Discrete X-ray Sourcesin the Small Magellanic Cloud— Tracing the Star Forming Activity —Jun Yokogawa
Department of Physics, Graduate School of Science,
Kyoto University, Sakyo-ku, Kyoto 606-8502, Japan

Submitted to the Department of Physics, Kyoto University
in partial fulfillment of the requirements
for the degree of Doctor of Science

January 7, 2002

This thesis will be available at
http://www-cr.scphys.kyoto-u.ac.jp/member/jun/job/phd/.

Abstract
The Advanced Satellite for Cosmology and Astrophysics (ASCA) has made 22 observations on the Small Magellanic Cloud (SMC) by the end of the mission. We performed
systematic analyses on all of these observation data and made comprehensive catalogues.
We detected 106 discrete sources with a criterion of S/N > 5. The source positions were
determined with an ∼ 4000 error radius (90% confidence) for sources detected in the central
200 radius of the GIS. We detected coherent pulsations from 17 sources. Among them, eight
were newly discovered during this study. We classified most of these pulsars as X-ray binary pulsars (XBPs) based on their properties such as flux variability. X-ray emission was
detected from eight supernova remnants (SNRs). Among them, five SNRs showed emission
lines in their spectra, hence we regarded the five as thermal. We find that XBPs and thermal SNRs in the SMC can be clearly separated by their spectral hardness ratio. Applying
this empirical law to faint (thus unclassified) sources, we find 20 XBP candidates and four
thermal SNR candidates. We find the spatial distributions of the XBPs and thermal SNRs
(including candidates) resemble the distribution of young stars with ages ∼ 2 × 107 yr. We
compare the source populations in the SMC and in our Galaxy and discuss the star forming
activity of the SMC.

Contents

1 Introduction

12 Review — the Small Magellanic Cloud

3

2.1

The Magellanic System . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

5

2.2Interaction between the Magellanic Clouds and our Galaxy . . . . . . . . . .

6

2.3

Structure of the Small Magellanic Cloud . . . . . . . . . . . . . . . . . . . .

7

2.4

Previous High Energy Observations . . . . . . . . . . . . . . . . . . . . . . .

10

2.4.1

X-ray surveys . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

10

2.4.2

GeV and TeV γ-ray observations . . . . . . . . . . . . . . . . . . . .

123 Instruments

15

3.1

Overview of the ASCA Satellite . . . . . . . . . . . . . . . . . . . . . . . . .

15

3.2

XRT: X-ray Telescope . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

15

3.3

GIS: Gas Imaging Spectrometer . . . . . . . . . . . . . . . . . . . . . . . . .

19

3.4

SIS: Solidstate Imaging Spectrometer . . . . . . . . . . . . . . . . . . . . . .

22

4 Observations and Data Reduction

27

4.1

Observation Fields . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

27

4.2

Data Reduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

27

i

ii

CONTENTS

5 Systematic Analyses on Discrete Sources
5.1

31

X-Ray Images . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

31

5.1.1

Images in the sky/detector coordinate systems . . . . . . . . . . . . .

31

5.1.2

Source extraction . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

32

5.1.3

Construction of mosaic images . . . . . . . . . . . . . . . . . . . . . .

32

Position Determination . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

33

5.2.1

Absolute accuracy of the position . . . . . . . . . . . . . . . . . . . .

36

5.2.2

Identification of sources detected in multiple ASCA observations . . .

39

5.3

Timing Analyses . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

42

5.4

Spectral Analyses . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

44

5.4.1

Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

44

5.4.2

Hardness ratio . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

44

5.4.3

SNRs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

45

5.4.4

X-ray pulsars and HMXBs . . . . . . . . . . . . . . . . . . . . . . . .

46

5.4.5

Remaining sources . . . . . . . . . . . . . . . . . . . . . . . . . . . .

46

Catalogue of Discrete X-Ray Sources in the SMC Region . . . . . . . . . . .

47

5.5.1

Catalogue – I. Source identification . . . . . . . . . . . . . . . . . . .

47

5.5.2

Catalogue – II. Names and comments . . . . . . . . . . . . . . . . . .

51

5.5.3

Catalogue – III. Spectral parameters . . . . . . . . . . . . . . . . . .

55

5.2

5.5

6 Comments on Specific Sources and Pulsar Statistics
6.1

59

X-Ray Pulsars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

59

6.1.1

Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

59

6.1.2

No. 17 — AX J0043−737 . . . . . . . . . . . . . . . . . . . . . . . .

60

6.1.3

No. 24 — AX J0049−729 . . . . . . . . . . . . . . . . . . . . . . . .

63

CONTENTS

iii

6.1.4

No. 26 — AX J0049−732 . . . . . . . . . . . . . . . . . . . . . . . .

65

6.1.5

No. 30 — AX J0049.5−7323 . . . . . . . . . . . . . . . . . . . . . . .

66

6.1.6

No. 32 — AX J0051−733 . . . . . . . . . . . . . . . . . . . . . . . .

68

6.1.7

No. 37 — AX J0051−722 . . . . . . . . . . . . . . . . . . . . . . . .

71

6.1.8

No. 40 — AX J0051.6−7311 . . . . . . . . . . . . . . . . . . . . . . .

72

6.1.9

No. 43 — RX J0052.1−7319? . . . . . . . . . . . . . . . . . . . . . .

73

6.1.10 No. 44 — XTE J0054−720 . . . . . . . . . . . . . . . . . . . . . . . .

75

6.1.11 No. 47 — 1WGA J0053.8−7226 = XTE J0053−724 . . . . . . . . . .

76

6.1.12 No. 49 — SMC X-2 . . . . . . . . . . . . . . . . . . . . . . . . . . . .

76

6.1.13 No. 51 — XTE J0055−724 = 1SAX J0054.9−7226 . . . . . . . . . .

79

6.1.14 No. 56 — AX J0057.4−7325 . . . . . . . . . . . . . . . . . . . . . . .

80

6.1.15 No. 61 — AX J0058−7203 . . . . . . . . . . . . . . . . . . . . . . . .

82

6.1.16 No. 67 — RX J0059.2−7138 . . . . . . . . . . . . . . . . . . . . . . .

84

6.1.17 No. 74 — RX J0101.3−7211 . . . . . . . . . . . . . . . . . . . . . . .

86

6.1.18 No. 78 — 1SAX J0103.2−7209. . . . . . . . . . . . . . . . . . . . .

87

6.1.19 No. 83 — AX J0105−722 . . . . . . . . . . . . . . . . . . . . . . . .

88

6.1.20 No. 90 — XTE J0111.2−7317 . . . . . . . . . . . . . . . . . . . . . .

90

6.1.21 No. 94 — SMC X-1 . . . . . . . . . . . . . . . . . . . . . . . . . . . .

92

6.1.22 Other pulsars not detected with ASCA — 2E 0050.1−7247, XTE J0052−723,
RX J0117.6−7330, and XTE J0050−732#1 and #2 . . . . . . . . . . 94
6.2

Period Distribution of XBPs . . . . . . . . . . . . . . . . . . . . . . . . . . .

96

6.3

SNRs – I: from SIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 101
6.3.1

No. 21 — 0045−734 (N19) . . . . . . . . . . . . . . . . . . . . . . . . 101

6.3.2

No. 25 — 0047−735 . . . . . . . . . . . . . . . . . . . . . . . . . . . 104

iv

CONTENTS

6.4

6.5

6.3.3

No. 66 — 0057−7226 (N66) . . . . . . . . . . . . . . . . . . . . . . . 104

6.3.4

No. 81 — 0102−723 . . . . . . . . . . . . . . . . . . . . . . . . . . . 106

6.3.5

No. 82 — 0103−726 . . . . . . . . . . . . . . . . . . . . . . . . . . . 106

SNRs – II: from GIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 109
6.4.1

No. 23 — 0046−735 . . . . . . . . . . . . . . . . . . . . . . . . . . . 109

6.4.2

No. 36 — 0049−736 . . . . . . . . . . . . . . . . . . . . . . . . . . . 109

6.4.3

No. 64 — 0056−725 . . . . . . . . . . . . . . . . . . . . . . . . . . . 110

Other Interesting Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 112
6.5.1

No. 20 — IKT1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 112

6.5.2

No. 22 — AX J0048.2−7309 . . . . . . . . . . . . . . . . . . . . . . . 113

6.5.3

No. 105 — AX J0128.4−7329 . . . . . . . . . . . . . . . . . . . . . . 114

7 Source Classification

117

7.1

Criteria for Source Classification . . . . . . . . . . . . . . . . . . . . . . . . . 117

7.2

Classification by Hardness Ratio . . . . . . . . . . . . . . . . . . . . . . . . . 119
7.2.1

XBPs and thermal SNRs . . . . . . . . . . . . . . . . . . . . . . . . . 119

7.2.2

Validity of the classification by hardness ratio . . . . . . . . . . . . . 121

7.2.3

Candidates for XBPs and thermal SNRs . . . . . . . . . . . . . . . . 122

8 Implication for the Star Forming Activity
8.1

8.2

125

Source Populations in the SMC and in our Galaxy . . . . . . . . . . . . . . . 125
8.1.1

Basic data . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 125

8.1.2

Comparison of the source populations . . . . . . . . . . . . . . . . . . 126

Spatial Distribution of Various Classes of Sources . . . . . . . . . . . . . . . 128

9 Conclusion

131

CONTENTS
A Two Color X-Ray Mosaic Image of the SMC

v
133

vi

CONTENTS

List of Figures
2.1

Photo of a Cerro Tololo sky . . . . . . . . . . . . . . . . . . . . . . . . . . .

3

2.2

Optical images of the Magellanic Clouds . . . . . . . . . . . . . . . . . . . .

4

2.3

H i distribution of the Magellanic System . . . . . . . . . . . . . . . . . . . .

5

2.4

Spatial distribution of stars with different ages . . . . . . . . . . . . . . . . .

8

2.5

H i distribution in the SMC . . . . . . . . . . . . . . . . . . . . . . . . . . .

9

2.6Distances to Cepheids in the SMC . . . . . . . . . . . . . . . . . . . . . . . .

9

2.7

X-ray view of the SMC in late 1970’s . . . . . . . . . . . . . . . . . . . . . .

10

2.8

X-ray view of the SMC with ROSAT PSPC . . . . . . . . . . . . . . . . . .

11

3.1

Schematic view of ASCA . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

16

3.2

Mechanism of an X-ray telescope . . . . . . . . . . . . . . . . . . . . . . . .

17

3.3

Effective area of XRT vs. X-ray energy . . . . . . . . . . . . . . . . . . . . .

17

3.4

Effective area of XRT vs. off-axis angle . . . . . . . . . . . . . . . . . . . . .

18

3.5

Point spread function of XRT . . . . . . . . . . . . . . . . . . . . . . . . . .

18

3.6

Stray light of XRT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

19

3.7

Schematic view of GIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

20

3.8

Effective area of GIS with XRT . . . . . . . . . . . . . . . . . . . . . . . . .

21

3.9

SIS chip configuration . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

23

vii

viii

LIST OF FIGURES
3.10 Effective area of SIS with XRT . . . . . . . . . . . . . . . . . . . . . . . . .

23

3.11 Grade definition for an SIS event . . . . . . . . . . . . . . . . . . . . . . . .

25

4.1

Observation fields of ASCA . . . . . . . . . . . . . . . . . . . . . . . . . . .

29

5.1

X-ray mosaic image of the SMC (0.7–2.0 keV) . . . . . . . . . . . . . . . . .

34

5.2

X-ray mosaic image of the SMC (2.0–7.0 keV) . . . . . . . . . . . . . . . . .

35

5.3

Separation angles between the selected ROSAT–ASCA counterparts . . . . .

37

5.4

Separation vs. off-axis for ASCA sources detected multiple times . . . . . . .

40

5.5

Examples of power spectrum densities. . . . . . . . . . . . . . . . . . . . .

43

5.6

Periodigram for XTE J0055−724 . . . . . . . . . . . . . . . . . . . . . . . .

43

6.1

History of pulsar discoveries in the SMC . . . . . . . . . . . . . . . . . . . .

60

6.2

Power spectrum densities of AX J0043−737 . . . . . . . . . . . . . . . . . .

62

6.3

GIS spectrum and pulse profile of AX J0043−737 . . . . . . . . . . . . . . .

63

6.4

GIS spectrum and pulse profiles of AX J0049−729 . . . . . . . . . . . . . . .

64

6.5

GIS spectrum and pulse profile of AX J0049−732 . . . . . . . . . . . . . . .

66

6.6

GIS spectrum and pulse profiles of AX J0049.5−7323 . . . . . . . . . . . . .

67

6.7

GIS spectrum and pulse profiles of AX J0051−733 . . . . . . . . . . . . . . .

69

6.8

Light curves of AX J0051−733 in obs. Q . . . . . . . . . . . . . . . . . . . .

70

6.9

GIS spectrum and pulse profiles of AX J0051−722 . . . . . . . . . . . . . . .

71

6.10 GIS spectrum and pulse profiles of AX J0051.6−7311 . . . . . . . . . . . . .

73

6.11 GIS spectrum of No. 43 (= RX J0052.1−7319?) . . . . . . . . . . . . . . . .

74

6.12 GIS spectrum and pulse profiles of XTE J0054−720 . . . . . . . . . . . . . .

75

6.13 GIS spectrum and pulse profiles of 1WGA J0053.8−7226 . . . . . . . . . . .

77

6.14 GIS spectrum and pulse profiles of SMC X-2 . . . . . . . . . . . . . . . . . .

78

LIST OF FIGURES

ix

6.15 ROSAT PSPC image around SMC X-2 . . . . . . . . . . . . . . . . . . . . .

79

6.16 GIS spectrum of XTE J0055−724 . . . . . . . . . . . . . . . . . . . . . . . .

80

6.17 GIS spectrum and pulse profiles of AX J0057.4−7325 . . . . . . . . . . . . .

81

6.18 GIS spectrum and pulse profiles of AX J0058−7203 . . . . . . . . . . . . . .

83

6.19 Power spectrum densities of AX J0058−7203 . . . . . . . . . . . . . . . . . .

84

6.20 GIS spectrum and pulse profiles of RX J0059.2−7138 . . . . . . . . . . . . .

85

6.21 GIS spectrum and pulse profiles of RX J0101.3−7211 . . . . . . . . . . . . .

86

6.22 GIS spectrum and pulse profiles of 1SAX J0103.2−7209 . . . . . . . . . . . .

88

6.23 GIS images around AX J0105−722 . . . . . . . . . . . . . . . . . . . . . . .

89

6.24 GIS spectrum and pulse profiles of AX J0105−722 . . . . . . . . . . . . . . .

90

6.25 GIS spectrum and pulse profiles of XTE J0111.2−7317 . . . . . . . . . . . .

91

6.26 GIS spectra and pulse profiles of SMC X-1 in obs. A and I . . . . . . . . . .

93

6.27 GIS spectrum and pulse profiles of SMC X-1 in obs. C . . . . . . . . . . . .

94

6.28 Period distribution of XBPs in the SMC and our Galaxy . . . . . . . . . . .

97

6.29 Period vs. minimum luminosity of XBPs in the SMC . . . . . . . . . . . . .

99

6.30 SIS spectrum and ROSAT HRI image of 0045−734 . . . . . . . . . . . . . . 102
6.31 Element mass and theoretical predictions for 0045−734 . . . . . . . . . . . . 104
6.32 SIS spectrum of 0047−735 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 105
6.33 SIS spectrum and ROSAT HRI image of 0057−7226 . . . . . . . . . . . . . . 106
6.34 SIS spectrum and ROSAT HRI image of 0103−726 . . . . . . . . . . . . . . 107
6.35 Element mass and theoretical predictions for 0103−726 . . . . . . . . . . . . 108
6.36 GIS spectrum of 0046−735 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 109
6.37 GIS spectrum of 0049−736 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110
6.38 GIS spectrum of 0056−725 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 111

x

LIST OF FIGURES
6.39 GIS spectrum of IKT1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 112
6.40 Light curves of IKT1 in obs. Q. . . . . . . . . . . . . . . . . . . . . . . . . 113

6.41 GIS spectra of AX J0048.2−7309 . . . . . . . . . . . . . . . . . . . . . . . . 114
6.42 GIS spectrum of AX J0128.4−7329 . . . . . . . . . . . . . . . . . . . . . . . 115
7.1

HR vs. Lobs for XBPs, thermal SNRs, Crab-like SNRs, and BH binaries . . . 120

7.2

HR vs. Lobs for the other sources in the SMC

7.3

HR vs. Lobs for all sources in the SMC . . . . . . . . . . . . . . . . . . . . . 123

8.1

Spatial distribution of various classes of sources . . . . . . . . . . . . . . . . 129

8.2

Spatial distribution of HMXBs and young stars . . . . . . . . . . . . . . . . 130. . . . . . . . . . . . . . . . . 121

A.1 Two color X-ray mosaic image of the SMC . . . . . . . . . . . . . . . . . . . 134

List of Tables
3.1

Basic properties of GIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

19

3.2

Bit assignment for GIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

21

3.3

Temporal properties of GIS . . . . . . . . . . . . . . . . . . . . . . . . . . .

22

3.4

Basic properties of SIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

24

3.5

Temporal properties of SIS . . . . . . . . . . . . . . . . . . . . . . . . . . . .

24

4.1

Observation fields of ASCA . . . . . . . . . . . . . . . . . . . . . . . . . . .

28

5.1

Selected ASCA counterparts for ROSAT sources . . . . . . . . . . . . . . . .

38

5.2

ASCA sources detected in multiple observations . . . . . . . . . . . . . . . .

41

5.3

Final ASCA catalogue of X-ray sources in the SMC (source identification) .

48

5.4

Final ASCA catalogue of X-ray sources in the SMC (names and comments) .

52

5.5

Final ASCA catalogue of X-ray sources in the SMC (spectral parameters) . .

56

6.1

X-ray pulsars in the SMC . . . . . . . . . . . . . . . . . . . . . . . . . . . .

61

6.2

Spectral parameters for AX J0043−737 . . . . . . . . . . . . . . . . . . . . .

63

6.3

Spectral parameters for AX J0049−729 . . . . . . . . . . . . . . . . . . . . .

64

6.4

Spectral parameters for AX J0049−732 . . . . . . . . . . . . . . . . . . . . .

66

6.5

Spectral parameters for AX J0049.5−7323 . . . . . . . . . . . . . . . . . . .

67

6.6

Spectral parameters for AX J0051−733 . . . . . . . . . . . . . . . . . . . . .

68

xi

xii

LIST OF TABLES
6.7

Spectral parameters for AX J0051−722 . . . . . . . . . . . . . . . . . . . . .

71

6.8

Spectral parameters for AX J0051.6−7311 . . . . . . . . . . . . . . . . . . .

73

6.9

Spectral parameters for No. 43 (= RX J0052.1−7319?) . . . . . . . . . . . .

74

6.10 Spectral parameters for XTE J0054−720 . . . . . . . . . . . . . . . . . . . .

75

6.11 Spectral parameters for 1WGA J0053.8−7226 . . . . . . . . . . . . . . . . .

76

6.12 Spectral parameters for SMC X-2 . . . . . . . . . . . . . . . . . . . . . . . .

78

6.13 Spectral parameters for XTE J0055−724 . . . . . . . . . . . . . . . . . . . .

80

6.14 Spectral parameters for AX J0057.4−7325 . . . . . . . . . . . . . . . . . . .

82

6.15 Spectral parameters for AX J0058−7203 . . . . . . . . . . . . . . . . . . . .

83

6.16 Spectral parameters for RX J0059.2−7138 . . . . . . . . . . . . . . . . . . .

85

6.17 Spectral parameters for RX J0101.3−7211 . . . . . . . . . . . . . . . . . . .

86

6.18 Spectral parameters for 1SAX J0103.2−7209 . . . . . . . . . . . . . . . . . .

87

6.19 Spectral parameters for AX J0105−722 . . . . . . . . . . . . . . . . . . . . .

90

6.20 Spectral parameters for XTE J0111.2−7317 . . . . . . . . . . . . . . . . . .

91

6.21 Spectral parameters for SMC X-1 . . . . . . . . . . . . . . . . . . . . . . . .

94

6.22 Minimum luminosities of XBPs . . . . . . . . . . . . . . . . . . . . . . . . . 100
6.23 Spectral parameters for 0045−734, 0057−7226, and 0103−726 . . . . . . . . 103
6.24 Plasma parameters for 0045−734, 0057−7226, and 0103−726 . . . . . . . . . 103
6.25 Spectral parameters for 0047−735 . . . . . . . . . . . . . . . . . . . . . . . . 105
6.26 Spectral parameters for 0046−735 . . . . . . . . . . . . . . . . . . . . . . . . 109
6.27 Spectral parameters for 0049−736 . . . . . . . . . . . . . . . . . . . . . . . . 110
6.28 Spectral parameters for 0056−725 . . . . . . . . . . . . . . . . . . . . . . . . 111
6.29 Spectral parameters for IKT1 . . . . . . . . . . . . . . . . . . . . . . . . . . 112
6.30 Spectral parameters for AX J0048.2−7309 . . . . . . . . . . . . . . . . . . . 114

LIST OF TABLES

xiii

6.31 Spectral parameters for AX J0128.4−7329 . . . . . . . . . . . . . . . . . . . 115
7.1

Nonpulsating HMXBs and candidates in the SMC . . . . . . . . . . . . . . . 118

8.1

Source populations in the SMC and in our Galaxy . . . . . . . . . . . . . . . 126

xiv

LIST OF TABLES

Chapter 1
Introduction
Bright X-ray sources such as supernova remnants (SNRs) and binaries containing a neutronstar (NS) or a black hole (BH) are “relics” of massive stars. Therefore, the number and spatial
distribution of each class of X-ray source carry key information on the star forming activity,
dynamics, and structure of the host galaxy. However, our knowledge of the populations ofbright X-ray sources in our Galaxy is highly incomplete due to Galactic absorption in the
low energy band, the large angular size that must be observed, and limited information on
the distance of the relevant objects.
The Small Magellanic Cloud (SMC), a satellite of our Galaxy, is the next nearest neighbor after the Large Magellanic Cloud (LMC). The proximity (60 kpc is assumed in thispaper; van den Bergh 2000), reasonable angular size (∼ 3◦ × 3◦ ), and low interstellar absorption in the direction of the SMC are all favorable for an unbiased survey of the X-ray source
populations throughout the entire galaxy. Surveys of soft X-ray sources (below ∼ 2 keV)have been carried out with the Einstein and ROSAT satellites. Haberl et al. (2000) presents
the most complete catalogue, which contains 517 sources detected with ROSAT PSPC (Position Sensitive Proportional Counter). Sasaki, Haberl, & Pietsch (2000) used ROSAT HRI(High Resolution Imager) to determine the most accurate positions of 121 sources.
X-ray sources belonging to different classes have different properties in their spectra
and temporal behaviour. X-rays from young SNRs are either due to a shock heated plasma(in shell-like SNRs) which preferentially emits soft (below ∼ 2 keV) line dominated Xrays. Hard and featureless spectra are observed from Crab-like SNRs (due to release of therotational energy of isolated NSs) and from some shell-like SNRs such as SN 1006 (due to1

2

CHAPTER 1. INTRODUCTION

synchrotron emission from shock-accelerated electrons). Unlike SNRs, the other classes of
X-ray sources emit time-variable X-rays. NSs with a low mass companion star (low mass
X-ray binaries: LMXBs) have hard spectra and often exhibit X-ray bursts. BH binaries
have bimodal behavior and show both high-soft and low-hard spectral states. NSs with ahigh mass stellar companion (high mass X-ray binaries: HMXBs) have even harder X-ray> 10–100. Many HMXBs
spectra and often exhibit long-term flux variations of a factor ∼also show coherent pulsations (X-ray binary pulsars: XBPs). Consequently, most bright
X-ray sources emit copious hard X-rays (above ∼ 2 keV). However, due mainly to a lackof imaging capability in the hard X-ray band, a systematic survey of the populations ofhard X-ray sources had not previously been made. The Advanced Satellite for Cosmology
and Astrophysics (ASCA: Tanaka, Inoue, & Holt 1994) is equipped with wide-band Xray imaging instruments having a large field of view, and hence is optimized for an X-ray
population study.
In this paper, we report the results from multiple observations of the SMC with ASCA.
The whole SMC region including the main body and the eastern wing has been completely
covered. A brief review of the SMC is given in chapter 2. The ASCA satellite and its
equipment are described in chapter 3. Observation fields and the methods of data reduction
are presented in chapter 4. In chapter 5, we describe the method of data analysis and present
the results in catalogues. We give comments on specific sources in chapter 6. We show asource classification method using the spectral hardness ratio in chapter 7. In chapter 8,
we combine our results with those of other studies to derive the source populations and thespatial distribution. Based on the results, we discuss the star forming activity of the SMC.
Finally, we summarize this study in chapter 9.

Chapter 2
Review— the Small Magellanic Cloud

Figure 2.1: Photograph of the southern hemisphere sky behind the 4-m telescope of the Cerro
Tololo Inter-American Observatory. Diffuse light from three galaxies is seen in the upper leftc(SMC), lower left (LMC), and to the right (our Galaxy). (°Roger
Smith, AURA, NOAO, NSF.)

3

4

CHAPTER 2. REVIEW — THE SMALL MAGELLANIC CLOUD

Figure 2.2: Optical images of the LMC (left) and the SMC (right), where north is up and east isleft. The “wing” is protruding out from the main body of the SMC to the east, connecting to thecMagellanic Bridge. (°Anglo-Australian
Observatory)

The Small and Large Magellanic Clouds (SMC and LMC) are conspicuous objects in thesouthern hemisphere sky, as seen in the photograph (figure 2.1). Magnified images of theseclouds are shown in figure 2.2: both have a barred shape with dimensions of ∼ 2.◦ 5 × 1◦
(SMC) and ∼ 3◦ × 1◦ (LMC).
These clouds have probably been known since ancient times, as seen in legends of many
tribes of Australian Aboriginals. However, these people produced few documents which arestill preserved. Therefore, it was the discovery by Ferdinand Magellan on his expedition in
1519 that brought the two clouds to the knowledge of people outside Australia. The clouds
are thus named after Magellan.
In this chapter, we present a brief review of the Magellanic Clouds with particularemphasis on the SMC.

5

Magellanic St
ream

2.1. THE MAGELLANIC SYSTEM

BridgeLeading Arm

0
..1

4M

6
...
Figure 2.3: Global H i distribution in the Magellanic System, taken from Mathewson(1985).
SAu
5PA
198

Coordinates are Galactic. The Leading Arm was discovered later (Putman et al. 1998).

2.1

The Magellanic System

The Magellanic Clouds, belonging to the local group1 , are the nearest satellite galaxies ofour own. The total masses of the Magellanic Clouds are about 2 × 109 M¯ (SMC) and2 × 1010 M¯ (LMC), which are about 1% and 10% of the mass of our Galaxy. The fractionof H i among the total mass is ∼ 30% (SMC) and < 8% (LMC), both of which are larger
than that of our Galaxy (∼ 1%). This fact indicates that there still remains much material
to produce stars, and thus the Magellanic Clouds are in a younger evolutional phase thanour Galaxy. This is also supposed from the metal-poor interstellar matter (ISM), with metalabundances of ∼ 20% (SMC) and ∼ 30% (LMC) of those in the solar vicinity (Russell &Dopita 1992).
1

Galaxies gravitationally bound to our Galaxy; 36 members are listed in van den Bergh (2000).

6

CHAPTER 2. REVIEW — THE SMALL MAGELLANIC CLOUD

The H i distribution shown in figure 2.3 indicates a close connection between the Magellanic Clouds and our Galaxy. The Magellanic Clouds themselves are connected by the H i“Bridge” and surrounded by a common H i envelope. In addition, our Galaxy and the Magellanic Clouds are connected by two streams of H i gas, the Magellanic Stream (Mathewson,
Cleary, & Murray 1974) and the Leading Arm (Putman et al. 1998). These structures are,as a whole, called the Magellanic System.

2.2

Interaction between the Magellanic Clouds and our
GalaxyThere are many pieces of evidence indicating interactions between the Magellanic Clouds andour Galaxy. The existence of the eastern wing of the SMC, extending out in the direction of
the LMC, and the H i Bridge is indicative of interactions between the two clouds. Stars in the
Bridge and around the tip of the SMC wing were found to be as young as 108 yr, suggesting
a common origin (Irwin, Demers, & Kunkel 1990). Vallenari et al. (1996) investigated theages of stars outside the main body of the LMC and concluded that the stars on the westernside (toward the Bridge) are the youngest (2–3 × 109 yr) among their survey fields. The
existence of young stars (108−9 yr) in and around the Bridge is a hint of tidal interactions
which influenced star-forming activity (Vallenari et al. 1996). The Leading Arm and theMagellanic Stream indicate interactions between the clouds and our Galaxy. The former isleading the motion of the clouds while the latter is a trailing stream. The idea that bothstreams were torn from the clouds by tidal interactions is favored because the existence ofboth the leading and trailing stream is naturally expected in tidal interactions (Putman etal. 1998).
Many numerical simulations of the orbits of the clouds and our Galaxy have been performed to reproduce the Bridge and the Magellanic Stream via such tidal interactions (e.g.
Murai & Fujimoto 1980; Gardiner, Sawa, & Fujimoto 1994; Gardiner & Noguchi 1996).
Murai & Fujimoto (1980) showed that the Magellanic Clouds have been a binary galaxy
system for a long time (∼ 10 Gyr), which is consistent with the existence of the common
H i envelope. They also showed that the Magellanic Stream was torn off from the SMC by
tidal force and that the two clouds experienced a close encounter about 2 × 108 yr ago, which
synchronizes with the ages of stars in the Bridge. Recently, two observational constraints
have been obtained. One is the direct measurement of the proper motions of the SMC and2.3. STRUCTURE OF THE SMALL MAGELLANIC CLOUD

7

the LMC with the Hipparcos satellite (Kroupa & Bastian 1997). The proper motion inequatorial longitude and in equatorial latitude is (1.23 ± 0.84, −1.21 ± 0.75) mas yr−1 for
the SMC and (1.94 ± 0.29, −0.14 ± 0.36) mas yr−1 for the LMC. The other constraint is thestudy of the age distribution of star clusters in the LMC by Girardi et al. (1995): the LMC
experienced enhanced cluster formation about 2 × 109 yr and 1 × 108 yr ago. Turbulence of
the interstellar matter triggered by two close encounters may have resulted in such burstsof cluster formation. Sawa, Fujimoto, & Kumai (1999) showed that some orbits with two
close encounters 3 × 109 yr ago and 2 × 108 yr ago could reproduce the Magellanic Stream
and were consistent with the Hipparcos measurement. It is open to question whether theseorbits can explain other features of the Magellanic System, such as the Leading Arm.

2.3

Structure of the Small Magellanic CloudAs seen from figure 2.2, most of the visible stars in the SMC are distributed along thenortheast-southwest geometry (main body), and some are located in the eastern wing.
Maragoudaki et al. (2001) carried out a wide-field (6◦ × 6◦ ) observation covering the SMC
in the U , V , and I wavebands. They determined ages of the stars according to their colorand magnitude, and derived spatial distributions of stars in seven ranges of ages (figure2.4). They found that older stars show a rather regular distribution which is typical for aspheroidal body. On the other hand, the distribution of younger stars is highly asymmetric with a bent-barred shape (main body) and the eastern wing. The similar conclusions
were drawn by Zaritsky et al. (2000) and Cioni, Habing, & Israel (2000) with observations
in different wavebands. Therefore, the bar and wing of the SMC cannot originate solely
from tidal effects; in other words, a scenario in which the bar and wing consist of stars
extracted from the central SMC by tidal forces is rejected. Zaritsky et al. (2000) suggested
that the origin of those shapes must be a relatively recent star formation triggered by ahydrodynamic interaction between gaseous components. The younger stars’ distribution isqualitatively very similar to the H i distribution (figure 2.5). A large velocity gradient along
the bar axis was found (Stanimirović et al. 1999), which might arise from a collision of two
gaseous components along the bar.
The three-dimensional structure of the SMC can be derived by measuring distances tothe member stars. Mathewson (1985) determined the distances to 161 Cepheids using theperiod-luminosity relation as shown in figure 2.6. The depth of the SMC was thus revealed

8

CHAPTER 2. REVIEW — THE SMALL MAGELLANIC CLOUD
(a)

(b)

(c)

(d)

(e)

(f)

(g)

Figure 2.4: Isodensity contour maps of stars with different ages taken from Maragoudaki et al.
(2001). Ranges of ages are (a) > 2 × 109 yr, (b) 3.4–4 × 108 yr, (c) 1.7–3.4 × 108 yr, (d) 0.3–
1.7 × 108 yr, (e) 1.2–3 × 107 yr, (f) 0.8–1.2 × 107 yr, and (g) < 8 × 106 yr. The central region of (a)was masked bacause of the high concentration of stars. The coordinates are J2000.

to be 20–30 kpc, which is much larger than the projected size of the SMC (i.e., the SMC has
a large inclination angle).

1985PASAu...6..

2.3. STRUCTURE OF THE SMALL MAGELLANIC CLOUD

9

Figure 2.5: H i column density image of the SMC taken from Stanimirović et al. (1999).

Figure 2.6: Distances to 161 Cepheids in the SMC plotted against their projected distancesalong the major axis of the SMC (taken from Mathewson, 1985). The filled/open circles represent Cepheids with periods longer/shorter than 10 days (or younger/older than 5 × 107 yr). The
solid lines envelop the younger Cepheids. The dashed line demarcates the low velocity H i gas
component (near) and high velocity component (far).

1978ApJ...221L..37C

10

CHAPTER 2. REVIEW — THE SMALL MAGELLANIC CLOUD

Figure 2.7: X-ray view of the SMC in late 1970’s taken from Clark et al. (1978). The positions of
the three X-ray sources SMC X-1, SMC X-2 and SMC X-3 are shown by squares. The equatorial
coordinates (1950B) are presented. Contours and dots represent the surface brightness of the 21-cm
radio emission and the positions of O–B2 stars, respectively.

2.4

Previous High Energy Observations

2.4.1

X-ray surveys

Searches for X-rays from the Magellanic Clouds have been carried out since the very beginning of the X-ray astronomy. The first attempt to detect X-rays from the SMC was madewith proportional counters and scintillation counters which were launched in 1970 by a Thor
missile. In a scan observation with these instruments, the first X-ray source in the SMC,
known as SMC X-1 today, was discovered (Price et al. 1971). The position was subsequently
localized with an error box of ∼ 70 × 20 by the Uhuru satellite (Giacconi et al. 1972). X-ray
pulsations with a 0.716 s period from SMC X-1 were discovered by Lucke et al. (1976). Clark
et al. (1978) discovered two new transient sources, SMC X-2 and SMC X-3, with the SAS-3
satellite, which had a 12◦ × 12◦ field of view and thus covered the whole area of the SMC.
These sources are all considered to be HMXBs today. In that era, the X-ray view of the
SMC was rather sparse as shown in figure 2.7.

2.4. PREVIOUS HIGH ENERGY OBSERVATIONS

11

Figure 2.8: X-ray mosaic image of the SMC obtained with ROSAT PSPC, taken from Haberl etal. (2000). The equatorial coordinates (J2000) are overlaid.

The Einstein satellite, which was sensitive to soft X-rays (below a few keV) and had
a nice imaging capability, discovered a number of sources in the SMC. Wang & Wu (1992)
carried out a comprehensive analysis of the Einstein data of the SMC and detected 70 sources.
Some of these sources were identified with radio SNRs, H ii regions, HMXBs, background
AGNs, and foreground stars, while over half were left unidentified.
The ROSAT satellite greatly enchanced the number of X-ray sources in the SMC. Haberl

12

CHAPTER 2. REVIEW — THE SMALL MAGELLANIC CLOUD

et al. (2000) analyzed the data from 31 ROSAT PSPC observations which cover the wholearea of the SMC, and catalogued 517 discrete sources. The mosaic image (figure 2.8) indicatesthat many sources are concentrated in the main body and the eastern wing. A smaller area
was also covered by 71 ROSAT HRI observations. Sasaki, Haberl, & Pietsch (2000) usedthose data and catalogued 121 sources with a good positional accuracy of ∼ 100 –1000 . Cowley
et al. (1997) and Schmidtke et al. (1999) found five Be/X-ray binaries based on the positional
coincidence of the ROSAT HRI sources with Be stars. Haberl & Sasaki (2000) found thatmany of the known Be/X-ray binaries have an Hα emission line object (Meyssonnier &Azzopardi 1993) as a counterpart. This fact indicates that X-ray sources correlated with
an emission line object are candidates for Be/X-ray binaries. Thus they made a crosscorrelation of the 517 ROSAT PSPC sources and a few other sources with the catalogueof Meyssonnier & Azzopardi (1993) and found ∼ 30 new candidates for Be/X-ray binaries.
X-ray binaries were also identified from the flux variability of the ROSAT sources (Kahabka& Pietsch 1996). Supersoft X-ray sources (SSSs) have a very soft spectrum with a blackbody
temperature of 15–80 eV (Kahabka & van den Heuvel 1997); thus, interstellar absorption
severely reduces the flux of SSSs. Therefore, the Magellanic Clouds, which are located athigh Galactic latitude are suitable sites for studies of SSSs. Five SSSs were discovered in
the SMC, whose spectral and temporal properties are presented by Kahabka, Pietsch, &Hasinger (1994).

2.4.2

GeV and TeV γ-ray observations

GeV γ-ray observations were carried out with CGRO EGRET (Sreekumar et al. 1993). The
result was negative, giving an upper limit of 0.5 × 10−7 photon cm−2 s−1 in the > 100 MeV
range. GeV γ-rays are produced via the interaction between cosmic ray nucleons and theinterstellar matter as follows:p + p → p + p + π0 + · · · ,
π 0 → 2γ.
Sreekumar & Fichtel (1991) calculated the γ-ray emission from the SMC based on
the H i measurement, non-thermal (synchrotron) radio observations and assumptions of
the dynamical state. They derived the cosmic ray density from the radio data, and thencalculated the γ-ray emission via the interaction between cosmic rays and the H i interstellar
matter. The flux above 100 MeV thus derived was ∼ 10−8 photon cm−2 s−1 . This is consistent2.4. PREVIOUS HIGH ENERGY OBSERVATIONS

13

with the EGRET result (flux < 0.5 × 10−7 photon cm−2 s−1 ). The amount of cosmic rays
per unit mass of the interstellar matter was found to be several times smaller than thatin our Galaxy. Sreekumar & Fichtel (1991) concluded that this could be attributed to the
dynamical state of the SMC: the interstellar matter in the SMC is not dynamically stable
and rather leaky for cosmic rays.
TeV γ-ray observations of the SMC were carried out with a Cerencov telescope (Allenet al. 1993). The intensity of γ-rays in the direction of SMC X-1 showed positive excess,although it was not statistically significant. They gave upper limits of γ-ray flux above
30 TeV from SMC X-1 and the whole SMC as 2.6 × 10−13 photon cm−2 s−1 and 19.3 ×
10−13 photon cm−2 s−1 msr−1 , respectively.

14

CHAPTER 2. REVIEW — THE SMALL MAGELLANIC CLOUD

Chapter 3
Instruments
3.1

Overview of the ASCA Satellite

ASCA (Advanced Satellite for Cosmology and Astrophysics: Tanaka, Inoue, & Holt 1994)
was the 4th Japanese satellite for X-ray astronomy, and was launched on Feb. 20, 1993.
After operating for eight years, ASCA re-entered the atmosphere on March 2, 20011 .
ASCA had four X-ray telescopes (XRTs, Serlemitsos et al. 1995) and four instruments
at the focal planes: two Gas Imaging Spectrometers (GISs, Ohashi et al. 1996) and two
Solidstate Imaging Spectrometers (SISs, Burke et al. 1994). This equipment made it possible
to simultaneously obtain X-ray images and spectra in the 0.5–10.0 keV band.

3.2

XRT: X-ray Telescope

X-rays can be totally reflected if their incident angle is smaller than a critical value on theorder of ∼ 1◦ . The critical value is smaller for higher energy X-rays. Mirrors of X-ray
telescopes which utilize this mechanism are usually placed in a Wolter type-I configuration:
a combination of parabolic and hyperbolic shells. Incident X-rays are reflected in series by
each of the shells, and focused on the focal plane (see figure 3.2). The effective area canbe increased by nesting the shells so that they have a common focus. The previous X-ray
1

Surprisingly, this is just one day after Prof. Minoru Oda passed away.

15

16

CHAPTER 3. INSTRUMENTS

S-band antennauzuki

X-band antenna
Shunt regulator


Extendable
optical bench

Sun shade

Solar paddle



NSAS

Y

SSAS

X-ray telescopeX

Z
Satellite coordinates

Figure 3.1: Schematic view of ASCA in orbit.
satellites Einstein and ROSAT had this type of X-ray mirror configuration and provided
high resolution images up to a few keV.
However, since this type of mirror is heavy, a telescope with a large effective area cannot
be launched. Therefore an approximation of the Wolter type-I configuration, which consistsof two conical foils, was adopted for the XRTs of ASCA. To get a large effective area forhigher energy X-rays, the incident angle of each foil must be smaller, which results in asmaller projected area. Nesting a large number of foils compensates for this. According tothis concept, each foil of the ASCA XRTs was made thin (125 µm) and 120 such foils were
nested in a coaxial configuration into one quadrant. As a result, the field of view (FOV) is
180 –240 (FWHM) and the effective area of each XRT is ∼ 300 cm2 below 2 keV and ∼ 50 cm2
at 10 keV. The sum of the effective areas of the four XRTs is about 5 times larger than thatof Einstein, or 2 times larger than that of ROSAT, at ∼ 1 keV. Plots of the effective area ofa single XRT vs. incident X-ray energy and off-axis angle are shown in figure 3.3 and figure
3.4, respectively.
However, utilizing an approximation of the Wolter type-I configuration and very thinfoils resulted in a poorer focusing ability. Figure 3.5 shows the point spread function (PSF)
at several incident positions. The apparent butterfly-like structure is due to the junctions
between the four quadrants of each XRT. The half power diameter (HPD) is ∼ 30 , and 80%of the incident energy is encircled in a 60 diameter for an on-axis source. Two point sources
can be distinguished if the separation is over ∼ 3000 .
Stray light (photons from outside of the FOV) is another weak point of the ASCA XRTs.
It is caused by irregular manners of reflection, such as a single reflection or more than three3.2. XRT: X-RAY TELESCOPE

parabola17

hyperbola

focal plane
Figure 3.2: Mechanism of an X-ray telescope which utilizes Wolter type-I optics.

Figure 3.3: Effective area of a single XRT as a function of the energy of the incident X-ray.
Dependence on the off-axis angle θ and azimuthal angle φ is also shown.

18

CHAPTER 3. INSTRUMENTS

Figure 3.4: Effective area of a single XRT as a function of the off-axis angle. The data forincident X-rays of 1.0 and 7.0 keV are shown.

Figure 3.5: Point spread function of the ASCA XRT.
reflections, or reflections at the back side of the foils. Figure 3.6 shows a GIS image of straylight from the Crab nebula, which is located 1◦ away from the center of the GIS.

3.3. GIS: GAS IMAGING SPECTROMETER

19

Figure 3.6: Stray light image obtained with GIS2. The Crab nebula is 1◦ from the center of
the GIS.

3.3

GIS: Gas Imaging Spectrometer

GISs (GIS2 and GIS3) are gas scintillation proportional counters which are placed on the
focal planes of XRTs. Compared to SIS, the characteristics of GIS are a larger FOV, higher
sensitivity in the hard X-ray band, and better timing resolution (see table 3.1 for basic
properties of GIS).
Figure 3.7 shows the operating mechanism of GIS. An incident X-ray photon is detectedby a gas mixture of Xe (96%) and He (4%), and a cloud of photo electrons is producedin the drift region. The cloud is made to drift by a weak electric field and enters the
scintillation region. It is then accelerated by a strong electric field there, but the Xe atoms
Table 3.1: Basic properties of GIS.
Size of one pixel
0.5 mm (' 00· 5)
FOV (diameter)
500
Energy band
0.7–10.0 keVEnergy resolution 8% at 5.9 keV (FWHM)
Timing resolution 62.5 ms for usual PH mode

20

CHAPTER 3. INSTRUMENTS

Figure 3.7: Schematic view of the operating mechanism of GIS.
are not ionized as in a gas proportional counter. Instead, they are excited by electrons with
sufficient energy, and when they de-excite, UV photons are emitted. Therefore the number
of UV photons is proportional to the number of electrons in the cloud, hence to the energyof the incident photon. Finally the UV photons are detected by an imaging phototube. Theenergy resolution is proportional to E −0.5 (here E is the incident X-ray energy) and is ∼ 8%
(FWHM) at 5.9keV. The effective area of a single GIS combined with an XRT is shown infigure 3.8.
GIS has two observation modes: pulse height (PH) mode and multi-channel pulse count(MPC) mode. Normally PH mode is used. In this mode, 31 bits are assigned to the pulse
height (PH), X position (RAWX), Y position (RAWY), rise-time (RT), spread of UV photons

3.3. GIS: GAS IMAGING SPECTROMETER

21

Figure 3.8: Dependence of the effective area of a single GIS combined with and XRT.
Dependence on the off-axis angle θ and azimuthal angle φ is also shown.
Table 3.2: Freedom of bit assignment for PH mode of GIS.
Possible numberof bits assigned
PH6, 8, 10∗
RAWX
6, 8∗
RAWY
6, 8∗
RT
0, 5∗ , 6, 8
SP
0∗ , 8
Timing bit 0∗ –10
Total
31
∗
: Default assignments.

(SP) and event time (“timing bit”); usually 10, 8, 8, 5, 0, and 0 bits, respectively. We canadjust the bit assignment within a limited degree of freedom as shown in table 3.2. MPC
mode, which has been rarely used, is for bright sources. In this mode, positional information
is not recorded. The timing resolution and maximum source intensity for various recording
rates are summarized in table 3.3. The gain is always monitored with a radio isotope 55 Fe
attached to each GIS, which emits monochromatic X-rays with an energy of 5.89 keV.

22

CHAPTER 3. INSTRUMENTS

Table 3.3: Timing resolution and maximum source intensity of GIS.
Mode Bit rate Timing resolution
Maximum(sec)source intensity∗
H
0.0625256 count/sec
PH
M
0.5∗
32 count/sec∗
L
2.0
8 count/sec
H
0.5
128 Crab†
MPC
M
4.0
16 Crab†L
16.0
4 Crab††: 1 Crab corresponds to ∼ 1000 count/sec.
∗: For the case when the timing bit is 0.

3.4

SIS: Solidstate Imaging Spectrometer

SISs (SIS0 and SIS1) are the first CCD cameras used for X-ray astronomy. As shown infigure 3.9, each SIS consists of four chips, called S0C0 (chip0 of SIS0), S0C1 (chip1 of SIS0),and so on to S1C3 (chip3 of SIS1). Compared to GIS, the characteristics of SIS are a higher
energy resolution, higher sensitivity in the soft X-ray band, and better position resolution.
The basic properties of SIS are given in table 3.4. The effective area of a single SIS combinedwith an XRT is shown in figure 3.10.
SIS has three observation modes: Faint, Bright and Fast mode. In the Faint mode,
positional information of the center pixel2 and the pulse-height of each 9 (3 × 3) pixels
surrounding the center pixel are recorded as one event. In the Bright mode, each event isclassified into eight grades (grades 0–7, see figure 3.11). The pulse heights are summed up
according to the determined grade. The sum of the pulse heights, the positional informationof the center pixel, and the grade are then recorded as one event. Events of grade 0, 2,
3 and 4 are usually regarded as X-ray events. In these two modes, we can use 1, 2 or 4
chips of each SIS according to the scientific goal of the observation3 . The timing resolution
is determined by the time necessary to read out all the chips being used: 4 s per chip (seetable 3.5). The Fast mode is for very bright sources and is operated using only one chip of2

The center pixel is the pixel where the pulse-height distribution is a maximum among the surrounding3 × 3 pixels.
3
We use the terms like “1-CCD Faint mode” or “4-CCD Bright mode.”3.4. SIS: SOLIDSTATE IMAGING SPECTROMETER

23

µ
4
1
1P

2
2P

1
2
4
1

P

2P

Figure 3.9: SIS chip configuration.

Figure 3.10: Dependence on energy of the effective area of a single SIS combined with an
XRT. Dependence on the source locations is also shown.

24

CHAPTER 3. INSTRUMENTS

Table 3.4: Basic properties of SIS.
Size of one pixel
0.027 mm (' 00· 027)
FOV
110 × 110 per one chip
Energy band
0.5–10.0 keVEnergy resolution 2% at 5.9 keV (FWHM)
Timing resolution 4/8/16 s for 1/2/4-CCD mode

Table 3.5: Timing resolution and maximum source intensity of SIS.
Mode Bit rate Timing resolution
Maximum source(sec)intensity (count/sec)∗
H
4/8/16128
FaintM
4/8/16∗16∗L
4/8/164∗H
4/8/16512BrightM
4/8/16∗64∗
L
4/8/16
16
H
0.016
1024
Fast
M
0.016
128
L
0.016
32
∗
: For 1/2/4-CCD modes, respectively.

each SIS. In this mode, a better timing resolution of 16 ms is obtained at the expense of
positional information.
The correction sequence of pixel dark levels is as follows. In orbit, the dark level averaged
over each 16 × 16 pixels (“dark frame”) is estimated by the onboard CPU and subtracted
from the PH of each pixel. The Dark Frame Error (DFE) is the difference between the real
average dark level and the dark frame. On the ground, the DFE averaged over the wholechip is estimated using the PH values of the corner pixels4 , and is subtracted again.
Only the “average” dark level is taken into account in these correction methods. Therefore, a difference between the dark level of each pixel and the DFE-corrected PH value still
exists. This is called Residual Dark Distribution (RDD). RDD is corrected by the RDD
4Corner pixels are the pixels on the corners of each array of 3 × 3 pixels around the center pixel.

3.4. SIS: SOLIDSTATE IMAGING SPECTROMETER[Definition]

25[Examples]

Grade 0= perfect single
Grade 1= single+ detouched corners
Grade 2= vertical single-sided split+ detouched corners
Grade 3= left single-sided split+ detouched corners
Grade 4= right single-sided split+ detouched corners
Grade 5= single-sided split
with touched corners
Grade 6
= L-shape or square-shape+ detouched corners

The center pixel.
A pixel whose PH level is larger than the split threshold andwhich is included when summing up the PHs.
A pixel whose PH level is larger than the split threshold andwhich is not included when summing up the PHs.

Figure 3.11: Grade definitions (0–6) for one event of SIS. All the other patterns are classified
as grade 7.

26

CHAPTER 3. INSTRUMENTS

map, which contains information on the dark level of each pixel.
Pixels damaged by charged particles show various values of dark currents, according to
the degree of damage. This fact is the main cause of RDD. Therefore, data taken in 4-CCD
mode are more affected than data taken in 2/1-CCD mode because of the longer exposure
time. In addition, later observations are more affected than earlier ones because the degreeof damage gradually increases from the time of the launch.

Chapter 4
Observations and Data Reduction
4.1

Observation Fields

ASCA had observed the SMC region 22 times by the end of the mission, as summarized intable 4.1. Although the observations made before 1999 had aimed at specific objects suchas SNRs, X-ray pulsars, and a supersoft X-ray source, the assembly of these observations
had already covered most of the main body and the eastern wing of the SMC. In order tocover all the blank area, the survey project (SMC 1–10 in table 4.1) including long exposure
observations (SMC SW N1 and N2 in table 4.1) were performed. These observations havecovered most of the SMC region as shown in figure 4.1. In this study, we used all of theseobservation data and carried out various analyses in a coherent manner.

4.2

Data Reduction

In each observation, X-ray photons were collected with four XRTs and detected separately
with two GISs and two SISs. We rejected any data obtained in the South Atlantic Anomaly,or when the cut-off rigidity was lower than 4 GV, or when the elevation angle was lower
than 5◦ . Particle events in the GIS data were removed by a rise-time discrimination method.
SIS data obtained when the elevation angle from the bright earth was lower than 25◦ , orwith hot and/or flickering pixels, were also rejected. The effects of RDD (Residual Dark
Distribution) on the SIS data were corrected with the method given in Dotani et al. (1997)27

28

CHAPTER 4. OBSERVATIONS AND DATA REDUCTION

Table 4.1: SMC fields observed with ASCA.
ID

A
B
C
D
E
F
G
H
I
J
K
L
M
N
O
P
Q
R
S
T
U
V

Observation Date
(UT)
Start
End
1993/04/26 22:23
1993/05/12 08:01
1995/10/18 18:32
1996/05/21 02:25
1997/05/21 13:01
1997/11/13 03:58
1997/11/14 06:51
1997/12/12 11:05
1998/11/18 17:37
1998/11/24 09:39
1999/05/10 10:31
1999/05/11 13:59
1999/05/12 23:11
1999/05/14 00:35
1999/05/28 13:24
2000/04/04 00:09
2000/04/11 18:50
2000/04/25 12:09
2000/04/26 14:20
2000/04/27 16:40
2000/04/28 19:01
2000/04/29 21:30

04/27 23:41
05/13 12:26
10/19 21:31
05/23 06:31
05/22 01:31
11/14 06:51
11/15 08:22
12/12 23:55
11/19 15:11
11/25 10:51
05/11 14:00
05/12 16:46
05/14 00:35
05/15 01:56
05/29 15:11
04/07 01:30
04/17 12:11
04/26 14:21
04/27 16:40
04/28 19:01
04/29 21:31
05/01 00:01

Target Name

SMC X-1
E0102−72.2
SMC X-1
0103−72.6
1E 0035.4−7230
N19
N66
SMC X-3
XTE J0111.2−7317
SMC Center
SMC 1
SMC 2
SMC 3
SMC 4
SMC 5
SMC SW N1
SMC SW N2
SMC 9
SMC 6
SMC 7
SMC 8
SMC 10

Pointed Direction(J2000)R.A.
Dec.
1h 17m 04s
1h 04m 02s
1h 17m 04s
1h 04m 52s
0h 37m 20s
0h 47m 16s
0h 59m 25s
0h 52m 06s
1h 11m 26s
0h 56m 05s
0h 42m 06s
0h 54m 44s
1h 03m 33s
1h 12m 08s
1h 28m 36s
0h 42m 35s
0h 47m 30s
0h 58m 20s
0h 33m 00s
0h 53m 30s
1h 00m 00s
1h 20m 00s−73◦ 260 3400
−72◦ 010 5500
−73◦ 260 3500−72◦ 230 1100
−72◦ 140 1400
−73◦ 080 3000−72◦ 100 1200
−72◦ 260 0600
−73◦ 250 0700
−72◦ 520 1200
−73◦ 430 5300
−73◦ 300 0000
−73◦ 000 0000
−72◦ 470 4700
−73◦ 300 0300
−73◦ 400 3000−73◦ 090 1300
−73◦ 350 0000
−73◦ 450 0000
−71◦ 540 0000
−71◦ 200 0000
−72◦ 500 0000

Expos.
(ksec)
GIS/SIS4/—
35/31∗
41/—
78/61∗
23/—
43/38†
39/34†
20/—
26/—
27‡ /—
40/—
32/—
30/—
41/—
29‡ /—
93/—
177/96†
27/—
45/—
44/—
45/—
38/—

∗†: SIS was operated in 1-CCD mode (∗) or 2-CCD mode (†).
‡: In these observations, GIS was operated with a special bit assignment of 8-6-6-5-6 (J) or
8-8-8-5-2 (O), in which the time resolution was increased by a factor of 64 (J) or 4 (O), at thesacrifice of a reduction of energy resolution by a factor of 4 (J and O) a reduction of spatial
resolution by a factor of 4 (J).

4.2. DATA REDUCTION

29

600.5
-70°30′-71°00′-71°30′-72°00′-72°30′-73°00′-73°30′-74°00′

h

m

1 30

h

m

1 20

h

m

h

1 10

1 00

m

h

m

0 50

h

0 40

m

-74°30′

0.5
0.5
600.5
Figure 4.1: All observation fields with ASCA GIS superimposed on the optical image of the SMC
CHBN SMC_DSS_large.fits0

region from the Digitized Sky Survey (DSS). Equatorial coordinates with an equinox of 2000 are
also shown. Each observation field is represented by a ∼ 500 -diameter circle, the thickness of which
is proportional to the exposure time. The discontinuity seen in the DSS image is an artifact.

30

CHAPTER 4. OBSERVATIONS AND DATA REDUCTION

for observations carried out later than 1996. After the screening, the total available exposure
time for two GISs was ∼ 2000 ks.
In order to study X-ray source populations in a coherent manner, the GIS is moresuitable than the SIS because of its larger field of view, larger effective area at high energy,
and better time resolution. Therefore, we mainly used the GIS data in this study, while the
SIS data were used for peculiar objects which need a fine energy resolution and/or a better
spatial resolution.

Chapter 5
Systematic Analyses on Discrete
Sources
5.1

X-Ray Images

5.1.1

Images in the sky/detector coordinate systems

Images in each observation were constructed in the sky and detector coordinate systems(hereafter “sky images” and “detector images”, respectively), with the XSELECT package.
In a detector image, the positions of photons are recorded using the coordinates peculiar tothe detector, and are converted to sky coordinates via the average attitude of the satellite
during the observation. In a sky image, on the other hand, the position of each photon isdetermined based on the satellite attitude at the incident time of the photon. Therefore, astrong advantage of sky images is that blurring of the image due to attitude fluctuation iscorrected. On the other hand, the advantage of a detector image (only for the GIS) is its
larger FOV with a radius of ∼ 250 , compared to ∼ 200 for a sky image.
We thus used sky images for sources within the central ∼ 200 radius of the GIS (hereafter “inner circle”) and detector images for the concentric region of ∼ 200 –250 (hereafter“outer ring”). Although the outer ring region suffers higher background, larger calibration
uncertainty, and larger distortion of the PSF than the inner circle region, it was still usefulfor studies of source populations which need a large FOV. For SIS analyses, we always used
sky images.
31

32

5.1.2

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES

Source extraction

For each observation, X-ray sources were identified from images in the soft (0.7–2.0 keV), hard
(2.0–7.0 keV), and total (0.7–7.0 keV) bands. We smoothed the images with a Gaussian filter(σ = 3000 ), and examined the significance of each local peak. Photons were extracted froma circle of 30 radius centered on the peak, in which 90% of incident photons were contained(Serlemitsos et al. 1995), or from an ellipse at larger off-axis angles because of the distortionof the PSF. In several cases, a smaller circle/ellipse was used to avoid contamination fromnearby peaks. For obs. J, a larger circle was used, because in this observation the spatial
resolution was reduced by a factor of 4 and thus the images were blurred. These photon eventswere also used in the subsequent analyses described in the following sections. Backgroundregions were selected from the sky near each peak. We then derived the S/N ratio definedasn(P ) − n(B)
S
q
≡
,
(5.1)
N
n(B)
where n(P ) and n(B) represent the photon counts in the circle/ellipse at the peak and in
the background region, respectively. The local peak was identified as an X-ray source if the
S/N ratio exceeded 5 in at least one of the soft-, hard-, or total-band images. Several sourceswere detected in multiple observations as described in section 5.2.2. In all, 106 sources weredetected as catalogued in tables 5.3, 5.4, and 5.5.
In obs. Q, No. 26 and No. 27 were resolved only in the SIS image with a separation of∼ 1.0 3. No. 85 was detected only in the hard band because of severe contamination from
No. 81 in the soft band. No. 88 and No. 89 are located near the calibration isotope of GIS3,
thus we used only GIS2 to estimate the significance. Since No. 88 and No. 89 are separated
by only ∼ 1.0 5, which caused severe mutual contamination, we used very small circles toestimate their S/N ratio. We found that the S/N of No. 89 well exceeds 5, while that of No.
88 is slightly less than 5. However, a local peak at No. 88 was evidently also found in GIS3
(although no quantitative estimation is possible), thus we regard No. 88 as an X-ray source.

5.1.3

Construction of mosaic images

Figure 5.1 and 5.2 show the mosaic images in the soft and hard bands, that were createdfrom images in each pointing observation. The method of creating a mosaic image fromASCA GIS images was developed by Sakano (2000), and thus we only give a brief summary

5.2. POSITION DETERMINATION

33

here.
1. We created images in the soft and hard bands for each observation, and cut off outerregions where non–X-ray background (NXB), contamination from the calibration isotope, and calibration uncertainty are large, according to the “GIS good mask” suppliedby the GIS team (Ueda 1996).
2. The GIS night-earth image provided by NASA/GSFC was used as the NXB image,for which the same cut-off was applied. The intensity of the NXB image was adjusted
according to the observation date, in consideration of the increasing rate of the NXB
reported by Ishisaki et al. (1997).
3. An effective exposure map was created for each observation. Differences in the detection efficiency at different points on the GIS, which were caused by the vignetting of
the XRT, the supporting grid structure of the GIS, and other factors, are included inthis map.
4. The images, the NXB map, and the effective exposure map were smoothed with aGaussian filter of σ = 2.5 pixels.
5. The smoothed NXB map was subtracted from the smoothed image.
6. The NXB-subtracted smoothed images from all the observations were put together
into a mosaic image, according to the pointing direction and the roll angle of eachobservation.
7. Finally, the mosaic image was divided by the smoothed exposure map. The exposurecorrected mosaic images thus obtained in two energy bands are given in figures 5.1 and5.2. These images are combined into a two color image and presented in figure A.1.

5.2

Position Determination

We first determined the position of each source simply by the coordinates of the peak pixelin the smoothed GIS images, and then performed a correction developed by Gotthelf et al.
(2000). This correction compensates for the positional uncertainty caused by the ASCA
attitude error, which depends on the temperature of the base-plate of the star-tracker and

34

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES

600-71°00′400-72°00′200-73°00′

0
-200
-400

-74°00′

h

m

1 20

h

m

1 00

h

m

0 40

-400
-200
0 obtained with200
600keV),
Figure 5.1: Mosaic
image of the SMC
ASCA GIS in400
the soft band (0.7–2.0
overlaid with equatorial coordinates (J2000). The effects of non–X-ray background, telescope vignetting, and difference of exposure time between observations were corrected as described in the
text. Contour levels are linearly spaced. Of the two observations centered on SMC X-1 (obs. A
and C), only obs. C was used, in which SMC X-1 was much fainter.

5.2. POSITION DETERMINATION

35600-71°00′400-72°00′200-73°00′

0
-200
-400

-74°00′

h

m

1 20

h

m

1 00

h

m

0 40

-400 Figure
-200
200
400
600
5.2: The same as0
figure 5.1, but for the hard band (2.0–7.0 keV).

36

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES

the geometry of the solar illumination. SIS images were used only for resolving No. 26 andNo. 27 in obs. Q.
According to this correction, the coordinates of some X-ray pulsars with an “AX J”
name (which have been included in previous publications) are now inconsistent with thesource name; for example, the coordinates of AX J0051.6−7311 (No. 40; Yokogawa et al.
2000b) are now (00h 51m 44.s 5, −73◦ 100 3400 ). In this paper, we do not rename these sources toavoid name confusion, and adopt the names used in the first publications for each pulsar.

5.2.1

Absolute accuracy of the position

The correction developed by Gotthelf et al. (2000) reduces the systematic positional uncertainty to 2400 (90% error radius) for sources detected in the central 100 radius of the GIS.
However, additional errors from the photon statistics and the method of position determination, and errors for the sources located out of the central 100 radius, are unknown. Therefore,
we examined the “practical” errors for sources detected anywhere in the GIS as follows.
So far the ROSAT HRI catalogue (Sasaki, Haberl, & Pietsch 2000) presents the mostaccurate positions for the SMC X-ray sources, with an error radius of ∼ 1–1000 . Several
sources in the ROSAT PSPC catalogue (Haberl et al. 2000) also have a small error radius< 1000 . Therefore, we investigated the separation angles between the ROSAT sourcesof ∼and their ASCA counterparts, which would represent the “practical” errors for the ASCA
sources.
We primarily selected ASCA counterparts for ROSAT sources that were within 9000 ofeach ROSAT source. In order to reject accidental coincidences and ambiguous counterparts,
we further added rather conservative criteria:
1. For sources catalogued in both the PSPC and HRI catalogues, only HRI sources wereused, which provide more accurate positions.
2. Only ROSAT sources with an error radius smaller than 700 were used. Since the error
radii for ASCA sources are > 2400 , the additional error from the ROSAT sources is< 100 when we take a root sum square of all the errors.
3. For a ROSAT source with a Be star or a supergiant companion, the ASCA counterpart> 1 s. The procedure for pulse
should exhibit coherent pulsations with a period of ∼

5.2. POSITION DETERMINATION

37

Figure 5.3: Plot of the separation angles between the ROSAT–ASCA counterparts as a function of
the off-axis angle of the ASCA sources (see table 5.1). Filled and open circles indicate the sources
detected with ROSAT HRI and PSPC, respectively.

detection is described in section 5.3. This criterion selects well-established X-ray binary
pulsars (hereafter, XBPs). Although ASCA source No. 94 in obs. C exhibited nosignificant pulsations, it entered an eclipse phase as the ephemeris predicts for SMC X1 (Wojdowski et al. 1998), thus we regard No. 94 in obs. C as SMC X-1.
4. For a ROSAT source at the position of a radio SNR, the ASCA counterpart should
exhibit a soft spectrum with emission lines from ionized atoms. The method used todetect emission lines is described in section 5.4.3. This criterion selects bright thermal
SNRs.
According to these criteria, we selected 19 pairs of ASCA–ROSAT counterparts assummarized in table 5.1. Although No. 67 is certainly the counterpart for RX J0059.2−7138(see section 6.1.16), this pair is not included in table 5.1 because No. 67 is detected at thevery edge of the GIS (or may be slightly outside of the GIS) and so the position determination
is not reliable. We show the separation angles as a function of the off-axis angle of the ASCA38

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES

Table 5.1: ASCA counterparts for ROSAT sources selected with the conservative criteria (seetext).
— ROSAT∗ — ———— ASCA† ————
No. Error (00 ) No. Obs. ID Off-axis (0 ) Sep‡ (00 ) Remarks
34
1.8
32
F
17.0
31.9
Be-XBP AX J0051−733 (P = 323 s)
34
1.8
32
Q
9.7
24.0
Be-XBP AX J0051−733 (P = 321 s)
37
2.6
37
H
17.2
45.5
Be-XBP AX J0051−722 (P = 91 s)
38
2.7
36
Q
14.3
40.2
SNR 0049−736
41
1.5
40
Q
11.7
32.6
Be-XBP AX J0051.6−7311 (P = 172 s)
46
3.0
44
T
8.1
19.2
Be-XBP XTE J0054−720 (P = 167 s)
51
2.4
47
H
10.3
10.7
Be-XBP 1WGA J0053.8−7226 (P = 46 s)
58
2.6
51
G
21.9¶
67.1
Be-XBP XTE J0055−724 (P = 58 s)
73
6.4
61
G
11.6
29.8
Be-XBP? AX J0058−7203 (P = 280 s)
82
4.0
66
G
3.5
4.8
SNR 0057−7226
101
3.8
78
D
9.4
11.1
Be-XBP 1SAX J0103.2−7209 (P = 348 s)
107
3.3
81
B
6.7
23.7
SNR 0102−723
107
3.3
81
D
16.0
14.3
SNR 0102−723
109
3.6
82
D
5.2
10.8
SNR 0103−726
118
0.7
94
A
6.9
33.6
XBP (supergiant) SMC X-1 (P = 0.70 s)
118
0.7
94
C
5.3
33.8
XBP (supergiant) SMC X-1 (eclipse)
¶
118
0.7
94
I
23.5
35.9
XBP (supergiant) SMC X-1 (P = 0.70 s)
§
§
§
413
6.8
21
F
3.4
27.4
SNR 0045−734
413§
6.8§
21
Q
6.4
36.8§
SNR 0045−734
Note — “Be-XBP” indicates an XBP with a Be star counterpart.
∗: Source number and error radius of the ROSAT source presented by Sasaki, Haberl, & Pietsch (2000;
HRI catalogue). The last two lines are from Haberl et al. (2000; PSPC catalogue).
†: Source number, observation ID, and off-axis angle of the ASCA counterpart (table 5.3).
‡: Separation angle between the ROSAT and ASCA sources.
§: The position of the ROSAT source was assumed to be the central point of No. 413 and No. 419 inHaberl et al. (2000), because SNR 0045−734 is catalogued as those two sources due to its extended
X-ray emission. The error radius of the ROSAT source was taken from No. 413, while that for No. 419
is smaller, 4.00 7.
¶: Detected in the outer ring of the GIS.

5.2. POSITION DETERMINATION

39

source in figure 5.3. No clear correlation between the separation angle and the off-axis angle
could be found. Out of 17 ASCA sources detected in the inner circle (off-axis < 200 ), 15
sources have separation angles less than 4000 . Therefore, we tentatively conclude that the“practical” error radius for GIS sources detected in the inner circle is 4000 at 90% confidence
level, although the statistics are rather limited. This is similar to the result obtained fromthe more elaborate analysis by Ueda et al. (1999). For sources detected in the outer ring, noconstraint could be applied due to the small number of sources.
From the ROSAT and Einstein catalogues (Haberl et al. 2000; Sasaki, Haberl, & Pietsch2000; Wang & Wu 1992), we selected the counterpart for each ASCA source within a circle
of a radius ∼ 6000 for sources detected in the inner circle, or within a circle of a radius ∼ 7000
for sources detected in the the outer ring. Radii larger than the 90% error radius (4000 ) wereused in order to simply avoid missing identification.

5.2.2

Identification of sources detected in multiple ASCA observations

As shown in figure 4.1, neighbouring ASCA observation fields more or less overlap each other.
Therefore, a pair of sources detected in two observations within the overlapped region may
be the same source. In order to examine whether those pairs are the same source or not,we primarily selected pairs of sources within 9000 of each other, and classified them into four
grades as follows:
Grade A Both of the sources exhibit coherent pulsations with nearly the same period, or
exhibit emission lines from the same elements and have the same radio SNR as acounterpart (see sections 5.3 and 5.4.3 for the relevant analyses).
Grade B Both of the sources have soft spectra and have the same radio SNR as a counterpart. Pairs of No. 94 in obs. A and C and that in obs. I and C (SMC X-1) are also
included in this grade. Grades A and B surely consist of pairs of XBPs and thermal
SNRs.
Grade C Both of the sources are located near the same pulsar and their spectral parameters
are consistent with those of the pulsar. Sources of grade C are likely to be X-ray pulsars.
Grade D The remainder.

40

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES

Figure 5.4: Plot of the separation angles as a function of the larger off-axis angle (off-axis1 in table5.2) of ASCA sources detected multiple times.

We regarded pairs of grade A–C as the same source, i.e., sources detected in multiple observations, and thus labeled them with the same source number in the ASCA catalogues (table
5.3, 5.4, and 5.5). We summarize the separation angle and the off-axis angles of the twosources in each pair in table 5.2, while in figure 5.4 we give a plot of the separation angle< 6000 if both of the two
vs. the larger off-axis angle. We found that the separation angle is ∼< 7000 (7300 ) if at least one of the two is detectedsources are detected in the inner circle, or ∼in the outer ring. Therefore, we regard pairs of grade D to be the same source if they satisfy
the above condition, and labeled them the same source number. Pairs thus selected are also
summarized in table 5.2 and plotted in figure 5.4. After this selection, we concluded that
ASCA detected 106 sources with no double count.

5.2. POSITION DETERMINATION

Table 5.2: ASCA sources detected in multiple observations.
No.

Obs. OffObs. OffSep Grade Remarks
00
ID1 axis1 ( ) ID2 axis2 ( ) (00 )
21
Q
6.4
F
3.4
56.9
A
SNR 0045−734
32
F
17.0
Q
9.7
19.1
A
Be-XBP AX J0051−733
56
L
11.5
R
7.3
31.0
A
XBP AX J0057.4−7325
81
D
16.0
B
6.7
26.8
A
SNR 0102−723
∗
94
I
23.5
A
6.9
32.4
A
XBP (supergiant) SMC X-1
25
F
9.7
Q
4.2
16.5
B
SNR 0047−735
∗
36
F
19.2
L
16.9
70.0
B
SNR 0049−736
36
F
19.2∗
Q
14.2
59.4
B
SNR 0049−736
36
L
16.9
Q
14.2
53.2
B
SNR 0049−736
∗
94
I
23.5
C
5.3
33.4
B
XBP (supergiant) SMC X-1
94
A
6.9
C
5.3
1.0
B
XBP (supergiant) SMC X-1
17
P
4.3
K
2.7
21.7
C
Pulsar AX J0043−737
24
F
22.1∗
H
21.7∗
72.1
C
Be-XBP AX J0049−729
24
F
22.1∗
Q
17.4
34.5
C
Be-XBP AX J0049−729
24
H
21.7∗
Q
17.4
57.3
C
Be-XBP AX J0049−729
∗
30
L
20.7
F
16.3
48.9
C
Be-XBP? AX J0049.5−7323
30
F
16.3
Q
13.7
53.8
C
Be-XBP? AX J0049.5−7323
30
L
20.7∗
Q
13.7
39.9
C
Be-XBP? AX J0049.5−7323
40
F
20.9∗
L
20.7∗
31.8
C
Be-XBP AX J0051.6−7311
40
F
20.9∗
Q
11.7
53.7
C
Be-XBP AX J0051.6−7311
∗
40
L
20.7
Q
11.7
53.7
C
Be-XBP AX J0051.6−7311
∗
61
T
22.6
G
11.6
69.0
C
Be-XBP? AX J0058−7203
74
B
18.4
D
17.0
47.1
C
Be-XBP RX J0101.3−7211
74
D
17.0
G
10.4
19.0
C
Be-XBP RX J0101.3−7211
74
B
18.4
G
10.4
54.5
C
Be-XBP RX J0101.3−7211
78
B
13.9
D
11.1
18.4
C
Be-XBP 1SAX J0103.2−7209
78
G
19.0
D
11.1
20.0
C
Be-XBP 1SAX J0103.2−7209
78
G
19.0
B
13.9
25.8
C
Be-XBP 1SAX J0103.2−7209
20
Q
7.5
F
2.7
19.6
D
22
F
6.0
Q
3.3
25.7
D
†
26
F
10.5
Q
2.6
38.7
D
28
L
21.2∗
F
21.0∗
48.6
D
28
L
21.2∗
Q
20.9∗
68.7
D
∗
52
T
21.9
G
15.6
63.0
D
53
R
10.9
L
4.7
55.5
D
76
G
16.1
D
14.3
14.1
D
Note — Columns 1–6 show the source number, the observation IDs and off-axis angles for multiple
detections, and the separation angle between the positions in the two observations. Off-axis angles
are sorted so that off-axis1 is larger than off-axis2 . “Be-XBP” indicates an XBP with a Be star
counterpart.
∗: Detected in the outer ring of the GIS.
†: Detected with the SIS.

41

42

5.3

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES

Timing Analyses

We performed a Fast Fourier Transform (FFT) analysis on all of the sources to search forcoherent pulsations. For sources with high count rates, we used only high-bit rate data inorder to utilize the maximum time resolution (up to 62.5 ms). Otherwise we used highbit and medium-bit data simultaneously to achieve better statistics at the sacrifice of areduction in time resolution to 0.5 s (7.8125 ms for obs. J and 125 ms for obs. O; see thecaption of table 4.1). We detected coherent pulsations from 17 sources, eight of which are
new discoveries from this study. We show the power spectrum densities (PSDs) obtained
from No. 94 (SMC X-1; in obs. A) and No. 56 (AX J0057.4−7325; in obs. R) in figure 5.5(a) and (b) as examples of unambiguous detection. No. 83 (AX J0105−722) exhibits rather
weak signs of pulsations as shown in figure 5.5 (c). Detection of pulses from AX J0105−722
and No. 26 (AX J0049−732) was not straightforward because of contamination from nearby
sources. Details are described in section 6.1.19 and section 6.1.4, respectively.
In any observation, photon events are originally counted with a time spacing of 1/16 of
the nominal resolution and stored in the temporary memory. The events are then collectively
sent to the telemetry with a time spacing equal to the nominal resolution. Therefore, if theevent rate is so low as not to fill the memory, the time resolution could be 1/16 of the nominal
value (Hirayama et al. 1996). For this reason, we carried out FFT analysis on several faintsources with a time resolution of 31.25 ms, using the high- and medium-bit data. The 87 ms
pulsations from AX J0043−737 were thus discovered (see section 6.1.2 for further details).
In order to determine the pulse period precisely, we performed an epoch folding search
for the 17 sources from which pulsations were detected by FFT analysis. The orbital Doppler
effect was corrected only for SMC X-1, using the ephemeris presented by Wojdowski et al.
(1998). Since the eccentricity is near zero (< 0.00004), the intrinsic period P0 was derived
as
P
P0 =
,
(5.2)
1 − 2πa sin i sin Φ/cPorb
where P , a, i, Φ, and Porb are the pulse period determined by the epoch folding search,
orbital radius, inclination angle, orbital phase at the midst of the observation, and orbital
period, respectively. The derived pulse periods are presented in table 6.1.
We detected no pulsations by FFT analysis from three sources which are positionallycoincident with known pulsars: No. 43 (RX J0052.1−7319), No. 51 (XTE J0055−724), andNo. 94 in obs. C (SMC X-1). Therefore, we performed the epoch folding search around the5.3. TIMING ANALYSES

(a)

43

(b)

(c)

Figure 5.5:

Examples of the PSDs from (a) SMC X-1, (b) AX J0057.4−7325, and (c)AX J0105−722. Data points with a power less than 5 are not plotted. Pulsations were detected unambiguously from most pulsars like (a) or (b), while AX J0105−722 exhibits rather weak evidence
as seen in (c).

Figure 5.6: Periodigram for XTE J0055−724 around the known period of ∼ 59 s.
known periods. Since SMC X-1 was in the 0.6-day eclipse phase during obs. C, we onlyused the data from noneclipse times. Consequently, we detected a weak peak only fromNo. 51 at the known period of ∼ 59 s (figure 5.6), which is the evidence that No. 51 and
XTE J0055−724 are the same source. This period is, however, not presented in table 6.1
because of the low significance of the pulse detection.
We also searched for burst-like activities by using light curves binned with various timescales from ∼ 1 s to ∼ 1 hr. Although no source exhibited the bursts typical of LMXBs, No.
20 (IKT1, in obs. Q) showed a flare with a decay time of ∼ 2 × 104 s. Details are given in
section 6.5.1.

44

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES

5.4

Spectral Analyses

5.4.1

Overview

We analyzed the spectrum of each source and derived various parameters as given in table
5.5: the hardness ratio (HR), photon index (Γ), temperature (kT ), column density (NH ),flux (FX ), and absorption-corrected luminosity (LX ).
The analyses were not performed for Nos. 85, 88, and 89 because of severe contaminationof these sources (see section 5.1.2). Spectra from GIS2 and GIS3 were co-added to increase
the statistics, except for sources detected near the calibration isotope of either GIS andsources detected in only a single GIS1 . SIS spectra (SIS0 + SIS1) were used for No. 26 andNo. 27 in obs. Q in order to spatially resolve these sources, and for SNRs 0045−734 (No.
21), 0047−735 (No. 25), 0057−7226 (No. 66), 0102−723 (No. 81), and 0103−726 (No. 82)
to perform high resolution spectroscopy2 .
The parameters were derived from fitting the spectra with spectral models: different
models were used according to the nature of each source, as described in the subsequent
sections. For sources detected in multiple observations (table 5.2), we first fitted the spectrum
from each observation separately. Except for SMC X-1, the spectral parameters (except flux)in each observation were found to be consistent with each other. We thus simultaneously
fitted all the spectra with parameters linked between the observations, in order to apply
more stringent constraints. However, the flux was not linked in the simultaneous fitting, in< 10–20% error in the flux oforder to examine the flux variability; note that there could be ∼most sources. Hardness ratios for those sources were derived after adding the spectra fromall observations.

5.4.2

Hardness ratio

Spectral hardness ratio (HR) was derived by the definition
HR ≡
1

H −S
,
H +S

(5.3)

Since the FOVs of the two GISs are pointed toward slightly different directions, it is possible for a source
to be located at the very edge of one GIS and outside of the other GIS.
2
For 0047−735, the SIS spectrum in obs. F was not used because the statistics were too poor.

5.4. SPECTRAL ANALYSES

45

where H and S represent background-subtracted GIS count rates in 2.0–7.0 keV and 0.7–
2.0 keV, respectively. HR is not given for No. 27 in table 5.5 because this source was only
resolved with SIS in obs. Q. For the same reason, HR of No. 26 was derived only from thedata of obs. F.

5.4.3

SNRs

X-rays were detected from eight SNRs3 , 0045−734 (No. 21), 0047−735 (No. 25), 0057−7226(No. 66), 0102−723 (No. 81), 0103−726 (No. 82), 0046−735 (No. 23), 0049−736 (No. 36),
and 0056−725 (No. 64). The former five were detected with SIS and the latter three were
only detected with GIS. Before fitting the overall spectra, we distinguished thermal SNRs
from the others. Direct evidence for a thermal spectrum is the presence of emission lines
from highly ionized atoms. The brightest SNR, 0102−723, has been known to exhibit athermal spectrum (e.g., Hayashi et al. 1994), thus we looked for the existence of emission
lines in the spectra of the remaining seven SNRs as follows. We fitted the spectra in a narrowenergy band (∼ 1.1–3 keV) with a bremsstrahlung continuum and three narrow Gaussianlines centered at 1.34 keV, 1.85 keV, and 2.46 keV (Kα lines from He-like Mg, Si, and S,
respectively). If the line intensity I and its 90% error δI satisfied I − δI > 0 for at least one
line, we regarded the SNR as a thermal SNR. As a result, 0045−734, 0057−7226, 0102−723,
0103−726, and 0049−736 were regarded as thermal SNRs. We therefore fitted their spectrawith thin-thermal plasma models as described in section 6.3. For the other SNRs, we first
fitted the spectra with both a power-law model and a thin-thermal model in a collisional
ionization equilibrium (CIE) state (Raymond & Smith 1977), and finally adopted a powerlaw for 0056−725 and the CIE thermal model for 0047−735 and 0046−735, for reasonsdescribed in section 6.3. When fitting with thermal models, the metal abundances wereprimarily fixed at 0.2 solar, which is the mean value for the SMC ISM (Russell & Dopita
1992; hereafter “SMC mean value”), unless otherwise mentioned.

3

No. 83 (AX J0105−722) was once identified with SNR DEM S128 (Yokogawa et al. 2000e), but now
the identification is questionable due to the improved position of the ASCA source and the high resolution
ROSAT study (Filipović, Pietsch, & Haberl 2000). See also section 6.1.19.

46

5.4.4

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES

X-ray pulsars and HMXBs

The spectra of X-ray pulsars (regardless of whether they are accretion-powered or rotationpowered) are generally described by a power-law in the ASCA bandpass. Therefore, we
adopted a power-law model for the 20 detected pulsars (summarized in table 6.1) and also
for No. 63 (a Be/X-ray binary, RX J0058.2−7231).
Several sources exhibited systematic deviation from the simple power-law. Since No. 49(SMC X-2) and No. 90 (XTE J0111.2−7317) showed bump-like residuals around 6–7 keV, we
added a narrow Gaussian line to the model. For No. 67 (RX J0059.2−7138), the power-law> 7 keV, thus we included a high-energy cutoff in the model.
model exceeded the data at ∼The brightest pulsars (RX J0059.2−7138, XTE J0111.2−7317, and SMC X-1) all exhibited< 2 keV, thus we added a blackbody component tolarge data excess over the power-law at ∼describe the soft excess. Details of the analyses and comments are given for each source insection 6.1.

5.4.5

Remaining sources

Although the nature of X-ray emission from the remaining sources is unclear at this moment,
we basically adopted a power-law model in the spectral fitting. Since No. 22 (AX J0048.2−7309)
showed weak evidence for an emission line at around 6–7 keV, we added a Gaussian line tothe model (see section 6.5.2). For No. 2 and No. 13, no constraint on the spectral parameters
could be obtained due to highly limited statistics, thus we do not present the parametersin table 5.5. For No. 39, the best-fit model (Γ = 10 and NH = 1.4 × 1023 cm−2 ) yielded avery high luminosity of LX ∼ 3 × 1039 erg s−1 . Such a high luminosity is unrealistic andis probably an artifact caused by the large Γ and NH , thus we do not present LX in table
5.5. For very soft sources (No. 6 and No. 45), we present the results both from a power-law
model and a CIE thermal model.

5.5. CATALOGUE OF DISCRETE X-RAY SOURCES IN THE SMC REGION

5.5

47

Catalogue of Discrete X-Ray Sources in the SMC
RegionWe present a catalogue of the 106 X-ray sources in the SMC region detected with ASCA. The
sources are presented separately in tables 5.3, 5.4, and 5.5. This catalogue is an upgraded version of that presented by Yokogawa et al. (2000e), including such improvements as enhanced
observation fields, better positional accuracy, and the source flux for each observation.

5.5.1

Catalogue – I. Source identification

Legend for each column of table 5.3 —1
Source number in the catalogue.
2–3 Source position in equatorial coordinates (J2000).
4
Indication whether the source was detected within the inner circle of the GIS (i) or in
the outer ring region (o). The error radius for the source position is ∼ 4000 for “i” sourcesand ∼ 7000 for “o” sources (see section 5.2).
5
IDs for observations in which the source was detected. For sources detected multiple
times, the first ID in this column was used to determine its position.
6
IDs for observations in which the source was not detected, although it was covered by
the FOV of the GIS.
∗: In these observations, nearby sources No. 32 and No. 94 were too bright, thus detectionof No. 34 and No. 95 was hampered.
7–9 Identification with sources in other X-ray catalogues, including ROSAT PSPC (Haberl etal. 2000), ROSAT HRI (Sasaki, Haberl, & Pietsch 2000), and Einstein IPC (Wang & Wu
1992). Parentheses indicate the separation angle from the ASCA source in arcseconds.
†: The separation angle from the central point of No. 413 and 419 in Haberl et al. (2000).
10
Source class defined in section 7.1 and section 7.2.3.
TS: Thermal SNRs. TSc is a candidate for TS.
RS: Radio SNRs which are not classified as TS.
BP: X-ray binary pulsars. BPc is a candidate for BP.
P: Pulsars which are not classified as BP.
NH: Non-pulsating HMXBs.
FS: Sources coincident with a foreground star.
AGN: Sources coincident with a background AGN.
UN: Unclassified sources.

48

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES
Table 5.3: Final ASCA catalogue of discrete X-ray sources in the SMC region– I. Source identification.
No.
1
2
3
4
5
6
7
8
9
10
11
12
13
14
15
16
17
18
19
20
21
22
23
24
25
26
27
28
29
30
31
32
33
34
35
36
37
38
39
40
41
42
43
44
45
46
47

Coordinates (J2000)
R.A.
Dec.
00h 33m 44.s 3 −73◦ 210 4100
00h 34m 05.s 4 −73◦ 310 3000
00h 34m 19.s 4 −73◦ 330 5700
00h 35m 18.s 3 −73◦ 320 4200
00h 35m 28.s 0 −72◦ 120 4300
00h 36m 09.s 2 −72◦ 210 0500
00h 38m 09.s 9 −73◦ 270 5400
00h 38m 56.s 5 −72◦ 040 5300
00h 39m 03.s 1 −72◦ 070 4900
00h 39m 27.s 3 −73◦ 410 5800
00h 39m 52.s 0 −73◦ 410 3600
00h 40m 46.s 4 −73◦ 360 5800
00h 41m 29.s 9 −73◦ 360 3700
00h 41m 37.s 0 −73◦ 260 4700
00h 41m 57.s 9 −73◦ 430 2200
00h 42m 04.s 8 −73◦ 440 5800
00h 42m 39.s 9 −73◦ 400 2500
00h 44m 06.s 6 −73◦ 370 0300
00h 45m 25.s 6 −73◦ 530 5500
00h 47m 22.s 9 −73◦ 120 0600
00h 47m 30.s 0 −73◦ 080 2500
00h 48m 14.s 0 −73◦ 090 3900
00h 48m 37.s 6 −73◦ 180 5900
00h 49m 01.s 6 −72◦ 510 4600
00h 49m 08.s 3 −73◦ 130 2100
00h 49m 18.s 5 −73◦ 120 0100
00h 49m 26.s 3 −73◦ 100 5100
00h 49m 32.s 0 −73◦ 300 170000h 49m 35.s 2 −73◦ 020 4700
00h 49m 42.s 8 −73◦ 220 4000
00h 50m 28.s 8 −72◦ 580 3500
00h 50m 40.s 5 −73◦ 150 4600
00h 50m 45.s 4 −72◦ 410 5700
00h 50m 50.s 7 −73◦ 170 160000h 50m 53.s 5 −73◦ 100 0700
00h 50m 58.s 1 −73◦ 200 5500
00h 51m 06.s 4 −72◦ 130 5800
00h 51m 25.s 4 −72◦ 270 2900
00h 51m 39.s 9 −73◦ 020 5800
00h 51m 44.s 5 −73◦ 100 3400
00h 51m 55.s 7 −72◦ 560 4900
00h 51m 56.s 2 −73◦ 190 3900
00h 52m 22.s 1 −73◦ 190 0500
00h 52m 58.s 4 −71◦ 570 5400
00h 53m 29.s 2 −73◦ 360 4000
00h 53m 42.s 1 −72◦ 520 120000h 53m 55.s 2 −72◦ 260 4200

i/o
i
i
i
i
i
i
i
i
i
i
i
i
i
i
i
i
i
i
o
i
i
i
i
o
i
i
i
o
i
i
i
i
i
o
i
i
i
i
i
i
i
i
o
i
i
i
i

Obs. ID—— ROSAT ——
Det.
N.D.
PSPC (sep00 ) HRI (sep00 )S
462 (14.0)
SSS
518 (42.6)
E
167 (30.6)
2 (30.7)E
211 (33.9)
4 (34.4)
PKE
127 (47.4)
9 (44.8)
E
P
K
K
P
P
K
P
K
P
K
481 (30.2)
KPP
K
11 (15.8)
KP
546 (16.7)
12 (17.8)
P
K
P
K
QF
434 (17.6)
QF
413(† 36.8)
QF
QF
454 (60.4)
FQH
351 (54.1)
QF
437 (44.8)
FQ
430 (33.2)
QF
427 (49.6)
LFQ
511 (67.6)
28 (52.4)
QF
QFL
468 (22.2)
Q
F
QF
L
444 (20.1)
34 (24.0)
H
35 (17.9)
L
F∗ Q∗
Q
F
421 (13.2)
36 (16.1)
QFL
461 (50.4)
38 (40.2)
H
37 (45.5)
H
QF
QFL
424 (30.6)
41 (32.6)
Q
Q
L
F
LQ
453 (36.1)
44 (37.5)
T
94 (19.3)
46 (19.2)
LR
J
H
242 (20.9)
51 (10.7)(Continuing to the next page.)

Einstein
IPC (sep00 )18 (55.3)
16 (37.9)
20 (46.5)
21 (49.4)
22 (46.7)

24 (43.0)

25 (31.4)
29 (46.5)
29 (70.9)
32 (47.5)

34 (15.0)

Class
UN
UN
UN
UN
UN
TSc
UN
UN
UN
UN
UN
UN
UN
BPc
UN
BPc
P
UN
UN
NHc,BPcTS
NHc,BPc
RS,TScBP
RS,TScP
UN
NHc,BPc
BPc
BP
UN
BP
FS
UN
NHc,BPc
TS
BP
BPc
BPc
BP
UN
UN
BP
BP
TSc
UN
BP

5.5. CATALOGUE OF DISCRETE X-RAY SOURCES IN THE SMC REGION
Table 5.3: Final ASCA catalogue of discrete X-ray sources in the SMC region– I. Source identification (continued from the previous page).
No.
48
49
50
51
52
53
54
55
56
57
58
59
60
61
62
63
64
65
66
67
68
69
70
71
72
73
74
75
76
77
78
79
80
81
82
83
84
85
86
87
88
89
90
91
92
93
94

Coordinates (J2000)
R.A.
Dec.
00h 54m 25.s 1 −73◦ 300 3200
00h 54m 36.s 2 −73◦ 400 3500
00h 54m 49.s 4 −72◦ 440 2300
00h 55m 01.s 6 −72◦ 250 4900
00h 55m 36.s 1 −72◦ 100 5600
00h 55m 51.s 6 −73◦ 300 3200
00h 55m 55.s 5 −72◦ 520 1500
00h 57m 27.s 2 −72◦ 250 0100
00h 57m 28.s 6 −73◦ 250 2900
00h 57m 28.s 7 −72◦ 500 170000h 57m 29.s 3 −72◦ 580 0800
00h 57m 37.s 4 −72◦ 190 0700
00h 57m 54.s 2 −71◦ 180 0400
00h 57m 54.s 7 −72◦ 020 2600
00h 57m 57.s 3 −73◦ 080 5500
00h 58m 15.s 8 −72◦ 300 2600
00h 58m 19.s 5 −72◦ 170 4000
00h 59m 24.s 2 −72◦ 220 3500
00h 59m 27.s 2 −72◦ 100 0300
00h 59m 35.s 3 −71◦ 380 0200
01h 00m 09.s 0 −71◦ 170 3700
01h 00m 09.s 6 −72◦ 570 3600
01h 00m 12.s 1 −71◦ 100 1500
01h 00m 32.s 8 −73◦ 400 1200
01h 00m 41.s 1 −72◦ 110 1600
01h 00m 58.s 1 −71◦ 200 0400
01h 01m 19.s 6 −72◦ 100 5900
01h 01m 27.s 8 −73◦ 350 0000
01h 01m 53.s 0 −72◦ 230 0100
01h 02m 14.s 6 −71◦ 170 2000
01h 03m 15.s 3 −72◦ 090 1900
01h 03m 27.s 3 −73◦ 010 2400
01h 03m 27.s 5 −72◦ 460 5500
01h 04m 01.s 7 −72◦ 010 4000
01h 05m 01.s 2 −72◦ 230 0600
01h 05m 07.s 6 −72◦ 100 3400
01h 05m 30.s 1 −72◦ 120 4700
01h 05m 52.s 1 −72◦ 030 4100
01h 06m 48.s 5 −72◦ 240 3200
01h 07m 13.s 5 −72◦ 340 3900
01h 07m 27.s 2 −72◦ 430 2600
01h 07m 39.s 5 −72◦ 420 1500
01h 11m 10.s 1 −73◦ 160 3200
01h 11m 45.s 8 −72◦ 490 5200
01h 12m 38.s 7 −72◦ 360 5100
01h 13m 05.s 5 −72◦ 460 2500
01h 17m 08.s 5 −73◦ 260 0700

i/o
i
i
i
o
i
i
i
i
i
i
i
i
i
i
i
i
i
i
i
o
i
i
i
i
i
i
i
i
i
i
i
i
i
i
i
i
i
i
i
o
o
o
i
i
i
i
i

Obs. ID—— ROSAT ——
Det.
N.D.
PSPC (sep00 ) HRI (sep00 )L
R
R
L
547 (30.5)
J
324 (50.4)
57 (52.1)
G
241 (63.5)
58 (67.1)
GT
157 (40.9)
59 (34.0)LR
508 (32.8)
65 (34.5)
J
66 (31.1)G
234 (47.6)
RL
476 (28.1)
70 (13.6)
JJGU
14 (28.3)
GT
114 (33.3)
73 (29.8)
JG
258 (32.7)
76 (28.0)
G
194 (28.8)
77 (28.9)
G
218 (44.0)
81 (46.8)
G
148 (10.6)
82 (4.8)
B
53(121.5)
U
13 (39.4)
M
84 (10.1)
URG
162 (18.5)
90 (18.1)
U
18 (21.2)
GDB
159 (22.0)
95 (18.6)
R
DG
220 (26.9)
97 (34.3)
U
DBG
143 (6.7)
101 (9.4)
M
384 (8.8)M
334 (34.4)
104 (32.8)
BD
107 (7.9)
107 (23.7)
D
217 (14.3)
109 (10.8)
DB
DB
172 (45.5)
112 (45.4)
DB
120 (17.9)
D
230 (41.4)
116 (44.2)
D
279 (58.3)
N
313 (9.1)
N
307 (71.6)
IA
446 (40.7)N
348 (42.2)
N
283 (55.2)
N
330 (25.8)
ACI
482 (28.4)
118 (33.6)(Continuing to the next page.)

Einstein
IPC (sep00 )35 (43.0)
36 (19.6)

39 (51.3)

41 (50.1)

42 (17.4)
44 (10.5)

45 (38.4)

50 (22.7)

51 (18.7)
52 (42.0)55 (65.7)
56 (46.0)

63 (13.4)

Class
UN
BP
NHc,BPc
BP
AGN
UN
UN
UN
BP
UN
UN
UN
UN
BP
BPc
NH
RS,BPc
BPc
TS
BP
UN
UN
BPc
UN
AGN
BPc
BP
BPc
NHc
UN
BP
UN
AGN
TS
TS
RS,P
AGN
NHc
UN
NHc
UN
UN
BP
AGN
AGN
BPc
BP

49

50

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES
Table 5.3: Final ASCA catalogue of discrete X-ray sources in the SMC region– I. Source identification (continued from the previous page).
No.
95
96
97
98
99
100
101
102
103
104
105
106

Coordinates (J2000)
R.A.
Dec.
01h 18m 37.s 9 −73◦ 250 2200
01h 19m 27.s 3 −73◦ 000 5100
01h 20m 01.s 0 −72◦ 510 4600
01h 20m 44.s 5 −72◦ 550 1300
01h 21m 16.s 2 −72◦ 270 4200
01h 21m 23.s 9 −72◦ 440 5300
01h 23m 44.s 6 −72◦ 340 2600
01h 24m 15.s 7 −72◦ 420 3200
01h 26m 44.s 2 −73◦ 040 2000
01h 27m 48.s 0 −73◦ 070 3900
01h 28m 28.s 6 −73◦ 290 4600
01h 31m 12.s 0 −73◦ 130 4400

i/o
i
i
i
i
i
i
i
i
o
i
i
i

Obs. ID
Det.
N.D.
C
A∗
V
V
V
V
V
V
V
O
O
O
O

—— ROSAT ——PSPC (sep00 ) HRI (sep00 )478 (4.0)
120 (5.0)385 (38.6)

Einstein
IPC (sep00 )65 (13.0)
66 (19.6)

316 (54.6)
275 (43.7)

67 (35.3)

Class
FS
UN
BPc
BPc
UN
UN
UN
UN
UN
BPc
UN
UN

5.5. CATALOGUE OF DISCRETE X-RAY SOURCES IN THE SMC REGION

5.5.2

Catalogue – II. Names and comments

Legend for each column of table 5.4 —1 Source number in the catalogue.
2 Source names usually used in the literature. References for the abbreviations are Davies,
Elliot, & Meaburn (1976) for DEM’s, Inoue, Koyama, & Tanaka (1983) for IKT’s, Henize(1956) for N’s, and Mathewson et al. (1983, 1984) for SNR’s.
3 Comments such as periods and companion stars of pulsars, and spectral information for
peculiar sources. “Be/X?” indicates that the ROSAT counterpart has an emission line
object in the error circle (Haberl & Sasaki 2000; here HS00). Excerpts of comments in
ROSAT catalogues (Haberl et al. 2000 and Sasaki, Haberl, & Pietsch 2000; here HFP00
and SHP00, respectively) are also included.

51

52

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES
Table 5.4: Final ASCA catalogue of discrete X-ray sources in the SMC region– II. Names and comments.

No.
1
2
3
4
5
6
7
8
9
10
11
12
13
14
15
16
17
18
19
20
21
22
23
24
25
26
27
28
29
30
31
32
33
34
35
36
37
38
39
40
41
42
43
44
45
46
47
48

Names

Comments

fg star (HFP00); AGN? (SHP00)

AX J0043−737

P = 87ms? (obs. K); XB? or AGN? (SHP00)

RX J0047.3−7312; IKT1
SNR 0045−734; N19; DEM S32; IKT2
AX J0048.2−7309
SNR 0046−735; N22; DEM S37; IKT4
AX J0049−729; RX J0049.1−7250
SNR 0047−735; DEM S49; IKT5
AX J0049−732; RX J0049.0−7314
RX J0049.4−7310
RX J0049.5−7331

Be/X? (HS00); flare (obs. Q)
old, ejecta-dominated, center-filled
Be/X? (section 6.5.2); Fe line? (obs. Q; EW=240eV)
P = 74s; Be/X

AX J0049.5−7323; RX J0049.7−7323

P = 755s; Be/X? (HS00)

AX J0051−733; RX J0050.8−7316

P = 323s; Be/X; flux increased gradually (obs. Q)
fg star SkKM 62 (SHP00)

RX J0050.9−7310
SNR 0049−736; IKT6
AX J0051−722

Be/X? (HS00)
P = 91s; Be/X

AX J0051.6−7311; RX J0051.9−7311; IKT7

P = 172s; Be/X

RX J0052.1−7319; IKT8XTE J0054−720; AX J0052.9−7157; RX J0052.9−7158; IKT9

P = 15s (not detected in this work); Be/X
P = 167s; Be/X

1WGA J0053.8−7226; XTE J0053−724; IKT10

P = 46s; Be/X

P = 9.13s
Be/X? (HS00); pos. and spec. are from SIS
Be/X? (HS00)(Continuing to the next page.)

5.5. CATALOGUE OF DISCRETE X-RAY SOURCES IN THE SMC REGION

53

Table 5.4: Final ASCA catalogue of discrete X-ray sources in the SMC region– II. Names and comments (continued from the previous page).
No.
49
50
51
52
53
54
55
56
57
58
59
60
61
62
63
64
65
66
67
68
69
70
71
72
73
74
75
76
77
78
79
80
81
82
83
84
85
86
87
88
89
90
91
92
93
94
95

Names
SMC X-2
RX J0054.9−7245
XTE J0055−724; 1SAX J0054.9−7226; IKT11

Comments
P = 2.37s; Be/X; Fe line (EW=400eV)
Be/X? (HS00)
P = 58s; Be/X
AGN? (HFP00)

IKT14
AX J0057.4−7325

P = 101s; probably an XBPAX J0058−7203

P = 280s; Be/X? (HS00)

RX J0058.2−7231
SNR 0056−725; IKT16
SNR 0057−7226; N66; DEM S103; IKT18
RX J0059.2−7138

Be/X (HS00)
hard spectrum
XB? (SHP00)
old, ISM-dominated, center-filled
P = 2.76s; Be/X; at the very edge of the GIS

IKT19[hard] (HFP00); AGN? (SHP00)

RX J0101.3−7211

P = 455s (not detected in this work); Be/X

RX J0101.8−7223

Be/X? (HS00)

1SAX J0103.2−7209; SNR 0101−724; IKT21

P = 348s; Be/X
AGN? (SHP00)

SNR 0102−723; N76; DEM S124; IKT22
SNR 0103−726; DEM S125; IKT23
AX J0105−722

RX J0105.9−7203
DEM S134
RX J0107.1−7235

XTE J0111.2−7317

SMC X-1

old, ejecta-dominated, center-filled + faint shellP = 3.34s?; resolved into multiple sources by
ROSAT[nonstar] (HFP00); AGN? (SHP00)
Be/X? (HS00); severe contamination from No. 81[hard] (HFP00)
Be/X? (HS00)[nonstar] (HFP00); contamination from No. 89
contamination from No. 88
P = 30s; Be/X; Fe line (EW=50eV)AGN z = 0.197 (HFP00)
AGN z = 1.376 (HFP00)
P = 0.70s; supergiant HMXB (B0I, Sk160)
fg star G5V, HD 8191
(Continuing to the next page.)

54

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES
Table 5.4: Final ASCA catalogue of discrete X-ray sources in the SMC region– II. Names and comments (continued from the previous page).

No.
96
97
98
99
100
101
102
103
104
105
106

Names

AX J0128.4−7329

Comments
cluster z = 0.0656 (HFP00)

in stray lightin stray lightat the center of a supergiant shell, SMC-1

5.5. CATALOGUE OF DISCRETE X-RAY SOURCES IN THE SMC REGION

5.5.3

55

Catalogue – III. Spectral parameters

Legend for each column of table 5.5 —1
Source number in the catalogue.
2
Spectral hardness ratio (HR) defined as HR ≡ (H −S)/(H +S), where H and S represent
background-subtracted GIS count rates in 2.0–7.0 keV and 0.7–2.0 keV, respectively. For
sources detected multiple times, spectra from all observations were added to derive theHR (except No. 94).
3
Error in the HR.
4
Observation ID. Flux and luminosity are given for each ID (in columns 8–10).
5–7
Photon index, temperature (keV), and column density (1021 cm−2 ) derived in the spectral
fitting (section 5.4). Parentheses indicate 90% confidence limits. For sources showing softexcess (Nos. 67, 90, and 94), a two-component model consisting of a power-law plus a
blackbody was used; thus the temperature in the table is for the blackbody component.
For No. 67, a power-law with a high-energy exponential cutoff was used. The cutoff ande-fold energies, Ecut and Efold , are presented in table 6.16.
8–10 Flux (10−12 erg s−1 cm−2 ), observed luminosity (1036 erg s−1 ), and absorption-corrected
luminosity (1036 erg s−1 ). Observed luminosity is defined as
Lobs = FX × 4πd2 ,where d = 60 kpc is the distance to the SMC. We also adopted d = 60 kpc for sources
coincident with background AGNs or foreground stars.
∗: The best-fit model yielded an unusually high luminosity of ∼ 3 × 1039 erg s−1 . Thisis probably an artifact due to the large values of Γ and NH , thus we omitted it from thetable.
†: The intrinsic intensity would be much higher (see section 6.1.21).

56

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES
Table 5.5: Final ASCA catalogue of discrete X-ray sources in the SMC region– III. Spectral parameters.

No.

HR

HRerr

1
2
3
4
5
6
6
7
8
9
10
11
12
13
14
15
16
17
17
18
19
20
20
21
21
22
22
23
24
24
24
25
25
26
26
27
28
28
28
29
30
30
30
31
32
32
33

−0.27
−1.00
0.64
−0.13
−0.27
−0.79
−0.79
0.02
0.16
−0.16
−0.02
−0.15
−0.09
−1.00
0.22
0.86
0.43
−0.16

0.14
2.00
0.09
0.18
0.24
0.35
0.35
0.12
0.23
0.25
0.20
0.16
0.24
1.15
0.23
0.25
0.10
0.15

0.15
0.13
0.39

0.10
0.15
0.03

−0.86

0.06

0.59

0.05

−0.62
0.36

0.36
0.03

−1.00

1.38

0.85

0.20

...
0.58

...
0.12

0.25
0.31

0.18
0.06

0.08
0.43

0.31
0.01

−0.16

0.16

Obs.
ID
S
S
S
S
E
E
E
P
E
E
P
K
P
P
P
K
P
K
P
P
P
F
Q
F
Q
F
Q
Q
F
H
Q
F
Q
F
Q
Q
F
L
Q
Q
F
L
Q
Q
F
Q
H

Γ

kT

NH

FX

Lobs

LX

1.4 (1.0-2.0)
1.8 (0.7-4.2)
2.2 (1.3-5.2)
1.2 (0.5-5.1)
2.0 (1.4-3.5)
2.1 (1.3-4.4)...
2.1 (0.8-4.2)
3.9 (2.2-7.0)
1.8 (0.9-2.9)
1.7 (1.3-2.4)

2 (<15)...
47 (17-100)
0 (<7)
3 (<22)2 (<37)0 (<20)0 (<0.5)
7 (<46)0 (<20)0 (<69)0 (<6)
0 (<31)...
9 (<31)57 (23-120)
18 (0.7-39)
0 (<4)

2.2 (1.3-3.6)
1.5 (0.9-2.5)
1.0 (0.9-1.2)

11 (0.6-28)
0 (<6)
4 (1-7)

0.53
...
0.72
0.21
0.28
0.25
0.23
0.52
0.19
0.12
0.20
0.24
0.066
...
0.13
0.16
0.31
0.20
0.11
0.27
0.48
0.67
1.4
0.29
0.29
0.20
1.1
0.058
12
1.6
0.71
0.14
0.11
0.64
0.37
0.23
0.82
1.1
0.84
0.10
1.1
1.5
1.7
0.077
2.9
6.8
0.39

0.23
...
0.31
0.091
0.12
0.11
0.099
0.22
0.082
0.052
0.086
0.10
0.029
...
0.056
0.069
0.13
0.086
0.048
0.12
0.21
0.29
0.60
0.13
0.13
0.086
0.48
0.025
5.2
0.69
0.31
0.060
0.048
0.28
0.16
0.099
0.35
0.48
0.36
0.043
0.48
0.65
0.73
0.033
1.3
2.9
0.17

0.29
...
0.65
0.091
0.15
0.16
0.099
0.22
0.11
0.052
0.086
0.10
0.029
...
0.086
1.2
0.22
0.086
0.048
0.20
0.21
0.32
0.65
1.3
1.3
0.099
0.52
0.025
5.5
0.76
0.33
0.060
0.048
0.33
0.19
0.17
0.39
0.51
0.40
0.043
0.48
0.65
0.73
0.033
1.3
3.0
0.20

2.5 (1.7-4.6)...
1.8 (1.0-3.0)
1.9 (1.2-3.7)
2.1 (1.0-5.3)
3.6 (>2.2)
1.1 (0.1-2.5)

0.37 (0.32-0.47)
0.8 (0.6-1.1)

10 (8-11)
9 (3-17)0.9 (0.1-1.8)

0 (<15)3 (2-5)0.86 (0.82-0.90)

0 (<0.4)

1.0 (0.8-1.1)

0.9 (0.3-1.6)

15 (6-28)

1.7 (0.8-2.9)
0.9 (0.3-1.7)

27 (11-58)
6 (<20)1.0 (0.3-1.7)
0.8 (0.7-1.0)

0 (<19)0 (<3)

1.3 (−0.1-4.5)
0.76 (0.69-0.83)

0 (<30)2.2 (1.2-3.3)

2.2 (1.5-3.6)

2 (<11)(Continuing to the next page.)

5.5. CATALOGUE OF DISCRETE X-RAY SOURCES IN THE SMC REGION

57

Table 5.5: Final ASCA catalogue of discrete X-ray sources in the SMC region– III. Spectral parameters (continued from the previous page).
No.

HR

HRerr

34
35
36
36
36
37
38
39
40
40
40
41
42
43
44
45
45
46
47
48
49
50
51
52
52
53
53
54
55
56
56
57
58
59
60
61
61
62
63
64
65
66
67
68
69
70
71

0.09
0.42
−0.92

0.20
0.04
0.07

0.55
0.58
0.53
0.33

0.04
0.20
0.38
0.04

0.05
0.67
0.83
0.42
−0.97
−0.97
−0.06
0.37
−0.06
0.36
0.20
0.51
−0.05

0.32
0.29
0.36
0.05
0.51
0.51
0.18
0.02
0.27
0.03
0.11
0.10
0.09

−0.01

0.11

0.11
−0.06
0.48

0.17
0.31
0.05

0.17
−0.02
0.15
−0.09
0.36

0.14
0.17
0.12
0.18
0.04

0.24
1.00
0.20
0.25
−0.68
0.34
0.07
1.00
0.44
−0.09

0.20
0.29
0.12
0.09
0.10
0.01
0.20
0.77
0.29
0.10

Obs.
ID
L
Q
F
L
Q
H
H
Q
F
L
Q
Q
Q
F
T
L
L
J
H
L
R
J
G
G
T
L
R
J
G
L
R
J
J
G
U
G
T
J
G
G
G
G
B
U
M
U
R

Γ

kT

NH

0.6 (−0.4-1.6)
1.0 (0.8-1.3)

0 (<11)4 (0.9-8)
6 (2-7)0.23 (0.18-0.47)

1.0 (0.7-1.3)
0.1 (−1.0-1.3)
10 (>2.7)
0.9 (0.8-1.0)

10 (5-16)
0 (<21)140 (95-200)
0 (<0.8)

1.8 (0.5-5.5)
3.3 (1.2-7.1)
2.9 (...)
0.7 (0.5-1.0)
5.1 (>2.4)
4.0 (1.8-8.9)
0.8 (0.7-0.9)
1.3 (0.2-4.1)
0.7 (0.5-0.9)
1.8 (1.1-2.8)
0.4 (0.1-1.0)
1.5 (1.2-1.8)

0 (<43)6 (1-16)53 (<290)
0.7 (<5)
9 (<45)10 (<31)27 (0.7-92)
1 (<2)
0 (<14)0 (<1)
6 (<19)0.02 (<10)
0 (<2)

2.1 (1.6-2.7)

0 (<2)

2.4 (1.3-4.1)
5.3 (>2.1)
1.1 (0.9-1.4)

11 (<36)
35 (5-110)
8 (4-14)

1.4 (0.9-2.4)
1.7 (1.2-3.2)
1.6 (0.9-2.6)
2.2 (1.5-3.5)
0.8 (0.6-1.0)

2 (<13)0 (<17)4 (<15)0 (<10)0.7 (<4)

1.8 (0.7-4.1)
2.6 (0.2-6.0)
1.5 (1.0-2.3)
1.0 (0.7-1.5)

7 (<37)89 (29-230)
2 (<12)0 (<5)
1 (<3)
0 (<1)
0.02 (<12)
28 (<910)
0 (<64)0 (<3)

0.6 (0.2-2.1)

0.9 (0.6-1.7)
0.51 (0.46-0.56)
0.15 (0.14-0.17)
1.7 (0.9-3.7)
1.4 (>−0.6)
0.8 (−0.1-3.4)
1.6 (1.3-2.1)(Continuing to the next page.)

FX

Lobs

LX

1.3
0.92
0.64
0.65
0.74
9.5
0.58
0.038
0.68
1.5
1.3
0.10
0.19
0.68
2.0
0.093
0.096
0.21
12
0.12
9.1
0.49
2.7
0.44
1.4
0.21
0.24
0.18
0.13
1.1
2.2
0.47
0.35
0.42
0.20
2.5
4.4
0.62
0.70
0.34
0.75
0.30
240
0.13
0.42
0.15
0.84

0.56
0.40
0.28
0.28
0.32
4.1
0.25
0.016
0.29
0.65
0.56
0.043
0.082
0.29
0.86
0.040
0.041
0.091
5.2
0.052
3.9
0.21
1.2
0.19
0.60
0.091
0.10
0.078
0.056
0.48
0.95
0.20
0.15
0.18
0.086
1.1
1.9
0.27
0.30
0.15
0.32
0.13
100
0.056
0.18
0.065
0.36

0.56
0.43
1.8
1.8
2.0
4.8
0.25
...∗
0.29
0.65
0.56
0.043
0.78
1.6
0.88
0.29
0.41
0.86
5.2
0.052
3.9
0.28
1.2
0.19
0.60
0.091
0.10
0.15
2.6
0.56
1.1
0.22
0.15
0.21
0.086
1.1
1.9
0.36
1.4
0.16
0.32
0.17
100
0.056
0.27
0.065
0.36

58

CHAPTER 5. SYSTEMATIC ANALYSES ON DISCRETE SOURCES
Table 5.5: Final ASCA catalogue of discrete X-ray sources in the SMC region– III. Spectral parameters (continued from the previous page).

No.

HR

72
73
74
74
74
75
76
76
77
78
78
78
79
80
81
82
83
84
85
86
87
88
89
90
91
92
93
94
94
94
95
96
97
98
99
100
101
102
103
104
105
106

−0.37
0.32
0.43

0.34
0.19
0.14
0.36

0.66
−0.37
−0.86
−0.88
−0.20
−0.28
...
0.16
0.68
...
...
0.33
0.31
−0.08
0.36
0.43
0.40
0.44
−0.43
−0.36
0.24
0.41
1.00
0.11
−0.22
0.63
0.18
0.20
−0.06
−0.21

HRerr

Obs.
ID
0.12
G
0.11
U
0.05
B
D
G
0.32
R
0.13
D
G
0.39
U
0.06
B
D
G
0.11
M
0.13
M
0.01
B
0.04
D
0.08
D
0.12
D
...
D
0.11
D
0.35
D
...
N
...
N
0.003
I
0.26
N
0.21
N
0.23
N
0.01
A
0.02
C
0.01
I
0.09
C
0.13
V
0.18
V
0.20
V
0.74
V
0.07
V
0.18
V
0.17
V
0.11
O
0.10
O
0.12
O
0.31
O

Γ

kT

NH

3.2 (2.4-4.7)
1.4 (0.9-2.1)
0.7 (0.5-0.9)

3 (<11)2 (<11)0.4 (<4)

0.4 (−1.1-2.5)
1.9 (1.2-2.7)

2 (<62)9 (2-19)

1.4 (0.0-5.4)
0.8 (0.6-1.0)

3 (<740)0 (<4)

2.4 (0.5-4.4)
2.6 (2.0-4.4)

2.2 (1.9-2.6)
2.6 (2.1-3.5)...
1.3 (0.9-2.0)
4.2 (1.5-9.7)...
...
0.79 (0.76-0.82)
2.0 (>1.0)
1.4 (0.6-2.7)
1.3 (0.6-3.1)
0.69 (0.61-0.76)
0.45 (0.35-0.53)
0.35 (0.27-0.44)
3.3 (2.8-3.9)
2.6 (2.0-4.0)
0.6 (−0.1-1.2)
1.2 (0.5-2.5)
5.2 (>1.9)
2.4 (1.8-3.1)
2.1 (1.3-4.3)
1.9 (0.8-3.6)
1.2 (0.9-1.6)
1.1 (0.8-1.8)
2.8 (1.7-4.8)
6.6 (>1.8)

110 (13-210)
0.4 (<9)— see section 6.3.4 —
1.0 (0.8-1.5)
0 (<0.3)
0 (<2)
0 (<3)...
0.3 (<8)
51 (12-150)...
...
0.144 (0.137-0.150)
2.3 (1.6-3.0)
2 (<160)
0 (<8)
0 (<20)0.13 (0.11-0.15)
4 (2-6)0.20 (0.13-0.29)
0 (<2)
0.16 (0.13-0.19)
2 (<4)
0 (<1)
1 (<10)0 (<12)3 (<21)110 (29-290)
8 (3-16)
1 (<14)24 (5-59)
0 (<4)
0 (<8)
7 (<20)29 (<88)

FX

Lobs

LX

0.36
0.37
1.5
1.6
1.7
0.52
0.46
0.36
0.12
1.3
1.1
0.79
0.67
0.59
14
1.0
0.30
0.16
...
0.38
0.19
...
...
370
0.11
0.23
0.17
220
8.9†
150
0.38
0.68
0.35
0.37
0.44
0.41
0.45
0.99
1.8
1.7
0.34
0.14

0.16
0.16
0.65
0.69
0.73
0.22
0.20
0.16
0.052
0.56
0.48
0.34
0.29
0.25
6.0
0.43
0.13
0.069
...
0.16
0.082
...
...
160
0.048
0.099
0.073
95
3.8†
65
0.16
0.29
0.15
0.16
0.19
0.18
0.19
0.43
0.78
0.73
0.15
0.060

0.23
0.17
0.65
0.69
0.73
0.23
0.28
0.22
0.056
0.56
0.48
0.34
1.3
0.27
7.3
0.43
0.13
0.069
...
0.16
1.8
...
...
170
0.052
0.099
0.073
110
3.8†
67
0.16
0.34
0.15
0.18
31
0.30
0.21
0.78
0.78
0.73
0.29
6.5

Chapter 6
Comments on Specific Sources andPulsar Statistics
6.1

X-Ray Pulsars

6.1.1

Overview

Figure 6.1 shows the history of discoveries of X-ray pulsars in the SMC. As is easily recognized, there have recently been a rush of pulsar discoveries, and the total number has alreadyexceeded 20. The contribution from this study to the discoveries and position determinationof pulsars (stars and circles in figure 6.1) is rather large, thanks to the high sensitivity in thehigh energy band, fine time resolution, and imaging capability of ASCA.
Here we give brief comments on all the pulsars, including their spectra, pulse profiles,
optical counterparts, and flux variability. Pulse profiles are shown in 0.7–2.0 keV and 2.0–
7.0 keV for most sources, while for the remainder different energy bands are chosen for
purposes peculiar to each source. We regard a pulsar as an XBP if the pulsar has a long pulse
period (∼ 1–1000 s), hard spectrum (Γ ∼ 1), flux variability, and/or an optical counterpart.
We ultimately classified at least 22 out of the 25 pulsars as XBPs1 . Basic properties are
summarized in table 6.1. In this section, pulse periods are designated with the error for thelast digit in parentheses. The significance of pulse detection with FFT analysis is defined by1

AX J0043−737, AX J0049−732, and AX J0105−722 are not included in the 22 because further confirmation is needed for them.

59

60 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Figure 6.1: History of discoveries of X-ray pulsars in the SMC as of December, 2001. The symbols
represent the pulsars discovered in this study (stars), the pulsars for which only the positions were
determined in this study (circles), and the remainder (triangles).

1 − PRND , where PRND is the probability that the maximum power is attributed to randomly
distributed events. For the maximum power Pmax obtained in a PSD, PRND is calculated by
PRND = NZ ∞Pmax

p2 (χ2 )dχ2 ,(6.1)

where N and p2 (χ2 ) represent the number of frequencies in the PSD and the χ2 probability
density for 2 degrees of freedom, respectively (e.g., Leahy et al. 1983).

6.1.2

No. 17 — AX J0043−737

Coherent pulsations of a 87.58073(4) ms period from No. 17 (AX J0043−737) were discoveredby Yokogawa & Koyama (2000) in obs. K, using a timing resolution of 31.25 ms (1/16 of thenominal value) as described in section 5.3. The PSD derived from the events in 0.5–7.7 keV
is shown in figure 6.2 (a). The maximum power of 47.3 indicates that the significance of thepulse detection is at a marginal level, ∼ 99.98%. The pulse profile in the same energy band

6.1. X-RAY PULSARS

61

Table 6.1: X-ray pulsars in the SMC.
No.

Name

17
24
26
30
...
...
32

AX J0043−737†
AX J0049−729
AX J0049−732
AX J0049.5−7323
XTE J0050−732#1
XTE J0050−732#2
AX J0051−733

Period (s)∗OpticalIdentification
...
Be
...‡
Be?
...
...
Be

References

0.08758073(4)174.68(2)2, 3, 4, 5
9.1320(4)6, 7, 8755.5(6)9, 10, 11[16.6]
12[25.5]
12323.2(5) - obs. F
5, 13, 14, 15
321.0(1) - obs. Q
37
AX J0051−722
91.12(5)Be
2, 4, 1640
AX J0051.6−7311
172.40(3)Be
15, 17, 18...
2E 0050.1−7247[8.8816(2)]
Be
4, 19
...
XTE J0052−723[4.782(1)]
...
20
43
RX J0052.1−7319[15.3(1)]Be21, 22, 23
44
XTE J0054−720
167.8(2)Be
15, 24, 25
47
1WGA J0053.8−7226
46.63(4)
Be
2, 26
49
SMC X-2
2.37230(4)Be
27, 28, 29, 30
51
XTE J0055−724[58.963(3)]Be
4, 5, 31, 32, 33
56
AX J0057.4−7325
101.45(7) - obs. R...
28, 34
101.47(6) - obs. L
61
AX J0058−7203
280.4(4)Be?
11, 13, 35
67
RX J0059.2−7138
2.763221(4)§Be
36, 37, 38
74
RX J0101.3−7211[455(2)]Be
39
78
1SAX J0103.2−7209
348.9(3)Be
5, 40, 41, 42
83
AX J0105−722†
3.34300(3)...‡43, 4490
XTE J0111.2−7317
30.9497(4)Be
22, 45, 46, 47
94
SMC X-1
0.70865(6) - obs. A
B (SG)¶
48, 490.7065(5) - obs. I...
RX J0117.6−7330[22.0669(1)]Be
50, 51
Note — Pulsars in this table are mostly regarded as XBPs, except for AX J0043−737, AX J0049−732, andAX J0105−722 (see text).
References — (1) Yokogawa & Koyama 2000; (2) Corbet et al. 1998; (3) Yokogawa et al. 1999; (4) Stevens, Coe, &Buckley 1999; (5) Coe & Orosz 2000; (6) Imanishi, Yokogawa, & Koyama 1998; (7) Ueno et al. 2000a; (8) Filipović,
Pietsch, & Haberl 2000; (9) Ueno et al. 2000b; (10) Yokogawa et al. 2000a; (11) Haberl & Sasaki 2000; (12) Lambet al. 2001; (13) Yokogawa & Koyama 1998a; (14) Imanishi et al. 1999; (15) Cowley et al. 1997; (16) Ozaki etal. 2000; (17) Torii et al. 2000a; (18) Yokogawa et al. 2000b; (19) Israel et al. 1997; (20) Corbet, Marshall, &Markwardt 2001; (21) Lamb et al. 1999; (22) Israel et al. 1999; (23) Kahabka 2000; (24) Lochner et al. 1998; (25)
Yokogawa et al. 2001a; (26) Buckley et al. 2001; (27) Corbet et al. 2001; (28) Torii et al. 2000b; (29) Yokogawa
et al. 2001b; (30) Murdin et al. 1979; (31) Marshall & Lochner 1998; (32) Santangelo et al. 1998; (33) Israel etal. 1998b; (34) Yokogawa et al. 2000c; (35) Tsujimoto et al. 1999; (36) Hughes 1994; (37) Kohno, Yokogawa, &Koyama 2000; (38) Southwell & Charles 1996; (39) Sasaki et al. 2001; (40) Israel et al. 2000; (41) Hughes & Smith
1994; (42) Ye et al. 1995; (43) Yokogawa & Koyama 1998b; (44) Filipović et al. 2000; (45) Chakrabarty et al.
1998a; Wilson & Finger 1998; (46) Yokogawa et al. 2000d; (47) Coe, Haigh, & Reig 2000; (48) Lucke et al. 1976;(49) Petro, Feldman, & Hiltner 1973; (50) Coe et al. 1998; (51) Macomb et al. 1999∗: The periods determined in this work are presented with the error for the last digit in parentheses. Correction
for the orbital Doppler effect was carried out only for SMC X-1. Periods in square brackets were determined byothers and not detected in this work.
†: Care must be taken for these pulsars. See section 6.1.2 and section 6.1.19.
‡: AX J0049−732 and AX J0105−722 were once respectively identified with ROSAT sources RX J0049.4−7310
and RX J0105.1−7211 (Filipović, Pietsch, & Haberl 2000; Filipović et al. 2000). Since these ROSAT sources haveemission line objects as counterparts, it was previously proposed that AX J0049−732 and AX J0105−722 were
Be/X-ray binaries. However, according to the improved positions determined in this study, the identification with
the ROSAT sources is now a little questionable (see table 5.3, section 6.1.4, and section 6.1.19).
§: Determined by a detailed pulse phase analysis (Kohno, Yokogawa, & Koyama 2000).
¶: A supergiant (B0I), Sk160.

62 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

(a)

(b)

Figure 6.2: PSDs of AX J0043−737 in (a) obs. K and (b) obs. P, obtained from the events in
0.5–7.7 keV. Data points with a power less than 10 are not plotted. The peak detected in (a) is
not found in (b).

is shown in figure 6.3. AX J0043−737 was also detected in a follow-up observation witha longer exposure time (obs. P). We performed the FFT on the events in the same energy
band. However, as shown in figure 6.2 (b), no significant peak was found as in the previous
observation. Although events in different energy bands were also used for the FFT analysis,
we again detected no significant peak in the PSDs.
The count rate, total count (without background), and background level in 0.5–7.7 keV
are 2.5 × 10−3 cnt s−1 , 200 cnt, and 64% in obs. K, and 1.0 × 10−3 cnt s−1 , 186 cnt, and
84% in obs. P. Although the total count is nearly identical in these two observations, thesmaller count rate and larger background level in obs. P may have caused the pulsations tobe hidden in the background. Therefore, confirmation of the pulsations by observations witha much higher S/N ratio is still needed.
The spectral shape (Γ and NH ) is consistent between the two observations, while FX
shows a slight decline (table 6.2). The photon index is softer than those for usual XBPs
(Γ ∼ 1; e.g., Nagase 1989). A short pulse period, soft spectrum, and luminosity far smaller
than the Eddington limit for a neutron star are also detected from SAX J0635+0533, which
has a Be star counterpart (P = 33.8 ms and Γ = 1.50; Kaaret et al. 1999; Cusumano et al.
2000). Cusumano et al. (2000) argue that SAX J0635+0533 may be an accretion-powered
Be/X-ray binary, and if so, the magnetic field should be weaker than 2 × 109 G for the6.1. X-RAY PULSARS

63

Figure 6.3: GIS spectrum (left) and pulse profile (right) of AX J0043−737. (Left) The crosses aredata points obtained in obs. K, while the solid line represents the best-fit model derived from thesimultaneous fitting. (Right) Pulse profile in obs. K obtained from 0.5–7.7 keV.

Table 6.2: Spectral parameters for AX J0043−737 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

−0.16

0.15

K
P

1.7 (1.3–2.4)

0 (< 4)

0.20
0.11

0.086
0.048

Note — See page 55 for legends.

accretion to occur against the centrifugal force at the magnetosphere.
Since no optical counterpart has been reported for AX J0043−737, we investigated thecatalogues of emission line objects (potential candidates for Be stars) by Meyssonnier &Azzopardi (1993) and Murphy & Bessell (2000), but no counterpart was found. Therefore,
a search for the optical counterpart is needed, in addition to confirmation of the 87-ms
pulsations.

6.1.3

No. 24 — AX J0049−729

Coherent pulsations with a 74.8(4) s period were first discovered with RXTE (Corbet etal. 1998) in the direction of SMC X-3. However, the very large positional uncertainty (∼64 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Figure 6.4: GIS spectrum (left) and pulse profiles (right) of AX J0049−729. (Left) The crosses aredata points obtained in obs. F, while the solid line represents the best-fit model derived from thesimultaneous fitting. (Right) Pulse profiles in obs. F obtained from 0.7–2.0 keV and 2.0–7.0 keV,and their ratio.

Table 6.3: Spectral parameters for AX J0049−729 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.36

0.03

F
H
Q

1.0 (0.8–1.1)

3 (2–5)

12
1.6
0.71

5.5
0.76
0.33

Note — See page 55 for legends.

2◦ ) hampered identification with other sources. During this study, we found that No. 24(AX J0049−729) was pulsating with a 74.68(2) s period in obs. F (Yokogawa et al. 1999)
and so determined its position much more accurately. The ROSAT counterpart of this pulsar,
RX J0049.1−7250, provides the most accurate position with a ±1300 error circle (Kahabka &6.1. X-RAY PULSARS

65

Pietsch 1998), in which one Be star has been discovered (Stevens, Coe, & Buckley 1999). Thispulsar has been included in 11 observation fields of Einstein, ROSAT, and ASCA. Yokogawa> 100 during those observations. Fromet al. (1999) found a large flux variability of a factor ∼all the information, we conclude that AX J0049−729 is an XBP with a Be star companion(hereafter, a Be-XBP).

6.1.4

No. 26 — AX J0049−732

Coherent pulsations with a 9.1320(4) s period from No. 26 (AX J0049−732) were first detected during this study in the following manner. Using the data of obs. F, we performed an
FFT on the events in 1.0–5.1 keV collected from a 30 -radius circle centered on AX J0049−732,
and discovered the coherent pulsations with 99.99% confidence (Imanishi, Yokogawa, &Koyama 1998; Ueno et al. 2000a). However, since AX J0049−732 is detected at only ∼ 1.0 8
from No. 25 (SNR 0047−735), mutual contamination was not negligible due to the poor
resolution of the ASCA XRT. Therefore, we also performed an FFT on the events in a 30 radius circle centered on SNR 0047−735, and detected no pulsations. These analyses clearly
indicate that the 9.13-s pulsations are attributable to AX J0049−732.
AX J0049−732 was also detected in obs. Q. However, in this observation another source,
No. 27, was detected at ∼ 1.0 3 from AX J0049−732, in addition to SNR 0047−735 at ∼ 1.0 8.
The SIS data were used to resolve these sources and determine the spectral parameters,
but were not used for an FFT analysis because the timing resolution (8 s) is insufficient
to detect the ∼ 9-s pulsations. We therefore used the GIS data for an FFT analysis, although No. 27 and AX J0049−732 were not resolved. We collected photons from a 30 -radius
circle centered on AX J0049−732 and performed the FFT analysis, but no significant pulsations were detected. This is probably due to the poor statistics caused by the reduced
flux of AX J0049−732 and large contamination from No. 27, which has the same flux level
as AX J0049−732 (see table 5.5). Therefore, a follow-up observation with better spatial
resolution is needed to confirm the pulsations from AX J0049−732.
Filipović, Pietsch, & Haberl (2000) found two ROSAT sources near AX J0049−732,
RX J0049.2−7311 and RX J0049.5−7310 (No. 430 and No. 427 in Haberl et al. 2000, respectively). Since RX J0049.2−7311 has an emission line object as a counterpart while
RX J0049.5−7310 does not, Filipović, Pietsch, & Haberl (2000) suggest that RX J0049.2−7311
is likely to be the counterpart for AX J0049−732. However, according to the updated po-

66 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Figure 6.5: GIS spectrum (left) and pulse profile (right) of AX J0049−732. (Left) The crosses aredata points obtained in obs. F, while the solid line represents the best-fit model derived from thesimultaneous fitting. (Right) Pulse profile in obs. F obtained from 1.0–5.1 keV.

Table 6.4: Spectral parameters for AX J0049−732 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.85

0.20

F
Q

0.9 (0.3–1.6)

15 (6–28)

0.64
0.37

0.33
0.19

Note — See page 55 for legends.

sition of the ASCA source, the separation between AX J0049−732 and RX J0049.2−7311
is ∼ 8000 , which is much larger than the improved positional accuracy for ASCA sources(see section 5.2.1). Therefore, the identification made by Filipović, Pietsch, & Haberl (2000)
is now questionable. On the other hand, the separation between RX J0049.5−7310 and
AX J0049−732 is ∼ 3300 , thus we regard this source as the counterpart.

6.1.5

No. 30 — AX J0049.5−7323

Coherent pulsations with a 755.5(6) s period from No. 30 (AX J0049.5−7323) were first
discovered during this study from the data of obs. Q (Ueno et al. 2000b; Yokogawa et al.
2000a). The energy resolved pulse profiles are shown in figure 6.6. The pulse shape is rather
spiky in 1.2–2.0 keV and nearly sinusoidal in 2.0–5.5 keV, while photons above 5.5 keV do6.1. X-RAY PULSARS

67

Figure 6.6: GIS spectrum (left) and pulse profiles (right) of AX J0049.5−7323. (Left) The crosses
are data points obtained in obs. Q, while the solid line represents the best-fit model derived fromthe simultaneous fitting. (Right) Pulse profiles in obs. Q obtained from 1.2–2.0 keV, 2.0–5.5 keV,and 5.5–10.0 keV.

Table 6.5: Spectral parameters for AX J0049.5−7323 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.31

0.06

F
L
Q

0.8 (0.7–1.0)

0 (< 3)

1.1
1.5
1.7

0.48
0.65
0.73

Note — See page 55 for legends.

not exhibit definite pulsations. The smaller pulsed fraction at higher energy is also seen in
RX J0101.3−7211 (section 6.1.17).
AX J0049.5−7323 was also detected in obs. F and obs. L with nearly the same flux level68 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS(see table 6.5), but the pulsations were not detected in those observations. It may be dueto the much better statistics of the data in obs. Q, which has a much longer exposure time(177 ks) than obs. F and obs. L (∼ 40 ks; see table 4.1). Otherwise, it may imply a smaller
pulsed fraction in obs. F and obs. L. To examine the former possibility, we divided the obs.
Q data into four segments with equal exposure times (∼ 44 ks), and performed an FFTanalysis and epoch folding on each segment. Since no significant pulsations were detectedfrom any of the segments, we conclude that the very long exposure time was essential forthe pulse detection.
AX J0049.5−7323 has been covered by 19 observations of ROSAT and Einstein. Yokogawa et al. (2000a) investigated these archival data and found that AX J0049.5−7323 has
shown a flux variation with a factor of at least ∼ 10. All the information indicates thatAX J0049.5−7323 is an XBP. In addition, it has been proposed that the ROSAT counterpart, RX J0049.7−7323, is a Be/X-ray binary because an emission line object exists in theerror circle of ∼ 1500 -radius (Haberl & Sasaki 2000). An optical spectroscopy of the emission
line object is encouraged to examine whether it is a Be star or not.

6.1.6

No. 32 — AX J0051−733

Coherent pulsations with a 323.2(5) s period from No. 32 (AX J0051−733) were first discovered during this study from the data of obs. F. and reported by Yokogawa & Koyama(1998a) and Imanishi et al. (1999). New results were obtained from a long follow-up observation (obs. Q): AX J0051−733 was detected at a larger flux (see table 6.6) and the pulsations
were detected again with a shorter period of 321.0(1) s. The GIS spectra may show a softexcess below ∼ 1 keV (figure 6.7). A blackbody model with a temperature of ∼ 0.1 keV could
describe the excess, although the blackbody component is not necessary from a statisticalpoint of view.
Table 6.6: Spectral parameters for AX J0051−733 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.43

0.01

F
Q

0.76 (0.69–0.83)

2.2 (1.2–3.3)

2.9
6.8

1.3
3.0

Note — See page 55 for legends.

6.1. X-RAY PULSARS

69

Figure 6.7: GIS spectrum (top) and pulse profiles (bottom) of AX J0051−733. (Top left) The
crosses are data points obtained in obs. Q, while the solid line represents the best-fit model derived
from the simultaneous fitting. (Bottom) Pulse profiles in obs. F (left) and obs. Q (right) obtained
from 0.7–2.0 keV and 2.0–7.0 keV, and their ratio.

Imanishi et al. (1999) investigated 16 observations of ROSAT and Einstein by which> 10. In addition,
AX J0051−733 was covered, and found a flux variability of a factor ∼70 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Figure 6.8: Light curves of AX J0051−733 in obs. Q with a bin time of 3000 s. The upper andmiddle panels present the light curves in the soft (0.7–2.0 keV) and hard (2.0–7.0 keV) bands,
respectively. The lower panel shows the ratio of the two bands.

gradual flux increase with a factor of ∼ 2 was found during obs. Q, as shown in figure 6.7.
No change was found in the spectral hardness. The ROSAT counterpart, RX J0050.8−7316,
has a Be star in its error circle (Cowley et al. 1997). All the information indicates thatAX J0051−733 is a Be-XBP.
From the empirical relation between the pulse and orbital periods of Be-XBPs (Corbet
1984), the orbital period is predicted to be ∼ 185 days for AX J0051−733. In fact, Imanishi
et al. (1999) found weak evidence for a periodicity of ∼ 185 days in the long-term X-ray
flux history which could be interpreted as the orbital period. On the other hand, Coe &Orosz (2000) clearly detected a 0.708 day period from the I band data of the companion
Be star. They showed that the orbital period would be 1.416 days in order for the Be starto fit within the Roche lobe. The large discrepancy between Corbet’s empirical law and theperiodicity in the I band is rather problematic, and must be solved by future works.

6.1. X-RAY PULSARS

6.1.7

71

No. 37 — AX J0051−722

Coherent pulsations with a ∼ 92 s period were first discovered in the direction of SMC X-3
during an RXTE observation on November 15, 1997 (Marshall, Lochner, & Takeshima 1997).
A TOO (Target Of Opportunity) observation with ASCA on December 12 (obs. H) revealed
that there were two new pulsars, No. 37 (AX J0051−722) and No. 47 (1WGA J0053.8−7226),

Figure 6.9: GIS spectrum (left) and pulse profiles (right) of AX J0051−722. (Left) The crosses
are data points obtained in obs. H, while the solid line represents the best-fit model. (Right) Pulse
profiles in obs. H obtained from 0.7–2.0 keV and 2.0–7.0 keV, and their ratio.

Table 6.7: Spectral parameters for AX J0051−722 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.55

0.04

H

1.0 (0.7–1.3)

10 (5–16)

9.5

4.8

Note — See page 55 for legends.

72 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS
with periods of 91.12(5) s and 46.63(4) s, respectively (Corbet et al. 1998). At first,
AX J0051−722 steadily faded after its discovery, and then was found to rebrighten at theends of March and July in 1998 (Lochner et al. 1998; Lochner 1998; Israel et al. 1998a). The
spacing of these two flares and the initial outburst on 1997 November is ∼ 120 days.
Stevens, Coe, & Buckley (1999) carried out optical spectroscopic observations and discovered a Be star counterpart for AX J0051−722. Thus, AX J0051−722 is classified as a
Be-XBP. Therefore, the ∼ 120 day spacing of X-ray flares could be interpreted as the orbital
period of the neutron star in this system (Israel et al. 1998a), which is supported by theempirical law between the pulse and orbital periods (Corbet 1984).

6.1.8

No. 40 — AX J0051.6−7311

Coherent pulsations with a 172.40(3) s period from No. 40 (AX J0051.6−7311) were first
discovered during this study from the data of obs. Q (Torii et al. 2000a; Yokogawa et al.
2000b). The energy resolved pulse profiles shown in figure 6.10 exhibit the energy dependence: the higher energy photons are mostly responsible for the pulsations. The GIS spectra
may show a soft excess below ∼ 1 keV (figure 6.10). A blackbody model with a temperature of ∼ 0.18 keV compensates for the excess, although the blackbody component is notnecessary from a statistical point of view.
AX J0051.6−7311 was also detected in obs. F and obs. L with a slightly different flux
(see table 6.8), but the pulsations were not detected. As we did for AX J0049.5−7323 (seesection 6.1.5), we divided the obs. Q data into four segments and performed an FFT analysis
and epoch folding on each segment, resulting in no detection of pulsations. Therefore, weagain conclude that the very long exposure time was essential for the pulse detection.
AX J0051.6−7311 has been covered by 17 observations of ROSAT and Einstein. Yokogawa et al. (2000b) investigated these archival data and found a flux variation with a factor> 20. In addition, the ROSAT counterpart, RX J0051.9−7311, has been identified with aof ∼Be star (Cowley et al. 1997). Therefore, all the information indicates that AX J0051.6−7311
is a Be-XBP.

6.1. X-RAY PULSARS

73

Figure 6.10: GIS spectrum (left) and pulse profiles (right) of AX J0051.6−7311. (Left) The crosses
are data points obtained in obs. Q, while the solid line represents the best-fit model derived from thesimultaneous fitting. (Right) Pulse profiles in obs. Q obtained from 0.7–2.6 keV and 2.6–8.6 keV,and their ratio.

Table 6.8: Spectral parameters for AX J0051.6−7311 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.33

0.04

F
L
Q

0.9 (0.8–1.0)

0 (< 0.8)

0.68
1.5
1.3

0.29
0.65
0.56

Note — See page 55 for legends.

6.1.9

No. 43 — RX J0052.1−7319?

Coherent pulsations with a 15.3 s period from RX J0052.1−7319 were discovered in contemporaneous ROSAT and CGRO observations in 1996 (Lamb et al. 1999). Kahabka (2000)

74 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Figure 6.11: GIS spectrum of No. 43 (= RX J0052.1−7319?). The crosses are data points obtainedin obs. F, while the solid line represents the best-fit model.

Table 6.9: Spectral parameters for No. 43 (= RX J0052.1−7319?) taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.83

0.36

F

2.9 (...)

53 (< 290)

0.68

1.6

Note — See page 55 for legends.

investigated the data of two ROSAT HRI observations in 1995 and 1996 and found a largechange of flux with a factor of ∼ 200. The unabsorbed luminosities in the ROSAT band
were determined to be 2.6 × 1035 erg s−1 and 5.2 × 1037 erg s−1 (in 0.1–2.4 keV; P. Kahabka
2001, private communication), assuming a photon index of 1.0 and a column density of
3 × 1021 cm−2 . Israel et al. (1999) searched for an optical counterpart in the 1000 error circle
and discovered a Be star, in addition to a fainter object with an unknown spectral type.
From all the information, RX J0052.1−7319 is thought to be a Be-XBP.
A faint ASCA source, No. 43, is positionally coincident with RX J0052.1−7319. The
ASCA spectrum puts almost no constraint on the parameters (see table 6.9) and is consistent
with Γ and NH assumed by Kahabka (2000). We detected no sign of coherent pulsations
from either an FFT analysis or an epoch folding search, which is probably due to the highly
limited statistics. Therefore, it is not clear whether No. 43 is really the counterpart of
RX J0052.1−7319.

6.1. X-RAY PULSARS

6.1.1075

No. 44 — XTE J0054−720

A new transient pulsar XTE J0054−720 was discovered with RXTE on December 17, 1997
(Lochner et al. 1998). The pulsations have persisted in observations performed throughJanuary 12, 1998, and a monotonic decrease of the pulse period from 169.30 s (December17) to 168.40 s (January 12) was found. The flux (in 2–10 keV) initially increased from

Figure 6.12: GIS spectrum (left) and pulse profiles (right) of XTE J0054−720. (Left) The crosses
are data points obtained in obs. T, while the solid line represents the best-fit model. (Right) Pulse
profiles in obs. T obtained from 0.7–2.0 keV and 2.0–7.0 keV, and their ratio.

Table 6.10: Spectral parameters for XTE J0054−720 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.42

0.05

T

0.7 (0.5–1.0)

0.7 (< 5)

2.0

0.88

Note — See page 55 for legends.

76 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS
6.0 × 10−11 erg s−1 cm−2 (December 17) to 8.3 × 10−11 erg s−1 cm−2 (December 20), andthen decreased to 3.3 × 10−11 erg s−1 cm−2 (January 12).
Since the RXTE error circle was rather large (100 radius), identification with othersources has been difficult. In fact, five ROSAT HRI sources (Sasaki, Haberl, & Pietsch 2000)
are located within the error circle. During this study, we detected coherent pulsations witha 167.8(2) s period from No. 44 (AX J0052.9−7157) and determined its position accurately(Yokogawa et al. 2001a). We found that AX J0052.9−7157 is located within the error circle of XTE J0054−720 and has a variable Be/X-ray binary, RX J0052.9−7158 (Cowley etal. 1997), as a counterpart. From the nearly equal pulse period and the positional coincidence, we conclude that the ASCA, ROSAT, and RXTE sources are identical, and thus
XTE J0054−720 is a Be-XBP.

6.1.11

No. 47 — 1WGA J0053.8−7226 = XTE J0053−724

Coherent pulsations from No. 47 (1WGA J0053.8−7226) with a 46.63(4) s period were discovered in obs. H, as described in section 6.1.7 (Corbet et al. 1998). Buckley et al. (2001)
investigated the archival data of 21 ROSAT observations and found a large flux variability.
They also carried out follow-up optical and infrared observations and discovered two Be stars
in the error circle. All the information indicates that 1WGA J0053.8−7226 is a Be-XBP.
Table 6.11: Spectral parameters for 1WGA J0053.8−7226 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.37

0.02

H

0.8 (0.7–0.9)

1 (< 2)12

5.2

Note — See page 55 for legends.

6.1.12

No. 49 — SMC X-2

SMC X-2 is a long-known transient Be/X-ray binary with a maximum luminosity of ∼1038 erg s−1 (Clark et al. 1978; Murdin et al. 1979). Large outbursts have been detectedwith SAS-3, HEAO1, and ROSAT (Clark et al. 1978; Marshall et al. 1979; Kahabka &Pietsch 1996).
Corbet et al. (2001) detected a large outburst (∼ 1038 erg s−1 ) in the direction of6.1. X-RAY PULSARS

77

Figure 6.13: GIS spectrum (left) and pulse profiles (right) of 1WGA J0053.8−7226. (Left) The
crosses are data points obtained in obs. H, while the solid line represents the best-fit model. (Right)
Pulse profiles in obs. H obtained from 0.7–2.0 keV and 2.0–7.0 keV, and their ratio.

SMC X-2 in January–April, 2000, with the All-Sky Monitor onboard RXTE. The position
was determined with an error radius of 30 , and SMC X-2 is located near the edge of theerror circle. Coherent pulsations were also discovered during the outburst, with periods of
2.371532(2) s on April 12 and 2.371861(3) s on April 22–23 (Corbet et al. 2001). We made afollow-up ASCA observation (obs. R) and detected No. 49 at the position of SMC X-2 witha better accuracy (∼ 4000 ; see section 5.2.1). We also detected pulsations with a 2.37230(4) speriod (Yokogawa et al. 2001b), which is in full agreement with the RXTE result, indicating
the RXTE pulsar and No. 49 are identical. Figure 6.15 shows a ROSAT PSPC image during
the outburst, overlaid with error circles of SMC X-2 (determined with SAS-3), the pulsar
discovered with RXTE, and the ASCA source No. 49. We find positional coincidence of thesesources; hence, we conclude that the 2.37-s pulsations are attributable to the long-known
transient SMC X-2.

78 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Figure 6.14: GIS spectrum (left) and pulse profiles (right) of SMC X-2. (Left) The crosses are data
points obtained in obs. R, while the solid line represents the best-fit model. Spectral components
are indicated by the broken lines. (Right) Pulse profiles in obs. R obtained from 0.7–2.0 keV,
2.0–6.0 keV, 6.0–7.0 keV, and 7.0–10.0 keV.

Table 6.12: Spectral parameters for SMC X-2 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.36

0.03

R

0.7 (0.5–0.9)

0 (< 1)9.1

3.9Note — See page 55 for legends.

When fitted to a simple power-law, the ASCA spectrum showed bump-like residuals at
6–7 keV, which were eliminated by a narrow Gaussian. Adding the Gaussian component
improved the χ2 value by 7.3, which corresponds to a significance level of 97% in the F test. The center energy and the equivalent width of the Gaussian were determined to be 6.3(6.1–6.5) keV and 400 (150–640) eV, respectively. The center energy is consistent with theK-shell emission from neutral or low-ionization Fe. Yokogawa et al. (2001b) carried out a6.1. X-RAY PULSARS

79-73°35′

170
165
160

RXTE

155
ASCA

150

-73°40′

145
140

SAS-3

star 5B

135-73°45′

130
125

0h56m

0h55m

0h54m

170 175 180 185 190 195 200 205 210 215
Figure 6.15: ROSAT PSPC image around SMC X-2 overlaid with error circles of SAS-3, ASCA(No. 49), and RXTE. The plus sign indicates the position of the proposed optical counterpart (“star5B” in Murdin et al. 1979)

pulse phase resolved spectroscopy and found marginal evidence for pulsations of the Fe line
intensity.

6.1.13

No. 51 — XTE J0055−724 = 1SAX J0054.9−7226

A scan observation made with RXTE on 1998 January 20 revealed a new X-ray pulsar,
XTE J0055−724, with a pulse period of ∼ 59 s (Marshall & Lochner 1998). Santangelo
et al. (1998) made a follow-up observation with BeppoSAX on January 28, and detectedpulsations with a 58.963(3) s period from 1SAX J0054.9−7226, which is located within the
100 error circle of XTE J0055−724. The agreement of the period and position indicatesthat XTE J0055−724 and 1SAX J0054.9−7226 are the same source. During this study,
we detected weak evidence for the ∼ 59 s pulsations from No. 51 (see figure 5.6), whichis positionally coincident with 1SAX J0054.9−7226; hence, we consider No. 51 to be thecounterpart for this pulsar.

80 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Figure 6.16: GIS spectrum of XTE J0055−724. The crosses are data points obtained in obs. G,
while the solid line represents the best-fit model.

Table 6.13: Spectral parameters for XTE J0055−724 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.51

0.10

G

0.4 (0.1–1.0)

0.02 (< 10)

2.7

1.2

Note — See page 55 for legends.

Israel et al. (1998b) investigated the archival data of 13 ROSAT observations covering
1SAX J0054.9−7226. They found a flux variation with a factor of > 30 between two observations of ROSAT in 1996 and RXTE in 1998. They also determined the position with a1000 error circle, in which a Be star was later discovered by Stevens, Coe, & Buckley (1999).
XTE J0055−724 = 1SAX J0054.9−7226 is thus a Be-XBP.

6.1.14

No. 56 — AX J0057.4−7325

Coherent pulsations with a 101.45(7) s period from No. 56 (AX J0057.4−7325) were first
discovered during this study from the data of obs. R (Torii et al. 2000b; Yokogawa et al.
2000c). This source was also detected in obs. L, and weak evidence for pulsations witha period of 101.47(6) s was detected. Yokogawa et al. (2000c) investigated six ROSAT
observations covering this pulsar, and found a flux variability with a factor of > 10. All theinformation indicates that AX J0057.4−7325 is an XBP.

6.1. X-RAY PULSARS

81

Figure 6.17: GIS spectrum (top) and pulse profiles (bottom) of AX J0057.4−7325. (Top left) The
crosses are data points obtained in obs. R, while the solid line represents the best-fit model derived
from the simultaneous fitting. (Bottom) Pulse profiles in obs. L (left) and obs. R (right) obtained
from 0.7–2.0 keV and 2.0–7.0 keV, and their ratio.

So far optical follow-up observations have not been carried out. As far as we have investigated, only one optical source, MACS J0057−734#010, is located within the ASCA error

82 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Table 6.14: Spectral parameters for AX J0057.4−7325 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.48

0.05

L
R

1.1 (0.9–1.4)

8 (4–14)

1.1
2.2

0.56
1.1

Note — See page 55 for legends.

circle, for which the spectral type and existence of the Hα emission line are not known. No
counterpart is found in the catalogues of emission line objects by Meyssonnier & Azzopardi(1993) and Murphy & Bessell (2000). This is a rare case in which an XBP in the SMC is not
associated with a Be star or an emission line object (see Haberl & Sasaki 2000). We notethat AX J0057.4−7325 is located at the edge of the SMC main body, fronting the easternwing. The fact that OB supergiant X-ray binaries (only SMC X-1 and EXO 0114.6−7361;
see table 6.1 and table 7.1) are both located in the eastern wing leads us to suspect thatAX J0057.4−7325 might be the third example. Therefore, deep and detailed optical observations around this pulsar are strongly encouraged.

6.1.15

No. 61 — AX J0058−7203

Coherent pulsations with a period of 280.4(4) s from No. 61 (AX J0058−7203) were first
discovered during this study (Yokogawa & Koyama 1998a), and the PSD derived from theevents in 1.0–6.0 keV is shown in figure 6.19 (a). In a new observation T, AX J0058−7203
was detected at a large off-axis angle of ∼ 230 . We performed an FFT analysis on the eventsin the same energy band (1.0–6.0 keV) and found no significant peak as shown in figure 6.19(b). The count rate, total counts (without background), and background level in 1.0–6.0 keV
are 1.3 × 10−2 cnt s−1 , 1014 cnt, and 39% in obs. G, and 7.3 × 10−3 cnt s−1 , 642 cnt, and 42%
in obs. T. The reduced total counts may cause the non detection of pulsations, although noquantitative estimation has been done.
Tsujimoto et al. (1999) investigated the archival data of 14 observations with Einstein> 10. Haberl
and ROSAT covering this pulsar, and found a flux variation with a factor of ∼& Sasaki (2000) found an emission line object as a counterpart for AX J0058−7203. All
the information indicates that AX J0058−7203 is an XBP, and probably has a Be starcompanion. The GIS spectra may show a soft excess below ∼ 1 keV (figure 6.18). A
blackbody model with a temperature of ∼ 0.06 keV compensates for the excess, although6.1. X-RAY PULSARS

83

Figure 6.18: GIS spectrum (left) and pulse profiles (right) of AX J0058−7203. (Left) The crosses
are data points obtained in obs. G, while the solid line represents the best-fit model derived from thesimultaneous fitting. (Right) Pulse profiles in obs. G obtained from 0.7–2.0 keV and 2.0–7.0 keV,and their ratio.

Table 6.15: Spectral parameters for AX J0058−7203 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.36

0.04

G
T

0.8 (0.6–1.0)

0.7 (< 4)

2.5
4.4

1.1
1.9

Note — See page 55 for legends.

the blackbody component is not necessary from a statistical point of view.

84 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

(a)

(b)

Figure 6.19: PSDs of AX J0058−7203 in (a) obs. G and (b) obs. T, obtained from the events in
1.0–6.0 keV (bin time = 10 s). Data points with a power less than 5 are not plotted. The peakdetected in (a) is not found in (b).

6.1.16

No. 67 — RX J0059.2−7138

A new transient source RX J0059.2−7138 was detected with ROSAT and ASCA in simultaneous observations of SNR 0102−723 (obs. B). Coherent pulsations with a period of 2.7632(2) swere discovered from the ROSAT data (Hughes 1994), and were confirmed with the ASCA
data (Kohno, Yokogawa, & Koyama 2000). The pulsed fraction is highly energy dependent
as seen from figure 6.20, and is larger at higher energy (Hughes 1994; Kohno, Yokogawa,& Koyama 2000). The possible optical counterpart proposed by Hughes (1994) was later
revealed to be a Be star (Southwell & Charles 1996), thus RX J0059.2−7138 is undoubtedlya Be-XBP.
< 2 keV and a cut-off above > 7 keV
The ASCA spectrum exhibits a soft excess below ∼∼as described in section 5.4.4. We thus included a blackbody model as the soft component
and a high energy exponential cut-off. The best-fit parameters are shown in table 6.16.
Kohno, Yokogawa, & Koyama (2000) used the ROSAT and ASCA data simultaneously andcarried out more elaborate analyses. They adopted various models to describe the soft excess
and found that the soft component could be described by a thin-thermal plasma model ora broken power-law combined with an oxygen overabundance in the absorption column.
They also found that the normalization of the soft component does not change (i.e., the6.1. X-RAY PULSARS

85

Figure 6.20: GIS spectrum (left) and pulse profiles (right) of RX J0059.2−7138. (Left) The crosses
are data points obtained in obs. B, while the solid line represents the best-fit model. Spectral
components are indicated by the broken lines. (Right) Pulse profiles obtained with ROSAT PSPC
(upper three panels) and ASCA GIS (lower three panels) in the contemporaneous observations.
Energy bands from top to bottom are 0.07–0.4 keV, 0.4–1.0 keV, 1.0–2.4 keV, 0.7–2.0 keV, 2.0–
4.0 keV, and 4.0–10.0 keV.

Table 6.16: Spectral parameters for RX J0059.2−7138.
HR

HRerr

Obs. ID

Γ

Ecut

Efold

kT

NH

FX

LX

0.34

0.01

B

0.51
(0.46–0.56)

6.4(5.8–6.9)

6.3(4.4–9.4)

0.15(0.14–0.17)

0(< 1.1)240100Note — See page 55 for legends. Units for Ecut and Efold are keV.

soft component exhibits no pulsations) during the pulse phase, which is consistent with theenergy dependent pulsed fraction (figure 6.20).

86 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

6.1.17

No. 74 — RX J0101.3−7211

RX J0101.3−7211 has been known to be a highly variable ROSAT source with an emission
line object (Haberl & Sasaki 2000), and was classified as an XBP candidate in our earlier
study (source No. 27 in Yokogawa et al. 2000e) because of its hard spectrum. Therefore, thissource is strongly suspected to be a Be-XBP, although no pulsations have been detected,
probably because of the limited statistics.
Sasaki et al. (2001) performed an XMM-Newton observation on SNR 0102−723 andserendipitously detected RX J0101.3−7211. Although the flux level was the lowest among
previous detections with ROSAT and ASCA, coherent pulsations with a 455(2) s period were

Figure 6.21: GIS spectrum (left) and XMM-Newton pulse profiles (right) of RX J0101.3−7211.
(Left) The crosses are data points obtained in obs. D, while the solid line represents the best-fit
model derived from the simultaneous fitting. (Right) Pulse profiles obtained with XMM-Newton
in 0.3–1.0 keV, 1.0–2.0 keV, and 2.0–10.0 keV (taken from Sasaki et al. 2001).

Table 6.17: Spectral parameters for RX J0101.3−7211 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.43

0.05

B
D
G

0.7 (0.5–0.9)

0.4 (< 4)

1.5
1.6
1.7

0.65
0.69
0.73

Note — See page 55 for legends.

6.1. X-RAY PULSARS

87

discovered thanks to the high S/N ratio achieved by XMM-Newton. The pulsed fraction
was found to be energy dependent and higher at lower energy, the same as AX J0049.5−7323(section 6.1.5). Sasaki et al. (2001) also carried out a spectroscopic optical observation on
the emission line object found by Haberl & Sasaki (2000), and revealed it to be a Be star.
Therefore, RX J0101.3−7211 is undoubtedly a Be-XBP.

6.1.18

No. 78 — 1SAX J0103.2−7209

Hughes & Smith (1994) and Ye et al. (1995) performed ROSAT HRI observations of theshell-like radio SNR 0101−724 and detected no X-rays from the radio shell. Instead, an
X-ray point source RX J0103.2−7209 (= 1SAX J0103.2−7209) having a Be star counterpart
was detected inside the SNR.
Coherent pulsations with a period of 345.2(3) s from 1SAX J0103.2−7209 were first
discovered in a BeppoSAX observation made on 1998 July 26–27 (Israel et al. 1998a; Israel
et al. 2000). Subsequently, pulsations with a period of 348.9(3) s were detected from the
ASCA source No. 78 on May 21–23, 1996 (obs. D; Yokogawa & Koyama 1998c). Israel et al.
(2000) also detected 343.5(5) s pulsations from Chandra data obtained on August 23, 1999,
and found a monotonic spin-up with a period derivative of −1.7 s yr−1 .
1SAX J0103.2−7209 has been detected in various observations with Einstein, ROSAT,< 1036 erg s−1 (Israel et al.
ASCA, and Chandra, with nearly the same luminosity level of ∼2000, and references therein). Therefore, this pulsar can possibly be classified as a “persistent” Be-XBP, such as X Per, RX J0146.9+6121, RX J0440.9+4431, and RX J1037.5−564(White, Swank, & Holt 1983; Mereghetti et al. 2000; Reig & Roche 1999). The 755-s pulsar
AX J0049.5−7323 may also belong to this class, although the upper limit of the flux variation
is not known.
Table 6.18: Spectral parameters for 1SAX J0103.2−7209 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.36

0.06

B
D
G

0.8 (0.6–1.0)

0 (< 4)

1.3
1.1
0.79

0.56
0.48
0.34

Note — See page 55 for legends.

88 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Figure 6.22: GIS spectrum (left) and pulse profiles (right) of 1SAX J0103.2−7209. (Left) The
crosses are data points obtained in obs. D, while the solid line represents the best-fit model derived
from the simultaneous fitting. (Right) Pulse profiles in obs. D obtained from 0.7–2.0 keV and
2.0–7.0 keV, and their ratio.

6.1.19

No. 83 — AX J0105−722

Since No. 83 (AX J0105−722) and No. 84 are located only ∼ 30 from each other, mutualcontamination is not negligible. We therefore used an oval-shaped region including both
AX J0105−722 and No. 84 (region 1 in figure 6.23 (a)) in the timing analysis, and detected
coherent pulsations with a 3.34300(3) s period at a marginal significance of ∼ 99.5% (Yokogawa & Koyama 1998b; figure 5.5 (c)). We then separately searched for pulsations fromregions 2 and 3 in figure 6.23 (b) and found weak evidence for the 3.34-s pulsations only
from region 2 (which includes AX J0105−722). Therefore, we conclude that the pulsations
are attributable to AX J0105−722. The spectrum of AX J0105−722 is rather soft (see table6.19) compared to those of typical XBPs (Γ ∼ 1; Nagase 1989).

6.1. X-RAY PULSARS
180 (a)89-72°05′
180

175

(b)

-72°05′
180

175

170

-72°05′175

170

region 1(c)170

region 2

165

165

165

160

-72°10′160-72°10′
160

155

155

155

150

AX J0105-722150

AX J0105-722150

P145
P147
P160

-72°10′

P163

AX J0105-722
P172

145

145

145

140

-72°15′140-72°15′
140

135

h

1 06

m

h

1 05

m

h

m

1 04 135

region 3
h

1 06

m

h

1 05

m

h

m

1 04 135-72°15′
h

1 06

m

h

1 05

m

h

1 04m

100 105 110 115 120 125 130 135 140 145 100 105 110 115 120 125 130 135 140 145 100 105 110 115 120 125 130 135 140 145

Figure 6.23: ASCA GIS contour images around AX J0105−722 in 0.7–7.0 keV, overlaid with theequatorial coordinates (J2000). Two ASCA sources are detected in this image, one of which is
AX J0105−722 and the other is No. 84, located to the south. Regions 1–3 from which event lists
were extracted for pulsation searches (see text) are indicated by the ellipse and circles in (a) and(b). The five dots marked with “Pn” in (c) are the ROSAT sources in Haberl et al. (2000) with an
ID number of n.

Filipović et al. (2000) used the data from high resolution X-ray and radio observations around this source made with ROSAT, ATCA (Australia Telescope Compact Array),and MOST (Molonglo Observatory Synthesis Telescope). They resolved the ASCA sources
AX J0105−722 and No. 84 into five sources as shown in figure 6.23 (c): AX J0105−722
is surrounded by three ROSAT PSPC sources, Nos. 145, 147, and 163 (here, PSPC-145,
PSPC-147, and PSPC-163) in Haberl et al. (2000). PSPC-163 (RX J0105.1−7211) exhibitsa hard spectrum typical of HMXBs and is identified with an emission line object cataloguedin Meyssonier & Azzopardi (1993), and thus is a likely candidate for a Be/X-ray binary (e.g.,
Haberl & Sasaki 2000). PSPC-145 is positionally coincident with the ratio emission fromSNR DEM S128, has a soft spectrum typical of SNRs, and exhibits no flux variation, thus islikely to be an X-ray emitting SNR. PSPC-147 has the hardest spectrum among the three,although the source nature of the source is unclear. Considering these facts, we propose thatthe 3.34-s pulsations can be attributed to RX J0105.1−7211, which is a Be-XBP, and thatthe mutual contamination of X-rays from RX J0105.1−7211, PSPC-145, and PSPC-147 in
the ASCA data caused the rather soft spectrum of AX J0105−722. Because of this situation and the rather marginal detection of pulsations, follow-up X-ray observations with highspatial resolution and good S/N ratio are needed.

90 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Figure 6.24: GIS spectrum (left) and pulse profiles (right) of AX J0105−722. (Left) The crosses
are data points obtained in obs. D, while the solid line represents the best-fit model. (Right) Pulse
profiles in obs. D obtained from 0.7–2.0 keV and 2.0–7.0 keV, and their ratio.

Table 6.19: Spectral parameters for AX J0105−722 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

−0.20

0.08

D

2.2 (1.9–2.6)

0 (< 2)

0.30

0.13

Note — See page 55 for legends.

6.1.20No. 90 — XTE J0111.2−7317

A new transient source, XTE J0111.2−7317, was serendipitously found with RXTE and
CGRO, and at the same time coherent pulsations with a period of ∼ 31 s were discovered(Chakrabarty et al. 1998a; Wilson & Finger 1998). In the TOO observation with ASCA
(obs. I), pulsations with a 30.9497(4) s period were unambiguously detected from No. 90.
Thus No. 90 is undoubtedly the counterpart for XTE J0111.2−7317, and the position was6.1. X-RAY PULSARS

91

Figure 6.25: GIS spectrum (left) and pulse profiles (right) of XTE J0111.2−7317. (Left) The
crosses are data points obtained in obs. I, while the solid line represents the best-fit model. Spectral
components are indicated by the broken lines. (Right) Pulse profiles in obs. I obtained from 0.7–
2.0 keV and 2.0–7.0 keV, and their ratio.

Table 6.20: Spectral parameters for XTE J0111.2−7317 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

kT

NH

FX

LX

0.33

0.003

I

0.79 (0.76–0.82)

0.144 (0.137–0.150)

2.3 (1.6–3.0)

370

170

Note — See page 55 for legends.

determined with better accuracy than RXTE (Chakrabarty et al. 1998b). Israel et al. (1999)
detected two stars in the error circle; the fainter one was revealed to be a Be star and was
proposed to be the optical counterpart for this pulsar. Coe, Haigh, & Reig (2000) carriedout a more elaborate optical spectroscopy on the optical counterpart, and confirmed the Be
nature of this star. Therefore, XTE J0111.2−7317 is undoubtedly a Be-XBP.

92 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS
The ASCA spectrum of this pulsar shows a soft X-ray excess above a power-law modeland bump-like residuals around 6–7 keV. Adding a blackbody component and a narrow
Gaussian gave a nice fit to the spectrum, and the best-fit parameters in table 5.5 andtable 6.20 are derived from this model. However, the pulse phase resolved spectroscopy
performed by Yokogawa et al. (2000d) revealed that the soft blackbody component shouldbe pulsating in this model, while the emission region of the blackbody is extremely large(∼ 800 km in radius) and so pulsation is impossible. Therefore, Yokogawa et al. (2000d)
proposed an alternative model, the “inversely broken power-law,” which is a power-law witha larger/smaller photon index below/above a break energy, to describe the whole continuum
and the pulsations in the low energy band. The pulsations of the soft component are in
striking contrast to the non-pulsating soft component of RX J0059.2−7138, although theoverall continuum shapes in the ASCA band resemble each other (see table 6.16 and table6.20).

6.1.21

No. 94 — SMC X-1

SMC X-1 (No. 94) is a well-established XBP with a pulse period of ∼ 0.71 s, having a Btype supergiant companion (e.g., Bildsten et al. 1997). SMC X-1 has been detected in three
ASCA observations. The 0.71-s pulsations were detected only from the data of obs. A andobs. I (see table 6.1). In obs. C, it was in a low-state and went into eclipse as predicted by
the ephemeris of Wojdowski et al. (1998). Since the low-state of SMC X-1 should be causedby an occultation of the X-ray emitter, probably by the tilted accretion disk (Wojdowski etal. 1998; Vrtilek et al. 2001), the intrinsic luminosity in the low-state should be as high asthat in the high-state. Therefore, the small FX and LX derived from the obs. C data (table6.21) appear to be an underestimate of the true intensity.
The soft excess above a power-law was detected and was described by a blackbody
component with parameters given in table 6.21. Paul et al. (2001) elaborately carried out
a phase resolved spectroscopy with the ASCA data (GIS and SIS) of obs. A and foundpulsations of the soft component. Therefore, for the same reason as described in section
6.1.20, attributing the soft excess to blackbody emission is not appropriate. Paul et al.
(2001) found that the inversely broken power-law model or two-component power-law model
could describe both the continuum shape and the pulsating nature of the soft excess.

6.1. X-RAY PULSARS

93

Figure 6.26: GIS spectra (top) and pulse profiles (bottom) of SMC X-1. (Top) The crosses aredata points obtained in obs. A (left) and obs. I (right), while the solid lines represent the best-fit
models. Spectral components are indicated by the broken lines. (Bottom) Pulse profiles in obs. A(left) and obs. I (right) obtained from 0.7–2.0 keV and 2.0–7.0 keV, and their ratio.

94 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Figure 6.27: GIS spectrum of SMC X-1 in obs. C. The crosses are data points, while the solid line
represents the best-fit model. Spectral components are indicated by the broken lines.

Table 6.21: Spectral parameters for SMC X-1 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

kT

NH

0.43
0.40
0.44

0.01
0.02
0.01

A
C
I

0.69 (0.61–0.76)
0.45 (0.35–0.53)
0.35 (0.27–0.44)

0.13 (0.11–0.15)
0.20 (0.13–0.29)
0.16 (0.13–0.19)

4 (2–6)0 (< 2)2 (< 4)FX

LX

220 1108.9∗ 3.8∗
150 67Note — See page 55 for legends.
∗: Intrinsic intensity is much higher (see text).

6.1.22Other pulsars not detected with ASCA — 2E 0050.1−7247,
XTE J0052−723, RX J0117.6−7330, and XTE J0050−732#1
and #2

2E 0050.1−7247 Coherent pulsations with a 8.8816(2) s period from 2E 0050.1−7247 =RX J0051.8−7231 were discovered in a ROSAT observation by Israel et al. (1997). Flux
variability of a factor 20 between two ROSAT observations and a Be star in the error circle
were found, hence this pulsar is a Be-XBP. ASCA covered the position of 2E 0050.1−7247 inobs. H. We did not detect any positive excess above the background level from this position.
The upper limit of its flux is estimated to be ∼ 1 × 10−13 erg s−1 cm−2 (0.7–10.0 keV),assuming a photon index of ∼ 1.

6.1. X-RAY PULSARS95

XTE J0052−723 A new transient pulsar XTE J0052−723 with a period of 4.782(1) s
was discovered in RXTE observations on December 27, 2000, and January 5, 2001 (Corbet,
Marshall, & Markwardt 2001). The transient nature and the pulse period imply that thissource is propably a Be-XBP. The position was determined with a 20 × 10 error box and iscovered by ASCA observation H. We detected no X-rays from this position and set the upper
limit to be ∼ 1 × 10−13 erg s−1 cm−2 (0.7–10.0 keV), assuming a photon index of ∼ 1.

RX J0117.6−7330 RX J0117.6−7330 was serendipitously discovered in a ROSAT PSPC
observation (Clark, Remillard, & Woo 1996). The luminosity was 2.3 × 1037 erg s−1 between
0.2–2.5 keV at that time (Clark, Remillard, & Woo 1997), and was found to diminish bya factor of over 100 within one year. Macomb et al. (1999) discovered coherent pulsations
with a ∼ 22.07 s period from the same data, with the aid of the archival data obtained by
BASTE onboard CGRO in the same epoch. Strong Balmer emission lines and infrared excess
were detected from the companion star in the error circle (Coe et al. 1998), indicating that
RX J0117.6−7330 is a Be-XBP. Although the position of RX J0117.6−7330 was covered in
ASCA observations A and C, no X-ray emissions were detected. It was difficult to estimate
the upper limits of the flux because of the contamination from SMC X-1, which is locatedonly ∼ 50 away from RX J0117.6−7330.

XTE J0050−732#1 and #2 Lamb et al. (2001) submitted a paper reporting on discovery of two pulsars with periods of 16.6 s and 25.5 s from archival data of RXTE. Although
they gave no names for these pulsars, we tentatively designate them as XTE J0050−732#1(16.6 s) and XTE J0050−732#2 (25.5 s) because the RXTE observation was centeredon (00h 50m 44.s 64, −73◦ 160 04.00 8). Their long pulse periods are typical of XBPs, althoughmore information (optical counterparts and flux variability) is required for further confirmation. We argue that their identification of XTE J0050−732#1 with a ROSAT source,
RX J0051.8−7310, is incorrect. They detected RX J0051.8−7310 in the ASCA archival dataof obs. Q and found very weak evidence (significance = 82%) for 16.6 s pulsations by performing an FFT only in a range of 0.02–0.2 Hz. However, the ASCA source at the positionof RX J0051.8−7310 is AX J0051.6−7311, which is a 172-s pulsar (see section 6.1.8) andis identified with RX J0051.9−7311. Non-detection of the 172 s pulsations by Lamb et al.
(2001) is simply due to the limited frequency range used for the FFT analysis. Note that
RX J0051.8−7310 and RX J0051.9−7311 are identical. The former name was used for No.
82 in Kahabka et al. (1999) while the latter was used for No. 41 in Sasaki, Haberl, & Pietsch

96 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS(2000). Therefore, in private we proposed that Lamb et al. should perform an FFT in awhole frequency range.

6.2

Period Distribution of XBPs

We have shown that at least 22 out of the 25 X-ray pulsars in the SMC are XBPs2 probably with a high mass star companion. Recently Liu, van Paradijs, & van den Heuvel(2000) compiled a catalogue of Galactic HMXBs in which 49 XBPs are included3 . In addition, we investigated the recent literature and found four new XBPs, AX J1740.1−2847,
XTE J1543−568, SAX J2239.3+6116, and AX 1841.0−0536 (Sakano et al. 2000; Finger &Wilson 2000; In ’t Zand et al. 2001; Bamba et al. 2001). In figure 6.28 we show the distribution of pulse periods of the 53 XBPs in the Galaxy and 22 in the SMC; the actual numbers
are given in (a), while the number of SMC pulsars is normalized (multiplied by 53/20) in(b). A weak hint that the Galactic XBPs tend to have longer pulse periods can be found,although the statistics are limited. However, the difference in the period distribution may
be attributable to an observational selection effect for the following reason.
Mass accretion onto a neutron star is characterized by three radii: (1) the accretion
radius ra , within which material is captured by the gravitation and accretes onto the neutronstar, (2) the magnetospheric radius rm , where the magnetic field pressure and the rampressure of the accreting material equate with each other, and (3) the corotation radius
rc , where the neutron star corotation velocity is equal to the Keplerian velocity. Thus thedefinitions are:ra = 4GM/v02 ,(6.2)

B(rm )2 /8π = ρ(rm )v(rm )2 , and(6.3)

rc = (GM P 2 /4π 2 )1/3 ,(6.4)

where M , B(r), and P are the mass, magnetic field strength at r, and spin period of theneutron star, while ρ(r), v0 , and v(r) are the density at r, terminal velocity, and inflowing
velocity at r of the stellar wind. If rc < rm < ra , the stellar wind material within ra accretes
onto the magnetosphere and is blown away by the centrifugal force, because the drag exerted
2

AX J0043−737, AX J0049−732, and AX J0105−722 were excluded, as further confirmation is needed.
In this paper we do not regard 4U 2206+543 as an XBP because of the non detection of pulsations in afurther study by Corbet & Peele (2001)
3

6.2. PERIOD DISTRIBUTION OF XBPS
(a)

97
(b)

Figure 6.28: Period distribution of XBPs in the SMC (solid lines; triangles) and our Galaxy(dotted lines; circles). Horizontal and vertical axes respectively indicate the pulse period and thenumber of XBPs in each period bin. The number of SMC pulsars are normalized by multiplying
the factor 53/22 in (b), while the actual (unnormalized) number is shown in (a).

by the magnetosphere is super-Keplerian (the so-called “propeller” effect). Therefore, by
imposing rc = rm , the minimum luminosity Lmin at which accretion is permitted isµ37

Lmin ≈ 2 × 10

R
6
10 cm

¶−1 Ã

M
1.4M¯

!−2/3 µµ30
10 G cm3

¶2 µ

P
1s

¶−7/3

erg s−1 ,(6.5)

where µ = BR3 /2 = 0.5×1030 (B/1012 G)(R/106 cm)3 G cm3 is the magnetic dipole moment
and R is the radius of the neutron star. From this equation, it is expected that XBPs withlonger periods are fainter than those with shorter periods. Therefore, the lack of long period
XBPs in the SMC (figure 6.28) may be due to the fact that such pulsars are too faint at the
SMC distance, and so pulsations have not been detected.
We used our ASCA data and various data in the literature to investigate the relation
between the pulse periods and minimum luminosities (Lmin ) of the SMC XBPs, which werecollected from observations in which the pulsations were detected. The luminosities were calculated for the 2–10 keV band in order to avoid the uncertainty from the absorption column
density and the soft excess in the spectra. For the sources of which Lmin is detected with
ASCA, we used the spectral parameters determined in this study (table 5.5) to calculate theluminosity. For SMC X-1, Lmin was obtained with ROSAT in its low-state (Kahabka & Li
1999). As mentioned in section 6.1.21, the intrinsic luminosity in the low-state should be98 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS
as high as that in the high-state; thus, we used Lmin = 1038 erg s−1 for SMC X-1, which
is a typical value for the high-state. XTE J0052−723 has been detected with only RXTE,
but the spectral parameters have not been reported. We thus assumed the photon index
Γ to be 0.8, which is the mean value for XBPs derived from this study, and converted thesource intensity (8 mCrab; Corbet, Marshall, & Markwardt 2001) to the luminosity. The
minimum luminosities from XTE J0055−724 and RX J0101.3−7211 were detected with BeppoSAX and XMM-Newton, respectively. Since these satellites cover the 2–10 keV band, weused the spectral parameters determined with these satellites (Santangelo et al. 1998; Sasaki
et al. 2001) to derive the 2–10 keV luminosities. For XTE J0050−732#1 and #2, only
pulsed flux is reported (Lamb et al. 2001) because of contamination from other sources in
the FOV of the RXTE observation. We followed the spectral model they assumed (Γ = 1.0
and NH = 1 × 1021 cm−2 ) and determined the pulsed luminosity. The minimum luminosities from 2E 0050.1−7247 and RX J0052.1−7319 were detected with ROSAT. In order toestimate the flux, Israel et al. (1997) determined that Γ was 1.1 from the ROSAT PSPC
spectrum of 2E 0050.1−7247, while Kahabka (2000) assumed Γ = 1 for RX J0052.1−7319.
Although the ROSAT bandpass is below ∼ 2 keV, we extrapolated those power-law models
to determine the 2–10 keV luminosities. ROSAT also detected the minimum luminosity from
RX J0117.6−7330 and found a very soft spectrum (Γ ∼ 2.7) in the ROSAT bandpass (Clark,
Remillard, & Woo 1997). Such a soft spectrum is also found from RX J0059.2−7138 (Γ ∼ 2.1
for the ROSAT spectrum; determined from the archival data in this study), which has arather hard spectrum above ∼ 2 keV (section 6.1.16; Kohno, Yokogawa, & Koyama 2000).
Therefore, we assumed that RX J0117.6−7330 also has a hard spectrum above ∼ 2 keV, andthat the flux ratio below and above ∼ 2 keV is the same as that of RX J0059.2−7138. We
then converted the ROSAT flux to the 2–10 keV luminosity.
Table 6.22 and figure 6.29 give the minimum luminosities and pulse periods. The threelines in figure 6.29 indicate the centrifugal limit (equation 6.5) with µ = 1028 , 1030 , and 1032 G cm3 .
Since no XBP can exist below the line corresponding to its magnetic moment, this plot provides an upper limit of the magnetic moment for each XBP. A loose anti-correlation between
Lmin and pulse periods is found in figure 6.29, which would roughly be explained by the centrifugal barrier mechanism (equation 6.5). The fact that the observed Lmin is less sensitive
to pulse periods than predicted by equation 6.5 may indicate a severe selection effect for
XBPs on the fainter side. It is notable that the faintest four pulsars (RX J0101.3−7211,
1SAX J0103.2−7209, AX J0051.6−7311, and AX J0049.5−7323) have rather long periods,
and their pulsations were detected in a high S/N observation with XMM-Newton or in very

6.2. PERIOD DISTRIBUTION OF XBPS

99

Figure 6.29: Plot of minimum lunimosities as a function of pulse periods of X-ray binary pulsars
in the SMC. Triangles and circles represent pulsed luminosity and total luminosity, respectively.
The three lines indicate equation 6.5 with µ = 1028 , 1030 , and 1032 G cm3 , below which accretionis inhibited by centrifugal force at the magnetosphere.

long observations (D and Q) with ASCA. Therefore, high sensitivity observations in the
future will probably discover long pulsations from fainter objects, which will compensates
for the current lack of long pulsars in the SMC.

100 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Table 6.22: Minimum luminosities and other relevant information of XBPs in the SMC. XBPs arearranged in order of Lmin .
No.

Name

Period(s)0.709103
15.3
30.9497
2.763221
22.06694.782
58.96346.63
74.68
2.37230
8.8816
91.12
16.6
25.5
323.2
101.45
280.4
167.8
755.5
172.40
348.9
455

L∗min
(erg s−1 )1 × 1038
1 × 1038
1 × 1038
9 × 1037
5 × 1037
3 × 1037
9 × 1036
5 × 1036
5 × 1036
4 × 1036
4 × 1036
4 × 1036†
2 × 1036†
2 × 1036
1 × 1036
9 × 1035
9 × 1035
8 × 1035
7 × 1035
5 × 1035
4 × 1035
2 × 1035

References

94 SMC X-1
1 (ROSAT)
43 RX J0052.1−7319
2 (ROSAT)
90 XTE J0111.2−7317
3 (obs. I)
67 RX J0059.2−7138
4 (obs. B)... RX J0117.6−7330
5 (ROSAT)... XTE J0052−723
6 (RXTE)
51 XTE J0055−724
7 (SAX)
47 1WGA J0053.8−7226
This work (obs. H)
24 AX J0049−729
8 (obs. F)
49 SMC X-2
9 (obs. R)... 2E 0050.1−7247
10 (ROSAT)
37 AX J0051−722
This work (obs. H)... XTE J0050−732#1
11 (RXTE)... XTE J0050−732#2
11 (RXTE)
32 AX J0051−733
12 (obs. F)
56 AX J0057.4−7325
13 (obs. R)
61 AX J0058−7203
14 (obs. G)
44 XTE J0054−720
15 (obs. T)
30 AX J0049.5−7323
16 (obs. Q)
40 AX J0051.6−7311
17 (obs. Q)
78 1SAX J0103.2−7209
This work (obs. D)
74 RX J0101.3−7211
18 (XMM)∗: In 2–10 keV.
†: Pulsed luminosity.
References — (1) Kahabka & Li 1999; (2) Kahabka 2000; (3) Yokogawa et al.
2000d; (4) Kohno, Yokogawa, & Koyama 2000; (5) Macomb et al. 1999 (6)Corbet, Marshall, & Markwardt 2001; (7) Santangelo et al. 1998; (8) Yokogawa
et al. 1999; (9) Yokogawa et al. 2001b; (10) Israel et al. 1997; (11) Lamb et al.
2001; (12) Imanishi et al. 1999; (13) Yokogawa et al. 2000c; (14) Tsujimoto etal. 1999; (15) Yokogawa et al. 2001a; (16) Yokogawa et al. 2000a; (17) Yokogawa
et al. 2000b; (18) Sasaki et al. 2001;

6.3. SNRS – I: FROM SIS

6.3

SNRs – I: from SIS

6.3.1

No. 21 — 0045−734 (N19)101

Yokogawa et al. (2001c) carried out a detailed study of three SNRs, No. 21 (0045−734), No.
66 (0057−7226), and No. 82 (0103−726), with ROSAT HRI images and ASCA SIS spectra.
We give a brief summary of each SNR separately in corresponding subsections.
We simultaneously used the SIS spectra of 0045−734 obtained in obs. F and obs. Q(figure 6.30). We fitted the spectra with an optically thin thermal plasma model in a nonequilibrium ionization (NEI) state, coded by Masai (1994). The physical parameters in the
NEI model are the electron temperature kT , metal abundances, and the ionization timescaleτ = ne t, where ne is the electron density and t is the elapsed time after the plasma washeated up.
We first fitted the spectra with all the metal abundances fixed to the SMC mean value(0.2 solar). This model, however, left bump-like residuals at ∼ 1 keV and ∼ 1.3 keV,
indicating the existence of emission lines from Ne-K/Fe-L and Mg-K. Therefore, we allowed
the Mg abundance to be free. Since the Ne-K and Fe-L lines may equally contribute to the
X-rays around ∼ 1 keV, we treated the abundances of Ne and Fe in the following three cases:(I) Ne is free, (II) Fe is free, and (III) both Ne and Fe are free. In model III, the derived
best-fit temperature is unusually large, kT ∼ 25 keV, thus we did not adopt it. Models I
and II, on the other hand, give equally better fits with no bump-like residuals as shown infigure 6.30. The best-fit parameters are given in table 6.23. In both models the abundancesof Mg and Ne/Fe are larger than those of the SMC mean value.
The ROSAT HRI image is shown in figure 6.30, overlaid with the radio contours at
843 MHz. The radio emission exhibits a partial shell in the northwest (NW) and weakdiffuse structure in the southeast (SE). No X-ray emissions are found at the radio shell;
instead, X-rays are concentrated within the radio shell and diffuse structure.
From the X-ray image, we estimated the volume V of the X-ray emitting plasma. We
simply assumed that the plasma in 0045−734 fills a cylinder with a height of 30 and adiameter of 10 , thus we obtained V = 3.6 × 1059 · β cm3 , where β represents the uncertaintyof the volume. From the emission measure EM determinedby the spectral fitting, we canqestimate the electron density of the plasma ne by ne = EM/V . Here and in the following,
the number density of nucleons is simply assumed to be the same as that of electrons. The

102 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS-72d23m-72d24m-72d25m-72d26m-72d27m00h46m00h45m30s

Figure 6.30: (Left) SIS spectrum of 0045−734 obtained in obs. Q (crosses) with the best-fit modelin case I (solid line). (Right) Gray-scale image around 0045−734 obtained with ROSAT HRI,
overlaid with contours of the MOST radio image at 843 MHz (Mills et al. 1982). The equatorial
coordinates are in equinox 1950.0.

age t was then determined from the ionization timescale τ by t = τ /ne . The total mass ofthe plasma Mtotal is estimated by Mtotal = ne V mH , where mH is the mass of a hydrogenatom. We then determined the mass of the overabundant heavy elements using Mtotal andthe best-fit abundances. These estimated parameters are given in table 6.24. The lower limit
for the ionization age of 0045−734 is 2.9 × 104 yr in both the models, which is consistent
with the age of 5.7 × 104 yr derived from an optical study of the kinematics (Rosado, Le
Coarer, & Georgelin 1994).
Since the best-fit abundances are much higher than the SMC mean value (table 6.23),
the ejecta should contribute largely to the element masses derived above. Therefore, in
figure 6.31, we compare the element masses with the theoretical predictions of chemicalcompositions produced with supernovae (SNe) with various progenitor masses (Tsujimoto
et al. 1995). A type Ia origin is rejected because the data do not show a large excess of Fe
as expected from the theory. On the other hand, the data roughly agree with a type II SN< 20M¯ .
origin with a progenitor mass of ∼

6.3. SNRS – I: FROM SIS

103

Table 6.23: Spectral parameters for 0045−734, 0057−7226, and 0103−726.
SNR . . . . . . . . . . . . . . . . . .
—— 0045−734 ——∗Model . . . . . . . . . . . . . . .
III+∞log τ (s cm−3 ) . . . . . . . . 13.00‡ +∞
11.74
−0.41
−1.20
+0.08
kTe (keV) . . . . . . . . . . . . . 0.37+0.10
0.96
−0.05
−0.05
NH (1021 H cm−2 ) . . . . 10+1
3+1
−2
−2
Flux† (erg s−1 cm−2 ) . 2.9 × 10−13 3.1 × 10−13– Abundances (solar) –O . . . . . . . . . . . . . . . . . . . 0.2 (fixed)
0.2 (fixed)+1.0
Ne . . . . . . . . . . . . . . . . . . 2.6−0.7
0.2 (fixed)
Mg . . . . . . . . . . . . . . . . . . 1.0+0.5
1.6+1.7−0.4−0.6
Si . . . . . . . . . . . . . . . . . . . 0.2 (fixed)
0.2 (fixed)
Fe . . . . . . . . . . . . . . . . . . . 0.2 (fixed)
0.9+0.5−0.3
Reduced χ2 /d.o.f. . . . . . 1.45/89
1.60/89
Note — Errors indicate the 90% confidence limits.
∗: Defined in section 6.3.1.
†: In 0.7–10.0 keV.
‡: Pegged by the upper limit of the plasma code.

0057−7226...
11.01+∞−0.44
0.9+0.8−0.3
1+2−1
3.0 × 10−13—— 0103−726 ——
I
III
+0.25
10.82
10.97+0.21
−0.26
−0.21
+0.5
0.67+0.13
1.0
−0.10
−0.2
0+0.5
0+0.3
1.1 × 10−12 1.0 × 10−12

0.2 (fixed)
0.2 (fixed)
0.2 (fixed)
0.2 (fixed)
0.2 (fixed)
0.91/220.5+0.3−0.1
1.0+0.2−0.2
0.9+0.3−0.2
0.4+0.2−0.2
0.2 (fixed)
1.47/64

1.7+3.0−0.7
1.9+1.7−0.6
2.0+1.8−0.8
0.7+1.1−0.4
1.0+1.1−0.5
1.41/63

Table 6.24: Plasma parameters for 0045−734, 0057−7226, and 0103−726.
SNR . . . . . . . . . . . . . .
——— 0045−734 ———0057−7226——— 0103−726 ———∗Model . . . . . . . . . . . .
I
II
...
I
III
3
59
59
60
V (β cm ) . . . . . . . .
— 3.6 × 10 —6.6 × 10
— 1.3 × 10 —ne (β −1/2 cm−3 ) . .
0.69
0.23
0.18
0.17
0.097
+∞
+∞
+1.0
+1.6
1/2
4
4
4
4
t (β
yr) . . . . . . . .
(> 2.9 × 10 )7.6−4.7 × 10
1.8−1.2 × 10
1.8−0.7 × 10
2.2−1.0 × 104
1/2
Mtotal (β M¯ ) . .
210
69
100
180
110
1/2
– Mass (β M¯ ) –−1
O. . . . . . . . . . . . . . . .
...†
...†
...†
5.6+3.3
1.2+2.0−1.1 × 10
−0.5
+0.4
+2.1
−1
†
†
−1
−1
Ne . . . . . . . . . . . . . . . 6.0+2.3
×
10
...
...
2.0
×
10
2.3
−1.6
−0.4
−0.7 × 10
−2
−2
−2
−2
Mg . . . . . . . . . . . . . . 8.6+4.4
4.6+4.9...†
6.7+2.3
9.1+7.9−3.4 × 10−1.7 × 10−1.5 × 10−3.6 × 10
+1.6
+5.5
†
†
†
−2
Si . . . . . . . . . . . . . . .
...
...
...
3.3−1.7 × 10
3.5−2.0 × 10−2
−2
−1
Fe . . . . . . . . . . . . . . .
...†6.9+4.1...†...†
1.2+1.4−2.3 × 10−0.6 × 10
Note — The uncertainty of the volume V is represented by β (which could be different for different
SNRs). Errors are derived from the 90% confidence limits for τ and elemental abundances in the spectral
fitting (table 6.23).
∗: Defined in section 6.3.1.
†: Fixed at 0.2 solar abundance in the spectral fitting.

104 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS- model I- model II

Figure 6.31: Nucleosynthesis products of various SNe overlaid with the element masses derived
from the spectral analyses (table 6.24) for 0045−734. Solid lines with asterisks and diamonds
represent products in type II and type Ia SNe, respectively (Tsujimoto et al. 1995). Progenitor
masses of type II SNe are also indicated. Dotted lines with circles and triangles represent models
I and II, respectively. Plots for elements fixed at 0.2 solar in the fitting are not presented.

6.3.2

No. 25 — 0047−735

Source and background regions for 0047−735 (No. 25) were carefully chosen to avoid contamination from the nearby pulsar AX J0049−732 (No. 26). As a result, the extracted SIS
spectrum had rather poor statistics, which probably caused non-detection of emission lines(section 5.4.3). We thus used the GIS spectra (in obs. F and Q) as well to compensate forthe poor statistics of the SIS spectrum, and carried out a simultaneous fit.
The spectra were well fitted to both the power-law and the CIE thermal model. However,
we adopt the thermal model because the best-fit temperature is reasonable for an SNR. This
fact would indicate that X-ray emission from this SNR has a thermal origin.

6.3.3

No. 66 — 0057−7226 (N66)

Although the SIS spectrum of 0057−7226 (figure 6.33) can be fitted well with an NEI modelwith free abundances, we obtained no constraint on the abundances due to the limited

6.3. SNRS – I: FROM SIS

105

Figure 6.32: SIS spectrum of 0047−735. The crosses are data points obtained in obs. Q, while thesolid line represents the best-fit model obtained by the simultaneous fitting.

Table 6.25: Spectral parameters for 0047−735 taken from table 5.5.
HR

HRerr

Obs. ID

kT

NH

FX

LX

−1.00

1.38

F
Q

0.86 (0.82–0.90)

0 (< 0.4)

0.14
0.11

0.060
0.048

Note — See page 55 for legends.

statistics. Therefore, we fixed the abundances at the SMC mean value (0.2 solar) and fitted
the spectrum again. The best-fit parameters are given in table 6.23 (page 103).
The ROSAT HRI image is shown in figure 6.33, overlaid with the contours which indicate
non-thermal radio structure. The radio emission exhibits a shell-like structure with a ∼ 30
diameter, while the X-rays are concentrated in the radio shell with a ∼ 20 diameter.
Assuming that the plasma fills a sphere with a diameter of ∼ 20 , we obtained thevolume V = 6.6 × 1059 · β cm3 . With the same procedure as described in section 6.3.1,
plasma parameters are obtained and given in table 6.24.

106 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS-72d24m-72d25m-72d26m-72d27m-72d28m00h58m00h57m30s

Figure 6.33: (Left) SIS spectrum of 0057−7226 obtained in obs. G (crosses) with the best-fit model(solid line). (Right) Gray-scale image around 0057−7226 obtained with ROSAT HRI. Overlaid
contours represent the difference map of Hα and MOST 843 MHz images reproduced from Ye,
Turtle, & Kennicutt (1991), which selectively reveals non-thermal radio emission. The equatorial
coordinates are in equinox 1950.0.

6.3.4

No. 81 — 0102−723

A detailed analysis of the SIS spectrum of SNR 0102−723 was carried out by Hayashi et al.
(1994). Strong emission lines from various elements were detected, which is consistent withthis SNR being young. We adopted the same model and determined its flux and luminosity
in table 5.5.

6.3.5

No. 82 — 0103−726

We first fitted the SIS spectrum of 0103−726 with an NEI model, in which all abundanceswere fixed to the SMC mean value. This model showed bump-like residuals at energies ofK-shell emission lines from O, Ne, Mg, and Si, and Fe-L lines. Therefore, we allowed theabundances of O, Mg, and Si to be free, while the abundances of Ne and Fe were treatedas described in section 6.3.1 (models I–III). We did not adopt model II because this model6.3. SNRS – I: FROM SIS

107-72d37m-72d38m-72d39m-72d40m-72d41m

01h04m01h03mFigure 6.34: (Left) SIS spectrum of 0103−726 obtained in obs. D (crosses) with the best-fit modelin case III (solid line). (Right) Gray-scale image around 0103−726 obtained with ROSAT HRI,
overlaid with contours of the MOST radio image at 843 MHz (Mills et al. 1982). The equatorial
coordinates are in equinox 1950.0.

gave a much worse fit than models I and III, with a reduced χ2 = 1.94 for 64 d.o.f. (see table6.23). In both models (I and III), abundances are larger than the SMC mean value. The
best-fit model is shown in figure 6.34.
The ROSAT HRI image is shown in figure 6.34, overlaid with the radio contours at
843 MHz. The radio emission has a faint shell with a diameter of ∼ 30 , with enhanced
emission at the SE part. The X-rays also show faint emission along the radio shell, in
addition to the more prominent emission near the center of the radio shell.
Assuming that the plasma fills a sphere with a diameter of ∼ 2.0 5, we obtained thevolume V = 1.3 × 1060 · β cm3 . With the same procedure as described in section 6.3.1,
plasma parameters are obtained and given in table 6.24.
For the same reason described in section 6.3.1, we compared the element masses withthe theoretical predictions as shown in figure 6.35. A type Ia origin is rejected and a type II< 20M¯ is preferred, similar to 0045−734.
SNe with a progenitor mass of ∼108 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS- model I- model III

Figure 6.35: Nucleosynthesis products of various SNe overlaid with the element masses derived
from the spectral analyses (table 6.24) for 0103−726. Solid lines with asterisks and diamonds
represent products in type II and type Ia SNe, respectively (Tsujimoto et al. 1995). Progenitor
masses of the type II SNe are also indicated. Dotted lines with circles and squares represent models
I and III, respectively. Plots for elements fixed at 0.2 solar in the fitting are not presented.

6.4. SNRS – II: FROM GIS

109

6.4

SNRs – II: from GIS

6.4.1

No. 23 — 0046−735

Both a power-law model and a thermal CIE model describe the GIS spectrum well due
to highly limited statistics. However, we adopted a thermal model because the derived
temperature is reasonable for an SNR; the thermal nature of this SNR is thus suggested.

Figure 6.36: GIS spectrum of 0046−735. The crosses are data points obtained in obs. Q, while
the solid line represents the best-fit model.

Table 6.26: Spectral parameters for 0046−735 taken from table 5.5.
HR

HRerr

Obs. ID

kT

NH

FX

LX

−0.62

0.36

Q

0.9 (0.1–1.8)

0 (< 15)

0.058

0.025

Note — See page 55 for legends.

6.4.2

No. 36 — 0049−736

In our earlier study (Yokogawa et al. 2000e), this SNR was regarded as a thermal SNR
because of its soft spectrum, although no emission lines were detected due to highly limited
statistics. Data from new observations (L and Q) allowed us to detect emission lines (seesection 5.4.3), and thus the thermal nature is now established. The GIS spectra were well110 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Figure 6.37: GIS spectrum of 0049−736. The crosses are data points obtained in obs. Q, while
the solid line represents the best-fit model obtained by the simultaneous fitting.

Table 6.27: Spectral parameters for 0049−736 taken from table 5.5.
HR

HRerr

Obs. ID

kT

NH

FX

LX

−0.92

0.07

F
L
Q

0.23 (0.18–0.47)

6 (2–7)

0.64
0.65
0.74

1.8
1.8
2.0

Note — See page 55 for legends.

fitted to a CIE thermal model, and the best-fit parameters are given in table 6.27. The
bump-like residuals are found at ∼ 1.9 keV (figure 6.37), which corresponds to the energyof the emission line from He-like Si. However, allowing the Si abundance to be free did not
improve the result within the statistical error.

6.4.3

No. 64 — 0056−725

Since fitting the spectrum with a thin thermal plasma model yielded an unusually hightemperature of ∼ 20 keV, we adopted a power-law model. There is a possibility that X-rays
from 0056−725 actually have a non-thermal origin.

6.4. SNRS – II: FROM GIS

111

Figure 6.38: GIS spectrum of 0056−725. The crosses are data points obtained in obs. G, while
the solid line represents the best-fit model.

Table 6.28: Spectral parameters for 0056−725 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.20

0.12

G

1.5 (1.0–2.3)

2 (< 12)

0.34

0.16

Note — See page 55 for legends.

112 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

6.5

Other Interesting Sources

6.5.1

No. 20 — IKT1

We tentatively consider No. 20, RX J0047.3−7312, and IKT1 to be identical simply becauseof the positional coincidence. Haberl & Sasaki (2000) proposed RX J0047.3−7312 as a Be/Xray binary candidate because this source exhibits a flux variation with a factor of 9 and has
an emission line object as a counterpart.
The GIS spectra in two observations (F and Q) are hard, having a photon index of ∼ 1,
which is typical of Be/X-ray binaries. In addition, a flare-like behaviour is found from thelight curve in obs. Q (figure 6.40). These results further strengthen the Be/X-ray binary
scenario.

Figure 6.39: GIS spectrum of IKT1. The crosses are data points obtained in obs. Q, while thesolid line represents the best-fit model derived from the simultaneous fitting.

Table 6.29: Spectral parameters for IKT1 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.39

0.03

F
Q

1.0 (0.9–1.2)

4 (1–7)

0.67
1.4

0.32
0.65

Note — See page 55 for legends.

6.5. OTHER INTERESTING SOURCES

113

Figure 6.40: Light curves of IKT1 in obs. Q with a bin time of 2000 s. The upper and middle
panels present the light curves in the soft (0.7–2.0 keV) and hard (2.0–7.0 keV) bands, respectively.
The lower panels present the ratio of the two bands. A flare-like activity is found at the time
∼ 2 × 105 s; the right panel shows the time zone around the flare.

6.5.2

No. 22 — AX J0048.2−7309

No. 22 (AX J0048.2−7309) was detected in two observations (F and Q) and shows a hard
spectrum (Γ ∼ 1) and a flux variability with a factor of ∼ 5 (table 6.30). In addition, we
found an emission line object, No. 215 in Meyssonnier & Azzopardi (1993), in the error circleof AX J0048.2−7309. All the information suggests that this source is probably a Be-XBP.
Follow-up observations for pulsation searches and optical identification are thus encouraged.
Although the GIS spectra were well fitted to a simple power-law (figure 6.41 (a)), there
remains a bump-like residual in the vicinity of 6–7 keV. In order to examine the existenceof an emission line, we used only the data of obs. Q, which has much better statistics. We
added a narrow Gaussian and fitted the spectrum (figure 6.41 (b)), and then determined114 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

(a)

(b)

Figure 6.41: GIS spectra of AX J0048.2−7309. The crosses are data points obtained in obs. Q,
while the solid lines represent the best-fit models described in the text; a power-law without aGaussian (a), and with a Gaussian (b). Spectral components are indicated by the broken lines in(b).

Table 6.30: Spectral parameters for AX J0048.2−7309 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

0.59

0.05

F
Q

0.8 (0.6–1.1)

9 (3–17)

0.20
1.1

0.099
0.52

Note — See page 55 for legends.

the center energy and the equivalent width of the Gaussian to be 6.6 (6.4–6.9) keV and 240(40–440) eV, respectively. However, the significance of the Gaussian is only ∼ 90% in the
F -test, thus the existence of the emission line is still not clear. Confirmation of the existenceof the emission line would further strengthen the X-ray binary nature of this source.

6.5.3

No. 105 — AX J0128.4−7329

No. 105 (AX J0128.4−7329) is located near the center of an expanding supergiant shell,
SMC-1, which has a diameter of ∼ 1◦ (Meaburn 1980). Wang & Wu (1992) reported that
the X-ray emission from AX J0128.4−7329 appears to be diffuse in an Einstein IPC image,and that the spectral hardness ratio indicates a temperature of ∼ 0.8 keV if the absorption

6.5. OTHER INTERESTING SOURCES

115

Figure 6.42: GIS spectrum of AX J0128.4−7329. The crosses are data points obtained in obs. O,
while the solid line represents the best-fit model.

Table 6.31: Spectral parameters for AX J0128.4−7329 taken from table 5.5.
HR

HRerr

Obs. ID

Γ

NH

FX

LX

−0.06

0.12

O

2.8 (1.7–4.8)

7 (< 20)

0.34

0.29

Note — See page 55 for legends.

column density is 3 × 1020 cm−2 .
Since the GIS spectrum (figure 6.42) has rather limited statistics, no rigid constraint
was obtained for the spectral parameters. When we fitted the spectrum with a CIE thermal
model, a temperature of 2.6 (0.9–10) keV was obtained, which is marginally consistent withthe argument of Wang & Wu (1992). Since half the GIS FOV is polluted by stray light thatis probably from SMC X-1, it is difficult to know the extent of AX J0128.4−7329.

116 CHAPTER 6. COMMENTS ON SPECIFIC SOURCES AND PULSAR STATISTICS

Chapter 7
Source Classification
7.1

Criteria for Source Classification

We classify the 106 ASCA sources into several source classes, which are given in the lastcolumn of table 5.3. In this section, we give the definitions of the classes.
XBPs (“BP” in table 5.3) are the 22 pulsars having a long pulse period (∼ 1–1000 s),
hard spectrum (Γ ∼ 1), flux variability, and/or an optical counterpart, as described in section
6.1; 17 of the 22 were detected with ASCA. Pulsars which are not definitely regarded as XBPs
are designated as “P”. Thermal SNRs (“TS”) are the five SNRs from which emission lines ofionized atoms were detected in section 5.4.3. The other SNRs with no significant emission
line are classified as radio SNRs (“RS”). Candidates for XBPs and thermal SNRs (“BPc”
and “TSc”) are defined in section 7.2.3.
Nonpulsating HMXBs (“NH”) and candidates (“NHc”) are defined by the optical counterparts and flux variability as summarized in table 7.1. Grades from A to E are assigned
according to the following criteria. Grade A and B sources are X-ray sources with a Be ora supergiant star companion. Flux variability has been known for grade A’s while not for
grade B’s. Grade C and D sources consist of X-ray sources having an emission line object (Bestar candidate) as a counterpart, which are catalogued in Haberl & Sasaki (2000). Again,
flux variability has been known for grade C’s while not for grade D’s. Grade E sources also
have an emission line object counterpart, but there are other possibilities (AGN or active
corona of a late type star) for the nature of the sources. In table 5.3, we designate grade A
117

118

CHAPTER 7. SOURCE CLASSIFICATION

Table 7.1: Nonpulsating HMXBs and likely candidates in the SMC.
NumberName
OpticalVariability†
Grade
ASCA
HS00
Identification∗...
1
RX J0032.9−7348
Be
Yes (KP96)
A...
22
SMC X-3Bea
Yes (CDL78)
A
...
...
1H 0103−762
Beb
Yes (KP96)
A...
42
RX J0101.0−7206
Be
Yes (KP96)
A
...
...
EXO 0114.6−7361
B (SG)cYes (WW92)
A...
24
2E 0051.1−7304
Be
Yes? (WW92)
B63
38
RX J0058.2−7231
Be
...
B
...
...
RX J0106.2−7205
Bed
...
B
20
6
RX J0047.3−7312 (IKT1)
Be?
Yes (HS00, ASCA, flare‡ )
C
§
e
22
...
AX J0048.2−7309
Be?
Yes (ASCA)C
50
30
RX J0054.9−7245
Be?
Yes (HS00)
C...
36
RX J0057.8−7207
Be?
Yes (HS00)
C...
44
RX J0101.6−7204
Be?
Yes (HS00)
C76
45
RX J0101.8−7223
Be?
Yes (HS00)
C...
50
RX J0103.6−7201
Be?
Yes (HS00)
C...
53
RX J0105.1−7211
Be?
Yes (FHPM00)
C
87
56
RX J0107.1−7235
Be?
Yes (HS00)
C...
2
RX J0041.2−7306
Be?
...
D
...
3
RX J0045.6−7313
Be?
...
D
...
8
RX J0048.5−7302
Be?
...
D
28
11
RX J0049.5−7331
Be?
Yes? (ASCA)D...
14
RX J0050.7−7332
Be?
...
D
35
15
RX J0050.9−7310
Be?
...
D
...
17
RX J0051.3−7250
Be?
...
D
...
21
No. 26 in WW92
Be?
...
D
...
26
RX J0053.4−7227
Be?
...
D
...
29
RX J0054.5−7228
Be?
...
D
...
32
2E 0054.4−7237 (IKT13)
Be?
...
D
...
34
2E 0055.8−7229 (IKT15)
Be?
...
D
...
37
RX J0057.9−7156
Be?
...
D
...
51
RX J0104.1−7243
Be?
...
D
...
52
RX J0104.5−7121¶
Be?
...
D
...
54
RX J0105.7−7226
Be?
...
D
85
55
RX J0105.9−7203
Be?
...
D
...
60
RX J0119.6−7330
Be?
...
D
...
5
RX J0047.3−7239
Be?AGN?∗∗
...
E
...
18
RX J0051.8−7159
Be?AC?††...
E
Note — In table 5.3, we designate grade A and B sources as NH and grade C and D sources as NHc.
∗: References are found in Haberl & Sasaki (2000), except for the five sources marked with a–e. (a) Crampton,
Hutchings, & Cowley 1978; (b) Whitlock & Lochner 1994; (c) Wang & Wu 1992 (SG=supergiant); (d) Hughes& Smith 1994; (e) this work (No. 215 in MA93 is located 2300 from the ASCA position).
†: For sources with a “Yes” sign, flux variability has been detected by KP96 (Kahabka & Pietsch 1996),
CDL78 (Clark et al. 1978), WW92 (Wang & Wu 1992), HS00 (Haberl & Sasaki 2000), ASCA (this work),
and FHPM00 (Filipović et al. 2000).
‡: Flaring behaviour is detected in obs. Q (see section 6.5.1).
§: Fe line is marginally detected in obs. Q (see section 6.5.2).
¶: This source is designated as RX J0104.5−7121 in HS00 (section 6.4), but probably should be
RX J0104.5−7221 according to its coordinate.
∗∗: A point-like radio source is also located in the error circle, thus the possibility that it is an AGN remains(HS00).
††: X-ray emission may be attributable to an active corona (AC) of a late type star (HS00).

7.2. CLASSIFICATION BY HARDNESS RATIO

119

and B sources as “NH” and grade C and D sources as “NHc”.
Foreground stars (“FS”) and AGN candidates (“AGN”) are defined merely based on
positional coincidence of ROSAT sources in these classes (see Sasaki, Haberl, & Pietsch 2000
and references therein). Sources which do not fall into any classes are designated as “UN”
(unclassified sources).

7.2

Classification by Hardness Ratio

7.2.1

XBPs and thermal SNRs

The source classes most rigidly defined in this study are XBPs and thermal SNRs. Fora statistical discussion, more samples of thermal SNRs are needed because there are only
five thermal SNRs in the SMC. X-ray sources in the LMC are suitable for inclusion for thereason described in section 7.2.2. Therefore, we included ASCA data of XBPs and thermal
SNRs in the LMC, which are systematically analyzed by Nishiuchi (2001). Four XBPs
in the LMC have been detected with ASCA: EXO 053109−6609.2, LMC X-4, A0538−67,
and 1SAX J0544.1−710. From these, flux variability and optical counterparts have been
detected (Bildsten et al. 1997; Burderi et al. 1998; Cusumano et al. 1998; Haberl, Dennerl,& Pietsch 1995; Corbet et al. 1997). Although pulsations were not detected from the ASCA
data of A0538−67, this source is very likely a Be-XBP (Corbet et al. 1997, and references
therein). ASCA has detected over 20 SNRs in the LMC. Among them, 10 SNRs analyzed by
Hughes et al. (1995) and Hughes, Hayashi, & Koyama (1998) are well-established thermal
SNRs (N103B, 0509−67.5, 0519−69.0, N23, N49, N63A, DEM71, N132D, 0453−68.5, andN49B). Two other classes of sources, Crab-like SNRs1 (0540−69 and N157B) and BH binaries(LMC X-1 and LMC X-3), have also been detected with ASCA, and are included in our
discussion.
Spectral parameters for the four XBPs, 10 thermal SNRs, two Crab-like SNRs, and two
BH binaries in the LMC were derived by Nishiuchi (2001), almost in the same way as we didfor the SMC sources. In figure 7.1, we show a plot of HR for these classes against observed
luminosity Lobs , defined as Lobs = Fx × 4πd2 , where d is the distance to the SMC (60 kpc)
or LMC (50 kpc). Sources detected multiple times are represented by single points, which1

We use “Crab-like SNRs” to refer to SNRs associated with a rotation-powered X-ray pulsar.

120

CHAPTER 7. SOURCE CLASSIFICATIONXBP region

Thermal SNR region

Figure 7.1: Plot of HR against Lobs (0.7–10.0 keV) for XBPs, thermal SNRs, Crab-like SNRs, andBH binaries. The symbols represent XBPs (stars), thermal SNRs (circles), Crab-like SNRs (crosses),and BH binaries (asterisks). Open and filled symbols are for the SMC and LMC, respectively.
Sources detected multiple times are represented by single points, which correspond to the largest
Lobs of each source. The source at HR = 0.83 is No. 43, which is positionally coincident with XBP
RX J0052.1−7319, but does not exhibit coherent pulsations in the ASCA data.

correspond to the largest Lobs of each source. Thanks to the LMC samples, the number ofthermal SNRs is enhanced, and so a clear split between XBPs and thermal SNRs could befound. We can safely say that in the Magellanic Clouds, XBPs fall in a narrow region of< HR < 0.6 (“XBP region”), while thermal SNRs fall in the region −1.0 < HR < −0.6
0.2 ∼
∼
∼
∼
(“thermal SNR region”). The source located above the XBP region (at HR = 0.83) is No. 43,
which is positionally coincident with XBP RX J0052.1−7319, but does not exhibit coherent
pulsations in the ASCA data. Because of the large HR error and the fact that there are no
pulsations, we ignored this source when defining the XBP region.
We also show the HR–Lobs plot for the other sources in figure 7.2, i.e., nonpulsating
HMXBs, radio SNRs (not regarded as thermal), pulsars (not regarded as XBPs), sources coincident with AGNs and foreground stars, and the remaining unclassified sources. Althoughsome sources have very large HR errors, we find that most of the sources tend to be located7.2. CLASSIFICATION BY HARDNESS RATIO121

XBP region

Thermal SNR region

Figure 7.2: Plot of HR against Lobs (0.7–10.0 keV) for the SMC sources not plotted in figure7.1. The symbols represent nonpulsating HMXBs and their candidates (triangles), radio SNRs notregarded as thermal (dotted circles), pulsars not regarded as XBPs (circles with a plus sign), sources
coincident with AGNs and foreground stars (diamonds), and other unclassified sources (squares).
Sources detected multiple times are represented by single points, which correspond to the largest
Lobs of each source.

in the XBP region and between the XBP and thermal SNR regions.

7.2.2

Validity of the classification by hardness ratio

The spectral hardness ratio has been used as an aid for classification of X-ray sources (e.g.,
Wang & Wu 1992; Haberl et al. 2000). We should note that the hardness ratios used intheir/our studies are defined only by the number of detected photons. Therefore, a strong
point of the classification by hardness ratio is that it is independent of any spectral models.
However, there could be uncertainty from interstellar absorption, which reduces the soft
photon flux and thus makes the spectra harder. Below we show that interstellar absorption
does not affect our HR classification method when applied to the Magellanic Cloud sources,
but does have an effect if applied to sources in our Galaxy.

122

CHAPTER 7. SOURCE CLASSIFICATION

We simulated HRs for simple spectra of thin-thermal plasma with a temperature of
0.5 keV, a typical value for a thermal SNR, and of a power-law with a photon index of 1.0, atypical value for an XBP. For an absorption column density of 3 × 1021 cm−2 , which roughly
corresponds to a distance of 1 kpc in our Galaxy, HRs are −0.89 and 0.36 for the thermal
and power-law model, respectively. For a larger absorption column density of 8 × 1022 cm−2 ,which is equivalent to the Galactic center region, the HR of the thermal model and thepower-law model increases to 0.5 and 0.96, respectively. We therefore conclude that in our
Galaxy, HR values can have a large scatter even for the same source class, depending on thesource distances. On the other hand, according to H i observations made by Luks (1994), we
can safely assume that the interstellar absorption is rather uniform towards the SMC and
LMC with a scale down to 150 (HPBW). In fact, the H i column density derived by Luks
(1994) was ∼ 1 × 1020 –7 × 1021 cm−2 . Within this NH range, HR values show little variation:−0.94 to −0.81 for thin-thermal spectra, and 0.25 to 0.48 for power-law spectra. Thus weconclude that the possible uncertainty in HR due to the NH spatial variation is smaller than
the width of the XBP region and thermal SNR region, hence the present classification by
HR value should be reliable for X-ray sources in the Magellanic Clouds.

7.2.3

Candidates for XBPs and thermal SNRs

In the HR–Lobs plot (figure 7.1), no thermal SNRs are found in the XBP region, nor viceversa. In addition, Crab-like SNRs and BH binaries are not found in these two regions.
Therefore, we can regard sources in the XBP (or thermal SNR) region as candidates for
XBPs (or thermal SNRs). In figure 7.3, we show the HR–Lobs plot for all sources in the
SMC and for the selected LMC sources. We find 20 XBP candidates and four thermal SNR
candidates, which are designated as BPc and TSc in table 5.3, respectively. We excluded
No. 91 from the XBP candidates because it is identified with an AGN. Although the HRs of
No. 2 and No. 13 are −12 , we do not classify them as thermal SNRs because of their large
HR errors. No. 25 also has a large HR error (HR = −1.00 ± 1.38), but we classify this source
as a thermal SNR because its SIS spectrum is typical of this class and because it has a radio
SNR counterpart (0047−735).
Nonpulsating HMXBs and their candidates (table 7.1) are all located within or near the
XBP region, thus are very promising targets for pulsation searches. Non-detection of pulsations from these sources is very likely due to the poor statistics. Detection of pulsations from
2

These sources are not plotted in figure 7.2 and 7.3 because Lobs is not derived (see section 5.4.5).

7.2. CLASSIFICATION BY HARDNESS RATIO123

XBP region

Thermal SNR region

Figure 7.3: Plot of HR against Lobs (0.7–10.0 keV) for all sources in the SMC. Legends for thesymbols are given in the right panel. Sources detected multiple times are represented by single
points, which correspond to the largest Lobs of each source. HR errors for the unclassified sources
are mostly ∼ ±0.2 (see figure 7.2).

low flux sources AX J0049.5−7323, AX J0051.6−7311, and RX J0101.3−7211 demonstrates
well the necessity of high sensitivity observations (see sections 6.1.5, 6.1.8, and 6.1.17).
Two radio SNRs (0046−735 and 0047−735) are found among the thermal SNR candidates. These SNRs have very limited statistics, and thus were not classified as thermal
SNRs in section 5.4.3. Another radio SNR, 0056−725, is located within the XBP region.
Therefore, X-rays from this source may be attributed to an unresolved XBP in the radio
SNR 0056−725, such as 1SAX J0103.2−7209.
Crab-like SNRs, BH binaries, and sources coincident with AGNs or foreground stars
are mostly located between the XBP and thermal SNR regions (except for No. 91 in the
XBP region, but it has a rather large HR error of ±0.26). We find 33 unclassified sources inthis regime; are they located within the SMC, or are they background/foreground sources?
Ueda et al. (1999) carried out the ASCA Medium-Sensitivity Survey and derived the log N –log S relation with a sensitivity limit of 6 × 10−14 erg s−1 cm−2 (0.7–7 keV). It is difficult

124

CHAPTER 7. SOURCE CLASSIFICATION

to obtain an accurate detection limit for our survey because of various uncertainties, suchas contamination from bright sources, and the position-dependent efficiency of the detector.
Therefore, we only make a rough estimation from figure 7.3 that the detection limit forbackground sources in our SMC survey is 10−13 –10−12 erg s−1 cm−2 (4×1034 –4×1035 erg s−1 ).
With this limit, several AGNs would be detected in the GIS FOV in each observation (Ueda
et al. 1999), thus we can expect that several tens of AGNs would be detected in our SMC
survey. Therefore, it is likely that most of the 33 unclassified sources in the medium HR
regime are background AGNs. This conclusion is independently supported by the spatial
distribution of these sources, as shown in section 8.2.
The very hard regime (HR > 0.6) contains seven unclassified sources. Even if a source
has a hard spectrum with Γ = 1.0, a large absorption column of NH > 1022 cm−2 is requiredfor HR to exceed 0.6. Therefore, we regard these seven very hard sources as highly absorbed
objects, although a possibility that they are XBP candidates could not be rejected due tothe large HR error (see figure 7.2).

Chapter 8
Implication for the Star Forming
Activity
8.1

Source Populations in the SMC and in our Galaxy

8.1.1

Basic data

We have shown that most of the SMC sources detected with ASCA are classified as XBPs
(mostly Be-XBPs), nonpulsating HMXBs, and SNRs, in addition to foreground stars andbackground AGNs. In order to make a comparison of source populations in the SMC and in
our Galaxy, we combined the results of our study with various catalogues compiled by otherauthors. The source classes included in this discussion are summarized in table 8.1: HMXBs(XBPs and nonpulsating HMXBs), LMXBs, and SNRs (Crab-like and others).
HMXBs in the SMC consist of 22 XBPs (section 6.1) and eight nonpulsating HMXBs(grades A and B in table 7.1). Candidates for HMXBs are a combination of the 20 XBPcandidates defined in this study (BPc in table 5.3) and the grade C and D sources in table 7.1.
Here we regard the XBPs (and candidates) with an unknown optical counterpart as HMXBs,
because most XBPs hitherto found are HMXBs (e.g., Bildsten et al. 1997). Evidence for
LMXBs has not been detected from any source in the SMC. SNRs in the SMC have beensurveyed by Mathewson et al. (1983; 1984) in the radio band and later by Filipović et al.
(1998a) in both the radio and X-ray bands; as a result 14 SNRs have been detected. Five
SNR candidates have been found by Filipović et al. (1998b) and Haberl et al. (2000). In
125

126

CHAPTER 8. IMPLICATION FOR THE STAR FORMING ACTIVITY

Table 8.1: Source populations in the SMC and in our Galaxy.

GalaxyGalaxy ×1/100∗
SMC———— HMXBs ————XBPs Nonpulse Candidates

LMXBs

53
0.53
22

130
1.3
0

30
0.3
8

...
...
41

———— SNRs ————
Crab-like Other Canditates
10
0.1
0

220
2.2
14

100
1
7

∗: Since the mass of the SMC is about 1/100 of that of our Galaxy, the source numbers in our
Galaxy should be divided by 100 for a simple comparison.

addition, two thermal SNR candidates defined in this study (No. 6 and No. 45) have no radio
SNR counterpart, thus are candidates for new SNRs. No evidence for Crab-like SNRs has
been found.
HMXBs and LMXBs in our Galaxy are catalogued by Liu, van Paradijs, & van den
Heuvel (2000, 2001; BH binaries are also included). We find 49 XBPs, ∼ 30 nonpulsating HMXBs, and ∼ 130 LMXBs in those catalogues. We also include four new XBPs,
AX J1740.1−2847, XTE J1543−568, SAX J2239.3+6116, and AX 1841.0−0536 (Sakano etal. 2000; Finger & Wilson 2000; In ’t Zand et al. 2001; Bamba et al. 2001). Green (2000)
has compiled a most complete catalogue of 225 radio SNRs in our Galaxy. Among them,
about 10 are associated with a rotation-powered X-ray pulsar, thus are regarded as Crab-like
SNRs. In addition, some new SNRs have been discovered in the X-ray band with ROSAT
and ASCA. We tentatively expect 100 SNR candidates to be discovered in the near future(e.g., Aschenbach 1996).

8.1.2

Comparison of the source populations

Based on table 8.1, we compare the source populations in the SMC and in our Galaxy fromvarious aspects.
Since the mass of the SMC is about 1/100th the mass of our Galaxy, the source numbers
in our Galaxy should be divided by 100 for a simple comparison (as in the second row of table
8.1). Accordingly, we find that the normalized number of HMXBs is much higher in the SMC.
On the other hand, the normalized numbers of LMXBs are comparable in the two galaxies.
Since the detection limit of our survey is ∼ 1035 erg s−1 (see figure 7.3), the typical luminosity
ranges of HMXBs and LMXBs has been covered. Hard X-rays from HMXBs brighter than8.1. SOURCE POPULATIONS IN THE SMC AND IN OUR GALAXY

127∼ 1035 erg s−1 can penetrate the thick interstellar matter in our Galaxy. Although LMXBs
have softer spectra, they are mainly distributed at high Galactic latitudes, thus the Galacticinterstellar absorption is again neglectable. Therefore, the discrepancy/similarity of thenormalized numbers of HMXBs/LMXBs are not attributable to a selection effect. LMXBs
are probably much older than HMXBs because they have properties such as weak magnetic
fields and spatial distribution outside active star forming regions (e.g., Lewin, van Paradijs,> 109 yr ago) the star
& Taam 1995). Therefore, we propose that (1) in a very old epoch (∼forming rate in the SMC was similar to that in our Galaxy, which resulted in the similarity ofthe (normalized) LMXB numbers, and then (2) more recently (∼ 106 –107 yr ago) there wasa bursting star formation only in the SMC, which resulted in the much larger (normalized)
HMXB number in the SMC.
According to this proposed star forming history, the normalized number of SNRs should
be higher in the SMC, similar to HMXBs. However, SNRs are rather lacking in the SMC
when compared with HMXBs; the number ratio of SNRs to HMXBs is 14/30 (21/71 when
the candidates are included), while it is 230/83 in our Galaxy (330/83 when the candidates
are included). This fact may indicate that many SNRs in the SMC remain to be discovered.
For example, as mentioned before, No. 6 and No. 45 are thermal SNR candidates with noradio SNR counterpart (the two squares in the thermal SNR region in figure 7.3). Therefore,
follow-up radio observations for these objects may prove fruitful.
Both HMXBs and Crab-like SNRs are direct descendants of massive stars. However, thenumber ratios of these classes in our Galaxy (10/83) and in the SMC (0/30; 0/71 when the
HMXB candidates are included) are significantly different. If the ratio is identical between
galaxies, at least several Crab-like SNRs should be found in the SMC. According to figure 7.3,
Crab-like SNRs should be located around HR ∼ 0, where those in the LMC are located. Since
the unclassified sources in that HR regime are relatively faint, non-detection of pulsations
may be due to limited statistics. Therefore, high sensitivity observations of those sources
with a good time resolution would be fruitful. In addition, high resolution radio surveys tosearch for plerionic emission are also encouraged. If, on the other hand, Crab-like SNRs arereally lacking in the SMC, it may be implied that the binary frequency of massive stars ismuch higher in the SMC because Crab-like SNRs and HMXBs are descendants of single andbinary massive stars, respectively. A higher binary frequency can in part be a cause of thelower number ratio of SNRs to HMXBs mentioned above.

128

CHAPTER 8. IMPLICATION FOR THE STAR FORMING ACTIVITY

8.2

Spatial Distribution of Various Classes of Sources

In order to explore spatial inhomogeneity of the star forming history in the SMC, we investigated the spatial distributions of HMXBs and SNRs. These sources are produced by SNe of
massive stars (except for SNRs from type Ia SNe), thus are good indicators of the recent star
forming history (106 –107 yr ago). In this section, “HMXBs” include the XBPs, nonpulsating
HMXBs, and HMXB candidates in table 8.1, while “SNRs” include the Crab-like and other
SNRs and SNR candidates in table 8.1.
In figure 8.1 (a) and (b), we show the spatial distributions of HMXBs and SNRs. We
find that most HMXBs are concentrated in the optical main body, and ∼ 10% are located in
the eastern wing. Maragoudaki et al. (2000) carried out an optical survey of the SMC andderived spatial distributions of stars in seven ranges of ages (figure 2.4; page 8). Comparing
figure 8.1 (a) and (b) with figure 2.4, we find that the distribution of the HMXBs wellresembles that of stars with ages of 1.2–3 × 107 yr (figure 8.2). This is naturally expected
because HMXBs are descendants of massive (young) stars with an age of ∼ 107 yr. The SNRs
have a distribution pattern similar to HMXBs, thus we suppose that most of the SNRs were
produced by SNe of massive stars.
We also investigated the distribution of the 33 unclassified sources with medium hardness
ratios (−0.6 < HR < 0.2; see section 7.2.3) and six AGN candidates (Nos. 52, 72, 80, 84, 91,
and 92) as shown in figure 8.1 (c). The distribution is relatively uniform, which exhibits aclear contrast with HMXBs and SNRs and is qualitatively consistent with the distributionof old stars (> 2 × 109 yr; see figure 2.4). This fact implies either that (1) most of thosemedium HR sources are unrelated to the SMC, i.e., background or foreground sources, orthat (2) they represent a population much older than the HMXBs. Scenario (1) is consistent
with the estimate from the log N –log S relation made in section 7.2.3. However, taking into
account the large uncertainty in the above estimate, there still remains a possibility thatsome of the medium HR sources represent an older population, probably LMXBs. Since thenumber of LMXBs has a relatively large impact on the star forming activity in the old epoch(see section 8.1.2), follow-up observations of these medium HR sources are encouraged.
The distribution of the seven very hard sources (HR > 0.6) shown in figure 8.1 (d) seems
to be as uniform as that of the AGN candidates and the medium HR sources, although thenumber of sources is highly limited. Therefore, the very hard sources may be unrelated to the
SMC, i.e., background or foreground sources. Considering the suggestion that these sources

8.2. SPATIAL DISTRIBUTION OF VARIOUS CLASSES OF SOURCES
600.5

600.5

(a)

h

m

1 20

h

m

1 00

h

0 40

129

(b)

-71°00′-71°00′-72°00′-72°00′-73°00′-73°00′-74°00′-74°00′

m

0.5
600.5
0.5(c)

h

m

1 20

h

1 00

m

h

0 40

m

0.5
600.5
600.5
0.5(d)

600.5

CHBN SMC_DSS_large.fits0

h

m

1 20

CHBN SMC_DSS_large.fits0

h

m

1 00

0.5
0.5

h

0 40

-71°00′-71°00′-72°00′-72°00′-73°00′-73°00′-74°00′-74°00′

m

h

m

1 20

h

1 00

0.5
600.5
0.5
CHBN SMC_DSS_large.fits0

m

h

0 40

m

600.5
CHBN SMC_DSS_large.fits0

Figure 8.1: Spatial distribution of HMXBs (a: indicated by stars), SNRs (b: circles), the six AGN
candidates and the 33 unclassified sources with medium HR (c: diamonds), and the seven very hardsources (d: squares), superimposed on the DSS image of the SMC. The large and small diamondsin (c) indicate the AGN candidates and the medium HR sources, respectively. The discontinuity
seen in the DSS image is an artifact. Equatorial coordinates in equinox 2000.0 are overlaid.

should be highly absorbed (see section 7.2.3), we propose that most of them are seyfert 2galaxies. We thus encourage optical follow-up observations of these sources.

130

CHAPTER 8. IMPLICATION FOR THE STAR FORMING ACTIVITY

600.5-71°00′-72°00′-73°00′-74°00′
h

1 20

m

h

1 00

0.5
0.5

m

h

m

0 40

600.5

Figure 8.2: Spatial distribution of HMXBs (indicated by star-marks) and young stars (∼ 2×107 yr;
CHBN SMC_DSS_large.fits0

contours) overlaid with the DSS image of the SMC. The discontinuity seen in the DSS image is an
artifact. Equatorial coordinates in equinox 2000.0 are overlaid.

Chapter 9
Conclusion
We performed systematic analyses on all the ASCA data of the SMC region. We detected
106 sources with a criterion of S/N > 5. Basic properties such as positions, correlation withother catalogues, and spectral parameters are compiled in the catalogue. We determined thesource positions with Gotthelf et al.’s correction, and found that the 90% error radius is inpractice ∼ 4000 for sources detected in the inner circle of the GIS.
We discovered eight new X-ray pulsars and also determined precise positions for four
already known X-ray pulsars. As a result, there are now 25 X-ray pulsars known in the
SMC. We regarded 22 pulsars among the 25 as XBPs because of their pulse period, hard
spectrum, flux variability, and, for some sources, the existence of an optical counterpart.
Among the 14 radio SNRs known in the SMC, eight were detected with ASCA. We
detected emission lines of highly ionized atoms from five SNRs, and thus regarded them asthermal SNRs.
We established a simple method of source classification using the spectral hardness< HR < 0.6 andratio (HR). XBPs and thermal SNRs respectively fall in a range of 0.2 ∼∼< HR < −0.6. Other source classes such as Crab-like SNRs, BH binaries, AGNs,−1.0 ∼∼
and foreground stars have medium hardness ratios (−0.6 < HR < 0.2). This method is a
powerful tool for classifying faint sources which remain unclassified. With this method, we
found 20 XBP candidates and four thermal SNR candidates.
We combined the results of this study with those of other catalogues to compare thesource populations in the SMC and in our Galaxy. Normalized by the galaxy mass, the
131

132

CHAPTER 9. CONCLUSION

number of LMXBs is consistent between the SMC and our Galaxy, while the number of
HMXBs is very much higher in the SMC. This contrast indicates that in a very old epoch
the star forming activity was comparable between the two galaxies, and that very recently
there was a bursting star formation only in the SMC. SNRs in the SMC are not as populous
as HMXBs, although the normalized number is still larger than that in our Galaxy. This
may indicate that there are many SNRs in the SMC that remain to be discovered.
We derived the spatial distributions of HMXBs and SNRs, including candidates. We
found that HMXBs are distributed mainly along the main body and partly in the eastern< 2 × 107 yr) stars. SNRs also have awing; this quite resembles the distribution of young (∼similar distribution, indicating that most of the SNRs are remnants of type II SNe. Medium
HR sources (−0.6 < HR < 0.2) have, on the other hand, a relatively regular distribution.
This implies that most of them are AGNs, which is also supported by the log N –log S relation
derived from the ASCA Medium-Sensitivity Survey.

Appendix A
Two Color X-Ray Mosaic Image of
the SMCThe mosaic images shown in figure 5.1 and 5.2 are combined into a two color image infigure A.1. Red and blue indicate X-ray photons in the soft (0.7–2.0 keV) and hard (2.0–
7.0 keV) bands, respectively. Most of the “blue” or “red” sources are XBPs or thermal
SNRs (including their candidates), respectively. Evidently many XBPs and the candidates
are found, which is in a strong contrast between the LMC color image presented by Nishiuchi(2001).

133

134

APPENDIX A. TWO COLOR X-RAY MOSAIC IMAGE OF THE SMC

Figure A.1: Two color X-ray mosaic image of the SMC. Electronic version of this image can beobtained from http://www-cr.scphys.kyoto-u.ac.jp/member/jun/job/phd/.

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Acknowledgment
I am deeply grateful to Prof. K. Koyama for his continuous guidance throughout the five
years of my graduate course and his leadership in the ASCA SMC survey project. I also
thank my collaborators, Mr. K. Imanishi, Mr. M. Tsujimoto, and Dr. M. Nishiuchi, for theirhelpful comments and sharing in data analyses throughout the survey project.
Some of new pulsars have not been discovered without efforts made by Dr. K. Torii, Mr.
M. Ueno, and Mr. T. Kohmura. Collaboration with Prof. F. Nagase, Dr. B. Paul, and Dr.
M. Ozaki in the analysis of XTE J0111.2−7317 was very good experience. Comments fromDr. N. Mizuno and Dr. K. Yoshita were helpful in the analysis of SNRs.
I thank all the members of the Cosmic-ray laboratory in Kyoto University for the daily
discussion and insightful comments. Of course I thank you, who are interested in this thesis
and have just finished reading it!

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