X-ray Study of the Intermediate Mass
Young Stars Herbig Ae/Be Stars
Hamaguchi Kenji
January 5, 2001

Abstract
We study the x-ray property of intermediate mass pre-main-sequence stars (PMSs) Herbig
Ae/Be stars (HAeBes), mainly using the observational and archival data of the ASCA andROSAT satellites. Among 39 HAeBes covered with 42 ASCA pointings, we detect 15 plausible x-ray counterparts. X-ray luminosities of the detected sources are in a high range of LX∼1030−32 ergs s−1 compared with those of low mass stars. This fact excludes the possibility
that HAeBe x-ray emission comes from a low mass companion. Most sources with moderate
photon counts (>1000 cts) show time variations. In particular, three HAeBes - MWC 297,
TY CrA and VY Mon - exhibit ﬂare-like time variations with long decay timescales (efolding time ∼ 10−50 ksec). These ﬂare shapes are quite similar to the ﬂares of low mass
protostars. In contrast, HR 6000 show a rapid variation with fast rise and fast decay (τ <2ksec). Each spectrum is successfully reproduced with an absorbed one or two temperaturethin-thermal plasma (MeKaL) model. The temperature ranges between kT ∼ 1−5 keV,which is signiﬁcantly higher than that of main-sequence OB stars (kT < 1 keV).
For their x-ray origin, we test a few models for emission mechanisms of stellar x-rays wind driven shock, wind-fed magnetosphere, accretion shock and magnetic activity. Thesemodels except for the magnetic activity cannot explain the plasma temperature, x-ray luminosity or time variability. Nevertheless, HAeBes are not thought to have solar type dynamo
mechanism to amplify the magnetic ﬁeld because they lack the convection zone on the stellar
surface. We cannot ﬁnd the LX dependence on v sini and ω sini, that are indicators of the solar type dynamo. To ﬁnd out the generating mechanism of the magnetic ﬁeld, we investigate
the evolutional variation of plasma temperature and x-ray luminosity. We then ﬁnd that theplasma temperatures have a clear rising trend with increasing the absorption column density(NH ), which traces the surrounding materials and hence the evolutional status of young stars.
This indicates that younger stars have more violent high-energy activity. Whereas, the x-ray
luminosity seems to depend more on stellar masses. We therefore plot the x-ray luminosity
in the HR diagram together with additional samples in previous papers. We then ﬁnd LX
drops after the age a few×106 yrs. Considering the circumstellar structure, our results favor
the mechanism of large scale magnetic ﬁelds linking a star and its circumstellar disk.
Stars earlier than spectral type B reach to the main-sequence phase before a few ×106
year-old. This indicates that young massive main-sequence stars still operate the magnetic
activity. From several pieces of supporting evidence, we propose that the x-ray emission
from B1−9 stars and hard tails seen in O stars originate in this magnetic activity. On theother hand, we ﬁnd the hard x-ray emission from the giant molecular cloud core Monoceros
R2. We suggest that it might come from progenitors of HAeBes. Finally, we draw a uniﬁed
picture of x-ray activity on young stellar objects in view of the magnetic activity.

Contents
1 Introduction

12 Review

3

2.1 Herbig Ae/Be Stars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

3

2.1.1

Classiﬁcation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

3

2.1.2

Circumstellar Structure . . . . . . . . . . . . . . . . . . . . . . . . . .

4

2.1.3

Evolutional Status . . . . . . . . . . . . . . . . . . . . . . . . . . . .

7

2.1.4

Other Types of Stars Related to Herbig Ae/Be Stars . . . . . . . . .

8

2.2 Stellar X-ray Emission . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

9

2.2.1

X-ray Emission from Main-Sequence Stars . . . . . . . . . . . . . . .

9

2.2.2

X-ray Emission from Pre-Main-Sequence Stars . . . . . . . . . . . . .

123 Instruments

15

3.1 The ASCA Satellite . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

15

3.1.1

Outline. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

15

3.1.2

X-ray Telescope (XRT) . . . . . . . . . . . . . . . . . . . . . . . . . .

17

3.1.3

Solid-state Imaging Spectrometer (SIS) . . . . . . . . . . . . . . . . .

18

3.1.4

Gas Imaging Spectrometer (GIS) . . . . . . . . . . . . . . . . . . . .

20

3.2 The ROSAT Satellite . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

21

3.2.1

Outline. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

21

3.2.2

X-ray Mirror Assembly (XMA) . . . . . . . . . . . . . . . . . . . . .

22v

CONTENTS

vi
3.2.3

Position Sensitive Proportional Counter (PSPC) . . . . . . . . . . . .

22

3.2.4

High Resolution Imager (HRI) . . . . . . . . . . . . . . . . . . . . . .

22

4 Observation & Data Reduction
4.1 Target Selection . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

23
23

4.1.1

Data Retrieval. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

23

4.1.2

Quantities of Fundamental Stellar Parameters . . . . . . . . . . . . .

24

4.2 Reduction & Screening Methods of the ASCA Data . . . . . . . . . . . . . .

29

4.2.1

RDD Correction. . . . . . . . . . . . . . . . . . . . . . . . . . . . .

29

4.2.2

Filtering Criteria . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

29

4.2.3

Telemetry Saturated Data . . . . . . . . . . . . . . . . . . . . . . . .

29

4.2.4

Event Extraction . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

30

4.3 Basic Procedure of Data Analyses . . . . . . . . . . . . . . . . . . . . . . . .

30

4.3.1

Imaging Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

31

4.3.2

Timing Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

32

4.3.3

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

33

5 Result I – Herbig Ae/Be Stars
5.1 Detected Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

35
35

5.1.1

V892 Tau . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

35

5.1.2

V380 Ori . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

39

5.1.3

VY Mon & CoKu VY Mon/G2 . . . . . . . . . . . . . . . . . . . . .

39

5.1.4

T Cha . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

42

5.1.5

HD 104237. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

43

5.1.6

IRAS 12496−7650 . . . . . . . . . . . . . . . . . . . . . . . . . . . .

45

5.1.7

HR 5999 / HR 6000 . . . . . . . . . . . . . . . . . . . . . . . . . . .

48

5.1.8

V921 Sco . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

52

5.1.9

MWC 297 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

53

CONTENTS

vii

5.1.10 EC 95 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

55

5.1.11 S CrA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

57

5.1.12 TY CrA / HD 176386 . . . . . . . . . . . . . . . . . . . . . . . . . .

59

5.1.13 HD 200775. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

67

5.1.14 MWC 1080 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

69

5.2 Non Detected Sources. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

71

5.2.1

BD+30◦ 549 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

71

5.2.2

CoKu Tau 1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

71

5.2.3

AB Aurigae . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

71

5.2.4

MR Ori / V361 Ori / T Ori . . . . . . . . . . . . . . . . . . . . . . .

72

5.2.5

BF Ori . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

73

5.2.6

Z CMa / LkHα 218 / LkHα 220 . . . . . . . . . . . . . . . . . . . . .

74

5.2.7

HD 97048 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

74

5.2.8

HD 97300 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

78

5.2.9

HD 147889. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

78

5.3 Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

79

6 Result II – Giant Molecular Cloud
6.1 Monoceros R2 Giant Molecular Cloud . . . . . . . . . . . . . . . . . . . . . .
7 Discussion I – X-ray Emission Mechanism

87
87
91

7.1 Could the X-ray Emission Come from a Low Mass Companion? . . . . . . .

91

7.2 X-ray Emission Mechanism . . . . . . . . . . . . . . . . . . . . . . . . . . . .

92

7.2.1

Stellar Wind . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

92

7.2.2

Accretion Shock . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

93

7.2.3

Wind-fed Magnetosphere . . . . . . . . . . . . . . . . . . . . . . . . .

95

7.2.4

Magnetic Activity on the Stellar Surface . . . . . . . . . . . . . . . .

95

7.3 X-ray Properties Related to the Magnetic Activity . . . . . . . . . . . . . . .

95

CONTENTS

viii
7.3.1

LX vs. Lbol Relation . . . . . . . . . . . . . . . . . . . . . . . . . . .

97

7.3.2

Standard Dynamo Mechanism . . . . . . . . . . . . . . . . . . . . . .

97

7.3.3

Prominent X-ray Variability on HAeBes . . . . . . . . . . . . . . . . 101

7.3.4

kT vs. E.M. Relation . . . . . . . . . . . . . . . . . . . . . . . . . . . 104

7.4 Evolution of Stellar X-ray Activity . . . . . . . . . . . . . . . . . . . . . . . 105
7.4.1

NH vs. kT and LX Relation. . . . . . . . . . . . . . . . . . . . . . . 105

7.4.2

Equivalent Width of Hα vs. LX Relation . . . . . . . . . . . . . . . . 107

7.4.3

Evolutional Trends of kT and LX in the HR Diagram . . . . . . . . . 107

7.5 Models to Generate the Magnetic Field on HAeBes . . . . . . . . . . . . . . 108
8 Discussion II – A Uniﬁed Picture of Stellar X-ray Activity

113

8.1 Evolution of the X-ray Activity to the Main-sequence Phase . . . . . . . . . 113
8.1.1

X-ray Emission from Late B Type Stars . . . . . . . . . . . . . . . . 113

8.1.2

X-ray Emission from O and Early B Type Stars . . . . . . . . . . . . 115

8.2 X-ray Emission from Giant Molecular Cloud Cores . . . . . . . . . . . . . . 116
8.3 A Uniﬁed Picture of X-ray Activity on YSOs . . . . . . . . . . . . . . . . . . 117
9 Conclusion & Future Problem

121

A Models of the Stellar Magnetic Activity

125

A.1 Hydrodynamic Decay, Sustained Heating Model . . . . . . . . . . . . . . . . 125
A.2 Shear Dynamo Model . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 127
A.3 Diﬀusion Time Scale of the Magnetic Field . . . . . . . . . . . . . . . . . . . 128
B Stellar Physical Parameters

129

List of Figures
2.1 Schematic picture of the HAeBe system . . . . . . . . . . . . . . . . . . . . .

4

2.2 SEDs of three HAeBes in diﬀerent groups . . . . . . . . . . . . . . . . . . . .

5

2.3 UV spectra around Lyman α emission line . . . . . . . . . . . . . . . . . . .

6

2.4 Schematic picture of deuterium burning in protostars . . . . . . . . . . . . .

7

2.5 Evolutional status of each group of stars . . . . . . . . . . . . . . . . . . . .

8

2.6 An x-ray ﬂare from the sun . . . . . . . . . . . . . . . . . . . . . . . . . . .

10

2.7 Log LX and log LX /Lbol vs. eﬀective temperature . . . . . . . . . . . . . . .

11

2.8 Shematic views of two x-ray emission models of massive MS stars . . . . . .

12

2.9 Schematic view of possible magnetic conﬁguration between a star and its disk

13

3.1 Schematic view of the ASCA satellite . . . . . . . . . . . . . . . . . . . . . .

16

3.2 Point spread function of the ASCA XRT for four diﬀerent oﬀ-axis angles. . .

17

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

19

3.4 Schematic view of the GIS sensor system . . . . . . . . . . . . . . . . . . . .

21

3.5 Schematic view of the ROSAT satellite . . . . . . . . . . . . . . . . . . . . .

22

5.1 ASCA images of V892 Tau . . . . . . . . . . . . . . . . . . . . . . . . . . . .

36

5.2 SIS light curves of V892 Tau . . . . . . . . . . . . . . . . . . . . . . . . . . .

37

5.3 SIS spectra of V892 Tau . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

37

5.4 ASCA GIS and ROSAT HRI images of V380 Ori . . . . . . . . . . . . . . .

38

5.5 GIS light curves & spectrum of V380 Ori . . . . . . . . . . . . . . . . . . . .

39

5.6 ASCA GIS and ROSAT PSPC images of VY Mon . . . . . . . . . . . . . . .

40ix

LIST OF FIGURESx

5.7 GIS light curves of VY Mon . . . . . . . . . . . . . . . . . . . . . . . . . . .

41

5.8 GIS spectra of VY Mon . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

42

5.9 ASCA images of T Cha . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

43

5.10 Light curves of T Cha . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

44

5.11 Spectra of T Cha . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

44

5.12 ASCA images of HD 104237 . . . . . . . . . . . . . . . . . . . . . . . . . . .

45

5.13 Light curves of HD 104237 . . . . . . . . . . . . . . . . . . . . . . . . . . . .

46

5.14 Spectra of HD 104237. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

46

5.15 ASCA GIS image of IRAS 12496-7650 . . . . . . . . . . . . . . . . . . . . .

47

5.16 Light curves & spectrum of IRAS 12496–7650 . . . . . . . . . . . . . . . . .

47

5.17 ASCA images of HR 5999 and HR 6000 . . . . . . . . . . . . . . . . . . . . .

49

5.18 Light curves of HR 5999 and HR 6000 . . . . . . . . . . . . . . . . . . . . .

50

5.19 Spectra of HR 5999 and HR 6000 . . . . . . . . . . . . . . . . . . . . . . . .

51

5.20 ASCA images of V921 Sco . . . . . . . . . . . . . . . . . . . . . . . . . . . .

52

5.21 Light curves and spectra of V921 Sco . . . . . . . . . . . . . . . . . . . . . .

53

5.22 ASCA images of the MWC 297 ﬁeld . . . . . . . . . . . . . . . . . . . . . . .

54

5.23 Light curves of MWC 297 . . . . . . . . . . . . . . . . . . . . . . . . . . . .

55

5.24 Spectra of MWC297 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

56

5.25 ASCA images of the EC 95 ﬁeld . . . . . . . . . . . . . . . . . . . . . . . . .

57

5.26 light curves of EC 95 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

58

5.27 Spectra of EC 95 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

58

5.28 ASCA GIS Images of the S CrA ﬁeld . . . . . . . . . . . . . . . . . . . . . .

59

5.29 GIS light curves of S CrA . . . . . . . . . . . . . . . . . . . . . . . . . . . .

60

5.30 Spectra of S CrA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

60

5.31 ASCA images of the TY CrA/HD 176386 ﬁelds . . . . . . . . . . . . . . . .

61

5.32 Light curves of TY CrA and HD 176386 . . . . . . . . . . . . . . . . . . . .

61

5.33 Spectra of TY CrA and HD 176386 . . . . . . . . . . . . . . . . . . . . . . .

62

LIST OF FIGURES

xi

5.34 Conﬁdence contour of the spectrum ﬁtting parameters of HD 176386 (1stobservation) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

63

5.35 GIS light curves of TY CrA + HD 176386 . . . . . . . . . . . . . . . . . . .

64

5.36 GIS spectra of TY CrA and HD 176386 . . . . . . . . . . . . . . . . . . . . .

65

5.37 ASCA and ROSAT images of the HD 200775 ﬁeld . . . . . . . . . . . . . . .

66

5.38 Light curves of HD 200775 . . . . . . . . . . . . . . . . . . . . . . . . . . . .

67

5.39 Spectra of HD 200775. . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

68

5.40 ASCA images of the MWC 1080 ﬁelds . . . . . . . . . . . . . . . . . . . . .

69

5.41 Light curves of MWC 1080 . . . . . . . . . . . . . . . . . . . . . . . . . . . .

70

5.42 Spectra of MWC 1080 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

70

5.43 ASCA SIS images of the AB Aur and T Ori ﬁelds . . . . . . . . . . . . . . .

72

5.44 ASCA GIS and ROSAT images of the BF Ori ﬁeld . . . . . . . . . . . . . .

73

5.45 Spectrum of an x-ray source in the BF Ori ﬁeld . . . . . . . . . . . . . . . .

74

5.46 ASCA images of the HD 97048 ﬁeld . . . . . . . . . . . . . . . . . . . . . . .

76

5.47 Spectra of HM16 in ChamI IV . . . . . . . . . . . . . . . . . . . . . . . . . .

77

5.48 ASCA and ROSAT images of the HD 147889 ﬁeld . . . . . . . . . . . . . . .

78

5.49 ASCA spectrum of the HD 147889 region . . . . . . . . . . . . . . . . . . . .

79

5.50 Number of detected sources in the ASCA survey . . . . . . . . . . . . . . . .

81

5.51 Distance vs. x-ray luminosity relation . . . . . . . . . . . . . . . . . . . . . .

81

6.1 ASCA images of the Monoceros R2 giant molecular cloud core . . . . . . . .

88

6.2 Hard band light curves of the Monoceros R2 core (4–10 keV) . . . . . . . . .

89

6.3 Quiescent and ﬂare state spectra of the Monoceros R2 core . . . . . . . . . .

89

7.1 Wind kinetic luminosity and x-ray luminosity relation . . . . . . . . . . . . .

94

7.2 Terminal velocity (v∞ ) vs. plasma temperature relation . . . . . . . . . . . .

94

7.3 LX vs. Lbol relation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

98

7.4 Teﬀ vs. LX relation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

98

7.5 v sin i vs. LX relation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

99

LIST OF FIGURES

xii

7.6 v sin i vs. log LX /Lbol relation . . . . . . . . . . . . . . . . . . . . . . . . . .

99

7.7 Ω sin i vs. log LX relation . . . . . . . . . . . . . . . . . . . . . . . . . . . . 100
7.8 Ω sin i vs. log LX /Lbol relation . . . . . . . . . . . . . . . . . . . . . . . . . . 100
7.9 Variation of emission measure during ﬂares . . . . . . . . . . . . . . . . . . . 102
7.10 Variation of plasma temperature during ﬂares . . . . . . . . . . . . . . . . . 102
7.11 Assumed geometry of the ﬂare magnetic ﬁeld . . . . . . . . . . . . . . . . . . 103
7.12 kT vs. E.M. relation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 104
7.13 kT dependence on NH . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 106
7.14 LX dependence on NH. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 106

7.15 W(Hα) vs. log LX relation . . . . . . . . . . . . . . . . . . . . . . . . . . . . 107
7.16 X-ray luminosity and plasma temperature of the HAeBe samples in the HR
diagram . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 109
7.17 X-ray luminosity of the young stars in the HR diagram . . . . . . . . . . . . 110
7.18 Radio continuum (λ1.3mm ) vs. LX relation . . . . . . . . . . . . . . . . . . . . 112
8.1 LX vs. Lbol relation of OB stars and Herbig Ae/Be stars. . . . . . . . . . . 114

8.2 kT vs. LX relation of MS massive stars and Herbig Ae/Be stars . . . . . . . 116
8.3 kT distribution of giant molecular cloud and Herbig Ae/Be stars . . . . . . . 117
8.4 Uniﬁed picture of the x-ray emission from stars . . . . . . . . . . . . . . . . 119
A.1 Evolution of the MWC 297 ﬂare parameters in the log√

E.M. - log T plane

126

A.2 Evolution of relative x-ray luminosity proposed by Tout et al. (1995). . . . . 127

List of Tables
2.1 X-ray properties of giant molecular cloud cores . . . . . . . . . . . . . . . . .

14

3.1 SIS and GIS detector performances . . . . . . . . . . . . . . . . . . . . . . .

18

4.1 Stellar parameters of the observed Herbig Ae/Be stars. . . . . . . . . . . .

25

4.2 ASCA observation log . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

26

4.3 Observing parameters of both SIS & GIS . . . . . . . . . . . . . . . . . . . .

27

4.4 Calibration ﬁles used for the RDD correction . . . . . . . . . . . . . . . . . .

28

4.5 Data selection criteria . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

30

4.6 ROSAT observation log . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

31

4.7 RDF parameters of the ROSAT data . . . . . . . . . . . . . . . . . . . . . .

32

4.8 Flux converting ratio from the ROSAT and Einstein band to the ASCA band33

4.9 Speciﬁc subjects related to observations and analyses . . . . . . . . . . . . .

34

5.1 ASCA detected sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

82

5.2 ASCA non-detected sources . . . . . . . . . . . . . . . . . . . . . . . . . . .

83

5.3 Selected regions for event extraction . . . . . . . . . . . . . . . . . . . . . . .

84

5.4 Result of ASCA timing analyses . . . . . . . . . . . . . . . . . . . . . . . . .

85

5.5 Result of ASCA spectral analyses . . . . . . . . . . . . . . . . . . . . . . . .

86

6.1 Spectral parameters of the x-ray emission from the Monoceros R2 core . . .

90

7.1 Physical Parameters of Herbig Ae/Be stars . . . . . . . . . . . . . . . . . . .

96

xiii

xiv

LIST OF TABLES
7.2 Other HAeBes . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

97

7.3 Fitting results of ﬂare light curves . . . . . . . . . . . . . . . . . . . . . . . . 101
7.4 Physical parameters of the MWC 297 ﬂare . . . . . . . . . . . . . . . . . . . 103
7.5 Young stars in the Orion cloud (Gagne et al. 1995) . . . . . . . . . . . . . . 111
8.1 X-ray characteristics of massive main-sequence stars . . . . . . . . . . . . . . 115
A.1 Physical parameters of the MWC 297 ﬂare in the hydrodynamic model . . . 126
B.1 Stellar parameters . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 130
B.2 ZAMS age & MS life time . . . . . . . . . . . . . . . . . . . . . . . . . . . . 130

Chapter 1
Introduction
Intermediate mass stars ranging from 2M to 10M do not exhibit x-ray activity. This isdue to the absence of physical processes which produce x-ray photons such as strong UV
acceleration to drive high-speed stellar winds working in high mass stars and the surface
dynamo to drive magnetic activity working in low mass stars. In the pre-main-sequencestage which is classiﬁed as Herbig Ae/Be stars (HAeBes), the UV ﬁelds are much weakbecause the temperature of the stellar surface is low, but the structure of the stellar inside
is considered to be almost the same as that of the main-sequence phase. No x-ray activity,
therefore, could be working on HAeBes according to the current theories.
Nevertheless the ROSAT satellite detected the x-ray emission from nearly half the members of HAeBes (Zinnecker and Preibisch 1994). Subsequent ASCA spectroscopy with hard
x-ray eﬃciency demonstrated that x-ray plasma have quite a high temperature (kT ∼3 keV)in certain objects (Skinner and Yamauchi 1996; Yamauchi et al. 1998; Hamaguchi et al.
2000). More surprisingly, Hamaguchi et al. (2000) discovered a large x-ray ﬂare from the
HAeBe, MWC 297. These latest results strongly indicate the presence of high-energy activity on or around HAeBes. The x-ray activity is thought to drive outﬂow and jet phenomena
seen around young stellar objects. X-ray study of HAeBes is thus crucial to grasp the originof the high-energy activity and its eﬀect on the circumstellar environments.
In this thesis, we perform a HAeBes survey using observational and archival data of the
ASCA satellite. From the ﬁrst systematic analyses of the spectral and timing information,
we restrict the x-ray emission mechanism probable in the HAeBe system. HAeBes are characterized by a dynamically evolving system to the main-sequence phase. We investigate theevolutional dependence of the x-ray activity as well, by merging results from the ROSAT and
Einstein satellites. This research gives us an important implication for the x-ray propertiesof post HAeBes, massive main-sequence stars, and pre HAeBes, protostars. By combining
all these results, we ﬁnally draw an overall picture of stellar x-ray activity from protostars
to the main-sequence phase.
1

2

CHAPTER 1. INTRODUCTION

This thesis consists of eight chapters. Chapter 2 describes a review of the HAeBe system
and x-ray property of known stars. Chapter 3 describes an instrument overview about ASCA
and ROSAT. The next two chapters contain the data analyzing procedure (Chapter 4) andthe results of all individual sources (Chapter 5). Chapter 6 & 7 discuss x-ray emission
mechanism of HAeBes, pre and post HAeBes. The conclusion of our survey is in the lastchapter (Chapter 8).

Chapter 2
Review
2.1

Herbig Ae/Be Stars

2.1.1

Classiﬁcation

Herbig (1960) made a catalogue of emission line objects which are candidate stars in thecontracting phase between 3−20 M . The selection criteria in Herbig (1960) are as follows:§1. spectral type A or B with emission lines,§2. locating in an obscured region,§3. accompanying a bright nebulousity.
The classiﬁed sources have been widely accepted as pre-main-sequence stars (PMSs) of intermediate mass, ie high mass counterparts of T-Tauri stars1 (TTSs). They are now called
Herbig Ae/Be stars. Recent sensitive surveys in near and far infrared wavelengths presented
a few more candidates of intermediate mass PMSs which are not associated with nebulousity.
To include these stars in the category, the criteria of Herbig Ae/Be stars were slightly modiﬁed as follows (cf. Waters and Waelkens 1998),§1. spectral type A or B with emission lines,§2. infrared (IR) excess due to hot or cool circumstellar dust or both,§3. luminosity class III to V.
These new criteria are convenient to exclude classical Be stars and supergiants, which have
similar photometric appearances to Herbig Ae/Be stars. Newly identiﬁed objects are especially called ’isolated Herbig Ae/Be stars’. We later abbreviate Herbig Ae/Be star to1

see Tsuboi (1999) for the evolution of low mass young stars

3CHAPTER 2. REVIEW

4

Herbig-Haro (HH)

Jet, Outflows
Gas Envelopes

Figure 2.1: Schematic picture of
the HAeBe system. This picture
does not hold relative scale size.

Accretion Flow

SecondaryStar’HAeBe’ according as an usual manner.

2.1.2

Circumstellar Structure

Number of HAeBes observable in the sky is relatively small. This is because the birth ratio of
HAeBes is not high and their PMS life time is about two orders as short as those of solar mass
stars (∼106 yrs). These constraints make it diﬃcult to ﬁgure out the circumstellar structureof HAeBes. However, infrared and radio observations have indicated several similarities
between HAeBes and TTSs.
A schematic picture of the HAeBe system is displayed in Figure 2.1. The HAeBe system
consists of a stellar main body, circumstellar disk and spherical gas envelope. Circumstellar
materials in the disk gradually accrete onto the stellar body, and some of them are eliminated from the system seen as jets or mass outﬂows, which take away most of the angularmomentum holding in the accreting gases. They collide with the ambient materials and form
bow-shocks called Herbig-Haro objects. Most of the HAeBes (P>85%, Pirzkal et al. 1997)
accompanies companion stars, some of which are also HAeBes. Here we describe the detailed
properties of each component.

2.1. HERBIG AE/BE STARS

5

Figure 2.2: SEDs of three HAeBes in diﬀerent groups (Hillenbrand et al. 1992).
Circumstellar Environments - Disks and Envelopes
Hillenbrand et al. (1992) classiﬁed spectral energy distribution (SED) of HAeBes into threegroups according to the spectral shapes above λ ∼ 2.2µm. By interpreting these SED shapes,
they suggested each group has the following blackbody components (see also Figure 2.2),
Group I: star + geometrically ﬂat, optically thick circumstellar accretion disk,
Group II: star + gas and dust which is not conﬁned to a disk (envelope),
Group III: star without any excess emission.
Group III stars contain the highest mass stars and group I stars contain the next highest
mass stars. They therefore concluded that the disks and envelopes of higher mass stars
have shorter lifetimes than those of lower mass stars. Nevertheless, Hartmann et al. (1993)
pointed out that the infrared excess regarded as the disk component can be explained byemission from spherical envelopes. Natta et al. (1993) also suggested the presence of envelope
structure even for the Group I sources in the far infrared. Recent understanding is that the
SEDs of Group I sources require the existence of both disks and envelopes. On the other
hand, with the 1.3 mm radio continuum, Henning et al. (1998) evaluated the average densityof the HAeBe cores to be 105−8 cm−3 and those of the extended envelopes to be 104−5 cm−3.
Fuente et al. (1998) found that the surrounding gases within 0.08 pc disperse in the timescale of 106 yrs.
Accretion Flows
Hillenbrand et al. (1992) deduced mass accretion rate of HAeBes to be high (Macc ∼ 10−4−−6M yr−1 ) from excess emission over the standard disk blackbody. Skinner et al. (1993) andHartmann et al. (1993), however, rebuted this result because the strong emission requiredfor high accretion rate was not observed in the radio wavelength. UV excess is also a good

6

CHAPTER 2. REVIEW

Figure 2.3: UV spectra around Lyman α
emission line of HD 104237 (swp 41589,
36800, 36782), HD 163296 (swp 28844,
8003), HR 5999 (swp 26370, 9974), AB
Aur (swp 4907) and HD 190073 (swp15499) (Blondel et al. 1993). The strongline on 1216Å come from the geocoronal
Lyman α emission. A redshift componentof Lyman α emission (λ ∼ 1218 Å) is seen
in younger HAeBes (upper) but it is absent in elder HAeBes (lower).

estimator of mass accretion rate (Meeus et al. 1998; Valenti et al. 2000; Gullbring et al.
2000), but the estimate is not easy for HAeBes because UV radiation in stellar blackbody
is dominant and moreover, the extinction of ISM strongly depends on AV and adopted
conversion ratio between AV and NH . Blondel et al. (1993) derived low accretion rate
(Macc ∼ 10−6−−8 M yr−1 ) for Herbig Ae stars from red-shifted Lyman α emission lines
using the IUE data (see Figure 2.3). This value is as many as those of T-Tauri stars but 3
orders of magnitude lower than those of active stars such as FU Ori and Z CMa. Assuming
accretion geometry, they also suggested the faster infall velocity (vacc ∼ 600−2000 km s−1 )than the free-fall velocity (vf f ∼ 600 km s−1 ). This requires a certain acceleration mechanismin the inﬂow process.
Outﬂows & Stellar winds
Herbig-Haro (HH) outﬂows, including jets and CO outﬂows, are usually seen around YSOs
(eg Levreault 1988). Outﬂows in HAeBes are qualitatively identical to, but a factor of two
to three faster than, those of TTSs (v ∼600−900 km s−1 , Mundt and Ray 1994). Outﬂow
activity is usually seen in young HAeBes in 105 -year-old, but collimated, bipolar outﬂow was
also found from the elder HAe star HD 163296, whose age (1.5−10 Myrs) is ∼10 times older
than the age of TTSs showing outﬂow activity (Devine et al. 2000). Most HAeBe outﬂows

2.1. HERBIG AE/BE STARS

7

Figure 2.4: Schematic picture of deuterium burning in protostars (Palla et al. 1990). (a)At low masses, accreted deuterium is fed in a steady-state fashion to the burning centerby convective eddies. (b) Eventually, a radiative barrier appears, cutting oﬀ the supply tothe center. (c) The central region soon exhausts its deuterium and reverts to a radiatively
stable state. (d) At higher mass, deuterium ignites as a shell source just outside the depleted
central region.
(≈70%) are highly, but the rest are poorly, collimated (Mundt and Ray 1994).
From P-Cygni proﬁles of Hα and NaD lines, Finkenzeller and Mundt (1984) indicatedthat the stellar wind was accelerated within a few stellar radii and decelerated to the terminal
velocity at larger distances (eg 20−50 R∗ ). Valenti et al. (2000) detected broad absorption
features in UV spectra (eg C, O and Si), which are originated in wind materials. Thesefeatures have a variety in strength even for stars with the same spectral type: some areblue-shifted and are not seen in early HBes. This absence on early HBe stars is interpreted
as the highly ionization of circumstellar materials. On the other hand, the radio emission
from HAeBes, which probably comes from stellar winds, is mostly thermal except for a few
sources (eg TY CrA and EC95, Skinner et al. 1993; Smith et al. 1999).

2.1.3

Evolutional Status

In theoretical models a star is born in the upper right of the HR diagram, ie low Teﬀ &large Lbol . It descends along nearly vertical track (Hayashi track) and then moves along
a radiative track at almost constant luminosity when a radiative core develops. Finally,hydrogen ignites in the core, contraction is halted, and a star reach to the main-sequencephase (zero-age-main-sequence: ZAMS). Iben (1965) made numerical calculations of youngstellar evolution in the HR diagram. He treated stars of masses ranging between 0.5−150M ,on the assumption that the initial state of stars is fully convective, with radii nearly twoorders of magnitude larger than that of ZAMS. However, such a large amount of gases is
not thought to be gathered in the initial phase. Palla and Stahler (1990) therefore set thestellar birthline much closer to the main-sequence branch. In their scenario, stars just below

CHAPTER 2. REVIEW

8
B

F

A

G

K Sp. Type

Protostars
HLYSOsPMS

ZAMS

HAeBes
ρO

ph

UXORs

S1

TTSs

βP

MS

ic

Ma

gne

tic

Bp

Ap

?

Age

Figure 2.5: Evolutional status of each group of stars.
the birthline are fully convective, but they soon ignite deuterium burning, which makes thebarrier to penetrate the convective eddy into the core and suppress the convection (see also
Figure 2.4). During this phase, a weak convection occurs at the stellar surface but not sostrong as to generate magnetic ﬁelds. The stellar inside soon turns to fully radiative andthen reaches to ZAMS. HAeBes are a few times as young as the age estimated by Iben (1965)
and stars larger than 8M already ignite the ordinary hydrogen burning in the core when
they appear on the birthline.

2.1.4

Other Types of Stars Related to Herbig Ae/Be Stars

We brieﬂy mention other stellar categories which are related to HAeBes. Their evolutional
status is illustrated in Figure 2.5.

Classical Be stars: Classical Be stars are deﬁned as the main-sequence or somewhat higher
luminosity B-type stars with the Balmer emission lines, originated from circumstellar
envelopes. This stellar category was ﬁrst discovered in γ Cas and later called as’classical’ for distinguishing them with the HAeBe group. The Balmer lines and thepolarization properties, called Be phenomena, are thought to relate with disks madeby rapid axial rotation (v sin i ∼250 km s−1 ). Be stars consist of 20% of the B-star
population and is largest in the spectral type range B3–4 (Slettebak 1991).

2.2. STELLAR X-RAY EMISSION

9β-Pic (Vega) type stars: The IRAS satellite found an excess emission from some mainsequence stars such as β Pictoris, Vega, and Fomalhaut (Walker and Wolstencroft1988). A high resolution imaging, masking a stellar photosphere, revealed that thisexcess emission is originated in cool debris disks, probably made from cometary clumps.
Though a part of them are later identiﬁed as isolated HAeBes, these stars now groupedin β-Pictoris type stars and thought to be in the phase between HAeBes and the mainsequence stars.
(Magnetic) Bp-Ap stars: Bp-Ap stars show chemically peculiar spectra and lie on the
main-sequence branch. Some of them have magnetic ﬁelds and are especially called
as magnetic Bp-Ap stars. From a HIPPARCOS survey of magnetic Ap stars, Hubrig
et al. (2000) concluded that magnetic ﬁelds appear only in stars that have completed
at least ∼30% of their main-sequence lifetime. Nevertheless, a few HAeBes also have
Bp-Ap characteristics (see Table 4.1) and the plausible magnetic Bp star S1 in ρ-Oph
has characteristics of its youth (Andre et al. 1991).
High Luminosity Young Stellar Objects: High luminosity young stellar objects have
Lbol ≥ 1000 L (HLYSOs), most of which are embedded infrared sources and thought
to be progenitors of HAeBes (cf Mundt and Ray 1994).
UX Ori variables: UX Ori stars (UXORs) are deﬁned as highly variable stars with deep
irregular minima. A UXOR phenomenon is thought to occur by an occultation ofpatchy circumstellar dust clouds. 1/3 of HAe stars may be UXORs but no HBe shows
UXOR phenomenon (see Natta et al. 2000).
A-shell type stars: A certain HAeBes coexist two types of line proﬁles in their spectra,
one originating in the stellar photosphere and the other in the cooler shell (see Hauck
and Jaschek 2000).

2.2

Stellar X-ray Emission

2.2.1

X-ray Emission from Main-Sequence Stars

Low mass Stars – Later than F5
X-ray emission from low mass MS stars like the sun exhibits occasional rapid x-ray ﬂares(see Figure 2.6), and therefore is originated in the magnetic activity. The magnetic ﬁeld isgenerated by the dynamo mechanism, powered by a coupled eﬀect on surface convection anddiﬀerential rotation (α − ω dynamo). The role of stellar diﬀerential rotation can be seen in
LX dependence on the projected rotational velocity, vrot sin i. LX increases with vrot sin iand reaches the saturation level at vrot sin i ∼ 15 km s−1 (log LX /Lbol < −3, see also Figure

CHAPTER 2. REVIEW

102.7). The level is also proportional to a square of a stellar radius (Fleming et al. 1989),
which implys that LX is proportional to the converting factor of magnetic loops on stellar
disks. Through MS lifetime, low mass stars decay the x-ray luminosity owing to a systematic
decline in the rotational speed by magnetic breaking (see Simon and Drake 1989). In the
Pleiades cluster, however, the dependence of the x-ray luminosity on stellar age is not clearlyseen (Stauﬀer et al. 1994).
Intermediate Mass Stars – Late B to F5 including β-Pic & Magnetic Bp-Ap Stars
Normal intermediate mass stars including β-Pic stars have the smallest x-ray luminosities.
− log LX (ergs s−1 ) ∼ 28−30, (Panzera et al. 1999). This is due to deﬁciency of highenergy activity. Even the above mentioned luminosity is suspected to come from a low masscompanion. Whereas He-strong (Bp) stars have the x-ray emission level of log LX /Lbol ∼ −7
and He-week (Bp) and Si (Ap) stars have the emission level up to log LX /Lbol ∼ −6, but
SrCrEu-type (Ap) stars emit no detectable x-rays (Drake et al. 1994). The He-strong star σ
Ori E and the plausible magnetic Bp stars Rho-Oph S1 have incredibly high x-ray luminosity
for B stars (LX ∼ 3×1030 ergs s−1 , Montmerle et al. 1983; Drake et al. 1987). Followup radio
observations of Rho-Oph S1 by Andre et al. (1988) and Andre et al. (1991) indicated thenon-thermal radio spectrum, which can be simulated by a pole-on bipolar magnetic ﬁeld of ∼2 kG on its stellar surface. A model proposed for magnetic Bp-Ap stars (Havnes and Goertz
1984) required that dense inner region with a chromospheric characteristic (n ∼ 1012 cm−3,
T ∼ 104 K) extended out to 10 R∗ . The radio and x-ray emission cannot penetrate sucha thick envelope and hence would be originated in the outer magnetosphere (Andre et al.
1988).
High Mass Stars – OB Stars & Be Stars
MS OB stars produce almost constant x-ray emission (kT < 1 keV, log LX (ergs s−1 ) ∼ 32
and log LX /Lbol ∼ −6 − −8, Seward et al. 1979; Harnden et al. 1979; Long and White

Figure 2.6: An x-ray ﬂare from the sun in the Yohkoh soft x-ray band

2.2. STELLAR X-RAY EMISSION

11

Figure 2.7: Log LX (left) and log LX /Lbol (right) vs. eﬀective temperature distribution in
the Orion molecular cloud (Gagne et al. 1995)
1980; Chlebowski et al. 1989). Coronal x-ray emission like the sun was ﬁrst proposed forthe origin, but the strong mass loss log Ṁ (M yr−1 ) ∼−5 – −7 with v ∼ 1500–3000 km s−1
derived from P-Cygni proﬁles (Lamers and Leitherer 1993; Berghoefer and Schmitt 1994)
required large NH , which is not observed. Nordsieck et al. (1981) searched for, but did notﬁnd, optical lines characterizing coronal emission from two OB stars.
Lucy and White (1980) proposed a wind-shock model for OB stellar x-rays. In thismodel, line absorption of stellar radiation accelerates circumstellar materials and produce
stellar winds. The winds conﬁne ambient gases by ram pressure and form shocks (see Figure
2.8 left), which thermalize the gases and ﬁnally causes soft x-ray emission. This theory
was strengthened by numerical calculations (Owocki and Rybicki 1984; Owocki and Rybicki
1985; Owocki and Rybicki 1986; Owocki et al. 1988) and the observational evidence ofsmall time variations in ζ-Orionis (Berghofer and Schmitt 1994). Nevertheless conﬂicting
results were reported as well, and therefore Usov and Melrose (1992) proposed an alternativeemission mechanism – magnetic reconnection on current sheets of large magnetosphere (seeFigure 2.8 right). The presence of magnetic ﬁelds on O stars was suggested from the nonthermal radio emission and the continuous variability in the blueshifted absorption of Si IV
doublet lines (ζ Pup, Howarth et al. 1995). The x-ray emission mechanism on OB stars
is under discussion, but it should be noted that latest high-resolution x-ray spectroscopy
with Chandra and XMM-Newton suggests that x-ray line emission from massive MS stars
originates from conﬁned high density plasma well above 1012 cm−3 , with a mean velocity ofv ∼700–1000 km/s (Schulz et al. 2000; Kahn et al. 2001).
Late B stars have diﬀerent x-ray characteristics from O and early B type stars. In
a survey of Grillo et al. (1992), 1545 B-type SAO stars with Einstein, early B-type stars
between B0−B3 emit x-rays at the level above LX >1030 ergs s−1 , while late B-type stars
tend to have fainter or no x-ray emission. Cassinelli et al. (1994) surveyed nearby MS B
stars with ROSAT and found that the LX /Lbol ratio sharply decreases later than B1 downCHAPTER 2. REVIEW

12

Figure 2.8: Shematic views of two x-ray emission models of massive MS stars. The left
panel displays successive propagating shocks in the stellar wind (Lucy 1982). The holizontal
and vertical axes show distance from a star and wind velocity, respectively. The rightpanel displays the conﬁguration of magnetic ﬁelds (Usov and Melrose 1992). The magnetic
reconnection fed by winds occurs on the equatorial plane.
to LX /Lbol ∼ 10−9 . On the other hand, Berghoefer et al. (1996) performed a complete
survey of OB stars with the RASS and follow-up HRI data. they demonstrated that the
canonical relation of LX /Lbol values ∼ −7 is valid for stars earlier than B1-1.5 (see also thediscussion 8.1), but such a clear relation is not seen later than B1-1.5. They also found thatthe detection rate decreases less than 10% below the spectral typoe.
The prototypical classical Be star, γ Cas, has normal x-ray luminosity of, but quite hightemperature agaist, MS OB stars (LX ∼1033 ergs s−1 , kT ∼10.6 keV). The x-ray emission
would therefore come from a white-dwarf companion (Kubo et al. 1998). It is not known as
whether the x-ray property is typical of classical Be stars.

2.2.2

X-ray Emission from Pre-Main-Sequence Stars

Low Mass Stars – T-Tauri Stars & Protostars
Low mass PMSs – TTSs and protostars – have more violent x-ray activity than low mass
MS stars (Feigelson and Decampli 1981; Montmerle et al. 1983 for T-Tauri stars; Koyama
et al. 1996; Kamata et al. 1997; Tsuboi et al. 2000 for protostars). Both types of stars
are highly variable and exhibit large ﬂares in the x-ray wavelength. From the similarity to
the low mass MS stars, the x-ray emission from TTSs is probably originated in the magnetic
activity enhanced by a coupled eﬀect of deep convection and rapid rotation. Actually, therelation between LX and vrot sin i, an indicator of the solar-type dynamo, is seen for the
T-Tauri membership of the Orion molecular cloud (Gagné et al. 1995), but it was not forthe membership of Chameleon I cloud (Feigelson et al. 1993). On the other hand, protostars

2.2. STELLAR X-RAY EMISSION

13

3
1

5

5
4

2

4
1

2

3

Figure 2.9: Schematic view of possible magnetic conﬁguration between a star and its disk
quoted from Montmerle et al. (2000). The solid lines describe a transition of magnetic ﬁelds
in a sequence of No. 1 – 5, driven by rotating diﬀerence between a star and its disk.
cannot be thought to operate solar-type dynamo because they are fully convective and do
not have an anchor point of the magnetic ﬁeld. One of the most prausible model - thestructure of the magnetic ﬁeld linking a star and its disk (Koyama et al. 1996; Tsuboi et al.
2000; Montmerle et al. 2000, see Figure 2.9) - might explain the x-ray activity. This model
assumes that magnetic ﬁelds perpendicular to the equatorial plane are inherently present.
probably through the process of gathering interstellar magnetic ﬁelds during the gas infall.
These magnetic ﬁelds are continuously twisted by the diﬀerential rotation between a starand its disk, and reconnect after one rotational diﬀerence, seen as an x-ray ﬂare. In the
numerical simulation by Hayashi et al. (1996), the reconnection accompanies a gas eruption,
which could be later seen as outﬂows or jets.

Intermediate Mass Stars – Herbig Ae/Be Stars
Damiani et al. (1994) detected x-ray emission from 11 HAeBes out of 31 observed sources
using Einstein (9/20 for Be stars, 1/8 for Ae stars and 1/3 for F0−F5 stars). They foundthat Be stars are more luminous than Ae stars. The luminosity function of Be stars is much
diﬀerent from, but that of Ae stars is similar to, low mass stars in Taurus-Aurigae. They
therefore suggested that x-ray emission from Be stars arise in the same mechanism of MS
OB stars while that of Ae stars might actually originate in low mass companions. The X-ray
luminosity depends on the bolometric luminosity, which, they concluded, only see a diﬀerent
expression of LX − spectral type relation because observed HAeBes are close to the MS track.
They suggested a correlation between the x-ray luminosity and the terminal wind velocity
but no correlation between the x-ray luminosity and the vrot sin i. On the other hand,CHAPTER 2. REVIEW

14

Table 2.1: X-ray properties of giant molecular cloud cores
kT
log LX † Band
Region(keV)(ergs s−1 ) (keV)
Monoceros R2
1
5.6 (4.7–7.0)
32.4
0.8–10 All GMC
NGC 63342
9.0 (7.2–12.3)
33.8
0.5–10 Five FIR cores
NGC 22643
4.4 (3.3–7.4)
32.2
0.7–10 IRAS 9 (Allen’s source) or IRS1
W3
4
5.6 (4.6–7.3)
32.2
2–10 core
Orion Trapeisum5
3.1 (2.9–3.2)——
Trapesium region† Absorption corrected luminosity for the energy band column. References: 1 this work (the subsection 6.1, see also Hamaguchi et al. 2000), 2 Sekimoto et al. (2000), 3 Nakano et al. (2000), values of
hard component in the 2C model, 4 Hofner and Churchwell (1997), 5 Yamauchi et al. (1996).
Targetref.

Zinnecker and Preibisch (1994) detected x-ray emission from 11/21 HAeBes using ROSAT.
They found a rising trend of LX with Lbol . In their estimation, there was no correlation withthe spectral type or vrot sin i, but a weak correlation with the wind velocity and a rather
good correlation with the mass loss rate. They suggested that x-ray emission of HAeBes
would originate in shocks of strong stellar winds to remnant circumstellar material.
Giant Molecular Cloud Cores – Massive Protostar Emission?
Massive stars go through their PMS age in the deep inside of GMC cores and reach to the MS
stage in a very short time (< 104 yrs). Owing to the strong absorption by the surroundinggases, the beginning phase of massive stars is not seen in optical or even in near infrared. A
theory predicts that massive stars larger than 10 M cannot form from mass contraction of
circumstellar gases but from the merging of smaller mass stars because strong radiation ﬁelds
disperse accreting gases. The existance of massive protostars is therefore still doubted, but
strong outﬂows and maser emission (OH, SiO) are generally observed around GMC cores.
These phenomena suggest violent activities on very young massive stars.
Several massive young stars in the PMS phase, whose extinction exceeds over AV ∼ 50
(NH > 1023 cm−2 ), were found in nearby GMC cores. Soft x-rays from massive PMSs
undergo strong absorption so that hard x-ray imaging is the only tool to probe the x-ray
activity. ASCA is the ﬁrst satellite holding such capability. Though the spatial resolution isstill limited to resolve crowded membership in GMC cores, ASCA have detected hard x-ray
sources from the GMC core (Table 2.1). The next generation x-ray satellite Chandra has
an excellent spatial resolution (∼0.5 ) in the ASCA band. Garmire et al. (2000) resolved
all x-ray sources in the Orion cloud with Chandra, which includes the on-going massive star
forming cite Becklin-Neugebauer/Kleinman-Low (BN/KL) region, but they did not detect
x-ray emission from neither the IR luminous massive YSOs BN object nor the source i,showing strong outﬂow activity.

Chapter 3
Instruments
The atmosphere of the earth completely blocks cosmic x-ray radiation. Study of celestial
x-ray sources thus requires an observation in the outer space. Fortunately, recent advanceof space technology enabled continuous launch of spacecrafts for x-ray astronomy, and ﬁvesatellites, ASCA, XTE, Beppo SAX, Chandra and XMM-Newton are now working (at Jan
5, 2001). In this chapter, we summarize the performances of two x-ray observatories ASCA
and ROSAT, which led x-ray astrophysics in 1990s. We only mention subjects related to ouranalyses.

3.1

The ASCA Satellite

3.1.1

OutlineAdvanced Satellite for Cosmology and Astrophysics - ASCA - is the 4th Japanese x-ray
astronomy satellite developed by the Institute of Space and Astronautical Science (ISAS)
in collaboration with domestic x-ray astronomy groups and American institutes MIT and
NASA. ASCA was launched on February 20 1993 at the Kagoshima Space Center (KSC)with the M-3S-II rocket and put into an approximately circular orbit (perigee 520 km, apogee
620 km, inclination 30.1◦ , orbital period ∼96 min). ASCA worked properly until July 1999,
but since then it has malfunctioned for an unexpected attack on air disturbance caused byan increase of solar activity.
ASCA carries four identical x-ray telescopes (XRT) and coupled focal plane detectors,
two of which are the Solid-state Imaging Spectrometers (SISs) and the other two are theGas Imaging Spectrometers (GISs). All these detectors run independently and output dataat the same time. The schematic view of the ASCA satellite is shown in Figure 3.1. An
introduction to ASCA is described in Tanaka et al. (1994).
15

CHAPTER 3. INSTRUMENTS

16

AAAAAAAAAAAAAAAAAAAAAAAAAAA

AA
AA
AA
AAAA
AAAA
AAAA

AAAA
AAAA
AAAA
AAAA
AAAA

S-band antennauzuki

X-band antenna
Shunt regulator

Extendable
optical bench

AAA
AAA

AAAAA
AAAAA
AAAAA

Sun shade

NSAS
SSAS

AAAA
AAAA
AAAAAAAAAAAAAAAAAAAAAAAAAAA

AAAA
AAAA
Solar paddle

Y

X-ray telescope

X

Z
Satellite coordinates

Figure 3.1: Schematic view of the ASCA satellite

The satellite attitude is controlled in three-axis with the stability less than 10 . Two star
trackers calibrate the absolute pointings through observations, whose pointing accuracy was
designed better than 12 but later turned out to be around 40 in onboard calibrations causedby thermal distortion of star trackers. From the constraint on power supply, orientation of
the spacecraft is limited to the normal direction of solar paddle within 30◦ from the sun.
This constraint limits the observable sky at one time and the entire sky is accessible every
half year. Harsh environs on the ASCA orbit constrains the satellite operations. At ﬁrst,
whether satellite is in the sunshine or not, ie satellite day or night, decides the temperature of
the spacecraft. Detector response, especially on SIS, is aﬀected by the satellite temperature.
Second, the radiation of cosmic particles not only degrades the quality of the data but also
can damage detectors. South Atlantic Anomaly (SAA) is a region where the particle density
is always high. During the passage of SAA, all detectors onboard ASCA are shut down.
The science data are ﬁrst stored in onboard memories (Bubble data recorder; BDR) andthen downloaded to the ground station at KSC and NASA deep space network (DSN). The
storage capacity of BDRs is limited and thus the memory space of the science data written
in a time is pre-programmed by each observing schedule. There are three memory modes;
bit HIGH mode secures 32K bit memory space per second, bit MEDIUM does 4K bit andbit LOW does 1K bit. Overﬂowed data are discarded and thus an observer has to prepare aproper program met for observing targets and operational constraints.

3.1. THE ASCA SATELLITE

17

Figure 3.2: Point spread function of the ASCA XRT for four diﬀerent oﬀ-axis angles. Innerboxes and outer circles represent the ﬁeld of view of SIS and GIS respectively.

3.1.2

X-ray Telescope (XRT)

X-ray imaging telescopes utilize the optics of Wolter type I geometry, which focus x-rays
through two mirrors of paraboloids and hyperboloids. ASCA XRTs are based on this optics
as well, but they approximate these shapes to two ﬂat conical surfaces made from thin aluminum foils. This simpliﬁcation makes the spatial resolution modest (Half Power Diameter:
HPD ∼3 ) but enables us to nest multiple mirrors (119 foils an XRT), which provides large
collective area especially in the hard x-ray band. ASCA is thus the ﬁrst satellite which holds
the imaging capability of astronomical objects up to ∼10 keV.
XRT has a cusped shape point spread function (psf) and concentrate 20% of parallel
x-ray beams within a circle of 1 diameter. Ideal psf is radially symmetrical, but spokes
connecting nested mirrors block part of x-ray beams and thus real psf looks like a butterﬂyshape (see Figure 3.2). As the characteristics of the Wolter I optics, psf is elongated to thedirection perpendicular to the radial vector of XRT fov, which makes the spatial resolution inouter fov poor. XRT eﬀective area decreases on the far side from the optical axis, dependingon x-ray photon energy. This characteristic is called vignetting eﬀect. See Serlemitsos et al.
(1995) about the details of ASCA XRT.

CHAPTER 3. INSTRUMENTS

18

Table 3.1: SIS and GIS detector performances
Unit
(keV)

SIS
Energy range
0.4–10Field of View
20 ×20
Energy resolution (5.9 keV) (eV)130
Maximum Time Resolution (sec)4†1†11CCD mode, †2 diameter, †3 PH mode

3.1.3

GIS
0.7–1050 †2460
60µ†2Solid-state Imaging Spectrometer (SIS)The Solid-state Imaging Spectrometers (SISs) are the ﬁrst x-ray CCD cameras aboard astronomical x-ray satellites. CCD, the abbreviation of charge coupled device, is so called aµm pixel array of solid-state detectors. The detection process is as follows. An x-ray photon
interacted in the depletion layer of a CCD generates large amount of electrons proportional
to its photon energy. This electron cloud is collected at an electrode of an incident pixelby electric ﬁelds and conveyed to a readout gate according as a CCD control clock. The
capacitor on the gate converts an electron charge to a signal voltage, which is ampliﬁed ina circuit. A readout system measures it with addressing the pixel position. A certain percentage of electron clouds, with many electrons or near pixel gaps, extend over plural pixels.
Because cosmic particle events also generate huge electrons through multiple pixels, eventsare classiﬁed into 7 groups according to the pattern of 3×3 pixel pulse heights and 4 groups(grade 0,2,3,4) are adopted as x-ray events. The thickness of the depletion layer is ∼30 µm
so that the SIS quantum eﬃciency drops above 6 keV. The soft x-ray eﬃciency (kT <1 keV)is better because insensitive layers above the depletion layer is thin. SIS sensitive band istherefore wider than GIS (0.4–10 keV: see Table 3.1).
Each of two SISs comprises four identical CCD chips (see Figure 3.3), which are cooled
down on −60 ◦ C by their Peltier coolers. One CCD chip covers a 11.1 ×11.1 fov and needs
4 second a readout cycle. The readout module of each SIS sensor processes the output dataof one CCD chip at one time so that the other CCD chips waits to be read and exposed on
the focal plane. Exposure time per CCD frame is thus multiplied according as the numberof working CCD chips. Because long exposure not only worsen the timing resolution but
also causes increases of pile-up events and dark current noise, observers can select number of
CCD chips to work, from 1, 2 and 4 CCD modes. Exposures of the selected modes are 4, 8
and 16 seconds, respectively. On the other hand, an observer can choose a data format fromtwo modes according as the BDR capacity. Both data formats consist of the information
about CCD ID, event pixel position and arrival time. The diﬀerence of these two modes is
described like this.

3.1. THE ASCA SATELLITE

19~22.5 arcmin

Gap: ~0.5mm~18.5 pixels~29.4 arcsec

S0C0

S0C3

S1C2

S1C1

x

X

y
x
Y

1
y

420 pixels~11.3 mm
~11.1 arcmin

Satellite coordinates

S0C1
S1C31P 1

S0C2
S1C0

4 0
2
2P

~22.8 arcmin

y
x

DETX

1 pixel= ( 27 µ m )2
= ( 1.59 arcsec )2x
DETY
Detector coordinatesy422 lines~11.4 mm
~11.2 arcmin

Gap: ~0.13 mm~4.8 pixels~7.6 arcsec

1 1 CCD mode default position
2 2 CCD mode default position
4 4 CCD mode default position

Sun

1

P

2 P Old 1,2 CCD mode default position
0

SIS0 Optical Axis

1

SIS1 Optical Axis

Figure 3.3: SIS chip conﬁguration
FAINT: All the pulse heights of 3×3 pixels are transmitted to the ground. This mode canbe only used for faint sources.
BRIGHT: An added pulse height of 3×3 pixels and its grade information are transmitted
to the ground. Though event number memorable in BDRs is about four times as
many as that of the FAINT mode, observers cannot re-process downloaded data on
the ground.
Several problems have been found through onground and onboard calibrations. Properties of each problems and how to deal with them are summarized below.
Echo: A certain fraction of a pulse height is added on its right-side pixel (Otani and Dotani
1994). This false pixel height, linear to the real pixel height, is made in the processingof the Analogue Electronics (Yamashita 1995). This eﬀect can be easily corrected if theevent data is in the FAINT mode, but it cannot be corrected in the BRIGHT mode.
Its response matrices correct the deviation from the real signal output.
Dark Frame Error (DFE): is to misestimate the mean pulse heights of CCD pixels without events (zero level of Dark frame; Otani and Dotani 1994). Onboard software
calculate the zero level by averaging pixel pulse heights between -40 ADU and 40
ADU, but PH is asymmetrically distributed around zero level by the integration of

CHAPTER 3. INSTRUMENTS

20

X-ray photons, charged particles and optical light leakages, which causes zero level tooverestimate. In an orbital night, the DFE is stable (∼3 ADU for 4CCD mode). Ina satellite day, the mean DFE value is almost as same as in the night, but the ﬂuctuation is typically doubled owing to the optical light leakage, and worsen when thebright Earth edge is within 25◦ from the SIS fov (∼30 ADU at maximum). DFE isalso severe during night-day terminators and the passage of SAA. DFE in the FAINT
mode can be corrected at the ground data processing by averaging all PHs of corner
pixels of an event in a certain time scale (Otani and Dotani 1994). There is no way tocorrect DFE for the BRIGHT mode and the eﬀect is included in its response matrices.
Residual Dark Distribution (RDD): The zero level still has large scatters after the DEF
correction, called “Residual Dark Distribution (RDD)” (Dotani et al. 1997). In acertain probability, charged particles on orbit damage some Si lattice. A CCD pixel
holding a damaged lattice always output large charge currents in proportion to numberof damaged lattice. The on-orbit frontend processor only stores the averaged zero-levelof each 16×16 pixels. The residuals against the average zero-level remain. The RDD
eﬀect has been gradually worsen after launch by accumulating radiation damages.
Charge currents stored in a pixel in an exposure is in proportion to the exposure time.
The eﬀect is still negligible in the 1-CCD mode, but it is signiﬁcant in the 4-CCD
mode data after 1995 (Yamashita 2000). The RDD correction method, which takes
zero-level maps of each pixel in a certain interval, have been developed.

3.1.4

Gas Imaging Spectrometer (GIS)

Gas Imaging Spectrometers (GISs) are imaging gas scintillation proportional counters. The
process of detecting x-ray photon is illustrated in Figure 3.4. An incident x-ray photon
through the entrance window generates a photo electron cloud in the drift region. This
cloud drifts to the scintillation region by weak electric ﬁelds, and then accelerates by thelarge voltage gap between the intermediate and ground meshes. In this process, electrons
constituting the cloud excite Xe atoms, which emit recombination UV lines (λ ∼ 170 nm).
The imaging photo-multiplier tube installed under the gas cell collects these photons andconverts anode signals to position data and signals of the last dynode to pulse-height andrise-time data. The drift region is thick enough and thus the quantum eﬃciency of hard xrays is high (∼100 %), while the entrance window absorbs soft x-rays. A method of rejecting
cosmic particle background, using the information of the rise time duration of event signals,
was established though onground calibrations. This method rejects ∼90% of background
events in 2–10 keV band. Among four system modes prepared for GIS operation, observers
use the Pulse Height (PH) mode, which assigns the information of the pulse height, event
position, rise time, event spread and timing data within 31 bit data length. The nominal bit
assign is 10, 8, 8 and 5 bits for pulse height, x axis, y axis and rise time data, respectively.
Details of the GIS detector are described in Ohashi et al. (1996) and Makishima et al.

3.2. THE ROSAT SATELLITE

21
HV

X-ray

GND
52mm

support grids- 6kV

Be entrance window

photo electrons
cloud(10um)

Drift Region

Xe + He

Gas cell

96%

4%

10mm
- 5kV

intermediate mesh

Fluorescence
UV light

Scintillation
Region15mm

ground mesh

2.5mm
3.0mm

Quartz

GND

Imaging PhotoMultiplier Tube

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

1400V

PH
RT

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

Pre-Amplifier

XP / YP

Figure 3.4: Schematic view of the GIS sensor system.
(1996).
GIS background levels occasionally ﬂare up out of the passage of SAA. These phenomena
are called “hard-ﬂares” and “soft-ﬂares”. They were carefully studied by Ishisaki et al.
(1997). Hard-ﬂares occur on the periphery of SAA, probably caused by particles trapped in
SAA. These events can be monitored with the H02 parameter including GIS house keeping
data. While soft-ﬂares do not accompany an increase of the H02 count. These events
frequently occur in the south west to SAA and above Hawaii, probably caused by electrons
entered through XRTs.

3.2

The ROSAT Satellite

3.2.1

OutlineROSAT (acronym of Röntgensatellit) is the German x-ray satellite mainly developed by
Max-Planck-Institut für extraterrestrische Physik (MPE) and launched into low Earth orbit
on June 1990. ROSAT has one x-ray telescope (XMA) consisting of four Wolter I mirrors andthree detectors, two identical position sensitive proportional counters (PSPC−B, PSPC−C)
and one high-resolution imager (HRI). These detectors are switched on and oﬀ the focal
plane of XMA. The sensitive energy bands of both detectors are between 0.1−2.4 keV, but

CHAPTER 3. INSTRUMENTS

22

WFC
XRT

Star sensors

Solar Panels
Figure 3.5: Schematic view of the ROSAT satellite.
the quantum eﬃciency between 0.3–0.4 keV is negligible owing to the absorption of windowmaterials. The eﬀective area of PSPC is about three times as large as that of HRI.

3.2.2

X-ray Mirror Assembly (XMA)

XMA is constructed from a zerodur block - a ceramic material - polished up to the ideal
Wolter type I geometry. The highly smoothed mirrors focus soft x-rays below 2.4 keV within2 radius circle on-axis. The psf however becomes less sharp with increasing oﬀ-axis angles.
(See the dependence of HPD on oﬀ-axis angles in page 92 of Briel et al. 1994).

3.2.3

Position Sensitive Proportional Counter (PSPC)

PSPCs are multiwire proportional counters. These detectors provide modest energy resolution (27% FWHM at 1.5 keV) and relative time resolution down to ∼130 µsec. The intrinsic
spatial resolution is ∼25 in the focal plane at 0.93 keV. At oﬀ-axis angles larger than about
20 , the position resolution is dominated by XMA.

3.2.4

High Resolution Imager (HRI)

HRI comprises two cascaded microchannel plates (MCPs). A channel size of HRI is enoughsmall to map XMA resolved x-ray sources. It has negligible energy resolution but provides
relative time resolution down to 60 µsec. HRI has UV bandpass below ∼4000Å, caused byenhanced transmission of the UV/ion shield (Zombeck et al. 1997; Berghöfer et al. 1999).

Chapter 4
Observation & Data Reduction
4.1

Target Selection

4.1.1

Data Retrieval

The et al. (1994) compiled a comprehensive catalogue of members and candidate membersof HAeBes in previous articles. This catalogue comprises ﬁve categories ; i) Be and Ae
stars historically recognized as true members or potential candidates of HAeBes, ii) Fe andemission line stars, iii) extreme emission line objects, iv) early type peculiar emission line
stars listed by Allen and Swings (1976) and Dong and Hu (1991), v) non-emission line objects
including β Pic and Vega type stars and vi) stars rejected as HAeBe candidates. We adoptedi), ii) and v) as our HAeBe samples. The reason for adopting the categories ii) and v) isthat these sources will ﬁll the gap between HAeBes and TTSs or MSs. Whereas we did not
adopt the categories iii) and iv). These members seem to be mainly on the Galactic planeand might be other types of stars such as classical Be stars or white dwarfs. In addition to i),ii) and v), we adopted few more samples from a few publications, which is lately recognized
as HAeBes (Preibisch 1999; van den Ancker et al. 1997; van den Ancker et al. 1998).
In the ASCA guest observing programs (Announce of Opportunity: AO 1–7), we haveobserved several HAeBes, MWC 297, MWC 1080, V892 Tau and HAeBes in the R CrA
cloud. For the other HAeBes observed with ASCA, we picked up the data from the ASCA
archive - DARTS database provided by ISAS and W3 database by HEASARC. The observed
HAeBes are shown in Table 4.1 and their observation logs are in Table 4.2, 4.3. On the other
hand, we added the Monoceros R2 GMC (see also the section 8.2). Its log is listed in thesame table of HAeBes.
23

24

4.1.2

CHAPTER 4. OBSERVATION & DATA REDUCTION

Quantities of Fundamental Stellar Parameters

The fundamental parameters of HAeBes could have some uncertainties. Eﬀective temperature of a HAeBe is measured from the blackbody component of the stellar core, but it
is usually merged with disk and envelope components or contaminating lights from associating reﬂection nebula. The temperature therefore, especially on embedded sources, canhave large uncertainties. Whereas the bolometric luminosity has another uncertain factor,
distance. Almost all HAeBes are more than a hundred parsec away from the sun. Groundbased telescopes do not have an accuracy to measure the distance from trigonometric parallax. Distances to most HAeBes are thus estimated from those of associating MS stars or
clouds, but it is almost no way to prove the relationship between HAeBes and these objects.
Some HAeBes have large uncertainties for distance in a factor to order (eg V921 Sco, see thesubsection 5.1.8). The astrometry satellite HIPPARCOS launched in 1989 had an ability tomeasure the absolute sky coordinate of all stars with an accuracy of 2 milli-arcsec (Perryman
et al. 1992; Walter and Sovers 2000). From the measurement of trigonometric parallax ofthe orbital diameter of the earth, HIPPARCOS can measure stellar distance up to ∼ 400 pc.
We thus ﬁrst rely on the results using the HIPPARCOS data (van den Ancker et al. 1997;
van den Ancker et al. 1998; Bertout et al. 1999) and, if there is no information, refer to thelatest articles.
Visible lights and X-rays are attenuated by the Interstellar medium (ISM) through different processes. Visible lights are absorbed or scattered by interstellar dusts. The extinction
is basically continuous and prominent in short wave-lengths although some molecules makeline absorptions. Degree of extinction is described with the attenuation ratio in V band(AV ). Whereas x-ray lights are absorbed by photo-ionization with heavy elements such ascarbon and oxygen (Morrison and McCammon 1983). The extinction is also continuousand eﬀective in soft X-rays, ie long wave-lengths, and it also has sharp absorption edges at
ionization energies of each absorbing element. Degree of extinction is described with the column density of hydrogen (NH ), usually assuming solar values for chemical abundances. The
relation between AV and NH depends on the ratio and size distribution of dust grains. A few
authors derived empirical relations of these parameters, using optical and x-ray observations.
We adopt the formula provided by Ryter (1996); NH = 2.2 × 1021 AV cm−2.
In this article, we also show the physical parameters of HAeBes in Table 4.1. All
parameters are referred to articles chosen by the above mentioned procedure. Teﬀ of some
HAeBes is estimated from the spectral type (Tokunaga 2000).

4.1. TARGET SELECTION

Table 4.1: Stellar parameters of the observed Herbig Ae/Be stars
R.A. (2000)
Dec. (2000)dRef. G
Sp. T
log Teﬀ
log Lbol
AV
Bin.
◦


h m
s
(pc)
(K)
(L )
BD+30 549
3 29 19.8
31 24 57
350
a
B8
4.08
1.1 †1
1.9
V892 Tau
4 18 40.6
28 19 17
160
i
II
A6
–
–
4.1
vis.
CoKu Tau 1
4 18 51.5
28 20 28
160
ii
M2
(3.55)–
2.1
vis.
AB Aur
4 55 45.8
30 33 4
144
i
I
A0s
3.98
1.7
0.5
IX Ori
5 34 40.8−5 22 43460 †2 ii
K––
0.5
V372 Ori
5 34 47.0−5 34 15460 †2 v
B9.5+A0.5
3.93
2.2
0.5
YZ Ori
5 34 54.0−5 3 30460 †2 ii
K5
3.64
0.0
0.2
HD 36939
5 34 55.3−5 30 22460 †2 v
B8-9
4.05
1.9
0.5
HD245185
5 35 9.6
10 1 52
400
i
A2
3.96
1.3
0.1
LP Ori
5 35 9.8−5 27 53
460 †2 v
B2
4.29
3.1
0.7
MR Ori
5 35 17.0−5 21 46460 †2 v
A2
3.93
1.9
1.6
V361 Ori
5 35 31.4−5 25 16460 †2 v
B1.5/B4
4.14
2.6
0.4
T Ori
5 35 50.0−5 28 42
460
i
I
A3
3.93
1.6
1.1
V380 Ori
5 36 25.4−6 42 58
460
i
I
B9-A0
3.97
2 †3
1.4
IR
BF Ori5 37 13.3−6 35 1
430
i
I
A5-6
3.90
0.1
0.3
MWC1205 41 2.3−2 43 1
500
i
A2p
3.95
1.3
0.03
VY Mon
6 31 7.0
10 26 5
800
i
O9−B8
4.05
2.9
7.4
VY Mon G2
6 31 8.2
10 26 1800v
A0s
(3.99)
4.0
1.6
V590 Mon
6 40 41.3
9 48 1
800
i
B7
4.09
2.6
0.6
LkHα 218
7 2 42.3−11 26 10 1150
i
I
B9
4.03
2.1
1.5
Z CMa
7 3 43.2−11 33 6
1150
i
F6
3.80
3.4
2.8
IR
LkHα 220
7 4 5.4−11 26 0
1150
i
B5
4.19
2.4
–
HD76534
8 55 8.7−43 28 0
830
i
B2n
4.34
3.9
1.2
HD97048
11 8 3.3−77 39 17
180iI
B9-A0ps
4.00
1.6
1.2
HD97300
11 9 50.0−76 36 48
188v
B9(4.03)
1.5
1.3
T Cha
11 57 13.5
−79 21 32
66
ii
G2
3.77
0.1
1.6
HD104237
12 0 5.1−78 11 35
116
i
A4s
3.93
1.6
0.3
IRAS12496
12 53 16.1
−77 7 2200i
A-F
(3.91)
1.7
12
HR599916 8 34.3−39 6 18
210
i
A5-7s
3.90
1.9
0.5
HR600016 8 34.6−39 5 34
241
a
B6
4.15
2.5
0.3
spec.
HD147889
16 25 24.3
−24 27 57
140
a
B2
4.34
3.3
3.3
Hen 3−1191
16 27 14.2
−48 39 28
–
i
B0
–
–
–
V921 Sco
16 59 6.9
−42 42 8
500
i
I
B4-5
(4.14)
3.0
7.1
MWC29718 27 39.6
−3 49 52
250
i
I
B1.5
(4.38)
4.5
8EC95
18 29 57.9
1 12 47310p
K0-K4(3.65)
1.8
36
S CrA
19 1 8.5−36 57 20
130ii
K6(3.62)−0.1
2.8
HD176386
19 1 38.9−36 53 27
122
v
B9
4.03
1.7
0.6
TY CrA
19 1 40.8−36 52 34
130
i
II
B7-9
(4.07)
1.8
1.0
R CrA
19 1 53.7−36 57 8
130
i
II
B8
(4.06)
2.1
1.9
spec.
T CrA
19 1 59.0−36 58 0
130
ii
II
F0
(3.86)
0.9
1.7
HD200775
21 1 36.9
68 9 48
429
i
I
B2.5
4.31
3.9
1.9
MWC1080
23 17 26.1
60 50 43 2200iI
A0-3
(3.96)
3.9
4.4
vis.
Ref.: References [i/ii/v – listed in Table 1/2/5 in The et al. (1994), a – van den Ancker et al. (1997) or van den
Ancker et al. (1998), p – Preibisch (1998)]. G: HAeBe Group deﬁned by Hillenbrand et al. (1992), Sp.T: Spectral
type [p: peculiar, s: shell-type], Bin.: binary [spec.: spectroscopic, vis.: visual, IR: infrared], Values in parenthesis is
typical values of the spectral type and group. †1 calculated from van den Ancker et al. (1998) assuming d ∼350 pc.
†2 distance to the Ori OB1 association †3 intermettent value of Hillenbrand et al. (1992).
Object25

CHAPTER 4. OBSERVATION & DATA REDUCTION

26

Table 4.2: ASCA observation log
Obs. ID

Abbr.

Seq. ID

Obs. TimeMod. Value
Start TimeOn Time
SIS
GIS
Date
Time
R.A.
Dec.
R.A.
Dec.
(◦ )
(◦ )
(◦ )
y m d
h m s
h m s
(◦ )
SVS13
SVS13
23021000
95 8 30 13 28 11 58 37 20–0.473 0.107–0.186 –0.001
IC359
IC359
23009000
95 9 20 8 11 29 3 28
–0.055 0.163
0.239 0.026
SU AurSU Aur23042000
95 2 25
4 43 23 31 48 16
0.750 –0.097
0.703 –0.080
ORION TRAPEZIUM
Ori Trap1
20004000
93 8 30 23 9 39 14 16 00.561 0.097
0.809 –0.128
Ori Trap2
25032000
97 9 23 17 0 11 27 10 56–0.104 0.184
0.111 –0.006
TRAPEZIUM I
Ori Trap3
25032010
97 9 30 13 3 21 30 37 52–0.333 0.217–0.128 0.041
TRAPEZIUM II
TRAP 2 S1Ori Trap4
25032020
98 9 20
5 41 12 31 34 40–0.008 0.167
0.240 0.006
Iota Ori1
25035000
97 9 21 23 37 39 41 23 12–0.049 0.169
0.189 0.000
IOTA ORIONIS 1
Iota Ori2
25036000
97 10 6
21 4 23 20 16 32–0.702 0.292–0.520 0.134
IOTA ORIONIS 2
L1641 NORTH
L1641N
21005000
94 3 13 23 6 38 27 4 32
–0.017 –0.153–0.063 –0.133
LAMBDA ORILamb Ori
21025000
94 3 123 56 3 14 14 56
0.220 –0.157
0.177 –0.135
Ori OB1B-3
Ori OB1
23019000
95 10 418 1 6 22 10 40–0.789 0.279–0.605 0.132
PSRJ0631+10
PSRJ0631
26007000
98 10 16 23 40 32 43 40 32–0.778 0.332–0.602 0.172
15 Mon15 Mon25015000
97 10 18 13 36 20 26 54 22–0.721 0.363–0.553 0.192
Z CMa
Z CMa
23014000
95 3 31 20 34 29 22 6 24
0.309 –0.295
0.278 –0.259
VELA SHRAPNEL BS
Vela Shrap
55030000
97 5 916 32 11 10 59 12
0.424 –0.259
0.378 –0.226
ChamI I
21009000
94 5 14
5 15 26 10 54 56
4.499 0.273
4.325 0.279
CHAMELEON I 1
ChamI II
21009010
94 5 14 16 8 14 16 2 30
4.432 0.248
4.251 0.254
CHAMELEON I 2
CHAMELEON I 3
ChamI III
21009020
94 5 15
8 11 26 12 58 40
4.661 0.230
4.475 0.237
ChamI IV
21009030
94 5 15 21 10 5 12 25 44.609 0.208
4.417 0.215
CHAMELEON I 4CHAIBROWNDWARF
ChamI V
27017000
99 8 7
4 15 28 71 55 44–0.045 –0.436–0.017 –0.390
ChamIII
27002000
99 3 820 0 10 44 9 36
0.802 0.313
0.081 0.074
CHAMELEON III
HD104237
HD104237
23003000
95 4 310 44 43 18 45 52
0.860 0.689
0.750 0.651
CHAMELEONIIChamII
24000000
96 3 10 16 4 7 19 11 28
0.297 0.304
0.025 0.035
LUPUS3
Lupus324008000
96 2 22 21 33 47 27 6 40
0.369 0.086
0.620 –0.160
RHO-OPHIUCHI
Rho-Oph
20015010
93 8 20
2 19 34 20 26 40
0.481 –0.221
0.440 –0.194
GALACTIC RIDGE 10-7
Gal R 155003060
97 9 320 19 44
8 51 440.200 –0.146
0.131 –0.128
Gal R 254004020
96 8 31 23 21 42
7 39 120.622 –0.148
0.566 –0.127
GAL RIDGE1 3
Gal R 354004030
96 9 17 0 508 30 240.658 –0.147
0.602 –0.126
GAL RIDGE1 4
MWC297 N1
MWC297 1 21007000
94 4 80 2 0 13 51 12–0.295 0.174–0.068 0.021
MWC297 2 21007010
94 4 123 31 36 55 12–0.585 0.180–0.374 0.043
MWC297 N2
MWC297 3 21007020
94 4 12 23 53 26 12 37 20–0.568 0.180–0.355 0.042
MWC297 N3
SERPENS
Serp
25021000
97 4 13 23 47 14 68 23 23–0.698 † 0.165 † –0.489 † 0.039 †R CrA1
R CrA 1
21001000
94 4 44 55 35 25 5 37
0.117 0.158
0.404 –0.037
R CrA2
R CrA 221002000
94 4 823 56 6 27 4 17
–0.139 0.184
0.129 –0.004
R CrA3
R CrA 324016000
96 4 5
1 39 40 44 30 90.192 0.150
0.481 –0.049
R CrA4
R CrA 424016010
96 10 18
3 25 59
8 33 20–0.034 –0.155–0.092 –0.136
R CrA5
R CrA 566017000
98 4 19
2 21 23 15 59 28–1.195 0.057–0.936 –0.034
R CrA6
R CrA 626018000
98 10 19 13 36 48 30 4 16
0.039 –0.150–0.022 –0.137
HD200775
HD200775
23029000
95 11 17
0 48 22 30 33 4
2.126 0.001
2.015 0.014
MWC1080
MWC1080
21006000
93 12 86 3 34 23 52 32
2.566 0.211
2.488 0.222
Monoceros R2
Mon R225019000
97 3 96 27 31 57 23 44
0.627 –0.155
0.587 –0.133
Obs. ID: Observation ID, Abbr.: Abbreviation of Obs. ID, Seq. ID: Sequence ID, On time: On source time, Mod value:
Modiﬁed value of an image coordinates presented by Gotthelf et al. (2000), † : the value in fa970413 2347.2359.

4.1. TARGET SELECTION

27

Table 4.3: Observing parameters of both SIS & GIS
Obs. ID

SIS
GIS
Memo
Format
LD
RC TS
Exp.
Mode
Exp.
Obs.
Ana.
(keV)
(ksec)(ksec)
SVS13
4
F/B
n.u.
PH
93.0
IC359
4
F/B
B2
0.55
y
n
27.4
PH
37.8
SIS: 0.55–10 keV
SU Aur
2
F/B
B2
n
n
40.8
PH
45.0
Ori Trap1
4
F/B
B
n
n
19.3
PH
18.0
PH
8.2
Ori Trap2
1
F/B
B
45.3 †1
Ori Trap3
1
F/B
B
n
n
7.9
PH
49.8
Ori Trap4
1
F/B
B
n
n
40.4
PH
35.9
Iota Ori1
1
F/B
n.u.
PH
72.5
Iota Ori2
1
F
n.u.
PH
24.6
L1641N
4
F/B
n.u.
PH
34.6
Lamb Ori
2/1
F
B2
n
n
21.8
PH
22.4
Ori OB1
4
F/B
n.u.
PH
28.6
PSRJ0631
1
F
n.u.
PH
81.6
GIS: 8-8-8-0-0-7
15 Mon
1
F
B2
0.48
n
n
36.6
PH
36.1
Z CMa
1
F
B2
n
n
30.5
PH
37.6
Vela Shrap
2
F
n.u.
PH
12.3
ChamI I
4
F/B
B
n
y
8.5
PH
15.2
ChamI II
4
F/B
n.u.
PH
13.9
ChamI III
4
F/B
n.u.
PH
16.8
ChamI IV
4
F/B
B2
n
y
12.6
PH
16.4
ChamI V
2
F
B2
y
n
82.4
PH
75.0
ChamIII
1
F
B2
n
n
62.5
PH
66.2
HD104237
1
F
B2
n
n
29.9
PH
27.1
ChamII
2
F/B
n.u.
PH
24.9
18.2
PH
41.3
SIS: 0.8–10 keV
Lupus3
4F/B
B2
0.55→0.7yy†2
Rho-Oph
4
F/B
n.u.
PH
41.1
Gal R 1
4
F/B
n.u.
PH
9.7
Gal R 2
4
F/B
B
y
n
12.2
PH
11.8
Gal R 3
4
F/B
n.u.
PH
8.6
MWC297 1
1
F
B2
n
n
10.5
PH
10.3
MWC297 2
1
F
B2
n
n
5.3
PH
5.0
MWC297 3
1
F
B2
n
n
12.9
PH
12.3
F
B2
a
y
y
87.4
PH
89.9
SIS: 0.8–10 keV
Serp
4/2 †3
R CrA 1
4
F/B
B
n
n
38.4
PH
39.4
R CrA 2
4
F/B
B
n
n
35.5
PH
36.9
R CrA 3
1
F
B2
n
n
65.2
PH
48.8
R CrA 4
1
F
B2
n
n
10.6
PH
11.7
R CrA 5
1
F
B2
n
n
17.1
PH
16.2
R CrA 6
1
F
B2
n
n
31.5
PH
34.1
HD200775
1FB2
0.45/0.55
n
n
46.0
PH
48.7
SIS: 0.55–10 keV
MWC1080
1
F
B2
n
n
36.1
PH
38.4
Mon R2
4/2
F
B2
b
y
y
62.5
PH
70.6
Chip: Chip mode, Format: SIS data format (Obs: observation, Ana: analysis, F: Faint, B: Bright, B2: Bright2, n.u.: notused), LD: Level Discri (a: s0[0.48→0.7], s1[0.55→0.69], b: 0.55→0.70/1.10), RC: RDD correction, TS: Telemetry Saturation,Exp.: Average exposure time, Mode: GIS mode, “subject1 / subject2” means “bit High / bit Medium & Low”, †1 Using anarchived image, see Table 4.9, †2 See Table 4.9. †3 non-standard chips
Chip

Det.

Temp.
◦C
−60.9

Chip
C0
rdd9704 s0c0m4 t61 62
rdd9702 s0c0m2 t61 62
rdd9703 s1c0m4 t61 62

Used RDD map
C1
C2
rdd9704 s0c1m4 t60 61 rdd9704 s0c2m4 t60 61

Serp

s0

C3
4
rdd9704 s0c3m4 t60 61
2
rdd9702 s0c3m2 t61 62
rdd9703 s1c1m4 t61 62 rdd9703 s1c2m4 t61 62 rdd9703 s1c3m4 t61 62
s1
−60.2
4
2
rdd9702 s1c1m2 t60 61 rdd9702 s1c2m2 t60 61
Gal R2s0−61.6
4
rdd9609 s0c0m4 t61 62 rdd9609 s0c1m4 t61 62 rdd9609 s0c2m4 t61 62 rdd9609 s0c3m4 t61 62
s1
−61.6
4
rdd9609 s1c0m4 t61 62 rdd9609 s1c1m4 t61 62 rdd9609 s1c2m4 t61 62 rdd9609 s1c3m4 t61 62
HD200775 † s0
−61.3
2
rdd9511 s0c0m2 t60 61 rdd9511 s0c1m2 t60 61
s1
−60.9
2
rdd9511 s1c2m2 t0 60
rdd9511 s1c3m2 t0 60
ChamI Vs0
−61.6
2
rdd9908 s0c1m2 t61 62 rdd9908 s0c2m2 t61 62
s1
−61.4
2
rdd9908 s1c0m2 t61 62
rdd9908 s1c3m2 t61 62
Lupus3s0−61.4
4
rdd9603 s0c0m4 t61 62 rdd9603 s0c1m4 t61 62 rdd9603 s0c2m4 t61 62 rdd9603 s0c3m4 t61 62
s1
−61.1
4
rdd9603 s1c0m4 t61 62 rdd9603 s1c1m4 t61 62 rdd9603 s1c2m4 t61 62 rdd9603 s1c3m4 t61 62
IC359
s0
−61.0
4
rdd9509 s0c0m4 t61 62 rdd9509 s0c1m4 t61 62 rdd9509 s0c2m4 t61 62 rdd9509 s0c3m4 t61 62
s1
−60.5
4
rdd9509 s1c0m4 t61 62 rdd9509 s1c1m4 t61 62 rdd9509 s1c2m4 t61 62 rdd9509 s1c3m4 t61 62
Mon R2s0−60.2
4
rdd9703 s0c0m4 t61 62 rdd9703 s0c1m4 t61 62 rdd9703 s0c2m4 t61 62 rdd9703 s0c3m4 t61 62
2
rdd9703 s0c0m2 t61 62 rdd9703 s0c1m2 t61 62 rdd9703 s0c2m2 t61 62 rdd9703 s0c3m2 t61 62
s1
−59.6
4
rdd9703 s1c0m4 t61 62 rdd9703 s1c1m4 t61 62 rdd9703 s1c2m4 t61 62 rdd9703 s1c3m4 t61 62
rdd9703 s1c2m2 t61 62 rdd9703 s1c3m2 t61 62
2
rdd9703 s1c0m2 t60 61
Det.: Detector, Temp.: CCD temperature, File names of maps are [map name].ﬁts., † : RDD correction is not performed. See Table 4.9.

Obs. ID

Table 4.4: Calibration ﬁles used for the RDD correction

28
CHAPTER 4. OBSERVATION & DATA REDUCTION

4.2. REDUCTION & SCREENING METHODS OF THE ASCA DATA

4.2

29

Reduction & Screening Methods of the ASCA Data

We used the ASCA archival data processed through the method, “revision2”1 . All theobserved quantities such as the arriving time, coordinate and pulse height of each event arelisted in rev2 ﬁles. We thus did not perform further data processing before event selection
unless the RDD correction is required. We processed events selection with the software
package FTOOLS 4.2.

4.2.1

RDD Correction

We dealt with the RDD correction for SIS data taken after 1995 in 2 and 4 CCD FAINT mode,using the software package “correctrdd” (see the section 3.1.3). A pixel-height distribution
on a RDD map depends on both observing date and CCD temperature, but RDD maps
in some condition was not taken owing to the operational constraints. We thus searched
for RDD maps taken under the same CCD temperature within 2 months before and afteran observation, and if no RDD map exists, we used RDD maps taken under diﬀerent CCD
temperature. Used RDD maps are shown in Table 4.4.
For the RDD corrected data, we dealt with the DFE correction using “correctdfe” andthen converted the data format to BRIGHT or BRIGHT2 modes using the softwares “f2b” or“f2b2” respectively, which simultaneously process the ECHO correction. These procedures
are automatically done with the script “gordd.alpha” provided by GSFC.

4.2.2

Filtering Criteria

We referred selection criteria to the ASCA Data Reduction guide (ASCA ABC guide, Arida
and ASCA team 2000) and applied the most lax values. We conﬁrmed that our loose selection
criteria do not aﬀect the data quality, by investigating background levels in diﬀerent cutoﬀ
energy. The criteria are listed in Table 4.5

4.2.3

Telemetry Saturated Data

Because dummy events from hot pixels increase in high CCD temperature (T > −60◦ C), wesometimes have to take telemetry saturation into account for SIS data These events have
low PH and therefore can be rejected by the level discrimination in DP at the expense of
low-energy photons during the observation. The ASCA calibration team set up the standard
1

Revision2, or rev2 in short, is an improved analyzing method which do away with a lot of raw ﬁles madeby observing mode changes and arrange in a few unﬁltered event ﬁles for each observing mode. See ASCA
Getting Started Guide (Pier and ASCA Processing Team ) in detailCHAPTER 4. OBSERVATION & DATA REDUCTION

30

Table 4.5: Data selection criteria
Unit
SAA
COE
SIS(GIS) RBMF
T SAA & T DY NT
SFR & HFRELVBR EARTH
ANG DIST[GeV][readout cycle][degree][degree][degree]

SIS
ex.
4/6
ex.
4/inc.
inc.
520
0.01/inc.

GIS
ex.
4/6
ex.
inc.
ex./inc.
5
inc.
0.01/inc.

Right values are criteria of R CrA 1–6, MWC 297 1–3, Serp and Mon R2
SAA: SAA passage, COE: Minimum of cutoﬀ energy, SIS(GIS) RBMF:
SIS(GIS) radiation belt monitor ﬂag, T SAA(DY NT): Time after passing through SAA (the day-night terminator), SFR: Soft ﬂare region(G2 H0+G2 H2+G3 H0+G3 H2<45 & RBM CONT<100), HFR: Hard ﬂare
region (G2 H0+G2 H2+G3 H0+G3 H2<0.45×COR2−13×COR+125) ELV:
Elevation from the rim of the earth, BR EARTH: Elevation from the brightearth, ANG DIST: Maximum of angular distribution, ex.: excluding, inc.: including

discrimination level and ASCA duty scientists at KSC monitors telemetry saturation at
ASCA contact passes, but some data cannot avoid being telemetry saturated. When thetelemetry saturation occurs, events on the far side from readout gates are ﬁrst discarded (seethe section 3.1.1). We thus investigated whether events on a source position were discard ornot, and only if these events were discarded, we removed telemetry saturated CCD frames.

4.2.4

Event Extraction

We removed SIS events associated with ’hot’ and ’ﬂickering’ pixels using “sisclean” and thenpicked up all events of the ASCA grade 0234. We collected GIS events satisfying the normal
criteria using “gisclean” (Ohashi et al. 1996).

4.3

Basic Procedure of Data Analyses

Event data are reduced to image, spectrum and light curve data with XSELECT. Further
analyses are based on the XANADU software package provided by GSFC. This package comprises three tasks, XIMAGE for image analysis, XSPEC for spectral analysis and XRONOS
for timing analysis. We used XIMAGE ver. 2.53, XSPEC ver. 9.0 and XRONOS ver. 4.02.

4.3. BASIC PROCEDURE OF DATA ANALYSES

31

Table 4.6: ROSAT observation log
TargetSeq. ID

CoKu Tau 1
HD245185
V380 Ori
BF Ori
MWC120VY MonRP200001A01
RH202047N00
RH900010A00
RH900010A01
RP900189N00
RP900355N00
RP900355A01
RP900355A02
RP900355A03
V590 Mon
RH200130A00
LkHa218
RP201011N00
LkHa220
RH202153N00
MWC 297
RH202048N00
HD200775
RH202319N00
Det.: P−PSPC, H−HRI,

y
91
95
91
92
91
92
93
93
94
91
93
96
95
97

Obs. Time
Start time
m d
h m s
8 20 23 30 52
3 16 13 07 53
3 22
5 57 39
3 14
2 24 42
9 19 20 05 57
10 6 13 14 2
3 14
1 34 50
9 16 22 43 30
3 29 23 8 27
3 20
5 43 32
4 22 17 38 19
4 6 15 42 06
9 28 21 28 37
8 11 19 51 22

Det.
Exp.
sec
25,591
9,380
11,037
4,835
24,154
5,564
4,598
5,354
5,117
19,567
19,66533,495
1,089
4,799

P
H
H
H
P
P
P
P
P
H
P
H
H
H

This package is convenient in performing standard analyses but less ﬂexible especially on
image analyses. We then utilized the Display45 software package (ver. 1.90), which was originally developed by PEP4/SLAC group and optimized for ASCA data analyses by ASCAteam members (Ishisaki et al. 1997).

4.3.1

Imaging Analysis

For each observation, we made whole x-ray band images of both SIS and GIS in the J2000
coordinate system (SIS: 0.5–10 keV, GIS: 0.8–10 keV). XSELECT automatically cuts oﬀ
outer rim of GIS images in the sky coordinate, but some HAeBes are located on this outerrim. In this case, we analyzed GIS data using Display45. Gotthelf (1996) estimated position
accuracy on SIS images at ∼ 40 (see the subsection 3.1.1). Gotthelf et al. (2000) found asystematic modulation of the ASCA absolute attitude by referencing AGNs, which depends
on temperature of supports of the star trackers. New method reduces the position uncertainty
to ∼12 for SIS. We used this latest method for all the imaging analyses.
We searched for an X-ray source above 5σ detection level around a optical position ofa HAeBe. We simultaneously made use of ROSAT images for source identiﬁcation because
low mass stars are usually present around HAeBes. We ﬁrst referred ROSAT images toZinnecker and Preibisch (1994) and Preibisch (1998). For HAeBes not in these lists, we
retrieved processed ROSAT images from HEASARC (rationalized data format: RDF, see
Table 4.7). For HRI, there is only one total band image, while for PSPC there are three
images energy-sliced in the total (im1), hard (im2) and soft band (im3). We basically used
total band images.
For non-detected HAeBes, we determined upper-limit ﬂuxes. We measured source

CHAPTER 4. OBSERVATION & DATA REDUCTION

32

Table 4.7: RDF parameters of the ROSAT data

Low
(ch)

PI
High
(ch)

PSPC
Total
Hard
Soft

(im1)
(im2)
(im3)

11
52
11

235
201
41

HRI
Total

(im1)

1

16

EnergyLow High(eV) (eV)
110
520
110

2360
2020
420

upper-limit counts in 99% conﬁdence levels (3σ) from 20 half box for SIS and 40 half
box for GIS, by taking background regions from a nearby source free region, using sosta
package in XIMAGE. This upper-limit counts were converted to ﬂux upper-limit assuming
an absorbed thin-thermal model kT = 1 keV, abundance = 0.3 Z and NH estimated fromAV . The assumption of plasma temperature, however, does not strongly aﬀect the results
because source ﬂux is less dependent on plasma temperature (less than factor of two fortypical temperature, kT ∼ 0.5−4 keV, see Table 4.8). We also saw LX values derived fromROSAT results in other articles and, if there is no article, estimate LX from the ROSAT
archival images with the ximage package. The method is basically the same as the method
for the ASCA data, but the background level is estimated from surrounding regions (12×12 ) with the ‘background‘ command.
For detected sources, we normally took wider source regions with 2.5 or 3 radius circles
for timing and spectral analyses (see Table 5.3). We selected background regions from nearby
source free region or, when contaminating sources are present, a symmetrical region withrespect to the contaminating sources.

4.3.2

Timing Analysis

We made source light curves for each detector by subtracting normalized background light
curves using the “lcmath” package. We then merged light curves of same type of detectors,
but we separately analyzed light curves of diﬀerent types of detectors to avoid artiﬁcial gaps
caused by the diﬀerent time selection. All light curves were basically made in the wholeenergy band. Default minimum time bin is 2048s, but it is longer for weak sources. We
tested time variability by the χ2 -test method.

4.3. BASIC PROCEDURE OF DATA ANALYSES

33

Table 4.8: Flux converting ratio from the ROSAT and Einstein band to the ASCA band
kT
ROSAT
Einstein(keV) 0.5–10 2–10
4–10
0.5–10 2–10
4–10
0.5
0.69
0.010 0.00015
0.75
0.011 0.00016
1.0
0.73
0.087 0.0089
0.78
0.093 0.0095
2.0
0.97
0.36
0.11
0.90
0.33
0.10
5.0
1.57
0.95
0.52
1.32
0.80
0.43
Values are ratio of ASCA ﬂux from a unit ROSAT (0.1–2.4 keV)or Einstein ﬂux (0.16–3.5 keV).

4.3.3

Spectral Analysis

Spectral ﬁttings were basically performed in 0.5–10 keV for SIS and 0.8–10 keV for GIS.
Response matrix generators were made with “ascaarf” (ver. 2.81), sisrmg (ver. 1.1) for SIS.
Whereas, standard GIS response matrices are used for GIS2. As for a plasma model, we adopt
an absorbed thin-thermal plasma model because some sources clearly showed line featuresin their spectra. We used the MeKaL plasma code, which well treats ﬁne structures of Fe−L
lines (Mewe et al. 1995, see also Matsumoto 1998 Appendix B). Chemical abundances in
the MeKaL model was ﬁxed to 0.3Z because there are not enough photons to determine
the chemical abundances. The stellar x-ray spectra taken by ASCA show low chemical
abundance around 0.3 Z (eg OB stars Kitamoto and Mukai 1996; Kitamoto et al. 2000;
low mass MSs Tagliaferri et al. 1997; low mass PMSs Yamauchi et al. 1996; Kamata et al.
1997; Tsuboi et al. 1998). Though the reason is not known, the chemical abundance 0.3Z
seems to be the most appropriate value for stellar x-rays. At ﬁrst, we tried one-temperaturethermal model, and when this model is rejected, we retried two-temperature model with
common NH .
SIS event data available through the level discrimination is complicated. Onboard frontend processors simply discard events of center pixel height under a programmed thresholdvalue. Splitting events whose total pulse-heights exceed this values are also discarded. Thisdegrades eﬀective quantum eﬃciency just above this threshold energy. These eﬀects aretaken into account for making SIS response matrices, but the RDD correction makes thisproblem more complicated. As a criterion, we accepted a 10% error in quantum eﬃciency
and brieﬂy set the lowest ﬁtting energy like this. We made a histogram of a used RDD
map and then measured ∆PH including 90% of the total pixels. By approximating that
the RDD distribution is symmetrical, we added ∆PH on the PH of level discrimination anddetermined it as the lower threshold energy.

2

gis[23]v4 0.rmf

34

CHAPTER 4. OBSERVATION & DATA REDUCTION

Table 4.9: Speciﬁc subjects related to observations and analyses
IC 359 ... 95% in the BRIGHT mode were operated during the satellite day. The data taken
under this condition is less reliable so that the FAINT data were only analyzed.
L1641N ... GIS3 accidentally outputed the spectral column with only 7bit length. V380 Ori isnear the calibration source in GIS3 so that the GIS2 data was only used.
PSRJ0631 ... The rise time bits were not assigned in the GIS event data format, in compensation
for the precise determination of event arrival time for pulse detection. The background levelis thus higher than that of ordinary observations.
Lupus3 ... The SIS telemetry saturation did not damage the data quality of HR 5999/6000. The
BRIGHT mode data taken in the former observation, when the CCD temperature is relatively
high, was terribly bad. The BRIGHT data was therefore discarded.
Gal R 2 ... For the ﬁrst 5 ksec in the satellite-day condition, SIS was operated in the BRIGHT
mode. XSELECT reported inconsistency in the OPTICS values for the SIS data.
Serp ... Background levels of both GIS detectors bursted up at 120 ksec. They were detected on
the outer side of GIS and did not damage the EC 95 data.
R CrA 1 ... GIS3 outputed wrong PH values in lower 3 bits. The spectral bins were thus summed
by 8 channels.
R CrA 4 ... The GIS2 front end processor stopped during the observation and more than half of
GIS2 data were lost. A data station on ground failed to contact with ASCA at the end of
the observation.
HD 200775 ... The observed exposure was shorter than the planned. The RDD correction was
not processed because the quantum eﬃciency improved less than 4%.
Ori Trap2 ... For SIS, XSELECT did not run properly in the process of cleansis. We downloaded
the screened event ﬁles from HEASARC. Its ﬁltering criteria are thus HEASARC basis.
Ori Trap3 ... The exposure was far shorter than the planned because the down-link from ASCA
failed by the antenna trouble at KSC.
ChamI IV ... SIS1 was saturated owing to hot CCD temperature (Ts1 ∼ −59.3◦ C) and completely lost the data around HD 97048, but SIS0 managed to keep them.
ChamI V ... Nonstandard CCD chips were used (S0/C1C2, S1/C3C0). The address discrimination of S0C2 for light leakage was mistakenly unset during the former one-thirds of theobservation, but it did not aﬀect the data analyses of HD 97048.
Gal 1 ... Stray light from a strong x-ray source covers the whole fov of both detectors
SVS13 ... During the former half of the observation, SIS data are saturated.

Chapter 5
Result I – Herbig Ae/Be StarsThis section describes the observing results of ASCA detected and non-detected HAeBes,
separately in the ascending order of right ascension. The data of the imaging, timing andspectral analyses are summarized in Table 5.1, 5.2, 5.4 and 5.5.

5.1

Detected Sources

5.1.1

V892 Tau

V892 Tau, often designated as Elias 3-1, is located in the L1495E dark cloud. The stellar
parameters are open to dispute; Strom and Strom (1994) quoted it as a B9 star of Lbol ∼64.3
L , while Zinnecker and Preibisch (1994) quoted it as an A6e star of Lbol ∼38 L . It showsan unusual infrared spectral with anomalous 3.43 µm and 3.53 µm features and an 11.2 µm
emission peak. (eg Whittet et al. 1984; Hanner et al. 1994). Millimeter interferometry by diFrancesco et al. (1997) certiﬁed the presence of a circumstellar disk with the mass of ∼0.1
M . It has a companion star (Elias I NE) at ∼4 north east from V892 Tau, which would
be a WTTS with Lbol ∼0.3 L (Skinner et al. 1993; Leinert et al. 1997; Pirzkal et al. 1997).
Strom and Strom (1994) ﬁrst reported the x-ray emission, but the LX value is about twiceas large as that estimated by Zinnecker and Preibisch (1994) (∼2×1030 ergs s−1 ) from thesame observing data. Zinnecker and Preibisch (1994) also measured the plasma temperature(kT = 2.3±0.4 keV) and the column density of hydrogen (NH = 4.8±1.1 ×1021 cm−2).
ASCA detect a strong x-ray source within the error circle of V892 Tau (Figure 5.1).
In the SIS image, two other x-ray sources are seen close to V892 Tau. The bright one at∼2 north-east from V892 Tau is V1023 Tau and the very faint one at 30 south-east isV410 x-ray 7. V410 x-ray 7 is a reddened M0.5 pre-main-sequence star (AV ∼5.9; Briceñoet al. 1998) probably belonging to the same nebulousity as V892 Tau, and it showed an x-ray
35

36

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.1: ASCA images of the V892 Tau/ CoKu Tau 1 ﬁeld (12 ×12 fov: upper left: SIS
total band, upper right: GIS total band, lower left: SIS soft band, lower right: SIS hardband, See the section 5.2.2 for CoKu Tau 1). These images are drawn with contours of total
photon counts and smoothed with a Gaussian function of σ ∼2 pixel. Solid-lines represent
photon counts per pixel above the background level and are drawn by 0.4 cts/pixel/obs
for SIS and 1.0 cts/pixel/obs for GIS. Dot-dash lines represent 5σ (SIS) and 3 & 5σ (GIS)
detection levels, which are calculated from the Poisson-distributed photon error of selectedbackground counts (see the table 5.3). Dash-lines in SIS images are CCD chip gaps and thatin the GIS images (not in this ﬁgure). Star marks are the HAeBe positions (center: V892
Tau, left: CoKu Tau 1). Plus marks are optical or infrared source positions listed in the
SIMBAD database. Each ﬁlled circle on the bottom right of each panel shows the typicalerror circle (20 for SIS, 40 for GIS). The coordinate system is J2000.

5.1. DETECTED SOURCES

37

Figure 5.2: ASCA SIS light curves of V892 Tau after the background subtraction (left: total
band, right: hard band). Vertical axes are averaged detector count rates and dotted lines
are the zero level. 1 bin is 2,048 sec.

Figure 5.3: SIS spectra of V892 Tau (left: 1T component, right: 2T components including
V410 x-ray 7). In each panel, the best-ﬁt model is drawn with a solid line. Components inthe 2T model is drawn with dotted lines. Each residual from the best-ﬁt value is displayedin its lower section.

38

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.4: ASCA GIS (left) and ROSAT HRI (right) images of the V380 Ori ﬁeld. Detailsof the ASCA image are the same as Figure 5.1, but the count rate per solid line is half
(0.2 cts/pixel/obs for SIS and 0.5 cts/pixel/obs for GIS). In the ROSAT image, lines andsymbols are used in the same manner. The image is smoothed with Gaussian function withσ ∼ HPD at the source position (>8 : binned pixel size). Solid-lines are drawn by 20 counts
per 16×16 pixels (0.078cts/pixel). 3 and 5 σ levels are drawn with Dot-dash lines.

ﬂare (Stelzer et al. 2000). Because GIS image cannot separate V1023 Tau from V892 Tau,
we only proceed the SIS analyses. The peak of V410 x-ray 7 is obscured in the total band
(0.55−10 keV), but it is clearly seen in the soft band (0.55−2 keV).
Figure 5.2 shows light curves of V892 Tau in the total band (0.55−10 keV). The source
count rate seems to ﬂuctuate within a factor of two, but it accepts a constant model within
90% conﬁdence level. The spectrum can be reproduced with an absorbed 1T thermal model,
but it obviously includes, at least, a component from V410 x-ray 7. The dominant source
is V892 Tau in the hard band (2−10 keV, see Figure 5.1), Because the high energy part in
the x-ray spectrum determines the plasma temperature, we ﬁx the plasma temperature atthe 1T best-ﬁt values by regarding it as the temperature of V892 Tau. In the soft band,
the peak counts of V892 Tau is 1.1 times as large as those of V410 x-ray 7. The soft bandﬂux reduces to be half when NH ∼1.8×1022 cm−2. We thus ﬁx these parameters and add
an absorbed 1T component. We then found an acceptable solution of kT. The luminosityof V410 x-ray 7, LX ∼5×1030 ergs s−1 (0.5−10 keV), is half of the ROSAT estimate in aquiescent state (Stelzer et al. 2000). We also made light curves in the hard band (2−10 keV)
to minimize the contribution of V410 x-ray 7 (<10% from the peak ﬂux). The light curves
accept a constant model as well.

5.1. DETECTED SOURCES

39

Figure 5.5: GIS light curves (left) and spectrum (right) of V380 Ori. Details can be seen in
Figure 5.2, 5.3

5.1.2

V380 Ori

V380 Ori belongs to the star forming cloud L1641 (eg Strom et al. 1989). From a studyof the multi-wave lengths in UV, optical and IR, V380 Ori is thought to be a B9–A0 starof 3.3−3.6 M , showing small ampliﬁcation variability (Hillenbrand et al. 1992; Boehm
and Catala 1995; Rossi et al. 1999; Herbst and Shevchenko 1999). It has a companion
star 0.15 apart from the primary (Corporon and Lagrange 1999). Recent study of V380
Ori is summarized in Rossi et al. (1999). The x-ray emission was detected with Einstein(LX ∼ 1.1×1031 ergs s−1 , Pravdo and Marshall 1981) and was conﬁrmed with ROSAT (LX∼ 2.3×1031 ergs s−1 , Zinnecker and Preibisch 1994). Damiani et al. (1994) evaluated theplasma temperature and absorption column density using the Einstein data (kT ∼ 1.1 keV,NH ∼ 7.4×1021 cm−2 ).
ASCA detect the x-ray emission from V380 Ori on the edge of GIS fov as a relatively
faint source (Figure 5.4). There seems to be a small peak on the north tail of V380 Ori. but
there is no source in ROSAT HRI. Light curves (Figure 5.5) show no apparent variability. In
the spectrum (Figure 5.5), the best-ﬁt temperature is relatively high (kT ∼3.2 keV), but it
has a large error range from 0.65 keV to 11 keV. The luminosity is nearly the same as thoseof Einstein and ROSAT.

5.1.3

VY Mon & CoKu VY Mon/G2

Both VY Mon and VY Mon G2 are located at the 85 south from the reﬂection nebular
IC 446, which is illuminated by the B2.5 star IC 446 No.1 (eg Casey and Harper 1990).

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

40

Figure 5.6: ASCA GIS (left) and ROSAT PSPC (right) images of the VY Mon region.
Symbols are the same as Figure 5.1, 5.4, but the count rate per solid-line in the ASCA
image is half (0.2 cts/pixel/obs for SIS and 0.5 cts/pixel/obs for GIS). The right star markin each image is VY Mon and the left is CoKu VY Mon/G2. Five plus marks surrounding
the star marks are VY Mon/G1, G4, G3, LkHα-274 and G5, anti-clockwise from right.
They form a small group of stars and are memberships of the Mon OB1 association (Cohen
and Kuhi 1979; Testi et al. 1998). VY Mon is a O9−B8 star in the very young phase,(∼5×104 yrs, Casey and Harper 1990; The et al. 1994; Testi et al. 1999). It shows Algoltype variability and a highly reddening feature with large IR excess (AV ∼7.4). CoKu VY
Mon/G2 is a A0 shell-type star without emission line (Cohen and Kuhi 1979; The et al.
1994). X-ray emission from this region was ﬁrst suggested with Einstein though the source
signiﬁcance was very small (S/N ∼2.5, Damiani et al. 1994). It is uncertain which was the
x-ray source, but the emission was only detected in the Einstein hard band, suggesting thesource is deeply embedded.
There are large discrepancy of the sky coordinate of VY Mon between the SIMBAD
database and recent papers (Testi et al. 1998; Henning et al. 1998; Yudin and Evans 1998).
This is probably because SIMBAD refers the coordinate to an old paper less reliable inabsolute coordinate. We thus refer the absolute coordinate of VY Mon to the recent papers
and determine the coordinates of other sources from the relative coordinates to VY Mon1 . A
strong x-ray source is detected in the VY Mon region (Figure 5.6). The separation from theoptical position is about 40 , which is larger than the typical error (∼24 ) but acceptable for
the GIS outer fov. ROSAT observed this region ﬁve times, but in each observation the oﬀaxis angle at VY Mon is more than 30 and the HPD is relatively large (∼ 90 ). No apparent
1

The relative position to VY Mon is referred to Weintraub (1990) for CoKu VY Mon/G2, /G3, /G4 andto Herbst et al. (1982) for IC 446 and LkHa 274. There is no reference for CoKu VY Mon/G1 except for a
ﬁgure in Cohen and Kuhi (1979). The relative coordinate is roughly measured from this ﬁgure.

5.1. DETECTED SOURCES

41

Figure 5.7: GIS light curves of VY Mon. See Figure 5.2 for details.

source appears in the ROSAT band and no signiﬁcant variation occurs between observations(Figure 5.6 shows a hard band image taken in March 1994). X-ray emission concentrate on
the stellar cluster in every observation, but we did not identify the counterpart even in the
ROSAT images.
The light curve observed with ASCA gradually decreases with small ﬂuctuations andbecomes almost constant after 6×104 sec (Figure 5.7). We split the observation into twophases: 6×104 sec – high state (before) and low state (after). The low state spectrum rejects
an absorbed 1T model in 90% conﬁdence level, which is caused by a line-shaped structure
at kT ∼5 keV. We added a Gaussian component and found an acceptable ﬁtting result. Theline center energy is 5.1 keV, which is close to highly ionized Ca line energies. We thus triedan absorbed MeKaL model with free Ca abundance, but the Ca line appears lower than 5.1
keV and this model is rejected again. The line center energy corresponds to the Kα line of
neutral vanadium, but this element is scarce in the universe. The x-ray luminosity in the low
state is log LX (ergs s−1 ) ∼32.2 (d ∼800 pc), which is one of the largest luminosity among
our HAeBe samples. On the other hand, the high state spectrum accepts an absorbed 1T
model in 90% conﬁdence level. The best-ﬁt temperature is very high, kT ∼6 keV. The 1T
model assumes that there is only one plasma blob around a star, which is very hot during thehigh state and cools down in the low state. The fact that the emission measure varies lessthan 20% between two states supports this model (see also Figure 7.9). Another situation is
also possible – two hot plasma components are present in the high state and one disappers
in the low state. We also found an acceptable solution for the model, by adding the low
state component. All ﬁtting results showed a relatively small NH . The converted AV value- ∼5×1021 cm−2 (AV ∼2.5) - corresponding with VY Mon/G2, G3, G4. VY Mon/G3 and/G4 are classiﬁed as spectral types K−M (Cohen and Kuhi 1979). LX of low mass stars are
typically less than log LX (ergs s−1 ) ∼30. VY Mon/G2 is the most plausible counterpart.

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

42

Figure 5.8: GIS spectra of VY Mon. Top panels show the spectra in the low state (left: anabsorbed 1T best-ﬁt model, right: adding a Gaussian component). Bottom panels show the
spectra in the high state (left: an absorbed 1T best-ﬁt model, right: adding the low statecomponent).

5.1.4

T Cha

T Cha is in the dark cloud DC 300.2–16.8 locating between the Chameleon I and II clouds on
the sky, but the distance measured with HIPPARCOS (d ∼66 pc) is closer than the distances
to the Chameleon I and II. T Cha is suspected to be a Weak-line T-Tauri star (Alcala et al.
1993). Hα lines show strong variability on a time scale of a day, sometimes accompanying
with inverse P Cygni features (Alcala et al. 1993).
T Cha was detected as a strong x-ray source in the error circle. There is no opticalcounterpart other than T Cha and therefore T Cha is the most probable counterpart ofthe x-ray source. The ﬂux is variable and seems to be gradually increasing through the5.1. DETECTED SOURCES

43

Figure 5.9: ASCA images (left: SIS, right: GIS) of the T Cha ﬁeld. Symbols are the same as
the captions in Figure 5.1, 5.4, but lines above 20 solid-lines in SIS are shown with dottedlines by 2cts/pixel/obs.
observation, with small ﬂare like events (t ∼60, 130 ksec, Figure 5.10). The spectra aresuccessfully reproduced with an absorbed 1T model (Figure 5.11). The log LX /Lbol value isthe largest among the observed sources.

5.1.5

HD 104237

HD 104237, located in the Chamaeleon III dark cloud, is in the vicinity of the sun (d ∼ 116pc) and hence the brightest HAeBe star in the sky (V =6.6 mag, Hu et al. 1991). It is aslightly old star (Age ∼6.3 yrs) with the spectral type A4IVe+sh (Knee and Prusti 1996;
van den Ancker et al. 1998). It shows P Cygni proﬁles of MgII lines but it does not have CO
molecular outﬂows (Knee and Prusti 1996). It also emit several emission lines, indicative of
the existence of a hot chromosphere (Hu et al. 1991). The X-ray emission was detected in the
ROSAT all-sky survey as a relatively bright source (Alcala et al. 1995). The spectroscopic
analysis with the ASCA deep observation have been performed by Skinner and Yamauchi(1996). Our data is the same as those of Skinner and Yamauchi (1996) so that the results
is basically the same except for adopting the distance with HIPPARCOS (van den Ancker
et al. 1998).
ASCA detect a strong x-ray source at the position of HD 104237 (Figure 5.12). Thereis no other counterpart within the error circle, and thus the x-ray source is identiﬁed as
HD 104237. The x-ray ﬂux ﬂuctuates at the middle of the observation (Figure 5.13). A
constant model is rejected in the 90% conﬁdence level. Spectral ﬁttings with an absorbed
1T model has large residual around 1keV (Figure 5.13). We thus tried an absorbed 2Tmodel with common NH . Then the reduced χ2 value greatly improves from 3.82 to 1.57,

44

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.10: Light curves of T Cha (top: SIS, bottom: GIS). See Figure 5.3 for details.

Figure 5.11: Spectra of T Cha (left: SIS, right: GIS). See Figure 5.3 for details.

5.1. DETECTED SOURCES

45

Figure 5.12: ASCA images of HD 104237 (left: SIS, right: GIS). For symbols, see Figure 5.1.

but the model is still rejected in the 90% conﬁdence level. We also try other three models;
2T thermal model of free abundance with common NH , 2T model with diﬀerent NH values,and 3T model with a common NH , but their reduced χ2 values do not improve. Skinner andYamauchi (1996) showed acceptable results in the 2T and 3T model for the SIS spectrum.
Bad reduced χ2 values in our ﬁttings could be responsible for the inconsistency between SIS
and GIS. However, we could not ﬁnd any systematic diﬀerence between these spectra.

5.1.6

IRAS 12496−7650

IRAS 12496−7650 is located in the Chameleon II Dark cloud. It is an embedded Herbig Aestar (Hughes et al. 1989) and is the most luminous far-infrared source in the Chameleon II
cloud. Due to its extremely large extinction, the spectral type is not well established (A−F,
Hughes et al. 1991). It is suggested to be the exciting source for the group of Herbig-Haro
objects, HH 52-54, which places this object at the very young stage of stellar evolution(Cohen and Schwartz 1987).
We discover a highly absorbed x-ray source within 1 error radius at the position of IRAS
12496−7650 (Figure 5.15). Because no other cataloged source is present around there, thex-ray emission is most probably due to this peculiar object. We do not ﬁnd any prominent
x-ray variation within the limited photon counts (Figure 5.16). One important discovery is
its very high absorption of about NH ∼ 1×1023 cm−2, which exceeds 5-times the estimatedoptical extinction of IRAS 12496−7650 (Figure 5.16, Hughes et al. 1991). The x-ray emitting
region might be localized near to the star, and would be in a selected geometry to have large
absorption, possibly due to the circumstellar disk.

46

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.13: Light curves (left: SIS, right: GIS) of HD 104237. See Figure 5.2 for details.

Figure 5.14: Spectra of HD 104237 (upper left: SIS 1T, lower left: SIS 2T, lower right: GIS
2T). See Figure 5.3 for details.

5.1. DETECTED SOURCES

47

Figure 5.15: ASCA GIS image of IRAS 12496-7650. Symbols are the same as the caption in
Figure 5.1. The count rate per solid-line is half (0.5 cts/pixel/obs).

Figure 5.16: Light curves (left) and spectrum (right) of IRAS 12496–7650. See Figure 5.2,
5.3 for details.

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

48

5.1.7

HR 5999 / HR 6000

HR 5999 and HR 6000 are located in the center of the Lupus 3 subgroup and close toeach other on the sky (∼45 ). Both distances are approximately the same ; 210 pc for
HR 5999 and 241 pc for HR 6000 (Bertout et al. 1999). HR 5999 and HR 6000 would
be thus physically related to each other (Andersen et al. 1984)2 . HR 5999 is a variable
A5-7III/IVe star in 5×105 -year-old (van den Ancker et al. 1998; Tjin A Djie et al. 1989). It
is thought to have an accretion disk (eg The et al. 1996), which might cause photometricaland spectral variations. It also has a close companion at 1.3 apart (Rossiter 3930, Stecklum
et al. 1995), probably in the T-Tauri phase. On the other hand, HR 6000 is known to
exhibit peculiar line features, which is recognized as abundance anomalies (Andersen et al.
1984; Castelli et al. 1985). These lines make it diﬃcult to determine the spectral types andevolutional stage. From a model calculation, the spectral type was measured as late-B toearly A (Castelli et al. 1985; Kurtz and Marang 1995). The detailed photometry and multiband SED by van den Ancker et al. (1996) suggested the existence of a T-Tauri companion.
Their paper interprets the abundance anomalous as a result of several atomic lines merging
from the primary and secondary and concluded that the primary is a B6 star in a young
main-sequence phase. However, the distance of HR 6000 was later modiﬁed to be 241 pc
with HIPPARCOS (Bertout et al. 1999). This makes Lbol large and then HR 6000 should bea pre-main-sequence star of about 4M . The x-ray emission of both HR 5999 and HR 6000
were ﬁrst detected in the ROSAT observation as bright sources (Zinnecker and Preibisch1994).
In both SIS and GIS images (Figure 5.17), HR 6000 was clearly seen as a strong source.
Because HR 5999 is on the outskirts of the psf of HR 6000, only SIS could manage to detecta faint x-ray peak of HR 5999.
The light curves and spectrum (top left in Figure 5.18 and Figure 5.19) at the position of
HR 5999 are made from the SIS data. We do not ﬁnd any systematic diﬀerence between the
SIS0 and SIS1 data. We think that the contamination from HR 6000 is satisfactory removed.
The light curves of HR 5999 show no variable feature. The spectrum is successfully ﬁt witha 1T thermal model. The best-ﬁt NH value is very small but its error range includes theconverted NH value from the relatively large AV (∼0.5). The error bar of the temperatureincludes the hot component in the ROSAT spectra (kT ∼1.7 keV, Zinnecker and Preibisch1994). The best-ﬁt temperature derived with the ASCA data gets closer to the ROSAT
value when we assume the same NH (NH ∼7.5×1020 cm−2 , 1Tf of HR5999 in Table 5.5).
The light curves of HR 6000 show no time variation (top right in Figure 5.18) as well. For
the spectrum (top right in Figure 5.19), we tried an absorbed 1-T thermal model but it isrejected because of having an excess in the hard energy band. We then analyzed the GIS
data because GISs have better eﬃciency in the hard band and have longer exposure (see the2

The distance of ∼200 pc is far larger than the distance of the Lupus 3 cloud (140±20 pc; Hughes et al.
1993). They could be out of membership of the Lupus 3 cloud.

5.1. DETECTED SOURCES

49

Figure 5.17: ASCA images (left: SIS, right: GIS) of HR 5999 (lower star mark) and HR 6000
(upper star mark). Top panels and a bottom left (3–10 keV) panel show images during thelow state. A bottom right panel shows an image during the high state. Solid-lines of the SIS
3–10 keV images are drawn by 0.1cts/pixel/obs and those of the GIS high state image aredrawn by 0.5cts/pixel/obs. See Figure 5.1 for symbols.

50

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.18: SIS light curves of HR 5999 (top left) and HR 6000 (top right) and the GIS
light curve of HR 5999 + HR 6000. See the caption in Figure 5.2 for details.
table 4.9) At the beginning of GIS light curves (bottom left in Figure 5.18), the x-ray ﬂux
suddenly increased (τinc < 1,000 sec) by a factor of two, had a long plateau (τpla ∼1×104sec) and then rapidly decreased (τdec <1,000 sec). The ﬂux level during the other period
were almost constant. We call the period during the high as “high state” and the others as“low state”. The low state includes the SIS exposure window. The x-ray peak position in
the high state image is the same as in the low state (bottom right in Figure 5.17). HR 6000
is hence the most probable candidate of the variation source.
We made GIS spectra during the high and low state, separately (Figure 5.19). The GIS
source region of HR 6000 inevitably includes HR 5999. We therefore made spectral ﬁttingsof GIS spectra adding the HR 5999 component in the model derived from the SIS spectrum
though the ﬂux of HR 5999 is almost negligible. The low state spectrum ﬁt with an absorbed
1T thermal model has a distinct hard residual above 3 keV, which is not clearly seen in the
SIS spectrum. We therefore checked the SIS hard band image above 3 keV. HR 6000 is thenwithin the error circle of this hard peak though the T-Tauri stars CD-39 10894B is another5.1. DETECTED SOURCES

51

Figure 5.19: SIS spectra of HR 5999 (top left) and HR 6000 (top right) and GIS spectra of
low state spectra (1Th : middle left, 2Thi : middle right) and high state spectra (1T: bottom
left, 1Th : middle right). See Figure 5.3 for details.

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

52

Figure 5.20: ASCA images (left: SIS, right: GIS) of V921 Sco. Meanings of symbols are thesame as in Figure 5.1. The count rate per solid-line is half (0.2 cts/pixel/obs for SIS and 0.5
cts/pixel/obs for GIS).

candidate. We then gave an additional hard components with a common NH and ﬁnd an
acceptable solution. We check the consistency of the derived model with the SIS spectrum
and ﬁnd that the SIS spectrum also has good agreement with the GIS spectrum above 1
keV. The reason for poor agreement below 1 keV is not sure but it could be due to the large
uncertainty of the SIS response matrices in the soft band (see the subsection 3.1.3). We
therefore adopt the ﬁtting results derived from the GIS spectrum. The GIS images suggestthat a source showing the ﬂux increase is or lies around HR 6000. An absorbed 1T thermal
model with the HR 5999 component has residue below 1 keV and above 4 keV, about thesame ﬂux during the low state. We therefore think that another plasma component is newly
made or appears during the high state and soon disappears at the end. We ﬁt the highstate spectrum again with adding the low state component and ﬁnd an acceptable solution.
The NH value is the same as that of HR 6000 during the low state, which indicates that thesource is HR 6000 itself.

5.1.8

V921 Sco

V921 Sco, often called as CoD −42◦ 11721, is located on the galactic plane [(l, b) = (343.35,−0.08)]. Some papers classiﬁed it as a supergiant (eg Hutsemekers and van Drom 1990),
but it is recently recognized as a young star. V921 Sco is embedded in a diﬀuse nebulousity,that spreads over 40 ×80 area. The distance to V921 Sco has large uncertainty from 200
pc (Pezzuto et al. 1997) to 2.6 kpc (Brooke et al. 1993). We adopt the distance of 500 pc
derived from a ISO−SWS spectral analysis (Benedettini et al. 1998).

5.1. DETECTED SOURCES

53

Figure 5.21: Light curves and spectra of V921 Sco. See Figure 5.2, 5.3 for details.
In the observation Gal R 2, a weak x-ray source is detected within both SIS and GIS
error circles. The source in the SIS image seems to split into two humps, but it is uncertain
as whether there are really two x-ray sources or simply made by limited photon statistics. In
Gal R 3, GIS also detects a plausible x-ray emission at the position of V921 Sco with signalto-noise ratio ∼7.0 (0.8−10 keV), but the data quality is bad because V921 Sco is on the edgeof GIS fov. We therefore only analyze the Gal R 2 data. The background subtracted light
curves (Figure 5.21) do not show any time variation. The GIS spectrum (Figure 5.21) canbe reproduced by an absorbed 1T model with an extremely large NH ∼5×1022 cm−2. Even
the lowest NH is 1.5 times larger than the converted NH from the largest AV measurement(7.1). It is however comparable to the column density derived from a millimeter survey (NH∼1023 cm−2 , Henning et al. 1998). The absorption corrected luminosity is quite large (LX∼1032 ergs s−1 ).

5.1.9

MWC 297

MWC 297 is a highly reddened star with extremely strong Balmer and silicate lines. Hillenbrand et al. (1992) reported that MWC 297 is a pre-main-sequence star with a spectral typeof O9, at a distance of 450 pc and the age of ∼3×104 years. Drew et al. (1997), however,predicted it to be a B1.5 zero-age main-sequence star at a distance of ∼ 250 pc. They also
argue that MWC 297 is a rapid rotator with v sini ∼ 350 km/s. The visual extinctionis accepted as 8−10 (McGregor et al. 1984; Hillenbrand et al. 1992; Porter et al. 1998)
though Berrilli et al. (1992) and Hou et al. (1997) measured lower AV (∼3). ROSAT had
an observation of this region with short exposure (∼1089 sec) but detected no x-ray source.
A pointlike x-ray source is detected near MWC 297. Within the typical ASCA position

54

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.22: ASCA images of the MWC 297 ﬁeld in the observation MWC 297 1 (left: SIS,
right: GIS). Images of the other observations (MWC 297 2,3) are almost the same as thisimage. Symbols are the same as the captions of Figure 5.1. The count rate per solid-line ishalf (0.2 cts/pixel/obs for SIS and 0.5 cts/pixel/obs for GIS).

error, two objects MWC 297 and SH 2-62 are found from the SIMBAD database at theposition of this x-ray source; SH 2-62 is a HII region surrounding the MWC 297 (Sharpless
1959; Chen et al. 1995). Figure 5.23 is the x-ray light curves in the full energy band. The ﬂux
in the ﬁrst observation is found to be constant by the χ2−test, but is increased by a factor
of ﬁve at the beginning of the second observation, then gradually decreases exponentially tonearly the same ﬂux of the pre-ﬂare level at the end of the third observation. This short time
variability indicates that the x-ray emitting region is compact, and hence is attributable to
Herbig Be star MWC 297 rather than an extended object SH2-62. Here and after, we regard
this source to be MWC 297. The time variability with an exponential-decay is similar tothat of a ﬂare from low mass stars. We accordingly regard that we observed a decay phase ofan x-ray ﬂare during the second and third observations, whereas the ﬁrst observation fell ona quiescent phase. The on-set of the ﬂare is between the ﬁrst and second observations. We
ﬁtted the light curve to a constant plus exponential decay model and obtained the best-ﬁt
e-folding time of ≈ 5.6 ×104 sec (see also Table 7.3).
The background subtracted SIS and GIS spectra for each observation are separately
displayed in Figure 5.24. Acceptable ﬁts with an absorbed 1T model are obtained for all
the spectra of the three observations. The column density of hydrogen is consistent with aconstant value during the three observations. The mean NH value of ≈2.4±0.9×1022 cm−2
is converted to the visual extinction of ∼10±4 (Cardelli et al. 1988; Ryter 1996), which isin excellent agreement with the observed value of AV = 8−10. The plasma temperature in
the quiescent phase of 2.7 keV is increased to 6.7 keV at the ﬂare and then decreases to
3.2 keV. This would be the ﬁrst detection of a ﬂare for early-type stars, accompanied with

5.1. DETECTED SOURCES

55

Figure 5.23: SIS light curves of MWC 297. The horizontal axis is longer than the light curvesof the other sources. See also Table 5.4.
a plasma heating in the rising phase, and cooling in the decaying phase. The ﬂare peak,
however, is missed; we only detected x-rays after the on-set of the ﬂare. Thus, the peakluminosity should be larger than 1.25×1032 ergs s−1 , which is the maximum value found at
the beginning of the second observation. Then, the total energy released in the x-ray bandis estimated to be more than 6.9×1036 ergs (0.4−10 keV).

5.1.10EC 95

EC 95 is a member of the SVS4 IR stellar cluster in the Serpens star forming core (Strom et al.
1976; Strom and Strom 1994), which is one of the densest clustering region known (Eiroa
and Casali 1989). Preibisch (1998) ﬁrst reported the constant x-ray emission (X3) from it
although they still leave the possibility of the other candidate EC 92. The very high visualextinction of both sources (AV ≈25−35 mag) made them suspect that the x-ray source hasa large x-ray luminosity of about 1033 ergs s−1 . Preibisch (1999) analyzed a near-infrared
spectroscopy with UKIRT and estimated that EC 95 is an extremely young intermediate
mass star (Age ∼2×105 yrs; M ∼4 M ), so called “proto-HAeBes”. Radio observations havedetected a point-like source around EC 95, which is named SH 68-2 (Rodriguez et al. 1980;
Snell and Bally 1986). Smith et al. (1999) conﬁrmed the source position of SH 68-2 within
the error circle of EC 95 with VLA. They also found a falling spectrum and a time variation
on a timescale of year, which suggests a gyrosynchrotron emission mechanism and hence a
presence of magnetosphere on or around EC 95. They searched for, but found no evidenceof, ﬂaring activity and circular polarization.
Figure 5.25 shows total band images around EC 95. A point-like source is detected
within the error circle of EC 95. In the SIMBAD database, there are 74 sources within 1.5

56

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.24: Spectra of MWC297 (left: SIS and right: GIS) from the ﬁrst(top), second(middle), and third(bottom) observations. See Figure 5.3 for details.

5.1. DETECTED SOURCES

57

Figure 5.25: ASCA images (left: SIS, right: GIS) of the EC 95 ﬁeld. See Figure 5.1 forsymbols.
centered on the EC 95 although they include clouds, nebulousities and so on. A ROSAT
HRI image with deep exposure of about 17 ksec (Preibisch 1998) displays only six sources
within 15 ×15 squares and the nearest x-ray source is 20 west from EC 95. We therefore
identiﬁed the ASCA source as EC 95. Flare-like small ﬂuctuation is clearly seen in SIS but
less prominent in GIS (Figure 5.26). However both light curves reject a constant model over
99.9% probability. Both SIS and GIS spectra (Figure 5.27) are consistent with each otherexcept for the model normalization, which is probably responsible for the CCD degradation(see the subsection 3.1.3). We thus vary each normalization parameter independently in thesimultaneous ﬁtting and derived its source ﬂux from GIS parameters. The NH value is oneof the largest among HAeBes but about one fourth of the converted value from AV ∼36.
The plasma temperature is very high (kT ∼ 3.9keV), and the luminosity log LX (ergs s−1 )∼31.7 is large, but it is about an order as small as the ROSAT estimation assuming AV ∼36
and kT ∼1−3keV. The count rate of ROSAT HRI in the ASCA observation is estimated at
2.8×10−3 cnts s−1 , which is about two times as large as the ROSAT result (Preibisch 1998).
The diﬀerence between ROSAT and ASCA is therefore not due to time variation but simply
due to the inconsistency of NH .

5.1.11S CrA

S CrA is a K6 star in the CTTS phase (Herbig and Bell 1988), but The et al. (1994) quoted
it at a potential candidate of a F-type PMS. It is thought to be extremely young, ∼7×104
yrs. The mass is smaller than the sun, ∼0.55 M (Olofsson et al. 1999). It has a companion
star 1.4 apart (eg Reipurth and Zinnecker 1993; Chelli et al. 1995; Prato and Simon 1997).
The x-ray emission was detected with ROSAT as relatively strong source log LX (ergs s−1 )58

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.26: light curves of EC 95 (top: SIS, bottom: GIS). See Figure 5.2 for details.

Figure 5.27: Spectra of EC 95 (left: SIS,right: GIS). See Figure 5.3 for details.

5.1. DETECTED SOURCES

59

Figure 5.28: ASCA GIS image of the S CrA ﬁelds on Oct 1998 (R CrA6). The bright sourceat the south east is CrA1. See Figure 5.1 for symbols.

∼30.1 (Carkner et al. 1998).
S CrA is on the ASCA fov in the 1st (SIS, GIS), 3rd(GIS), 4th(GIS) and 6th(GIS)
observations. It locates near, and suﬀers strong contamination of, the bright soft x-ray
source CrA1, but we ﬁnd the signiﬁcant x-ray emission at the position of S CrA (Figure5.28) in the 1st, 3rd and 6th observations. The ﬂux during the 1st observation is weak andtherefore only the SIS can manage to detect the x-ray emission. A ﬂare like variation is seen
during the 6th observation though the photon count is too small. The plasma temperatureis slightly high through the observations (kT ∼2 keV). The average ﬂux goes up in an order
between the 1st and 3rd observation, but it is constant after that.

5.1.12

TY CrA / HD 176386

TY CrA and HD 176386 belong to the R CrA cloud and illuminate the reﬂection nebula
NGC 6726/7 (Graham 1991). TY CrA is 1 .1 southwest northeast of HD 176386. The
parallaxes of HD 176386 measured with HIPPARCOS is consistent with the distance of the
R CrA cloud measured from associating early type stars. (Marraco and Rydgren 1981; deZeeuw et al. 1999). TY CrA is a detached Herbig Be eclipsing binary with an orbital periodof 2.88878 days (Kardopolov et al. 1981). The photometric variability associated with theeclipsing phenomena has been intensively studied by Casey et al. 1993,Casey et al. 1995 andCasey et al. 1998 3 . From the orbital motion and the eclipsing feature, the primary star isestimated at 3.16M located near the zero-age-main sequence stage and the secondary is at
1.64M near the base of the Hayashi track. Both stars are consistent with an age of 3 Myr.
3

The detailed photometry also suggests the presence of one more companions (Casey et al. 1998)

60

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.29: GIS light curves of S CrA in the 3rd and 6th observations. See Figure 5.2 fordetails.

Figure 5.30: Spectra of S CrA in the 1st, 3rd and 6th observations. See Figure 5.3 for details.

5.1. DETECTED SOURCES

61

Figure 5.31: ASCA images (left: SIS, right: GIS) of the TY CrA/HD 176386 ﬁelds in Oct
1998 (R CrA6). See Figure 5.1 for symbols.

Figure 5.32: SIS light curves of TY CrA(left) and HD 176386(right) in the 1st(top) and6th(bottom) observations. See Figure 5.2 for details.

62

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.33: SIS spectra of TY CrA(left) and HD 176386(right) in the 1st(top), 4th(middle)
and 6th(bottom) observations. See Figure 5.3 for details.

5.1. DETECTED SOURCES

63

Figure 5.34: Conﬁdence contour of the spectrum ﬁtting parameters of HD 176386 (1st observation)
Whereas the spectral lines of the primary star is extremely narrow, indicative of an almost
non-rotating star (v sini ∼10 km s−1 ). Spectral lines of the secondary is equivalent to v
sini ∼36 km s−1 , supporting the synchronous rotation (Casey et al. 1993; Lagrange et al.
1993). The time variations observed in the non-eclipsing phase are thought to originate in the
dust obscurity of the circumbinary materials (Vaz et al. 1998), but the lack of the infraredexcess indicates the absence of an optically thick disk (Casey et al. 1993). TY CrA is a nonthermal radio source (Skinner et al. 1993). Both Einstein and ROSAT detected the x-ray
emission from TY CrA. Einstein also measured the plasma temperature and the absorption
column density (kT ∼1.04 keV, NH ∼ 5.5×1021 cm−2, Damiani et al. 1994). On the other
hand, HD 176386 is a B9.5 star with the age of 2.8 Myrs (eg Grady et al. 1993, Bibo et al.
1992). HD 176386 is thought to be close to ZAMS from these observational results; constantbrightness in optical and near IR, extended mid-infrared emission indicating a presence oflarge inner cavity (Prusti et al. 1994), excess radiation only in far-IR (Bibo et al. 1992) andintermediate values of the accreting gas column densities and maximum accretion velocity(Grady et al. 1993).
ASCA detect the strong x-ray emission from TY CrA and HD 176386 and ﬁnd no
signiﬁcant variation between observations. We show images in the 6th observation (Figure5.31). Both sources are clearly seen in the SIS image but they are not in the GIS image.
Because SIS only covers TY CrA and HD 176386 in the 1st, 4th and 6th observations, weﬁrst discuss the SIS results in detail and then mention the GIS results.
Figure 5.32 displays SIS light curves of TY CrA and HD 176386. The light curveof the 4th observation is excluded because of the very short exposure. We do not ﬁnd any
signiﬁcant time variation and the constant model is accepted. All SIS spectra of both sourcesin Figure 5.33 are reproduced with an absorbed 1T thermal model. The best-ﬁt values ofthe spectrum in HD 176386 had two local minima (Figure 5.34). A local peak with smallerCHAPTER 5. RESULT I – HERBIG AE/BE STARS

64

Figure 5.35: GIS light curves of TY CrA + HD 176386 in the 2nd (top left), 3rd (bottom)
and 5th (top right) observations. The vertical scale of the bottom panel(3rd observation) islarger than the other panels. See Figure 5.2 for details.
NH is consistent with the AV 4 . Because of the limitation of photon counts, statistical errorsof the 4th spectra are quite large. NH values of both sources in the 1st and 6th observations
are about the same so that we also try the spectral ﬁttings of 4th spectrum with ﬁxing NH
value on the average of the 1st and 6th observations. Then the plasma temperature of TY
CrA and HD 176386 are about the same as those in the other observations.
Figure 5.35 shows the GIS light curves of all observations, normalized by an eﬀective
area. The ﬂux is almost constant except for the 3rd observation when the x-ray count rate
began to increase at 100 ksec and is maximum at 120 ksec. The x-ray peak in the GIS
4Cardelli and Wallerstein (1989) measured the RV value to the direction of TY CrA to be considerably
higher than normal and derived relatively large AV ∼ 3.1, compared with the Hillenbrand et al. (1992)’s
estimation (AV ∼1.0). In the derivation of the AV −NH relation (Ryter 1996), however, RV is assumed tobe the normal value, 3.1. We adopt the Hillenbrand et al. (1992)’s estimation.

5.1. DETECTED SOURCES

65

Figure 5.36: GIS spectra of TY CrA + HD 176386 in the 2nd(top left), 3rd in the low state(topright), 3rd in the ﬂare state(middle left), 4th (middle right) and 5th(bottom) observations.
See Figure 5.3 for details.

66

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.37: Images (top left: ASCA SIS, top right: ASCA GIS, bottom right: ROSAT HRI)of the HD 200775 ﬁeld. See Figure 5.1, 5.4 for symbols.

image do not change during the ﬂare state and thus the ﬂaring source is probably TY CrA.
Unfortunately, the data for about 20 ksec after the maximum is lost owing to the operational
failure (see Table 4.9) so that we do not know the variable feature after the maximum. At
150 ksec, the ﬂux decreases by half the maximum. All GIS spectra are reproduced with anabsorbed 1T thermal model (Figure 5.36), but they really consist of two spectral components
of TY CrA and HD 176386. The spectral parameters therefore do not reﬂect the plasma
temperature itself but they represent the hardness of the merged x-rays. TY CrA is far
luminous than HD 176386 above 1 keV. We thus think that GIS spectra is representative of
the TY CrA component. The plasma temperature is about 2 keV, but it is higher during
the ﬂare state, which is clearly seen by adding the GIS2 data.

5.1. DETECTED SOURCES

67

Figure 5.38: Light curves of HD 200775. See Figure 5.2 for details.

5.1.13

HD 200775

HD 200775 is a B2.5IVe star (see van den Ancker et al. 1998) illuminating the reﬂection
nebula NGC 7023, The distance to HD 200775 was precisely measured with HIPPARCOS
and determined at about 430 pc (van den Ancker et al. 1998; Benedettini et al. 1998).
HD 200775 is in a biconical cavity that would have been swept away by energetic bipolar
outﬂows, but there is no evidence of ongoing outﬂow activity (Fuente et al. 1998). From thecomparison of the position in the HR diagram to an evolutional model, van den Ancker et al.
(1997) estimated HD 200775 to be a very young star (∼2×104 year). They however suggestthat it could be in a post pre-main-sequence stage in 20M-year-old. It has a companion star
2.25 apart (∆K = 4.9 mag, Li et al. 1994; Pirzkal et al. 1997) and only eight surrounding
sources within ∼50 (0.21 pc at 430 pc, Testi et al. 1999). Einstein detected relatively strong
x-rays from HD 200775 and measured its plasma parameters at kT ∼0.8 keV, NH ∼5×1021
cm−2 (Damiani et al. 1994).
In both SIS and GIS images (Figure 5.37), there is a pointlike source around HD 200775.
The error circle includes one more optical source IRAS 21009+6758 to the north from HD
200775, but the ROSAT HRI image clearly shows the x-ray emission coming from HD 200775.
IRAS 21009+6758 is thought to be a dusty cloud heated by UV radiation from HD 200775(Whitcomb et al. 1981) and is not appropriate for an x-ray source.
The light curves (Figure 5.38) show no evidence of time variability. A ﬂare-like event isseen at about 9.3 ×104 sec in the SIS light curve, but this is probably instrumental origin
because it is only seen in SIS0. Figure 5.39 shows the spectra. In the spectral ﬁttings weonly use data bins above 0.73 keV because the level discrimination is set on a diﬀerent value
on each bit-rate data (see the subsection 4.3.3). We perform simultaneous ﬁtting of both
SIS and GIS spectra and ﬁrst try the ﬁtting with an absorbed 1T model, but it is rejected
over the 99.9% conﬁdence. The bad χ2 values seems to be originated in several bumpy

68

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.39: Spectra of HD 200775. Top panels show the 1T best-ﬁt model of SIS (left) andGIS (right). Bottom panels show the SIS spectra of a 1T model with free abundance (left)
and a 2T model (right). See Figure 5.3 for details.

5.1. DETECTED SOURCES

69

Figure 5.40: ASCA Images (left: SIS, right: GIS) of the MWC 1080 ﬁeld. Details of symbols
are the same as the captions of Figure 5.1,5.4, but the count rate per solid-line for the ASCA
image is half (0.2 cts/pixel/obs for SIS and 0.5 cts/pixel/obs for GIS).
features at 4.2 and 5−6 keV (SIS) and 1.9, 3.1, 4.2 and 5.3 keV (GIS). We retry a spectral
ﬁtting with free abundance, but the reduced χ2 value does not recover. The bumpy structure
may be originated in several emission lines emanating from a diﬀerent thermal component.
We added one more thermal component with the common NH (2T model in Table 5.5).
Then, the reduced χ2 value improves much though the ﬁtting result is still rejected in the
90% conﬁdence level. The NH value in this 2T model is largely consistent with the visualextinction (AV ∼1.92). We thus think that the excess of the χ2 value is due to systematic
errors. We also try the other 2T thermal models (i) with diﬀerent NH and (ii) setting theabundance parameter free. However these reduced χ2 values are worse − 1.187 for (i) and
1.190 for (ii).

5.1.14

MWC 1080

MWC 1080 is an embedded star in the dark cloud Lynds 1238. The distance is measured at
2.2–2.5 kpc from the UVBR photometry of neighboring MS stars (Grankin et al. 1992) andthe kinematic distance of the 13 CO molecular line (Canto et al. 1984). The spectral type isclassiﬁed as B0e by Cohen and Kuhi (1979) using the He I emission lines, but the latter study
using the other lines optimized for early-type stars indicated the spectral type to be early
A with AV ∼4.4 mag (Yoshida et al. 1992). MWC 1080 has a hierarchical multiple system.
The primary star is an eclipsing binary with an orbital period of 2.8869 days (Grankin et al.
1992). It has a sub-arcsec companion (0.7 ) which can also be a HAeBe, and another closecompanion 4.7 apart (Leinert et al. 1997; Pirzkal et al. 1997). The binary system itself issurrounded by a small group of near-infrared stars (Testi et al. 1997). The CO molecular

70

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.41: Light curves (left: SIS , right: GIS) of MWC 1080. See Figure 5.2 for details.

Figure 5.42: Spectra (left: SIS , right: GIS) of MWC 1080. See Figure 5.3 for details.

outﬂow were mapped by Canto et al. (1984) and Leinert et al. (1997). The blue- and redshift components are concentric and centered on the star. It is thought to be a sphericallysymmetric appearing outﬂow. The spectrum shows the strong P-Cygni proﬁle and the windspeed can be reached up to 1000 km s−1 (eg Yoshida et al. 1992). The spectrum also
indicates the presence of ionizing winds (Benedettini et al. 1998). The X-ray emission was
detected by Zinnecker and Preibisch (1994). The count rate detected by ROSAT was verysmall, but the x-ray luminosity is the largest among their samples log LX (ergs s−1 ) ∼32.
ASCA detect an x-ray source within the error circle of MWC 1080 as a relatively faintsource (Figure 5.40). A fainter peak at the north east of this source can be another x-ray
source. The source intensity (Figure 5.41) seems to oscillate periodically in a factor of four by
about 40 ksec, but the constant ﬁtting is also acceptable. The spectra could be reproduced

5.2. NON DETECTED SOURCES

71

with an absorbed 1T MeKaL model (Figure 5.42). The temperature is extremely high (kT>3.8 keV) and the absorption column density is modest (NH ∼1022 cm−2), which is consistent
with AV . A line-like excess is seen around 6.5 keV. When we put an additional componentof Gaussian line with σ ∼0, the center energy is determined at 6.38 keV, corresponding to
the neutral Fe. The total luminosity exceeds log LX (ergs s−1 ) ∼ 32 .

5.2

Non Detected Sources

5.2.1

BD+30◦549

BD+30◦ 549 is located in the NGC 1333 star forming region. In an ASCA GIS image,BD+30◦ 549 suﬀered strongly contamination from GSC 2342359. The contamination is lesssigniﬁcant in the SIS image, but the data quality is bad at the position of BD+30◦ 549 by
the telemetry saturation.

5.2.2

CoKu Tau 1

CoKu Tau 1 is classiﬁed as a potential HAeBe candidate in The et al. (1994). They did
not quote its spectral type, but Strom and Strom (1994) estimated the spectral type at M2eso that it might be rather a low mass YSO. It exhibits bipolar outﬂows and accompanies
a companion star (Padgett et al. 1999). ASCA and ROSAT PSPC detect no x-ray source
around CoKu Tau 1 (Figure 5.1).

5.2.3

AB Aurigae

AB Aurigae is the brightest Herbig Ae star in the northern hemisphere (A0e: Bohm andCatala 1993). It is relatively close to the sun (d ∼144 pc, van den Ancker et al. 1998) andtherefore have been well studied as a prototype of HAeBes. It has the photometric variabilityon the magnitude of one tenth, which could be caused by the changes of accretion rate orphotospheric extinction (Skrutskie et al. 1996). Spectral lines, some of which have P-Cygni
features, also exhibit high level variability with a periodic velocity modulation (Catala et al.
1999). It is surrounded by a circumstellar disk inclining of ∼76◦ (Mannings and Sargent
1997; Grady et al. 1999). There is no evidence of the presence of a companion star in theintensive binary searches (Leinert et al. 1997; Pirzkal et al. 1997; Testi et al. 1997). X-ray
emission was ﬁrst detected with ROSAT as a very faint source, log LX (ergs s−1 ) ∼29.5(Zinnecker and Preibisch 1994).
AB Aurigae was fortunately within the ASCA fov of the SU Aurigae pointing (Figure

72

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.43: ASCA SIS images of the AB Aur (left) and T Ori (right) ﬁelds. For the T
Ori ﬁeld, Star marks are T Ori, V 361 Ori and MR Ori from left. Meanings of symbols are
the same as Figure 5.1, 5.4, but lines above 20 solid-lines are shown with dotted lines by
2cts/pixel/obs. See Figure 5.1 for symbols.
5.43). SU Aurigae is ∼4 east from AB Aur and the brightest x-ray source in this region(Skinner and Walter 1998). The AB Aurigae region seems to have excess x-ray emission, but
it is less signiﬁcant merged with the emission of SU Aurigae – signal-to-noise ratio ∼1.5σ(0.5−10 keV), 2.0σ (0.5−2 keV).

5.2.4

MR Ori / V361 Ori / T Ori

V361 Ori is a B1.5 star in ZAMS or the β-Pic phase. It shows a large amplitude variable
and has absorption line features in Hα and Mg, and inverse P-Cygni features in CIV (Grady
et al. 1996). Corporon and Lagrange (1999) did not conclude its binarity. ROSAT detected
relatively strong x-rays, whose spectra suggest low-temperature plasma with a moderate
absorption (NH ∼1.24×1022 cm−2 , Gagné et al. 1995; Geier et al. 1995). MR Ori is an
A2 star with AV ∼1.63 (Gagné et al. 1995). ROSAT deep pointings determined the upper
limit log LX (ergs s−1 ) <29.93. T Ori is located at about 10 south-east from the Orion
Trapezium. Shevchenko and Vitrichenko (1994) noticed that T Ori would be an eclipsingbinary with the period of 14.3 days, but Corporon and Lagrange (1999) did not ﬁnd any
eclipsing phenomenon. There is a plausible companion star at 7 .7 from the primary star
(Hillenbrand 1995; Leinert et al. 1997). The LX upper limit was determined with ROSAT
log LX (ergs s−1 ) <29.89 (Gagné et al. 1995).
In the ASCA SIS image (August 1993, Figure 5.43), x-rays from the Orion Trapezium
strongly contaminate MR Ori and V361 Ori regions. On the other hand, we ﬁnd no x-ray

5.2. NON DETECTED SOURCES

73

Figure 5.44: ASCA GIS(left) and ROSAT (right) images of the BF Ori ﬁeld. See Figure 5.1
and 5.4 for symbols.

emission at the position around T Ori though there is small peak (S/N ∼4.3) at south-east
from T Ori, which we identify as 648 in Ali and Depoy (1995) from a ROSAT image.

5.2.5

BF Ori

BF Ori is an isolated Herbig Ae star (A5−6IIIe) and showing β-Pic phenomenon (van den
Ancker et al. 1998; de Winter et al. 1999). It is photometric variable with a large amplitude(Shevchenko et al. 1993). The distance is determined to be more than 210 pc with HIPPARCOS and estimated at ∼430 pc assuming a membership of the Orion OB1 c association(Warren and Hesser 1977; van den Ancker et al. 1998). A wing feature, inverse P-Cygni
proﬁles and absorption lines indicate the presence of infalling materials (Welty et al. 1992;
de Winter et al. 1999).
In the ASCA GIS image, we detect an x-ray source on BF Ori, but unsharp image on theedge of GIS allow another three candidates of the optical counterpart; [PC87b] 12, [VSS88] 3
and V865 Ori (error circle ∼ 1 ). In the ROSAT image, we ﬁnd an x-ray source at the positionof V865 Ori and not BF Ori. Pravdo and Angelini (1995) reported plausible x-ray detectionof BF Ori with the same data, but the coordinate of BF Ori in Pravdo and Angelini (1995) isabout 1 south west from the HIPPARCOS coordinate. We thus estimate the source upperlimit ﬂux with the ROSAT data (See Table 5.2). We check the ASCA spectrum because
BF Ori might be a hard x-ray source. However we ﬁnd no enhancement above the ROSAT
band so that we conclude that the x-ray emission only comes from V865 Ori.

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

74

Figure 5.45: Spectrum of an x-ray source in the BF Ori ﬁeld. See Figure 5.3 for details.

5.2.6

Z CMa / LkHα 218 / LkHα 220

Z CMa, LkHα 218 and LkHα 220 belong to the star forming region CMa R1-I (Shevchenko
et al. 1999). Z CMa is classiﬁed as a member of the FU Ori class of variables. It is a binary
with a separation of 0.1 , and both the primary and secondary are thought to be HAeBes(eg Leinert et al. 1997). It has a unique spectrum, dominated by its gaseous accretion disk(see intro. of Shevchenko et al. 1999) and showing a jet phenomenon. LkHα 218 showed thepolarimetry in U and V (variable) band (Yudin and Evans 1998). ROSAT detected x-ray
emission from Z CMa (Zinnecker and Preibisch 1994).
The x-ray emission is enhanced on a large region around Z CMa, LkHα 218 and LkHα220, probably because weak x-ray emission from clustering sources form diﬀuse like image.
A hard x-ray peak (2–10 keV) seems to be at Z CMa. Clear x-ray peaks are present around
LkHα 218 and LkHα 220 as well, but these peaks are obscured and have another candidates.
We therefore check the ROSAT data. The x-ray peak near LkHα 218 computed by the detectcommand (α, δ) = (07h 02m 43s .3, –11◦ 27 17 ) has better coincidence with the variable star
FZ CMa (sp. type: B2.5IV-V). On the contrary, there is no x-ray source around LkHα 220.
Though Damiani et al. (1994) say that Einstein detect the x-ray emission from LkHα 218
of log LX (ergs s−1 ) ∼31.59, the x-ray source is probably PPM 713454 about 20 south west
from LkHα 218, which is also detected in the ROSAT image.

5.2.7

HD 97048

HD 97048 is an embedded B9−A0 star located in the Chameleon I dark cloud and associatedwith the reﬂection nebula Ced 111. Figure 2 in Comerón et al. (1999) helps us to grasp

5.2. NON DETECTED SOURCES

75

the surrounding environs. HD 97048 shows remarkable spectrum variability suggesting theintermittent phases of mass loss (Irvine and Houk 1977). It is one of the few stars withunusual 3.43 and 3.52 µm emission features. HD 97048 itself has no evidence of a binary
star (Corporon and Lagrange 1999). Feigelson et al. (1993) surveyed the Chameleon I cloudwith ROSAT, but they do not detect x-rays from HD 97048 due to the short exposure (∼6ksec). Zinnecker and Preibisch (1994) ﬁrst detected its weak x-ray emission with a ROSAT
deep pointing (∼32 ksec).
ASCA observed the Chameleon I dark cloud two times. In each observation, we did
not ﬁnd clear evidence of the x-ray emission from HD 97048. In the GIS image of ChamI
IV (May 1994, top panel in Figure 5.46), there is an x-ray source close to HD 97048, but
we have another candidate HM16, which is a TTS surrounded by a dust shell (Feast andGlass 1973). In the SIS0 image, HD 97048 unfortunately drop on a chip gap. We, however,ﬁnd a faint peak at ∼30 north from HD 97048. This can be the peak of HM 16 or an
artiﬁcial peak made by the tail of HD 97048. We thus compared the SIS spectrum with the
GIS spectrum. Both spectra accept an absorbed 1T thin-thermal model (Figure 5.47). The
best-ﬁt parameters are exactly alike; large absorption column density (NH ∼4×1022 cm−2)and relatively high temperature (kT ∼2 keV). The SIS ﬂux is 75% of the GIS one but theerror bars also include the GIS value. If the SIS peak is artiﬁcially made by the tail of
HD 97048, the source region of SIS covers less than a quarter of the total eﬀective area andthe ﬂux should be less than a quarter of the GIS one. We therefore conclude that the x-rays
are not coming from HD 97048 but from probably HM16. Cambresy et al. (1998) estimatedthe AV of HM16 at 5.3, which corresponds to NH ∼ 1.1×1022 cm−2, almost a quarter of theobserved NH . The x-ray counterpart might be hidden another PMS nearby HM16. As forthe timing analysis, a constant model is acceptable for both light curves.
In the GIS image of ChamI V (bottom panel in Figure 5.46), we ﬁnd no bright source
around HD 97048. In the SIS image, we see a small excess halfway between HD 97048 andHM16 though it is on a tail of SZ19. This source shows no evidence of time variation. The
spectrum (Figure 5.47) is reproduced by an absorbed 1-T thermal model with relatively hightemperature and moderate NH , which is not so large as the NH of the source in ChamI IV
and almost the same as the converted NH of HM16, but larger than the converted NH of
HD 97048. It suggests that the x-ray emission mainly comes from HM16.
We do not ﬁnd the clear evidence of x-ray emission from HD 97048, but we ﬁnd the
x-ray emission from an embedded source probably associated with HM16. The properties
of HM16 is not well known, but Feast and Glass (1973) reports that the optical spectrum
resembles the Herbig Be star R Mon. Intensive observations are needed for HM16.

76

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Figure 5.46: ASCA images of the HD 97048 ﬁeld. Top panels are of ChamI IV (left: SIS,
right: GIS) and bottom panels are of ChamI V (left: SIS, right: GIS). Names of surrounding
sources are written in the upper right panel. Details of symbols are the same as the captionsof Figure 5.1, but the count rate per solid-line of ChamI IV is half (0.2 cts/pixel/obs for SIS
and 0.5 cts/pixel/obs for GIS).

5.2. NON DETECTED SOURCES

77

Figure 5.47: Spectra in ChamI IV(top left: SIS0, top right: GIS) and in ChamI V (bottom
left: SIS). See Figure 5.3 for details.

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

78

Figure 5.48: ASCA (left) and ROSAT HRI (right) images of the HD 147889 ﬁeld. Meaningsof symbols are the same as the captions of Figure 5.1, 5.4, but the count rate per solid-line
for the ASCA image is half (0.2 cts/pixel/obs).

5.2.8

HD 97300

HD 97300 is located in the Chameleon I dark cloud at d ∼188 pc (Bertout et al. 1999). It
is a B9 star near on the ZAMS in the HR diagram (Whittet et al. 1997). Feigelson et al.
(1993) detected the weak x-ray emission corresponding to log LX (ergs s−1 ) ∼29.0. ASCA
do not ﬁnd any apparent x-ray source at the position of HD 97300. It is near on a bright
x-ray source, which is just on a SIS chip gap. We at least conﬁrm that no hard source is
detected at HD 97300.

5.2.9

HD 147889

HD 147889 is located at a bright side of the Rho-Ophiucus dark cloud. van den Ancker
et al. (1997) determined the distance at ∼140 pc using the HIPPARCOS data. They quoted
that it might be a very young star (∼0.1 Myr). It has a double-lined feature in the optical
spectrum, which suggests a spectroscopic binary of two B2 stars orbiting every ∼5 days
(Haﬀner and Meyer 1995). The X-ray emission was not detected and only the upper limit
is determined assuming NH ∼0, log LX (ergs s−1 ) < 27.5 (Casanova et al. 1995).
ASCA detect an X-ray source at the position of HD 147889 (Figure 5.48). Because it
is on the edge of GIS fov where the position uncertainty is large, we check ROSAT images
in three observations from March 1991 to August 1995, which include a HRI deep pointingfor ∼50 ksec. In all observations, we do not ﬁnd any X-ray source in the HD 147889 region
but two X-ray sources, the A2 star 2E 1622.3-2420 located at 100 north-west and the Ae
star V852 Oph located at 100 south. The spectrum accepts an absorbed 1T model with

5.3. SUMMARY

79

Figure 5.49: ASCA spectrum of the HD 147889 region. See Figure 5.3 for details.
slightly absorbed cool plasma (NH ∼5×1021 cm−2 , kT ∼0.75 keV). Though a small residue isseen in the hard band, the ﬂux is dominant in the ROSAT band (0.8−2 keV). It is probably
originated in 2E 1622.3-2420 and V852 Oph.

5.3

Summary

With the ASCA satellite, 16 sources among 40 targets are detected above the 5 σ threshold(Table 5.1, 5.2). The detection rate (∼ 38%) is smaller than the detection rate of a ROSAT
survey (∼ 52%, Zinnecker and Preibisch 1994) because the detection upper-limit is high due
to source contamination from nearby bright sources (eg AB Aur, HD 97048 and Z CMa).
Nevertheless, a few embedded sources (eg MWC 297 and V921 Sco) are detected with ASCA
for the ﬁrst time, thanks to its hard x-ray imaging capability. Number of detected sourceson each spectral type is summarized in Figure 5.50. To see whether there is any bias eﬀectagainst distance, we make a plot of distant vs. x-ray luminosity (Figure 5.51). There is nobright samples near from the sun, but when we see sources with same distance, the x-ray
luminosity scatters in a wide range. We therefore do not think that samples are seriously
aﬀected by a bias eﬀect against distance.
8 sources among 16 detected sources exhibited time variation (50%, Table 5.4). Because the sensitivity of the temporal time variation strongly depends on source intensity, we
also see variability of the sources obtained above total 1000 photon counts through observations. Though the sample number decrease by eight, six sources are x-ray variable and thepercentage of variable sources increases up to ∼75%.
Every spectrum is successfully reproduced with an absorbed one or two temperature

80

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

thin-thermal (MeKaL) model (Table 5.5). The plasma temperature ranges between kT∼1−5 keV, which is signiﬁcantly higher than main-sequence OB stars (kT <1 keV).

5.3. SUMMARY81
Detected Source

25
20
15
10
5
0
Detection
Non Detection

Total

Above B

A

F

Below G

Unknown

16
24

8
11

3
6

0
2

3
3

2
2

Figure 5.50: Number of detected sources in the ASCA survey. Sources
with uncertain range in spectral type are grouped into the intermittent spectral type, but a few sources with wide uncertain range (egIRAS 12496-7650) are grouped into the unknown. Spectral type below G includes the HAeBe progenitor EC95, but the others might begenuine low mass stars.

Figure 5.51: Distance vs. x-ray luminosity relation. Circles show theLX values of the ASCA detected sources. Plus marks and arrows(upper-limit) show the values of the non-detected sources – LX values
are of the ROSAT estimates except for HD 76534, converted to the
ASCA band assuming kT = 2 keV.

Obs. ID

Detected Position
R.A. (2000)
Dec. (2000) Det.
◦


h m s
4 18 39.8
28 19 16S
5 36 23.8−6 43 18 G2
6 31 3.1
10 26 46
G
11 57 19.1 −79 21 30
S
12 0 6.2 −78 11 29
S
12 53 16.5 −77 7
6G
16 8 34.9 −39 6 12S
16 8 33.4 −39 5 25
SG
16 59 8.7 −42 42 30
G
18 27 37.9
−3 49 36
S

Net count
SIS
GIS
Total
cnts %
cnts %
cnts
V892 Tau
IC359
722 (59)– (–)
722
V380 Ori
L1641N– (–)
109 b(33)109
VY Mon (/G2)
PSRJ0631– (–)
338 (36)
338
T Cha
ChamIII
4248 (76) 2509 (74) 6757
HD104237
HD104237
2427 (84)916 (74) 3343
IRAS12496-7650 ChamII– (–)
100 (31)100
HR5999Lupus3143 (48)– (–)
143
HR6000
Lupus31420 (95)
– (–)
1420
Lupus3– (–)
3146 (78) 3146
2HR†
V921 Sco
Gal R 2– (–)
151 (38)
151
MWC297
MWC297 1
353 (62)257 (67)
610
MWC297 21023 (79)904 (87) 1927
MWC297 3
528 (64)349 (70)
877
EC95
Serp
18 29 58.3
1 12 43
S
3216 (75) 3679 (70) 6895
S CrA
R CrA 1
19 1 7.9 −36 57 11
S
122 (34)
– (–)
122
337
R CrA 3
– (–)
337 b(65)
R CrA 6– (–)
421 (71)421
TY CrAR CrA 1
19 1 40.8 −36 52 31S
524 (67)
– (–)
524
– (–)
75
R CrA 475a(56)
R CrA 6437 (71)– (–)
437
HD176386
R CrA 1
19 1 38.7 −36 53 18
S
227 (49)
– (–)
227
– (–)
45
R CrA 445a(52)
R CrA 6228 (49)
– (–)
228
TYHD†
R CrA 1
– (–)
1488 (78) 1488
R CrA 2– (–)
676 (64)676
R CrA 3
– (–)
1585 c(78) 1585
R CrA 4– (–)
510 (73)510
R CrA 5– (–)
543 (91)543
R CrA 6
– (–)
1029 (68) 1029
HD 200775
HD200775
21 1 34.9
68 10
1
S3361 (85) 1641 (69) 5002
MWC 1080
MWC1080
23 17 24.0
60 50 43S
273 (33)
185 (27)
458
Det.: Detector measuring source positions (S: SIS, G: GIS), Net count: background subtracted source
counts [Parentheses – source event ratio deﬁned as “source event / total event”, Total – Photon countsof SIS + GIS], † 2HR/TYHD: GIS data merged with HR 5999 (TY CrA) and HR 6000 (HD 176386),
a: s0 data, b: g2 data, c: g3 data

Objects

Table 5.1: ASCA detected sources

82
CHAPTER 5. RESULT I – HERBIG AE/BE STARS

Objects

Obs. ID

ASCA
ROSAT
L
Intensity Upper limit
Det.
log LX
Note
Ref.
Intensity
log LX
log L X
Note
bol
Counts
Flux
CountsFlux
BD 30◦ 549
SVS13
2.1 10−2
1.1 10−12
G
<31.2
con.
d
<29.8
< −4.9
CoKu Tau 1
IC359
4.4 10−2
1.9 10−12
S
<30.8
con.
f
<1.5 10−3<6.9 10−14<29.3< −4.3con.
AB AurSU Aur
3.1 10−2
7.4 10−13
S
<30.3
con.
a
29.5
−5.7
9.1 10−13
G
<31.4
con.
b
30.3
−
IX OriOri Trap1
2.5 10−2
V372 Ori
Ori Trap1
3.5 10−2
1.3 10−12
G
<31.5
b
30.3
−5.5
6.2 10−13
G
<31.2
b
30.7
−2.9
YZ OriOri Trap1
1.8 10−2
HD 36939
Ori Trap1
4.6 10−2
1.7 10−12
G
<31.6
b
<29.9
< −5.6
HD245185
Lamb Ori
1.6 10−2
3.3 10−13
S<30.8
f<2.4 10−3<9.7 10−14<30.3< −4.69.2 10−12
S
<32.4
con.
b
30.1
−6.6
LP OriOri Trap1
3.6 10−1
MR OriOri Trap1
2.7
9.4 10−11
S
<33.4
con.
b
<29.9
< −5.2
2.3 10−11
S
<32.8
con.
b
30.9
−5.4
V361 Ori
Ori Trap1
9.8 10−1
T Ori
Ori Trap1
6.5 10−2
1.6 10−12
S
<31.6
con.
b
29.9
−5.3
BF Ori
L1641N
—
5.8 10−13
G
<31.1
ano.
f
<2.2 10−3<6.7 10−14<30.2< −3.5
3.1 10−13G<31.0
f<2.7 10−3<3.3 10−14<30.0< −4.9
MWC120Ori OB1
9.5 10−3
V590 Mon
15 Mon4.8 10−2
1.2 10−12S<32.0
f4.2 10−3
2.4 10−13
31.3−4.8
4.4 10−12
G<32.8
f<3.7 10−2<1.2 10−12<32.3< −3.4con.
LkHα 218Z CMa
9.9 10−2
Z CMa
Z CMa
1.9 10−2
8.9 10−13
S
<32.1
a
31.1
−5.8
LkHα 220Z CMa
8.5 10−3
2.8 10−13†
G
<31.6
con.
f
<7.0 10−4<2.2 10−14<30.5< −5.4
HD76534
Vela Shrap 1.0 10−2
4.3 10−13
G
<31.6
n
–
–
S
<29.8
ano.
a
29.0
−6.2
HD97048
ChamI V—
1.8 10−13
HD97300
ChamI I
5.1 10−2
2.3 10−12
G
<31.0
con.
c
29.0
−6.1
G
<30.6
ano.
e
<27.5
< −9.4
HD147889
Rho-Oph
—1.7 10−12
Hen 3-1191
Gal R 1
2.3 10−2
7.5 10−13†
G
–
n
–
–
R CrA
R CrA 1
1.4 10−1
5.2 10−12
S
<31.0
con.
a
<28.7
< −7.0
1.5 10−12
S
<30.5
con.
a
<28.5
< −6.0T CrA
R CrA 1
4.0 10−2
ASCA data – Intensity upper limit [Counts : cnts s−1 , Flux: ergs cm−2 s−1 ]: 3σ level or count rate of detected events within selected regions
for contaminated sources (Counts), estimated from 20 half-box for SIS and 40 for GIS. Unabsorbed ﬂux upperlimit in the ASCA band (Flux,
0.5–10 keV) caluculated using the PIMMS package, assuming thin-thermal plasma of kT ∼1.3keV and NH converted from AV ., Log LX [ergs s−1]: unabsorbed luminosity in the ASCA band, Det.: Detector measured count rate (S: SIS, G: GIS)., Note: sources marked “con.” suﬀer x-ray
contamintion from nearby sources., sources marked “ano.” are identiﬁed as the other objects (BF Ori: V865 Ori, HD 97048: 11 h 7m 58s .5, −77◦39 13 , HD 147889: 2E 1622.3-2420 and V852 Oph)., † assuming NH = 0.0.
ROSAT data – Ref.: Reference of x-ray luminosity, [a: Zinnecker and Preibisch (1994), b: Gagné et al. (1995), c: Feigelson et al. (1993), d:
Preibisch (1997), e: Casanova et al. (1995), f: this work, n: no ROSAT pointing observation.], Intensity [Counts: cnts s−1 , Flux: ergs cm−2 s−1 ]:3σ upper limit or count rate of detected events using the PIMMS package (Counts), and unabsorbed ﬂux in the ROSAT band (Flux, 0.1–2.4 keV)
calculated using the PIMMS package, assuming thin-thermal plasma of kT ∼1.3keV and NH converted from AV ., Log LX [ergs s−1 ]: unabsorbed
luminosity in the ROSAT band (0.1–2.4 keV)., Note: “con.” suﬀer x-ray contamintion from nearby sources.

Table 5.2: ASCA non-detected sources

5.3. SUMMARY
83

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

84

Table 5.3: Selected regions for event extraction
Target
V892 Tau
V380 Ori
VY MonT Cha

Det. Source
BackgroundS
c3
sym V1023 Tau
sr includes V410 x-ray 7.
Gc3
a36
sf (sym optical axis)G
e 4.5 2.5
sf
S
c3
a36
G
c3
HD 104237
S
c 2.5
whole chip ex. sr
a36
G
c3
a36
IRAS 12496
G
c3
sym HR 6000
optimized to reduce HR 6000
HR 5999
S
e .66 .25
sf
ex. of the HR 5999 peak
HR 6000S
e 2 .66
sym GSC 07851-00426 †HR 5999+6000 G
c 2.5
V921 ScoG
c3
a36
a 3 4.4
MWC 297
S
c3
a36
G
c3
b93
EC 95
S
c3
a36
G
c3
sym HD 176386
TY CrAS
e .74 .53
sym TY CrA
HD 176386
S
e .64 .42
sym the protostars
sf br for 6th obs.†2 .
TY CrA+HD
G
c 2.5
sym CrA 1
S CrA
S
c 1.5
sf
HD 200775
S
c 2.5
sym the SE sourceGc3
sf
MWC 1080
S
c 2.5
a36
G
c3
sym SU Aur
AB AurS
c1
S
c 0.5
b 1.27 1.27
br the other tail of SZ19
HD 97048†3
ex. of the SZ19 peakG
e 1.25 1.63 symax SZ19 SZ23
†4
†5
S
c .67
c .67
small sr to reduce SZ19
HD 97048
e 3 7.5
sym DoAr21
HD 147889
G
c3
sf
Mon R2G
c5
c num: num arcmin radius circle centered on a target, e num1 num2: ellipse with num1 arcmin
long and num2 arcmin short axes. a num1 num2: annular circle with num1 arcmin inner andnum2 arcmin outer radius centered on a target, b num1 num2: num1 arcmin × num2 arcmin box
sym: symmetrical region with respect to sth, symax: symmetrical region with respect to the axis
between sth1 and sth2, nu: not used, sr: source region, br: background region, sf: source free
region† GSC 07851-00426: the x-ray luminous T-Tauri star located at the east from HR 6000, †2 to avoid
the ﬂare contamination of CrA 3 and V702 CrA, †3 CHAI III, †4 CHAIBROWN, †5 There is thebrown dwarf Hα3 at the another tail of SZ19 of the same CCD chip, whose x-ray emission ROSAT
detected by Neuhauser and Comeron (1998). The background is therefore selected from another
CCD chip (S0C2/S1C0).

5.3. SUMMARY

85

Table 5.4: Result of ASCA timing analyses
Objects

St.

V892 Tau
V380 Ori
VY Mon (/G2)
T Cha
HD104237
IRAS12496
HR5999HR60002HR
V921 Sco
MWC297

1
1
2
2
3
3

EC95
S CrA

TY CrA

HD176386

TYHD

1
3
6
1
4
6
1
4
6
1
2
3
4
5
6

HD200775
MWC1080

Det.

Bin

S
SH
G2
G
S
G
S
G
G
G
S
S
G
G
S
G
S
G
S
G
S
G
S
G
G
S
S
S
S
S
S
G
G
G
G
G
G
S
G
S
G

(sec)
2048
2048
2048
2048
2048
2048
2048
2048
2048
4096
2048
2048
2048
2048
2048
2048
2048
2048
2048
2048
2048
2048
2048
4096
4096
2048
2048
2048
2048
2048
2048
2048
2048
4096
2048
2048
2048
2048
2048
4096
4096

Constant Fittings
Mean
∆χ2
Var.
(cnts s−1 )1.5 10−2 1.42 (27)n
5.2 10−3 1.13 (27)n
2.4 10−3 0.94 (33)n
8.9 10−3 1.69 (70)yy
2.9 10−2 3.88 (69)
1.7 10−2 2.75 (74)y
4.0 10−2 1.38 (30)ny
1.6 10−2 1.58 (29)
1.6 10−3 1.00 (30)nn
1.7 10−3 1.22 (16)
3.4 10−3 1.08 (21)n
3.8 10−2 2.39 (21)yy
3.6 10−2 3.70 (39)n
6.3 10−3 0.51 (11)
1.7 10−2 0.81 (10)n
1.2 10−2 0.87 (10)n
5.0 10−2 1.84 (7)ny8.7 10−2 3.11 (6)1.5 10−2 1.18 (13)ny
1.0 10−2 5.68 (13)
1.7 10−2 2.13 (88)y
2.0 10−2 1.81 (93)yy
1.1 10−3 2.40 (35)
3.8 10−3 2.52 (33)yy5.2 10−3 2.11 (25)
6.4 10−3 0.79 (34)n
5.8 10−3 1.09 (12)nn
6.7 10−3 0.90 (39)
2.8 10−3 0.83 (33)nn
4.4 10−3 0.40 (12)2.1 10−3 1.21 (39)n
1.4 10−2 1.67 (36)yy
1.3 10−2 1.42 (37)
4.0 10−2 3.14 (33)yy
1.5 10−2 1.79 (14)1.8 10−2 0.84 (21)n
6.7 10−3 0.90 (39)nn
3.6 10−2 1.32 (45)
1.6 10−2 1.04 (47)nn
3.6 10−3 0.64 (19)
2.3 10−3 0.78 (18)n

X-ray sources identitiﬁed as other objects
G2048
1.8 10−3 1.30 (33)
n
S0
2048
1.9 10−3 0.71 (13)
n
n
G
2048
3.9 10−3 1.14 (18)S
4096
1.1 10−3 0.86 (47)nS
8192
1.1 10−3 0.78 (24)nnHD147889G
2048
4.5 10−3 1.44 (32)St.: Number of observation ID., Det.: detector (S: SIS, G: GIS, H hardband [2-10 keV])., Bin.: binning time scales of light curves., Mean:
average count rates derived from constant ﬁttings., ∆χ2 : parentheses
show the degree of freedom., Var.: sources showing time variation
above 90% conﬁdence level. 2HR/TYHD: GIS data merged with HR
5999 (TY CrA) and HR 6000 (HD 176386),
BF Ori
HD970482
2
3
3

CHAPTER 5. RESULT I – HERBIG AE/BE STARS

86

Table 5.5: Result of ASCA spectral analyses
Objects
V892 Tau

St.

V380 Ori
VY Mon (/G2)

LS
LS
HS
HS

T Cha
HD104237

IRAS12496
HR5999HR60002HR

LS
LS

Det. ML
S
1T
S
2T
G2
G2
G
G
G
G
SG
SG
SG
G
S
S
S
G
G

HS
HS

V921 Sco
MWC297EC95
S CrA

TY CrA

HD176386

TYHD

HD200775

G
G
G
1
SG
2
SG
3
SG
SG
1
S
3
G
6
G
1
S
4
S
6
S
1
S
1
S
4
S
6
S
6
S
1
G
2
G
3LS G3
3HS G3
4
G
5
G
6
G
SG
SG

MWC1080SGkT
2.0 (1.6–2.5)
2.0 (ﬁx)
0.5 (0.3–0.7)
1T
3.2 (0.7–11.0)
1Tf
2.4 (1.3–6.1)
1T
2.9 (2.0–4.6)
1Tg†2 2.2 (1.6–3.3)
1T
6.0 (3.8–12.3)
1Tls 20.3 (3.0–)
1T
2.0 (1.9–2.2)
1T
2.1 (—)2T
0.7 (0.6–0.7)
4.0 (3.3–4.8)
1T
2.0 (0.9–7.2)
1T
0.9 (0.6–1.7)
1.6 (1.2–2.4)
1Tf
1T
1.4 (1.2-1.4)
1Th
1.2 (0.9–1.3)
2Th
0.7 (0.5–0.9)
3.2 (2.4–5.0)
1T
1.8 (1.5–2.3)
1Th
4.0 (2.5–8.2)
1T
2.8 (1.1–9.1)
1T
3.4 (2.4–5.1)
1T
6.5 (4.7–9.6)
1T
3.4 (2.6–4.8)
1T
3.9 (3.5–4.4)
1T
2.9 (1.2–13.5)
1T
2.4 (1.4–4.0)
1T
4.7 (2.4–45.2)
1T
1.8 (0.8–2.3)
1.8 (1.2–3.4)
1Tf
1T
1.2 (1.0–1.4)
1T
1.3 (1.0–1.8)
1Tf
1.1 (0.7–1.4)
1Tf
1.0 (0.7–1.5)
1T
1.9 (1.3–2.8)
1Tf
1.6 (1.2–2.3)
1T
2.1 (1.9–2.4)
1T
1.7 (1.2–2.0)
1T
2.6 (2.1–3.1)
1T
3.0 (2.4–3.8)
1T
2.5 (2.0–3.3)
1T
2.1 (1.3–2.6)
1T
1.7 (1.4–1.9)
1T
1.7 (1.6–1.9)
2T
0.5 (0.4–0.6)
2.7 (1.9–4.1)
1T
10.0 (3.8–)

E.M.
NH
54.0 (53.9–54.1)
1.2 (0.9–1.6)
54.0 (ﬁx)1.8 (ﬁx)54.3 (53.9–55.4)
1.7 (1.2–2.8)
54.1 (53.9–54.3) †0 0.0 (0.0–1.5)
54.2 (54.0–54.4)
0.3 (ﬁx)
55.1 (55.0–55.3)
0.6 (0.2–1.0)
55.2 (55.0–55.4)
0.8 (0.4–1.2)
55.2 (55.1–55.3)
0.4 (0.1–0.8)
54.8 (54.7–55.2)
0.6 (0.0–2.5)
53.2 (53.2–53.3)
1.1 (1.0–1.3)
53.3 (—)
0.0 (—)
52.9 (52.8–53.1)
0.1 (0.0–0.2)
53.1 (53.0–53.1)
c54.1 (53.4–55.3)
13 (6.2–24)
53.9 (53.5–54.3)
0.7 (0.0–1.2)
53.5 (53.4–53.6)
0.1 (ﬁx)54.2 (54.2–54.3)
0.0 (0.0–0.1)
54.2 (54.2–54.4)
0.0 (0.0–0.3)
54.1 (54.0–54.3)
0.0 (0.0–0.2)
53.9 (53.7–54.1)
c54.5 (54.4–54.5)
0.0 (0.0–0.1)
54.0 (53.9–54.1)
0.0 (0.0–0.3)
55.0 (54.6–56.1)
4.9 (2.3–14)
54.2 (54.1–54.4)
2.4 (1.8–3.1)
54.9 (54.8–55.0)
2.5 (2.1–2.8)
54.2 (54.1–54.3)
1.9 (1.5–2.4)
54.5 (54.5–54.6) †3 2.3 (2.0–2.5)
52.5 (52.3–52.6) †0 0.6 (0.2–1.3)
53.4 (53.2–53.6)
0.2 (0.0–0.7)
53.1 (52.9–53.4)
1.0 (0.2–1.9)
53.4 (53.3–53.5)
0.3 (0.1–0.5)
53.5 (53.3–53.7)
0.4 (ﬁx)53.7 (53.5–53.8)
0.6 (0.3–0.9)
52.8 (52.7–52.9)
0.0 (0.0–0.2)
52.9 (52.8–53.0)
0.1 (ﬁx)53.0 (52.8–53.2)
0.0 (ﬁx)52.9 (52.8–53.9)
0.0 (0.0–0.2)
53.0 (52.9–53.1)
0.1 (ﬁx)53.6 (53.6–53.6)
0.0 (0.0–0.1)
53.6 (53.6–53.8)
0.0 (0.0–0.5)
53.7 (53.6–53.7)
0.0 (0.0–0.1)
53.8 (53.8–53.9)
0.0 (0.0–0.2)
53.6 (53.5–53.6)
0.0 (0.0–0.2)
53.7 (53.6–53.9)
0.0 (0.0–0.4)
53.5 (53.5–53.6)
0.0 (0.0–0.2)
54.5 (54.5–54.6)
0.1 (0.0–0.2)
55.2 (55.1–55.4)
0.9 (0.8–1.1)
54.3 (54.1–54.5)
c55.1 (55.0–55.3)
1.0 (0.3–2.2)∆χ2
0.99 (51)0.87 (54)1.31 (14)1.25 (15)1.40 (42)1.23 (40)
1.29 (45)1.28 (45)0.94 (183)
3.82 (89)1.57 (87)0.93 (14)1.15 (11)1.27 (12)1.53 (51)1.63 (62)0.74 (61)1.70 (37)0.96 (37)0.96 (18)0.92 (43)1.04 (105)0.89 (63)1.20 (205)1.68 (14)1.69 (22)1.37 (31)1.02 (28)1.33 (5)0.87 (25)0.90 (16)1.07 (17)0.56 (4)1.93 (13)1.91 (14)1.87 (26)0.66 (41)0.79 (65)0.88 (22)1.18 (30)0.98 (25)1.03 (64)1.39 (211)
1.18 (210)1.02 (56)Flux
LX
LX /Lbol
1.22
—
—
1.02
31.0
–4.4
0.20
31.3 †1
—
0.59
31.2
–4.4
0.68
31.2
–4.3
1.21
32.2
–5.4 †4
1.18
32.2–5.4 †4
2.49
32.4–4.1 †4
1.24
32.1–4.4 †4
1.24
30.2
–3.5
—
—
—
0.43
29.9
–5.2
0.93
30.2
–4.9
0.29
31.1
–4.2
0.37
30.9
–4.6
0.44
30.4
–5.1
2.01
31.1
–5.0
—
—
—
1.64
31.1
–5.0 †5
1.22
30.9–5.2 †5
4.14
31.5–4.6 †51.87
31.1–5.0 †5
1.14
32.1
–4.5
1.18
31.3
–6.8
9.23
32.1
–6.0
1.31
31.3
–6.7
1.97 †3 31.7
–3.7
0.11
29.5
–3.9
0.99
30.4
–3.1
0.59
30.3
–3.2
0.76
30.4
–5.1
0.89
30.5
–4.9
0.82
30.6
–4.8
0.32
29.8
–5.5
0.39
29.8
–5.4
0.50
30.0
–5.3
0.39
29.8
–5.4
0.48
29.9
–5.3
1.86
—
—
1.74
—
—
2.43
—
—
3.87
—
—
1.92
—
—
2.26
—
—
1.49
—
—
1.11
31.5
–6.0
0.59
32.1
–5.3
0.47
31.3
–6.2
0.26
32.3
–5.1

X-ray sources identitiﬁed as other objects
BF Ori
HD97048G 1T
0.8 (0.5–1.9)
54.5 (54.0–55.0)
0.0 (0.0–0.9) 0.85 (17)
0.45
—
—
S0 1T
1.4 (0.4–12.2)
54.2 (53.2–56.1)
4.2 (0.9–10)
1.14 (6)
0.42
—
—
G 1T
2.2 (1.2–5.2)
53.9 (53.5–54.5)
3.5 (1.5–7.0) 0.34 (12)
0.55
—
—
S
1T
2.8 (1.7–6.0)
53.1 (52.9–53.4)
0.7 (0.3–1.5) 0.85 (13)0.23——
HD147889
G 1T
0.8 (0.6–1.6)
53.8 (53.5–54.1)
0.5 (0.0–0.7) 1.13 (36)0.82——
St.: Number of observation ID & status (HS: high or ﬂare state, LS: low state), Det.: Detector (S: SIS, G: GIS, SG: SIS+GIS),
ML: Model [1T: 1-temperature, 2T: 2-temparature, h/g/ls : adding the components of HR5999/Gaussian/low-state, f : ﬁxing
NH value, i : two temperature model with independent NH , kT [keV], E.M. [log, cm−3 ], NH [1022 cm−2 ] Parentheses show
90% conﬁdence level, ∆ χ2 : Parentheses denote number of degree of freedom (d.o.f), Flux[10 −12 ergs cm−2 s−1 ]: observed
ﬂux in the ASCA band (0.5–10keV), LX [log, ergs s−1 ]: absorption corrected xray luminosity in the ASCA band, LX /Lbol[log], 2HR/TYHD: GIS data merged with HR 5999 (TY CrA) and HR 6000 (HD 176386), c: common NH values with theabove column, †0 NH is ﬁxed in the error estimate. †1 The distance of V410 X-ray 7, †2 Best-ﬁt parameters of the gaussiancomponent [center energy, 5.10keV (4.86-5.30), σ ∼0.0, ﬂux includes the gaussian component], †3 GIS value, †4 Lbol value of
VY Mon G2, †5 Lbol value of HR 6000,
2
2
3

Chapter 6
Result II – Giant Molecular Cloud
6.1

Monoceros R2 Giant Molecular Cloud

The Monoceros R2 giant molecular cloud (d = 950 pc) shows recent high mass star formation
activities, including a compact HII region (Downes et al. 1975), H2O and OH masers (Knappand Brown 1976) and a well-developed molecular outﬂow (Loren 1977). In the core region,
there are three luminous compact near-infrared objects with PMS characteristics (Beckwith
et al. 1976). Infrared spectra of IRS2 and IRS3 show deep absorption in the water-ice
band (eg Willner et al. 1982), which is a signature of young embedded sources such as the
Becklin-Neugebauer object. ROSAT observed the Monoceros R2 cloud and detected 41 Xray reliable sources (Gregorio-Hetem et al. 1998) but ROSAT found no X-ray source fromthe core.
Figure 6.1 shows ASCA images of the Monoceros R2 GMC core. Gray scale images show
the SIS medium band (left: 2−4 keV) and the SIS hard band (right: 4−10 keV). Though
the ASCA spatial resolution is limited, we can see at least three x-ray peaks in these images,
which correspond to IRS1 and source d (medium band) and IRS 2 (hard band). No peaks
have soft x-ray counterparts, suggesting them to be deeply embedded in the cloud core.
In hard band light curves, we found a hard X-ray ﬂare during the latter half of the observation (Figure 6.2). The light curves in the soft band show no variable feature, which meansthat the ﬂaring source is deeply embedded. The ﬂare shows a fast-rise and an exponentialdecay with the e-folding time of ∼1.8×104 sec. The peak count rate reaches almost threetimes as large as the quiescent level. The strongest source during the ﬂare is identical to thehard source during the quiescent state, ie probably IRS 2.
The left panel of Figure 6.3 shows a GIS spectrum during the quiescent state. This
spectrum has a line like feature at ∼6.5 keV and can be reproduced with an absorbed
bremsstrahlung model with a Gaussian line (Table 6.1). The plasma temperature is very87

88

CHAPTER 6. RESULT II – GIANT MOLECULAR CLOUD

Figure 6.1: ASCA images of the Monoceros R2 giant molecular cloud core (left: 2−4 keV,right: 4−10 keV). Massive YSO candidates IRS1, 2, 3 and source d are shown by star marks.
White contours represent the ROSAT PSPC ﬂux (soft band: 0.1–2.4keV).
high, kT ∼5.6±1.0 keV. Line center energy is 6.58±0.20 keV in the 90% conﬁdence level.
This error range leaves possibility of both neutral iron (6.4 keV) and Helium-like iron (6.7
keV) origins. The x-ray luminosity is typical of GMCs, LX (0.8−10 keV, ergs s−1 ) ∼32.6 (seealso Table 2.1). The right panel of Figure 6.3 shows a spectrum of the ﬂare component, whichtakes the quiescent spectrum as its background. The spectrum seems to have Fe, Ca andmarginal Ar lines and is successfully reproduced with an absorbed thermal bremsstrahlung
model with three Gaussian lines (Table 6.1). Line emissions from Ar and Ca materials has
never been detected in YSO spectra, and such a clear detection implys the richness of thesematerials. The ﬂare temperature is also very high, kT >5.3 keV. The absorption column
density is very large (NH ∼5.6×1022 cm−2), and almost equal to that of low mass protostars.
The ﬂare luminosity, log LX (0.8−10 keV, ergs s−1 ) ∼32.4 on average, is one of the largest
among previously detected YSO ﬂares.
The ﬂare spectrum shows a certain peculiar feature. The quantities of line center energy
for both Fe and Ca are signiﬁcantly lower than those calculated in a thin thermal plasma
model. We do not ﬁnd any systematic error in calibration data (55 Fe radioactive source),
which is obtained during observations. We tentatively ﬁt the ﬂare spectrum with an absorbed
Raymond-smith model with free redshift parameter, and then ﬁnd an acceptable result whenz = 0.072. One possibility is a ﬂare from an hidden AGN. Certainly, log LX (ergs s−1 ) ∼43.5
at z = 0.072 is typical of AGN ﬂares, but the ﬂare source is strongly associated with thecloud core and its chance probability is extremely low. We thus try another model, assuming
non-ionization equilibrium plasma (NEI). We adopt an absorbed masai89 model with freeabundances for Fe and Ca materials and one solar ﬁxed for the other materials. This model6.1. MONOCEROS R2 GIANT MOLECULAR CLOUD

Figure 6.2: Hard band light curves of the Monoceros R2 core (4–10 keV).

Figure 6.3: Quiescent and ﬂare state spectra of the Monoceros R2 core.

89

90

CHAPTER 6. RESULT II – GIANT MOLECULAR CLOUD

Table 6.1: Spectral parameters of the x-ray emission from the Monoceros R2 corequiescent state
ﬂare state
temperature kT[keV]
5.62 (4.67–6.95)(5.3–)
absorption
NH[1022 cm−2 ] 1.31 (1.13–1.50) 5.59 (2.9–14.2)
Fe lineEnergy [keV]
6.58 (6.37–6.72) 6.41 (6.21–6.64)
Ca lineEnergy [keV]
3.85 (3.76–3.93)
Ar lineEnergy [keV]2.90
2χ /d.o.f
1.06
0.78
Notes. Errors listed in parentheses are quoted for 90% conﬁdence level

well reproduce Fe and Ca line features when log nt = 10.09, but Ca is extremely overabundant against the solar value. Shu et al. (1996) suggest that calcium-aluminum-richinclusions (CAIs) condense near the X-point. We might see an x-ray ﬂare occurred in sucha region.

Chapter 7
Discussion I – X-ray Emission
Mechanism
This chapter discusses the relation of x-ray properties of HAeBes to stellar parameters andstudies the characteristic behavior on each individual HAeBe. The ﬁnal goal is to reveal the
x-ray emission mechanism.

7.1

Could the X-ray Emission Come from a Low Mass
Companion?

Large fraction of HAeBes is thought to be visible or spectroscopic binaries (P >85%, Pirzkal
et al. 1997). Though not a few companions are higher mass stars (eg MWC 1080), lowmass companions are generally much fainter than HAeBes in the optical wavelength and are
not easy to be detected. Therefore almost all HAeBes could possess a low mass companion.
Because the x-ray luminosity of a low mass star constitutes a large portion in the bolometricluminosity compared with a high mass MS star (low mass: log LX /Lbol <–3, High mass: log
LX /Lbol ∼ –6 – –8, see the section 2.2), a detected x-ray source cannot be easily identiﬁed
with a HAeBe itself even if no companions are found. Zinnecker and Preibisch (1994) havealready mentioned that the luminosity of unresolved T-Tauri companions, log LX (ergs s−1 )∼28–30, cannot explain the ROSAT samples, but we discuss the possibility in detail.
We estimate the low mass contribution of x-ray emission if a HAeBe has a low masscompanion. Binary formation theories say that the primary and secondary stars are simultaneously form in a molecular cloud, which suggests that low mass companions of HAeBes
would be in protostar or TTS phases (Age ∼105 −107 yrs). This picture is supported by theobservational evidence of the eclipsing binary TY CrA by Casey et al. (1998), in which they
estimated both of primary and secondary at ∼3M-year-old. We then refer to the x-ray lumi91

CHAPTER 7. DISCUSSION I – X-RAY EMISSION MECHANISM

92

nosity function of optically selected low mass PMSs. Neuhaeuser et al. (1995) performed acomplete survey of TTSs in the Taurus Aurigae region using the RASS data. They investigated the x-ray luminosities of optically selected TTSs and provided luminosity functions ofboth CTTSs and WTTSs. In their survey, they found signiﬁcant diﬀerence of the functions
between CTTSs and WTTSs, but deep pointings of the Chameleon I cloud with ROSAT(Feigelson et al. 1993) and the ρ-Ophiuchi with Chandra (Imanishi et al. 2001) indicatedthat the x-ray luminosity does not evolve from protostars to WTTSs. This discrepancy
might be due to the limited detection eﬃciency for CTTSs with relatively large absorption.
We thus adopt the luminosity function of WTTSs in Figure 2 of Neuhaeuser et al. (1995)1 .
We measured 90% of optically selected WTTSs below LX ∼1030 ergs s−1 (0.1−2.4 keV). For
the typical WTTS temperature (kT ∼0.5−2 keV), the x-ray ﬂux in the ASCA band is almost
as same as that in the ROSAT band (see Table 4.8). The detection rate of our survey is ∼40% and almost all the detected samples exceed LX ∼1030 ergs s−1 . We again cannot explain
the HAeBe luminosity with a T-Tauri low mass companion.
Nevertheless we should cope with a complex problem for ASCA results, especially on
the GIS data. A clustering survey of HAeBes by Testi et al. (1999) indicated that early Bestars generally associated several low mass stars, while late Be and Ae stars did not. ASCA
psf would cover many low mass stars around a HAeBe. We therefore evaluate possible ﬂux
contribution of these members. Testi et al. (1998) estimated that HAeBes accompanied atmost 30 TTSs within 0.21 pc, corresponding to from 17 (d ∼2.5 kpc) to 270 (d ∼160 pc).
There could be ∼10 low mass stars within the typical ASCA resolution (r ∼20 ). From
Neuhaeuser et al. (1995), we estimated the average LX of WTTSs to be log LX (ergs s−1 )∼29.52 , which is largely as same as the average LX value in the Chameleon I cloud log LX(ergs s−1 ) ∼29.2 (Feigelson et al. 1993). The estimated luminosities from low mass clusters
are up to log LX (ergs s−1 ) <30.5. This is less than 10% of the x-ray luminosity of HBes.

7.2

X-ray Emission Mechanism

We ﬁnd that x-ray emission from a HAeBe probably comes from HAeBes itself. We thendiscuss four plausible x-ray emission mechanisms on stars. This procedure traces the method
by Skinner and Yamauchi (1996).

7.2.1

Stellar Wind

The stellar wind model is proposed by Lucy and White (1980) and Lucy (1982) to explain theobserved x-ray features of O stars (see the section 2.2.1). This model assumes a complicated1
2

The graph shows the value in the Kaplan-Meier estimator. The detection rate of WTTS is 37%
LX ∼ 1029.7 ergs s−1 at Kaplan-Meier estimator 0.5, and the WTTS detection rate ∼66%

7.2. X-RAY EMISSION MECHANISM

93

situation; the presence of successive shocks in the wind. Here we simplify this problem as
whether the observed stellar wind has enough kinetic energy to generate the x-ray emissionof HAeBes. We investigate this possibility in two ways; i). LX vs. total kinetic energy instellar winds, ii). thermal energy vs. kinetic energy in a unit volume.
2.
For the estimate of i), we adopt the wind kinetic luminosity deﬁned as Lkin =1/2Ṁ Vterm
Figure 7.1 shows Lkin vs. LX . All LX are far below Lkin so that the wind can constantly
produce the observed luminosity.

For the estimate of ii), we made a graph of the wind velocity vs. the plasma temperature(Figure 7.2). we refer the wind velocity to Nisini et al. (1995) and Blondel et al. (1993).
All observed kT values are far larger than maximum temperatures produced from the windvelocity. The wind velocity has to exceed v ∼ 1000 km s−1 to heat gases up to ∼ 2 keV. Sucha high velocity wind is only seen in O stars. The stellar wind, therefore, cannot produce theobserved HAeBe temperature.

7.2.2

Accretion ShockWe estimate the maximum plasma temperature heated up by the mass accretion process,which is the origin of the UV excess emission from YSOs (eg Gullbring et al. 2000). We
assume that accreting gases have free-fall velocity on stellar surface. Then the equation is
as follows,
3nkT =

GM∗ nmp
3R∗

where n: electron (proton) density, mp : proton mass, M∗: stellar mass, R∗ : stellar radius.
The derived temperature kT ∼0.5−1 keV does not explain the typical temperature of HAeBes(kT ∼2 keV). Moreover the maximum temperature should be smaller than the derived value
because gases fall from the inner edge of a circumstellar disk and the gas velocity should be
slower than the free-fall velocity. The accretion shocks therefore cannot explain the plasma
temperature of HAeBes.
Blondel et al. (1993) reported that Lyman α emission from a few HAeBes is red-shiftedby v ∼600−900 km s−1 . They suggested accretion velocity can be up to v ∼ 800 − 2000
km s−1 , assuming the accreting geometry(see also the subsection 2.1.2). This is much fasterthan the typical free-fall velocity and these accreting gases can be heated up by kT ∼7 keV
when v ∼2000 km s−1 and can explain the temperature of HAeBes. If this accreting velocity
is real, a certain gas accelerating mechanism is needed.

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CHAPTER 7. DISCUSSION I – X-RAY EMISSION MECHANISM

Figure 7.1: Wind kinetic luminosity and x-ray luminosity relation.
Square plots mean high temperature components in the 2T models.
Lkin of HD 104237 and HR 5999 are intermittent values.

Figure 7.2: Terminal velocity (v∞ ) vs. plasma temperature relation.
Square plots mean high temperature components in the 2T models.
The line denotes the plasma temperature when whole kinetic energy
is converted to the thermal energy (3kT= 12 mpv2∞ , mp: proton mass).

7.3. X-RAY PROPERTIES RELATED TO THE MAGNETIC ACTIVITY

7.2.3

95

Wind-fed MagnetosphereA model by Usov and Melrose (1992) assumes successive occurrence of magnetic reconnection on the equatorial plane by dragging the magnetic ﬁeld out with stellar winds (see thesubsection 2.2.1 and Figure 2.8). This model does not assume the presence of a circumstellar disk and therefore is not appropriate for the structure of HAeBes. An x-ray emission
model, matching the circumstellar structure of HAeBes, is developed by Babel and Montmerle (1997). This model suggests that wind-shock in the magnetosphere can eﬃciently heatup gases. Because simulated stellar parameters are optimized in the magnetic Ap star, IQ
Aur, the wind velocity is faster than, and the mass loss rate is much smaller than, those of
HAeBes. In this simulation, the plasma temperature mainly depends on the wind velocity
and the x-ray luminosities depends on the mass-loss rate. HAeBes satisfy the latter condition, but plasma temperature produced by the HAeBes wind (v <500 km s−1 ) is below kT∼1 keV, which is lower than the observed HAeBe temperature. This model does not seemto account for the x-ray emission from HAeBes as well.

7.2.4

Magnetic Activity on the Stellar Surface

Magnetic activity on or near stars seems to be the most plausible mechanism to explain the
HAeBe x-ray emission. High percentage of our observed sources were x-ray variable (see thesubsection 5.3) and a few sources, MWC 297, HR 6000, VY Mon and TY CrA, show ﬂare liketime variations, similar to TTS ﬂares. The presence of an azimuthal magnetic ﬁeld structure
on a HAe is also suggested by time covariation of emission lines with stellar rotational period
in the optical wavelength (AB Aur: MgII 2795Å, Praderie et al. 1986; HeI 5876Å, Boehm
et al. 1996). In the next section, we discuss the magnetic activity on HAeBes in details.

7.3

X-ray Properties Related to the Magnetic Activity

After this section, we add HAeBe samples observed with ROSAT (Zinnecker and Preibisch1994). We also add the x-ray emitting PMS candidate SSV 63E+W discovered with ASCA
(Ozawa et al. 1999). The x-ray source SSV 63E+W would be either the Class I protostar
candidates SSV 63E or SSV 63W, which are located in the HH 24-26 region in the Orioncloud. Zealey et al. (1992) mention the evolutional age to be less than 106 yrs, estimatingthat SSV 63E to be in the pre-main-sequence stage and the mass 1.5M <M<5M from its
reddening factor.

CHAPTER 7. DISCUSSION I – X-RAY EMISSION MECHANISM

96

Table 7.1: Physical Parameters of Herbig Ae/Be stars
R∗
Mass
Wind Vel.
log Ṁ
log Lkin
v siniW(Hα) F1.3mm
(R )
(M )
(km s−1 )(M yr−1 ) (ergs s−1 ) (km s−1 )
(Å)
(mJy)
BD+30 549
0.8
V892 Tau
4.0
250
20.6
490
CoKu Tau 1
2.7
AB Aur
2.5
2.4
260
−8.0
32.5
140
26.3
430
IX Ori–
15
V372 Ori
6.0
3.3
125
2.5
YZ Ori
1.7
HD 36939
2.4
275
1HD245185
1.9
150
LP Ori
3.1
3.6
100
0.5
MR Ori
3.9
V361 Ori
3.6
50
T Ori
3.0
−8.0
32.8
100
130
V380 Ori
3.8
3.3
260
−7.9
32.7
200
71.0
<20
BF Ori
0.6
2.0
−8.8
32.1
100
3.7
<70
MWC1202.0
120<12VY Mon7.0
50
VY Mon G2
36.3
21V590 Mon
4.2
−8.5
32.3
47
LkHa218
3.2
3.8
31.9
Z CMa
40.0
<130
24.8
489
LkHa2202.1
4.0
HD76534
6.2
110
14.3
HD97048
2.1
2.5
140
51.1
451
HD97300
1.7
2.5
<91
T Cha
1.1
1.1
48
HD104237
2.7
2.3
500
−8.9, −6.7
32.3, 34.5
66
IRAS12496
3.5
522
HR59994.9
3.2
100−8.4, −6.8
31.4−33.0
180<40†HR60003.0<40†HD147889
3.1
7.5
Hen 3−1191–
V921 Sco
5.6
500
−5.5
35.7
270
MWC297
9.9
10
380
350
102.6
110
EC95
12.7
4
S CrA
1.6
0.6
328
HD176386
2.0
2.3
TY CrA
2.0
3.2
10
2
130
R CrA
2.9
3.0
−8.0
T CrA
1.8
1.5
−8.6
40
HD200775
7.0
10
280
−6.6
34.1
60
17.7
<20
MWC1080
35.1
400
−5.3
35.7
200
94.2
Reference: Radius– Caliculated from Teﬀ and Lbol (see the subsection 7.3.2). Mass– Hillenbrand et al.
(1992), Bibo et al. (1992), Boehm and Catala (1995), Casey et al. (1998), Drew et al. (1997), van den
Ancker et al. (1997), van den Ancker et al. (1998), de Winter et al. (1999), Shevchenko et al. (1999),
Preibisch (1999), Olofsson et al. (1999). Wind Vel.– Nisini et al. (1995), Benedettini et al. (1998). Log
Ṁ, log Lkin – Damiani et al. (1994), Benedettini et al. (1998). v sini– Grady et al. (1993), Boehm
and Catala (1995), Drew et al. (1997), Casey et al. (1998).van den Ancker et al. (1998), Corporon and
Lagrange (1999). W(Hα)– Damiani et al. (1994), Reipurth et al. (1996), Corcoran and Ray (1998).
F1.3mm – Wendker (1995) (22,500 or 23,000MHz). † HR5999 and HR6000 are not resolved.
Targets

7.3. X-RAY PROPERTIES RELATED TO THE MAGNETIC ACTIVITY

97

Table 7.2: Other HAeBes
Teﬀ
Lbol
LX
LX /LbolKLergs s−1
LkHα 101
−
4.1
<31.4
<−6.3
LkHα 349
4.15
3.1
30.9
−5.8
AFGL 4029
−
4.3
<31.8
<−6.2
LkHα 198
4.28
2.1
<30.5
<−5.2
V376 Cas
3.86
2.3
<30.0
<−5.9
2.1
31.2
−4.5
BD +46◦ 3471 3.92
LkHα 257
−
−
<30.4
—
GU CMa
4.32
3.6
<30.4
<−6.8
LkHα 234
4.14
2.7
<30.7
<−5.6
ω Ori
4.28
4.3
<29.6
<−8.3
LkHα 215
4.11
2.7
31.1
−5.2
All values are in logarithmic scales and referred toZinnecker and Preibisch (1994). Teﬀ is the typicalvalues of class V stars.
Objects

7.3.1

LX vs. Lbol Relation

Figure 7.3 shows the LX vs. Lbol relation. About half sources are in the range of −4< log
LX /Lbol <−6. This is an intermediate range between massive MSs and the saturation levelof low mass stars. We also plotted the Teﬀ vs. LX /Lbol relation in Figure 7.4. Here weomit sources whose Teﬀ values are not known. Sources near the saturation level of low massstellar x-rays (log LX /Lbol < −3) are concentrated in the region of log Teﬀ (K) <3.8, whichcorresponds to later than K type stars. One of them S CrA is really a low mass star, but
the other three are the promising HAeBes EC95, T Cha and SSV63E+W. This fact implys
that the x-ray emission mechanism during the very young phase is similar to low mass stars.
The LX /Lbol values gradually decrease as stellar eﬀective temperature increases.

7.3.2

Standard Dynamo Mechanism

The x-ray luminosities of low mass stars have a clear relation to the projected rotational
velocity (vrot sin i), that is the characteristics of the standard dynamo activity (see thesection 2.2.1). The same relation is expected for HAeBes as well if the standard dynamo
works on HAeBes on contrary to star forming theories. Figure 7.5, 7.6 show the dependenceof log LX and log LX /Lbol values on vrot sini. We cannot see any clear correlation between
these quantities. The LX /Lbol values saturate at the level of log LX /Lbol ∼ −3 above vrot
sini > 15 km/s in low mass MSs (Stauﬀer et al. 1994). All our samples rotate faster than
this boundary, but the quantities of log LX /Lbol scatter in the range between −4 and −6.
There seems to be no saturation level.
The standard (solar) dynamo, is believed to be a type of αΩ dynamo, which couples the
α eﬀect and diﬀerential rotation Ω. The LX and LX /Lbol values should be really dependenton Ω and not on vrot . We do not know the exact value of diﬀerential rotation Ω, but Ω could

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CHAPTER 7. DISCUSSION I – X-RAY EMISSION MECHANISM

Figure 7.3: LX vs. Lbol relation. Circles represent the ASCA detectedsources. Plus marks and arrows (upper-limit) represent the ASCA
non-detected sources, whose LX values are from the ROSAT estimates
except for HD 76534, and the ROSAT survey sources (Zinnecker andPreibisch 1994). LX values are converted to the ASCA band assuming
kT = 2 keV. Solid lines show log LX /Lbol = −3, −5 and −7 from thetop, respectively

Figure 7.4: Teﬀ vs. LX relation. Marks mean the same as the above.

7.3. X-RAY PROPERTIES RELATED TO THE MAGNETIC ACTIVITY

Figure 7.5: v sin i vs. LX relation. Marks mean the same as Figure
7.3.

Figure 7.6: v sin i vs. log LX /Lbol relation. Marks mean the same as
Figure 7.3.

99

100

CHAPTER 7. DISCUSSION I – X-RAY EMISSION MECHANISM

Figure 7.7: Ω sin i vs. log LX relation. Marks mean the same as
Figure 7.3.

Figure 7.8: Ω sin i vs. log LX /Lbol relation. Marks mean the same as
Figure 7.3.

7.3. X-RAY PROPERTIES RELATED TO THE MAGNETIC ACTIVITY

101

Table 7.3: Fitting Results of Flare Light Curves
τris
τdec
(104 s)(104 s)
MWC 297
MWC297 2-3 SG–
5.6 (4.5–7.0)
HR 6000
Lupus3
G
< 0.2
< 0.2
VY MonPSRJ0631G–
3.6 (2.0–7.4)
TY CrAR CrA3G∼1
2.4 (1.3–8.0)
Monoceros core Mon R2
G
< 0.8
1.8
SSV 63E+WaG
2.0 (1.6–2.6)
1.8 (1.6–1.9)
1.7
1.2 (0.9–1.5)
Det.: Detector(S:SIS, G:GIS, SG: summed lightcurve of both detectors), Iqui : count rate of a constantlevel, Ipeak : count rate of a ﬂare peak, τris (and τdec of HR 6000): a period between the constant level tothe ﬂux peak, τdec : e-folding time of a ﬂare decay, a: Ozawa et al. (1999).
SourceObs. ID

Det.

Iqui(10−2 cnts s−1 )2.3 (2.1–2.4)
3.3 (3.1–3.5)
0.25 (0.17–0.29)
3.6 (3.4–3.8)

Ipeak(10−2 cnts s−1 )>12
5.8 (5.3–6.2)> 0.38∼ 7.5

have a relation to the angular velocity ω. The stellar radius of low mass stars do not diﬀerent
so much, and threrefore the relation to vrot sini is a good approximation of the relation toω sini. However the stellar radius of HAeBes is diﬀerent up to an order (see Table 7.1) andthe relations could be much diﬀerent. We therefore estimate ω from the following equationand see the relation to LX and LX /Lbol .
4
Lbol = 4πR2 σTeffv = Rω
where R is a stellar radius and σ is the Stefan-Boltzmann constant. The result is shown in
Figure 7.7 and 7.8. We again do not ﬁnd any correlation between these parameters.

7.3.3

Prominent X-ray Variability on HAeBes

The quiescent x-ray luminosity does not show any characteristics of the standard dynamo
activity. We thus investigate x-ray variable features, which most indicates the magnetic
activity on HAeBes. Among our HAeBe samples, MWC 297, HR 6000, VY Mon and TY
CrA exhibited prominent time variations. All light curves but HR 6000 have similar time
proﬁles to low mass stellar ﬂares, though the rising phases of MWC 297 and VY Mon werelost. We ﬁt these light curves with a conventional model of constant and exponential decay.
The result is shown in Table 7.3. Most sources exhibited long decay timescales (a few ×105
sec). This is much longer than the typical decay timescales of low mass MSs and TTS ﬂares
but comparable to those of protostars (2−7 ksec for MSs and TTSs and 10−30 ksec for
protostars, Tsuboi et al. 2000; Stelzer et al. 2000).
Assuming the radiative cooling model (the quasi-static cooling model), the plasma volumes of low mass stellar ﬂares are extremely large, sometimes far exceed the stellar size.
Though another model (see the appendix A.1) estimates smaller plasma volumes than the

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CHAPTER 7. DISCUSSION I – X-RAY EMISSION MECHANISM

Figure 7.9: Variation of emission measure during ﬂares. Dots, squares
and triangles are the quiescent, ﬂare and ﬂare decay state, respectively.

Figure 7.10: Variation of plasma temperature during ﬂares. Symbols
mean the same as the above.

7.3. X-RAY PROPERTIES RELATED TO THE MAGNETIC ACTIVITY

103

Table 7.4: Physical parameters of the MWC 297 ﬂarene
4.3 1010 cm−3
V
5.0 1033 cm3L (= V1/3) 1.7 1011 cm
8.6 1011 cm
Lα=0.1
α = d/L. See Figure 7.11.

L/2

dFigure 7.11: Assumed geometry of the ﬂare magnetic ﬁeld
quasi-static model in some cases, the quasi-static model is convenient for the order estimateof plasma volume. We thus apply this model to the x-ray ﬂare on MWC 297. In this model,
ﬂare cooling timescales and luminosity is given byτcool =

3ne kB T
n2e Λ(T )LX = n2eΛ(T )V
where τcool : the cooling time scale of the plasma, ne : electron density, Λ(T ): the emissivity
in a unit volume. The derived parameters are shown in Table 7.4. The ﬂare size L >1011−12
cm is signiﬁcantly larger than those of solar radius and the typical ﬂare size of low mass
protostars. Because the radius of MWC 297 is about 10 times as large as the solar radius,a ﬂare loop can be conﬁned in just above the stellar surface of MWC 297. However, weonly saw the decay phase of the MWC 297 ﬂare. The maximum luminosity and plasma
temperature could be much higher than the observed value. It is, therefore, too early toconclude that a magnetic loop size on MWC 297 ﬂare is smaller than the stellar size.
We also investigated the kT and E.M. variation in quiescent, ﬂare and ﬂare decay phases(Figure 7.9, 7.10). kT values of MWC 297 and VY Mon/G2 rise and E.M. of MWC 297
increases during the ﬂares. This supports that some heating mechanism like solar-type ﬂares
is working during the ﬂare state. Other sources, however, do not show the prominent time
variation within the limited photon statistics.

CHAPTER 7. DISCUSSION I – X-RAY EMISSION MECHANISM

log(Emission Measure [cm3 ])

104

58
56
54
52
50
48
46
44
106

107
108
Temperature [K]Figure 7.12: kT vs. E.M. relation (black), superimposed on the plots of solar andstellar x-rays (grey). Parameters in the quiescent state are shown by dots (1T or
2T low-temperature) or squares (2T high-temperature), and those in the ﬂare stateor high state are shown by triangles. Grey asterisks, pluses and diamonds are plotsof a few stellar ﬂares. Oblique hatched region (upper right) are solar ﬂares andhorizontal hatched regions are of solar micro ﬂares (bottom right) and solar corona(bottom left). The solid line is the E.M.-T relation curve of stellar ﬂares for B =50
gauss. The dotted bar line is the relation of stellar corona. See details in Shibata
and Yokoyama (2001) and Shibata and Yokoyama (1999).

7.3.4

kT vs. E.M. Relation

Shibata and Yokoyama (1999) derived the universal correlation between ﬂare temperatureand emission measure for solar and stellar ﬂares, which is based on a magnetic reconnection
model with heat conduction and chromospheric evaporation. They extended this relation forcoronal emission from a basic equation of heat generation, conduction and radiation (Shibata
and Yokoyama 2001). In this theory, the kT vs. E.M. plots of quiescent and ﬂare states are ona single line when strength of the magnetic ﬁeld is given. We plot the observed kT and E.M.
values on this graph (Figure 7.12). All components in the ﬂare state are clearly on the E.M.–
T relation curve for B =50 gauss. 2T high-temperature components in the quiescent state
are also near this curve. These tendency suggests that HAeBe ﬂares and hot components
in the quiescent state are originated in the same types of magnetic activity as solar x-ray

7.4. EVOLUTION OF STELLAR X-RAY ACTIVITY

105

ﬂares. 2T high-temperature components do not exhibit time variations so that they might
be an assembly of microﬂares. 1T or 2T low-temperature components are distributed aroundthe E.M.-T relation curve of coronal emission, and hence might be originated in a stellar
corona. However, these values scatter in a very large range and some plots are on the ﬂarecurve. Because some sources are too weak to identify a time variation, a ﬂaring state might
be misidentiﬁed as a quiescent state.

7.4

Evolution of Stellar X-ray Activity

To ﬁnd out the alternative mechanism to generate magnetic ﬁelds, we investigate the dependence of the quantity LX and kT on the evolutional status of HAeBes. The stellar evolution
is generally simulated in the HR diagram, but the spectral types and bolometric luminosities
could have large uncertainties for a few HAeBes. We thus ﬁrst see the relation with NH , andsecond we investigate the LX dependence on the equivalent width of Hα lines. Finally, we
see the kT and LX dependence on the position in the HR diagram.

7.4.1

NH vs. kT and LX Relation

NH is a good indicator of stellar evolutional status because the surrounding materials aregradually exhausted with age by gathering parts of these circumstellar materials and dissipating the rest by the UV radiation or stellar winds. NH dependence on stellar masses or
inclination to the line of sight is not clear, and hence this is a rough estimate of tracing the
HAeBes evolution. Sources whose best-ﬁt values is NH = 0 is ﬁxed on the AV converted
NH values because they have too large error bars. Then, kT values of 1T or low temperature components seem to steadily rise from 2 keV to 5 keV as NH increases (Figure 7.13).
Only MWC 1080 (NH ∼1×1022 cm−2 ) lies above this line, but this might be the eﬀect on aclose HAeBe–HAeBe binary system (see the section 5.1.14). In contrast, high temperaturecomponent, which have small contribution to the x-ray luminosity, do not change aroundkT ∼2−4 keV. Low temperature plasma components of Large NH sources are diﬃcult to be
detected due to suﬀer strong absorption. It is, however, obvious that HAeBes with smaller
NH , ie elder HAeBes, tend to have lower plasma temperature.
Figure 7.14 shows the relation between NH and LX among the ASCA detected sources.
Square plots means stars earlier than B and dotted plots later than A. LX error bars are
assumed to be the Poisson distribution of source counts. All sources with log LX (ergs s−1 )>32 are B-type stars and have relatively large NH (>1022 cm−2 ). A-type stars are in theluminosity range between log LX (ergs s−1 ) = 30.5−31.5, being independent of NH values.
B-type stars with small NH seem to have low x-ray luminosities, but the luminosities havelarge scatters between log LX (ergs s−1 ) = 30−32.

106

CHAPTER 7. DISCUSSION I – X-RAY EMISSION MECHANISM

Figure 7.13: kT dependence on NH . High temperature components inthe 2T ﬁtting models are shown by squares. Two observing data of
TY CrA and HD 176386 (Obs. ID: RCrA 1 & RCrA 6) are plotted.
S CrA is omitted because it would be a genuine low mass star.

Figure 7.14: LX dependence on NH (square: earlier than B, dot: later
than A). VY Mon is classiﬁed as spectral type B while V892 Tau is as
A. Two observing data of TY CrA and HD 176386 (Obs. ID: RCrA
1 & RCrA 6) are plotted. S CrA is omitted because it would be agenuine low mass star.

7.4. EVOLUTION OF STELLAR X-RAY ACTIVITY

107

Figure 7.15: W(Hα) vs. log LX relation. Marks mean the same as
Figure 7.3.

7.4.2

Equivalent Width of Hα vs. LX Relation

Equivalent width of Hα emission is used for the classiﬁcation of low mass PMSs. This value
is thought to trace the mass accretion rate, and hence the evolutional stage. We thus see thedependence of LX on the equivalent width of Hα, W(Hα) (Figure 7.15). We refer equivalent
width to the Corcoran and Ray (1998) and Damiani et al. (1994). Some objects haveinconsistent W(Hα) values between articles, which could be the time variation or observing
uncertainty. In this case, we applied the value in Corcoran and Ray (1998). We do not see
any clear correlation between W(Hα) and log LX . This might mean that the x-ray emission
is irrelevant to the mass accretion, but it is too early to conclude it because the continuum
levels are much diﬀerent between stars and the Hα line luminosity might be a better to tracestellar evolution.

7.4.3

Evolutional Trends of kT and LX in the HR Diagram

We estimated the kT and LX evolution from the evolutional model provided by Palla andStahler (1993) in the HR diagram. Figure 7.16 shows absorption corrected LX and kT valuesof each HAeBe in the HR diagram. Teﬀ , Lbol , kT and LX values of each source are seen

108

CHAPTER 7. DISCUSSION I – X-RAY EMISSION MECHANISM

in Tables 4.1 and 5.5. In this ﬁgure, circle size is proportional to logarithmic values ofabsorption corrected x-ray luminosity and the shade of circles means plasma temperature.
Samples below the threshold luminosity log LX (ergs s−1 ) = 30, which is determined from the
WTTS LX function (see the section 7.1) are shown by crosses. We determine the positionof SSV 63E+W in the HR diagram by combining the information that Lbol ∼22.4 L , age<106 yrs and the mass 1.5 M < M <5 M . The position in the HR diagram is then close
to EC95.
We clearly see that samples on the upper right of Figure 7.16 has high temperatureplasma kT >4 keV and samples closer to the main-sequence tend to have lower plasma
temperature. The upper right in the HR diagram corresponds to the younger stars (seethe subsection 2.1.3). These results are consistent with the relation with NH values in thesubsection 7.4.1. The similar trend seems to be present in LX . The luminosity on the upper
right in the diagram is nearly LX ∼ 1032 ergs s−1 and gradually decreases as they close to themain-sequence. Nevertheless, the ASCA samples are limited to trace the overall picture ofthis LX trend. We thus omit the information on temperature and add samples from ROSAT
results (HAeBe:Zinnecker and Preibisch 1994, the Orion molecular cloud:Gagné et al. 1995).
We collect source samples whose detection limit is log LX (ergs s−1 ) <30, We also include
late-type samples with log Teﬀ (K) >3.65 (spectral type G) listed in Gagné et al. (1995)
because some would be progenitors of A or late-B type stars. These samples do not havekT information for the lack of spectral resolution or the constraints on photon statistics.
We calculate the ASCA band luminosity assuming the average kT ∼2 keV of the ASCA
detected sources, but the conversion factor only changes less than factor of two within therange of stellar plasma temperature, kT ∼0.5−5 keV (see Table 4.8). Figure 7.17 shows the
LX map of HAeBes in the HR diagram. The LX trend in the ASCA data is clearly seen inthis picture. The most remarkable feature is that most sources elder than 106 -year-old (3.8
< log Teﬀ < 4.1, 1.3 < log Lbol < 1.9) emit faint or no x-rays, ie log LX (ergs s−1 ) <30. We
might say that the age 106 -year-old would be at the end point of HAeBe x-ray activity.
Damiani et al. (1994) found a correlation between LX and spectral type or Lbol but
Zinnecker and Preibisch (1994) did not (see the subsection 2.2.2). We now see that thisdiﬀerence is simply due to the selection eﬀect. Damiani et al. (1994) collected HAeBe
samples close to the main-sequence so that they found a similar x-ray characteristics to MS
stars. While Zinnecker and Preibisch (1994) collected young A-type samples such as V380
Ori and HR5999. Their x-ray luminosity is comparable to early B-type stars and thus no
LX dependence was seen.

7.5

Models to Generate the Magnetic Field on HAeBes

The large scale magnetic ﬁelds connecting a star and its disk seems to account for the x-ray
activity on HAeBes (see the section 2.2.2 and the appendix A.3). In previous sections, we ﬁnd

7.5. MODELS TO GENERATE THE MAGNETIC FIELD ON HAEBES

109

log LX (ergs/s)
32
30
below
kT (keV)> 3.5
1.5 − 3.5
< 1.5

Figure 7.16: X-ray luminosity and plasma temperature of the HAeBe samples in the HR
diagram. Circle size is proportional to log LX and ﬁlled gray color indicates the plasma
temperature. The threshold luminosity is set at log LX (ergs s−1 ) =30.0. Sources below thisthreshold are shown with crosses. The diagram is superimposed on a evolutional model by
Palla and Stahler (1993). A dotted line means the birthline when accretion rate is Ṁ = 10−4M yr−1 .

110

CHAPTER 7. DISCUSSION I – X-RAY EMISSION MECHANISM

log LX (ergs s-1)
32
30.0
below

Figure 7.17: X-ray luminosity of the young stars in the HR diagram, including the ROSAT
and Einstein samples. Details are the same as the previous ﬁgures but the information ofplasma temperature is omitted.

7.5. MODELS TO GENERATE THE MAGNETIC FIELD ON HAEBES

111

Table 7.5: Young stars in the Orion cloud (Gagne et al. 1995)
Teﬀ
Lbol
v sin i
LX
LX /Lbol
Objects
Teﬀ
Lbol
v sin i
LX
LX /Lbol
log
log
log
log
log
log
log
log
K
L
km s−1ergs s−1
KL
km s−1ergs s−1
P1491
4.03
2.43
–
30.6
−5.4
P1799
3.77
0.91
123
30.5
−4.0
JW373
3.68
0.64
−
30.1
−4.2
P1623
3.93
1.79
171
<29.8
<−5.6
P1671
3.92
1.55
−
<29.9
<−5.2
P1828
3.68
0.69
29
31.1
−3.2
P1817
3.70
1.02
35
31.7
−2.9
P1712
4.03
2.00
75
30.6
−5.0
P1708
4.14
2.85
−
<29.9
<−6.5
P1910
3.68
1.25
−
31.4
−3.4
P1937
3.77
0.93
−
30.3
−4.2
P1795
4.03
1.87
−
29.9
−5.6
P1955
3.74
1.44
109
31.5
−3.5
P1798
4.25
2.93
50
31.2
−5.3
P1950
3.74
0.49
27
31.0
−3.0
P1891
4.58
5.01
140
32.7
−5.9
P1889
4.49
4.44
145
31.9
−6.2
P1972
3.68
0.71
<20
30.2
−4.1
P2001
3.72
0.16
<20
31.1
−2.7
P2058
3.99
1.59
170
<29.7
<−5.4
P1393
3.72
0.49
<20
31.0
−3.0
P2033
3.72
0.80
56
31.1
−3.3
P1374
3.80
0.15
77
30.7
−4.1
P2032
3.75 −0.03
−
30.2
−3.7
P2057
3.71 −0.20
−
30.4
−3.0
P1409
3.72
0.80
29
30.9
−3.5
P2069
3.68
0.60
38
31.3
−2.9
P1425
3.74
0.80
26
30.9
−3.5
P1455
3.78
1.03
21
30.6
−4.0
P2086
3.81
1.33
72
<29.9
<−5.0
P2084
3.68
0.88
<20
30.7
−3.7
P1510
3.70
0.27
38
31.2
−3.3
P1518
3.70 −0.06
38
30.6
−2.9
P2100
3.72
0.83
72
30.8
−3.6
P2167
3.72
1.10
57
31.3
−3.4
P1505
3.70
0.27
<20
31.0
−2.9
P1541
3.71
0.55
<20
31.0
−3.2
P2152
3.66
0.35
−
31.1
−2.9
P2181
3.70
0.58
−
31.1
−3.0
P1553
3.70
1.03
28
31.3
−3.4
P1587
3.70
0.88
<20
31.3
−3.2
P2244
3.68
0.87
51
31.1
−3.3
P2253
3.72 −0.05
−
30.5
−3.0
P1659
3.70
1.16
<20
31.5
−3.3
P2252
3.74
0.90
33
30.8
−3.7
P1736
3.75
1.57
41
30.9
−4.2
P1773
3.68
0.33
−
30.9
−3.1
P2257
3.72
0.58
47
31.2
−2.9
P2347
3.71
0.03
−
30.6
−3.0
P1760
3.68
0.21
−
<29.9
<−3.9
P1785
3.72
0.80
<20
31.0
−3.4
P2368
3.66
0.34
47
30.6
−3.3
Objects in the right and below P1393 in the left are young stars in the Orion cloud (Gagné et al. 1995) – above cross:early type, below cross: late type.
Objects

the similarity of the x-ray property with low mass protostars (eg plasma temperature, ﬂares).
This fact lets us suppose that the x-ray emission mechanism is also the same as protostars, ie
the magnetic activity linking a star and its disk. This model reconciles the lack of mechanism
to generate the magnetic ﬁeld on stellar surface. The magnetic ﬁelds are originated in
the fossil ﬁeld, remaining from the parent molecular cloud. The presence of circumstellar
disks around HAes is widely accepted by the observation of mm radio continuum survey.
The model also supports the origin of outﬂows. Hayashi et al. (1996) simulate that thereconnection accompanies a gas eruption, which might be seen as a jet or outﬂows. A certain
HAe continues the outﬂow activity till a few Myrs (Devine et al. 2000), which correspondwell with the decaying age of x-ray activity (see the subsection 7.4.3). The long ﬂare durationof HAeBes might be due to the large plasma volume, which needs long propagation time ofshocks. Moreover, this model explains the high velocity of accreting materials (Blondel et al.
1993, see also the subsection 7.2.2). Magnetic reconnection, which occurs at a mid-point
between a star and its disk, accelerates gases and forces them to accrete onto a star.
A problem lies in that there seems no apparent change in the structure of circumstellar
disks around 106 -year-old. For example, AB Aur and HD 97048, older than 106 -year-old, still
have relatively massive disks. We see but do not ﬁnd any correlation between the luminosity

112

CHAPTER 7. DISCUSSION I – X-RAY EMISSION MECHANISM

Figure 7.18: Radio continuum (λ1.3mm ) vs. LX relation. Marks are
the same as Figure 7.3.
of 1.3mm dust continuum emission, which is a good tracer of dense and cold circumstellar
envelopes (Figure 7.18). These stars might already dissipate inner most of their disks, whichcannot be resolved. On the contrary, even though there is no clear diﬀerence in the x-ray
property between HAes and HBes, the circumstellar environment seems to be much diﬀerent.
HBes lack a compact continuum emission expected from a circumstellar disk and form large
empty cavities. There are, however, a few pieces of supporting evidence for the star-diskmagnetic ﬁeld model. Drew et al. (1997) reports that even the very massive HBe MWC 297
shows the presence of an outﬂow phenomenon though they interpret its origin as radiative
driven disk winds (Drew et al. 1998). Young massive MS star θ1 C Ori has an inverse PCygni proﬁle in its spectrum, indicating the accretion ﬂows (Conti 1972). Moreover, someearly B-type MS stars are classiﬁed as β Pictoris type stars, having debris disks (eg Coulson
et al. 1998, see the section 2.1.4).

Chapter 8
Discussion II – A Uniﬁed Picture of
Stellar X-ray Activity
8.1

Evolution of the X-ray Activity to the Main-sequence
Phase

One of the main conclusions in the previous section is that the x-ray activity of intermediate
mass stars decays after 106 -year-old. Massive stars larger than 4M reach ZAMS before thisthreshold age. It suggests that the massive stars could sustain the magnetic activity during
the former phase of MS. Here, we demonstrate that this picture can solve most problems
posed in the x-ray emission from high mass MS stars.

8.1.1

X-ray Emission from Late B Type Stars

From a RASS data survey of OB stellar x-rays, Berghoefer et al. (1997) reported that x-rays
and bolometric luminosity relation of LX /Lbol ∼ 10−7 was valid for massive stars down toa spectral type B1-1.5. Stars below this spectral type had no clear relation and have large
scatters between these quantities. They also mentioned that the detection rate of these stars
are very low (P < 10%). These results suggest that the x-ray emission mechanism of massive
stars, which is believed to be the stellar wind origin, changes at the spectral type B1-1.5.
They invoked the x-ray origin of late B-type stars on extremely young low mass companion
like CTTS, but some sources seem to have too large x-ray luminosities compared with low
mass stellar x-rays. We therefore hypothesis that most sources detected in this survey areyoung and still possess the magnetic activity continuing from the PMS phase. There is the
supporting evidence for the magnetic activity on young B-type MSs.
113

114CHAPTER 8. DISCUSSION II – A UNIFIED PICTURE OF STELLAR X-RAY ACTIVITY

A0

B5

B0

O5

33
32

log(Lx)

31
30
29
28
27

detections
upper limits

-7

~ 10

35

36

37

38

39

40

log(LBol)

Figure 8.1: LX vs. Lbol relation of Herbig Ae/Be stars superimposedon the plots of OB stars provided by Berghoefer et al. (1997). Plotsof HAeBes are shown with black crosses while ones of OB stars areshown with grays.
LX vs. Lbol Relation We superimposed HAeBe data on the LX vs. Lbol graph of OB stars
provided by Berghoefer et al. (1997) (see Figure 8.1). All HAeBe plots are just above
the higher end of OB star. Considering the stellar evolution, LX /Lbol values decrease
with age and the x-ray ﬂux ﬁnally drop below the detection level. This is the reason
why LX /Lbol values scatter in a wide range.
Detection Rate Berghoefer et al. (1997) reported that the detection rate continuously
decreases toward later O and early B-type stars and drops to <10% between B2−B9.
The transition age a few×106 -year-old is roughly 10% of the life time of B2 stars(2.4×107 yrs, see also Table B.2).
Time Variability Smith et al. (1993) detected a giant x-ray ﬂare on the MS B star λElidani (B2e) with ROSAT PSPC. They concluded that the ﬂare arises on or related
to λ Elidani itself because of no evidence of existing a companion star in IUE andoptical spectra. They also suggested λ Elidani would not be a classical Be star. The
fact that λ Elidani shows emission line then demonstrates its youth. It then has acharacteristics of youth,8.1. EVOLUTION OF THE X-RAY ACTIVITY TO THE MAIN-SEQUENCE PHASE115

Table 8.1: X-ray characteristics of massive main-sequence stars
Temperature
Emission Measurelog LX(keV)(cm−3 )
(ergs s−1 )δ Ori.
O8+B0.5
1
Cool
0.25 (0.21–0.30) 55.20 (55.03–55.37)
32.1
Hot
0.60 (0.45–0.79) 54.73 (54.59–54.87)
31.6
λ Ori
O9.5+B01
Cool0.29 (0.25–0.34) 55.09 (54.97–55.21)
32.1
Hot
2.27 (1.14–4.52) 54.24 (54.01–54.47)
31.2
τ Sco. B0V
2Cool
0.62 (0.62–0.64) 54.54 (54.50–54.57)
31.8
Hot
3.67 (2.59– )53.77 (53.48–53.96)
30.9
Hot2
1.06 (0.88–1.24) 53.91 (53.75–54.01)
31.0
ζ Oph. O9.5V(e)31T
0.54 (0.52–0.56) 54.20 (54.08–54.30)
31.2
ζ Ori.
O9.7Eb+O9.5IV
3
1T
0.43 (0.41–0.45) 54.34 (54.15–54.48)
31.2
The energy band of LX is between 0.5 – 10 keV. References: 1.Corcoran et al. (1994), 2.Cohenet al. (1997), 3.Kitamoto et al. (2000)
Target

Sp.T

ref.

Comp.

8.1.2

X-ray Emission from O and Early B Type Stars

In Figure 8.1, stars earlier than B1-1.5 has a relation of log LX /Lbol ∼ −7, and the UV driven
stellar wind seems to be the dominant activity. Nevertheless MS O-type stars are generally
younger than B-type stars, and hence should possess the same kind of magnetic activity as
late-type stars. We hypothesis that this component is seen as the hard tails in the x-ray
spectra. The hard tails were believed to be emerged from the inverse Compton mechanism
when the stellar UV photons were scattered by relativistic electrons, which is accelerated by
the 1st order Fermi shocks (Chen and White 1991). However, the x-ray spectra of some Otype stars including the super massive star ζ Pup (Corcoran et al. 1993) can be reproduced
by a simple one-component plasma model (see the discussion of Kitamoto and Mukai 1996).
The inverse Compton model does not seem to explain the variety of this tail component.
We thus compared LX and kT values of hot and cool components in O stars with thoseof HAeBes (Table 8.1, Figure 8.2). The reference papers only show E.M. and kT values sothat we calculated the x-ray ﬂuxes using xspec 9.0. We clearly see that the hard tails isquite similar to the values of HAeBes. In contrast, the soft x-ray component is on the otherregion, high LX and low kT.
PMS natures in massive MS stars have never been predicted. However, a recent Chandra
observation demonstrated the evidence of magnetic activity on O stars (see the subsection
2.2.1). On the other hand, the hottest star in the Orion cloud θ1 Ori C showed a long term xray variation (Caillault et al. 1994). Because θ1 Ori C also shows an inverse P-Cygni proﬁle,this phenomenon might be associated with the accretion process linking to the magnetic
activity.

116CHAPTER 8. DISCUSSION II – A UNIFIED PICTURE OF STELLAR X-RAY ACTIVITY

Figure 8.2: kT vs. LX relation of MS massive stars and Herbig Ae/Bestars. Open squares and circles are soft and hard components inmassive stellar spectra respectively. Filled squares and circles arethose in HAeBes.

8.2

X-ray Emission from Giant Molecular Cloud Cores

GMC cores are the birth place of HAeBes (see the subsection 2.2.2). When we trace backthe evolution of magnetic activity on HAeBes, we need to study the x-ray emission from
GMC cores. These regions are generally very crowded and ASCA cannot clearly identify the
x-ray sources, which probably include HAeBe progenitors, ie intermediate and high mass
protostars. We ﬁnd the presence of extremely high plasma components in Monoceros R2
region (see the section 6.1), which might be originated in high mass protostars.
Table 2.1 summarizes ASCA observations of ﬁve GMC cores. These x-ray properties are
smeared over slightly large areas including surrounding soft x-ray sources, most of which areprobably high mass MSs and low mass young stars. Therefore these NH values are averageon including x-ray sources. However, high kT values would represent x-ray emission from
HAeBe progenitors because soft x-ray sources do not contribute the hard x-ray ﬂux which
determines kT. We thus compare these kT values with the HAeBe temperature (Figure 8.3).
We omit the data of Orion trapezium where the x-ray emission from the massive MS star
θ1 Ori C is dominant. Then, the rising trend of kT values seen in HAeBe samples smoothly
connects with kT of GMC cores. This result not only suggests that the hard x-ray ﬂux
comes from HAeBe progenitors but also implys that very young HAeBes have more violent
x-ray activity. Of course, observations of GMCs with high spatial resolution are essential

8.3. A UNIFIED PICTURE OF X-RAY ACTIVITY ON YSOS

117

GMC

Figure 8.3: kT distribution of giant molecular cloud and Herbig Ae/Bestars. kT values of GMC are plotted on the right side of Figure 7.13.
to address this problem. Chandra will reveal the nature of intermediate and high mass
protostars.

8.3

A Uniﬁed Picture of X-ray Activity on YSOs

We ﬁnd that x-ray emission from intermediate to high mass YSOs is originated in themagnetic activity. This means that magnetic activity plays an important role on all types
of YSOs from low mass to high mass stars. Here we describe the evolution of x-ray activity
on YSOs.
Stellar x-ray emission begins in very young phase when stars are deeply embedded inmolecular cloud cores, ie protostar phase. In this phase, stars possess ’fossil’ magnetic ﬁelds
inheritant from the parent clouds and exhibit violent magnetic activity linking a star andits disk, showing occasional x-ray ﬂares. In this process, a part of infalling gases goes out
from the system with large amount of angular momentum, which is later seen as molecular
outﬂows or optical jets. The rest of infalling gases streams down to the stellar surface,accelerated by reconnection shocks. Accreting shocks heat gases up to the level of x-ray
emission.
The link between a star and its disk fades around ∼106 yrs. After decoupling the link,
low mass stars continues magnetic activity on the stellar surface by operating αΩ dynamo,

118CHAPTER 8. DISCUSSION II – A UNIFIED PICTURE OF STELLAR X-RAY ACTIVITY
but stars earlier than F5 do not have such mechanism and the x-ray activity rapidly decays.
Because stars between A – F5 gradually evolve to the MS, ’dead zone’ of the x-ray activity
appears in the HR diagram. In contrast, earlier than B type stars reaches to the mainsequence before their x-ray activity ceases. Young MS stars - several percentage of all MS
stars - still possess the x-ray activity. Stars earlier than B1-1.5 have enough UV photons tooperate the wind driven x-ray activity. The x-ray emission by the magnetic activity becomesa minor component, but it is seen as hard tails.

8.3. A UNIFIED PICTURE OF X-RAY ACTIVITY ON YSOS

5

Star & Disk Dynamo
Solar Type Dynamo
Stellar Wind

Figure 8.4: Uniﬁed picture of x-ray emission from stars.

119120CHAPTER 8. DISCUSSION II – A UNIFIED PICTURE OF STELLAR X-RAY ACTIVITY

Chapter 9
Conclusion & Future Problem
The x-ray emission from HAeBes has long been puzzled. This is partly because the proper
x-ray emission mechanism is not known for these intermediate mass young stars and partly
because the satisfying observing results have not been obtained. By combining our ASCA
analyses with previous ROSAT and Einstein results, we ﬁnd that the x-ray emission of
HAeBes has its roots in the earliest phase of stars and continue to the former stage of theMSs. We list the results and implications below.
1. In the analyses of the ASCA data, 16 x-ray counterparts among 40 HAeBe candidates
are detected. The x-ray luminosity is ranging between log LX (ergs s−1 ) = 30 – 32,
which is signiﬁcantly larger than the typical luminosity of low mass stars. The x-ray
emission from low mass companions is unlikely.
2. 8 sources among the 16 detected sources (50%) exhibited time variation. Among 8
sources which have enough photon counts (above 1000 photons), 6 source showed time
variation.
3. All x-ray spectra are successfully reproduced with absorbed one or two temperaturethin-thermal (MeKaL) models. The plasma temperature ranges between kT ∼1−5
keV, which is signiﬁcantly higher than that of main-sequence OB stars (kT <1 keV).
4. Several x-ray emission models are investigated for HAeBe x-rays. Stellar wind shock,wind-fed magnetosphere and accretion shock models cannot explain the HAeBe x-ray
emission. Magnetic activity is the most probable x-ray emission mechanism.
5. There is no clear correlation between LX and the rotational velocity (v sin i) or the
angular velocity (ω sin i). The solar type dynamo mechanism does not seem to work
on HAeBes.
121

122

CHAPTER 9. CONCLUSION & FUTURE PROBLEM

6. Time proﬁles of sources showing prominent variability are investigated. All decay timescale is a few × 105 sec. This timescale is much longer than low mass TTSs and nearly
the same as the timescale of low mass protostars. The ﬂare size of an large x-ray ﬂarefrom MWC 297, estimated with the radiative cooling model, is far larger than the solar
radius but can be smaller than the radius of MWC 297.
7. The plasma temperature kT decreases as the observed NH decreases whereas LX depends more on stellar masses. This result indicates that the x-ray activity decays asstars evolve. By combining previous ROSAT and Einstein results, we ﬁnd clear tendency that the x-ray activity of young stars decays after the stellar age ∼ 106 -year-old.
Considering that the similarity of x-ray property and the circumstellar environment
to low mass protostars, the x-ray emission from HAeBes seems to be originated in themagnetic activity linking a star and its accretion disk.
8. PMS x-ray activity continues to the former of the MS lifetime earlier than B-type stars.
The x-ray emission driven by the stellar wind shocks seems to be dominant in stars
earlier than spectral type B1-1.5, but the x-rays originated in the magnetic activity is
still seen as hard tail component.
9. ASCA detected hard x-ray sources in the Monoceros R2 GMC core. The plasma
temperature is very high kT ∼ 5 keV. All GMC cores observed with ASCA have thesimilar temperature plasma. The rising trend of kT values seen as tracing back the
HAeBe evolution smoothly connects to the kT of GMC cores. This suggests that hard
x-rays from GMC cores are originated in HAeBe progenitors, ie high mass protostars.
In this study, a universal picture of the x-ray emission from higher mass YSOs are
obtained for the ﬁrst time. We think these three subjects as the next main problems.
• Can we obtain the deﬁnite evidence of magnetic activity on higher mass PMS? HAeBe
samples showing time variation is still limited and the data quality is not high enough.
A couple of x-ray variable sources with high signal-to-noise ratio are needed to address
the dynamics of hot plasma. On the other hand, high resolution spectroscopy addresses
the plasma density and temperature from atomic satellite lines. The information provides us the plasma geometry, for example, as whether the plasma is near or far fromthe stellar surface. This tool has already been utilized in the spectra of massive MSs(see the subsection 2.2.1). This will be a powerful tool for the analysis of PMSs.
• How is the structure of hot plasma around stars? There is no direct evidence toprove the structure of magnetic ﬁelds connecting a star and its disk. As we mentioned
before, the x-ray ﬂare could be strongly related to the jet formation. If the Doppler shiftcomponent in the PMS x-ray ﬂare is detected, it will provide the important informationon the jet acceleration mechanism and the structure of large scale magnetic ﬁelds.

123• When does the x-ray activity begin in high mass young stars? Each x-ray source in
the GMC core is not still resolved. By investigating x-ray properties of very young
massive stars, the early evolution of x-ray activity, and hence magnetic activity, will
be clear.
For statistical analyses of young higher mass stars, Chandra will provide sophisticateddata without source contamination. Time proﬁles with high photon statistics can be obtainedby XMM-Newton with very large eﬀective area. On the other hand, high resolution grating
spectroscopy onboard both Chandra and XMM-Newton is limited to soft x-rays around 1
keV. Micro calorimeter onboard Astro-E II will provide a powerful tool for performing highresolution spectroscopy of embedded sources. The x-ray emission from a young star provides
important information close to the stellar body, but the study is now on the starting-point. In
this ten years, the study will be greatly advanced with the new generation x-ray telescopes.

124

CHAPTER 9. CONCLUSION & FUTURE PROBLEM

Appendix A
Models of the Stellar Magnetic
Activity
A.1

Hydrodynamic Decay, Sustained Heating Model

Stellar x-ray ﬂares provide the nature of the magnetic activity on stars. The Algol system
is known to exhibit frequent x-ray ﬂares and has been intensively studied. Flare decay isconventionally regarded as the cooling phase of thermalized plasma - the radiative coolingmodel. On this assumption, a ﬂare volume becomes a few times as large as the stellar
size (eg White et al. 1986). This result is supported by a detailed analysis containing aneﬀect of heat conduction (the quasi-static cooling model, van den Oord and Mewe 1989).
Schmitt and Favata (1999) however indicated that the real ﬂare volume directly measured
from an eclipsing event was much smaller than the volume in the conventional estimate.
This suggests the presence of continuous heating during the ﬂare. Reale and Micela (1998)
applied a model of continuous heating during ﬂares to a stellar x-ray ﬂare with the ROSAT
data. Favata et al. (2000) developed this method into an ASCA result and derived the ﬂare
volume is much smaller than the result of the radiative cooling model. According to these
papers, the presence of additional heating is investigated by a slope of the locus of the ﬂaredecay in the log E.M. – log T plane. This model describes ﬂare parameters as the following
equations.
τLC= F (ζ) = cae−ζ/ζa + qa
(A.1)
τth
√
τLC TmaxL=(A.2)
αF (ζ)
where τth : the cooling time scale of plasma, τLC : a decay time scale of a light curve, and ζ:a slope of a ﬂare decay in the log E.M. – log T plane, and α = 3.7 × 10−4 cm−1s−1 K 1/2 .
The values ca, ζa and qa depend on a response of observed detectors, and ca = 10.9, ζa =
125

126

APPENDIX A. MODELS OF THE STELLAR MAGNETIC ACTIVITY

FigureA.1: Evolution of the MWC 297 ﬂare parameters in the log√
E.M. – log T plane. The data points of the 2nd and 3rd observations are plotted in the plane. The solid line shows the slope of theseplots.
0.56 and qa = 0.6 for GIS. Tmax is the maximum temperature of a ﬂaring loop and Favata
et al. (2000) derived an empirical relation to the observed maximum temperature Tobs .
1.176
Tmax = 0.085 × Tobs

(A.3)

The data of MWC 297 (the 2nd and 3rd observations) are plotted on the log E.M. –log T
plane (Figure A.1 and Table A.1). Because our plots have large errors, the uncertainty of ζis very large. The semi-length L derived from the best-ﬁt ζ value is as same as Lα=0.1 (seethe subsection 7.3.3) in the radiative cooling model though the uncertainty of ζ is too large.
Table A.1: Physical parameters of the MWC 297 ﬂare in the hydrodynamic model
ζ
Tmax
F(ζ)
τth
L

0.82 (0.32–1.32)
1.5 ×108 ◦K
3.1
1.78 ×104 ◦ K
5.9 ×1011 cm

A.2. SHEAR DYNAMO MODEL

127

Figure A.2: Evolution of relative x-ray luminosity proposed by Toutet al. (1995).

A.2

Shear Dynamo Model

Tout and Pringle (1995) proposed a dynamo mechanism on HAeBes which is driven by ashear component holding in very young stars. In this model, HAeBe x-ray emission comesfrom a corona heated by this magnetic activity. It predicts that x-ray luminosity varies withage as follows (see also Figure A.2),

where

LX (t) = LX0 (1 + t/t0 )−3

(A.4)

LX0
∆Ω0 3 γ 3 M∗ 5/2
R∗∼ 8 × 102 ε() ( −4 ) (
) (
)−5/2
L
ΩK
10
10M
3 × 1011(A.5)

t0 ∼ 4.8 × 104 (γ∆Ω0 −1 k 2
M∗ −1/2
R∗) ( )( −4 )−3 ()()3/2 yrΩK
0.1 1010M
3 × 1011(A.6)

and ∆Ω0 /ΩK : ratio of the initial value of the rotational shear in the star to its breakup
velocity, γ: eﬃciency of ﬁeld generation, and k: proportionality factor between the stellar
moment of inertia and M∗ R2∗ . In this formula, there are several unknown parameters. Toutand Pringle (1995) suggest that this luminosity variation is consistent with the HAeBeluminosity provided by Zinnecker and Preibisch (1994) if these parameters are ∆Ω0 /ΩK ∼1, γ ∼ 3 ×10 and / ∼ 10−3 .

APPENDIX A. MODELS OF THE STELLAR MAGNETIC ACTIVITY

128

This model seems to have a few weak points. It does not predict the time evolutionof the plasma temperature. The result in the subsection 7.4.1 indicates that the plasma
temperature decreases as stars evolve, but this model assumes that energy is constantly supplied to a corona, and does not seem to predict a systematic change of plasma temperature.
Second, x-ray luminosity only depends on stellar mass, radius and age except for its dynamo
parameters (see the equation A.4, A.5). Some sources however have signiﬁcantly small x-ray
luminosity compared with sources on a similar evolutional tracks (see Figure 7.17) although
the diﬀerence might be due to the uncertainty of the distance.

A.3

Diﬀusion Time Scale of the Magnetic Field

The star-disk magnetic ﬁeld model assumes the presence of ﬁssile magnetic ﬁelds inheritedfrom the parent molecular cloud so that we check the dissipating time scale of the magnetic
ﬁeld on stars (Ref. Kato 1989).
When plasma gases are stable (v = 0), 4th Maxwell equation can be written as∂B= η∇2 B∂t

(A.7)

This corresponds to the diﬀusion equation. Magnetic ﬁeld therefore dissipates without
plasma motion. The diﬀusion time scale (τD ) approximates asτD ∼ L2 /η(A.8)

where η is the magnetic diﬀusivity. The convection zone on the sun (L ∼1000 km) is about
108 yrs. Because stellar magnetic ﬁelds are generally larger than this scale, the diﬀusionof the magnetic ﬁeld is negligible during PMS phases. This also suggests that the decay of
x-ray activity is not caused by the dissipation of magnetic ﬁelds.

Appendix B
Stellar Physical Parameters
We list the typical stellar parameters and timescales of a stellar evolutional model.

129

APPENDIX B. STELLAR PHYSICAL PARAMETERS

130

Table B.1: Stellar parameters
Sp. type

Teﬀ

B−V

Mass

Radiusvrot
Type V

Type III

(K)
−
42000
−

Type I(M ) (R )(km/s) (km/s) (km/s)
O3
−
120
15
O5
−0.31
60
12
O6
−
37
10
O8
23
8.5
200
125
B0
30000 −0.30
17.5
7.4
170
120
102
B3
−
−
7.6
4.8
B5
15200 −0.17
5.9
3.9
240
130
40
B8
11400 −0.11
3.8
3.0
A0
9790 −0.02
2.9
2.4
180
100
40
A5
8180 +0.15
2.0
1.7
170
38
F0
7300 +0.30
1.6
1.5
100
30
F5
6650 +0.44
1.4
1.3
30
< 25
G0
5940 +0.58
1.05
1.1
10
< 25
G5
5560 +0.68
0.92
0.92
< 10
< 20
< 25
Case of class V (main sequence) except for vrot (Drilling and Landolt 2000).

Table B.2: ZAMS age & MS life time
Mass ZAMS age MS life time
M
(yrs)
(yrs)
30.0
−
5.9×106
15.0−
1.1×107
9.0−
2.4×1075
6.0
4.00×10−5.0
2.32×10 5
7.8×107
4.0
8.16×10 5−
3.0
2.02×10 6−
2.0
8.35×10 6
8.0×108
1.5
1.20×10 7
1.8×1097
1.0
3.16×101.0×1010
ZAMS age referred to a model provided by Palla andStahler (1993). Life time on the MS is referred toa model, assuming the initial condition of the main
sequence phase with the typical chemical compositionof Population I (H: 70%; He: 27%; Heavy elements
3%, Audouze 1994).

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Acronyms & Abbreviations
ASCA
CCD
CTTS
DFE
GIS
GMC
HAEBE
HEASARC
HRI
IR
IRAS
ISAS
ISM
IUE
KSC
MS
PI
PMS
PSF
PSPC
RASS
ROSAT
RDD
RDF
SAA
SED
SIS
UKIRT
UV
WTTS
XRT
YSO
ZAMS

Advanced Satellite for Cosmology and Astrophysics
Charge Coupled Device
Classical T-Tauri Star
Dark Frame Error
Gas Imaging Sepctrometer (onboard ASCA)
Giant Molecular Cloud
Herbig Ae/Be starHigh Energy Astrophysics Science Archive Research Center
High Resolution Imager (onboard ROSAT)
Infra-RedInfrared Astronomycal Satellite
Institute of Space and Astronautical Science
InterStellar Medium
International Ultraviolet Explorer
Kagoshima Space Center
Main Sequence
Pulse height Invariant
Pre Main Sequence
Point Spread Function
Position Sensitive Propotional Counter (onboard Einstein & ROSAT)ROSAT All Sky SurveyRöntgensatellit (Germany)
Residual Dark Distribution
Rationalized Data Format
South Atlantic Anomaly
Spectral Energy Distribution
Solid-state Imaging Spectrometer (onboard ASCA)
United Kingdom InfraRed Telescope
Ultra Violet
Weak-line T-Tauri Star
X-Ray Telescope
Young Stellar Object
Zero Age Main SequenceAcknowledgment
I should be grateful to Prof. Koyama K. for guidance of my research and ﬁnancial supports throughout my doctor course. His research into the x-ray activity of low-mass youngstars have attracted me and drove me to the study of higher-mass stars. This research could
not be successful without his warm advice and severe criticism. I should also be gratefulto Assoc. Prof. Awaki H. and Tsuru T. of their useful suggestion and technical support. I
learned how to research from them. I thank to other staﬀs Prof. Tanimori T., Prof. MaiharaT., Dr. Tazawa Y. and Dr. Iwamuro F. Discussions with them helped me much to proceed
my research. I also wish to thank all my colleagues. I was sometimes encouraged by them,and above all always very enjoyed chatting with them.
I especially express my thanks for persons directly supporting this research. Assoc.
Prof. Yamauchi S. gave me much advice about the methods to analyze the ASCA dataand information about latest results of stellar astrophysics. Prof. Kitamoto S. told me therecent understandings of the x-ray emission from massive MS stars. Assc. Prof. Nakano
M. told me the information about stars in the Orion cloud. Prof. Feigelson E.D. and Prof.
Montmerle T. kindly discussed the subject of star formation with me many times. I was
always realized of the weakness of my hypothesis from them. Dr. Terada H. told me quick
results of the ASCA data of a few HAeBes. Dr. Tsuboi Y. told me the ASCA results ofmany low-mass star forming regions. Mr. Imanishi K. told me the latest result of ρ-Ophcloud with Chandra. Dr. Ozawa H. told me the ASCA result of HH63E+W. Prof. Shibata
K. and Dr. Yokoyama T. told me the basic mechanism of the magnetic reconnection andtopics of solar physics. Assoc. Prof. Sakao T. told me the result of the Yohkoh satellite. Dr.
Nishi R. told me problems of models of star formation. I would also like to thank others
supporting me and all the members of the ASCA team.
This work is supported by the grants JSPS, research fellowships for young scientists, andmade use of data obtained through the High Energy Astrophysics Division Online Service,provided by Institute of Space and Astronautical Science (ISAS) and the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA/Goddard
Space Flight Center.
Finally, I would like to thank my family and friends.