster. We consider thatthe thermal component in the XIS spectrum is from the ensemble of point sources plus thethermal diffuse emission at the cluster center, while the ∼ 6.4 keV and the accompanying

5.4. ARCHES CLUSTER131∼ 7.1 keV lines as well as the power-law continuum are from the larger diffuse emission.
In order to compare directly the best-fit spectral model parameters between the ACIS
and XIS spectra, we reduced the Chandra ACIS data and obtained the best-fit parameters
for the composite of the three brightest point sources and the extended emission. The
parameters from XIS and ACIS spectral analysis agree with each other with overlapping
90% confidence ranges. For the thermal component, the XIS spectrum has a temperature of
1.9 (1.5–2.7) keV and the APEC normalization of 3.2 (1.2–7.8) × 10−3 s−1 cm−2 keV−1 , while
the ACIS spectrum has 2.0 (1.9–4.4) keV and 2.5 (1.8–3.0) × 10−3 s−1 cm−2 keV−1 . For the
power-law component, the photon index in the XIS spectrum is 0.6 (0.3–1.3), while that in
the ACIS diffuse spectrum is 1.2 (0.80–1.8). For the Gaussian lines, the photon flux of the∼ 6.4 keV line is 2.4 (2.1–2.7) ×10−5 s−1 cm−2 in XIS and 2.3 (2.0–2.9) ×10−5 s−1 cm−2 in
ACIS. The equivalent width of the line against the power-law component is ∼ 1.24 keV in
XIS and ∼ 1.25 keV in ACIS, indicating that the power-law flux as well as the line flux is
consistent between the two.

5.4.3

Diffuse Medium

The high S/N spectra obtained with the XIS give a stringent constraint on the center
energy of the ∼ 6.4 keV line, which is an increasing function of the ionization stage ofiron. The fluorescent line energy is also affected by whether and how the atom is bound inmolecules, but the resultant shifts are negligible of ∼1 eV (Paerels, 1998). The best-fit value
is consistent with the Kα line from neutral iron (6.40 keV). This is in agreement with thegeneral understanding that most of iron in the ISM is in the form of dust (Sofia et al., 1994).
However, slightly low ionization stages are also allowed up to about fourteenth (Fe XV; House
1969), including the systematic uncertainty of the XIS energy gain of ∼ 6 eV.

5.4.4

Cause of Line and Power-law Emission

From the XIS spectrum (figure 4.40), we found that the power-law component extends
up to ∼ 12 keV and dominates the spectrum above ∼ 8 keV. From the band-limited XIS
images (figure 4.38), we noticed a similarity in the spatial distribution between the 7.5–
10.0 keV and the 6.4 keV emission. This indicates that the power-law emission is related to
the 6.4 keV line in the underlying physical process. Therefore, the power-law component andthe Fe I lines need to be explained simultaneously. The equivalent width of the lines and thenormalization of the power-law comprise two major observational tests to discriminate ideason the origin of the emission. We examine two mechanisms (photoelectric ionization andelectron impact ionization) and derive the conditions for the primary ionizing beam. In the

132

CHAPTER 5. DISCUSSION

former, the Thomson scattering continuum and the fluorescence are respectively responsible
for the power-law and the line emission. In the latter, non-thermal electron bremsstrahlung
and the K shell vacancy filling after the electron ionization are considered.
Photoelectric Ionization
We first note that the lack of a prominent absorption edge feature at ∼ 7.11 keV alone
does not constitute evidence against photoionization (Revnivtsev et al., 2004). This is expected when the reflecting matter is optically-thin to the primary radiation. In fact, theoptically-thin approximation is justified by the Chandra result on the Arches cluster, wherein
the power-law emission suffers an extinction accountable only by the ISM extinction to the
Galactic center region (∼ 6×1022 cm−2 ; Wang et al. 2006).
We assume that the observed continuum emission does not include the direct X-rays.
<180 eV
If it does, the equivalent width of the iron Kα fluorescent line is expected to be ∼(Reynolds & Nowak, 2003; Tsujimoto et al., 2005) for the solar abundance, which contradicts
the observed value. The equivalent width of the fluorescent line against the Thomson scattering continuum is expressed as a function of the ratio between the photoelectric and Thomson
scattering cross sections, the ratio of the target (electron and iron atom) densities, and thegeometry; the Thomson scattering is angle-dependent while the fluorescence is sphericallysymmetric (Liedahl, 1998). Assuming that the incident X-ray spectrum is I(E), the electron and iron densities of the reflecting medium are uniform (ne and nFe , respectively), thephotoelectric absorption cross section by iron is σPFe (E), the differential Thomson scatteringcross section is (dσT /dΩ)(θ) where θ is the angle between the incident and scattered X-rays,
the Kα fluorescence yield is YKα , and the K edge energy is χ, the equivalent width of theKα line is given by

EWKα

!−1

dσT
nFe
= YKα
4π
(θ)
ne
dΩ
!
R∞
Fe
χ dEI(E)σP (E)×.
I(E = 6.4 keV)




(5.1)

By substituting YKα ∼ 0.34 (Kortright, 2001), nFe /ne ∼ nFe /nH ∼ 3 × 10−5 (Däppen, 2000)
where nH is the hydrogen density, dσT /dΩ(θ) ∼ 4.0 × 10−26 (1 + cos2 θ) cm2 , I(E) ∝ E −Γ , andσPFe (E) = 2 × 10−20 (E/7.1 keV)−3 cm2 (Gullikson, 2001), we obtain
1
6.4 Γ
1
EWKα ∼ 3
keV.
(5.2)
Γ+2
7.1
1 + cos2 θWe can determine both EWKα and Γ observationally, resulting in θ ≈ 50 degrees using
the best-fit values. In practice, though, higher accuracy in the parameter determination is




 



5.4. ARCHES CLUSTER

133

required for a meaningful geometrical constraint. The observed EW Kα is easily explained by
the photoionization interpretation. For the amplitude of continuum emission, the incident
flux of I(E) ∼ 3 × 107 (ne /102 cm−3 )−1 s−1 cm−2 keV−1 is required at 10 keV. Here, thediffuse medium is assumed to have a spherical shape of a 3 pc diameter. The estimate of
the electron density is from Rodrı́guez-Fernández et al. (2001).
To summarize, the photoionization interpretation requires the following conditions: (1)the primary source is an external source, (2) the incident spectrum has a power of Γ ∼ 1 and(3) a photon flux of ∼ 107 s−1 cm−2 keV−1 at 10 keV, (4) the direct emission is optically thin
to the reflecting medium (i.e., nH < 2 × 105 cm−2 ), and (5) the reflecting geometry should
satisfy the constraint on θ.
No source is found for the primary source in the vicinity at present. The point sourcesin the Arches cluster can be excluded. This is simply because the continuum emission by
the Thomson scattering cannot exceed the illuminating thermal emission above the iron Kedge energy, which is contrary to the XIS spectrum showing stronger non-thermal emission
than thermal emission at the band (Fig. 4.40). Wang et al. (2006) also discussed that Arches
point sources are too dim for the fluorescent lines.
Electron Ionization
Another favored interpretation for the cause of the ∼ 6.4 keV line is the vacancy filling
after the iron K shell ionization by low energy electrons with an energy of 10–100 keV. Thiswas proposed to explain the Galactic ridge diffuse X-ray emission (Valinia et al., 2000), in
which they claimed that the low energy electrons with a density of ∼ 0.2 eV cm−3 and thepower-law spectrum of an index of ∼ 0.3 contribute significantly to both the observed line
and continuum emission. Yusef-Zadeh et al. (2002b) and Wang et al. (2006) employed thismodel to account for the ∼ 6.4 keV line observed in G0.13–0.13 and the Arches cluster in
the Galactic center region, respectively.
The calculation of expected line and continuum intensity of this process requires numerical treatments because of the complex dependence of the cross sections on the electron
and X-ray energy. A detailed computation is given in Tatischeff (2002). Two important
things have to be considered. First, the electron beam is stopped at the surface of thehydrogen column. The stopping range for 10 and 100 keV electrons is ∼ 7×1019 mH and∼ 4×1021 mH g cm−2 where mH is the hydrogen mass (Tatischeff, 2002). Second, the energy
conversion rate from the electron beam to the X-ray bremsstrahlung is quite small with anorder of ∼ 10−5 .
With these kept in mind, we can compare the expected and observed values. The
equivalent width of the iron Kα line against the bremsstrahlung continuum is expected to be

134

CHAPTER 5. DISCUSSION∼ 290 eV (Tatischeff, 2002), which is smaller than the observed value. An iron abundanceof 4–5 times larger than the solar value needs to be introduced to reconcile the discrepancy,
though no such indication is present for the Arches cluster. From the spectral fit of the XIS
spectrum, we found that a solar abundance for iron accounts for the observed Fe XXV Kα line
intensity over the continuum. Wang et al. (2006) suggested that the iron abundance is ∼ 2from the X-ray spectral fits of point sources in the cluster. Najarro et al. (2004) conductednear-infrared spectroscopy of the photospheric emission of Wolf-Rayet stars in the Arches
cluster and found that the abundance is consistent with solar.
For the level of continuum emission, we can consult Figure 7 in Tatischeff (2002), which
illustrates the X-ray production rate by the low energy electron impact. The injected electrons are assumed to have a total energy rate of 1 erg s−1 in an energy range of 10–100 keV
and the power-law number density with an index of 2. At 10 keV, a continuum photon production rate of ∼ 8 × 10 s−1 keV−1 is expected. We observed ∼ 1.1 × 10−5 s−1 cm−2 keV−1 at
10 keV in the XIS spectrum, which can be converted to ∼ 9.5 × 1040 s−1 keV−1 at a distanceof 8 kpc. Therefore, an electron injection rate of ∼ 1×1039 erg s−1 is required. Assumingthat the diffuse medium is “optically-thick” to the electron beam, we can derive the required
flux of the injected electrons to be ∼ 2 × 10 erg s−1 cm−2 regardless of the target density.
Integrating over the entire mass of the medium (Yusef-Zadeh et al., 2002b; Wang et al.,2006) may lead to a significant overestimation.
To summarize, the electron ionization interpretation requires the following conditions:(1) the electron beam has an incident flux of ∼ 2 × 101 erg s−1 cm−2 and (2) an index of ∼ 2
for the injected power-law number density, and (3) 4–5 times larger abundance than solar
for iron in the diffuse medium.
We do not have a list of electron accelerators with an estimated energy injection rate in
the Galactic center region and further studies must be undertaken to elucidate the primarysource. However, we can at least claim that the accelerated electrons by wind-wind collisions
in the Arches cluster (Wang et al., 2006) are unlikely the case. This is simply because thetotal kinetic energy rate via stellar winds is smaller by an order than the required electronenergy injection rate. Here the total energy rate is calculated from an assumed wind velocityof ∼ 103 km s−1 and a mass loss rate of ∼ 10−4 M yr−1 estimated by the free-free emission
intensity at the centimeter continuum (Lang et al., 2005).
Protons accelerated via wind-wind collisions to have an MeV energy may cause boththe bremsstrahlung and the iron line emission in a similar manner as electrons with several
tens of keV. However, because the bremsstrahlung conversion rate for protons is as inefficientas that of electrons (Uchiyama et al., 2002), the protons ionization will encounter the sameenergy budget deficit.
Finally, heavy ions with an MeV amu−1 energy are also capable of producing the con-

5.4. ARCHES CLUSTER

135

tinuum emission by inverse bremsstrahlung and the iron line emission by K shell ionization.
In this case, however, the resultant lines are broader and bluer than the case of proton orelectron ionization by ∼ 50 eV (e.g., Burch et al. 1971), which disagrees with the iron Kαline center determined from the XIS spectrum. Therefore, the ionization by heavy ions canbe excluded regardless of whether they are accelerated by wind-wind collisions in the Arches
cluster or somewhere else in the Galactic center region.

Chapter 6
Conclusions
In this chapter, we summarize our findings in each studied field and draw general
conclusions from cross-regional points of view.

Contents
6.1

6.2

6.3

A New Wolf-Rayet Star in the Galactic Center . . . . . . . . . 135
6.1.1

Summary of the Results . . . . . . . . . . . . . . . . . . . . . . . . 135

6.1.2

Conclusion — Contribution to the Galactic Center Diffuse Emission 136

Discrete Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . 137
6.2.1

Summary of the Results . . . . . . . . . . . . . . . . . . . . . . . . 137

6.2.2

Conclusion — Mechanism for the Hard X-rays from Massive Stars 138

Diffuse Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 140
6.3.1

Summary of the Results . . . . . . . . . . . . . . . . . . . . . . . . 140

6.3.2

Conclusion — Various origins of Diffuse Emission in Star-forming
Regions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 141

6.1

A New Wolf-Rayet Star in the Galactic Center

6.1.1

Summary of the Results

From the Galactic center survey, we found that a point-like and bright X-ray source withan intense Fe XXV Kα emission line is a new WR star. The spectrum was well reproduced
137

138

CHAPTER 6. CONCLUSIONS

by an optically-thin thermal plasma model with a temperature of ∼ 4 keV attenuated by anabsorption column (NH ) of ∼ 2×1023 cm−2 and a luminosity of ∼ 3×1034 ×(d/8 kpc)2 erg s−1
in the 2.0–8.0 keV band. The source has a very red counterpart in the near- and midinfrared bands. The SED was fitted by a blackbody emission of ∼ 103 K and an extinctionof AV ∼ 31 mag. The high temperature and large luminosity in the X-ray bands suggest thatthe source is a wind-wind colliding Wolf-Rayet binary. The infrared SED is very similar to
those of eponymous Quintuplet cluster members, indicating that this source is a WC type.

6.1.2

Conclusion — Contribution to the Galactic Center Diffuse
Emission

The intense 6.7 keV line found in our analysis and the large predicted number of undiscovered WR stars in the Galactic center region (section 2.2) lead us to speculate how muchfraction of the Galactic center diffuse emission is accountable by the ensemble of WR stars.
The 6.7 keV line flux from CXOGC J174645.3–281546 is 2.5 × 10−13 erg s−1 cm−2 . thetotal flux of iron K-lines in the Galactic center region is ∼ 1.3×10−10 erg s−1 cm−2 (Yamauchi
et al., 1990). Since this value includes the 6.4 keV and 6.97 keV lines, it overestimates the
6.7 keV flux by a factor of 2–3. The detected point sources account for ∼ 10% of the total
flux of the 6.7 keV line (Wang et al., 2002; Muno et al., 2004a). Therefore, the 6.7 keV lineof CXOGC J174645.3–281546 alone accounts for ∼ 4% of that of the detected point sourcesand ∼ 0.4% of that of the total diffuse flux.
The estimated number of WR stars in the inner 1 degree of the Galaxy is ∼ 70, and
the WR binary frequency is ∼ 39% in the solar neighborhood (van der Hucht, 2001). We
hence expect ∼ 27 WR binaries in the Galactic center region. If 27 WR binaries have thesame iron K line flux with CXOGC J174645.3–281546, ∼ 11% of the iron K line flux inthe Galactic center diffuse X-ray emission is explained by WR binaries. It has long beenargued that CVs are the main contributor to the Galactic center diffuse X-ray emission(e.g., Wang et al. 2002). However, we propose that binaries of two massive stars such as
CXOGC J174645.3–281546 would comprise a substantial fraction of 6.7 keV emission line in
the Galactic center.

6.2. DISCRETE SOURCES

6.2

Discrete Sources

6.2.1

Summary of the Results

139

M 17
Although Suzaku did not resolve individual stars of the core of M 17, the major fractionof the X-ray emission from the OB association is due to the central O4–O5 binaries (CEN
1a and CEN 1b). The spectrum is composed of two optically-thin thermal plasma models
of kB Tlow ∼ 0.6 keV and kB Thigh ∼ 4 keV. Suzaku J182132–1613.4 is a newly found source.
The infrared brightness is consistent with a B0 V star. The X-ray spectra was well fittedwith a high temperature (kB T ∼ 10 keV) plasma. OI 352 is an isolated O8 star (Ogura& Ishida, 1976). The X-ray spectrum was well reproduced by two thermal components(kB Tlow ∼ 0.6 keV and kB Thigh ∼ 4 keV).
M17 North is comprised of ∼ 10 point-like sources. The spectrum was well fitted by heavily absorbed (NH ∼ 2 × 1022 cm−2 ) optically-thin thermal plasma model with a temperatureof ∼ 3 keV. These spectral values are consistent with embedded protostars. Suzaku J182132–
1610.4 shows flare-like time variability, but the IR brightness is inconsistent with low-mass
young stellar object located at the distance of M 17. The source is probably a nearby object.
The X-ray spectrum of Suzaku J182140–1604.1 is exceptionally hard. The prominent iron K
line with an EW of ∼ 2 keV suggests a thermal nature of the emission. When fitted by anabsorbed optically-thin thermal plasma model, the absorption column is ∼ 1024 cm−2 andthe temperature is ∼ 13 keV. A possible origin is an embedded protostar.
The Carina Nebula
Hard X-rays with the 6.7 keV line were also found in some massive stars in the
Carina region. The X-ray emission of WR 25 (WN6+O4f) in the Carina Nebula is bright(LX ∼ 1033.3 erg s−1 in the 1.0–8.0 keV band), hard (kB Thigh ∼ 2.3 keV), and shows a phaselocked flux variation of a factor of ∼ 2.
The spectrum of HD 93205 (O3 V+O8 V) was well fitted by a single component witha temperature of ∼ 0.25 keV. The flux showed no significant variability among the four
observations.
HD 93250 (O3 V, no known companion) has two thermal components of different temperatures (kB T low ∼ 0.6 keV and kB T high ∼ 3 keV) with a luminosity of ∼ 1032.7 erg s−1 inthe 1.0–8.0 keV band.

140

CHAPTER 6. CONCLUSIONSThe Arches cluster
The Arches cluster is one of the most densely packed young massive star clusters in our
Galaxy. Among the 17 WR stars in the cluster, only three sources (A1N, A1S, and A2) weredetected from the Chandra observations. The other sources, which are almost identical ininfrared and radio properties, are dim in hard X-rays by a factor of more than 100. The threesources have hot thermal components with a temperatures of ∼ 2 keV. The luminosities are1033.4 –1034.2 erg s−1 in the 2.0–8.0 keV band.

6.2.2

Conclusion — Mechanism for the Hard X-rays from Massive
Stars

All discrete massive sources except HD 93205 presented in this thesis have a hard thermal
component (kB T > 2 keV) with a prominent Fe XXV Kα emission line at ∼ 6.7 keV.
In WR 25, we saw that the binary orbital separation is a key parameter in hard Xray production. The flux decreases as the separation increases. The fact that the X-ray
flux modulates with the orbital motion clearly indicates that the X-ray production is aconsequence of the binarity.
HD 93205 and the OB association of M 17 (mostly CEN 1a and CEN 1b) are also confirmed to be binaries of two massive stars by optical or NIR observations (Gamen et al.,2006; Conti & Walborn, 1976; Hoffmeister et al., 2008). HD 93205 is the only source that isa binary of two massive stars but lacks hard X-ray emission. The peculiarity of HD 93205
is the small separation of the two massive stars. These facts suggest that the stellar winds
collide before the velocity exceed the sound velocity in HD 93205, and the distance between
two massive stars should be larger than several times the stellar radius to produce hard
X-rays.
The hard X-ray emission from massive stars presented in this thesis and other work is
in sharp contrast with the conventional soft emission from single O stars. We deductively
speculate that other massive stars with a hot component (A1S, A1N, A2, OI 352, HD 93250,
CEN 1a, CEN 1b, and Suzaku J182132–1613.4) are also colliding-wind binaries. Indeed,
CEN 1a and CEN 1b are found to have a binary nature. However magnetically-confined
wind shock model cannot be excluded (section 2.4), hence monitoring spectroscopy of theseobjects in the optical and infrared bands is highly encouraged to confirm this hypothesis.
We also consider that not all the X-ray emission is a consequence of colliding winds.
In WR 25, we expect a factor of ∼ 3.5 in the X-ray flux variation assuming that all X-rays
are phase-locked with the orbital modulation of e = 0.56, which is actually ∼ 2 (table 4.6

6.2. DISCRETE SOURCES

141

and figure 5.6). The discrepancy is also seen in other colliding-wind WR star, WR 140, a
smaller amplitude (figure 2.11) than that expected by the eccentricity determined by opticalobservations. We consider that this discrepancy stems from constant emission, which isprobably intrinsic emission of the individual component.
The Arches cluster provides a uniform sample of 17 WR stars: They are located at
almost identical distance from us and have similar ages. Among the 17 sources, only threesources are exceptionally bright in X-rays, although the NIR and radio properties are notspecial among the 17 sources. This indicates that ∼ 20% of WR stars, in general, arecolliding-wind binaries and emit hard X-rays.

142

CHAPTER 6. CONCLUSIONS

6.3

Diffuse Sources

6.3.1

Summary of the Results

M 17
We clearly resolved emission lines from highly ionized O, Ne, Mg, and Fe for the firsttime. The line intensity ratios well constrain the plasma temperature to be ∼ 0.25 keV. The
absorption column of ∼ 4 × 1021 cm−2 is typical to M 17. The temperature and chemicalcomposition show no spatial variation across the studied field of ∼ 10 pc.
The Carina Nebula

Metal abundance ([Z/Fe]/[Z /Fe ])In contrast to M 17, the Carina diffuse emission shows a spatial variation. The emission
in the west side of η Carinae shows a low temperature of ∼ 0.2 keV, while that in thesouth side shows a higher temperature of ∼ 0.5 keV. The 0.2 keV emission is absorbedwith a typical value to this region (NH ∼ 4 × 1021 cm−2 ), while the 0.5 keV emission suffersa small absorption (NH < 6 × 1020 cm−2 ). Figure 6.1 shows the metal abundances of the
diffuse emission in M 17, “diffuse 1 and 2” (south to η Carinae), and “diffuse 3” (west to η
Carinae).


3



2

1

0

5

10

15
20
Atomic number

25

30

Figure 6.1: Metal abundance relative to Fe in the unit of solar abundances. Data for ”Diffuse 1 and
2” and ”Diffuse 3” in the Carina Nebula and M 17 are shown in red, green, and black, respectively.

143

6.3. DIFFUSE SOURCES
The Arches Cluster

The Arches cluster exhibits a hard diffuse emission with a power-law index of ∼ 0.7
and a Gaussian line from Fe I. The 7.5–10.0 keV image is similar to the 6.4 keV image,including the local excess at the Arches cluster, which suggests that the hard continuum andthe Gaussian line are physically related.

6.3.2

Conclusion — Various origins of Diffuse Emission in Starforming Regions

Wind-blown bubble
We consider that the diffuse X-ray emission in M 17 and the west to η Carinae areclassical examples of stellar wind bubble. The spectroscopic characterization is presented forthe first time for X-ray bubbles in star-forming regions, which has a temperature of ∼ 0.2 keV< 1 solar.
and the O/Fe ratio of ∼Table 6.1: Plasma parameters of M 17 and western Carina

Par.
kB T
ne
l
LX
SX
V
M
Ṁ
Eth
tsys
trad
tcon

Unit
M 17
(keV)
0.25
−3
(cm )
1
(pc)
10
33
−1
(10 erg s )
5
−13
−1
−2
−2
(10
erg s cm arcmin )
2
3
(pc )
100
(M )
2
−5
−1
(10 M yr )348
(10 erg)
3(year)
106(year)
107(year)
30

Carina
0.18
5
3
16
2
8
1.5
4
1
105
107
100

Table 6.1 shows the plasma temperature (kB T ), the electron density (ne ), the scale
size (l), the X-ray luminosity (LX ), the X-ray surface brightness (SX ), the plasma volume(V ), the plasma mass (M ), the thermal energy (Eth = 3ne kB T V ), the age of the system(tsys ), the radiative cooling time scale (trad = Eth /LX ), and the conduction cooling time scale

144

CHAPTER 6. CONCLUSIONS(tcon = Eth /κ, where κ = 1.5×10−17 ×(dT /dx)×l2 erg s−1 is the thermal conductivity; Spitzer
1962) for the diffuse emission of M 17 and the west to η Carinae.
In M 17, if the plasma expands with the sound speed of ∼ 250 km s−1 and the age of thesystem of 106 years, the size would become ∼ 250 pc, which is larger than the observed value
by more than ten-fold. This means that the plasma is suppressed or cooled by some sortof mechanism. We see a clear edge at the west end of the diffuse emission with a thicknessof ∼ 1 pc in M 17 (figure 4.24). This sharp darkening can be attributed to the thermal
conduction between the hot plasma and cool ISM accumulated by the winds from OB stars
at the core of M 17. Because the thermal conduction time scale is very short (∼ 30 year,
table 6.1), the cooling scale length is determined by sound speed of the conduction front.
Multiplying the age of system (∼ 106 year), we expect the thickness to be ∼ 3 pc, which isroughly consistent with the observed value. The uniform temperature is also justified by thefact that the thermal conduction works only near the contact surface.
It is unclear which star made the diffuse emission in the west to η Carinae. We here
assume that WR 25, the closest object, made the stellar wind bubble. Since the mass-loss rateof WR 25 is comparable to that of OB association of M 17, we speculate that the differenceof the plasma size between M 17 and west to η Carinae is due to the system age.
Supernova remnants
The other two diffuse sources (Diffuse 1 and 2) show very small absorption column
towards the 0.5 keV emission. We argue that the emission is not physically related to the
Carina Nebula but originates from a foreground object. A possible origin is a type Ia SNR,
which is favored by the high O/Fe value of ∼ 5.
Reflection by cold matter
The diffuse emission in the Arches cluster has a hard spectrum, and is obviously different
in origin from diffuse emission in M17 and the Carina Nebula. We examined two ideas
to account for both the power-law and Gaussian components: (1) X-ray photoionizationthat produces fluorescence lines and the Thomson scattering continuum and (2) electron
impact ionization with subsequent K-shell vacancy filling and bremsstrahlung continuum.
We concluded that whichever scenario is adopted, the photon or particle flux from theArches cluster is too low to yield observed flux.

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Acknowledgments
My deepest appreciation goes to Katsuji Koyama who introduced me to the X-ray
astronomy and has been patiently encouraging me throughout my graduate course. I would
like to express my gratitude to Masahiro Tsujimoto for his continuous mentoring on this work.
I am grateful to Takeshi Go Tsuru and Hironori Matsumoto whose meticulous comments
were of enormous help.
I would also like to thank to Hideki Uchiyama and Midori Ozawa who informed me of
the latest information on the XIS calibration.
For the NIR study, I am indebted to Tetsuya Nagata and Shogo Nishiyama who offeredme the SIRIUS data and gave me constructive comments. I appreciate Itsuki Sakon for his
invaluable comments on the analysis of the MIR data.
My gratitude is also for Makoto Sawada and Masayoshi Nobukawa for making figures inthis work and checking the manuscript. Aya Bamba, Hiroya Yamaguchi, Yoshitomo Maeda,
Yohko Tsuboi, Kenji Hamaguchi, Yuichiro Ezoe, Shunji Kitamoto, and Hiroshi Murakamigave me insightful comments and suggestions.
I acknowledge financial support from the Japan Society for the Promotion of Science.
The work is financially supported by the grants-in-aid for a 21st century center of excellence
program ”Center for Diversity and Universality in Physics” and the Global COE Program”The Next Generation of Physics, Spun from Universality and Emergence” by the Ministryof Education, Culture, Sports, Science and Technology of Japan.
This research has made use of data obtained from the Data ARchive and Transmission
System at ISAS/JAXA, the High Energy Astrophysics Science Archive Research Center at
the NASA/Goddard Space Flight Center, and the Two Micron All Sky Survey, which is a
joint project of the University of Massachusetts and the Infrared Processing and Analysis
Center/California Institute of Technology, funded by the National Aeronautics and Space
Administration and the National Science Foundation. MSX was funded by the Ballistic
Missile Defense Organization with additional support from NASA Office of Space Science.
This research has also made use of the NASA/IPAC Infrared Science Archive, which is
161

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operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract
with the National Aeronautics and Space Administration.