X-ray Study of Low-mass Young Stellar Objects
in the ρ Ophiuchi Star-forming Region with
ChandraKensuke Imanishi
Department of Physics, Graduate School of Science,
Kyoto University, Sakyo-ku, Kyoto 606-8502, Japan

Submitted to the Department of Physics, Kyoto University
in partial fulfillment of the requirements
for the degree of Doctor of Science
January 6, 2003

Abstract
We present the results of a systematic X-ray study of low-mass young stellar objects, using two
deep exposure Chandra observations of the main region of the ρ Ophiuchi star-forming cloud. Fromthe central 1 pc × 2 pc region, we detect 195 X-ray sources with the detection limit of the X-ray
luminosity of ∼1028 ergs s−1 . About 60 % of them (110/195) are identified with optical and/or
infrared counterparts, including significant numbers of class I, II, III sources, and some young browndwarfs. Most of the X-ray spectra are well fitted by a single temperature thin thermal plasma
model with the elemental abundance of 0.3 solar, although some sources show unusual features;
overabundances (GY21, S2, and WL22), abundance variation among elements (DoAr21, ROXs21,
and ROXs31), multi-temperature spectra (ROXs21 and ROXs31), 6.4 keV line emission (YLW16A),
and a possible non-thermal spectrum with the absorption edge of Ca or Ar (S1). The absorption
column density (NH ) is spread around the average value through the cloud (∼1022.5 cm−2 ) and is
proportional to the visual extinction (AV ). The time-averaged plasma temperature (< kT >) is in
the range of 0.2–10 keV. The time-averaged X-ray luminosity (< LX >) is well correlated with thebolometric luminosity (Lbol ), stellar mass, and radio (cm) luminosity. We detect 71 X-ray flares.
Most of them have the typical profile of those of solar and stellar flares (fast rise and slow decay),while some bright flares show unusually long rise timescales. We find that (1) class I–II sources
tend to have higher < kT >, which sometimes exceeds 5 keV, (2) the distribution of < LX >during flares is nearly the same for all the classes from ∼1029.5 to ∼1031.5 ergs s−1 , although there
is a marginal hint of higher < LX > distribution for class I than class II–III, and (3) positive
and negative log-linear correlations are found between the rise (τr ) and decay (τd ) timescales, and< kT > and τr . In order to explain these relations, we use the framework of magnetic reconnection
model with heat conduction and chromospheric evaporation (Shibata & Yokoyama 1999, 2002) toformulate the observational parameters (τr , τd , and < kT >) as a function of the pre-flare (coronal)
electronic density (nc ), the half-length of the reconnected magnetic loop (L), and magnetic fieldstrength (B). The observed correlations are well reproduced if loop lengths are nearly the samefor all classes, regardless of the existence of an accretion disk. The estimated loop length is almost
comparable to the typical stellar radius of these objects (1010 –1011 cm), which indicates that theobserved flares are triggered by solar-type loops rather than larger ones (∼1012 cm) connecting thestar with its inner accretion disk. The higher < kT > observed for class I sources may be explainedby a slightly higher magnetic field strength (≈500 G) than for class II–III sources (200–300 G).

Contents

1 Introduction

12 Review of Low-mass Young Stellar Objects

3

2.1

Evolution of Low-mass Stars. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

3

2.2

Molecular Clouds . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

7

2.3

X-ray Observations of Low-mass Stars . . . . . . . . . . . . . . . . . . . . . . . . . .

8

2.3.1

X-rays from low-mass main-sequence stars . . . . . . . . . . . . . . . . . . . .

8

2.3.2

X-rays from low-mass young stellar objects . . . . . . . . . . . . . . . . . . . 10

2.3.3

X-rays from brown dwarfs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 13

3 Review of the ρ Ophiuchi Cloud

15

3.1

The ρ Ophiuchi Molecular Cloud . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15

3.2

Young Stellar Objects in ρ Oph . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17

3.3

Previous X-ray Observations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17

3.4

γ-ray Emission . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 18

4 Instrumentation

21

4.1

Overview of the Chandra X-ray Observatory . . . . . . . . . . . . . . . . . . . . . . . 21

4.2

HRMA – High Resolution Mirror Assembly . . . . . . . . . . . . . . . . . . . . . . . 23

4.3

ACIS – Advanced CCD Imaging Spectrometer . . . . . . . . . . . . . . . . . . . . . 25i

ii

CONTENTS
4.4

HRC, HETG, and LETG . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 30

5 Chandra Observations of the ρ Ophiuchi Cloud

31

5.1

Observations and Data Reduction

5.2

Image and Source Extraction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 32

5.3

Identification and Offset Correction . . . . . . . . . . . . . . . . . . . . . . . . . . . . 33

5.4

Source Classification . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 35

5.5

Timing Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 36

5.6

Spectral Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 36

5.7

The Catalogue of X-ray Sources in ρ Oph . . . . . . . . . . . . . . . . . . . . . . . . 37
5.7.1

Chandra X-ray sources in ρ Oph . . . . . . . . . . . . . . . . . . . . . . . . . 38

5.7.2

Identifications of the X-ray sources . . . . . . . . . . . . . . . . . . . . . . . . 47

5.7.3

Detected X-ray flares . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 54

6 Individual Sources
6.1

6.2

. . . . . . . . . . . . . . . . . . . . . . . . . . . . 31

57

Protostars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 57
6.1.1

A-49 – GY91 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 57

6.1.2

BF-26 – Elias29 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 58

6.1.3

BF-50 – WL6 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 58

6.1.4

BF-61 – YLW15A . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 58

6.1.5

BF-64 – YLW16A . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59

6.1.6

VLA1623 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 61

X-ray Brightest T Tauri Stars (DoAr21, ROXs21, and ROXs31) . . . . . . . . . . . . 62
6.2.1

Previous knowledge . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 62

6.2.2

Long-term observations with Chandra and ASCA . . . . . . . . . . . . . . . . 62

6.2.3

Time variability . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 63

6.2.4

X-ray spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 64

CONTENTS

iii

6.2.5

Temperature structure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 68

6.2.6

Chemical composition of X-ray emitting plasma – the FIP and IFIP effects . 69

6.2.7

Giant flare from ROXs31 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70

6.3

6.4

6.5

6.6

Other T Tauri Stars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 72
6.3.1

A-20 – GY12 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 72

6.3.2

A-23 – GY21 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 73

6.3.3

A-24 – S2 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 73

6.3.4

A-34 – GY51 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 74

6.3.5

A-41 – S1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 74

6.3.6

A-79/BF-17 – GY195 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 76

6.3.7

BF-10 – WL22 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 77

Brown Dwarfs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 78
6.4.1

X-ray detection from brown dwarfs in ρ Oph . . . . . . . . . . . . . . . . . . 78

6.4.2

X-ray spectra and luminosities . . . . . . . . . . . . . . . . . . . . . . . . . . 80

6.4.3

Time variability . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 80

6.4.4

X-ray detection rate . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 81

New YSO Candidates – A-29, A-31, and BF-36 . . . . . . . . . . . . . . . . . . . . . 82
6.5.1

Source selection and X-ray images . . . . . . . . . . . . . . . . . . . . . . . . 82

6.5.2

Time variability . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83

6.5.3

X-ray spectra and luminosities . . . . . . . . . . . . . . . . . . . . . . . . . . 84

6.5.4

Possible origin of A-29, A-31, and BF-36 . . . . . . . . . . . . . . . . . . . . . 85

6.5.5

Ionization effect to circumstellar materials . . . . . . . . . . . . . . . . . . . . 87

Unidentified Sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 89
6.6.1

Overall feature . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 89

6.6.2

A-48 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 91

6.6.3

A-75/BF-11 – LFAM21 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 92

iv

CONTENTS
6.6.4

BF-92 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 93

6.6.5

BF-S7 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 93

6.6.6

BF-H3 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 94

7 Overall Feature of X-ray Emission from YSOs

95

7.1

X-ray Detection Rate . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95

7.2

Luminosity Function . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 97

7.3

Distribution of X-ray Parameters . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 98

7.4

Correlation between X-ray and Other Physical Parameters . . . . . . . . . . . . . . . 102
7.4.1

NH vs AV . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 102

7.4.2< LX > vs Lbol . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 103

7.4.3< LX > vs Mass . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 103

7.4.4< LX > vs LR . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 104

8 Systematic Study of YSO Flares

107

8.1

Flare Rate . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 107

8.2

Correlation between the Flare X-ray Parameters . . . . . . . . . . . . . . . . . . . . 109

8.3

Magnetic Reconnection Model . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110
8.3.1

τr vs τd – Implication of the pre-flare density . . . . . . . . . . . . . . . . . . 111

8.3.2< kT > vs τr – Loop length and magnetic field strength . . . . . . . . . . . . 111

8.4

Comparison with the kT -EM scaling law . . . . . . . . . . . . . . . . . . . . . . . . 112

8.5

Effect of the Quiescent X-rays on the Flare Analysis . . . . . . . . . . . . . . . . . . 113

8.6

Evolution of YSOs and Their Flare Activity . . . . . . . . . . . . . . . . . . . . . . . 115

8.7

Comment on the Giant Flares and Quiescent X-rays . . . . . . . . . . . . . . . . . . 116

9 Conclusion

117

A Flare Light Curves

119CONTENTSv

B Physical Parameters of YSOs

125

C Modeling of the Flare131

C.1 Estimation of the Flare Parameters . . . . . . . . . . . . . . . . . . . . . . . . . . . . 131
C.2 Predicted Correlations between the Flare Parameters . . . . . . . . . . . . . . . . . . 133
C.3 The kT –EM Scaling Law . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 133
C.4 Possible Error for the Derived Parameters . . . . . . . . . . . . . . . . . . . . . . . . 134

List of Figures
2.1

The H-R diagram . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

4

2.2

SEDs for low-mass YSOs

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

6

2.3

Molecular cloud map in our Galaxy . . . . . . . . . . . . . . . . . . . . . . . . . . . .

7

2.4

The distribution of the Galactic molecular clouds . . . . . . . . . . . . . . . . . . . .

7

2.5

Multiwavelength light curves of a solar flare . . . . . . . . . . . . . . . . . . . . . . .

9

2.6

A Yohkoh SXT image of the solar flare occurred on February 21, 1992 and schematic
view of the magnetic reconnection model . . . . . . . . . . . . . . . . . . . . . . . . .

9

2.7

Schematic view of the α-ω dynamo mechanism . . . . . . . . . . . . . . . . . . . . . 10

2.8

Schematic view of the star-disk arcade model and quasi-periodic flares from YLW15A 11

2.9

HST and Chandra images of HH2 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 12

3.1

1.3 mm radio continuum map of the ρ Oph main body region . . . . . . . . . . . . . 16

3.2

850 µm map of the ρ Oph main body region . . . . . . . . . . . . . . . . . . . . . . . 16

3.3

ASCA/SIS images of ρ Oph . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 18

4.1

Schematic view of Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 21

4.2

The HRMA mirrors . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 23

4.3

Effective area of HRMA . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 24

4.4

The shape of PSF and fraction of encircled energy of HRMA at the on-axis position

24

4.5

The shape of PSF and fraction of encircled energy of HRMA at the off-axis position

25

4.6

The ACIS array . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 25
vii

viii

LIST OF FIGURES
4.7

Energy resolution of ACIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 26

4.8

Quantum efficiency and effective area of ACIS . . . . . . . . . . . . . . . . . . . . . . 27

4.9

Schematic for determining the ACIS grade . . . . . . . . . . . . . . . . . . . . . . . . 28

5.1

X-ray false-color image of the ρ Oph cloud . . . . . . . . . . . . . . . . . . . . . . . . 32

5.2

Offset between the Chandra and 2MASS sources (before the offset correction) . . . . 34

5.3

Offset between the Chandra and 2MASS sources (after the offset correction) plottedagainst their off-axis angle . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 35

5.4

Histogram of the best-fit abundance for bright sources . . . . . . . . . . . . . . . . . 37

6.1

ACIS image around A-49 = GY91 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 57

6.2

Light curve and spectrum (during the second flare) of BF-64 = YLW16A . . . . . . 60

6.3

Chandra light curves of DoAr21, ROXs21, and ROXs31 . . . . . . . . . . . . . . . . 64

6.4

ASCA light curves of DoAr21, ROXs21, and ROXs31 . . . . . . . . . . . . . . . . . . 64

6.5

Spectra of DoAr21, ROXs21, and ROXs31 . . . . . . . . . . . . . . . . . . . . . . . . 65

6.6

Elemental abundances of DoAr21, ROXs21, and ROXs31 . . . . . . . . . . . . . . . 65

6.7

Close-up view of the spectrum of ROXs21 around 1.2 keV . . . . . . . . . . . . . . . 67

6.8

X-ray luminosity plotted against (a) the plasma temperature and (b) the mean abundance of DoAr21 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 68

6.9Light curve and time evolution of the spectral parameters during the giant flare of
ROXs31 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 71

6.10 Light curve and time evolution of kT and EM of GY12 in obs-BF . . . . . . . . . . 72
6.11 Time-averaged spectrum of A-23 = GY21 . . . . . . . . . . . . . . . . . . . . . . . . 73
6.12 Time-averaged spectrum of A-24 = S2 . . . . . . . . . . . . . . . . . . . . . . . . . . 73
6.13 Time-averaged spectra of S1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 75
6.14 Light curve and time evolution of kT and EM of A-79 = GY195 . . . . . . . . . . . 77
6.15 Spectrum of BF10 = WL22 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 77
6.16 Spectra of BD and BDc . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 80

LIST OF FIGURES

ix

6.17 Light curves of BD and BDc . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 81
6.18 ACIS images around SM1 and B1-MM4 . . . . . . . . . . . . . . . . . . . . . . . . . 82
6.19 Light curves of A-29, A-31, and BF-36 in the 2.0–9.0 keV band . . . . . . . . . . . . 83
6.20 Spectra of A-29, A-31, and BF-36 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 85
6.21 12 CO (J=1–0) map around VLA1623. . . . . . . . . . . . . . . . . . . . . . . . . . 86

6.22 The relation between the X-ray ionization rate and the distance from A-29, A-31,
and BF-36 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 88
6.23 Composite spectra of unidentified sources . . . . . . . . . . . . . . . . . . . . . . . . 89
6.24 Spatial distribution of unidentified sources . . . . . . . . . . . . . . . . . . . . . . . . 90
6.25 ACIS image around the ρ Oph core A region in 0.5–2.0 keV . . . . . . . . . . . . . . 91
6.26 ACIS image around A-48 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 92
6.27 Spectra of LFAM21 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 92
6.28 ACIS image around BF-92 overlaid on the H2 contour map . . . . . . . . . . . . . . 93
6.29 Close-up ACIS image around BF-S7 and CRBR85 . . . . . . . . . . . . . . . . . . . 94
7.1

Relation between Lbol and AV. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 96

7.2

Normalized X-ray luminosity functions . . . . . . . . . . . . . . . . . . . . . . . . . . 97

7.3

Histogram of NH . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 100

7.4

Histograms of < kT > . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 100

7.5

Histograms of τr and τd . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 101

7.6

Relation between NH and AV . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 102

7.7

Relation between the bolometric and X-ray luminosities . . . . . . . . . . . . . . . . 103

7.8

Relation between the stellar mass and X-ray luminosity . . . . . . . . . . . . . . . . 104

7.9

Relation between the radio and X-ray luminosities . . . . . . . . . . . . . . . . . . . 105

8.1

Relation between the rise and decay timescales of flares . . . . . . . . . . . . . . . . 109

8.2

Relation between the temperature and rise timescale of flares . . . . . . . . . . . . . 110

x

LIST OF FIGURES
8.3

Plot of < kT > and < EM > of flares . . . . . . . . . . . . . . . . . . . . . . . . . . 113

8.4

Histograms of < kT > for the flares, when the quiescent X-rays are used as background spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 114

8.5

Relation between < kT > and τr , when the quiescent X-rays are used as background
spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 114

8.6

Schematic view of the evolution of X-ray flares of low-mass objects . . . . . . . . . . 115

A.1 Flare light curves of class I sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . 119
A.2 Flare light curves of class II sources. . . . . . . . . . . . . . . . . . . . . . . . . . . 120

A.3 Flare light curves of class III+IIIc sources . . . . . . . . . . . . . . . . . . . . . . . . 122
A.4 Flare light curves of unclassified NIR and unidentified sources . . . . . . . . . . . . . 123

List of Tables
3.1

Multiwavelength source catalogues in the ρ Oph cloud . . . . . . . . . . . . . . . . . 19

4.1

Basic characteristics of the Chandra scientific instruments . . . . . . . . . . . . . . . 22

4.2

Relation between the ASCA/SIS and ACIS grades . . . . . . . . . . . . . . . . . . . 28

4.3

Telemetry saturation limits of ACIS . . . . . . . . . . . . . . . . . . . . . . . . . . . 30

5.1

Log of the Chandra ACIS-I observations of the ρ Oph main body

5.2

Chandra X-ray sources in the ρ Oph region . . . . . . . . . . . . . . . . . . . . . . . 38

5.3

Identifications of the X-ray sources . . . . . . . . . . . . . . . . . . . . . . . . . . . . 47

5.4

Detected X-ray flares

6.1

Best-fit parameters of BF64 = YLW16A

6.2

ASCA/GIS observation log . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 63

6.3

Best-fit parameters of DoAr21, ROXs21, and ROXs31 . . . . . . . . . . . . . . . . . 66

6.4

Best-fit parameters of S1

6.5

Bona-fide and candidate brown dwarfs in the ρ Oph cloud cores

6.6

Spectral parameters of A-29, A-31, and BF-36

6.7

Best-fit parameters of the composite spectra of unidentified sources

6.8

Best-fit parameters of LFAM21 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 93

7.1

The number of X-ray sources and detection rate for each class. . . . . . . . . . 31

. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 54
. . . . . . . . . . . . . . . . . . . . . . . . 61

. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 76

xi

. . . . . . . . . . . 79

. . . . . . . . . . . . . . . . . . . . . 84
. . . . . . . . . 91

. . . . . . . . . . . . 96

xii

LIST OF TABLES
7.2

Results of two-sample tests for the differences of the observed parameters amongclasses . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 101

8.1

Flare rate sorted by several physical parameters

8.2

The estimated mean values of the flare physical parameters . . . . . . . . . . . . . . 112. . . . . . . . . . . . . . . . . . . . 108

B.1 Physical parameters of YSOs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 125

Chapter 1

Introduction
Star formation process is one of the most important issues in astronomy. Young stellar objects(YSOs) are born in dense molecular cloud cores and evolve through the Hayashi and Henyey tracks
in the Hertzsprung-Russell (H-R) diagram to the main-sequence stage. However, observational information about YSOs is still limited because the dense circumstellar envelope absorbs and scattersinfrared–optical photons emitted from YSOs.
Recent observations have discovered X-rays from many low-mass YSOs (Feigelson & Montmerle, 1999, for a review). Their characteristics of high temperature spectra (∼1 keV) and strong
time variability with occasional rapid flares are basically the same as the solar X-rays, hence their
origin would be solar-type activity attributable to magnetic dynamo processes. These pioneering
discoveries demonstrate the unique capability of X-ray observations to probe YSOs themselves
through the circumstellar envelope, using their high penetrating power.
Thus, the next step should be a systematic X-ray study of YSOs and to approach physicalconditions of these initial stages of stars, which gives extremely important information to the starformation theory. The ρ Ophiuchi cloud (ρ Oph) is the best site for this purpose, because it isamong the closest star-forming regions (distance = 145 pc1 ), well studied in the radio–infrared
bands, and contains a wide variety of samples in different physical conditions (age and mass). The
previous X-ray satellites such as Einstein, ROSAT, and ASCA detected X-rays from some low-mass
YSOs in ρ Oph (Montmerle et al., 1983; Koyama et al., 1994; Casanova et al., 1995; Kamata et al.,
1997; Grosso et al., 2000). However, those samples are too limited because of low sensitivity at thehard X-ray band (>2 keV) and/or poor spatial resolution. The Chandra X-ray Observatory has
the wide energy band sensitivity (0.5–9.0 keV) coupled with the unprecedented spatial resolution(∼0.00 5), hence is the best observatory for X-ray studies of YSOs.
1

The distance corresponding to a parallax angle of 100 . 1 pc = 3.1×1018 cm.

1

2

CHAPTER 1. INTRODUCTION

In this thesis, we report results of two deep exposure Chandra observations of ρ Oph. We first
give brief reviews of low-mass YSOs and ρ Oph in chapters 2 and 3. Details of the Chandra satellite
and its instruments are shown in chapter 4. Chapter 5 describes the method of data analysis andgives source catalogues. Comments on some interesting sources are shown in chapter 6. Then we
discuss overall X-ray properties of YSOs in chapter 7. We further put particular emphasis on theirflare activity and predict the evolution of X-ray activity on YSOs as shown in chapter 8. Finally,
chapter 9 gives a brief summary.

Chapter 2

Review of Low-mass Young Stellar
Objects
2.1

Evolution of Low-mass Stars

The evolution of low-mass stars (mass = 0.1–1.5 M¯ 1 ) is well traced by the locus on the HertzsprungRussell (H-R) diagram (Figure 2.1: Hayashi, 1966). Through a dynamical collapse, interstellar
medium starts condensing. The density soon becomes so large that the shock wave emerging at
its center propagates towards the outer region, which is followed by a dynamical equilibrium statewith the timescale of ∼104 yr. This phase is called “protostars”, the first stage of stars. Protostars are born at the top-right side in the H-R diagram (higher bolometric luminosity and lower
surface temperature), although we can not see the central sources because of large extinction atdense circumstellar envelope. The main mechanism of energy transportation in the stellar interior
is convection (fully convective structure). They evolve along the track of reducing the luminositywith nearly constant temperature (Hayashi track). While moving down on this track, the envelope
gradually disappears and the central sources become optically visible (∼106 yr), which is called“T Tauri stars (TTSs)”. After 107 yr later, the radiative energy transportation becomes dominant
and the temperature gradually increases with a constant luminosity (Henyey track: Henyey, Lelevier, & Levée, 1955). During the hydrogen burning (∼108 yr), the stars stay at a certain point in
the H-R diagram, which is called “main-sequence stars (MSs)”. Protostars and TTSs are called“pre-main-sequence stars (PMSs)”, or “young stellar objects (YSOs)”.

1

The stellar mass normalized by the solar value. 1 M¯ = 2.0×1033 g.

3

4

CHAPTER 2. REVIEW OF LOW-MASS YOUNG STELLAR OBJECTST Tauri Stars
Protostars
Young
Brown Dwarfs

Fig. 2.1.— The H-R diagram (D’Antona & Mazzitelli, 1997). The horizontal and vertical axes are
the surface (effective) temperature and bolometric luminosity normalized by the solar value (1 L¯= 3.8×1033 ergs s−1 ). Thin solid and dashed lines represent the evolutional tracks (Hayashi andHenyey tracks) with the mass of 0.02–3.0 M¯ and isochrones with the ages of 106 and 107 yrs, while
thick solid and dashed lines are the main-sequence track and the birth line, respectively. Ellipses
show the approximate positions of protostars, T Tauri stars, and young brown dwarfs.
Low-mass YSOs are classified into four types based on their spectral energy distributions(SEDs) from the near-infrared (NIR) to sub-millimeter (sub-mm) wavelengths (Figure 2.2; Lada
1991). Class 0 sources (André, Ward-Thompson, & Barsony, 1993; Barsony, 1994) are at the earlyphase of protostars (t ∼104 yr). Although the central source is hidden by the dense material, we canobserve emission from cold (∼30 K) circumstellar envelope at the far-infrared (FIR) to millimeter(mm) bands. In the class 0 phase, much angular momentum is transferred into the central source
through dynamical mass accretion. Part of them escape outward via bipolar outflows, which is
confirmed by some observational evidence such as molecular (CO, H2 , etc.) and atomic (HerbigHaro object) line emissions with a scale of .104 AU2 and centimeter (cm) continuum emissions
2

The average distance from the earth to the sun (astronomical unit). 1 AU = 1.5×1013 cm.

2.1. EVOLUTION OF LOW-MASS STARS

5

which may be produced by a shock of the jet from a putative protostar on the interstellar medium(e.g., Rodriguez et al., 1989, and references theirin). About 105 yr later, most of the materials haveaccreted onto the central source and the emission from the accretion disk reprocessing incident
starlight starts to be observed from the NIR to mid-infrared (MIR) bands, which is called class I
sources (late phase of protostars). They still show outflow activity. Class II sources are at the earlyphase of TTSs (classical T Tauri stars = CTTSs). Although the emission from the central source
becomes dominant, they show excess emissions from K (2.2 µm) to MIR bands and strong opticalHα lines, indicating the presence of an accretion disk. Soon the disk disappears and the SEDs
become a single temperature blackbody. We call them class III sources, or weak-lined T Tauri stars(WTTSs).
If sources have a mass below ≈0.08 M¯ , dynamical contraction stops before the hydrogenburning starts, hence they never arrive at the MS stage, which is called “brown dwarfs (BDs)”(Basri, 2000). They emit faint optical–IR photons using gravitational energy and consequently are
relatively bright at a younger phase (Figure 2.1). Hereafter, we call those BDs still in the masscontraction phase “young BDs” (i.e., on the Hayashi and Henyey tracks, Figure 2.1).
The evolution of high- and intermediate-mass stars (>1.5 M¯ ) would be the same as low-massones; they also evolve along the Hayashi and Henyey tracks. However, there are very few samples ofhigh/intermediate-mass YSOs (e.g., Waters & Waelkens, 1998; Howard, Pipher, & Forrest, 1994),
because they quickly start hydrogen burning then reach the MSs during the mass contraction phase.
Detailed understanding of these classes therefore has not been obtained. In this thesis, we basicallydo not discuss them. For the results of X-ray observations of high/intermediate-mass YSOs, see
Hamaguchi et al. (2000); Hamaguchi (2001); Kohno et al. (2002); Takagi et al. (2002); Nakajima
et al. (2003).

6

CHAPTER 2. REVIEW OF LOW-MASS YOUNG STELLAR OBJECTS

Fig. 2.2.— Evolutionary sequence of the spectral energy distributions for low-mass YSOs (André,
1994; Bachiller, 1996).

2.2. MOLECULAR CLOUDS

2.2

7

Molecular Clouds

Most of nearby YSOs are clustered at several regions in our Galaxy (mainly on the Galactic plane)
with a size of 1–200 pc. These regions are generally invisible in the optical light but show strong
radio continuum (mm) and atomic line (CO, HCO+ , etc.) emissions, hence are called “dark clouds”
or “molecular clouds” (Figure 2.3: Dame, Hartmann, & Thaddeus, 2001). Molecular clouds containa large amount of cold gas and dust with a mass of 103 –106 M¯ . In the smaller scale, there areseveral “molecular cloud cores” with a size and density of 103 –104 AU and 104 –107 cm−3 . YSOs areborn in those molecular cloud cores. The molecular clouds are therefore often called “star-forming
regions”. Many low-mass star-forming regions exist near the solar system (50–200 pc), such asρ Ophiuchi, Taurus, Chamaeleon, and Lupus. On the other hand, high- and intermediate-massstar-forming regions (Orion, Monoceros R2, etc.) are generally distant (0.5–1 kpc; Figure 2.4).

Fig. 2.3.— Molecular cloud map in our Galaxy obtained with velocity-integrated CO intensity(Dame et al., 2001). A dashed ellipse represents the position of the ρ Ophiuchi cloud.

Fig. 2.4.— The distribution of the Galactic
molecular clouds (Dame et al., 1987). The
circle radius and shading indicate the cloudmass and angular distance from the Galactic
plane, respectively.

8

CHAPTER 2. REVIEW OF LOW-MASS YOUNG STELLAR OBJECTS

2.3

X-ray Observations of Low-mass Stars

2.3.1

X-rays from low-mass main-sequence stars

Low-mass MSs emit significant X-rays. The most famous sample of this class is, of course, our sun.
There are two origins of the solar X-rays as follows.
Flare
A “solar flare” means a phenomenon in which a large amount of energy (1029 –1032 ergs) suddenly
arises in a surface magnetic loop with timescales from several minutes to hours. First, particles areaccelerated at the loop top, then move down along the loop and penetrate into a cold chromosphere.
During this phase, the non-thermal particles emit radio (synchrotron) and hard X-ray – γ-ray(bremsstrahlung and nuclear line emission) photons. Cold gas at the chromosphere is abruptly
heated to high temperature, then expands and flows up into the loop, which is called “chromosphericevaporation”. High temperature (0.1–1 keV) and density (∼1010 cm−3 ) plasma fills the loop, thenemits thermal X-rays. The plasma gradually cools down due to radiative and conductive energy
losses, hence X-ray emission disappears about several hours later (Figure 2.5).
It is obvious that the flare activity has a tight correlation with magnetic field. However, themechanism to release such large magnetic energy with a short timescale had been a long-standing
issue. In 1990’s, the Japanese solar X-ray satellite Yohkoh brought us key observational evidence.
Figure 2.6 (left) shows a Yohkoh SXT (soft X-ray telescope) image of the solar flare occurred on
February 21, 1992 (Tsuneta et al., 1992). This shows a beautiful cusp-like shape, which is in good
agreement with the magnetic reconnection models (Sweet, 1958; Parker, 1957; Petschek, 1964, seealso Figure 2.6 right); when a pair of anti-parallel lines of magnetic field crosses each other at theloop top, a reconfiguration of the magnetic field occurs and a large amount of magnetic energy issuddenly released. Now it is widely accepted that magnetic reconnections trigger the solar flares.
Another important issue is the origin of surface magnetic field. The main mechanism for thesun is the standard α-ω dynamo (Parker, 1955, 1970, see also Figure 2.7); poloidal magnetic field
makes toroidal field with the effect of differential rotation in a surface convection zone (ω term),and vice versa with the effect of turbulent motion (α term). Theories predict that the boundary
between the surface convection zone and the radiative core acts as the anchor point of magnetic
field, hence the α-ω dynamo mechanism may not work for sources with fully convective structuresuch as protostars and brown dwarfs, although a small-scale turbulent dynamo (α2 dynamo) would
be conceivable (Durney et al., 1993).

2.3. X-RAY OBSERVATIONS OF LOW-MASS STARS

9

Fig. 2.5.— Multiwavelength light curves of the solar flare occurred on August 17, 1994 (Silva et al.,1996). The top panel shows those in the soft (SXT: 0.25–4.0 keV) and hard (HXT LO: 14–23 keV,
M1: 23–33 keV) X-ray bands obtained with Yohkoh, while the lower panel shows those in the radio
17 (NRO) and 86 (BIMA) GHz bands.

Fig. 2.6.— (left) A Yohkoh SXT image of the solar flare occurred on February 21, 1992 (Tsuneta
et al., 1992) and (right) schematic view of the magnetic reconnection model (Yokoyama & Shibata,2001).

10

CHAPTER 2. REVIEW OF LOW-MASS YOUNG STELLAR OBJECTS

Fig. 2.7.— Schematic view of the α-ω dynamo mechanism (Parker, 1970). Solid line
and ribbon arrows represent a magnetic lineof force and fluid motion, respectively.

CoronaThe sun is covered by hot corona, which also emits X-rays. However, its X-ray properties are quite
different from the solar flare X-rays; steady emission from relatively low temperature (∼0.1 keV) anddensity (∼109 cm−3 ) plasma. Although the heating mechanism of the corona is still controversial,
it is widely accepted that the magnetic field also plays a key role.

2.3.2

X-rays from low-mass young stellar objects

X-rays from T Tauri stars
In 1980’s, the Einstein and ROSAT satellites discovered X-ray emission from many TTSs (Feigelson& DeCampli, 1981; Feigelson & Kriss, 1981; Montmerle et al., 1983; Feigelson & Montmerle, 1999)
using the imaging capability at the soft X-ray band (<2 keV). Their X-rays are highly variablewith timescales of days or longer and occasionally show flare activities like the sun. X-ray spectra
are thermal in origin and are well fitted by a single- or multi-temperature plasma model withtemperatures of 0.1–3 keV (Preibisch, 1997). These characteristics are quite similar to the solar
X-rays (§2.3.1), hence the origin of X-rays from TTSs would be the same as that of the sun, i.e.,magnetic activity at the stellar surface. The X-ray luminosity (LX ) of TTSs is, however, muchlarger than that of the sun; 1029 –1031 ergs s−1 and 1030 –1032 ergs s−1 for the quiescent and flarephases, respectively (cf., ∼5×1026 ergs s−1 for the solar corona and 1027 –1031 ergs s−1 for thesolar flares: Vaiana & Rosner, 1978). Also, LX of TTSs is well correlated with their bolometricluminosity, stellar mass (Preibisch & Zinnecker, 2002), and possibly with their rotation period(Stelzer et al., 2000), although the underlying physics is still controversial.

2.3. X-RAY OBSERVATIONS OF LOW-MASS STARS

11

X-rays from class I protostars

Fig. 2.8.— (left) Schematic view of the star-disk arcade model (Montmerle et al., 2000). Numerals(1–5), solid lines, and thick gray line indicate time sequence, magnetic field, and flare loop aftera reconnection. An arrow shows a reconnection point. (right) Quasi-periodic flares from a class I
YLW15A in the ρ Ophiuchi cloud obtained with ASCA (Tsuboi et al., 2000).
Protostars are deeply embedded in molecular cloud cores and shrouded by dense circumstellar
gas and dust, hence we can not observe them even with the soft X-ray band due to large extinction.
The Japanese X-ray satellites Temma and Ginga, however, discovered hard X-ray emission (>2keV) in the direction of star-forming regions (Koyama, 1987; Koyama et al., 1992). Although
they have no imaging capability, X-ray characteristics of flare-like light curves and highly absorbed
spectra (NH ∼ 1022 cm−2 ) indicate that X-rays are coming from discrete sources inside (or behind)the dense cloud cores.
The forth Japanese X-ray satellite ASCA (Advanced Satellite for Cosmology and Astrophysics:
Tanaka, Inoue, & Holt, 1994) is the first X-ray mission to have the X-ray telescopes (XRTs: Serlemitsos et al., 1995) with the imaging capability for the hard X-ray photons (2–10 keV), which−2.5
are less absorbed as the extinction cross-section decreases as EX(EX : X-ray energy). Koyama
et al. (1994) first revealed that the hard X-rays from the ρ Ophiuchi dark cloud regions would
come from class I sources (WL6 and WL15 = Elias29) with ASCA/GIS. Using ASCA/SIS, which
has better spatial resolution than GIS, Kamata et al. (1997) confirmed the X-ray emission fromthese class I sources. Independently, Casanova et al. (1995) obtained the hints of X-ray detection
from class I in the same star-forming area with ROSAT/PSPC. Then the successive observations
with ASCA further revealed that many class I protostars are X-ray emitters (Koyama et al., 1996;
Ozawa et al., 1999; Tsuboi et al., 2000) thanks to its higher sensitivity at the hard X-ray band.
ROSAT observations also detected X-rays from class I (Grosso et al., 1997; Neuhäuser & Preibisch,
1997; Grosso, 2001). They share the same characteristics of thermal emission with the plasma

12

CHAPTER 2. REVIEW OF LOW-MASS YOUNG STELLAR OBJECTS

temperature of 0.5–5 keV and strong variability with occasional rapid flares, which are consistent
with the scenario of enhanced solar-type activity, attributable to magnetic dynamo processes.
Class I protostars have fully convective structure, so the standard (α-ω) dynamo mechanism
may not work (§2.3.1). As an alternative scenario, Hayashi, Shibata, & Matsumoto (1996) andMontmerle et al. (2000) have proposed a scenario of star-disk connection by one magnetic loop,
where the central forming-star rotates faster than the inner edge of its accretion disk (Figure 2.8left). This well explains characteristic quasi-periodic three flares from a class I YLW15A (Figure
2.8 right: Tsuboi et al., 2000). If this is a real case, the rotation period of YLW15A is surprisingly
high of ∼ a day (e.g., ∼28 day for the sun and 3–10 day for TTSs).
X-rays from class 0 protostars
Recently, the Chandra and XMM-Newton satellites discovered X-rays from some class 0 systems (or candidates),
MMS 2 and 3 in the Orion Molecular Cloud 3 (OMC 3:
Tsuboi et al., 2001), IRS5 in L1551 (Favata et al., 2002)
and HH2 in the Orion Nebula (Pravdo et al., 2001). Their
X-ray properties, however, are different from those of class
I–IIIs; (1) there is no significant X-ray variability, and (2)X-ray shapes of IRS5 and HH2 seem to be diffuse and are
not located at the central source but near the HerbigHaro objects, which are shock induced optical emissions(Figure 2.9). As for MMS 2, Tsujimoto et al. (2002a)
and Tsujimoto (2003) made follow-up observations with
the Subaru and VLA telescopes and found that the Xray emission is associated with an H2 outflow of about
500 AU-long emanating from the VLA source. Their Xrays therefore may not be attributable to central sources
with the solar-like magnetic activity but are probably
associated with the violent outflows of more than 200–
1000 km s−1 (Pravdo et al., 2001; Favata et al., 2002; Tsujimoto, 2003), although the limited samples and statistics Fig. 2.9.— (top) HST and (bottom)
Chandra images of HH2 (Pravdo et al.,hamper further analyses and discussions.
2001). The contour in the bottom panel
is the same as the image in the top
panel.

2.3. X-RAY OBSERVATIONS OF LOW-MASS STARS

2.3.3

13

X-rays from brown dwarfs

ROSAT detected X-rays from several young BDs and their candidates in star-forming regions;
Chamaeleon, ρ Ophiuchi, and Taurus (Neuhäuser & Comrerón, 1998; Neuhäuser et al., 1999).
Also, an old BD LP944−20 showed a weak X-ray flare discovered with Chandra (Rutledge et al.,2000). Their X-rays are similar to those of low-mass stars and are likely to be magnetic origin.
However, the standard (α-ω) dynamo to produce the magnetic activity may not present, because
BDs are fully convective like class 0–I protostars.

Chapter 3

Review of the ρ Ophiuchi Cloud
3.1

The ρ Ophiuchi Molecular Cloud

The ρ Ophiuchi molecular cloud (hereafter, ρ Oph) is one of the nearest and most well-known
star-forming regions in our Galaxy, which is located about 15◦ above the Galactic center (Figure
2.3).
The distance to ρ Oph is somewhat uncertain. Chini (1981) made optical–infrared photometricobservations and estimated the distance to be 165±20 pc from the mean distance modulus of threebright and heavily reddened stars. Alternatively, Knude & Høg (1998) calculated optical reddening(EB−V ) of MS stars and giants using the Hipparcos and T ycho catalogues, then found an abrupt
increase of EB−V for stars with the distance larger than 120 pc. Recently, however, de Zeeuw etal. (1999) proposed the distance of 145±2 pc from the Hipparcos positions, proper motions, andparallaxes. In this thesis, we adopt the distance of 145 pc.
The total size of the ρ Oph complex is ∼10◦ × 20◦ (∼= 25 pc × 50 pc). The ongoing starformation is less active in the north-east region but is concentrated in the south-west region withthe size of ∼1 pc × 2 pc called “ρ Oph main body”. The radio observations of the ρ Oph main body
region have been frequently made since 1980’s. The C18 O J=1–0 (Wilking & Lada, 1983; Tachihara,
Mizuno, & Fukui, 2000) and 13 CO J=1–0 (Loren, 1989) mappings revealed the distribution of low
density molecular gas and were used to estimate the total mass and average density of the main
body to be 500–1000 M¯ and 104 –105 cm−3 , respectively. Loren, Wootten, & Wilking (1990)
observed higher density gas with DCO+ emission lines and divided the main body into sevensmaller scale cores called ρ Oph A, B1, B2, C, D, E, and F. The scale, mass, and density of thesecores are ∼0.1 pc, 10–100 M¯ , and 105 –106 cm−3 , respectively. The millimeter-submillimeter radio
continuum emission (850 µm: Johnstone et al. 2000; 1.3 mm: Motte et al. 1998; see also Figures
15

16

CHAPTER 3. REVIEW OF THE ρ OPHIUCHI CLOUD

3.1, 3.2, and Table 3.1) revealed further smaller scale structure and detected about 60 clumps withthe scale and mass of ∼5000 AU (∼=0.02 pc) and 0.1–6 M¯ . Within these clumps, new YSOs are
generated.

Fig. 3.1.— 1.3 mm radio continuum map of the ρ Oph main body region (Motte et al., 1998).
Squares are the fields of view of Chandra observations (chapter 5).

Fig. 3.2.— 850 µm map of the ρ Oph main body region (Johnstone et al., 2000). Circles show thedetected clumps. Squares are the fields of view of Chandra observations (chapter 5).

3.2. YOUNG STELLAR OBJECTS IN ρ OPH

3.2

17

Young Stellar Objects in ρ Oph

The search for YSOs in ρ Oph has been made with several instruments (Table 3.1), using the
NIR (Grasdalen, Strom, & Strom, 1973; Vrba et al., 1975; Elias, 1978; Wilking & Lada, 1983;
Barsony et al., 1989; Greene & Young, 1992; Strom, Kepner, & Strom, 1995; Barsony et al., 1997)
and MIR/FIR (Young, Lada, & Wilking, 1986; Wilking, Lada, & Young, 1989; Bontemps et al.,2001) photometry. Several hundreds of YSOs have been found. Bontemps et al. (2001) madeMIR observations with the ISO satellite and divided IR sources into 16 class I, 123 class II, and77 class III sources, which is based on the spectral index from the NIR to MIR bands. Such asensitive survey at the MIR band has been done for limited star-forming regions, which makes theunambiguous classification of YSOs possible. Luhman & Rieke (1999) obtained K-band spectra ofabout 100 NIR-bright cloud members and estimated the age and mass ranges of the ρ Oph YSOs
to be roughly 0.1–1 Myr and 0.02–3 M¯ .
In addition, there are two class 0 protostars in ρ Oph; VLA1623 and IRAS 16293−2422(André & Montmerle, 1994). They show strong sub-mm emission and CO bipolar outflow (Mundy,
Wootten, & Wilking, 1990; André, Ward-Thompson, & Barsony, 1993), which is firm evidence for
their class 0 nature. Unfortunately, however, IRAS 16293−2422 is located outside of the Chandra
observed field, hence we do not mention it hereafter.

3.3

Previous X-ray Observations

The first X-ray observation of ρ Oph was made with the Einstein satellite (Table 3.1: Montmerle etal., 1983). In the soft X-ray (<2 keV) band image, many TTSs showed X-ray flares. Montmerle etal. (1983) expressed such phenomena as “the X-ray Christmas Tree”. The X-ray spectra were wellfitted by a thin-thermal plasma model with a moderate temperature of kT ∼1 keV and absorptionof NH .1022 cm−2 . The ROSAT satellite further detected soft X-rays from about 50 TTSs andsome class I sources (Casanova et al., 1995; Grosso et al., 2000; Grosso, 2001). However, X-ray
emission from class I was still controversial due to the large positional uncertainty and limited
statistics.
As for the hard X-ray (>2 keV) band, the Japanese Tenma and Ginga satellites detectedsignificant emission toward the ρ Oph cloud (Koyama, 1987; Koyama et al., 1992). Based on thesepioneering results, the ASCA satellite observed ρ Oph and made the first hard X-ray image. Some
class I sources were identified with hard X-ray emitters (Koyama et al., 1994; Kamata et al., 1997;
Tsuboi et al., 2000, Figure 3.3). In particular, a class I source YLW15A showed three quasi-periodic
flares with the interval of ∼1 day (Tsuboi et al., 2000), which may be evidence for the star-disk

18

CHAPTER 3. REVIEW OF THE ρ OPHIUCHI CLOUD

arcade scenario and fast rotation of protostars (Montmerle et al., 2000, see also §2.3.2).

Fig. 3.3.— ASCA/SIS images of the central region of ρ Oph in the 0.7–2 (left) and 2–10 keV (right)
bands (Tsuboi et al., 2000). The contours are C18 O (1–0) density (Wilking & Lada, 1983). Crosses
show the position of class I sources. Arrows indicate X-ray detected class I sources (1: YLW15A,
2: IRS51, 3: Elias29), WTTSs (5: ROXs21, 6: ROXs31), and another class I source detected in
the previous ASCA observation (4: WL6). Circles are source and background regions used for the
spectral analysis in Tsuboi et al. (2000).

3.4

γ-ray Emission

The COS-B satellite found a γ-ray source 2CG353+16 near the ρ Oph cloud (Swanenburg et al.,1981). The position (J2000) is (16h 30m 05s , −24◦ 500 0600 ) with the error radius of 1.5◦ and flux(>100 MeV) of 1.1×10−6 photons cm−2 s−1 . One plausible idea about its emission mechanism isthat γ-rays are made by the interaction between high-energy cosmic ray nucleons and the interstellar
matter in the cloud via pion decay (π 0 → 2γ). Xu et al. (1992) proposed that those high-energy
cosmic rays were generated by the supernova remnant exploded in the 14th century B.C in thiscloud, which was reported in an ancient Chinese record. EGRET further found a γ-ray source
3EG J1627−2419, although Hartman et al. (1999) suggested that it may be a possible artifact.

3.4. γ-RAY EMISSION

19

Table 3.1: Multiwavelength source catalogues in the ρ Oph cloud
Instruments

AbbreviationArea(arcmin2 )Bandpass

VLA
VLA
VLA
VLA
IRAM
SCUBA

ROS
SFAM
ROC
LFAM
···
SMM

14000
14000
2500
340
480
700

1.4 GHz
1.4 GHz
1.4/5 GHz
5 GHz
1.3 mm
850 µm

IRAS
IRAS
ISO

YLW
WLY
···

620
1800
2500

12–100 µm
12–100 µm
6.7/14.3 µm

Resolution(arcsec)

Sensitivity

Numberof Sources

Ref.

5
5
5
0.5
15
3

1 mJy
1 mJy
1 mJy (@1.4 GHz)
0.09 mJy
30 mJy0.1 Jy

58
213
59
35
58*
55

1
2
1
3
4
5

45
45
3

1 Jy (@12 µm)
0.15 Jy (@12 µm)
10 mJy (@6.7 µm)

18
64
425

6
7
8

35
36
120
12
0.78
7.5
1.8
1.3
1.3
1.2
0.1

9 mag
12.0 mag (@K)
7.5 mag (@K)
12.1 mag (@K)
14.2 mag
14 mag (@K)
13 mag (@K)
16 mag (@K)
14.2 mag (@K)
14.5 mag (@K)
15 mag (@Ks)

41
31
50
20
61
58
481
91
119
4495
∼600‡

9
10
11
12
13
7
14
15
16
17
18

···
a few
··· #

···
··· k
··· #

38
86
47

19
20
21

45
15
2
30
0.5

1029.5 ergs s−1
1029 ergs s−1
1029 ergs s−1
1030 ergs s−1
1028 ergs s−1

47
55
63
19
195

22
23
24
25
26

— Radio —— MIR —— NIR —
KPNO
KPNO
CIT
†

KPNO
IRTF
Steward Obs.
Steward Obs.
KPNO
KPNO
2MASS

GSS, S
VSSG
Elias
WL
BBRCG
IRS
GY
CRBR
SKS
BKLT
···

230
720
65000
100
140
2000
650
200
1200
3600
···

K
JHKL
JHKLM N Q
JHKK
JHKLM N Q
JHK
HK
JHK
JHK
JHKs

···
···
ROXs

···
140000··· #U BV RIHα··· #— Optical —§CTIO
La Silla— X-ray —Einstein
ROSAT/PSPCROSAT/HRIASCA
Chandra

ROXs
ROXR1
ROXR
ROXA
···

14000
1200
800
2000
600

0.5–2 keV
0.1–2.4 keV
0.1–2.4 keV
0.5–10 keV
0.5–9 keV20

CHAPTER 3. REVIEW OF THE ρ OPHIUCHI CLOUD(footnotes for Table 3.1)*

The number of starless cores.

†

University of Arizona 1.54 m telescope.

‡

The number of sources within ∼700 arcmin2 region of the ρ Oph main body.

§

The 123 cm telescope in Cala Alto/Spain and the 50 cm telescope on the Gamsberg in South West Africa. No

observational information was given in the paper.
k

Relative Hα intensity was only given.

#

Pointing spectroscopic (3700–7500 Å) and photometric (U BV RI) observations on Einstein X-ray sources (Mont-merle et al., 1983).
References – (1) André et al. (1987), (2) Stine et al. (1988), (3) Leous et al. (1991), (4) Motte et al. (1998), (5) Johnstone et al. (2000), (6) Young et al. (1986), (7) Wilking et al. (1989), (8) Bontemps et al. (2001), (9) Grasdalen
et al. (1973), (10) Vrba et al. (1975), (11) Elias (1978), (12) Wilking & Lada (1983), (13) Barsony et al. (1989),(14) Greene & Young (1992), (15) Comerón et al. (1993), (16) Strom et al. (1995), (17) Barsony et al. (1997),(18) http://ipac.caltech.edu/2mass/, (19) Chini (1981), (20) Wilking et al. (1987), (21) Bouvier & Appenzeller(1992), (22) Montmerle et al. (1983), (23) Casanova et al. (1995), (24) Grosso et al. (2000), (25) Kamata et al.
(1997), (26) This thesis and Imanishi et al. (2003).

Chapter 4

Instrumentation
In this chapter, we show the basic properties of the Chandra satellite and its instruments. Moredetails are found in the Chandra Observatory Guide1 .

cFig. 4.1.— Schematic view of the Chandra satellite (°NASA/CXC/SAO).

4.1

Overview of the Chandra X-ray Observatory

The Chandra X-ray Observatory (CXO: Weisskopf et al., 2002, Figure 4.1) was launched by NASA’s
Space Shuttle Columbia on July 23, 1999. The total weight and length are ∼5,000 kg and ∼10 m.
The orbit is elliptical with the perigee and apogee altitudes of ∼16,800 km and ∼132,000 km,
1

http://asc.harvard.edu/udocs/docs/docs.html .

21

22

CHAPTER 4. INSTRUMENTATION

respectively. The orbital period is ∼63.5 hours, about 75 % of which are used for observation
times.
Chandra is equipped with one large X-ray telescope (HRMA = High Resolution Mirror Assembly; §4.2) and the focal-plane Science Instruments (SIs), consisting of two CCD (Charge CoupledDevice, ACIS = Advanced CCD Imaging Spectrometer; §4.3) and microchannel plate (HRC = High
Resolution Camera; §4.4) arrays. Also two transmission grating systems (HETG = High Energy
Transmission Grating and LETG = Low Energy Transmission Grating; §4.4), which are locatedbetween HRMA and SIs, offer observations with high energy resolution. The satellite pointing iscontrolled by the Pointing Control and Aspect Determination (PCAD) system. Table 4.1 shows
basic characteristics of each instrument. The most superior point is the highest spatial resolution(∼0.00 5) among the X-ray satellites operated by today.

Dithering
Since HRMA gives the excellent spatial resolution (∼0.00 5) comparable to the pixel size of SIs, it
sometimes happens that most X-ray photons fall into a certain bad pixel or the gap between the SI
chips. In order to avoid them, the spacecraft is dithered during observations in a Lissajous shape.
The peak-to-peak amplitude of the dithering is 1600 (ACIS) and 4000 (HRC) and the nominal period
is 1000.0/707.1 s (ACIS) and 1087.0/768.6 s (HRC) in the Y/Z axes, respectively. We can correct
the dither pattern in ground processing with the PCAD data.

Table 4.1: Basic characteristics of the Chandra scientific instruments
Instruments
Field of view . . . . . . . . . . . .

ACIS-I∼170 ×170

ACIS-S∼80 ×500

Bandpass . . . . . . . . . . . . . . . .
On-axis Effective area . . .

0.5–10.0 keV
600 cm2@1.5 keV
0.00 5130 eV†@6.0 keV
3.2 s‡

0.1–10.0 keV*
800 cm2*@1.5 keV
0.00 5200 eV*@6.0 keV
3.2 s‡

On-axis spatial resolution
Energy resolution . . . . . . .
Timing resolution . . . . . . .
*HRC∼300 ×300 (I)∼60 ×990 (S)
0.08–10.0 keV
200 cm2@1 keV
0.00 4
1 keV@1 keV
16 µs

HETG
···

LETG
···

0.4–10 keV
60 cm2@1 keV···
1 eV@1 keV
16 µs

0.08–2 keV
10 cm2@1 keV···0.1 eV@0.1 keV
16µs

Values for the backside illuminated (BI) CCD chips.
Values near the readout node. The energy resolution becomes worse at the further position from the node (seeFigure 4.7 right).
‡
Values for the Timed Exposure mode with full frames.
†4.2. HRMA – HIGH RESOLUTION MIRROR ASSEMBLY

4.2

23

HRMA – High Resolution Mirror Assembly

In spite of their high penetrating power, most X-rays can be reflected if the incident angle is smaller
than a critical value of . 1◦ . HRMA consists of four pairs of Wolter Type-I mirrors, a combinationof parabolic and hypabolic shells, which realize small incident angles (Figure 4.2). The mirrors are
made of Zerodur glass polished and coated with iridium. The diameters of the mirrors are from0.65 to 1.23 m and the focal length is 10.066 m. In order to keep high accuracy of the mirror
surface, the glass is very thick with the total mass of ∼1,500 kg.

Fig.
4.2.— (left) Schematic view of The HRMA mirrors.
c(°NASA/CXC/SAO).

(right) Wolter type-I optics

Effective area
Figure 4.3 shows the HRMA effective area, which is ∼800 cm2 at the photon energy of 1 keV and∼100 cm2 at 8 keV. It gradually decreases as the source off-axis angle becomes larger. These values
are almost comparable to those of ASCA and slightly smaller than those of XMM-Newton.

Point spread function and encircled energy fraction
Figure 4.4 shows the shape of the point spread function (PSF) of HRMA and the fraction of encircled
energy at the on-axis position. The point source image is nearly circular and very sharp. The half
power diameter (HPD) is ∼0.00 5 (at 1.49 keV) and about 90 % of the photons are concentrated in
a 200 diameter. The PSF gets elongated in an elliptical shape as the off-axis angle becomes larger(Figure 4.5), which is due to the aberrations of Wolter type I optics and the different shapes in
focal surfaces among the four mirrors. The HPD thus becomes large as a function of the off-axis

24

CHAPTER 4. INSTRUMENTATION

angle (Figure 4.5 right).

Fig. 4.3.— The effective area of HRMA plotted against (left) X-ray energy and (right) a source
off-axis angle. The sudden drop-off seen around 2 keV (left) is due to the M -edge of iridium used
for the mirror coating. The values multiplied by the quantum efficiency of ACIS and HRC are also
shown in the left panel.

Fig. 4.4.— (left) On-axis PSF of HRMA with the photon energies of (top-left) 1.49, (top-right)
4.51, (bottom-left) 6.40, and (bottom-right) 8.63 keV. (right) The fraction of encircled energy of
HRMA as a function of the angular radius at the on-axis position.

4.3. ACIS – ADVANCED CCD IMAGING SPECTROMETER

25

Fig. 4.5.— (left) Off-axis PSF with the photon energy of 1.49 keV. The difference between the top
and bottom panels is the off-axis of 50 and 100 , while that between the left and right panels arethose of HRMA only and of HRMA/HRC-I, respectively. (right) The fraction of encircled energyof HRMA as a function of source off-axis angle.

4.3

ACIS – Advanced CCD Imaging Spectrometer

ACIS FLIGHT FOCAL PLANE~22 pixels ~11"I0
w203c4r
0

22 pixels= 11"

constant with Z

I1
w193c2
1x

I2
w158c4r}

I3

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

w215c2r

2

3
330 pixels = 163"

S0

S1

S2

w168c4r
4

w140c4r
5

w182c4r
6

S3+
w134c4r
7

S4S5

w457c48

w201c3r
9}

ACIS-S(aimpoint on S3 = (252, 510))
  +∆Y
18 pixels = 8".8.

columnBottomthree

Image Region

two

Top

one

BI chip indicator

node zero

+Z

.

row

OffsetCoordinates+Z
Pointing+∆Z

Coordinates-Z

Pixel (0,0)+Y

Sim Motion

Frame StoreCCD Key

NodeDefinitions

Row/Column
Definition

Coordinate
Orientations

Fig. 4.6.— (left) photographic and (right) schematic view of the ACIS array. The legend of theterminology (right panel) is given in the lower left. The x and cross positions represent the nominal
aimpoints of the ACIS-I and ACIS-S arrays.

26

CHAPTER 4. INSTRUMENTATIONThe ACIS array (Figure 4.6) consists of ten CCDs (Charge Coupled Devices) with 1024 ×
1024 pixels each. ACIS-I (chip I0–I3) is laid out in 2 × 2 and is used for imaging observations,while ACIS-S (chip S0–S6) is in 1 × 6, which is optimized for both imaging and grating-spectroscopic
observations. The pixel size is 24 µm × 24 µm, corresponding to the angular size of 0.00 5 × 0.00 5.
Two CCDs in the ACIS-S array (chip S1 and S3) are back-illuminated (BI) and the others are
front-illuminated (FI). For each observation, we can use any six CCDs simultaneously.

Energy resolution
Figure 4.7 (left) shows the energy resolution (∆E) of ACIS estimated by ground calibrations. The
FI chips have a good resolution near the theoretical limit (∼120 eV @ 5.9 keV), while that of the
BI chips is slightly poor. After the launch, however, the ACIS-I chips were significantly damaged
by low energy protons encountered during radiation belt passages. These protons made many
charge traps at buried channels, which cause the increase of charge transfer inefficiency (CTI).
Consequently, ∆E becomes worse as a function of source positions from the readout nodes; ∆E ∼130 eV (@ 5.9 keV) near the readout nodes, while ∆E ∼ 330 eV (@ 5.9 keV) at the edge of thechip. Since the gate structure of the BI chips (S1 and S3) is at the opposite side to HRMA, ∆E
of the BI chips remains at their pre-launch value and has little dependence on the source position.
Now ACIS is not operated and not left at the focal position during the radiation belt passages,
hence no further degradation of CTI is expected. Furthermore, in order to reduce the effect of CTI,
the CCD temperature is lowered from −90 ◦ C to −120 ◦ C.

Fig. 4.7.— (left) The pre-launch energy resolution of ACIS. (right) The energy resolution of chip
S3 and I3 as a function of row number, after the proton damage on orbit. The data were taken at
the CCD temperature of −120 ◦ C.

4.3. ACIS – ADVANCED CCD IMAGING SPECTROMETER

27

Quantum efficiency and effective area
Figure 4.8 shows the quantum efficiency of ACIS (left) and the combined effective area of HRMA
and ACIS (right). A large fraction of low energy X-rays is absorbed at the optical blocking filter(OBF) and the gate structure for FI chips, hence the efficiency becomes worse at the lower energy.
Since the gate of BI chips is at the opposite side to HRMA, the efficiency of BI chips at the low
energy band is about two times better than that of FI chips. Due to the effect of CTI, the quantum
efficiency at the further side from the readout node is about 5–15 % worse than that near thenode. The energy ranges with the effective area of larger than 10 cm2 are ∼0.5–10.0 keV and∼0.1–10.0 keV for FI and BI chips, respectively.

Fig. 4.8.— (left) The quantum efficiency of ACIS with the transmission of OBF. (right) The
combined effective area of HRMA and ACIS. The sudden drop-off seen around 1.8 keV is due tothe K-edge of silicon.

Event grade
X-ray events are selected with the criterion that the bias-subtracted pixel value exceeds the event
threshold (38 ADU ∼= 140 eV and 20 ADU ∼= 70 eV for FI and BI chips) and is the highest
among the neighboring 9×9 pixels. In order to distinguish real events from background ones and to
estimate the total number of electrons including those leak into the neighboring pixels, the grade
method is used. Its algorithm is basically the same as the ASCA/SIS grade method (Burke et al.,1991), but ACIS gives the specified numbers for each pixel around the event pixel (Figure 4.9) anddetermines the grade by summing up those pixel numbers that have a pixel value larger than the

28

CHAPTER 4. INSTRUMENTATION

split threshold (13 ADU; if the neighboring pixel has larger pixel value than the split threshold,it is regarded as the leaking pixel). Table 4.2 gives the relationship of the ASCA/SIS and ACIS
grading schemes. Usually, we regard the grades corresponding to the ASCA/SIS grades 0, 2, 3, 4,
and 6 as real events.


















Fig. 4.9.— Schematic for determining the ACIS grade. For example,
an event that caused all pixels to exceed their threshold is grade 255.
A single pixel event is grade 0.

Table 4.2: Relation between the ASCA/SIS and ACIS grades
ASCA/SIS grade

ACIS grade

0 ...............
1 ...............
2 ...............
3/4 . . . . . . . . . . . .
5 ...............

0
1, 4, 5, 32, 33, 36, 37, 128, 129, 132, 133, 160, 161, 164, 165
2, 34, 64, 65, 68, 69, 130, 162
8, 12, 16, 17, 48, 49, 136, 140
3, 6, 9, 13, 20, 21, 35, 38, 40, 44, 52, 53, 96, 97, 100, 101, 131, 134,
137, 141, 144, 145, 163, 166, 168, 172, 176, 177, 192, 193, 196, 197
10, 11, 18, 22, 50, 54, 72, 76, 80, 81, 104, 108, 138, 139, 208, 209
All other grades

6 ...............
7 ...............

Operating mode
Two operating modes are available for ACIS. In this thesis, we only use the data of Timed Exposure
mode with full frames and the frame time of 3.2 s.
• Timed Exposure (TE) mode
The 1024 × 1024 active region of the CCD is exposed for a selected time interval (frametime), then the collected charge quickly (∼41 ms) transfers to the framestore region. For theentire CCD (full frames), the frame time is 3.2 s, which is the nominal timing resolution of
ACIS. In order to obtain better timing resolution or to avoid the pileup effect (see below),
we can select subarray mode; the active region is limited to 1/2, 1/4, or 1/8 subarray, which
shortens the frame time to 0.4 s at the minimum.

4.3. ACIS – ADVANCED CCD IMAGING SPECTROMETER

29• Continuous Clocking (CC) mode
CC mode achieves high timing resolution (3 ms) but has only one dimension of spatial resolution. Data are continuously clocked through the CCD and framestore to minimize thereadout time, with one-dimensional (1 × 1024 pixels) images.

Telemetry format
Three telemetry formats are available as follows. Table 4.3 lists the operating modes, associatedformats, and approximate event rates at the telemetry saturation. The faint and very faint formats
were used for obs-BF and obs-A, respectively (§5.1).

• GradedThis provides event positions in detector coordinates, event amplitudes, arrival times, andevent grades.
• Faint
This provides the same information as the Graded format other than the event grades. Alternatively, this supplies all information about 3 × 3 pixels around the event pixel. The bias
map (zero level of each pixel) is telemetered separately.
• Very FaintThis is basically the same as the Faint format, but provides all information about 5 × 5 pixels
around the event pixel. This format is used for the TE mode only and offers the advantage
for minimizing the background after ground processing, but is limited for low-count sources(Table 4.3).

Event pileup
For bright sources, two or more X-ray photons are often registered in a single ACIS pixel within one
readout cycle, which is called pileup. It causes an artificial spectral hardening and underestimateof the X-ray counts. In fact, the brightest two sources in the ρ Ophiuchi cloud (A-2 and BF-64,
see chapter 5) suffered the pileup effect. It can be reduced by selecting a proper observing mode(e.g., shorten the exposure time).

30

CHAPTER 4. INSTRUMENTATION

Table 4.3: Telemetry saturation limits of ACIS

*

4.4

Mode

FormatBits/event

Events/s*

CC
CC
TE
TE
TE

Graded
Faint
GradedFaint
Very Faint

58
128
58
128
320

375.0
170.2
375.0
170.2
68.8

The number of events per 1 s at the telemetry saturation.

HRC, HETG, and LETG

We do not use these three instruments in this thesis, hence only show their basic characteristics.

HRC – High Resolution Camera
HRC is a microchannel plate instrument consisting of two detectors like ACIS; HRC-I for imaging
and HRC-S for imaging and grating spectroscopy. HRC has larger field of view (300 × 300 ) andbetter time resolution (16 µs) than ACIS, although the energy resolution (∆E/E ∼1 @ 1 keV) is
largely limited.

HETG – High Energy Transmission Grating
HETG is designed for grating spectroscopy. The transmission grating inserted between HRMA
and the focal plane detector (ACIS-S and HRC-S) disperses incident X-ray photons. Although thetwo dimensional image is not available, spatial information along the detector can be converted
to the photon energy because the dispersion distance from the aimpoint is a function of energy(wavelength) of X-ray photons, which attains high energy resolution (∆E/E ∼ 1000 @ 1 keV).
HETG consists of two gratings; MEG (= Medium Energy Grating; energy range = 0.4–5.0 keV)and HEG (= High Energy Grating; 0.8–10.0 keV).

LETG – Low Energy Transmission Grating
The mechanism of LETG is basically the same as HETG, but is designed to cover the lower energy
range (0.08–2 keV). Usually, LETG is used with HRC-S as the focal plane detector.

Chapter 5

Chandra Observations of the ρ
Ophiuchi Cloud
5.1

Observations and Data Reduction

Table 5.1: Log of the Chandra ACIS-I observations of the ρ Oph main bodyObs.ID

Sequence ID

Observation Date (UT)
Start
End

Pointing Direction (J2000)*
R.A.
Decl.

BF . . .
2000602000 Apr 13 18:33
Apr 14 23:14 16h 27m 18.s 1
A ....
200062
2000 May 15 23:35 May 17 03:18 16h 26m 35.s 3*
The position of the detector aimpoint (the telescope optical axis).
†
The live time corrected from deadtime.

−24◦ 340 21.00 9−24◦ 230 12.00 9

Exposure†
(ks)
100.6
96.4

Chandra made two long-exposure observations on ρ Oph with the ACIS-I array (Table 5.1).
The fields of view for the ACIS observations are shown in Figure 3.1, which cover the almost entire
1 pc × 2 pc region of the ρ Oph main body. The first observation (here and after, obs-BF) is
aimed at the south-east 17.0 4 × 17.0 4 area including the millimeter cores B, C, E, and F, while thesecond observation (obs-A) covered the north-west area centered on core A (Loren et al., 1990;
Motte et al., 1998). Although some of the ACIS-S chips were simultaneously in operation, we do
not use these data except for some bright sources (chapter 6), because large off-axis angles cause
degeneration of sensitivity and position determination. From the Chandra X-ray Center (CXC),
we retrieve the level 2 data, in which the data degradation caused by the increase of CTI in orbit
was corrected. The X-ray events were selected with the ASCA grades 0, 2, 3, 4, and 6 (Table 4.2).
31

32

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD

The afterglow events were also removed. After the processing, each observation yields a ≈100 ks
live time (Table 5.1).

5.2

Image and Source Extraction

Fig. 5.1.— False-color ACIS image of the ρ Oph cloud. Red and blue colors represent photons inthe soft (0.5–2.0 keV) and hard (2.0–9.0 keV) bands, respectively.
Figure 5.1 shows the ACIS-I false-color image of the ρ Oph cloud. Red and blue colors represent

5.3. IDENTIFICATION AND OFFSET CORRECTION

33

photons in the soft (0.5–2.0 keV) and hard (2.0–9.0 keV) X-ray bands, respectively. First, we
conduct the source detection analysis for 2048 pixel × 2048 pixel images with the pixel size of
0.00 5 using the wavdetect command (Freeman et al., 2002) in the CIAO 2.2.1 package. We set the√
wavelet scales ranging from 1 to 16 pixels in multiples of 2 and the significance criterion at 10−6 ,which means that the expected false identification of the event pixel is one out of 106 pixel image.
The wavdetect analysis is separately applied for the total (0.5–9.0 keV), soft (0.5–2.0 keV), andhard (2.0–9.0 keV) X-ray bands.
We then detect 195 X-ray sources, nine of which are found in both observations. Thirteen andnine sources are exclusively found only in the soft and hard bands, respectively (hereafter, the “softband sources” and “hard-band sources”). Table 5.2 (column [1]–[4]) shows the name, backgroundsubtracted ACIS-I counts, and coordinate for each source (after the offset correction, see §5.3). The
soft-band and hard-band sources are indicated by the prefix “S” and “H”, respectively. The X-ray
photons are extracted from a circle of 1.00 2–23.00 3 radius depending on the point spread function(PSF) radius for 1.49 keV photons, which is a function of the angular distance from the opticalaxis of the telescope (psfsize20010416.fits1 in the CIAO 2.2.1 package). Typically, in the source
radius, ≈95 % X-ray photons are included. In order to obtain a high signal-to-noise ratio for somefaint sources, we use a circle of a half radius of the PSF. For A-48, we use a larger circle (500 ) than
the PSF because of its diffuse-like feature (§6.6.2). Background counts are extracted from circles
respectively of 38 arcmin2 and 57 arcmin2 source-free regions in the ACIS-I fields of obs-A andobs-BF.
In order to check possible false sources with the wavdetect procedure (Feigelson et al., 2002),
we estimate the confidence level (CL) of the X-ray counts using the Poisson statistics (see §6.4).
Then 14 sources have the CL smaller than 99.9 %. Although A-48 and A-H2 have significantly
high CL, these may also be possible false sources because of its larger source size than the PSF
radius (A-48) and of the severe contamination from A-2 (A-H2). We hence note that 16 sources
are marginal detections, and label them with (m) in column [2] of Table 5.2.

5.3

Identification and Offset Correction

We search for NIR counterparts of the X-ray sources and estimate the absolute positional offsetsof the Chandra observations. For reference of NIR sources, we use the Point Source Catalog in the
2MASS Second Incremental Data Release (Cutri et al., 2000) not in Barsony et al. (1997) used in
Imanishi et al. (2001a), because the former has better position accuracy (σ ∼ 0.00 1) than the latter
1

This file was constructed by using the SAOsac raytrace code (Jerius et al., 1995).
See also http://asc.harvard.edu/chart/.

34

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD(Barsony et al. 1997, σ ∼ 1.00 2) with comparable sensitivity (K . 14.5 mag).
We find the NIR counterparts for about 60 % of the X-ray sources within a 500 radius circle.
Using these Chandra-2MASS pairs, we determine and shift the absolute positions of the Chandra
sources so that the mean values of the Chandra-2MASS offset in the direction of right ascension
and declination become zero. The systematic positional offset of the Chandra frame is (∆R.A.,∆Decl.) = (+0.00 004, +0.00 948) for obs-A and (−0.00 045, +0.00 802) for obs-BF, respectively (Figure5.2). After the offset correction, we repeat the same procedure for the 2MASS counterpart search.
Figure 5.3 shows the resultant offset between Chandra and 2MASS sources plotted against theiroff-axis angle. For the sources within the off-axis angle of 50 , all Chandra-2MASS pairs have smaller
positional offset than 100 . However, it becomes larger as a function of the off-axis angle due to thedegeneration of position determination. Here, we pick up the Chandra-2MASS pairs located below
the dashed line in Figure 5.3, which represents the half-power radius for 1.49 keV X-ray photons(Figure 4.5). For remaining sources, we further search for counterparts using other NIR catalogues,
then identify an X-ray source BF-90 with GY322 (Greene & Young, 1992). We also regard thecounterparts of A-34 to be the nearest NIR source in spite of larger separations than that of thestandard criterion (Figure 5.3, see also §6.3.4). Finally, we conclude that 110 out of 195 X-ray
sources have NIR counterparts. The offset between the X-ray and NIR sources is shown at column[2] in Table 5.3. We also search for radio (cm) and X-ray counterparts in published catalogues(columns [3]–[7] in Table 5.3).

Fig. 5.2.— Offset between the Chandra and 2MASS sources for (left) obs-A and (right) obs-BF,
before the offset correction (see text).

5.4. SOURCE CLASSIFICATION

35

Fig. 5.3.— Offset between the Chandra and
2MASS sources (after the offset correction)plotted against their off-axis angle. Squares
and triangles represent the X-ray sources in
obs-A and obs-BF, respectively. The dashedline is the half-power radius of the PSF for1.49 keV X-ray photons (Figure 4.5).

5.4

Source Classification

Column [9] in Table 5.3 shows the IR classification based on the spectral indices from the NIR to
MIR band (ISOCAM survey at 6.7 and 14.3 µm; Bontemps et al. 2001). Imanishi et al. (2001a)used the terminology class Ic (class I candidate) for sources previously classified as class I (or flat
spectrum source). In this thesis, however, we regard all class Ic s as class II, following Bontemps etal. (2001). Furthermore, we define some additional classes as follows.
1. Class III candidates (IIIc ) are selected among unclassified NIR sources, which are locatedwithin the CS contours (Table 5 in Bontemps et al., 2001).
2. Brown dwarfs (BD) and candidates (BDc ) mean sources with the upper limit of the mass
below the hydrogen-burning limit of 0.08 M¯ and with the lower limit of the mass below
0.08 M¯ , respectively (Wilking, Greene, & Meyer, 1999; Cushing, Tokunaga, & Kobayashi,
2000).
3. Foreground stars (F) are selected with the optical color-color diagram and photometric parallaxes (Festin, 1998).
4. Remaining NIR sources are represented by “?”, which we call “unclassified NIR sources”.
5. Sources with no NIR counterpart are called “unidentified sources”, which are labeled by nodata point (...). Most of the unidentified sources are thought to be background extragalactic
sources (§6.6), although some fractions would be faint cloud members (§6.5).
We then list 8 class Is, 58 class IIs, 17 class IIIs, 9 class IIIc s, 2 BDs, 3 BDc s, 1 F, 12 unclassified
NIR sources, and 85 unidentified sources. Hereafter, we use the terminology “class III+IIIc ” bycombining the class III and class IIIc for brevity.

36

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD

5.5

Timing Analysis

We make two X-ray light curves in 0.5–9.0 keV for all Chandra sources (background is not subtracted) with the respective time bin of 2000 and 4000 s. The source regions are basically thesame as those used in the estimation of the source counts (§5.2). Unfortunately, the brightest twosources (A-2 and BF-64) suffer the photon pileup effect (§4.3), hence we use annulus of 2.00 5–12.00 5
and 2.00 5–7.00 5 radii for A-2 and BF-64, respectively. The light curves show many flare-like events.
We define a “flare time bin” by the following criterion:
N − N0
≥ 2,
∆N

(5.1)

where N , ∆N , and N0 are the X-ray counts, the statistical 1 σ error, and average counts in a
20000 s interval including the relevant time bin, respectively. If the above criterion is satisfied forboth of the two light curves (2000 s and 4000 s time bin), we define the event to be a “flare”. Thenwe pick up 71 flares (Table 5.4). The light curves of the flaring sources are shown in Appendix
A. Most flares have the typical profile of those of low-mass main-sequences and YSOs; fast-riseand slow-decay, while some sources show unusual flares having slow rise timescale; e.g., BF-64 =YLW16A (Figure A.1). We fit the rise and decay phases of the flare by a simple (exponential +constant) model using the QDP command in the LHEASOFT 5.02 , and derive respective e-folding
times and quiescent level (τr , τd , and Q; Table 5.4) with their 90 % errors, where the flare peaktime (tp ) is fixed to be the maximum time bin except for the second flare of BF-64. Evidently, thissimple model can not reproduce the unusually giant flare from BF-64 even if we relax tp to be free.
However, we do not intend to make a model of the flare light curve and only have interested in the
typical timescales, hence do not discuss this discrepancy in further detail.

5.6

Spectral Analysis

X-ray spectra are made for all Chandra sources. For the flaring sources, we make the spectra
separately in the quiescent and flare phases. The flare phase is the time from (tp − τr ) to (tp + τd ).
The quiescent phase is that outside of the flare phase. Background regions are the same as used in§5.2. We then fit them by a thin-thermal plasma model (MEKAL: Mewe et al. 1985; Kaastra 1992;
Liedahl et al. 1995) with the photoelectric absorption (WABS: Morrison & McCammon 1983). The
metal abundance (Anders & Grevesse, 1989) is well determined only for bright sources. Figure 5.4
shows the histogram of the best-fit abundance for sources with the total X-ray photons larger than
1000 counts. We exclude A-2, BF-64, BF-46, and BF-96 in this figure, because of the pileup effectfor the first two sources and their complex spectral shapes for the last two sources (§6.2). The
2

see http://heasarc.gsfc.nasa.gov/docs/software/lheasoft/.

5.7. THE CATALOGUE OF X-RAY SOURCES IN ρ OPH

37

abundance is spread around 0.3 solar, which is consistent with the previous ASCA results (Koyama
et al., 1994; Kamata et al., 1997). We therefore fix the abundance to be 0.3 solar unless otherwise
noted, because photon statistics are generally limited. The free parameters are the time-averaged
plasma temperature (< kT >), emission measure (< EM >), and absorption column (NH ). If thetemperature is not constrained, which is often the case for very faint sources, we fix the temperature
at two representative values of 1 keV (typical of the quiescent state of TTSs; see Figure 7.4) and5 keV (typical of the flare state of protostars), then estimate respective absorptions and luminosities.
For the nine sources detected in both observations, we derive the parameters of each spectrum
assuming the same absorption. This simple model is generally acceptable. BF-46 (ROXs21) and
BF-96 (ROXs31), however, need multi-temperature spectra with unusual abundances (§6.2). A-23,
A-24 and BF-10, on the other hand, require overabundances (§6.3.2, §6.3.3, and §6.3.7). In Table
5.2 (columns [5]–[9]), we summarize the best-fit parameters; < kT >, < EM >, NH , flux, andtime-averaged luminosity (< LX >) in 0.5–9.0 keV.

Fig. 5.4.— Histogram of the best-fit abundance
for bright sources with total X-ray photons larger
than 1000 counts.

5.7

The Catalogue of X-ray Sources in ρ Oph

Here, we present catalogues of 195 Chandra X-ray sources in ρ Oph and detected X-ray flares.
These are an upgraded version of those in Imanishi et al. (2001a) and also will be presented in
Imanishi et al. (2003). Table 5.2 shows the positions and basic X-ray parameters, while Table 5.3
shows identifications of each X-ray source with the previous IR, radio, and X-ray observations. The
detected X-ray flares and their fitting parameters are listed in Table 5.4.

38

5.7.1

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD

Chandra X-ray sources in ρ Oph
Table 5.2:: Chandra X-ray sources in the ρ Oph region

A-1#

Counts* R.A.†(J2000)13.0(m) 25 55.22

Decl.†(J2000)25 37.3< kT >‡(keV)
1.0(fixed)
5.0(fixed)log(< EM >)‡(cm−3 )
53.6(51.9–56.3)
51.6(50.7–53.2)

A-2**

79359.2††26 03.02

23 36.1

A-3

829.9

26 07.05

27 24.4

A-4
A-5

23.7
994.2

26 07.36
26 07.62

25 31.4
27 41.6

A-6
A-7

193.1
83.1

26 10.34
26 12.42

20 54.6
28 49.0

54.8(54.8–54.9)
54.8(54.8–54.8)
53.1(53.0–53.3)
53.3(53.1–53.5)
51.4(51.0–59.7)
53.1(53.0–53.2)
53.2(53.1–53.4)
53.2(53.0–53.3)
53.0(52.6–53.4)
52.1(51.8–52.8)

A-8
A-9#

16.2
12.9

26 13.30
26 14.34

28 23.0
22 28.1

2.9(2.8–3.0)
2.8(2.7–2.9)
2.1(1.8–2.7)
2.5(1.8–4.3)
19.4(>0.2)
2.1(1.8–2.4)
3.5(2.5–4.9)
2.6(1.9–4.0)
2.3(1.4–4.3)
3.3(1.4–
10.3)
2.0(>0.2)
1.0(fixed)
5.0(fixed)

51.4(50.4–51.9)...
3.7(>0.7)
1.2(1.1–1.6)
4.3(>0.9)
5.2(>0.5)
1.7(1.5–1.9)
2.6(2.2–3.1)
1.3(1.2–1.3)
1.4(1.3–1.5)
1.5(1.3–1.7)
1.5(1.4–1.6)
0.9(0.6–1.2)
1.8(0.6–
21.8)
2.7(>0.6)
2.2(1.7–2.9)
2.0(1.8–2.2)
2.4(2.1–2.8)
2.2(1.7–3.0)
2.4(2.0–2.8)
1.0(fixed)...
52.0(51.6–54.2)
52.6(52.3–52.9)
52.8(52.4–53.3)
51.4(50.9–55.4)
53.3(53.2–53.3)
53.5(53.4–53.6)
53.1(53.1–53.1)
53.1(53.1–53.2)
53.4(53.3–53.4)
53.2(53.2–53.3)
52.4(52.2–52.8)
52.4(51.4–55.0)

5.0(fixed)

51.7(50.8–52.3)

2.6(2.1–3.2)
2.8(2.1–4.1)
4.1(2.5–8.6)

52.6(52.6–52.7)
52.9(52.8–53.0)
52.8(52.5–53.1)

No.

A-10‡‡
A-11
A-12

10.9
59.8
180.3

26 15.06
26 15.71
26 15.81

25 48.1
27 49.3
19 22.2

A-13
A-14

16.2
2068.526 16.32
26 16.87

18 43.0
22 23.0A-15

3741.726 17.06

20 21.6

A-16
A-17

126.2
38.3

26 17.13
26 17.47

12 38.9
18 08.8

A-18
A-19
A-20

54.2
405.8
3435.026 18.01
26 18.88
26 22.41

15 17.6
28 19.922 52.8

A-21#

11.2

26 23.22

26 36.8

A-22

986.4

26 23.37

20 59.8

A-23§§

478.9

26 23.61

24 39.6

51.4(50.6–56.1)
53.0(51.4–54.1)

51.9(51.3–53.8)
52.6(52.4–52.7)
53.7(53.7–53.8)
53.7(53.7–53.8)
53.8(53.6–53.9)
53.8(53.7–53.9)
53.6(52.2–54.6)

NH ‡(1022 cm−2 )45.5(>6.7)
15.0(1.6–
129.3)
1.1(1.0–1.1)...
3.7(3.2–4.2)...
4.6(1.9–44.5)
2.8(2.5–3.1)...
...
5.8(4.2–8.5)
5.5(3.5–10.2)

Flux§ LX §Red-χ2k0.5
0.6

0.40(2)0.31(2)1257.2 602.5
1016.0 493.6
13.7 13.6
26.1 22.4
0.8
0.4
13.4 11.7
34.5 20.8
22.0 16.0
8.2
9.5
2.0
1.6

1.51(436)...
0.81(76)...
0.34(7)0.72(114)...
...
0.74(17)0.57(24)

3.0(0.8–10.2)
26.6(3.4–
92.8)
10.8(0.8–
46.5)...
6.6(3.3–25.6)
2.6(2.0–3.0)...
5.8(1.6–43.6)
1.9(1.8–2.1)...
0.5(0.5–0.6)
...
...
...
1.8(1.4–2.2)
11.1(3.0–
37.3)
5.0(1.6–14.4)
1.9(1.6–2.2)
3.6(3.5–3.9)
...
...
...
52.1(9.7–
130.8)
23.3(2.7–
71.4)
1.0(0.9–1.1)...
11.6(9.4–
14.7)

0.3
0.2

0.3
8.9

0.95(2)0.57(2)0.4

0.3

0.39(2)

...
1.5
1.8
15.1
0.5
18.6
54.0
19.0
22.5
41.4
28.4
1.1
0.9

...
1.3
3.1
8.1
0.3
16.3
33.4
10.7
12.1
20.7
14.3
2.5
2.3

...
0.56(14)0.73(15)...
0.01(1)1.14(165)...
1.29(181)
...
...
...
1.26(16)0.71(6)

1.0
5.3
49.2
65.2
60.3
72.7
0.4

1.0
3.6
51.5
56.2
57.4
63.9
37.9

1.06(13)
1.07(37)0.95(178)
...
...
...
0.48(2)0.5

0.7

0.33(2)

9.4
19.4
11.5

4.7
9.3
10.8

1.13(51)...
1.71(50)

38.4
0.6

Comment

Flare1
Flare2
QuiescentFlare
QuiescentFlare1
Flare2

QuiescentFlare
QuiescentFlare
QuiescentFlare1
Flare2
Flare3

QuiescentFlare1
Flare2
Flare3

QuiescentFlare
Quiescent5.7. THE CATALOGUE OF X-RAY SOURCES IN ρ OPH

39

Table 5.2 (Continued)
No.

Counts* R.A.†(J2000)Decl.†(J2000)< kT >‡(keV)
12.1(>5.4)
4.2(>1.8)
2.2(2.0–2.6)
5.7(4.5–7.7)
3.3(2.8–3.9)
12.7(>4.6)
3.3(2.6–4.4)
7.9(5.1–
19.2)
2.2(1.6–3.4)
2.7(1.7–5.3)
3.4(>1.6)
1.9(1.3–3.0)
1.1(0.3–6.8)
1.4(1.0–4.0)
1.7(1.2–4.1)

log(< EM >)‡(cm−3 )
53.3(51.4–53.6)
53.2(51.4–53.7)
53.4(53.3–53.5)
53.9(53.8–54.0)
53.1(53.0–53.1)
53.1(53.0–53.2)
53.1(53.1–53.2)
53.2(53.2–53.3)

NH ‡(1022 cm−2 )...
...
4.2(3.9–4.5)...
1.4(1.3–1.5)
...
...
...

52.8(52.5–53.1)
53.5(53.2–53.8)
53.0(52.7–53.4)
53.0(52.8–53.3)
51.7(50.8–55.0)
51.1(50.9–51.3)
51.8(51.6–51.9)

...

...

2.4(>0.7)
1.8(1.5–2.3)
2.1(1.6–2.7)
3.3(2.1–6.9)
1.0(fixed)

51.4(50.8–52.4)
52.7(52.6–52.9)
52.7(52.5–52.9)
53.0(52.8–53.2)
52.8(52.4–53.4)
51.5(51.2–51.8)
51.5(50.5–55.5)
52.2(49.8–53.1)
51.0(49.8–51.3)
51.2(50.8–51.7)
52.7(52.2–53.0)

A-24§§

2604.926 24.07

24 48.1

A-25

2036.426 24.08

16 13.1

A-26

380.8

26 25.26

23 23.8

A-27
A-28

22.3
88.4

26 25.31
26 25.72

24 44.9
14 27.1

A-29kkA-30‡‡
A-31kk
A-32
A-33
A-34

25.0
10.2
12.2
28.5
336.3
396.5

26 27.47
26 27.83
26 27.89
26 28.43
26 29.69
26 30.55

24 17.8
26 41.8
23 59.2
15 40.9
19 05.3
22 56.8

A-35#

26.8

26 30.62

20 23.4

A-36
A-37#

14.8
15.2

26 31.04
26 31.32

31 04.0
18 33.0

A-38
A-39#

36.6
21.6

26 31.39
26 31.55

25 30.3
23 18.1

5.0(fixed)1.5(>0.2)
1.0(fixed)
5.0(fixed)
4.2(>1.4)
1.0(fixed)

A-40#

13.8

26 33.26

19 51.2

5.0(fixed)
1.0(fixed)

51.4(51.0–51.7)
52.6(51.9–53.5)

5.0(fixed)

51.2(50.8–51.6)

A-41##
A-42#

2249.639.6

26 34.20
26 34.51

23 28.2
22 29.0

1.0(fixed)

53.6(53.1–53.9)

A-43#

14.9

26 36.77

19 33.4

5.0(fixed)
1.0(fixed)

51.6(51.4–51.8)
52.8(48.2–53.5)

A-44
A-45#

105.7
8.4

26 36.81
26 36.85

15 51.4
19 00.6

A-46
A-47‡‡
A-48#

63.6
26 37.16
12.2
26 37.90
14.2(m) 26 38.90

15 59.8
19 42.923 28.1

5.0(fixed)
2.3(1.3–3.8)
1.0(fixed)
5.0(fixed)
0.8(0.6–1.2)...
1.0(fixed)
5.0(fixed)

51.4(51.0–52.0)
52.3(51.9–52.9)
51.5(50.6–52.4)
50.7(50.0–51.0)
51.8(51.3–52.2)...
52.2(49.7–55.1)
50.4(49.8–52.1)

Flux§ LX §Red-χ2k

Comment

87.8
18.7
29.2
253.4
27.1
60.6
32.7
67.1

52.4
19.8
29.2
139.3
13.5
22.6
16.3
26.5

...
...
1.25(114)...
0.99(116)
...
...
...

Flare1
Flare2
QuiescentFlare
QuiescentFlare1
Flare2
Flare36.0(5.0–7.4)
4.9
...
32.7
...
15.1
...
6.8
2.5(0.5–11.1)
0.2
0.4(0.1–0.7)
0.2
...
1.4
– see Table 6.6 –...
...
– see Table 6.6 –1.5(0.2–4.2)
0.4
3.2(2.7–3.9)
4.8
3.6(3.0–4.4)
4.9
...
16.4
14.4(8.9–
0.4
32.5)
6.9(3.8–16.1)
0.6
3.1(0.5–13.2)
0.2
8.0(>35.2)
0.2
3.1(>13.3)
0.2
1.1(0.4–2.7)
0.5
12.3(7.6–
0.3
21.6)
5.8(2.9–13.1)
0.5
12.7(5.8–
0.2
46.7)
5.3(2.0–18.7)
0.4– see Table 6.4 –29.6(18.4–
0.7
43.3)
4.7(3.0–10.3)
0.8
16.4(8.3–
0.3
41.2)
8.0(3.2–25.4)
0.5
5.4(3.4–8.7)
2.0
3.3(0.9–14.2)
0.1
1.3(0.1–5.2)
0.2
1.0(0.8–1.4)
0.4
...
...
10.5(<218.8)
0.1
0.8(<73.9)
0.1

6.0
33.3
12.7
10.3
0.5
0.1
0.6

0.66(64)
...
...
...
1.03(2)
1.29(32)...

QuiescentFlare1
Flare2
Flare3

...

...

0.2
5.1
4.8
11.4
6.1

0.70(14)0.78(29)0.75(69)...
1.50(3)

0.5
0.3
1.3
0.1
0.2
4.2

0.77(3)0.19(3)
1.12(2)0.95(2)0.95(4)1.74(2)

0.4
3.3

1.39(2)1.01(1)0.2

0.03(1)32.1

1.60(6)0.5
5.7

0.88(6)0.13(1)

0.4
2.2
0.3
0.1
0.6
...
1.6
0.1

0.02(1)0.98(19)0.52(1)0.07(1)0.22(11)...
1.82(2)
1.12(2)

QuiescentFlare

Quiescent
Flare

40

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD
Table 5.2 (Continued)

A-49

Counts* R.A.†(J2000)186.0
26 40.50

Decl.†(J2000)27 14.4

A-50
A-51
A-52
A-53

19.0
51.2
123.9
1211.326 41.81
26 42.44
26 42.57
26 42.89

16 05.2
26 25.4
26 31.7
20 29.9

A-54‡‡
A-55
A-56
A-57‡‡
A-58‡‡
A-59
A-60
A-61
A-62#
A-63

11.2
20.8
608.5
11.6
5.9
73.9
81.7
55.9
19.5
1372.526 42.93
26 43.38
26 43.80
26 43.91
26 44.77
26 45.07
26 46.79
26 48.46
26 49.18
26 49.27

22 59.2
24 34.7
16 32.8
27 17.1
22 57.0
23 07.8
19 07.9
28 38.3
21 43.4
20 02.8

A-64
A-65***

7.7(m)
27.6

26 51.94
26 52.51

A-66
A-67#
A-68

33.6
15.5
25.1

A-69

461.7

No.

log(< EM >)‡(cm−3 )
52.4(52.0–52.9)
52.9(52.7–53.3)

NH ‡(1022 cm−2 )5.5(3.8–8.2)...

Flux§ LX §Red-χ2k

Comment

2.1
15.8

2.3
11.3

0.53(32)...

QuiescentFlare

30 39.9
31 33.9< kT >‡(keV)
2.3(1.4–3.6)
4.2(2.4–
15.2)
4.0(>0.5)
1.2(0.6–2.3)
2.6(1.5–5.9)
2.1(1.8–2.4)
2.6(1.9–4.2)...
9.6(>0.2)
1.9(1.5–2.4)...
...
6.7(>1.3)
1.9(1.1–3.4)
1.2(1.0–1.8)
5.0(fixed)
1.7(1.5–1.9)
2.2(1.9–2.4)
2.8(>0.0)
1.8(>0.5)

51.3(50.6–54.4)
52.6(51.9–53.6)
52.4(52.0–52.8)
53.0(53.0–53.1)
53.1(53.0–53.3)...
51.4(51.0–58.1)
52.7(52.6–52.9)...
...
51.7(51.5–52.4)
51.8(51.4–52.2)
52.6(52.4–52.9)
51.1(50.7–51.3)
53.1(53.0–53.2)
53.2(53.1–53.3)
51.1(<55.7)
52.7(52.4–54.1)

0.4
0.8
2.4
14.9
24.7
...
0.7
6.6
...
...
1.6
0.8
0.8
0.3
11.7
22.0
0.1
0.8

0.3
3.1
2.4
10.7
15.1
...
0.4
4.7
...
...
0.8
0.6
3.3
0.2
10.9
16.1
0.1
4.8

0.66(3)0.43(8)1.00(21)
1.13(70)...
...
1.68(1)1.47(52)...
...
0.86(11)1.05(15)0.84(24)
1.38(3)0.94(115)...
0.53(12)1.06(18)

26 53.38
26 54.12
26 54.12

32 36.2
23 25.6
28 40.2

0.9(0.2–3.0)
5.0(fixed)
3.4(>0.2)

53.0(51.7–58.0)
51.0(50.4–51.5)
51.8(51.0–60.4)

0.5
0.2
0.6

8.3
0.1
0.8

0.87(13)
1.36(2)0.95(9)26 54.48

26 20.5

1.8(1.5–2.2)

52.7(52.6–52.8)

3.9(0.7–14.3)
6.3(4.0–10.1)
5.7(4.0–8.2)
1.9(1.6–2.0)...
...
7.6(2.9–32.8)
1.6(1.3–2.0)...
...
3.1(1.7–7.1)
1.6(1.2–2.7)
7.5(5.6–9.4)
2.4(1.1–5.7)
2.1(1.9–2.3)...
5.9(4.0–10.9)
30.6(15.1–
68.5)
9.9(4.4–34.2)
3.1(0.8–25.6)
12.7(4.1–
47.5)
2.5(2.2–2.8)

4.7

4.4

0.75(139)

2.1(1.9–2.4)

52.9(52.7–53.0)

...

8.4

6.9

...

1.1(>0.1)
1.7(1.3–2.3)
2.0(0.8–
10.5)
2.3(>0.4)...

52.0(50.2–57.5)
52.7(52.5–53.0)
51.8(51.2–53.0)

5.6(0.0–65.7)
3.3(2.7–4.1)
3.4(1.5–7.9)

0.2
4.1
0.7

0.9
4.9
0.7

0.81(5)
1.13(25)0.83(10)51.4(50.8–52.7)
2.1(0.7–4.3)...
...
– see Table 6.8 –51.4(51.0–51.9)
0.7(0.5–1.3)
51.9(51.6–52.8)
...

0.4
...

0.3
...

0.14(2)0.85(2)

0.4
1.8

0.2
0.8

0.92(17)...

52.4(51.5–53.6)

0.7

2.2

1.17(19)

19.1

63.7

...

A-70
A-71
A-72

17.1
304.3
48.8

26 54.49
26 55.02
26 58.42

15 55.0
22 29.9
21 29.9

A-73
A-74‡‡
A-75† † †
A-76

17.3
5.6(m)
200.8
90.9

26 58.79
26 58.87
27 00.07
27 00.57

18 34.1
22 15.6
26 40.3
16 25.3

A-77

55.7

27 03.04

26 14.0

A-78

A-79

40.2

4959.927 04.11

27 04.63

28 29.0

27 15.4

1.8(0.7–4.5)
2.4(0.6–
15.5)
2.1(>0.3)
2.0(>0.6)

53.8(52.9–54.9)

20.1(7.6–
38.0)
...

12.1(>1.5)

51.3(51.1–52.1)

3.0(1.9–6.2)

0.8

0.3

1.22(26)

6.3(>1.1)

52.4(52.0–53.3)

...

7.3

3.6

...

2.7(2.4–3.1)

53.5(53.4–53.5)

3.9(3.7–4.1)

39.1

32.0

0.95(239)

QuiescentFlare= BF-2
QuiescentFlare= BF-4= BF-5=
BF-7,
Quiescent
=
BF-7,
Flare

= BF-11
QuiescentFlare= BF-15,
Quiescent= BF-15,
Flare= BF-16,
Quiescent= BF-16,
Flare= BF-17,
Quiescent5.7. THE CATALOGUE OF X-RAY SOURCES IN ρ OPH

41

Table 5.2 (Continued)
No.

A-80
A-81

Counts* R.A.†(J2000)

231.3
291.1

27 05.23
27 06.02

Decl.†(J2000)20 07.1
26 18.6

A-82
A-83
A-84
A-S1‡‡
A-S2#

1192.4
55.5
549.4
5.5
8.5

27 10.34
27 13.72
27 22.98
26 16.94
26 21.55

19 12.6
18 16.3
18 00.0
23 31.0
26 01.0

A-S3‡‡
A-S4‡‡
A-S5‡‡
A-H1‡‡
A-H2

4.4
2.8
2.7(m)
16.4
31.2(m)26 23.51
26 47.24
26 58.20
26 05.71
26 06.52

29 58.4
25 04.0
24 26.7
27 54.8
23 53.0

A-H3‡‡
A-H4‡‡
A-H5‡‡
BF-1
BF-2
BF-3

2.7(m)
6.9(m)
8.3
132.3
45.9
26.0

26 08.18
26 37.17
26 48.72
26 44.19
26 48.46
26 48.91

21 11.4
17 36.4
23 27.4
34 48.2
28 38.3
38 23.7

BF-4

23.4

26 51.96

30 40.2

BF-5***
BF-6#

26.7
21.7

26 52.56
26 53.96

31 29.7
39 24.7

BF-7
BF-8

526.4
34.3

26 54.46
26 57.34

26 20.5
35 38.7

BF-9
BF-10§§

32.8
128.1

26 59.07
26 59.16

35 57.2
34 58.3

BF11† † †BF-12
BF-13‡‡
BF-14‡‡
BF-15
BF-16

229.3

26 59.98

26 40.3

14.1
9.5
5.3(m)
12.6(m)81.6

27 02.32
27 02.46
27 02.83
27 03.04
27 04.10

37 26.7
28 39.9
31 51.5
26 11.5
28 28.9< kT >‡(keV)
5.1(4.2–6.6)

log(< EM >)‡(cm−3 )
54.1(54.0–54.1)

NH ‡(1022 cm−2 )...

Flux§ LX §3.0(2.7–3.4)

54.1(54.1–54.2)...

194.8 147.4 ...

2.5(1.7–4.0)
2.1(1.7–4.1)

52.2(52.0–52.4)
52.5(52.2–52.7)

1.6(1.2–2.2)
3.7(2.7–4.9)

2.8
2.9

1.7
2.9

0.77(47)0.95(70)

4.6(2.3–
19.2)
0.7(0.6–0.8)
1.0(0.5–1.8)
1.5(1.2–1.9)...
1.0(fixed)
5.0(fixed)
...
...
...
2.9(>0.7)...

53.1(52.8–53.4)

...

26.7

15.7

...

53.4(53.3–53.4)
52.2(51.7–53.2)
53.2(53.1–53.4)...
51.8(50.7–52.2)
50.7(50.2–51.1)
...
...
...
51.5(50.9–53.2)...

1.4(1.3–1.4)
2.6(1.7–3.9)
3.3(2.9–3.9)...
3.7(0.7–6.6)
0.6(<2.6)
...
...
...
3.7(1.4–15.8)
...

7.1
0.6
10.9
...
0.2
0.2
...
...
...
0.5
...

22.1
1.5
15.2
...
0.6
0.1
...
...
...
0.4
...

1.17(85)0.70(16)0.77(50)...
1.07(2)0.78(2)...
...
...
0.54(6)
...

...
...
...
3.2(1.6–9.4)
1.3(0.6–1.9)
1.8(0.3–
11.2)
1.5(0.3–
16.0)
1.8(>0.5)
1.0(fixed)

...
...
...
52.5(52.1–53.0)
52.6(51.9–54.1)
51.3(50.8–53.4)

...
...
...
8.3(5.7–12.6)...###
1.4(0.9–4.0)

...
...
...
3.2
0.9
0.3

...
...
...
3.4
3.1
0.2

...
...
...
0.97(25)...###
0.44(5)52.0(51.1–54.2)...###

0.4

0.9...###= A-6452.6(51.5–54.4)
53.3(52.9–53.7)

0.6
0.5

3.8
17.3

...###
1.25(6)= A-655.0(fixed)

51.8(51.5–52.1)

0.9

0.9

1.19(6)

1.6(1.3–1.9)
1.8(0.6–
39.5)
2.8(>0.4)
0.9(0.2–2.6)
2.1(1.3–3.1)

52.9(52.8–53.1)
52.0(51.2–54.2)...###
24.5(16.3–
45.6)
14.0(7.9–
31.2)...###
6.9(3.1–14.3)

6.8
0.5

7.6
0.9

...###
0.73(4)= A-69Quiescent

8.1
0.2
2.8

8.7
0.5
5.1

...
0.51(5)
1.16(20)

Flare

1.3(>0.2)...
...
0.9(0.4–9.8)
2.4(1.2–6.2)

52.9(52.1–55.9)
51.8(51.1–54.1)
52.8(52.3–53.4)...
1.8(0.7–3.8)
12.0(7.8–
19.6)– see Table 6.8 –51.4(50.2–55.5)...
...
53.1(51.2–55.6)
52.0(51.7–52.7)

2.6(0.1–19.3)...
...
...###...###

Red-χ2k

283.2 164.1 ...

Comment= BF-17,
Flare1= BF-17,
Flare2= BF-19,
Quiescent= BF-19,
Flare

Contaminationfrom A-2= A-61= A-75
0.1
...
...
0.3
1.3

0.2
...
...
12.2
1.0

0.13(1)...
...
...###...###= A-77= A-78

42

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD
Table 5.2 (Continued)

No.
BF-17

Counts* R.A.†(J2000)106.5
27 04.59

Decl.†(J2000)27 14.3< kT >‡(keV)
1.3(1.1–1.9)

log(< EM >)‡(cm−3 )
52.3(52.0–52.6)

NH ‡(1022 cm−2 )...###

Flux§ LX §Red-χ2k

Comment

0.9

1.7...###

2.0(1.1–8.1)

53.3(52.9–53.6)

...

16.6

17.8

...###=
A-79,
Quiescent
=
A-79,
Flare

0.8

91.6

1.35(6)

0.8
0.6
0.4
1.6
0.5

0.6
1.2
0.9
5.1
32.2

0.81(6)...###
0.88(12)0.69(18)
0.16(3)0.6

0.7

0.30(3)0.3

21.2

0.90(4)0.4

0.3

1.06(4)

0.8
0.6
22.4
71.8
2.4

0.7
8.2
20.4
52.5
2.2

0.94(11)
1.55(10)1.14(126)...
0.73(26)

QuiescentFlare
QuiescentBF-18#

30.6

27 05.44

41 14.4

1.0(fixed)

54.0(52.5–54.6)

BF-19
BF-20
BF-21
BF-22#

38.8
58.6
106.5
16.6

27 05.95
27 06.60
27 06.78
27 07.66

26 17.7
41 50.0
38 15.2
41 04.1

5.0(fixed)
1.2(0.6–3.3)
0.9(0.5–1.3)
1.2(0.8–2.0)
1.0(fixed)

51.6(51.3–52.2)
52.1(51.6–52.8)
52.0(51.6–53.0)
52.8(52.2–53.5)
53.6(52.9–54.1)

5.0(fixed)

51.7(51.1–52.1)

1.0(fixed)

53.4(52.1–53.9)

5.0(fixed)

51.4(50.6–51.9)

2.0(1.3–4.8)
0.8(0.4–1.3)
3.8(2.9–5.2)
5.8(3.8–9.5)
3.1(1.5–
22.7)
5.0(fixed)
3.8(2.0–
11.2)
2.3(1.2–4.5)
1.0(fixed)

51.8(51.4–52.3)
52.9(52.1–54.2)
53.2(53.1–53.3)
53.5(53.4–53.7)
52.3(51.9–52.9)

56.7(9.0–
102.3)
8.6(3.4–36.4)...###
1.8(1.1–2.5)
5.8(4.0–8.6)
40.1(21.4–
70.9)
17.0(4.5–
37.0)
38.4(9.9–
70.1)
10.9(1.9–
38.7)
2.2(1.5–3.0)
6.3(4.0–9.6)
8.2(7.4–8.8)...
6.6(4.4–9.8)

52.9(52.5–53.1)
52.4(52.1–52.8)...
6.5(4.8–9.1)

14.7
3.7

10.4
3.0

...
0.64(37)

Flare
Quiescent53.6(53.2–54.2)
53.0(52.4–53.5)

40.9
9.8

...
0.96(2)Flare

5.0(fixed)...
2.5(2.1–3.0)
2.8(2.3–3.0)...
9.5(>0.5)
1.0(fixed)

51.3(50.7–51.7)...
53.0(52.9–53.1)
54.0(53.9–54.1)...
51.5(51.2–53.9)
53.2(52.7–53.7)

0.3
...
10.4
110.8
...
0.5
15.3

1.90(2)...
1.08(125)...
...
0.35(3)
1.13(4)5.0(fixed)
3.3(2.6–4.2)
5.9(4.1–8.0)

51.8(51.5–52.0)
53.1(53.0–53.2)
53.6(53.5–53.7)

0.8
16.0
62.6

0.64(4)0.86(94)...

...
5.0(fixed)...
52.2(51.3–52.8)

...
2.4

...
1.66(2)

6.3(>2.6)
1.9(1.6–2.3)
3.4(2.5–5.6)...
1.2(1.0–1.4)

52.2(52.0–52.6)
53.6(53.5–53.8)
53.8(53.7–54.0)...
53.1(52.9–53.3)

...
32.7
24.1(13.1–
0.3
43.5)
9.7(3.2–28.3)
0.3
...
...
3.9(3.6–4.3)
12.0
...
137.2
...
...
3.8(1.9–15.3)
1.0
19.4(11.0–
0.6
32.6)
7.8(4.6–15.2)
1.1
3.0(2.6–3.1)
24.9
...
124.2
– see Table 6.6 –...
...
41.8(5.4–
1.1
95.2)
6.2(4.0–9.3)
4.0
8.0(7.3–9.0)
23.7
...
87.2
...
...
3.2(2.8–3.7)
5.2

2.6
41.5
84.2
...
11.1

0.75(27)1.00(89)...
...
0.97(52)

BF-23#

13.6

27 08.40

39 40.3

BF-24
BF-25
BF-26

74.0
42.9
1141.227 09.12
27 09.32
27 09.44

34 08.2
40 22.1
37 19.0

BF-27

122.8

27 11.17

40 46.9

BF-28

217.8

27 11.73

38 32.4

BF-29#

12.6

27 12.60

40 52.5

BF-30‡‡
BF-31

4.5(m)
1400.527 15.25
27 15.46

34 10.2
26 39.5

BF-32‡‡
BF-33
BF-34#

4.5
29.3
33.3

27 15.46
27 15.51
27 15.55

41 21.4
30 54.0
33 01.8

BF-35

1086.927 15.92

38 43.6

BF-36kk
BF-37‡‡
BF-38#

105.4
9.1
19.0

27 16.41
27 16.98
27 17.47

31 14.7
36 21.2
36 27.1

BF-39
BF-40

155.4
1299.227 17.56
27 18.17

28 56.3
28 52.5

BF-41‡‡
BF-42

8.6
514.0

27 18.34
27 18.48

39 14.3
29 05.7#= A-81

QuiescentFlare

QuiescentFlare

QuiescentFlare
Quiescent5.7. THE CATALOGUE OF X-RAY SOURCES IN ρ OPH

43

Table 5.2 (Continued)
No.

Counts* R.A.†(J2000)Decl.†(J2000)

BF-43‡‡
BF-44
BF-45#

4.5
677.1
55.0

33 44.3
28 43.8
31 30.7

BF46‡ ‡ ‡BF-47‡‡
BF-48#

27 18.99
27 19.21
27 19.34

8664.427 19.53

41 40.7

6.4
13.8

27 19.91
27 20.01

35 31.8
30 13.6

BF-49
BF-50

144.7
121.0

27 21.44
27 21.82

41 43.1
29 53.4

BF-51

56.7

27 21.98

29 39.7

BF-52#

11.8(m) 27 22.42

39 53.0

BF-53‡‡
BF-54‡‡
BF-55

4.5
11.2
54.2

27 22.64
27 23.64
27 24.23

34 14.9
31 57.3
29 29.5

BF-56#
BF-57
BF-58‡‡
BF-59

20.7
55.4
4.3
647.3

27 24.38
27 24.63
27 25.02
27 26.51

41 49.4
29 35.1
34 35.4
39 23.1

BF-60§§§

12.7

27 26.62

25 53.7

BF-613055.027 26.95

40 50.9

BF-62

1103.927 27.09

32 17.6

BF-63

480.4

27 27.41

31 16.4

BF-6421904.0††27 28.04

39 33.5

#

BF-65
BF-66

26.2
309.3

27 28.46
27 28.49

29 42.2
27 20.7

BF-67

18.5

27 28.82

38 06.3< kT >‡(keV)
1.7(1.3–2.4)...
1.6(1.4–1.8)
1.0(fixed)log(< EM >)‡(cm−3 )
53.1(53.0–53.3)...
53.4(53.2–53.5)
53.8(53.4–54.3)

5.0(fixed)

52.1(51.8–52.3)...
1.0(fixed)...
54.1(53.1–54.7)

5.0(fixed)

51.9(51.3–52.4)

1.1(0.8–1.5)
3.9(1.9–
11.1)
1.2(0.5–7.7)
5.1(>1.0)
2.6(>0.3)
1.0(fixed)
5.0(fixed)...
...
2.4(>1.0)
5.0(fixed)
1.0(fixed)2.2(>0.9)...
2.7(2.1–3.5)
6.4(3.9–
14.7)
0.1(0.04–
0.9)
2.3(2.0–2.6)
4.2(3.6–5.1)
2.1(1.8–2.6)
3.5(2.3–5.0)
2.8(2.3–3.6)
1.9(1.5–2.3)
3.4(2.0–7.6)
3.1(2.8–3.6)
7.2(5.5–
10.4)
9.9(8.5–
11.8)
2.6(>0.6)
1.6(1.4–1.9)
1.4(0.6–6.5)
3.0(>0.1)

NH ‡
Flux§
22
−2
(10 cm )
...
10.8
...
...
5.2(4.6–5.8)
11.2
36.0(23.2–
1.0
65.1)
15.7(10.2–
1.5
25.3)
– see Table 6.3 –LX §Red-χ2k

Comment

12.4
...
19.8
57.0

...
...
0.83(58)
1.72(9)

Flare

1.6

1.21(9)

...
0.5

...
...
101.7 0.25(2)

0.6

1.1

0.27(2)52.8(52.4–53.3)
52.3(52.0–52.9)...
74.0(31.1–
129.1)
30.2(10.6–
74.0)
4.7(3.5–6.5)
8.8(6.0–13.4)

1.6
3.1

5.9
2.8

0.90(34)0.65(21)52.1(51.1–53.5)
52.6(52.2–53.6)
52.6(52.0–53.8)
49.9(<51.6)
50.0(<54.6)...
...
52.1(51.6–52.9)
52.9(52.3–53.3)
51.8(51.2–52.1)
52.3(51.7–53.2)...
52.7(52.6–52.9)
53.2(53.1–53.4)

3.6(1.5–7.4)...
...
0.0(<7.0)
0.0(<820.4)...
...
3.2(1.3–6.1)...
3.0(1.8–4.8)
3.0(0.7–6.0)...
3.4(3.0–3.7)
...

0.5
10.4
5.6
0.03
0.06
...
...
1.6
20.9
0.2
2.5
...
7.1
50.8

1.0
5.9
4.6
0.008
0.02
...
...
1.3
11.4
0.5
2.1
...
5.5
26.2

0.65(8)...
...
0.95(4)1.02(4)...
...
0.59(8)...
1.43(5)0.88(8)...
0.85(79)...

53.4(50.3–56.4)

0.9(0.5–1.8)

0.2

2.2

0.73(3)53.5(53.4–53.6)
54.1(54.0–54.2)
53.0(52.9–53.2)
53.0(52.8–53.2)
53.7(53.5–53.8)
52.4(52.3–52.5)
52.7(52.6–52.9)
53.5(53.4–53.6)
53.8(53.7–53.9)

4.8(4.5–5.1)...
4.5(4.0–5.1)...
...
1.2(1.0–1.5)...
5.1(4.9–5.3)
...

30.3
248.5
10.2
16.8
60.9
3.7
13.2
45.8
173.9

31.9
167.0
11.4
12.3
52.4
2.4
6.2
37.8
98.7

0.94(132)...
1.01(105)...
...
0.79(72)...
1.18(424)...

54.8(54.8–54.8)...

1912.2 1008.0 ...

Flare251.9(51.1–54.1)
52.7(52.5–52.8)
52.9(52.3–54.5)
51.4(50.7–54.8)

8.8(3.5–33.3)
2.7(2.4–3.0)...
4.8(1.1–60.9)

0.6
3.5
5.2
0.4

Quiescent
Flare

0.8
4.1
7.4
0.3

0.41(5)0.84(59)...
0.65(2)

QuiescentFlare1
Flare2

QuiescentFlare

QuiescentFlare

QuiescentFlare
QuiescentFlare1
Flare2
QuiescentFlare
QuiescentFlare1

44

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD
Table 5.2 (Continued)

BF-68#Counts* R.A.†(J2000)25.1
27 29.42

Decl.†(J2000)39 16.0< kT >‡(keV)
5.0(fixed)log(< EM >)‡(cm−3 )
51.6(51.2–52.1)

BF-69

95.1

27 29.99

33 36.7

52.2(51.8–52.7)

BF-70#

12.9

27 30.13

38 03.7

3.2(1.6–
10.1)
1.0(fixed)

BF-71
BF-72

233.7
194.8

27 30.15
27 31.09

27 43.6
34 03.4

BF-73
BF-74#
BF-75
BF-76
BF-77#

74.7
9.8
357.5
189.2
14.3

27 31.33
27 32.12
27 32.72
27 32.89
27 33.04

31 10.1
29 43.3
33 24.3
32 35.0
31 48.3

5.0(fixed)
1.6(0.9–2.6)
1.6(1.3–2.3)
1.2(>0.4)
3.1(>1.2)
1.0(fixed)
1.9(1.4–2.8)
1.7(1.2–2.5)
1.0(fixed)

51.0(50.5–51.6)
52.6(52.2–53.2)
52.0(51.8–52.4)
52.6(52.3–53.9)
52.1(51.6–53.0)
51.9(51.3–52.3)
52.7(52.5–53.0)
52.4(52.2–52.7)
52.8(51.8–54.8)

5.0(fixed)

51.4(50.7–52.3)

2.1(1.9–2.3)
4.0(3.2–5.0)
0.7(0.6–0.8)
5.0(fixed)
1.0(fixed)

53.4(53.3–53.4)
53.6(53.5–53.7)
51.6(51.3–51.9)
51.5(50.8–51.8)
53.5(52.4–54.8)

5.0(fixed)

51.7(51.0–52.8)

1.6(0.6–5.1)
1.0(fixed)

51.9(51.2–53.1)
52.7(51.9–53.2)

No.

BF-78

2853.727 33.15

41 15.4

BF-79
BF-80#
BF-81#

96.4
27.9
11.1

27 35.29
27 36.01
27 36.08

38 33.7
32 18.034 30.6

BF-82
BF-83#

26.4
18.8

27 37.26
27 37.83

42 37.4
36 04.7

BF-84

175.7

27 38.16

30 42.6

BF-85
BF-86
BF-87

1162.0264.6
764.1

27 38.35
27 38.67
27 38.98

36 58.4
38 39.2
40 20.5

BF-886909.327 39.47

39 15.5

BF-89

908.5

27 39.88

43 15.3

BF-90
BF-91
BF92kkk

52.2
125.9
229.0

27 41.53
27 42.79
27 43.63

35 37.6
38 51.0
31 26.5

5.0(fixed)
1.1(0.9–1.5)
1.2(0.8–1.8)
1.7(1.5–1.9)
1.6(0.9–2.0)
3.2(2.3–4.4)
4.5(3.8–5.9)
2.2(2.0–2.5)
3.2(2.9–3.7)
2.9(2.6–3.1)
2.2(1.9–2.7)
1.8(1.6–2.1)
5.0(fixed)
2.3(1.6–3.0)
2.5(1.5–5.9)
1.4(0.7–3.7)
0.9(0.7–1.1)>5.8

BF-93#

15.9

27 45.51

38 08.9>1.5
5.0(fixed)#

NH ‡(1022 cm−2 )10.5(4.2–
38.0)
6.1(4.0–9.6)

Flux§ LX §Red-χ2k0.6

0.5

1.47(3)2.1

1.8

0.71(14)0.2

4.2

0.73(2)

0.3
2.6
1.2
3.7
1.5
0.1
4.8
2.1
0.2

0.2
3.6
1.0
3.8
1.6
0.7
5.3
2.4
2.0

0.19(2)
1.47(46)
1.11(31)...
1.14(13)1.30(1)1.13(60)
1.11(31)1.95(1)0.4

0.2

0.81(1)

32.2
105.0
0.4
0.6
0.4

23.1
52.8
0.4
0.4
26.4

1.60(117)...
1.06(19)
1.21(5)0.53(1)0.6

0.7

0.19(1)

0.5
0.3

0.7
4.8

0.77(9)0.71(3)51.2(50.7–51.6)
52.3(52.1–52.4)
52.7(52.4–53.1)
53.1(53.0–53.2)
52.1(51.9–52.5)
52.9(52.7–53.1)
54.0(53.8–54.2)
53.6(53.5–53.6)
53.7(53.7–53.8)
53.8(53.7–53.8)
53.7(53.6–53.7)
53.3(53.2–53.5)
53.6(53.6–53.7)
53.5(53.3–53.7)
53.8(53.5–54.1)
51.9(51.4–52.8)
52.6(52.0–53.0)
52.3(52.3–52.4)

18.4(3.0–
78.4)
3.9(1.1–27.1)
3.5(2.3–5.3)
1.3(0.9–1.5)...
7.9(4.3–15.1)
5.1(2.6–9.6)
3.8(3.2–4.8)
3.1(2.4–4.1)
22.7(7.0–
173.4)
11.4(1.8–
81.4)
1.9(1.8–2.0)...
0.6(0.3–0.8)
6.3(0.8–15.2)
40.8(13.1–
113.5)
18.5(5.2–
78.6)
3.7(1.9–8.1)
15.6(5.9–
34.5)
4.5(1.2–17.6)
1.9(1.6–2.2)...
1.9(1.7–2.1)
0.8(0.6–1.5)
8.8(7.5–10.6)...
1.8(1.7–1.9)
...
...
...
5.4(4.9–6.1)
...
...
...
2.6(1.9–4.3)
2.9(1.7–3.9)
6.0(4.9–7.5)

0.4
1.1
3.0
12.3
1.7
8.2
148.7
54.7
119.8
111.1
66.9
12.6
93.2
27.3
64.5
0.5
0.9
6.5

0.2
1.7
4.2
10.7
1.0
8.7
126.1
36.5
64.8
64.1
44.7
18.5
61.2
30.5
65.1
0.7
3.3
3.6

0.50(3)0.92(31)...
1.25(48)
1.19(46)
1.04(71)...
1.18(290)
...
...
...
0.84(115)
...
...
...
1.18(9)1.41(21)1.03(22)52.8(52.5–52.9)
51.9(50.1–56.8)...
29.1(<721.3)

17.9
0.6

10.0
0.8

...
0.66(10)52.7(51.3–53.8)

Comment

QuiescentFlare

QuiescentFlare

QuiescentFlare

QuiescentFlare
QuiescentFlare1
Flare2
Flare3
QuiescentFlare1
Flare2
Flare3

QuiescentFlare5.7. THE CATALOGUE OF X-RAY SOURCES IN ρ OPH

45

Table 5.2 (Continued)

BF-94

Counts* R.A.†(J2000)29.8
27 49.16

Decl.†(J2000)39 06.9< kT >‡(keV)
2.7(>0.2)log(< EM >)‡(cm−3 )
51.9(51.2–60.5)

BF-95

162.6

27 50.42

31 49.5

1.4(0.8–2.2)

53.1(52.5–54.4)

BF96‡ ‡ ‡BF-97#

4289.327 52.01

40 50.0

36.4

27 54.71

37 14.8

1.0(fixed)

52.9(52.6–53.2)

BF-98
BF-S1‡‡
BF-S2‡‡
BF-S3‡‡
BF-S4#
BF-S5‡‡
BF-S6‡‡
BF-S7‡‡
BF-S8‡‡
BF-H1#

24.8
8.9
8.4
6.4
7.4
6.4(m)
4.4
4.4
2.9
9.8(m)

27 55.34
26 49.06
26 51.24
26 51.53
27 04.57
27 07.72
27 17.57
27 24.87
27 32.88
26 43.55

28 39.3
37 31.7
32 42.4
28 54.8
24 57.6
27 33.0
29 22.8
41 04.4
34 47.4
33 00.5

5.0(fixed)
0.9(0.2–2.2)
...
...
...
1.0(fixed)
...
...
...
...
5.0(fixed)

51.6(51.3–51.8)
51.5(50.8–53.3)
...
...
...
50.0(<50.6)
...
...
...
...
52.4(51.3–53.2)

BF-H2‡‡
BF-H3#

9.3(m)
11.1

26 49.38
27 32.06

36 18.7
38 33.1...
1.0(fixed)...
52.3(51.0–53.2)

BF-H4#

6.6(m)

27 50.01

36 49.1

5.0(fixed)
1.0(fixed)
5.0(fixed)

50.9(50.8–57.5)
51.5(49.6–54.5)
50.6(49.7–52.1)

No.

NH ‡
Flux§ LX §22−2
(10 cm )10.7(4.2–
0.6
0.9
26.9)
11.1(6.9–
2.8
11.4
23.2)
– see Table 6.3 –Red-χ2k12.8(8.8–
19.9)
6.6(4.1–12.1)
1.2(0.3–2.5)
...
...
...
0.0(<0.8)
...
...
...
...
55.1(14.8–
164.6)...
10.8(2.2–
49.1)
3.5(<857.9)
3.5(<139.8)
1.4(<69.9)

1.57(7)
1.23(39)0.5

7.6

1.08(10)

0.8
0.2
...
...
...
0.04
...
...
...
...
1.1

0.6
0.3
...
...
...
0.01
...
...
...
...
3.4

1.29(10)0.70(5)...
...
...
1.86(1)
...
...
...
...
1.15(4)

...
0.1

...
1.6

...
1.08(1)

0.2
0.1
0.1

0.1
0.3
0.1

1.00(1)1.56(1)1.08(1)

Comment

46

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD(footnotes for Table 5.2)*

Background-subtracted X-ray counts in 0.5–2.0 keV, 2.0–9.0 keV, and 0.5–9.0 keV for soft-band, hard-band, and theother sources, respectively. (m) denotes sources with marginal detections (the confidence level < 99.9 %, see §5.2).
Although the confidence levels of A-48 and A-H2 are significant enough, we regard them as marginal sources becauseof the larger source size (A-48) and severe contamination from A-2 (A-H2).
†

Right ascension and declination for all sources are 16h and −24◦ .

‡Parentheses indicate the 90 % confidence limits.

§

Observed flux (10−14 ergs s−1 cm−2 ) and absorption-corrected X-ray luminosity (1029 ergs s−1 ; “< LX >” in thetext) in 0.5–9.0 keV.
kReduced-χ2 for the spectral fittings. Parentheses indicate the degrees of freedom.

#

We determine spectral parameters with fixed temperatures of 1 keV and 5 keV (see text). For sources which only

show the parameters for kT = 1 or 5 keV, no good fitting is obtained for the other temperature.
**

We assume the same abundances as the “F2” phase in Table 6.3. The quiescent spectrum is not obtained because

the decay phases of the two flares occupy all of the light curve.
††

The pileup effect is not corrected.

‡‡

No spectral fit is done due to the limited statistics.

§§

Abundances are free parameters (§6.3.2, §6.3.3, and §6.3.7).

kk

We make the flare spectra with a bit larger timescale in order to obtain as good statistics as possible. Errors of< EM > for A-29 in the quiescent are not determined because of the limited statistics.
##

The spectra show the possible absorption edge of neutral Ca or warm Ar. The non-thermal model also well

reproduces the spectra (§6.3.5, see also Hamaguchi et al. 2002).
***

We assume the same temperature because of the limited statistics.

†††

These show non-thermal spectra (§6.6.3).

‡‡‡

These show two temperature model with an unusual abundance pattern (§6.2).

§§§

Foreground star. The distance is 60 pc (Festin, 1998).

kkk

The best-fit value of < kT > is not determined (larger than 10 keV), hence we assume 10 keV temperature forthe estimation of the other parameters.
###

NH and reduced-χ2 are estimated by the simultaneous fittings with the identical sources in obs-A.

5.7. THE CATALOGUE OF X-RAY SOURCES IN ρ OPH

5.7.2

47

Identifications of the X-ray sources
Table 5.3:: Identifications of the X-ray sources
Radio†A-1
A-2

Offset*
(00 )
...
0.33

...
R8, S7—- X-ray —-‡P SP C HRI ASCA IKT
...
...
...
...
13
A11 3
...

A-3
A-4
A-5
A-6
A-7
A-8
A-9
A-10
A-11
A-12
A-13
A-14

0.40
7.54(R)
0.60
0.23
...
...
...
...
...
0.47
...
0.11

...
R9?
...
...
...
...
...
...
...
...
...
Lp1

...
...
15?
16
...
...
...
...
...
...
...
18

...
...
A13
...
...
...
...
...
...
...
...
A14

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

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

A-15

0.44

...

17

A15

...

...

A-16
A-17
A-18
A-19
A-20
A-21
A-22

1.08
...
...
0.38
0.03
...
0.19

...
...
...
...
L2
...
...

...
...
...
19
...
...
24(22?)

...
...
...
...
A17
...
A18

...
...
...
...
...
...
4

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

A-23
A-24

0.23
0.14

L3
...

...
25

...
A19

...
C3

...
...

A-25
A-26
A-27
A-28
A-29
A-30
A-31
A-32
A-33
A-34

0.15
0.09
0.12
0.55
...
0.05
2.66(R)
0.92
0.06
0.72

...
...
...
...
...
...
L6?
...
L8
L9

26
...
...
...
...
C15
...
...
27
...

A20
...
...
...
...
...
...
...
...
...

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

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

A-35
A-36
A-37
A-38
A-39
A-40

...
1.14
1.42(R)
0.24
...
...

...
...
L10
...
...
...

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

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

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

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

No.

Other names§

Classk...
DoAr21, ROXs8, YLW26, GSS23,
Elias14, SKS1-5, ISO10
ISO13
...
ISO14
GSS26, CRBR5, SKS1-6, ISO17
...
...
...
...
...
CRBR9, SKS1-7, ISO18...
ROXs11, GSS29, Elias18, CRBR10,
SKS1-8, ISO19
DoAr24, ROXs10A, GSS28, Elias19,
CRBR11, SKS1-9, ISO20
2MASSI J1626172−241238...
...
YLW31, VSSG1, Elias20, ISO24
GY12, VSSG12, ISO34...
ROXs10B, GY20, GSS31, Elias22,
CRBR29, SKS1-15, ISO36
GY21, CRBR30, SKS1-16, ISO37
S2, GY23, GSS32, Elias23, CRBR32,
SKS1-17, ISO39
YLW32, Elias24, ISO40
GY31, CRBR34, SKS3-15
GY29, CRBR35, SKS1-18, ISO41
2MASSI J1626257−241427...
GY37, VSSG29, CRBR37, ISO42
...
ISO44
SKS1-19, ISO45
GY51, VSSG27, CRBR38, SKS1-20,
ISO46
...
YLW35, GY54, IRS14, ISO47...
GY59, CRBR39, SKS1-21
...
...

...
III
II
...
III
II
...
...
...
...
...
IIIc
...
II
II
?
...
...
II
III
...
II
II
II
II
?kk
II
?
...
BDc (IIIc )
...
IIIc
III
II
...
IIIc
...
BDc
...
...

48

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD
Table 5.3 (Continued)
Radio†

A-41

Offset*(00 )0.26

R15, S13, L11—- X-ray —-‡P SP C HRI ASCA IKT
28
A21 5
...

A-42
A-43
A-44
A-45
A-46
A-47
A-48
A-49
A-50
A-51
A-52
A-53

...
...
0.60
0.89
0.07
...
...
0.06
...
0.53
0.32
0.23

...
...
...
...
...
...
...
...
...
R16, L15
...
...

...
...
...
...
...
...
...
...
...
...
...
29

...
...
A22
...
...
...
...
...
...
...
...
A23

...
...
...
...
...
...
...
...
...
...
...
C4

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

A-54
A-55
A-56
A-57
A-58
A-59

0.25
...
0.44
...
...
0.31

L16?
...
R17, S15
...
...
Lp3

C19
...
30
...
...
...

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

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

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

A-60
A-61
A-62
A-63
A-64
A-65
A-66
A-67
A-68
A-69
A-70
A-71
A-72
A-73
A-74
A-75
A-76
A-77
A-78

0.32
0.67
...
0.26
0.52
...
1.51
...
...
0.33
...
0.51
0.16
2.05
...
1.95(R)
0.20
0.71
0.93

...
...
...
L18
...
...
...
...
...
...
...
...
...
...
...
R19, S20, L21
...
...
...

...
...
...
32
C21?
...
...
...
...
33
...
C22?
C23?
...
...
...
...
...
...

...
...
...
A25
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...

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

...
2
...
...
4
5
...
...
...
7
...
...
...
...
...
11
...
12
13

A-79
A-80
A-81
A-82

0.54
0.88
0.45
0.73

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

C24?
...
C25
38

...
...
...
A26

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

14
...
16
...

A-83
A-84
A-S1

0.57
2.82
...

...
R21
...

...
42
...

...
...
...

...
...
...

...
...
...

No.

Other names§

Classk

S1, ROXs14, YLW36, GY70, GSS35,
Elias25, CRBR41, SKS1-22, ISO48
...
...
ISO51
ISO50
2MASSI J1626371−241559...
...
GY91, CRBR42, ISO54...
GY101, CRBR43, ISO60
GY103, CRBR44, ISO61
GY110, VSSG2, GSS37, Elias26,
ISO62
GY109, ISO63...
YLW39, VSSG11, IRS19, ISO64...
...
GY116, VSSG28, GSS39, Elias27,
ISO67
2MASSI J1626467−241908
GY128, WL2, BBRCG6, ISO70...
ROXs17, GY135, VSSG3, ISO73
GY144, BBRCG13, CRBR47, ISO75...
GY146, BBRCG14, CRBR48, ISO76...
...
GY153, VSSG5, ISO78...
GY156, VSSG6, GSS40, ISO80
YLW1C, GY171, IRS26, ISO86
ISO87
...
...
2MASSI J1627005−241625
GY188, ISO93
YLW6, GY192, WL1, VSSG31,
BBRCG20, CRBR55, IRS29, ISO95
GY195, ISO98
GY199, VSSG21
YLW10A, GY203, IRS30, ISO101
SR21, YLW8, VSSG23, Elias30,
ISO110VSSG24, ISO116
ROXs23, VSSG22, ISO135
...

III
...
...
II
IIIc
?
...
...
I
...
III
III
II
II
...
III
...
...
II
?
II
...
III
II
...
II
...
...
II
...
III
II
II
...
...
?
II
II
II
?
III
II
II
III
...

5.7. THE CATALOGUE OF X-RAY SOURCES IN ρ OPH

49

Table 5.3 (Continued)
Radio†

A-S2
A-S3
A-S4
A-S5
A-H1
A-H2
A-H3
A-H4
A-H5
BF-1

Offset*
(00 )
0.10
...
...
...
...
...
...
...
...
0.25

BF-2
BF-3
BF-4
BF-5
BF-6
BF-7
BF-8

0.72
1.74
0.73
...
...
0.05
0.04

(= A-61)...
(= A-64)(= A-65)...
(= A-69)
...

BF-9

0.21

BF-10

No.

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

—- X-ray —-‡P SP C HRI ASCA IKT
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
1

Other names§

Classk(GY5, CRBR21, SKS1-13, ISO30)‡‡
...
...
...
...
...
...
...
...
YLW2, GY111, WL12, VSSG30,
BBRCG4, ISO65

BDc (II)
...
...
...
...
...
...
...
...
I

...

...

...

3

GY129, WL18, BBRCG8, ISO72

II

...

...

...

6

...

...

34

...

...

8

II

...

...

...

...

9

0.64

L23

...

...

...

10

YLW4B, GY164, WL21, BBRCG17,
CRBR50, SKS3-23, IRS23, ISO84
YLW4A, GY172, WL14, BBRCG18,
CRBR52, SKS3-25, ISO89
YLW4C, GY174, WL22, CRBR53,
IRS27, ISO90

BF-11
BF-12

0.97(R)
0.71(= A-75)
...

...

...

...

...

BF-13
BF-14
BF-15
BF-16
BF-17
BF-18
BF-19
BF-20
BF-21

...
...
3.16
1.11
1.35
...
0.98
1.17
0.33

...
...
(= A-77)(= A-78)(= A-79)...
(= A-81)
...
...

...
...

...
...

...
...

...

...

...
...

BF-22
BF-23
BF-24

1.79
...
0.10

...
...
...

BF-25
BF-26

0.73
0.27

BF-27
BF-28
BF-29
BF-30

II
II

II

...
...

YLW5, GY182, WL16, BBRCG19,
CRBR54, SKS1-25, ISO92
...
...

...

15

...

...

...
...

...
C5

17
18

II
II

...
...
37

...
...
...

...
...
7

19
20
21

...
L27

...
...

...
...

...
8

22
23

0.47
0.45

...
...

...
...

...
...

...
...

24
25

...
...

...
...

...
...

...
...

...
...

...
...

GY204, SKS1-26, ISO102
GY205, WL17, BBRCG22, CRBR58,
SKS3-27, ISO103
2MASSI J1627075−244104...
GY211, WL10, BBRCG24, CRBR59,
SKS1-27, ISO105
GY213, SKS3-28, ISO107
YLW7, GY214, WL15, Elias29,
BBRCG23,
CRBR60,
SKS1-28,
ISO108
GY224, SKS3-30, ISO112
YLW11A, GY227, WL19, BBRCG28,
CRBR65, SKS3-31, ISO114
...
...

...
...

?
...
II
II
I

II
II
...
...

50

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD
Table 5.3 (Continued)
Radio†

BF-31

Offset*
(00 )
0.29

...

—- X-ray —-‡P SP C HRI ASCA IKT
...
...
...
26

BF-32
BF-33

...
0.52

...
...

...
...

...
F15

...
...

...
27

BF-34
BF-35

...
0.28

...
L30

...
39

...
F16

...
...

28
29

BF-36
BF-37
BF-38
BF-39

...
5.71(R)
...
0.23

...
L32?
...
...

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

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

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

30
31
32
33

BF-40

0.19

R20, S22

40

...

...

34

BF-41

0.77

...

...

...

...

35

BF-42

0.24

...

...

...

C6

36

BF-43
BF-44

...
0.31

...
...

...
...

...
...

...
...

37
38

BF-45
BF-46

...
0.26

...
Lp6

...
41

...
F17

...
9A

39
40

BF-47
BF-48
BF-49

...
...
0.49

...
...
...

...
...
...

...
...
...

...
...
...

41
42
43

BF-50

0.19

...

...

...

10

44

BF-51

0.16

...

...

...

...

45

BF-52
BF-53
BF-54
BF-55
BF-56
BF-57
BF-58
BF-59

...
...
...
0.20
1.88
0.38
...
0.19

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

...
...
...
...
...
...
...
45?

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

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

...
46
47
48
49
50
51
53

BF-60
BF-61

0.73
0.25

...
R24, L33

...
43

...
F21

...
T††

52
54

BF-62

0.20

...

...

...

...

55

No.

Other names§

ClasskYLW10B, GY239, SKS3-36, IRS34,
ISO120...
GY238, BBRCG32, CRBR70, SKS335, IRS35, ISO119...
YLW11B, GY240, WL20, BBRCG31,
CRBR68, SKS1-31, ISO121
...
...
...
YLW12A,GY244,
VSSG26,BBRCG36,
CRBR71,
SKS1-32,
IRS37, ISO124
ROXs22, YLW12B, GY246, WL5,
VSSG26, BBRCG37, CRBR73, SKS133, IRS38, ISO125
GY245, BBRCG39, CRBR72, SKS337, ISO127
YLW12C, GY247, WL4, VSSG26,
BBRCG40,
CRBR75,
SKS1-34,
IRS39, ISO128...
YLW12D, GY249, WL3, VSSG26,
BBRCG41,
CRBR77,
SKS3-38,
IRS41, ISO129...
SR12, ROXs21, YLW13A, GY250,
CRBR76, SKS1-35, IRS40, ISO130...
...
YLW13B, GY252, CRBR79, SKS1-36,
IRS42, ISO132
YLW14, GY254, WL6, BBRCG38,CRBR80, SKS3-40, ISO134
GY256, BBRCG43, CRBR81, SKS341
...
...
...
GY257, CRBR83, SKS3-43
GY258, CRBR84, SKS3-42, ISO136
GY259, CRBR86, SKS3-45...
YLW16C,
GY262,BBRCG45,CRBR88, SKS3-47, ISO140
GY264, SKS1-37
YLW15A, GY265, CRBR89, SKS3-49,
IRS43, ISO141
GY266, SKS3-50

II
...
II
...
II
...
...
...
II

III

II
II

...
II

...
III
...
...
II
I
II
...
...
...
?
IIIc
?
...
II
F
I
?

5.7. THE CATALOGUE OF X-RAY SOURCES IN ρ OPH

51

Table 5.3 (Continued)
Radio†

BF-63

Offset*
(00 )
0.24

...

—- X-ray —-‡P SP C HRI ASCA IKT
44
F22
...
56

BF-64

0.49

L35

45

G**

...

57

BF-65
BF-66

...
0.88

...
...

...
...

...
...

...
...

58
59

BF-67
BF-68

...
0.25

...
...

...
45

...
...

...
9B

...
60

BF-69
BF-70
BF-71

0.52
...
0.30

...
...
...

...
...
...

...
...
...

...
...
...

61
62
63

BF-72
BF-73
BF-74
BF-75
BF-76
BF-77
BF-78
BF-79
BF-80
BF-81
BF-82
BF-83
BF-84
BF-85

0.51
...
0.54
0.69
0.53
...
0.56
0.42
...
...
0.49
...
0.82
0.65

...
...
...
Lp7
...
...
...
...
...
...
...
...
...
...

47
...
...
47?
48
...
49
...
...
...
...
...
...
50

...
...
...
...
...
...
F27
...
...
...
...
...
...
F29

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

64
65
...
66
67
68
69
70
71
...
72
73
74
75

BF-86
BF-87
BF-88
BF-89

0.58
0.39
0.50
0.72

...
...
Lp8
...

...
...
51
...

...
...
F30
...

...
...
...
T††

76
77
78
79

BF-90
BF-91
BF-92
BF-93
BF-94
BF-95
BF-96
BF-97
BF-98
BF-S1
BF-S2
BF-S3
BF-S4
BF-S5
BF-S6
BF-S7

(0.62)#
1.41
...
...
...
...
0.85
...
1.44
...
0.98
...
...
...
...
...

...
...
...
...
...
...
R32
...
...
...
...
...
...
...
...
...

...
...
...
...
...
...
54
...
...
...
...
...
...
...
...
...

...
...
...
...
...
...
F33
...
...
...
...
...
...
...
...
...

...
...
...
...
...
...
T††
...
...
...
...
...
...
...
...
...

80
81
82
...
83
84
85
86
87
...
...
...
...
...
...
...

No.

Other names§

Classk

GY267, WL13, VSSG25, Elias31,
BBRCG42, SKS1-38, ISO142
YLW16A, GY269, BBRCG49, SKS351, IRS44, ISO143...
YLW17A, GY273, VSSG18, Elias32,
SKS1-39, IRS45, ISO144...
YLW16B, GY274, BBRCG52, SKS354, IRS46, ISO145
GY278, BBRCG53, ISO146...
YLW17B, GY279, VSSG17, Elias33,
SKS1-40, IRS47, ISO147
GY283, BBRCG56, ISO148...
GY287, SKS3-56
GY289, BBRCG60, ISO152
GY291, BBRCG61, ISO154...
GY292, CRBR90, ISO155
GY295, SKS1-44, ISO156...
...
GY301, SKS3-59, ISO161...
GY306, SKS1-47, IRS50
YLW47, GY308, SKS1-48, IRS49,
ISO163
GY310, SKS1-49, ISO164
GY312, SKS3-61, ISO165
ROXs27, GY314, SKS1-50, ISO166
YLW45, GY315, SKS3-63, IRS51,
ISO167
GY322, ISO169
GY326, SKS1-52, ISO172
...
...
...
...
ROXs31, GY380, IRS55, ISO184...
GY397, ISO185...
GY141, CRBR46
...
...
...
...
...

II
I
...
II
...
I
IIIc
...
II
IIIc
...
?
III
II
...
II
IIIc
...
...
II
...
III
II
BD(II)
II
II
I
IIIc
IIkk
...
...
...
...
III
...
II
...
BD
...
...
...
...
...

52

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD
Table 5.3 (Continued)

No.
BF-S8
BF-H1
BF-H2
BF-H3
BF-H4

Offset*
(00 )
...
...
...
...
...

Radio†
...
...
...
...
...

—- X-ray —-‡P SP C HRI ASCA IKT
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...
...

Other names§

Classk

...
...
...
A2§§
...

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

5.7. THE CATALOGUE OF X-RAY SOURCES IN ρ OPH

53(footnotes for Table 5.3)*

Offset between the Chandra and nearest 2MASS sources. For sources having a radio counterpart only, the offset

between the Chandra and radio sources (André et al., 1987; Leous et al., 1991) are shown, which is indicated by (R).
†Radio sources. S, R, and L indicate sources listed in André et al. (1987) (ROC), Stine et al. (1988) (SFAM), andLeous et al. (1991) (LFAM), respectively. Lp denotes possible radio sources in LFAM.
‡

X-ray sources detected with ROSAT/PSPC (Casanova et al., 1995), ROSAT/HRI (Grosso et al., 2000), and ASCA

(Kamata et al., 1997). “IKT” denotes Chandra sources already listed in Imanishi et al. (2001a).
§

Source names used in the literature. Abbreviations for names are SR (Struve & Rudkjobing, 1949), DoAr (Dolidze& Arakelyan, 1959), ROXs (Bouvier & Appenzeller, 1992), S (Grasdalen et al., 1973), YLW (Young et al., 1986), GY(Greene & Young, 1992), WL (Wilking & Lada, 1983), VSSG (Vrba et al., 1975), GSS (Grasdalen et al., 1973), Elias(Elias, 1978), BBRCG (Barsony et al., 1989), CRBR (Comerón et al., 1993), SKS (Strom et al., 1995), IRS (Wilking
et al., 1989), and ISO (Bontemps et al., 2001). For sources having NIR counterpart only in the 2MASS catalog, we
alternatively show the 2MASS source names.
k

Source classifications; I: class I, II: class II, III: class III, IIIc : class III candidate (Table 5 in Bontemps et al. 2001),

BD: brown dwarf, BDc : brown dwarf candidate (Imanishi et al., 2001b), F: foreground star (Festin, 1998). Unclassified
NIR sources are indicated by “?”, and sources with no NIR counterpart are called “unidentified sources”, labeled nodata point (...) in this column. For BD/BDc s, we show the available IR classification in Bontemps et al. (2001) inthe parentheses.
#

Offset from a nearest NIR source in Greene & Young (1992).

**

Grosso (2001).

††

Tsuboi et al. (2000).

‡‡

The position of GY5 (Greene & Young, 1992) is slightly shifted (∼2.00 3) from the 2MASS source.

§§

A candidate of HH object (Gómez et al., 1998, see also §6.6.6).

kk

From NIR spectroscopy, Wilking et al. (1999) derived for these M dwarfs masses higher than the hydrogen burning

limit.

54

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD

5.7.3

Detected X-ray flares
Table 5.4:: Detected X-ray flares
No.

tp
(ks)

A-49
BF-26
BF-61
BF-64

58
18
82
22
87.3(87.1–87.6)§
4
88
...‡

†

BF-89

A-3
A-14
A-15

A-22
A-23
A-24
A-25

A-34
A-53
A-69
A-77
A-78
A-79
BF-8
BF-17
BF-27
BF-28
BF-31
BF-35
BF-42

24
8
2
60
86
20
56
74
72
10
38
88
36
82
70
98
94
42
62
66
32
86
18
54
61
76

τr
(ks)

τd
(ks)
— Class I —0.9*
12.2(9.2–16.0)
7.3(6.0–9.0)
5.5(4.4–6.7)
2.9(2.6–3.2)
14.0(12.7–15.6)
2.5(2.0–3.2)
14.6(13.2–16.1)*
5.1
27.7(24.3–31.5)
1.2(0.8–1.6)
3.4(2.9–4.1)
4.0(2.4–6.2)
3.8(2.1–7.5)
2.3(1.5–3.4)...‡— Class II —6.1(4.3–8.6)
3.7(1.8–9.3)
4.6(3.3–6.7)
12.8(10.6–15.7)...‡
11.7(4.6–28.2)
3.4(2.3–5.2)
4.1(3.0–5.8)
6.5(3.7–12.1)
15.1(8.1–48.0)
4.2(2.8–6.1)
10.3(6.3–16.5)
0.8*
4.0(3.3–4.8)
1.9(0.8–3.4)
3.8(2.1–6.3)
2.9(2.7–3.1)
12.3(11.3–13.3)
0.9*
1.3*
7.5(4.2–12.3)
3.0(0.7–10.1)
1.1(0.5–1.7)
5.6(3.8–9.1)
3.0(1.5–5.8)
8.1(5.6–11.8)
2.2(1.2–3.6)
5.1(2.8–10.0)
26.9(11.0–121.3)
4.7*
3.5(2.2–5.5)
...‡
1.5*
3.1*
2.8(2.5–3.0)
9.0(20.8–16.3)
8.1(6.9–9.2)
17.3(35.9–33.4)
1.0*
1.3*
1.4(0.8–2.1)
1.7(0.8–2.9)
1.5(0.6–2.5)
2.1(1.1–3.4)
0.9(0.2–1.3)
2.7(1.9–4.0)
2.5(2.3–2.7)
9.1(8.4–10.0)
6.0(5.2–6.9)
15.8(13.7–18.2)
7.5(3.8–15.1)
4.4(1.7–7.6)

Q(cts ks−1 )0.8(0.5–1.0)k7.0(6.4–7.5)
10.7(10.0–11.3)
11.4(9.7–13.0)...
6.2(5.6–6.7)...
...
7.0(6.4–7.5)
14.6(13.6–15.6)
28.1(22.4–30.8)...
...
7.5(6.7–8.2)
2.0(1.7–2.3)...
6.4(5.9–7.0)
17.0(15.9–18.0)...
...
2.4(2.0–2.8)
10.6(10.0–11.3)
3.0*
0.7(0.5–0.8)k
1.0(0.8–1.1)
2.4(1.7–3.1)...
0.6(0.3–0.8)k
1.0(0.8–1.1)
0.9(0.7–1.1)
1.1(0.9–1.3)
1.1(0.8–1.3)
2.8(2.3–3.2)
3.8(3.4–4.2)

Comment

Flare1
Flare2
Flare1
Flare2
Flare3

Flare1
Flare2
Flare3
Flare1
Flare2
Flare1
Flare2
Flare3

Flare1
Flare2

5.7. THE CATALOGUE OF X-RAY SOURCES IN ρ OPH

55

Table 5.4 (Continued)
No.
BF-51
BF-59
BF-63
BF-66
BF-78
BF-87
BF-88

A-2†

tp
(ks)
5
61
88
54
98
4
84
14
46
84

A-41
A-63
A-81
BF-40
BF-46
BF-84
BF-96

...‡
70
54
76
22
62
76
90
52
17
88
72
72
32

A-12
BF-72

20
...‡

A-26

2
18
76
50
72
30
20
53

A-5
A-20

A-28
A-76
BF-55
BF-62

A-29

80

τr
τd
(ks)
(ks)
*
1.0
1.0*
*
1.5
1.3*
3.2(2.7–3.6)
7.6(6.2–9.5)...‡
6.9(5.3–8.9)
3.0(1.2–5.8)
0.5*
3.7(2.5–5.6)
7.9(6.8–9.2)
3.4(3.1–3.8)
6.0(5.4–6.7)
21.5(10.7–56.0)
7.0(3.6–15.6)
31.1(17.9–52.5)
22.6(14.7–32.7)
4.2(1.9–20.2)
3.0(1.3–9.1)— Class III —
...‡
131*
3.7(2.1-5.8)
55.6(30.8–117.5)
2.8(1.8–4.1)
6.1(4.0–9.5)*
2.3
7.0(2.6–22.9)
17.0*
8.3(3.4–31.3)
1.4*
2.7(2.7–2.7)*
4.2
9.9(5.6–20.8)
2.0*
...‡
2.6(1.3–4.3)
66.4*
1.8(1.4–2.2)
5.3(4.2–6.7)
1.8(1.2–2.7)
4.9(3.8–6.4)
3.6(2.8–4.7)
10.8(8.1–14.2)
5.5(2.7–12.3)
10.2(3.3–24.9)
59.2(41.9–96.7)
13.8(10.8–17.9)— Class IIIc —
0.9*
0.9*
...‡
4.1(2.1–6.8)— Unclassified NIR sources —3.1(1.0–57.9)
4.0(3.1–5.1)
1.3*
4.6(2.3–10.9)
...‡
20.1*
8.5(4.3–17.8)
23.8(9.1–55.4)*
1.9
24.0(12.2–129.1)
*
1.5
0.8*
3.7(2.6–5.0)
6.9(4.9–9.9)
5.7(5.0–6.5)
10.5(9.3–12.0)— Unidentified sources —
1.0*
2.1*

Q
(Cts ks−1 )0.6(0.4–0.8)k
...
2.0(1.7–2.3)
2.9(2.6–3.2)
2.4(2.1–2.7)
21.8(20.9–22.8)
1.0(0.8–1.2)
26.0(14.7–32.7)...
...

Comment...‡...
8.0(7.2–8.7)
...
27.3*
...
...
21.6(20.7–22.4)
10.3(9.5–11.1)
1.4(1.2–1.7)
9.7(9.1–10.2)
76.5(74.6–78.3)
1.0(0.7–1.2)
27.2(24.8–29.2)

Flare1
Flare2
Flare1
Flare2
Flare1
Flare2
Flare3

Flare1
Flare2

Flare1
Flare2
Flare3

1.3(1.1–1.5)
1.2(1.0–1.4)
1.4(1.0–1.7)...
...
0.8(0.3–1.1)
1.0(0.8–1.2)
0.7(0.5–0.9)k
1.2(0.5–1.8)
...
0.5*k

Flare1
Flare2
Flare3

Flare1
Flare2

56

CHAPTER 5. CHANDRA OBSERVATIONS OF THE ρ OPHIUCHI CLOUD
Table 5.4 (Continued)
No.
BF-36
BF-92
*

tp
(ks)
30
20

τr
(ks)
4.4(1.7–9.4)
0.5*

τd
(ks)
...‡
2.6(1.1–4.4)

Q(Cts ks−1 )0.8(0.6–0.9)
1.9*

Comment

We can not determine errors because of the limited statistics and/or strong coupling of the otherparameters.
†
Fitting is obtained for light curves in 2.00 5–12.00 5 (A-2) and 2.00 5–7.00 5 (BF-64) radius circles in order
to avoid the pileup effect.
‡
No fitting is possible because of the limited statistics and/or the boundary of the observations.
§
We relax the flare peak time (tp ) to be free in order to obtain as good results as possible.
kThis should be the upper limit because the fitting does not include the time bin with zero counts.

Chapter 6

Individual Sources
In this chapter, we show details of some interesting sources.

6.1

Protostars

6.1.1

A-49 – GY91

Fig. 6.1.— ACIS image around A-49 = GY91 in 0.5–
9.0 keV, smoothed with a Gaussian of σ = 4 pixels.
The solid rectangle represents the region for the X-ray
data of jet-like emission (see text). Many cloud cores(not given in this figure) are aligned along the direction
given by the dashed line (see Figure 3.1).

Figure 6.1 shows the ACIS image of a class I source GY91 (A-49). Besides significant Xrays from the central source, we can see weak diffuse jet-like emission to the north-east (the solid
rectangle in Figure 6.1) perpendicular to the cloud core alignment; many cloud cores are aligned
along the dashed line in Figure 6.1 (see also Figure 3.1). This X-ray “jet” is likely due to an outflow
from GY91, although the significance level is still low (∼2.3 σ). We make the spectrum from therectangle region and fit it by the MEKAL model. Then we constrain the plasma temperature to be57

58

CHAPTER 6. INDIVIDUAL SOURCES

higher than ∼1 keV. This corresponds to the shock velocity of > 500 km s−1 (Favata et al., 2002),
indicating surprisingly high jet activity. We note that such an extreme phenomenon has also beensuggested by the ASCA detection of Doppler-shifted Fe lines from a class I in the R CrA cloud(Koyama et al., 1996). To confirm this jet-like X-rays with deeper exposure observations would be
important to understand the outflow activity at the protostellar phase.

6.1.2

BF-26 – Elias29

A class I BF-26 = Elias29 exhibits an X-ray flare with a timescale of ∼10 ks (Figure A.1). < kT >and < LX > during the flare are 5.8 keV and 5.3×1030 ergs s−1 , while those in the quiescent phaseare 3.8 keV and 2.0×1030 ergs s−1 , respectively (Table 5.2). The ASCA satellite observed Elias29
twice with respective exposure times of 40 and 100 ks. In the first observation, it exhibited an
X-ray flare similar to the present observation (Kamata et al., 1997). However, it became quiescentwith < LX > = 3×1030 ergs s−1 in the second observation (Tsuboi et al., 2000). It therefore infers
that Elias29 shows the quiescent luminosity of ∼1030 ergs s−1 and exhibits frequent X-ray flares.

6.1.3

BF-50 – WL6

Kamata et al. (1997) claimed the detection of hard X-rays from a class I WL6 with ASCA, althoughnearby GY256, another class II source, may contaminate the X-rays from WL6. The ASCA light
curve shows a sinusoidal shape of about 20 hour period. Since they found no temperature variation,
which is typical of normal flares, they claimed that the modulation is due to stellar rotation. In
the present observation, we clearly detect hard X-rays not only from WL6 (BF-50) but also from
GY256 (BF-51). Two faint flares are found from GY256, with higher temperature than that in thequiescent phase (Table 5.2 and Figure A.1). However, we can see no sinusoidal light curve from
WL6 nor from GY256.

6.1.4

BF-61 – YLW15A

ROSAT/HRI detected an extraordinary large X-ray flare from a class I YLW15A (Grosso et al.,1997). The inferred luminosity, however, depends largely on the assumed spectrum. We hence reestimate the luminosity using our best-fit < kT > (4.2 keV) and NH (4.8×1022 cm−2 ) of YLW15A= BF-61 at a flare (Table 5.2). The HRI count rate at the flare peak of ∼17 cts ks−1 is thenconverted to the peak luminosity of ∼1033 ergs s−1 , which is about ten times brighter than thatof the giant flare seen in YLW16A (see §6.1.5) but may not be extraordinary, comparable to thelargest YSO flares detected with ASCA (V773 Tau: Tsuboi et al. 1998; ROXs31: Imanishi et al.

6.1. PROTOSTARS

59

2002a, see also §6.2). The subsequent ASCA observation discovered quasi-periodic three X-ray
flares with an interval of ∼20 hours (Tsuboi et al., 2000). Montmerle et al. (2000) interpreted
these flares with a star-disk arcade conjecture in contrast to the X-ray emission from that of TTSs,
which may have magnetic arcades at the stellar surface. Although Chandra confirms hard X-ray
emission and finds a typical flare from YLW15A (Figure A.1), no multiple flares are found. Thus
quasi-periodic flares are not always present but rather occasional phenomena.
Recently, VLA (Girart et al., 2000) and Subaru/COMICS (M. Honda, private communication)
resolved YLW15A into a binary system with the separation angle of ∼0.00 5. The position of BF-64
favors that the origin of X-rays is the radio-faint and MIR-bright source VLA2 rather than VLA1.

6.1.5

BF-64 – YLW16A

The light curve of a class I BF-64 = YLW16A comprises two flares; the first is rather complex withspike-like structure and the second is a giant flare having an unusual profile (Figure 6.2 left). We
examine spectral evolution by slicing the data in the time intervals as shown in Figure 6.2 (left). In
the phases 7–9, we extract the spectra from the same region (a 2.00 5–7.00 5 radius circle) as the light
curve data, because these phases suffer the pileup effect (see §5.5). We first fit them by the MEKAL
model allowing the abundance and NH to be free for all the phases, then find no significant variation∼ 0.3 solar) and NH (=∼ 5.2×1022 cm−2 ). We hence fix
from phase to phase both in abundance (=the abundance to be 0.3 solar and fit them assuming that NH is the same in all phases. Table 6.1
shows the resultant best-fit parameters for each time interval. As are seen in Figure 6.2 (left), thetemperature (kT ) and emission measure (EM ) increase from phase 2 and reach their peak values
at phase 3, then gradually decrease. The peak luminosity is estimated to be ∼1×1031 ergs s−1 .
These phenomena of the first flare are similar to flares found in other YSOs. The second flare shows
unusual time profiles in the flux and emission measure (EM ). The temperature profile is, however,more typical of normal flares; at phase 7 it increases rapidly and stays almost constant or shows
gradual decay. Such a property is similar to the giant flare of ROXs31 detected with ASCA (§6.2)
and we interpret that it is due to a large loop length (§8.3).
Figure 6.2 (right) is the time-averaged spectrum during the large flare. A remarkable findingis an additional emission line near the 6.7 keV line of highly ionized (He-like) irons. The best-fit
line center energy is 6.43±0.08 keV, which is attributable to neutral or low ionized irons. The mostplausible origin is fluorescence from cold irons in the circumstellar matter. Assuming the spherical
distribution of circumstellar gas around YLW16A, the equivalent width of iron (EW ) is estimatedto beNH
EW ≈ 10ZFe ( 22
) [eV],(6.1)
10 cm−2
where ZFe and NH are the abundance of iron and column density of the fluorescent medium,

60

CHAPTER 6. INDIVIDUAL SOURCES

respectively (Inoue, 1985). If we regard that most of the observed column density (5.2×1022 cm−2 )originates from the circumstellar matter and use the iron abundance of 0.3 solar, we predict the
maximum equivalent width to be ∼15 eV, which is significantly lower than the observed value of∼200 eV. We therefore require non-spherical geometry; a larger amount of matter should be present
out of the line-of-sight. A possible scenario is that YLW16A has a face-on geometry (Sekimoto etal., 1997) and the disk is responsible for the fluorescent 6.4 keV line. Since we see no time delay(reflection timescale) with & 104 s between the flare on-set and 6.4 keV iron line appearance, thedistance of the star and the reflecting region should be less than . 20 AU, which is typical of the
distance between the stellar surface and the inner edge of the disk (Feigelson & Montmerle, 1999).

Fig. 6.2.— (left): (top) Light curve of BF-64 = YLW16A and (middle and bottom) time profile ofthe best-fit temperature and the emission measure. (right): Spectrum of YLW16A in phases 7–9.
The upper panel shows data points (crosses) and the best-fit model (solid line), while the lower
panel shows residuals from the best-fit model. Line structure exists at ∼6.4 keV and ∼6.7 keV.

6.1. PROTOSTARS

61

Table 6.1: Best-fit parameters of BF64 = YLW16A
Phase*kT(keV)

log(EM )(cm−3 )

NH †(1022 cm−2 )

EW ‡
(eV)

Flux§

LX k

1 .....
2 .....
3 .....
4 .....
5 .....
6 .....
7 .....
8 .....
9 .....

3.2(2.3–4.8)
4.7(3.1–9.2)
5.4(4.1–7.7)
4.3(3.3–6.5)
3.1(2.6–4.1)
2.6(2.2–3.3)
8.1(6.4–11.2)
11.2(8.9–15.4)
8.6(7.1–10.7)

53.3(53.1–53.4)
53.5(53.4–53.6)
53.9(53.8–54.0)
53.8(53.7–53.9)
53.6(53.5–53.7)
53.4(53.3–53.5)
54.5(54.5–54.6)
54.9(54.9–54.9)
54.9(54.9–54.9)

5.2(5.0–5.4)
5.2
5.2
5.2
5.2
5.2
5.2
5.2
5.2

<217
<435
<59
<99
<97
<1466
181(78–284)
80(17–143)
<86

0.3
0.7
1.8
1.2
0.6
0.3
9.7
26.7
24.5

0.2
0.5
1.1
0.8
0.5
0.3
5.3
13.8
13.4

*

See Figure 6.2.
NH is assumed to be the same value in all the phases.
‡
Equivalent width for the 6.4 keV line.
§
Observed flux (10−12 ergs s−1 cm−2 ) in 0.5–9.0 keV.
kAbsorption-corrected luminosity (1031 ergs s−1 ) in 0.5–9.0 keV.
†

6.1.6

VLA1623

Since there is no significant X-rays from a class 0 prototype VLA1623 (see also §6.5), we estimate
its upper limit of < LX > assuming the MEKAL spectrum with the temperature of 1, 5, and 10 keV
and with the abundance of 0.3 solar. From the radio data (André et al., 1993), the column density
to VLA1623 is estimated to be NH = (0.5–4)×1023 cm−2 . We then calculate the upper limit of< LX > to be (0.78–78)×1029 ergs s−1 . Its uncertainty mainly comes from the difference of theassumed temperature (1, 5, and 10 keV). The upper limit is similar to that of faint class I sourcesin their quiescent phase (Figure 7.2a). The ratio of < LX > and radio intensity (388 µJy at 6 cm;
Leous et al. 1991) is estimated to be <1013 –1015 , which is also comparable to that of other YSOs(§7.4.4). It may therefore be conceivable that VLA 1623 is an X-ray source but was in a quiescentphase during the Chandra observation. Further deep X-ray observations would be fruitful.

62

6.2

CHAPTER 6. INDIVIDUAL SOURCES

X-ray Brightest T Tauri Stars (DoAr21, ROXs21, and ROXs31)

Imanishi et al. (2002a) showed the Chandra/ACIS and ASCA/GIS results for a series of four longterm observations of DoAr21 (A-2), ROXs21 (BF-46) and ROXs31 (BF-96), the X-ray brightest
TTSs in the ρ Ophiuchi cloud, then studied their long-term time variability and detailed structureof X-ray spectra. Here we show details.

6.2.1

Previous knowledge

The Einstein satellite first recognized these sources as bright X-ray sources (Montmerle et al.,1983). Subsequent optical observations revealed that they are K–M type stars having weak Hαemission (Bouvier & Appenzeller, 1992), which are classified as WTTSs (class III). Simon et al.
(1995) observed them by using the lunar occultation in the infrared wavelengths and found that
ROXs21 and ROXs31 are binary systems with separation angles of 0.00 3 and 0.00 48, respectively.
Shevchenko & Herbst (1998) reported the rotation period of ROXs21 to be 1.39 days. Since theyshow no continuum emission at 1.3 mm (André & Montmerle, 1994), the circumstellar envelope
would have already disappeared, while DoAr21 may still have an accretion disk as suggested by
the detection of NIR polarized emission (Ageorges et al., 1997). Using the theoretical evolutionary
tracks in the H-R diagram, the age of DoAr21 is estimated to be ∼105 yr, which is younger thanthat of ROXs21 and ROXs31 (∼106 yr: Nürnberger et al., 1998). Finally, DoAr21 and ROXs31
show variable centimeter emission, which is probably due to gyro-synchrotron mechanism inducedby surface magnetic field (Stine et al., 1988).

6.2.2

Long-term observations with Chandra and ASCA

We use the data of obs-BF for ROXs21 and ROXs31. Although DoAr21 suffers the pileup effect in
obs-A, obs-BF detected X-rays from DoAr21 at the ACIS-S2 chip, which mitigates photon pileup
because of the large off-axis angle (= large PSF size). We hence use the ACIS-S2 data for thefollowing analysis. The background region for DoAr21 is taken from a source-free 19 arcmin2 region
on ACIS-S2, while that for ROXs21 and ROXs31 is the same as used in the previous analyses.
In order to examine long-term variability and spectral evolution of these sources, we additionally use the previous ASCA observations. Three observations (obs-A1, A2, and A3) were carriedout with the two Gas Imaging Spectrometers (GISs: Ohashi et al., 1996) and the two Solid-state
Imaging Spectrometers (SISs: Burke et al., 1991) onboard ASCA, which are at the foci of the X-ray
telescopes (XRTs: Serlemitsos et al., 1995) sensitive to photons in 0.4–10 keV. However, they wereoutside of or at the edge of the SIS’s field of view in the majority of the observations, hence we use

6.2. X-RAY BRIGHTEST T TAURI STARS (DOAR21, ROXS21, AND ROXS31)

63

the GIS data only. We retrieve the ASCA unscreened data from the HEASARC online service1 ,then filtered out the data taken at a geomagnetic cutoff rigidity lower than 4 GV, at an elevation angle less than 5◦ from the Earth, and during passage through the South Atlantic Anomaly.
Particle events were also removed using the rise-time discrimination method. The total available
exposure times were ≈38 ks (obs-A1), 93 ks (A2), and 75 ks (A3) (Table 6.2).

Table 6.2: ASCA/GIS observation log

6.2.3

Obs.ID

Sequence ID

A1 . . .
A2 . . .
A3 . . .

20015010
25020000
96003000

Observation Date (UT)
Start
End
1993 Aug 20 02:20
1997 Mar 2 07:32
1998 Aug 13 06:35

Aug 20 22:46
Mar 4 21:21
Aug 15 21:16

Exposure
(ks)
37.7
93.1
74.7

Time variability

Figures 6.3 and 6.4 show the X-ray light curves (no background subtraction) of (a) DoAr21, (b)ROXs21, and (c) ROXs31 obtained with Chandra/ACIS (obs-BF) and ASCA/GIS (obs-A1, A2,
and A3), respectively. The vertical axis is photon counts normalized by the effective area at 1 keV
for comparison with the different instruments and observation epochs. The time intervals of eachdata point are 103 s and 104 s for ACIS (Figure 6.3) and GIS (Figure 6.4), respectively. Theenergy bands of the ACIS and GIS data for DoAr21 are 0.5–9.0 keV, while the GIS light curves for
ROXs21 and ROXs31 are limited in the soft energy band of 0.5–1.5 keV in order to avoid possible
contamination from a nearby variable hard X-ray source YLW15A = BF-61 (Tsuboi, 1999; Tsuboi
et al., 2000). To verify whether or not the contamination from YLW15A is significant in <1.5 keV,
we check the amount of contamination from YLW15A and find that the best-fit absorption of NH= 4.8×1022 cm−2 (Table 5.2) reduces the X-ray emission below 1.5 keV to less than 3 % of the
1.5–9 keV flux. Possible contamination from YLW15A to the GIS source circles of ROXs21 and
ROXs31 is therefore insignificant in the 0.5–1.5 keV band. No data is shown in the light curve of
ROXs21 (Figure 6.4b) from MJD = 50511.55–50512.20 due to significant contamination from an
exceptionally giant flare of ROXs31 (see Figure 6.4c).
From Figures 6.3 and 6.4, we detect six, three, and two flares from DoAr21, ROXs21 and
ROXs31, respectively. These are labeled as “F”, “F1”, “F2”, and “F3” with arrows in Figure 6.3
and 6.4, while the quiescent phases are indicated by “Q”. Among them, we find an extremely large1

http://legacy.gsfc.nasa.gov/cgi-bin/W3Browse/w3browse.pl.

64

CHAPTER 6. INDIVIDUAL SOURCES

flare from ROXs31 with the maximum flux of ∼100 times higher than the quiescent level (Figure
6.4c), which is separately treated in §6.2.7.

F1

Q

F2

Q

F

F

Q

Fig. 6.3.— Light curves of (a) DoAr21, (b) ROXs21, and (c) ROXs31 in 0.5–9.0 keV obtained with
Chandra/ACIS (obs-BF). The vertical axis is the X-ray count rate normalized by the effective areaof 100 cm2 at 1 keV. The lower arrows indicate the times of the flare and quiescent phases (seeTable 6.3).

A1

A2

F1

F2

A3

Q F3

Q

A1

F

A2

F1

F2

A3

Q

A1

A2

Q

A3

F

Fig. 6.4.— Light curves of (a) DoAr21 in 0.5–9.0 keV and of (b) ROXs21 and (c) ROXs 31 in
0.5–1.5 keV, obtained with ASCA/GIS (obs-A1, A2, and A3). The vertical axis is the X-ray countrate normalized by the effective area of 100 cm2 at 1 keV. The lower arrows indicate the times ofthe flare and quiescent phases.

6.2.4

X-ray spectra

Figure 6.5 shows the background-subtracted ACIS spectra using all data of obs-BF. The spectra
show several emission lines from highly ionized elements such as O, Ne, Na, Mg, Si, S, Ar, Ca, andFe, which is characteristic of a thin thermal plasma. We hence fit them by the MEKAL modelallowing the abundance of each element with prominent lines (indicated by arrows in Figure 6.5)
to be free. For the study of time variation of the spectra, we separately fit the flare and quiescent
spectra as shown in Figure 6.3. A single-temperature (1-T) model is not acceptable for all spectra6.2. X-RAY BRIGHTEST T TAURI STARS (DOAR21, ROXS21, AND ROXS31)

65

except for DoAr21 and the flare phase of ROXs31. We hence fit them by a two-temperature (2T) plasma model for ROXs21 and ROXs 31 and find acceptable fits. Table 6.3 shows the best-fit
parameters. All elemental abundances are below the solar photospheric values (Anders & Grevesse,
1989) except for Ne in DoAr21 and Na in ROXs21. Figure 6.6 presents the best-fit abundanceswith the order of their first ionization potential (FIP).

Mg
Ne
Fe-L

Si
S

O
Ar

Ne
Fe-L

Na

Ne
Fe-L

Ca

Fe

Fig. 6.5.— Time-averaged X-ray spectra of (a) DoAr21, (b) ROXs21, and (c) ROXs31. Arrows
indicate prominent spectral features with the corresponding elements.

low-FIPhigh-FIP

low-FIPhigh-FIP

low-FIPhigh-FIP

Fig. 6.6.— Elemental abundances of (a) DoAr21, (b) ROXs21, and (c) ROXs31. Errors indicate
the 90 % confidence limits. The horizontal axis shows corresponding elements with their FIP values(eV). Gray squares, black-open, and black-filled circles represent abundances in the quiescent, firstand second flare phases, respectively.

66

CHAPTER 6. INDIVIDUAL SOURCES

Table 6.3: Best-fit parameters of DoAr21, ROXs21, and ROXs31
Source*Phase†< kT1 >
(keV)
< kT2 >
(keV)
NH
(1022 cm−2 )Abundances(solar)O ...............
Ne. . . . . . . . . . . . . . .
Na . . . . . . . . . . . . . .
Mg . . . . . . . . . . . . . .
Si . . . . . . . . . . . . . . .
S................
Ar . . . . . . . . . . . . . . .
Ca . . . . . . . . . . . . . .
Fe . . . . . . . . . . . . . . .
Others. . . . . . . . . . .

< LX > k(1030 ergs s−1 )reduced-χ2 (d.o.f.)
*

Q‡
2.5
(2.4–2.6)···— DoAr21 —F1
2.9(2.7–3.1)
···

F2‡
2.9
(2.8–3.1)···

1.0(0.97–1.1)

1.0(0.96–1.1)

1.10(1.06–1.13)

···

§

···

§

···

§

1.0
(0.58–1.8)
··· §

0.99
(0.45–1.8)
··· §

1.5
(0.98–2.2)
··· §

0.13
(<0.39)0.18(0.078–0.35)
0.37(0.22–0.65)
0.13(<0.43)0.17(<0.55)0.11(0.057–0.17)
0.13(<1.1)

0.39(0.11–0.71)
0.12(<0.26)0.24(0.071–0.42)
0.69(0.27–1.3)
0.24(<0.73)0.10(0.051–0.18)
0(<0.93)0.23(0.050–0.45)
0.34(0.23–0.50)
0.32(0.19–0.50)
0.54(0.27–0.88)
0.41(0.10–0.74)
0.14(0.097–0.22)
0.13(<0.84)28

34

0.953(175)0.910(150)— ROXs21 —
Q
F
0.69
0.86
(0.65–0.74)(0.78–4.5)
1.8
2.6(1.6–2.1)(2.3–3.1)
0.13
0.081
(0.079–0.16) (0.038–0.11)— ROXs31 —
Q
F
0.23
0.27
(0.15–0.31)(0.21–0.41)
1.5
2.0(1.2–1.7)(1.8–2.5)
1.7
1.7(1.1–2.1)(1.0–2.5)

0.21(0.086–0.37)
0.75(0.53–1.1)
3.7(1.4–8.2)
··· §

···

0.24
(0.058–0.52)
0.48(0.22–0.96)
2.6(<7.6)
··· §

§

0.081
(0.015–0.22)
··· §

···

§

0.22
(0.056–2.1)
··· §

···

§

···

§

···

§

···

§

···

§

···

§

···

§

···

§

···

§

···

§

···

§

···

§

···

§

···

§

···

§

···

§

···

§

···

§

0.092
(0.056–0.16)
0.14(0.065–0.22)

0.12(0.067–0.20)
0.080(<0.21)0.019(<0.064)0.0080(<0.049)

0.049(<0.18)0.11(0.022–0.34)

38

1.5

1.9

22

30

1.03(275)

1.12(96)1.07(88)

1.11(67)

0.931(85)

The 1-T model for DoAr21 and the 2-T model for ROXs 21 and ROXs31 are used (see text). Parentheses indicate
the 90 % confidence limits.
†
See Figure 6.3.
‡
The boundary between phases Q and F2 is not clear because of the existence of a rising part (Figure 6.3), however
there are no significant differences in parameters when the boundary is changed.
§
These are assumed to be the same as “Other” elements.
k
Absorption-corrected luminosity in 0.5–9.0 keV.

6.2. X-RAY BRIGHTEST T TAURI STARS (DOAR21, ROXS21, AND ROXS31)

67

We investigate the Na overabundance seen in ROXs21 in further detail (Figure 6.6b), because
Na is a relatively rare element compared with the major 4N nuclei (12 C, 16 O, etc) and because the
Na Lyman-α line lies near the Lyman-β line of the more abundant element Ne. Figure 6.7a shows
the best-fit result when the Na abundance is collectively varied with the other elements. There issignificant data excess at ≈1.2 keV, near the Lyman-β line of Ne (1.211 keV). The intensity ratioof Lyman-β to -α lines of Ne should increase as a function of kT for kT < 1 keV and stay almost
constant at ≈ 7 for kT > 1 keV (Mewe et al., 1985), hence 1/7 of the Lyman-α flux should be an
upper-limit to the Lyman-β line flux. Since the flux of the Ne Lyman-α line is fairly strong, theLyman-β line flux is well constrained as seen in Fig 6.7a. Thus the excess flux should originate
from another element; the Lyman-α line of Na (1.236 keV) is the best candidate. Allowing the Naabundance to be a free parameter improves the fit as demonstrated in Figure 6.7b. The reductionof χ2 (d.o.f.) from 114 (97) to 108 (96) is significant with ∼98 % confidence level for the F -test2(Bevington & Robinson, 1992).

Ne
He-like Kα

Ne LyαNe LyβNa LyαFig. 6.7.— Close-up view of the fitting results of ROXs21 in phase Q, in which the Na abundance(a) is varied collectively with the “other” elements and (b) is treated as a free parameter. Arrows
and dashed squares show the positions of K-shell lines from (a)Ne/(b)Na and the prominent featurenear the Lyman-α line of Na, respectively.
2

A statistical test to check the difference between the fitted and parent functions. See §11.4 in Bevington &Robinson (1992).

68

CHAPTER 6. INDIVIDUAL SOURCES

Using the GIS spectra of DoAr21, we investigate the long-term (∼yr) behavior of the spectral
parameters. Since the energy resolution does not allow us to separate the emission lines from the
key elements, we varied the abundances assuming the same value for all elements. The best-fit
luminosity shows long-term variability in the range of (2–8)×1031 ergs s−1 . The best-fit temperatures and abundances also vary from 2.5 to 4 keV and 0.15 to 0.5 solar, respectively. Both arepossibly correlated with the X-ray luminosity as shown in Figure 6.8(a) and 6.8(b). The NH value
is consistent with being constant within a factor of ≈ 2.

< >

(b)

< >

(a)

< >

Fig. 6.8.— X-ray luminosity plotted against (a) the plasma temperature and (b) the mean abundance of DoAr21. Circles and triangles represent the data of ACIS and of GIS. Errors indicate the
90 % confidence limits.

6.2.5

Temperature structure

The high-quality spectra of the long Chandra/ACIS exposure reveal that the 2-T model gives abetter fit than the 1-T model for ROXs21 and ROXs31. This supports previous 2-T model fits forsome fractions of other bright TTSs (Carkner et al., 1996; Yamauchi et al., 1996; Preibisch, 1997;
Costa et al., 2000; Ozawa, 2000; Tsujimoto et al., 2001). DoAr21, on the other hand, displays asimple 1-T spectrum with < kT > ∼ 3 keV; it has no additional soft component. Since the age of
DoAr21 (∼105 yr) is younger than those of ROXs21 and ROXs31 (∼106 yr: Nürnberger et al., 1998),
coupled with the result of Tsujimoto et al. (2001) that 2-T spectra are found more often in older
TTSs than in younger protostars, we speculate that the soft component (∼0.5 keV), probably a
relatively steady corona, appears gradually as the system increases its age, finally reaching solar-like
corona. In this scenario, the hard component (2–3 keV) would be the sum of unresolved flares.

6.2. X-RAY BRIGHTEST T TAURI STARS (DOAR21, ROXS21, AND ROXS31)

6.2.6

69

Chemical composition of X-ray emitting plasma – the FIP and IFIP effects

Generally, the abundances of a stellar corona of MS stars may be different from those in a photosphere, depending on the FIP of the relevant elements. For instance, in the solar corona, elementswith low FIP are overabundant relative to the solar photospheric values (“the FIP effect”, Feldman
1992). Recent observations, however, showed that in some of stellar coronae the abundances ofhigh-FIP elements are larger than the others (the Inverse FIP = IFIP effect). The IFIP effectappears in X-ray active stars, while the FIP effect is generally seen in less active stars (Audard
et al., 2001; Brinkman et al., 2001; Güdel et al., 2001a,b,c; Huenemoerder et al., 2001; Kastner etal., 2002). The dependence of abundance on X-ray activity is also existent within individual stars;
some sources show the enhancements in abundance from quiescent to flare phases, or even within
a single flare phase (e.g., HR1099, Audard et al. 2001; V773 Tau, Tsuboi et al. 1998). Due to thelimited statistics and samples, however, it is unclear whether or not abundance variations between
quiescent and flare phases are common and the FIP/IFIP effects are present in low-mass YSOs.
The coronal abundances of DoAr21, ROXs21, and ROXs31 are sub-solar, consistent with theaverage values of low-mass YSOs (0.3 solar: Figure 5.4). From Figure 6.6(a) and (b), we see thatboth high-FIP (Ne and Ar) and low-FIP (Na, Mg, and Ca) elements show higher abundances thanthe other elements (the IFIP and FIP effects).
For the abundances in the solar corona, the FIP effect appears in elements with FIPs below
10 eV. A possible interpretation for the FIP effect is as follows (Feldman, 1992). The low-FIP
elements (FIP < 10 eV) are collisionally ionized in the chromosphere at 6000–7000 K temperatureand electric fields would preferentially transfer them to the upper coronal region. DoAr21 andROXs21 are a K0 star and a binary of K4 with M2.5, respectively, hence have surface temperaturesof 4000–5000 K, which is 0.6–0.8 times of that of the sun (Nürnberger et al., 1998). If we simply
assume that the FIP-effect energy-limit is a linear function of the surface temperature, it should
be shifted from 10 eV to 6–8 eV, which is near Mg (FIP = 7.8 eV) and Ca (6.1 eV), but well above
Na (5.1 keV). The abundance enhancements of Mg, Ca and Na provide independent evidence
supporting the universality of the scenario proposed by Feldman (1992).
The X-ray luminosities of our TTS samples are a few × 1030−31 ergs s−1 , which is significantly
higher than typical FIP sources of ≤ 1029 ergs s−1 (such as the sun, π 1 UMa, and χ1 Ori; Güdel etal. 2001c). The high luminosity may also be responsible for the IFIP effect we observed. A related
issue would be the dramatic abundance increase in the giant flare of ROXs31 (§6.2.7).
Güdel et al. (2001c) proposed a possible scenario for the IFIP effect. During the X-ray active
phase, the high-energy non-thermal electrons (§2.3.1) produce a “downward-pointing” electric field,
which suppresses the transfer of the ionized low-FIP elements into the corona. The electrons also

70

CHAPTER 6. INDIVIDUAL SOURCES

lose the kinetic energy at the photosphere by collisional process, then heat up cool elements to aplasma temperature, making the abundance increase during the flare. The radio gyro-synchrotron
emission from DoAr 21 and ROXs31 further supports this idea, because this is the direct evidence
for the presence of the high-energy electrons. Since the abundance increase should be correlatedwith the non-thermal radio flux variations, simultaneous radio and X-ray observations would help
to constrain physical models.
Another possibility for the IFIP effect (ionization mechanism of high-FIP elements) is photoionization by coronal X-rays, which is proposed to explain a solar flare in which Ne is enhanced
relative to O (Shemi, 1991; Schmelz, 1993). Since sources with the IFIP effect have large X-ray
luminosity of < LX > = 1029 –1032 ergs s−1 (HR 1099, Brinkman et al. 2001; AB Dor, Güdel et al.
2001a; YY Gem, Güdel et al. 2001b; II Pegasi, Huenemoerder et al. 2001; and TW Hydrae, Kastneret al. 2002), they should suffer large X-ray photo-ionization at the photosphere, which increases
the ionization fraction of high-FIP elements. This may also well explain the abundance increase in
the X-ray active phase (Figure 6.8b). We should note, however, that this simple logic can work forall (both FIP and IFIP) elements, hence does not cause the enhancement of the ionization fractionof high-FIP elements. In order to solve the IFIP effect considering the effect of photo-ionization,other constraints (e.g., geometrical effect) should be considered.

6.2.7

Giant flare from ROXs31

For the detailed study of spectral evolution during the flare, we make and fit time-sliced GIS spectraof the giant flare from ROXs31 in obs-A2 (Figure 6.4c). The time intervals are shown in Figure
6.9 (left). The spectra show strong Kα emission of He-like Fe at ≈ 6.7 keV, hence we let the Feabundance to vary. The flare spectra are well fitted by a 1-T model in contrast to the quiescentphase (2-T model). Since the flux of the soft component in the quiescent phase is 10–100 times
lower than that of the flare flux, it is highly possible that the soft component does exist during theflare but is hidden by the more luminous harder component.
Figure 6.9 (right) shows the best-fit parameters for each time interval. As shown in the figure,this flare has an unusual time profile. The light curve increases slowly over ∼15 ks (see also Skinner2000). The temperature reaches its peak value before the flux maximum and stays nearly constantduring the decay phase.
The absorption (NH ) and abundances dramatically change from quiescent to flare; NH in theflare is ∼7.1×1021 cm−2 , about two times smaller than that of the quiescent phase (obs-BF) of
1.7×1022 cm−2 , while the abundances increase by a factor of 3–30 from quiescent (0.01–0.1 solar)
to the giant flare (0.26–0.29). They do not change during the flare. This could be due to thechromospheric evaporation by high-energy electrons, which may be produced just after magnetic

6.2. X-RAY BRIGHTEST T TAURI STARS (DOAR21, ROXS21, AND ROXS31)

71

reconnection. Evaporation by the flare X-rays may not be important, because the abundances andabsorption do not change much during the flare. Thus, in the case of the IFIP effect in the giant
flare of ROXs31, the high-energy electron origin is more likely than X-ray irradiation scenario.

1

2

3

4

5

6

7

Lx (1032ergs s-1) Others (solar)

Phase

Fe (solar)

NH (1022cm-2)

kT (keV)The peak luminosity of ∼1033 ergs s−1 and temperature of ∼10 keV are comparable to thegiant flare of another TTS binary system V773 Tau (Tsuboi et al., 1998). However, the time profileof the plasma temperature is different from that of V773 Tau. The rise timescale of ROXs31 ismuch longer (∼3 ks). This may be due to a large magnetic loop (§8.7), although the effect of thepartial occultation of the flare due to the stellar rotation may also be conceivable. The temperaturearrives at the peak value preceding the flare peak and stays nearly constant at 4–5 keV in the decay
phase. These imply the existence of successive energy release due to the magnetic reconnection
and that the generation and heating of the plasma occur even in the decay phase, as suggested by
Tsuboi et al. (1998); Schmitt & Favata (1999); Favata & Schmitt (1999). Such a phenomenon is
also seen in long duration event (LDE) flares of the sun (Isobe et al., 2002).

Fig. 6.9.— (left) A magnified view of the light curve of ROXs31 during the giant flare in 0.5–9.0 keV(Figure 6.4c). (right) Time profiles of the best-fit temperature, absorption column, abundances ofFe and other elements, and X-ray luminosity of ROXs31 during the giant flare in obs A2. Errors
indicate the 90 % confidence limits.

72

CHAPTER 6. INDIVIDUAL SOURCES

6.3

Other T Tauri Stars

6.3.1

A-20 – GY12

A class III source GY12 (GSS30-IRS2) shows three flares in obs-A (Figure A.3). GY12 was also
detected at the ACIS-S2 chip in obs-BF, hence we make light curve and time-sliced spectra as
22 cm−2 is the same as A-20, the X-ray emitting
shown in Figure 6.10. Since NH of 3.6+0.1−0.2 ×10
source in obs-BF is surely GY12 regardless of the large positional uncertainty. In Figure 6.10, we
detect four flares with our criterion (§5.5). We therefore conclude that GY12 is an X-ray active
source showing frequent flare activity, 3–4 flares per 100 ks.
Near GY12, two class I sources GSS30-IRS1 and LFAM1 = GSS30-IRS3 are located with therespective separation angles of ∼1700 and ∼1000 from GY12. Also, CO observations revealed the
expanding shell around them with the radius and expanding velocity of ∼1500 (∼= 2000 AU) and−1∼0.7 km s (Zhang et al., 1997), which is probably due to cloud disruption as a result of youngstar formation. Although we see no significant X-rays from the two class I sources and from theshell for both observations, it may be conceivable for future works whether or not such complex
systems affect the high flare frequency of GY12.

Fig. 6.10.— (top) Light curve in 0.5–9.0 keV
and (middle and bottom) time evolution ofkT and EM of GY12 in obs-BF. Flare timebins (§5.5) are represented by arrows.

6.3. OTHER T TAURI STARS

6.3.2

73

A-23 – GY21

A class II A-23 = GY21 shows two flares in obs-A (Figure A.1). The spectra show strong He-like Fe emission at∼6.7 keV both in the flare and quiescent phases (Figure
6.11). We hence fit them by the MEKAL model with theabundance to be free and obtain significantly higher values than the other sources; 1.0 (0.5–1.6), 3.5 (>0.6), and<15 solar for the quiescent, first flare, and second flarephases, respectively. GY21 thus shows not only overabundance but also a possible abundance increase during the flare like HR1099 (Audard et al., 2001), V773
Tau (Tsuboi et al., 1998), and ROXs31 (Imanishi et al.,Fig. 6.11.— Time-averaged spectrum of
2002a, see also 6.2.7), although uncertainties are large.
A-23 = GY21. Prominent line emission
The expected magnetic field of the first flare (see §8.3.2)
is seen at ∼6.7 keV.
is among the largest values (&1000 G). Since it has been
considered that the abundance anomaly is related to the activity at the stellar surface (e.g., Güdelet al., 2001c), it would be fruitful for future works to investigate whether or not such a large magnetic field has any correlation with its abundance anomaly. Possible binary nature of GY21 andS2 = A-24 (Haisch et al., 2002) would also be conceivable (see also §6.3.3).

6.3.3

A-24 – S2

A class II source S2 = A-24 exhibits a clear flare with
about ten times larger flux than that in the quiescentphase (Figure A.2). Similar to GY21 (§6.3.2), the spectra of S2 show strong He-like Fe emission (Figure 6.12).
We fit them by the MEKAL model with the abundance
to be free and obtain 0.6 (0.4–0.8) and 0.7 (0.5–0.9) solar
for the quiescent and flare phases, respectively. S2 wasalso detected on ACIS-S2 in obs-BF. Although the light
curve in obs-BF shows no flare, the time-averaged < kT >30 ergs s−1 ) are signif(2.9+0.4−0.3 keV) and < LX > (9.1×10
icantly larger than those in the quiescent phase of obs-A,
Fig. 6.12.— Time-averaged spectrum of
while the abundance is lower (0.2±0.1 solar). Since theA-24 = S2. Prominent line emission is
22 cm−2 ), it isabsorption is nearly the same (4.2+0.2−0.3 ×10
seen at ∼ 6.7 keV.
less possible that another source emits X-rays in obs-BF

74

CHAPTER 6. INDIVIDUAL SOURCES

despite of large positional uncertainty. S2 therefore has the large variation of < LX > = 1030 –
1031 ergs s−1 together with that of the abundance = 0.2–0.7 solar, although its mechanism is an
open question.
Interestingly, both members of a candidate wide binary system, GY21 and S2 (Haisch et al.
2002; see also Figure 6.18) show overabundance. This may indicate anomaly of global chemicalcomposition of this region, although this interpretation is insufficient to explain the abundance
variations of GY21 and S2.

6.3.4

A-34 – GY51

The offset between A-34 and the nearest class II source GY51 is ∼0.00 72, which is slightly larger thanthe positional uncertainty of the Chandra observation (∼0.00 5) at its small off-axis angle (∼1.0 1, see
Figure 5.3). However, the NH /AV and < LX >/Lbol ratios between A-34 and GY51 are respectively∼1.7×1021 cm−2 mag−1 and ∼10−3 , which are typical of YSOs (Figure 7.6 and 7.7). We hence
regard that A-34 is surely associated with GY51. A detection of a solar-like flare (Figure A.1) isalso consistent with out idea.

6.3.5

A-41 – S1

A class III S1 is among the brightest and hottest sources in ρ Oph (Lbol ∼ 1000 L¯ , Teff ∼20000 K), which is classified as “Herbig Ae/Be stars” (HAeBes = intermediate-mass YSOs: André
et al., 1988). André et al. (1988) detected polarized non-thermal radio emission, which probably
comes from gyro-synchrotron particles. This indicates that S1 has relatively strong magnetic field.
In obs-A, we detect significant X-rays (A-41) with a small-amplitude flare near the end of theobservation (Figure A.3). S1 was also detected at the ACIS-S2 chip in obs-BF (Hamaguchi et al.,2002). The light curve of S1 in obs-BF shows no apparent flare, although it seems to decrease
throughout the observation with the e-folding time of ∼40 ks.
Figure 6.13 shows time-averaged spectra of S1 in (left) obs-A and (right) obs-BF. Except for a
marginal hump at ∼6.5 keV in obs-BF, clear emission lines can not be seen, indicating either low
metal abundance or non-thermal nature. Also, there is possible edge absorption feature around4 keV in obs-A. We hence fit them by the power-law model, as well as by the MEKAL model with
abundance to be free. For obs-A, we separately make and fit the spectra in the flare and quiescentphases. Table 6.4 summarizes the fitting results. Both the thermal and non-thermal models well
reproduce the spectra. The derived abundance of the thermal model is relatively low of 0–0.2 solar.
For obs-A, we additionally apply the fitting with an edge absorption, which slightly reduces the6.3. OTHER T TAURI STARS

75

χ2 values. This reduction is significant with ∼96 (thermal) and ∼99 (non-thermal) % confidence
levels for the F -test (Bevington & Robinson, 1992).

Fig. 6.13.— Time-averaged spectra of S1 in (left) obs-A and (right) obs-BF. The upper panels
show data points (crosses) and the best-fit MEKAL model (solid line), while the lower panels
show residuals from the best-fit model. Arrows indicate (left) possible edge absorption and (right)
marginal hump (see text).

If X-rays are thermal in origin, < kT > (∼2–3 keV) and the < LX > /Lbol ratio (∼10−6 ) aresimilar to those of HAeBes (Hamaguchi, 2001). Flare-like time variability is also consistent. If nonthermal X-rays are the case, on the other hand, most plausible scenario is bremsstrahlung emission
from non-thermal electrons accelerated just after the magnetic reconnection event (§2.3.1), whichis the same mechanism as hard (>10 keV) X-rays of the sun (§2.3.1). This is further supportedby the radio observation which suggests the presence of MeV electrons around S1 (André et al.,1988). However, this requires that thermal plasma generated via magnetic reconnection shouldhave lower temperature than those of the sun, because the non-thermal X-rays are dominant over
the thermal component above ∼1 keV. To verify which is the real case, more sensitive observations(for example, with XMM-Newton) would be needed.
The threshold energy of the edge absorption is ∼3.9 keV, which corresponds to neutral Ca(4.04 keV) or ionized Ar (Ar XVI: 3.95 keV) with the temperature of ∼106.6 K. However, theformer is less possible, because it requires the extremely large Ca abundance of ∼500 solar. In
the case of Ar, light elements below N are fully ionized and do not affect the absorption, hence
the unusually large abundance of Ar does not needed. It is therefore probable that the warm
absorber makes Ar edge feature. Considering that the edge feature is not seen at the spectrum in
obs-BF, the absorber should have much smaller size than the stellar radius and be located near thestellar surface, if we interpret the disappearance of the edge feature in obs-BF as the effect of the

76

CHAPTER 6. INDIVIDUAL SOURCES

occultation. However, its detailed physical condition is still in discussion. More details about S1
are discussed in Hamaguchi et al. (2002).

Table 6.4: Best-fit parameters of S1
Model*< kT > /Γ†‡(keV/–)

Abundance†(solar)< EM >†(1053 cm−3 )

1.9(1.7–2.1)
2.7(2.1–3.0)
2.2(2.1–2.4)
3.4(2.9–4.1)
3.3**(3.1–3.5)···
2.9**(2.6–3.0)···

0.20**(0.05–0.39)
···

···

2.6
(2.2–3.2)
2.3(1.9–2.9)
2.3(2.2–2.3)
2.0(1.9–2.3)
···

···
···

···
···

···

···

NH †
(1022 cm−2 )Eedge †§
(keV)

τ †k

< LX >#

χ2 /d.o.f.

2.2**(2.1–2.4)
···

···

···

2.2

138.6/126

···

···

2.3

···

2.1**
(2.0–2.2)
···

3.9**
(3.8–4.0)···

0.47**(0.29–0.66)
···

2.0

126.8/124

2.3

···

2.8**
(2.6–3.0)···
2.6**(2.2–2.7)
···

···

···

5.8

142.2/128

···
3.9**
(3.8–4.1)······
0.56**(0.37–0.76)
···

6.6
4.1

···
127.3/126

4.8

···

···

···

2.6

117.6/105

···

···

7.2

116.3/106– obs-A –Mekal(Q)(F)
Mekal*Edge(Q)(F)
Power(Q)(F)
Power*Edge(Q)(F)

0.16**(0.03–0.33)···– obs-BF –
Mekal
Power

1.6
(1.5–1.9)
3.4(3.1–3.6)

0.0(<0.11)···

4.1(3.4–5.0)···

1.9(1.7–2.0)
2.4(2.2–2.6)*

Applied models; Mekal: thin-thermal plasma model, Power: power-law model, Edge: edge absorption. “Q” and“F” for obs-A denote the spectra in the quiescent and flare phases, respectively.
†
Parentheses indicate the 90 % confidence limit.
‡
Best-fit temperature and photon index for the Mekal and Power model, respectively.
§Threshold energy of the edge absorption.
k
Absorption depth at Eedge .
#Absorption-corrected X-ray luminosity (1030 ergs s−1 ) in 0.5–9.0 keV.
**
Parameters are derived by assuming the same values for the quiescent and flare spectra.

6.3.6

A-79/BF-17 – GY195

A class II source GY195 shows significant X-rays in both observations (A-79 and BF-17), withthree flares (Figure A.2). In particular, GY195 shows unusual flares in obs-A. First, a sharp flareoccurs at t ∼ 35 ks, then reaches its peak about 50 times larger than the pre-flare flux. During
the decay phase (t ∼ 50 ks), a second flare appears but the rise and decay timescales are larger.
Gagné (2001) interpreted this series of flares as a result of a re-heating event of the single loop, like

6.3. OTHER T TAURI STARS

77

quasi-periodic flares from YLW15A (Tsuboi et al., 2000). We hence make and fit the time sliced
spectra of A-79 as shown in Figure 6.14, however no significant sign of increase of kT is found at
the second flare but it keeps almost constant of 3–4 keV (phases 3–5 in Figure 6.14), which is a
sharp contrast to the case of YLW15A (Figure 2 in Tsuboi et al., 2000).

6.3.7

BF-10 – WL22

The X-ray spectrum of a class II BF-10 = WL22 exhibits
clear line structure at ∼3.2 and ∼3.9 keV (Figure 6.15).
The line energies correspond to He-like Ca and Ar. We
hence relax the abundances of these elements to be free
and fit it by the MEKAL model. The best-fit < kT >and NH are 2.1 keV and 12×1022 cm−2 , respectively (Table 5.2), while the abundances of Ar and Ca are 2.7+8.3−2.7
and 5.0+7.9
solar.
ThusCa
isextremely
overabundant,−4.1
although the mechanism is an open question.

Fig. 6.14.— (top) Light curve and(middle and bottom) time evolution ofkT and EM of A-79 = GY195.
Fig. 6.15.— Spectrum of BF10 = WL22. The
upper panel shows data points (crosses) and thebest-fit MEKAL model (solid line), while thelower panel shows residuals from the best-fit
model. The line structure at ∼3.2 keV and ∼3.9
keV is due to He-like Ar and Ca.

78

CHAPTER 6. INDIVIDUAL SOURCES

6.4

Brown Dwarfs

Imanishi et al. (2001b) did a systematic X-ray search from bona-fide and candidate brown dwarfs
in ρ Oph. Here, we show details.

6.4.1

X-ray detection from brown dwarfs in ρ Oph

In the fields of view of obs-A and obs-BF, 18 late M dwarfs have been reported based on the water
vapor absorption at λ = 2.4–2.5 µm (Wilking et al., 1999; Cushing et al., 2000). Among them,eight sources have the upper limit of the mass less than 0.08 M¯ , which we call “bona-fide browndwarfs” (here and after, BDs). Other seven sources have the lower limit of the mass less than
0.08 M¯ , which is called “candidate brown dwarfs” (BDc s)3 . Table 6.5 shows their names andspectral types.
In the X-ray source catalogue (Table 5.2), we found X-ray counterparts for four BDs and BDc s(A-30 = GY37, A-38 = GY59, BF-86 = GY310, and BF-S2 = GY141). Also, a BDc GY5 may
be associated with a soft-band source A-S2, although Imanishi et al. (2001b) reported no X-ray
emission from GY5 by using the NIR position of Barsony et al. (1997), which is ∼200 away fromthe 2MASS source. Since no apparent X-ray sources are found in the error radius of the other 11
catalogued BDs and BDc s, we define a circle with a half radius of the PSF size around each IR
position, then manually count X-ray photons. We note that a rather small source radius is selectedso as to increase the signal-to-noise (S/N) ratio, particularly for faint X-ray sources. Nevertheless,
as is demonstrated in the bright sources, the position error between IR and X-ray is always smaller
than the source radius, because both have generally similar dependence on a source off-axis angle;
both have the smaller values for sources with smaller off-axis angle. Therefore, most of the X-ray
photons for the relevant BDs and BDc s, if any, may fall in the source circles. Table 6.5 shows the
X-ray counts thus counted. The mean background counts are estimated from the same regions usedin §5.6. The soft band (0.5–2.0 keV) background counts (in units of 10−2 counts arcsec−2 ) are 2.3
and 2.2 for obs-BF and obs-A, respectively, while those in the hard (2.0–9.0 keV) band are about
three times larger, 6.8 and 6.5 for obs-BF and obs-A. The background counts in each source area
are given in Table 6.5. We then separately calculate the confidence level (CL) of X-ray detection
for the soft, hard, and total (0.5–9.0 keV) bands. Based on the Poisson statistics, we define the CL
as
0
NX
0 −1
NN
−Nbg bg
CL =
e
,
(6.2)
N 0!
0
N =0

3

Imanishi et al. (2001b) regarded GY31, GY163, and GY326 as BDc s. Now, however, the mass of these M dwarfs
is regarded to be higher than the hydrogen burning limit (Wilking et al., 1999), hence we do not include them in the
BDc sample.

6.4. BROWN DWARFS

79

where N0 and Nbg are the detected and the background counts in the source circle, respectively(see e.g., Gehrels, 1986, and references therein). We set a detection (or an upper limit) criterion
that the CL should be larger than 0.999 (= 3.3 σ) in any of the three energy band data. In Table
6.5, we show the maximum CL value for each source. For sources located in both fields of view
(GY141 and GY202), we also estimate the CL for the combined data as well as the separate data.
However, no higher CL value is obtained from all the sources. In addition to the seven X-ray
sources, a BDc GY84 shows significant X-rays, although it may be possible that these X-rays arepart of diffuse X-ray emission A-48 (see §6.6.2). We hence regard that the X-ray detection from
GY84 is marginal.

Table 6.5: Bona-fide and candidate brown dwarfs in the ρ Oph cloud cores
X-ray countsk
1 − CL#LX **log(LX /Lbol ) Ref‡‡
Soft
Hard BGD— Brown Dwarf (BD) —CRBR14
M7.5
A
···
5.7
4
0
0.5
ND (S)<0.55< −4.3 1
GY10
M8.5
A
···
3.2
2
0
0.2
ND (S)<0.48< −4.5 1
GY11
M6.5
A
···
3.3
1
1
0.2
ND (T)<0.12< −3.8 1, 2
CRBR31
M6.7
A
···
7.1
0
2
0.9
ND (H)<0.56< −3.6 2
GY64
M8
A
···
3.7
2
0
0.2
ND (S)<0.93< −3.5 1, 2
GY141
M8
BF
BF-S2
9.0
10
0
1.5
3.3×10−6 (S)
0.39−3.4 2
A
···
22.3
9
16
8.6
ND (S)<0.33< −3.5
GY202
M7
BF
···
9.4
2
5
1.6
ND (S)<1.2< −3.7 1
A
···
16.8
3
5
4.9
ND (S)<2.0< −3.4
GY310
M8.5
BF
BF-86
9.2
190
56
1.5
<1.0×10−15 (T)10−3.4 1— Candidate Brown Dwarf (BDc ) —CRBR15
M5
A
···
4.2
2
2
0.3
ND (S)<1.1< −4.2 1
GY5
M7
A
A-S2?
4.7
7
4
0.4
1.7×10−7 (S)0.42−4.8 1
GY37
M6
A
A-30
4.2
7
5
0.3
3.7×10−8 (S)0.59−4.2 1
GY59
M6
A
A-38
1.9
15
17
0.1
<1.0×10−15 (T)
2.6−3.9 1
GY84††
M6
A
···
1.9
1
3
0.1
1.3×10−4 (T)
0.85−4.7 1
GY218
M4.0
BF
···
7.2
3
5
0.9
ND (S)<0.45< −3.6 2
CRBR67
M4.4
BF
···
3.2
0
0
0.2
ND (–)<0.31< −3.5 2* Spectral type determined by the water vapor absorption.
† Observation identification (see Table 5.1).
‡ Identified X-ray sources picked up with the wavdetect command (see Table 5.2).
§ Radius of the point-spread function.
k X-ray counts in the soft (0.5–2.0 keV) and hard (2.0–9.0 keV) bands, and predicted background counts in the soft bandwithin the source circle. Background counts in the hard band are ∼3 times larger than those in the soft band.
# CL is the confidence level of the source detection (Eq.6.2). Parentheses indicate the energy band in which the maximum
CL value is obtained (S: 0.5–2.0 keV, H: 2.0–9.0 keV, T: 0.5–9.0 keV). ND represents non X-ray detection with CL < 0.999(i.e., 1−CL > 10−3 ).
** Absorption-corrected luminosity (1028 ergs s−1 , “< L
X >” in the text) in 0.5–9.0 keV.
†† X-ray emission may be part of diffuse emission A-48 (see §6.6.2).
‡‡ References: (1): Wilking et al. (1999); (2): Cushing et al. (2000)
Name

Sp. type*

ID†

X-ray‡

PSF§
(00 )

80

6.4.2

CHAPTER 6. INDIVIDUAL SOURCES

X-ray spectra and luminosities

For the brightest two sources —, one BD (GY310) and one BDc (GY59) —, we could analyze the
X-ray spectra (§5.6). Figure 6.16 shows their spectra. The best-fit plasma temperature (Table 5.2)
is in the range of 1–4 keV.

Fig. 6.16.— Time-averaged spectra of (a) GY310 and (b) GY59. The upper panels show data
points (crosses) and the best-fit MEKAL model (solid line), while the lower panels are the data
residuals from the best-fit model.
For the faint or non-X-ray sources, we assume the spectrum to be a thin-thermal plasma of
2 keV temperature (the mean value of the bright sources), while the absorption column is estimatedfrom the empirical relation, NH = 1.38×1021 AV cm−2 (§7.4.1), where AV is the visual extinction
derived with IR photometry (Wilking et al., 1999; Cushing et al., 2000). We then calculate the
X-ray luminosities or those of the 99.9 % upper limits in the soft band, which is summarized inTable 6.5.

6.4.3

Time variability

Figure 6.17 shows the light curves of (a) GY310 and (b) GY59 in 0.5–9.0 keV, where the backgroundlevels are represented by dashed lines. The light curve of GY59 is flare-like, although our criterion(Eq.5.1) rejects this as a flare. This light curve is typical of those of low-mass main-sequence andpre-main-sequence stars; fast rise and slow decay. The flare-decay timescale (∼104 s) is also thesame. GY310, on the other hand, shows no clear flare, but exhibits aperiodic variability of fluxchange by a factor of ∼2 within 100 ks exposure. Since we see no large intensity variation in the
background flux, we conclude that its variability is highly significant.

6.4. BROWN DWARFS

81

Fig. 6.17.— Light curves of (a) GY310 and (c) GY59 in 0.5–9.0 keV. Dashed lines show thebackground levels.

6.4.4

X-ray detection rate

We detect X-ray emission from two out of eight BDs and three out of seven BDc s. The X-ray
detection rate is thus 25 % (BD) and 43 % (BDc ). Neuhäuser et al. (1999) claimed the X-ray
detection of only one BD (GY202) with ROSAT/PSPC in the same field. Our present observations
find no X-rays from GY202 with the upper limit of 1.2×1028 ergs s−1 . Instead, we find a class
II WL1 (A-78 = BF-16) within the PSPC error circle of GY202. This source emits X-rays withcomparable luminosity to that of GY202 reported with ROSAT/PSPC. Conversely, Neuhäuser etal. (1999) found no X-rays with an upper limit of ∼1028 ergs s−1 from all Chandra detected BD
and BDc s. They, however, estimated the X-ray luminosity assuming < kT >= 1 keV and Av = a
few mag (i.e., NH = a few ×1021 cm−2 ), which may be significantly lower than the real case. We
hence re-estimate the upper limit with the same assumption adopted in this paper (§6.4.2), thenfind that the ROSAT upper limit could be nearly one or two order magnitude higher than those of
the original paper. If these newly estimated upper limits are adopted, we find no significant fluxchange from the ROSAT to the Chandra era.
In Table 6.5, we display the ratio of the X-ray (< LX >) and bolometric (Lbol ) luminosities.
The < LX >/Lbol value lies between 10−3 –10−5 , which is similar to low-mass YSOs (Figure 7.7)
and dMe stars (Giampapa et al., 1996), but is significantly higher than that of the quiescentsolar corona; ∼5×10−7 (Vaiana & Rosner, 1978). We should also note that most of the upper
limits of LX /Lbol of X-ray non-detected BDs and BDc s are scattered around 10−3 –10−5 , which
are comparable with or slightly lower than those of X-ray detected BDs and BDc s. We therefore
suspect that these X-ray non-detected BDs and BDc s may emit X-rays below (but possibly near)
the current sensitivity limit of Chandra.

82

CHAPTER 6. INDIVIDUAL SOURCES

6.5

New YSO Candidates – A-29, A-31, and BF-36

6.5.1

Source selection and X-ray images

Imanishi et al. (2001a) suggested that most of unidentified sources are likely background active
galactic nuclei (AGNs; see also §6.6). However, some fractions would be IR-faint cloud members,possibly brown dwarfs and/or very young stellar objects like class 0–I protostars. To search for
new YSO candidates, we select the unidentified sources located near 850 µm clumps (Johnstone etal., 2000), then find three candidates (A-29, A-31, and BF-36: Imanishi et al., 2002b) in the dense
millimeter/sub-millimeter cores of SM1/SMM J16264−2424 and SMM J16264−2423 in ρ Oph-A,
and B1-MM4/SMM J16272−2430 in ρ Oph-B1 (Motte et al., 1998; Johnstone et al., 2000).
SMMJ16264-2422b

unID

GY31

Dec.(J2000)-24:23:30

A-31

24:00

SM1
KSHK-A

SMMJ16264-2423

A-29

24:30
VLA1623GY21

SMMJ16264-2424

GY29

S2

26:28
26:26
R.A.(J2000)

26:24

25:00
16:26:32

26:30

Fig. 6.18.— ACIS images around the mm cores (left) SM1 and (right) B1-MM4 in the 2.0–9.0 keVband. The contour is the (left) 3 mm and (right) 1.3 mm radio intensities (Motte et al., 1998;
Kamazaki et al., 2001). Crosses, Large solid and dotted circles represent the center of each mmcore, size of 850µm clumps (Johnstone et al., 2000) and field of view of the 3 mm observation,
respectively. A square shows the position of the class 0 source VLA 1623. Other X-ray sources are
known YSOs (GY21, S2, GY29, GY31, and GY238) and unidentified sources (designated as unID).

Figure 6.18 (left) shows the ACIS image of obs-A around SM1 in the hard X-ray band (2.0–
9.0 keV), which is overlaid on a 3 mm contour map (Kamazaki et al., 2001). Large circles represent
the 850 µm clumps (Johnstone et al., 2000). We see two point-like sources (A-29 and A-31) near
the center of the 850 µm clump SMM J16264−2424 and at only ∼200 away from the center of the
3 mm core A (hereafter, KSHK-A). Since there is no significant soft X-ray (0.5–2.0 keV) from A-29
and A-31, they are deeply embedded in the cores. Another point source located near the core,a class II source GY29 (A-27), on the other hand, is clearly visible in the soft band because of6.5. NEW YSO CANDIDATES – A-29, A-31, AND BF-36

83

relatively small absorption (NH = 2.5×1022 cm−2 ; Table 5.2), hence may be located in front of thecore.
The hard band (2.0–9.0 keV) ACIS image near ρ Oph B1 is shown in Figure 6.18 (right). Thereis a hard X-ray source (BF-36) about 3000 south of the 1.3 mm peak B1-MM4. A class II source
GY238 (BF-33) is also appeared near the center of SMM J16272−2430, which would be located in
the clumps, as supported by the large absorption of NH = 4×1022 cm−2 (Table 5.2).

6.5.2

Time variability

Figure 6.19 shows the X-ray light curves (background is not subtracted) of A-29, A-31, and BF-36
in the 2.0–9.0 keV band. We have detected solar-like flares from A-29 and BF-36 (§5.5). The light
curve of A-31 also seems to be time variable; A-31 has only 12 photons and about 80 % of them
(10/12) are concentrated in a limited time interval of ∼30 ks. We estimate the chance probabilityof such a photon clustering from a constant-flux source to be less than 1 % with the KolmogorovSmirnov test (Press et al., 1992). Thus the time variability of A-31 is also significant. The X-ray
active time intervals are shown with dotted vertical lines in Figure 6.19 and are referred as the flarephase hereafter. The hard band counts in the quiescent phase (time periods outside of the flarephase) for A-29, A-31, and BF-36 are 5, 2 and 81, while those of the expected background counts
are 0.7, 0.5, and 1.8, respectively. This means that A-31 has no significant quiescent flux with ≥99.9 % of the confidence level (Poisson statistics; see §6.4.1).
A-29

A-31
Flare

Flare

BF-36
Flare

Fig. 6.19.— Light curves of (left) A-29, (middle) A-31, and (right) BF-36 in the 2.0–9.0 keV band.
The time bins of each data point are 2000 (A-29), 4000 (A-31), and 4000 (BF-36) s. The arrows
indicate the times of the flare phase (see text).

84

6.5.3

CHAPTER 6. INDIVIDUAL SOURCES

X-ray spectra and luminosities

Figure 6.20 shows the X-ray spectra of A-29, A-31, and BF-36 during the flare phases and thatof BF-36 in quiescent. We fit them by the MEKAL model with the fixed abundance of 0.3 solar.
However, reasonable constraint can not be obtained on < kT > and NH with the limited statistics
for A-29 and A-31. We hence fix the temperature at three typical values (1, 5, and 10 keV) of theother X-ray sources. The 1 keV temperature is typical of class II–III sources in quiescent, whilesome flares from class I–II sources have the 5 keV temperature and extremely high temperatureof 10 keV is found in some giant flares; e.g., YLW16A = BF-64 (§6.1.5). Although all fittings
are acceptable, the assumed temperature of 1 keV gives uncomfortably high luminosity of (4–7)×1031 ergs s−1 , hence we discard the case of the 1 keV temperature. Table 6.6 shows the best-fit
parameters, while the best-fit models (5 keV models for A-29 and A-31) are shown in Figure 6.20.
For the quiescent phase of A-29 and A-31, we assume the same parameters as those of the flare.
Since no significant X-ray is found in the quiescent phase for A-31, we estimate the 99.9 % upper
limit of the X-ray luminosity (Table 6.6). The quiescent luminosities are 2–30 times lower thanthose in the flare phases.

Table 6.6: Spectral parameters of A-29, A-31, and BF-36
No.
A-29† . . . . . . . . . . .

Cts*< kT >(keV)< EM >‡(1052 cm−3 )NH ‡(10 cm−2 )23< LX >(flare)‡§(1029 ergs s−1 )< LX >(quiescent)§(1029 ergs s−1 )21(5)5
7.5(3.1–16)
2.5(1.3–4.7)
11
0.37
···
10
5.0(2.1–10)
2.2(1.1–4.4)
8.1
0.27
A-31† . . . . . . . . . . . 10(2)5
5.0(0.53–110)
5.4(1.5–20)
6.9
<1.3
···
10
2.9(0.35–82)
5.0(1.2–20)
4.9
<0.90
k
BF-36(flare) . . . .
24
7.6(>1.7)
3.1(1.6–20)
0.92(0.53–1.6)
5.1···
BF-36(quiescent)80
2.7(1.1–43)
2.7(0.75–21)
···
···
2.9
*
X-ray counts in 2.0–9.0 keV during the flare phase. For A-29 and A-31, those during the quiescent phase are
given in the parentheses.
†
Spectral parameters are obtained with fixed temperatures of 5 and 10 keV (see text).
‡
These parameters for A-29 and A-31 are derived from the flare spectra. Parentheses indicate the 90 % confidence
limits.
§Absorption-corrected X-ray luminosity in 0.5–9.0 keV. < LX > in the quiescent phase for A-29 and A-31 isderived with the assumption of the same absorption as that in the flare phase.
kDerived by the simultaneous fitting of the flare and quiescent spectra.

6.5. NEW YSO CANDIDATES – A-29, A-31, AND BF-36

A-29

A-31

85

BF-36

Fig. 6.20.— X-ray spectra of (left) A-29, (middle) A-31, and (right) BF-36. Black and gray lines
indicate spectra of the flare and quiescent phases, respectively. The upper panels show the data
points (crosses) and the best-fit MEKAL models (solid line). For A-29 and A-31, we show the
best-fit models with the assumed 5 keV temperature. The lower panels are the data residuals fromthe model.

6.5.4

Possible origin of A-29, A-31, and BF-36

A-29
A-29 is found near the center of the 850 µm core SMM J16264−2424 with the separation angleof ∼1100 . This is larger than the position accuracy of Chandra but is comparable to that of thebeam size of the 850µm observation (∼1400 , SCUBA). Assuming a uniform density core, the column
density of SMM J16264−2424 is estimated to be ∼1023 cm−2 from the radio flux (Johnstone et al.,2000), which is consistent with the X-ray-determined NH of A-29 (∼ 2×1023 cm−2 ) within large
uncertainty for both the values. It is hence highly possible that A-29 is located in or behind SMM
J16264−2424. Using the LogN –logS relation by Mushotzky et al. (2000), we estimate the chanceprobability to detect a background AGN within the ∼1100 radius to be ∼5 %. Although this value
is relatively large, A-29 shows a clear X-ray flare typical of low-mass stars but not seen in usual
AGNs (Figure 6.19). No NIR counterpart of K < 15 mag, the association with the sub-millimetercore SMM J16264−2424, and the large absorption lead us to conclude that the origin of A-29 is anembedded star younger than class I.

A-31
A-31 is located within ∼200 from the center of the 3 mm clump KSHK-A (Kamazaki et al., 2001),
the central source of millimeter/sub-millimeter clump SM1/SMM J16264−2423 (Johnstone et al.,2000). Taking into account of the physical size of KSHK-A (∼400 ) and position uncertainty of
Chandra and the Nobeyama Millimeter Array interferometer observation (∼500 ), we can conclude

86

CHAPTER 6. INDIVIDUAL SOURCES

that A-31 is surely associated with KSHK-A. The chance probability to detect a background AGN
in the KSHK-A region is only ∼0.2 %. Furthermore, A-31 exhibits variable X-rays typical of
low-mass YSOs. The column density of KSHK-A is ∼1024 cm−2 (Kamazaki et al., 2001), which
is about twice of the X-ray determined NH of A-31 (∼ 5×1023 cm−2 ). Thus A-31 would be a
YSO embedded near the center of the compact core KSHK-A. KSHK-A is spatially coincident
with a 6 cm continuum source LFAM6 (Leous et al., 1991), which may be produced by a shockof outflow from a putative protostar (e.g., Rodriguez et al.1989, and references theirin). No NIR
counterpart with K < 15 mag, its close location at the compact core center KSHK-A, the large
absorption (∼5×1023 cm−2 ), and the presence of the 6 cm enhancement indicate that A-31 is astrong candidate of X-ray emitting class 0.
A northwest part of the CO molecular outflow from the class 0 VLA1623 (designated as E–H
in Yu, & Chernin 1997) is not aligned to the central part (B–D) but is slightly shifted to the northdirection pointing to A-29 and A-31 (Figure 6.21). We speculate that the origin of the outflow
E–H is A-29 or A-31 (not from VLA1623), although there is no CO emission at their closer vicinity(Kamazaki et al., 2001).

Fig. 6.21.— 12 CO (J=1–0) map around
VLA1623 (Yu, & Chernin, 1997). Filledsquares represent the positions of A-29 and
A-31.

BF-36
BF-36 is located in the 850 µm clump SMM J16272−2430 with about 3000 separation from its center.
It lies in an elongated 1.3 mm core chain with local peaks of B1-MM2, 3, and 4 (Motte et al., 1998).
The source is about 3000 south of the peak of B1-MM4 (Figure 6.18). The chance probability ofbackground AGN to be associated within 3000 radius from the core center is ∼20 %. However,
BF-36 is among the brightest unidentified sources and is very unique which exhibits a clear flare.
We therefore suggest that BF-36 is also a cloud member embedded near the dense core B1-MM4.
Since millimeter/sub-millimeter peak fluxes of B1-MM4/SMM J16272−2430 are about ten times
fainter than those harbor A-29 and A-31, although the clump size is comparable (6200 AU), the
2–6 times smaller NH of BF-36 (9×1022 cm−2 ) than those of A-29 and A-31 is consistent with the6.5. NEW YSO CANDIDATES – A-29, A-31, AND BF-36

87

embedded star scenario.
Considering that NH of BF-36 is rather typical of those of class I sources (see Figure 7.3), it
is located at rather off-center position, and a class II source GY238 (BF-33) is closer to the centerof SMM J16272−2430, BF-36 would be a more evolved star than the class 0 stage, possibly class
I or II. Assuming the average ratio of < LX >/Lbol to be 10−4.0 in the quiescent phase for YSOs(§7.4.2) and using NH of 9×1022 cm−2 , Lbol and AV of BF-36 are expected to be ∼0.8 L¯ and∼60 mag (§7.4.1). These are above the current detection limit; some NIR-detected sources havenearly the same values (see Figure 7.1). One possibility of no NIR counterpart is that BF-36 is a
YSO with very active X-rays. Assuming < LX >/Lbol of ≈10−3.5 , we then expect Lbol of 0.08 L¯ ,which is below or near the current detection limit (Figure 7.1).

6.5.5

Ionization effect to circumstellar materials

The flare luminosity and duty ratio of A-29, A-31, BF-36 are ∼1030 ergs s−1 and roughly one per
100 ks, respectively. Due to the large absorption, a significant fraction of X-ray energy is supplied
into circumstellar materials. Using the parameters in Table 6.6, we estimate the X-ray ionization
rate (ζX : Lorenzani & Palla 2001); the fraction of protons to be photo-ionized per unit time, which
is defined asR−n σ̃r( ν )−ndν −1< LX > σ̃ Jν e H νXRζX = 1.7
[s ],
(6.3)
4πr2 ∆²Jν dν
where σ(E) = σ̃(E/1keV)−n is the total photoelectric cross section with n = 2.51 and σ̃ =2.16×10−22 cm2 , ∆E (≡ 35 eV) is the mean energy to make an ion-electron pair with cosmic
abundance, nH is the proton mass, Jν is monochromatic X-ray flux (ergs s−1 Hz−1 ), and r is the
distance from the X-ray source. We assume nH of A-31 and BF-36 to be 1.5×108 and 1.3×106 cm−3 ,following Motte et al. (1998) and Kamazaki et al. (2001). Since NH of A-29 is about a half of A-31,
nH is assumed to be 7.5×107 cm−3 for A-29. Then we obtain the “ionization radius”, in which ζXexceeds that of cosmic rays, to be ∼102.5 , 102.4 , and 103.1 AU, respectively (Figure 6.22). These
sizes are significantly smaller than the “parent” core size of typically 103.8 AU, hence have no large
impact in global structure. In a smaller scale of close vicinity of the YSOs (KSHK-A, for example),however, the X-ray ionization may still have significant effect on the circumstellar materials. Since
the mass accretion is thought to be coupled to the magnetic field (Feigelson & Montmerle, 1999),
our result may put significant influence on the growth of YSOs.

88

CHAPTER 6. INDIVIDUAL SOURCES

A-29

A-31

BF-36

1keV
1keV

Flare

10keV

Quiescent
5keV

5keV
10keV

SMMJ16264-2424

KSHK-A

SMMJ16267-2430

Fig. 6.22.— The relation between the X-ray ionization rate and the distance from the source for(left) A-29, (middle) A-31, and (right) BF-36. For A-29 and A-31, we represent three lines for theflare phase with the assumed temperatures (1, 5, and 10 keV). For BF-36, the flare and quiescentphases are separately estimated. Dashed and dash-dotted lines represent the ionization rate fromcosmic rays and the core radii.

6.6. UNIDENTIFIED SOURCES

6.6

Unidentified Sources

6.6.1

Overall feature

89

In order to examine the average feature of unidentified sources, we co-add and make composite
spectra, separately for total-band, soft-band, and hard-band sources (hereafter, we call them “totalband unidentified sources” and so on). We do not include A-29, A-31, BF-36, A-75 = BF-11, BF-92,
BF-S7, and BF-H3 because they have different nature from the other unidentified sources. A-H2 isalso excluded because of severe contamination from A-2 (DoAr21). Figure 6.23 shows the composite
spectra. First we fit them by the MEKAL model. However, uncomfortably high temperature (>20 keV) is obtained for the total- and hard-band unidentified sources. We hence re-fit them by apower-law model and find acceptable fits. Figure 6.23 (solid lines) and Table 6.7 show the best-fit
models and parameters, respectively.

Fig. 6.23.— Composite spectra of (a) total-band, (b) soft-band, and (c) hard-band unidentified
sources. The upper panels show data points (crosses) and the best-fit non-thermal (a and c) andthermal (b) model (solid line), while the lower panels are the data residuals from the best-fit model.
Circles and triangles represent the data for obs-A and obs-BF, respectively.

The best-fit photon index (∼1.5) of the total-band unidentified sources is similar to that of
the canonical value of active galactic nuclei (AGNs: ∼1.7), while the best-fit NH of 5.3×1022 cm−2
is nearly equal to the average value through this cloud (Tachihara et al., 2000). Inversely, if the
background sources have a mean photon index of 1.5 and are absorbed by the cloud gas of NH =5.3×1022 cm−2 , the log N – log S relation by Mushotzky et al. (2000) predicts the detectable source
numbers to be several tens, which is consistent with the total number of unidentified sources (∼80).
Also, the spatial distribution of the total-band unidentified sources seems to be anti-correlated withthe molecular cloud (Figure 6.24). Most of the total-band unidentified sources therefore are likely
to be background AGNs, as well as the radio bright unidentified source A-75/BF-11 = LFAM21(§6.6.3).

90

CHAPTER 6. INDIVIDUAL SOURCES

The average temperature of the soft-band unidentified sources is ∼0.2 keV, which is significantly lower than those of YSOs, although the data excess around 1–2 keV (Figure 6.23a) indicatesthe existence of higher temperature component of < kT > = 1–2 keV. Considering the low temperature plasma and small absorption column (∼ 2×1021 cm−2 ), one plausible idea is coronal emissionof MS stars in front of the ρ Oph cloud. However, it would be less possible because there is no
optical counterpart (Chini, 1981). Another idea is that an outflow from YSOs in ρ Oph collides at
a dense interstellar matter and makes a low temperature (0.1–0.5 keV) plasma, like X-ray emission
from Herbig-Haro objects (HHs: Pravdo et al., 2001; Favata et al., 2002). In particular, A-S1 islocated in the direction of CO outflow from the class 0 source VLA1623 (Figure 6.25), hence is astrong candidate of jet-induced X-rays. The average flux (absorption-corrected) per one source of
1.4×10−15 ergs s−1 cm−2 is converted to < LX > of 3.5×1027 ergs s−1 assuming the distance of
145 pc, which is about 100 times lower than those of HHs (∼ 1029 ergs s−1 ). This indicates thatthe emitting region of soft-band unidentified sources has smaller volume and/or density than thatof HHs.
The hard-band unidentified sources show nearly the same spectral parameters as the totalband sources, although the average flux is about four times lower, hence are also thought to be
background AGNs.

Fig. 6.24.— Spatial distribution of unidentified sources overlaid on the 1.3 mm continuum map(Motte et al., 1998). Crosses, squares, and diamonds represent total-band, soft-band, and hardband unidentified sources, respectively.

6.6. UNIDENTIFIED SOURCES

91

Fig. 6.25.— ACIS image around the ρ Ophcore A region in 0.5–2.0 keV. Solid, dotted,
and dashed contours show 1.3 mm radio continuum, blueshifted and redshifted CO (2–1)
intensities, respectively (André et al., 1993).
Arrows indicate the positions of the class 0
source VLA1623 and mm cores (SM1, SM1N,
and SM2). The position of A-S1 is shown by
the circle.

Table 6.7: Best-fit parameters of the composite spectra of unidentified sources
Class*

N†< kT >/Γ‡§(keV/–)

NH §(1022 cm−2 )Fluxk
Fluxc #
(10−15 ergs s−1 cm−2 )

Total
Soft
Hard**60
107

1.5(1.1–1.7)
0.16(0.042–0.26)<2.2

5.3(4.2–6.6)
0.20(<0.59)
4.0(2.0–10)

441.7(7.4)
3.5(0.35)
14.8(2.1)

905.3(15.1)
14.4(1.4)
27.8(4.0)*

See text.
The total number of sources added in the spectra.
‡
Best-fit temperature for soft-band unidentified sources and photon index for the others.
§
Parentheses indicate the 90 % confidence limit.
kObserved flux in 0.5–9.0 keV. Parentheses indicate the average value per one source.
#
Absorption-corrected flux in 0.5–9.0 keV. Parentheses indicate the average value per one source.
**
The best-fit value of Γ is not determined, hence we assume Γ = 1.5 (best-fit value for the totalband unidentified source) for the estimation of the other parameters.
†

6.6.2

A-48

The shape of this unidentified source seems to be diffuse with a radius of ∼500 ∼= 103 AU (Figure6.26), which is significantly larger than the default PSF size (∼1.00 3), although we can not reject thepossibility that A-48 is a possible false source. We also note that a BDc GY84 is located at the edgeof the diffuse emission. We previously reported significant X-ray emission from GY84 (Imanishi etal., 2001b), which may be part of the diffuse emission, although any relation between GY84 andA-48 would be conceivable.

92

CHAPTER 6. INDIVIDUAL SOURCES

Fig. 6.26.— ACIS image around A-48 in 0.5–9.0 keV
smoothed with a Gaussian of σ = 6 pixels. Large andsmall circles represent the position of A-48 and a BDc
GY84, respectively.

6.6.3

A-75/BF-11 – LFAM21

A-75 = BF-11 has no NIR counterpart but is spatially identified with the radio-bright source
LFAM21, a candidate of a compact extragalactic AGN (Stine et al., 1988). Both X-ray spectra of
LFAM21 seem to be non-thermal typical of AGNs and show line-like structure at ∼5 keV (Figure6.27). We hence fit them by a (power-law + gaussian) model assuming the same absorption. The
best-fit models and parameters are shown in Figure 6.27 and Table 6.8. The best-fit photon index
of 1.2–1.4 and NH of ∼6×1022 cm−2 are reasonable to consider that the typical AGN spectrum isabsorbed at the ρ Oph cloud (Tachihara et al., 2000). The line center energy is 4.8–5.0 keV. If thisis a redshifted neutral iron (6.4 keV) frequently seen in the AGN spectra (Awaki et al., 1991), a
redshift (z) of LFAM21 is estimated to be ∼ 0.3.

Fig. 6.27.— X-ray spectra of LFAM21 in obs-A (A-75, left) and obs-BF (BF-11, right). The upper
panel shows data points (crosses) and the best-fit (power-law + gaussian) model (solid and dashedlines), while the lower panel shows residuals from the best-fit model.

6.6. UNIDENTIFIED SOURCES

93

Table 6.8: Best-fit parameters of LFAM21
No.

Γ*

NH *
(1022 cm−2 )Egauss *†
(keV)

EW *‡
(eV)

A-75
BF-11

1.2(1.1–1.5)
1.4(1.1–1.8)

5.9(5.2–7.0)
5.9k

5.0(4.9–5.7)
4.8(4.7–4.8)

380(130–640)
290(120–510)

Flux§
Fluxc §
(10−14 ergs s−1 cm−2 )
6.4
7.3

11.5
14.1

*

Parentheses indicate the 90 % confidence limit.
The gaussian center energy.
‡
Equivalent width for the gaussian line.
§
Observed and absorption-corrected fluxes in 0.5–9.0 keV.
kDerived by the simultaneous fitting of the spectra of A-75 and BF-11.
†

6.6.4

BF-92

An unidentified source BF-92 shows a weak flare (Figure A.4) similar to the new YSO candidates (A-29 and BF-36; §6.5). Although
there is neither mm core nor cm emission at the position of BF-92,
Grosso et al. (2001) detected jet-like H2 emission 20–4000 away from
BF-92 as shown in Figure 6.28. They concluded that the origin of
the jet is the class I source YLW15A (BF-61) located ∼90 away fromthe H2 jet, however we infer that the origin may be BF-92, a new
YSO (probably class 0–I protostar) having violent jet activity. The
flare light curve and highly absorbed spectrum (NH ∼6×1022 cm−2 )also indicate the YSO nature of BF-92. Followup deep NIR and radio observations would be fruitful.

6.6.5

BF-S7

A soft-band source BF-S7 is located near a class I CRBR85, although the offset (3.00 7) is significantly large considering its off-axis
6.28.— ACIS image
angle (6.0 9). Also, the detection only in the soft band is inconsis- Fig.
tent with the X-ray emission from class I, because class I sources around BF-92 overlaid on thegenerally suffer large absorption (§7.3) and consequently are hard H2 contour map (Grosso etto see at the soft band. One possible scenario is that an outflow al., 2001).
from the central source collides at dense circumstellar envelope andmakes low temperature (0.1–0.5 keV) plasma, like X-ray emission from Herbig-Haro objects (HHs:

94

CHAPTER 6. INDIVIDUAL SOURCES

Pravdo et al. 2001; Favata et al. 2002; see also Figure 2.9). Although CRBR85 has shown no evidence for outflow activity, bipolar nebulosity was detected, indicating the existence of an edge-on
circumstellar disk (Brandner et al., 2000). We therefore suspect that jet-induced X-rays with alarge offset from the central source are possible because of the large inclination angle.

BF-S7
CRBR85

Fig. 6.29.— Close-up ACIS image around BF-S7
and CRBR85. Circles represent the position of
BF-S7 and CRBR85. Filled squares correspondto one X-ray photon.

5arcsec

6.6.6

BF-H3

A hard-band source BF-H3 has no NIR counterpart, but is associated with an optical [SII ] source
A2, a candidate of HH (Gómez et al., 1998). Hence BF-H3 may be the third sample of X-ray
emitting HHs following HH2 (Pravdo et al., 2001) and L1551 IRS5 (Favata et al., 2002), although
there is no other evidence for its HH nature and for the origin of A2. Interestingly, BF-H3 shows
X-ray emission only in the hard band (2.0–9.0 keV), hence would have much higher temperatureplasma (&2 keV) than that of L1551 IRS5 and HH2 (0.1–0.5 keV). This indicates high-velocity
outflow activity, which is the similar case to GY91 (§6.1.1).

Chapter 7

Overall Feature of X-ray Emission
from YSOs
7.1

X-ray Detection Rate

We detected 195 X-ray sources from the ∼1 pc × 2 pc region of the ρ Oph main body. About
60% (110) of them are identified with NIR sources. Table 7.1 summarizes the number of X-ray
sources and detection rate for each class. High detection rate (≥ 50%) is obtained for class I
sources as well as class II and III sources, while that is slightly lower for class IIIc , BD, and BDc
sources (< 40%). Figure 7.1 shows the relation between the bolometric luminosity (Lbol ) and visualextinction (AV ) for the ρ Oph sources (Table B.1). We see the tendency that most sources with no
X-ray detection (gray symbols in Figure 7.1) are distributed at the upper-left area in this figure,i.e., having relatively lower Lbol and higher AV . Considering the almost constant < LX >/Lbol
and NH /AV correlations (§7.4.1 and §7.4.2), this means that the X-ray non-detected sources havelower < LX > and higher NH even if they emit X-rays. It hence would be conceivable that the
X-ray non-detected sources may emit X-rays below the current sensitivity limit of Chandra. Lower∼ lower < LX >) of
detection rate of class IIIc , BD, and BDc sources is simply due to lower Lbol (=these classes. On the other hand, X-ray non-detected class I sources have relatively high Lbol withmoderate AV (e.g., GSS30, G304, and GY378; Figure 7.1). However, this could be partly due toshort time variability. For example, WL21 (BF-8) and GY256 (BF-51) are hardly visible except
when they exhibit X-ray flares (see Figure A.1). Also long-term variability, like the 11-year cycleof the solar X-ray activity, may exist. We therefore expect that virtually all YSOs, at least in ρ
Oph, are potential X-ray emitters regardless of their classes.
95

96

CHAPTER 7. OVERALL FEATURE OF X-RAY EMISSION FROM YSOS

Table 7.1: The number of X-ray sources and detection rate for each classClass
I ...............
II . . . . . . . . . . . . . .
III . . . . . . . . . . . . . .
IIIc . . . . . . . . . . . . .
Unclassified NIR
BD&BDc . . . . . . .
F ...............
unidentified . . . .

Source*
15
70
17
27
···
15
3
···

X-ray†
8
58
17
9
12
5
1
85

Ratio‡
53 %
83 %
100 %
33 %
···
33 %
33 %
···

*

The number of sources within the ACIS field of view.
The number of X-ray detected sources.
‡
The X-ray detection rate.
†

Fig. 7.1.— Relation between the bolometric luminosity (Lbol ) and visual extinction (AV : Bontemps
et al., 2001, see also TableB.1). Circles, triangles, squares, and stars represent class I, class II, class
III+IIIc , and BD/BDc sources, respectively. Gray symbols are sources with no X-ray detection.
Three X-ray non-detected class I sources with relatively high Lbol are also indicated (GSS30, GY304,
and GY378).

7.2. LUMINOSITY FUNCTION

7.2

97

Luminosity Function

Fig. 7.2.— Normalized X-ray luminosity functions of class I (dashed), class II (dash-dotted), andclass III (solid) sources in the (a) quiescent and (b) flare phases. The mean value (Mean) andstandard deviation (σ) of log[< LX >] in the unit of ergs s−1 for each class are shown in the figures.
Parentheses indicate errors of the mean values.
We calculate the X-ray luminosity function for the detected X-ray sources of class I, II, and
III+IIIc (Figure 7.2), using the ASURV statistical software package (rev.1.2)1 based on the maximum likelihood Kaplan-Meier estimator in order to examine the differences along the evolutional
stages. The luminosity function of class I seems to be shifted toward the higher values than thoseof class II and III+IIIc both in the quiescent and flare phases.
To be more quantitative, we estimate the significance level using two nonparametric two-sample
tests in ASURV; the Gehan’s generalized Wilcoxon test (GW) and the logrank test (Feigelson &Nelson, 1985). The results are given in Table 7.2, where blank (...) indicates that the significancelevel is less than 90 %. From Table 7.2, we see that both the GW and logrank tests for the < LX >difference between class I and the others in the flares show the marginal significance level of ∼94 %.
We further should note that the effect of “undetected” faint flares of class I can not be ignored; if
we assume an undetected class I has < LX > smaller than 1030 ergs s−1 which is near the detection
threshold of a class I flare, and add this source to the GW and logrank test sample, then therelevant significance level is largely reduced to < 90 %. We hence conclude that the higher < LX >of class I than class II+III during flares is marginal.
On the other hand, from Table 7.2, we see no significant difference between class II and
III+IIIc source in the flare phase. This is contrary to the results of the Taurus-Auriga-Perseus
1

see http://www.astro.psu.edu/statcodes.

98

CHAPTER 7. OVERALL FEATURE OF X-RAY EMISSION FROM YSOS

samples (Stelzer et al., 2000), for which there was a significant difference between the distributionof flare < LX > of class II (CTTS) and class III+IIIc (WTTS) sources.
Also, we see no significant difference of < LX > in the quiescent phase among all the classes.
This is consistent with the previous estimation for this region with ROSAT derived by Grosso etal. (2000). They tried to estimate the significances rather strictly by considering the upper limit
of undetected sources. Our samples give more severe constraint because the detection thresholdis largely reduced to give the mean luminosities of 1029.5 –1029.8 ergs s−1 , which is about 10 times
lower than that of ROSAT.

7.3

Distribution of X-ray Parameters

In Figures 7.3–7.5, histograms of NH , < kT >, τr , and τd are shown for each class separately,
together with the mean values and standard deviations. We exclude the samples whose best-fit
parameters and/or those errors are not determined due to the limited statistics. For sources withmulti-temperature plasma (BF-46 and BF-96), we regard < kT > of the soft and hard components as the quiescent and flare values, respectively, following the discussion that the former would
be steady coronal emission while the latter is flare activity (§6.2). Like the case of LX , we estimate the significance level for the differences of these parameters among all the classes with two
nonparametric two-sample tests in ASURV. The results are also given in Table 7.2.
Absorption (NH )NH is distributed from 1021.5 to 1023.5 cm−2 (Figure 7.3). The mean value of NH ∼ 1022.5 cm−2
is roughly consistent with the average value through this cloud derived by the C18 O radio observations (Tachihara et al., 2000). Also, the distribution of NH of class I sources is larger than thatof the others. Using ASURV, we estimate that the significance level of this tendency is rather highof &97 %. This supports the idea that X-rays from class I sources come near the central source
deeply embedded in denser envelopes and suffer larger intrinsic absorption.
Plasma temperature (< kT >)Most of the < kT > values in the quiescent phase are in the range of 0.2–5 keV, while it
becomes systematically higher (1–10 keV) in the flare phase, indicating the plasma heating during
the flare. Although the flare temperature of all classes is distributed around 3–4 keV, some flares ofclass I and II sources show higher temperature than 5 keV, while all flares of class III+IIIc sources
have < kT > less than 5 keV. The mean values of < kT > during flares for class I and II (4–5 keV)is therefore larger than that of class III+IIIc (2.7 keV).
From Table 7.2, we see that both the GW and logrank tests for the difference of < kT >

7.3. DISTRIBUTION OF X-RAY PARAMETERS

99

during flares between class I+II and class III+IIIc sources show high significance level of 94 %and 98 %. One may concern that systematically large absorption of class I and II sources makes
artificial spectral hardening of these classes. We however argue that the observed tendency ofhigher < kT > during flares for younger sources is not changed by the extinction differences withthe following reasons;(1) If the intrinsic distribution of the flare < kT > in class III+IIIc is the same as classes I and II,some of the class III+IIIc flares should show higher < kT > than 5 keV, but there is no such flarein class III+IIIc .
(2) We re-estimate the mean < kT > of class I+II and class III+IIIc during the flare for the limited
sources of NH ≤ 5×1022 cm−2 (most of class III+IIIc sources have lower NH than this value).
Although the mean NH is nearly the same (2.5±0.2 and 2.2±0.3 × 1022 cm−2 for class I+II andclass III+IIIc ), the mean < kT > is still higher for class I+II (3.7±0.4 keV) than for class III+IIIc(2.4±0.3 keV).
In the quiescent phase, on the other hand, ASURV shows the significance level of >95% for
the tendency of higher < kT > of class I than the others (Table 7.2). However, we suspect thismay be a bias effect because the quiescent temperatures are systematically low (. 5 keV), hence
would be significantly affected by the absorption compared with the flare temperatures.
Flare timescales (τr and τd )The rise and decay timescales of the flares (τr and τd ) are distributed around 103.5 s and
104 s, respectively. Although the mean value of τr indicates that younger sources have shorter risetimescale, the difference of τr among these classes is not statistically significant (<90 %, Table 7.2).
The larger mean values for class III+IIIc sources are primarily due to the flares with unusuallylong timescales of &50 ks (A-2, A-63, and BF-96); such flares are not seen in class I and II sources(≈30 ks at the maximum). Since we can not reject the possibility that these long timescales are
simply because of the composition of unresolved two (or more) flares, we further estimate the mean
τr and τd values of class III+IIIc sources without these flares to be 3.1±0.6 and 6.9±0.8 ks, whichis comparable to those of the other classes.

100

CHAPTER 7. OVERALL FEATURE OF X-RAY EMISSION FROM YSOS

Fig. 7.3.— The histogram of log(NH ) for each class, with their mean values (Mean) and standard
deviation (σ). Parentheses indicate errors of the mean values.

Fig. 7.4.— Same as Figure 7.3, but for < kT > in the (a) quiescent and (b) flare phases for eachclass, with their mean values (Mean) and standard deviation (σ) in the unit of keV for each class.
Parentheses indicate errors of the mean values.

7.3. DISTRIBUTION OF X-RAY PARAMETERS

101

Fig. 7.5.— Same as Figure 7.3, but for flare rise and decay timescales (τr and τd ), with their meanvalues (Mean) and standard deviation (σ) in the unit of ks. Parentheses indicate errors of the meanvalues.

Table 7.2: Results of two-sample tests for the differences of the observed parameters among classes*

Parameter— Sample size* —I II
III+IIIc— Probability for the null hypothesis† —I vs II+III+IIIc I+II vs III+IIIc II vs III+IIIc
GW logrank GW logrank GW logranklogLX (Flare)
logLX (Quiescent)logNH< kT > (Flare)< kT > (Quiescent)
τr
τd

8
8
8
7
7
6
7

0.059
...
0.005
0.079
0.008
...
...

36
61
56
35
61
28
28

16
26
24
16
24
8
12

0.063
...
0.028
...
0.022
...
...

...
...
0.031
0.060
0.046
...
...

...
...
0.015
0.017
0.037
...
...

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

...
...
0.086
0.047
0.026
...
...

The number of data points used for the two-sample tests.
Probability that the hypothesis of two distributions being the same is true, which is derived with the
Gehan’s generalized Wilcoxon test (GW) and the logrank test (Feigelson & Nelson, 1985). Blank (...)
indicates that the probability is larger than 0.1.
†

102

CHAPTER 7. OVERALL FEATURE OF X-RAY EMISSION FROM YSOS

7.4

Correlation between X-ray and Other Physical Parameters

7.4.1

NH vs AVThe relation between the X-ray and visual extinction (NH and AV ) is a key to understand the gasto-dust ratio of the interstellar medium, since X-ray and optical–IR photons are mainly absorbed
and scattered by interstellar gas and dust, respectively. We find a good log-linear correlation asshown in Figure 7.6. The best-fit linear model (dashed line) is

(

AV
NH
) = (1.38 ± 0.03)(
).
21
−2
10 cm
mag

(7.1)

The NH /AV ratio of 1.38 is roughly consistent with that of other star-forming clouds; OMC 2/3(1.49: Tsujimoto et al., 2001) and Mon R2 (1.32: Kohno, Koyama, & Hamaguchi, 2002), but slightly
lower than that in the average interstellar space of our Galaxy (1.79: Predehl & Schmitt, 1995).
This indicates that the gas-to-dust ratio is lower in star-forming clouds than that in the interstellar
medium.

Fig. 7.6.— Relation between the absorption column (NH ) and visual extinction (AV : Bontemps etal., 2001, see also TableB.1). Circles, triangles, squares, diamonds, and stars represent class I, class
II, class III+IIIc , unclassified NIR, and BD/BDc sources, respectively. The dashed line shows thebest-fit linear model (see text).

7.4. CORRELATION BETWEEN X-RAY AND OTHER PHYSICAL PARAMETERS

7.4.2

103

< LX > vs Lbol

The ratio of the X-ray (< LX >) and bolometric (Lbol ) luminosities is often used to estimate
relative activity level of X-ray emission of various classes. We show the relation of < LX > andLbol for the ρ Oph X-ray sources with NIR counterpart in Figure 7.7. In the quiescent phase, mostsources have < LX > /Lbol of 10−5 –10−3 , which is comparable to those for other star-forming
clouds (Feigelson et al., 2002; Preibisch & Zinnecker, 2002) and is significantly larger than that for
the quiet sun (∼5×10−7 : Vaiana & Rosner, 1978). < LX >/Lbol for the flare X-rays is also large
and some flares exceed the saturation level of coronally active stars (< LX >/Lbol = 10−3 : Fleming
et al., 1995). These results directly indicate that YSOs have much higher X-ray activity level thanMSs. On the other hand, the magnetic B star S1 (= A-41) shows much lower < LX > /Lbol value of∼10−6 than, but almost comparable < LX > of ∼1029 ergs s−1 to the other (low-mass) YSOs. Thisis consistent with the ASCA results for intermediate-mass YSOs (Herbig Ae/Be stars: Hamaguchi,
2001), hence suggests that the X-ray emission mechanism for intermediate-mass YSOs is nearly thesame as that of low-mass YSOs, although the relative X-ray activity level normalized by the total
stellar emission is significantly low. It is also conceivable that S1 shows peculiar spectral features,indicating possible non-thermal emission (see §6.3.5).

Fig. 7.7.— Relation between the bolometric (Lbol ) and X-ray (< LX >) luminosities for the (left)
quiescent and (right) flare phases. Symbols are the same as Figure 7.6. Solid lines represent therelations of < LX >/Lbol = 10−1 , 10−3 , and 10−5 . The position of the magnetic B star S1 (=A-41) is indicated in the figure.

7.4.3< LX > vs Mass

Figure 7.8 shows the relation between < LX > and their stellar mass (M ). Within the rangeof M < 2M¯ , < LX > seems to be roughly proportional to M , which is the same as another

104

CHAPTER 7. OVERALL FEATURE OF X-RAY EMISSION FROM YSOS

star-forming region IC348 (< LX >∝ M 2 : Preibisch & Zinnecker, 2002). Since the stellar mass isroughly as a function of Lbol (Figure 2.1), this is basically the same as the constant < LX > /Lbol
relation (§7.4.2).

Fig. 7.8.— Relation between the stellar mass and X-ray luminosity (< LX >) for (left) quiescentand (right) flare phases. Symbols are the same as Figure 7.6. Dashed lines are the < LX >–mass
relation derived for X-ray sources in IC 348 (Preibisch & Zinnecker, 2002). The position of themagnetic B star S1 (= A-41) is indicated in the figure.

7.4.4< LX > vs LR

Figure 7.9 shows the relation between < LX > and monochromatic radio luminosity at 1.5 GHz(LR ergs s−1 Hz−1 ) for sources detected at this wavelengths (André et al., 1987; Stine et al., 1988;
Leous et al., 1991, see column 3 in Table 5.3). Possible log-linear correlation can be seen for boththe quiescent and flare X-rays except for A-75/BF-11; most sources have log(< LX > /LR ) of 13.5–
15.5, which is almost comparable to those of MS stars (Güdel & Benz, 1993). Although it is notclear that cm emission of the ρ Oph sources is thermal or non-thermal, the < LX > vs LR relation
indicates that the majority of its origin is basically the same as that of X-rays. One plausible idea
is that accelerated non-thermal electrons (see §2.3.1) emit synchrotron radio continuum, then heat
ambient matter which gives rise to the thermal X-ray flux. We also note that the background AGN
candidate LFAM21 (A-75 = BF-11, §6.6.3) has much lower log(< LX > /LR ) of ∼11 than theothers. This is the typical value of AGNs (Smith et al., 1999), hence further supports its AGN
nature. On the other hand, other unidentified sources detected at the cm wavelengths (A-4, A-31,
A-37, and BF-372 ) have large < LX > /LR comparable to YSOs and MSs. We therefore expectthat these are possible candidates of new YSOs, as already proposed for A-31 (§6.5).
2

BF-37 is not plotted in Figure 7.9 because we can not determine < LX > due to the limited statistics. However,BF-37 = LFAM32 has relatively small LR of ∼1.5×1014 ergs s−1 Hz−1 , hence would be in this class.

7.4. CORRELATION BETWEEN X-RAY AND OTHER PHYSICAL PARAMETERS

105

Fig. 7.9.— Relation between the radio (LR ) and X-ray (< LX >) luminosities for the (left)
quiescent and (right) flare phases. Symbols are the same as Figure 7.6. A shaded region represents
the approximate position of MS stars (Güdel & Benz, 1993). Four unidentified sources detected
at the cm wavelengths are also indicated (A-4, A-31, A-37, and A-75 = BF-11, see also §6.5 and§6.6.3).

Chapter 8

Systematic Study of YSO Flares
We detected 71 X-ray flares from the ρ Oph X-ray sources. Using this enormous sample, we makea systematic study on flare activity of low-mass YSOs in this chapter.

8.1

Flare Rate

Table 8.1 shows the number of sources and detected flares sorted by several physical parameterssuch as class, age, mass, and binarity. We should note that our flare detection method (§5.5) largely
depends on the quiescent flux level; the higher quiescent count rate results in a higher sensitivity
to smaller amplitude flares due to small statistical errors. We hence show the mean quiescent flux
for each class in the last column of Table 8.1. Following Stelzer et al. (2000), the flare rate (F ) isdefined asq√
στ2r + στ2d NF(τ¯r + τ¯d )NF
F =
±
,
(8.1)
N Tobs
N Tobswhere (τ¯r , τ¯d ) and (στr , στd ) are the mean value and uncertainty of (τr , τd ), N and NF are thenumber of sources and detected flares, and Tobs is the duration of the observations (≈100 ks). All
derived values of F are much higher than that derived with ROSAT (F ∼ 1 %: Stelzer et al. 2000),
which may be primarily due to the extended sensitivity in the hard X-ray band of the Chandra
observations, because the flare activity (flux increase) is clearer in the harder X-ray band.
We see larger flare rate of class III+IIIc than that of class II. However, this is simply due tothe effect of the flares with unusually long timescales (A-2, A-63, and BF-96). In fact, if we regard
those unusual flares are affected by other unknown components (unresolved flares, for example) andexclude them, F of class III+IIIc becomes ∼5.0 % (parentheses in Table 8.1), which is significantly
smaller than those of class I–II sources. Such a tendency is consistent with the ROSAT results (F
107

108

CHAPTER 8. SYSTEMATIC STUDY OF YSO FLARES= 1.1±0.4 % and 0.7±0.2 % for class II and III). Also, there is a relation that lower mass sources
show lower flare rate (∼5 %). However, this would be merely due to lower quiescent flux (∼= lowerluminosity) for lower mass sources (Figure 7.8 left), hence it is still controversial whether the flare
rate is affected by the stellar mass or not. In the last row of Table 8.1, we show the data of sources
previously revealed as binary systems (Simon et al., 1995; Haisch et al., 2002); A-22, A-23/A-24(see §6.3.2 and §6.3.3), A-41, A-78, A-82, BF-46, BF-61, and BF-96. Although the flare rate ofbinary sources is significantly larger than the others like the Taurus-Auriga-Perseus sources (Stelzer
et al., 2000), the effect of the flare detection sensitivity can not be negligible; the quiescent flux ofthe binary sources is ∼4–10 times larger than the others.

Table 8.1: Flare rate sorted by several physical parameters*

Parameter range

N*

NF †

Class IClass IIClass III+IIIc

8
63
27

8
36
16

t ≤ 0.1

11

11

0.1 < t ≤ 0.5
0.5 < t

20
12

15
7

M ≤ 0.2
0.2 < M ≤ 0.4

11
23

4
20

0.4 < M

99Binary#99Others

102

57

τ¯r ‡
τ¯d ‡
(ks)
(ks)
– Class –3.4±0.8
11.6±3.0
6.2±1.4
7.9±1.0
10.1±6.5
11.5±3.9(3.1±0.6)(6.9±0.8)– Age (Myr) –3.6±0.629.2±12.1(3.7±0.8)(8.6±2.0)
7.9±2.4
9.4±1.617.8±10.8
5.4±1.6(8.2±6.3)(4.0±1.0)– Mass (M¯ ) –3.3±1.2
9.7±3.78.4±2.2
11.4±3.1(8.7±2.3)(8.5±2.5)
9.3±6.728.7±15.8(2.5±0.5)(6.2±2.2)– Binarity –8.4±6.0
8.9±1.6(2.5±0.4)(8.3±1.7)
4.9±0.8
11.9±2.6(4.9±0.9)(7.9±0.9)F§(%)

log[Flux(Q)]k15.0±1.1
8.0±0.2
12.6±1.1(5.0±0.1)

−13.2
−13.6
−13.6

32.8±3.6(10.9±0.7)
13.0±0.6
13.5±2.4(6.6±1.4)−13.2

4.7±0.717.2±0.7(14.8±0.7)
38.0±5.7(7.4±0.8)

−14.2
−13.1

17.3±2.1(10.8±0.6)
9.8±0.2(7.2±0.1)

−12.9

−13.3
−13.4

−12.9

−13.7

The number of X-ray sources.
The number of detected flares (§5.5).
‡
Mean values of τr and τd . Parentheses indicate the mean values when we exclude the flares withunusually long timescales (A-2, A-63, and BF-96).
§
Flare rate defined by Eq.(8.1) (Stelzer et al., 2000). Parentheses indicate the mean values when
we exclude the flares with unusually long timescales (A-2, A-63, and BF-96).
k
Mean value of quiescent X-ray flux (ergs s−1 cm−2 ).
#Sources previously revealed as binary systems (Simon et al., 1995; Haisch et al., 2002).
†

8.2. CORRELATION BETWEEN THE FLARE X-RAY PARAMETERS

8.2

109

Correlation between the Flare X-ray Parameters

Figures 8.1 and 8.2 show the correlations between the derived parameters of flares; τr vs τd and< kT > vs τr . In these figures, we exclude those flares that could not determine errors mainly due
to the limited statistics. Possible positive and negative log-linear correlations can be seen in τr vsτd and < kT > vs τr , respectively. We hence check the significances of these correlations using
the Cox’s proportional hazard model (Isobe et al., 1986) in ASURV. This gives the significances of∼99 % (τr vs τd ) and ∼94 % (< kT > vs τr ). We further estimate the best-fit log-linear models
with ASURV to beτdτr) = 102.5±0.4 ( )0.4±0.1 ,
s
s

(8.2)

τr
< kT > −0.4±0.3) = 103.8±0.1 (
)
.
s
keV

(8.3)

(

(

These models are shown by the solid lines in Figures 8.1 and 8.2.

Fig. 8.1.— Relation between the rise and decay timescales of flares (τr and τd ). Symbols are
the same as Figure 7.6. The solid line represents the best-fit log-linear correlations derived with
ASURV (Eq.8.2). The dashed line is the best-fit model of Eq.(C.12).

110

CHAPTER 8. SYSTEMATIC STUDY OF YSO FLARES

Fig. 8.2.— Relation between the temperature and rise timescale of flares (kT and τr ). Symbols are
the same as Figure 7.6. The solid line represents the best-fit log-linear correlations derived with
ASURV (Eq.8.3). The dash-dotted and dashed lines in (b) are constant B and L lines derived by
Eqs.(C.14) and (C.15) with the assumption of nc = 1010.48 cm−3 .

8.3

Magnetic Reconnection Model

In order to explain the observed features of the flare parameters, we formulate τr , τd , and < kT >as a function of pre-flare (coronal) electronic density (nc ), half-length of the reconnected magnetic
loop (L) and strength of magnetic field (B), which are based on the idea of a standard magnetic
reconnection model (Petschek, 1964) and the balance between reconnection heating and conductive
cooling (Shibata & Yokoyama, 2002). For simplicity, we assume that τd is equal to the radiativeloss timescale (τrad ), although it would be possible that τd is slightly larger than τrad (see Appendix
C.4). We also assume that the reconnection rate MA (ratio of the reconnection interval and Alfvéntimescales; Eq.C.7) is 0.01, which is the mean value of MA for the solar observations (0.001–0.1,
Table 2 in Isobe et al. 2002). Details of the formulae are given in Appendix C.

8.3. MAGNETIC RECONNECTION MODEL

8.3.1

111

τr vs τd – Implication of the pre-flare density

From Eq.(C.12), the reconnection model predicts the positive correlation between τr and τd as
1/2
τd = Aτr . The slope of 1/2 shows good agreement with the observed value of 0.4±0.1 (Eq.8.2),
hence supports our assumption that the decay phase is dominant by radiative cooling. From
Eq.(C.12), we re-estimate log(A) with the fixed slope of 1/2 to be 2.06±0.04 (the dashed line in
Figure 8.1). We also calculate log(A) separately for each class, then obtain 2.1±0.1, 2.0±0.1, and
2.0±0.1 for class I, II, and III+IIIc , respectively. We see no significant difference between theclasses, hence A is assumed to be the same for all sources. Since A is a function of nc (Eq.C.13),
we can determine nc to beMAnc = 1010.48±0.09 () [cm−3 ].
(8.4)
0.01
Considering the possible range of MA (0.001–0.1), this favors higher pre-flare density for YSOs
(109 –1012 cm−3 ) than that for the sun (∼109 cm−3 ). In fact, Kastner et al. (2002) derived theplasma density of a CTTS (TW Hydrae) using the ratio of NeIX and OVII triplets and foundextremely high density of ∼1013 cm−3 even in their quiescent phase. Hence larger nc value for
YSOs than that for the sun would be common feature.

8.3.2< kT > vs τr – Loop length and magnetic field strength

The predicted correlations between < kT > and τr are shown in Eqs. (C.14) and (C.15). Positive(τr ∝< kT >7/2 ) and negative (τr ∝< kT >−7/6 ) correlations are expected for the constant B
and L values, respectively. The dash-dotted and dashed lines in Figure 8.2 show the constant B
and L lines, in which we uniformly assume nc = 1010.48 cm−3 (Eq.8.4). Our observed negative
correlation therefore predicts that flares with higher temperature is mainly due to larger magnetic
field. However, the observed slope of −0.4±0.3 is significantly flatter than −7/6 (Eq. C.15), hence
the effect of L would not be negligible. We estimate the mean B and L for each class (Table 8.2)
using the mean values of < kT > and τr (Figures 7.4b and 7.5a), which is summarized in Table
8.2. Younger sources tend to show larger B values; the mean values of B are ∼500, 300, and 200 G
for class I, II, and III+IIIc sources, respectively. This is consistent with the observed result ofhigher plasma temperature for younger sources. Also, the estimated values of L (1010 –1011 cm)
suggest that L is nearly the same among these classes, comparable to the typical radius of YSOs.
These results indicate that the flare loops for all classes are localized at the stellar surface, whichis in sharp contrast with the idea of the flares for class I sources triggered by the larger flare loops
connecting the star and the disk (Montmerle et al., 2000).
It is conceivable that the class III flares with unusually long timescales (A-2, A-63, and BF-96)
mainly affect the systematically lower B for class III sources. Since it is possible that two or more

112

CHAPTER 8. SYSTEMATIC STUDY OF YSO FLARES

flares make such unusually long timescales, we re-estimate B and L using the mean value of τrwithout these flares (§7.3). The results are shown by the parentheses in Table 8.2. Again, we
confirm the same results; class III+IIIc sources have lower B than, and comparable L (but slightly
lower) to class I and II sources.

Table 8.2: The estimated mean values of the flare physical parametersClasslog(nc )†‡(cm−3 )— This work —log(B)†
log(L)†
(G)
(cm)

I
II
III+IIIc

10.48±0.09...
...

Equation*(8.4)

2.7±0.1
2.5±0.1
2.3±0.1(2.4±0.1)k(C.14)

10.3±0.1
10.4±0.1
10.4±0.2(10.0±0.1)k(C.15)— kT vs EM —log(BSY )§ log(LSY )§(G)(cm)
2.6±0.1
2.5±0.1
2.2±0.1

10.4±0.2
10.4±0.1
10.7±0.1(C.18)(C.19)*

Equations in the text used for the estimation of each parameter.
The values are derived by assuming MA = 0.01‡
We assume the same values of nc for all classes (see text).
§
These values are estimated using the derived values of nc (= 1010.48 cm−3 , column 2).
k
The mean values for class III+IIIc when we exclude the flares with unusually long timescales (A-2,
A-63, and BF-96).
†

8.4

Comparison with the kT -EM scaling law

Shibata & Yokoyama (2002) showed that L and B of flares can be estimated from the kT –EM plot(see Appendix C.3). This method is completely independent of the flare timescales (τr and τd ),hence can be used as the consistency check of our estimation derived in §8.3. The < kT >–< EM >relation of the ρ Oph flares is shown in Figure 8.3. Using ASURV, we determine the mean valuesof log(< EM >) to be 53.77±0.20, 53.33±0.07, and 53.48±0.16 for class I, II, and III+IIIc sources,
respectively. We then calculate the mean values of the magnetic field strength and loop length (BSYand LSY ) by Eqs.(C.18) and (C.19) using the estimated value of nc = 1010.48±0.09 cm−3 (§8.3.1).
The results are shown in Table 8.2. The values of BSY (100–1000 G) are much higher than thosederived in Shibata & Yokoyama (2002) and Paper I (50–150G). This discrepancy is caused by two
different assumptions of nc value; 109 cm−3 for Shibata & Yokoyama (2002) and Paper I (typicalvalue of the solar corona), and 1010.48 cm−3 for this paper. BSY and LSY show good agreementwith those derived in §8.3.2 (B and L in Table 8.2); higher BSY values for younger sources and the8.5. EFFECT OF THE QUIESCENT X-RAYS ON THE FLARE ANALYSIS

113

same order of LSY comparable to the typical stellar radius. Therefore our estimation of nc , L, andB using the τr –τd and < kT >–τr relations is cross-checked.
The same conclusion of §8.3 and this section also indicates the propriety of our assumption of
MA = 0.01. If we assume the smaller (or larger) value of MA , nc (Eq.8.4) becomes smaller (larger),which causes smaller (larger) estimations of BSY (Eq.C.18) and makes larger discrepancy between
B and BSY (Table 8.2).

Fig. 8.3.— Plot of < kT > and < EM > in theflare phases. Symbols are the same as Figure 7.6.
The dashed-dotted and dashed lines are constant
B and L lines derived by Eqs.(C.18) and (C.19)
with the assumption of nc = 1010.48 cm−3 .

8.5

Effect of the Quiescent X-rays on the Flare Analysis

It is conceivable that the quiescent X-rays put significant influence on the results of our flare
analysis. Hence, in this subsection, we make and fit the flare spectra by subtracting the quiescent
spectra as a background level and compare the results with those in the previous section. For the
spectral fittings, the absorption column (NH ) is fixed to be the best-fit value in Table 5.2.
Figure 8.4 shows the histogram of the newly estimated < kT > in the flare phases. Since thestatistics are highly limited due to the higher background level than that for the previous analysis,
many faint flares can not constrain the 90 % confidence limit of < kT >. Hence, we only use thedata whose errors are well determined. Although the mean values of < kT > are systematicallyhigher and lower than the previous results, overall properties are the same; higher < kT > for class
I and II sources than class III+IIIc . Using ASURV, the significance level of higher < kT > for class
I and II sources is estimated to be & 92 %. This value is a bit smaller than the previous analysis(Table 7.2), which would be simply due to the limited samples for this analysis.
The plot of < kT > and τr is shown in Figure 8.5. Similar to Figure 8.4, we only use thedata whose 90 % confidence limit is well determined. The mean values of L are ∼2.6, 2.7, and3.1 ×1010 cm, while those of B are ∼600, 300, and 200 G for class I, II and III+IIIc , respectively.

114

CHAPTER 8. SYSTEMATIC STUDY OF YSO FLARES

Although there are slight differences between these values and the previous results, the relative
tendency is the same; higher B for younger sources and almost comparable L to the stellar size.
We hence conclude that, even if we use the quiescent spectra as the background level for theflare spectra, the main results in the previous section are not affected so seriously.

Fig. 8.4.— Same as Figure 7.4(b), but for the fitting results when the quiescent X-rays are used as
background spectra.

Fig. 8.5.— Same as Figure 8.2, but for the results
when the quiescent X-rays are used as background
spectra. Symbols are the same as Figure 7.6. The
dashed and solid lines are the constant L and B
lines derived by Eqs.(C.15) and (C.14).

8.6. EVOLUTION OF YSOS AND THEIR FLARE ACTIVITY

8.6

115

Evolution of YSOs and Their Flare Activity

Combining all the results discussed in the previous sections, we propose a simple view of theevolution of flare activity on low-mass objects as follows. In their earlier stage (class I), sources haverelatively strong magnetic field (≈500 G) and show frequent X-ray flares with higher temperature(≈5 keV). As stars evolve (class II and IIIs), the magnetic field gradually decreases (200–300 G)
and makes moderate temperature plasma (2–4 keV) via X-ray flares. During these phases (class
I–III), the length of the flare loop does not change significantly (1010 –1011 cm). As approaching
to the main-sequence stage, the magnetic field and length of the flare loop are becoming weak
(50–150 G for the sun) and short (108 –109 cm), which causes lower plasma temperature (0.1–1keV) and shorter flare timescales (10–100 s) as seen in the sun.

PMS

B(G)
binary?

1000

MS

log[L(cm)]RS CVn11
10

single

9
sun

100

4

5

6

7

8

8

log[Age(yr)]Fig. 8.6.— Schematic view of the evolution of X-ray flares of low-mass objects. Thick solid anddashed lines represent the predicted evolution of the magnetic field strength (B) and loop length(L), respectively.
We should note that the flare from the class 0 candidate A-29 (§6.5) has relatively short rise anddecay timescales (1–2 ks). Although we could not plot the data of the A-29 flare in the < kT >–τr
plane because < kT > is not determined due to the limited statistics, the small τr value predicts
that A-29 has a large magnetic field; if < kT > is larger than 10 keV, the expected B value is&1000 G, which may be the initial value of magnetic field in the YSO evolution.
Generally, RS CVn systems show smaller flare timescales than YSOs (∼100 s), while < kT >is comparable (∼5 keV). Based on the above idea, this may be due to larger B value. In fact,

116

CHAPTER 8. SYSTEMATIC STUDY OF YSO FLARES

Donati et al. (1990) estimated B of RS CVns as ≈1000 G, which is as large as the maximum value
for YSOs (Figures 8.2 and 8.3).

8.7

Comment on the Giant Flares and Quiescent X-rays

Figure 8.3 indicates that the giant flares (< EM > & 1055 cm−3 ) previously detected with ASCA
such as V773 Tau (class III: Tsuboi et al. (1998)) and ROXs31 (class III: §6.2), as well as BF-64= YLW16A (class I; < EM > ∼1055 cm−3 ), should have either large L or B, if the sustained
heating is negligible. Since the flares of ROXs31 and YLW16A show relatively large rise timescale(∼104 s), they may have large L values (1011 –1012 cm). The rise time of V773 Tau, on the other
hand, is among the shortest (∼103 s), then the large LX would be primarily due to the large B
value. In fact, the detailed flare decay time analysis for V773 Tau predicted extremely large B
value of &1000 G (Favata et al., 2001). We suspect that the binary nature of V773 Tau (Welty,1995) causes such an exceptionally large B, which would be the same case as the main-sequence
RS CVn systems.
Although the mechanism of making the quiescent X-ray emission from YSOs is still controversial, one possible idea is the composition of multiple unresolved faint flares. If we assume thatthe magnetic field with the same strength as those of the flares makes quiescent X-rays, unresolved flares are expected to have smaller loop length, considering the systematically lower < kT >and < LX > (∼= < EM >) values (Figures 7.2 and 7.4), which causes shorter flare timescales(Eq.C.8 and Eq.C.11). Such flares are hard to be detected because of the limited statistics. Future
observations with much higher sensitivity would reveal what makes quiescent X-rays.

Chapter 9

Conclusion
We summarize the main results of the systematic analysis of two Chandra observations of ρ Ophwith ≈100 ks exposure as follows;
1. From the central ∼1 pc × 2 pc region of ρ Oph, we detect 195 X-ray sources. About 60 % ofthem (110) are identified with optical and/or infrared counterpart, including class I, II, andIII sources and some young brown dwarfs. The X-ray detection rate is surprisingly high of∼60 %, hence we suspect that virtually all YSOs in ρ Oph may emit X-rays.
2. A total of 71 flares are detected. Most of the flares show the typical profile of the solar andstellar flares, while some bright flares have unusually long rise timescale.
3. All the X-ray spectra are well fitted by a single temperature thin plasma model of 0.3 solar
abundances except for some sources. The plasma temperature is 0.6–5 keV in the quiescentphase, then becomes systematically higher (1–10 keV) in the flare phase. Class I–II sources
show higher flare temperature than class III sources.
4. We find positive correlations of NH vs AV , < LX > vs Lbol , < LX > vs mass, and < LX >vs LR .
5. There is a positive correlation between τr and τd , which is well explained by the standard
magnetic reconnection model. Shorter τr and τd are due to the smaller L and/or larger B
values.
6. We find a negative correlation between < kT > and τr , which indicates the same order of theflare loop length regardless of their classes. Larger < kT > values for class I and II sources
are due to larger B values.
117

118

CHAPTER 9. CONCLUSION

7. The expected loop length is comparable to the stellar size (1010 –1011 cm), indicating that the
scenario of the star-disk arcade magnetic loop is unlikely.

Appendix A

Flare Light Curves

Fig. A.1.— Light curves of detected flares in 0.5–9.0 keV for class I sources, binned in 2000 sinterval. Solid lines are the best-fit exponential model. Source name for each flare is shown at theupper-left of the panels. The time axis starts at MJD = 51679.9944 (obs-A) and 51647.7848 (obsBF). Since the light curve of BF-64 suffer the photon pileup during the second flare, we show thetwo flares separately (the former is extracted from the 7.00 5 radius and the latter is from 2.00 5–7.00 5).
119

120

APPENDIX A. FLARE LIGHT CURVES

Fig. A.2.— Same as Figure A.1, but for class II sources.

121

Fig.A.2 (Continued)

122

APPENDIX A. FLARE LIGHT CURVES

Fig. A.3.— Same as Figure A.1, but for class III+IIIc sources. The light curve of A-2 is extracted
from a 2.00 5–12.00 5 radius circle in order to avoid the pileup effect.

123Fig. A.4.— Same as Figure A.1, but for unclassified NIR and unidentified sources.

Appendix B

Physical Parameters of YSOs
Here we summarize the physical parameters of YSOs. We estimate the stellar mass and age
for sources whose Lbol and Teff are well determined with the infrared spectroscopy (Luhman &Rieke, 1999), by comparing the position in the H-R diagram to the theoretical evolutionary model(D’Antona & Mazzitelli, 1997).
Table B.1:: Physical parameters of YSOs
No.

AV *
(mag)

A-49
BF-1
BF-26
BF-50
BF-61
BF-64
BF-68
BF-89
BF-S7

···
47k
57k
57k
56k
52k
37k
48k
···

A-3
A-6
A-14
A-15
A-19

19.5
22.6
9.4
1.8
17.1

Lbol †(L¯ )Teff ‡(K)— Class I —
0.17
···
2.6
···
26
···
1.7
···
6.7
···
8.7
···
0.62
···
1.1
···
0.36
···
— Class II —
0.22
···
0.95
3955
1.4
4205
0.72
3850
0.97
···
125

Sp.‡

Age§
(Myr)

Mass§
(M¯ )

···
···
···
···
···
K8
···
···
···

···
···
···
···
···
···
···
···
···

···
···
···
···
···
···
···
···
···

···
K8
K6
M0
···

···
0.5
0.3
0.5
···

···
0.40
0.35
0.30
···

126

APPENDIX B. PHYSICAL PARAMETERS OF YSOS
Table B.1 (Continued)
No.

AV *
(mag)

Lbol †(L¯ )Teff ‡
(K)

Sp.‡

Age§
(Myr)

Mass§
(M¯ )

A-22
A-23
A-24
A-25
A-27
A-34
A-44
A-53
A-54
A-59
A-61
A-64
A-66
A-69
A-72
A-73
A-77
A-78
A-79
A-82
A-83
BF-2
BF-3
BF-4
BF-7
BF-8
BF-9
BF-10
BF-12
BF-15
BF-16
BF-17
BF-20
BF-21

6.1
14.5
11.8
10.0
19.1
21.6
7.6
8.5
14.6
16.3
38.6
26.8
43.5
19.7
19.9
14.2
20.5
20.9
20.6
3.5
10.7
(= A-61)
10.4(= A-64)(= A-69)
14.2
17.2
···
25k
(= A-77)(= A-78)(= A-79)
1.5
22.5

5.9
0.25
3.7
4.5
0.13
0.11
0.28
1.6
0.064
1.1
1.1
0.023
0.17
0.55
0.12
0.050
0.034
0.15
0.087
4.0
0.57

5080
3850
3955
···
3010
···
···
3850
···
3955
···
···
···
3850
···
···
···
···
···
6590
···

K1
M0
K8
···
M5
···
···
M0
···
K8
···
···
···
M0
···
···
···
···
···
F4
···

0.5
3
0.1
···
0.3
···
···
0.07
···
0.2
···
···
···
0.7
···
···
···
···
···
30
···

1.40
0.50
0.50
···
0.12
···
···
0.30
···
0.30
···
···
···
0.35
···
···
···
···
···
1.40
···

0.19

···

···

···

···

0.007
0.059
29
44

···
3180
···
···

···
M4
···
···

···
3
···
···

···
0.16
···
···

0.039
0.12

···
···

···
···

···
···

···
···

127
Table B.1 (Continued)
No.

AV *
(mag)

Lbol †(L¯ )Teff ‡
(K)

Sp.‡

Age§
(Myr)

Mass§
(M¯ )

BF-24
BF-25
BF-27
BF-28
BF-31
BF-33
BF-35
BF-39
BF-41
BF-42
BF-44
BF-49
BF-51
BF-59
BF-63
BF-66
BF-71
BF-76
BF-78
BF-82
BF-85
BF-87
BF-88
BF-91††
BF-98

12.5
19.4
36.2
73.4
34.8
45.7
16.5
36.9
24.7
19.5
42.2
27.7
31k
23.7
9.8
19.5
26.8
23.0
10.8
44.9
10.1
14.8
6.4
11
4.1

0.5
···
···
···
···
···
0.5
···
···
0.07
···
···
···
0.5
1
1
0.1
···
0.07
···
0.2
···
0.5
0.3
···

0.35
···
···
···
···
···
0.35
···
···
0.25
···
···
···
0.30
0.40
0.50
0.35
···
0.30
···
0.30
···
0.35
0.15
···

A-2#
A-5
A-20
A-33
A-41
A-51
A-52
A-56

6.0
20.6
19.5
25.2
11.7**
55.5
48.5
15.2

0.78
3955
K8
0.075
···
···
1.1
···
···
52
···
···
1.1
···
···
0.34
···
···
0.86
3850
M0
0.99
···
M4
0.023
···
···
1.1
3595 M1.5
1.3
···
···
3.1
···
···
···
···
···
0.71
3680
M1
0.42
3850
M0
0.66
4132 K6.5
1.9
3955
K8
0.29
···
···
1.6
3955
K8
1.5
···
···
1.1
3955
K8
0.027
···
···
0.80
3850
M0
0.20
3090
M4
0.045
···
···
— Class III —
30
4800
K1
0.70
···
···
0.89
3955
K8
0.34
3595 M1.5
1100 18700
B3
2.0
···
···
1.1
···
···
0.83
3850
M0

0.08
···
0.5
1
0.2
···
···
0.5

1.6
···
0.40
0.35
7.5
···
···
0.35

128

APPENDIX B. PHYSICAL PARAMETERS OF YSOS
Table B.1 (Continued)
No.

AV *
(mag)

Lbol †(L¯ )Teff ‡
(K)

Sp.‡

Age§
(Myr)

Mass§
(M¯ )

A-63
A-71
A-81
A-84
BF-19
BF-40
BF-46#
BF-75
BF-84
BF-96#

15.7
22.5
36.5
17.1
(= A-81)
59.8
1.2
27.5
11.5
6.0

2.8
0.84
1.6
1.2

4205
3955
···
4205

K6
K8
···
K6

0.07
0.5
···
0.5

0.40
0.35
···
0.50

39
···
···
···
0.63
3800
M0
0.6
0.90
···
···
···
0.60
3180
M4
0.07
0.50
3890
K7
1
— Class IIIc —
17.4
0.64
3595 M1.5
0.2
21.7
0.45
···
···
···
16.6
0.21
···
···
···
14.0
0.025
···
···
···
10.8
0.004
···
···
···
47.8
1.3
···
···
···
13.1
0.32
···
···
···
4.8
0.32
3180
M4
0.7
16.3
0.43
···
···
···
— Unclassified NIR sources —
···
···
···
···
···
56
0.71
2840 M5.5 <0.1
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
k
38
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
···
— BD —

···
0.37
···
0.18
0.47

A-12
A-32
A-36
A-45
BF-56
BF-69
BF-72
BF-79
BF-90
A-16
A-26††
A-28
A-46
A-60
A-76
A-80
BF-22
BF-55
BF-57
BF-62
BF-74

0.25
···
···
···
···
···
···
0.20
···
···
0.10
···
···
···
···
···
···
···
···
···
···

129
Table B.1 (Continued)
No.

AV *
(mag)

Lbol †(L¯ )Teff ‡
(K)

Sp.‡

Age§
(Myr)

Mass§
(M¯ )

BF-86††
BF-S2‡‡

5.7
0.8

M8.5
M8.5

<0.1
···

0.04
<0.02

A-30††
A-38††
A-S2††

6.3
11
4.5

0.055 2350
0.008 2590
— BDc —
0.024 2760
0.051 2760
0.063 2590

M6
M6
M7

2.0
1.0
0.2

0.06
0.08
0.06

*

Visual extinction derived from an IR color (Bontemps et al., 2001).
Bolometric luminosity estimated from an intrinsic NIR–MIR band magnitude. (Bontemps et al., 2001).
‡
Effective surface temperature and its spectral type derived by an NIR spectroscopy (Luhman & Rieke, 1999).
§
Stellar age and mass derived from the locus in the H-R diagram (Figure 2.1).
kCasanova et al. (1995).
#Nürnberger et al. (1998).
**
Chini (1981).
††
Wilking et al. (1999).
‡‡
Cushing et al. (2000).
†

Appendix C

Modeling of the FlareIn this appendix, we show details for the estimation of the flare parameters (τr , τd , and kT ), which
are used in chapter 8. These formulations are based on the standard magnetic reconnection model(Petschek, 1964), as well as the observed results and analysis of solar flares.

C.1

Estimation of the Flare Parameters• Plasma temperature (< kT >)Fisher & Hawley (1990) showed that the flare maximum temperature, Tmax [K], is determined by the balance between the heating rate (Q ergs s−1 cm−3 ) and conduction cooling
7/2(∝ Tmax /L2 ), then obtained an equation,
2

2QL 2/7) ,Tmax ∼
=(
κ0

(C.1)

where L and κ0 are the semi-length of the flare loop and Spitzer’s thermal conductivity (≡9.0×10−7 ; Spitzer 1956). For the magnetic reconnection, Q is equal to the Poynting flux
divided by L, i.e.,c(B × v) × B
B2v
Q∼
E×B=
=
,
(C.2)
=
4πL
4πL
4πL
where E, B, and v are electric field strength, magnetic field strength, and fluid velocity,respectively. v is assumed to be the Alfvén velocity;Bv = vA ≡ √
,
4πmnc

(C.3)

where m and nc are the proton mass (≡ 1.67×10−24 g) and pre-flare (coronal) density, respectively. Numerical simulations (Yokoyama & Shibata, 1998) further showed that the observed
131

132

APPENDIX C. MODELING OF THE FLARE
maximum temperature (T ) is a factor of ∼3 lower than Tmax . Finally, Shibata & Yokoyama(1999, 2002) obtained the following relation;
1
nc
L
B 6/7
T ∼
)−1/7 ( 11
)2/7 ()[K],= Tmax ∼= 6.75 × 107 ( 9
−3
3
10 cm
10 cm
100 G

(C.4)

As for the usual stellar flares, we can only determine the time-averaged temperature (< kT >)due to the limited statistics. However, van den Oord et al. (1988) showed that the behaviorof the flare temperature is exponential, hence we obtain the relation
T = R0−τr

τr + τdR τdet/τr dt +0e−t/τd dt· < T >∼
= 1.6 < T >(C.5)

Using Eqs. (C.4) and (C.5), < kT > is determined as< kT >∼= 3.64(L
B 6/7
nc)−1/7 ( 11
)2/7 ()[keV].
9
−3
10 cm
10 cm
100 G

(C.6)

• Rise timescale (τr )On the basis of a standard reconnection model, τr is equal to the interval of magnetic reconnection. Petschek (1964) showed that the reconnection timescale is proportional to the
Alfvén time, τA ≡ L/vA . Hence τr is,√
τA
4πmnc L
τr =
=
(C.7)
MA
MA B
The correction factor MA , sometimes referred as the reconnection rate, is estimated to be
0.01–0.1 (Petschek, 1964), while the observations for the solar flares show that MA is in therange of 0.001–0.1 (Table 2 in Isobe et al. 2002), regardless of their flare size. τr is therefore,τr ∼= 1.45 × 104 (

nc
B −1
MA −1
L
) ( 9
)(
) [s].
)1/2 ( 11
−3
0.01
10 cm
10 cm 100 G(C.8)• Decay timescale (τd )We assume that τd is nearly the same as the radiative cooling timescale (τrad ), i.e,
3nkT
τd ∼
,
= τrad ≡ 2n Λ(T )

(C.9)

where n and Λ(T ) are the maximum plasma density and radiative loss function given by
10−24.73 T 1/4 for T > 20 MK (Mewe et al., 1985). Furthermore, we assume that magnetic
pressure is comparable to plasma pressure at the flare peak;
2nkT =

B2.
8π(C.10)

Using Eqs.(C.4), (C.9), and (C.10), we obtain τd asτd ∼= 7.75 × 104 (nc
L
B −1/2)−1/4 ( 11
)1/2 (
)
[s].
9
−3
10 cm
10 cm
100 G

(C.11)

C.2. PREDICTED CORRELATIONS BETWEEN THE FLARE PARAMETERS

C.2

133

Predicted Correlations between the Flare Parameters• τr vs τdUsing Eqs.(C.8) and (C.11), we obtain a relationτrτd) = A( )1/2 ,ss(C.12)

MA 1/2nc)−1/2 (
) .
9
−3
10 cm
0.01

(C.13)(
whereA∼= 640(

Hence the positive correlation is expected between τr and τd .
• < kT > vs τr
From Eqs.(C.6), (C.8) and (8.4), we obtainnc
MA −1 B −4 < kT > 7/2
) (
) (
)
[s],
)(
109 cm−3 0.01
100 G
keV

(C.14)nc
L< kT > −7/6
MA −1
) ( 11)4/3 ()[s].
)1/3 (
9
−3
10 cm
0.01
10 cm
keV

(C.15)

τr ∼= 102.20 (τr ∼= 104.82 (These equations indicate that the positive and negative log-linear correlations of < kT > vsτr are expected if B and L are equal for all sources, respectively.

C.3

The kT –EM Scaling Law

Using Eqs.(C.4), (C.10), and assuming V ∼= L3 (V : plasma volume), Shibata & Yokoyama (2002)
derived the relation between the flare maximum temperature (T ) and emission measure (EM , Eqs.
5–6 in Shibata & Yokoyama 2002). Similar to Eq.(C.5), the time behavior of EM is also exponential(van den Oord et al., 1988), hence we assume EM ∼= 1.6 < EM >. We thus replace Eqs.(5) and(6) in Shibata & Yokoyama (2002) asnc
B −5 < kT > 17/2
)3/2 () ()[cm−3 ],
109 cm−3
100 G
keV(C.16)nc
L< kT > 8/3)2/3 ( 11)5/3 ()[cm−3 ].
9
−3
10 cm
10 cm
keV

(C.17)< EM >∼= 3.7 × 1048 (< EM >∼= 7.0 × 1051 (We also derive the following two equations
B∼
= 52(

< EM > −1/5 < kT > 17/10
nc)3/10 ( 50
)
(
)
[G],
9
−3
10 cm
10 cm−3keVL∼= 7.8 × 109 (nc< EM > 3/5 < kT > −8/5)−2/5 ( 50
) (
)
[cm].
9
−3
10 cm
10 cm−3
keV(C.18)(C.19)134

C.4

APPENDIX C. MODELING OF THE FLARE

Possible Error for the Derived Parameters

In the discussions of the main text (chapter 8), the assumption of τd (Eq.C.9) includes relatively
large uncertainty. Reale et al. (1997) proposed that τd becomes about a factor of . 10 larger than
the radiative loss timescale (τrad ), if the sustained heating exists during the decay. The dependencesof the relevant parameters on τd are;(C.20)
nc ∝ τd−2−1/2

B ∝ n1/4∝ τdc.

L ∝ n−1/2∝ τd1c(C.21)(C.22)

Accordingly, nc has large dependence on τd ; larger nc value than ∼1010.5 cm−3 may be conceivable.
The dependence of L is also relatively large. However, this uncertainty makes L much smaller,
hence larger flare loops is still less possible.
Another uncertainty is in the kT –EM scaling law. Shibata & Yokoyama (2002) showed thatthe effect of the filling factor f (V = f L3 ) may not be negligible. This gives possible errors for thevalues in columns [5]–[6] of Table 8.2. The dependences of the relevant parameters on f are;
BSY ∝ EM −1/5 ∝ f 1/5(C.23)

LSY ∝ EM 3/5 ∝ f −3/5

(C.24)

Hence slightly smaller and larger values of BSY and LSY would be conceivable, although the effectis only a factor of < 5 (if f = 0.1).

Bibliography
Ageorges, N., Eckart, A., Monin, J. L., & Ménard, F. 1997, A&A, 326, 632
Anders, E., & Grevesse, N. 1989, Geochimica et Cosmochimica Acta, 53, 197
André, P., Montmerle, T., & Feigelson, E. D. 1987, AJ, 93, 1182
André, P., Montmerle, T., Feigelson, E. D., Stine, P. C., & Klein, K. 1988, ApJ, 335, 940
André, P., Ward-Thompson, D., & Barsony, M. 1993, ApJ, 406, 122
André, P. 1994, in The Cold Universe, ed. T. Montmerle, C. J. Lada, I. F. Mirabel, J Trân ThanhVân (France: Gif-sur-Yvette), 179
André, P., & Montmerle, T. 1994, ApJ, 420, 837
Audard, M., Güdel, M., & Mewe, R. 2001, A&A, 365, L318
Awaki, H., Koyama, K., Inoue, H., Halpern, J. P. 1991, PASJ, 43, 195
Bachiller, R. 1996, ARA&A, 34, 111
Barsony, M., Burton, M. G., Russell, A. P. G., Carlstrom, J. E., & Garden, R. 1989, ApJ, 346, L93
Barsony, M. 1994, ASP Conf. Ser. 65, Clouds, Cores, and Low Mass Stars, D. P. Clemens & R.
Barvainis (San Francisco: ASP), 197
Barsony, M., Kenyon, S. J., Lada, E. A., & Teuben, P. J. 1997, ApJS, 112, 109
Basri, G. 2000, ARA&A, 38, 485
Bevington, P. R., & Robinson, D. K. 1992, Data Reduction and Error Analysis for the PhysicalScience (2d ed.; Boston: WCB/McGraw-Hill)
Bontemps, S. et al.2001, A&A, 372, 173
Bouvier, J., & Appenzeller, I. 1992, A&AS, 92, 481
135

136

BIBLIOGRAPHY

Brandner, W. et al. 2000, A&A, 364, L13
Brinkman, A. C. et al.2001, A&A, 365, L324
Burke, B. E., Mountain, R. W., Harrison, D. C., Bautz, M. W., Doty, J. P., Ricker, G. R., Daniels,
P. J. 1991, IEEE Trans. ED-38, 1069
Carkner, L., Feigelson, E. D., Koyama, K., Montmerle, T., Reid, I. N. 1996, ApJ, 464, 286
Casanova, S., Montmerle, T., Feigelson, E. D., & André, P. 1995, ApJ, 439, 752
Chen, H., Myers, P. C., Ladd, E. F., & Wood, D. O. S. 1995, ApJ, 445, 377
Chen, H., Grenfell, T. G., Myers P. C., & Hughes, J. D. 1997, ApJ, 478, 295
Chini, R. 1981, A&A, 99, 346
Comerón, F., Rieke, G. H., Burrows. A., & Rieke, M. J. 1993, ApJ, 416, 185
Costa, V. M., Lago, M. T. V. T., Norci, L., and Neurs, E. J. A. 2000, A&A, 354, 621
Cushing, M. C., Tokunaga, A. T., & Kobayashi, N. 2000, AJ, 119, 3019
Cutri, R., et al. 2000, http://www.ipac.caltech.edu/2mass/releases/second/doc/
Dame, T. M., et al., 1987, ApJ, 322, 706
Dame, T. M., Hartmann, D., & Thaddeus, P. 2001, ApJ, 547, 792
D’Antona, F. & Mazzitelli, I. 1997, in Mem. Soc. Astron. Italiana, 68, 807
de Zeeuw, P. T., Hoogerwerf, R., de Bruijne, J. H. J., Brown, A. G. A., & Blaauw, A. 1999, AJ,
117, 354
Dolidze, M. V. & Arakelyan, M. A. 1959, Astronomicheskii Zhurnal, 36, 444
Donati, J. F., Semel, M., Rees, D. E., Taylor, K., & Robinson, R. D. 1990, A&A, 232, L1
Durney, B. R., De Young, D. S., & Roxburgh, I. W. 1993, Sol. Phys., 145, 207
Elias, J. H. 1978, ApJ, 224, 453
Favata, F., & Schmitt, J. H. M. M. 1999, A&A, 350, 900
Favata, F., Micela, G., & Reale, F. 2001, A&A, 375, 485
Favata, F., Fridlund, C. V. M., Micela, G., Sciortino, S., & Kaas, A. A. 2002, A&A, 386, 204

BIBLIOGRAPHY

137

Feigelson, E. D., & DeCampli, W. M. 1981, ApJ, 243, L89
Feigelson, E. D., & Kriss, G. A. 1981, ApJ, 248, L35
Feigelson, E. D., Nelson, P. I. 1985, ApJ, 293, 192
Feigelson, E. D., & Montmerle, T. 1999, ARA&A, 37, 363
Feigelson, E. D. et al.2002, ApJ, 574, 258
Feldman, U. 1992, Phys. Scr, 46, 202
Festin, L. 1998, A&A, 336, 883
Fisher, G. H. & Hawley, S. L. 1990, ApJ, 357, 243
Fleming, T. A., Schmitt, J. H. M. M., & Giampapa, M. S. 1995, ApJ, 450, 401
Freeman, P. E., Kashyap, V., Rosner, R., Lamb, D. Q. 2002, ApJS, 138, 185
Gagné, M. 2001, in ASP Conf. Ser., High Energy Universe at Sharp Focus: Chandra Science, ed.
E. M. Schlegel & S. Vrtilek, 4Gehrels, N. 1986, ApJ, 303, 336
Giampapa, M. S., Rosner, R., Kashyap, V., Fleming, T. A., Schmitt, J. H. M. M., & Bookbinder,
J. A. 1996, ApJ, 463, 707
Girart, J. M., Rodriguez, L. F., & Curiel, S. 2000, ApJ, 544, L153
Gómez, M., Whitney, B. A., & Wood, K. 1998, AJ, 115, 2018
Grasdalen, G. L., Strom, K. M., & Strom, S. E. 1973, ApJ, 184, L53
Greene, T. P. & Young, E. T. 1992, ApJ, 395, 516
Grosso, N., Montmerle, T., Feigelson, E. D., André, P., Casanova, S., & Gregorio-Hetem, J. 1997,
Nature, 387, 56
Grosso, N., Montmerle, T., Bontemps, S., André, P., & Feigelson, E. D. 2000, A&A, 359, 113
Grosso, N. 2001, A&A, 370, L22
Grosso, N., Alves, J., Neuhäuser, R., & Montmerle, T. 2001, A&A, 380, L1
Güdel, M. & Benz, A. O. 1993, ApJ, 405, L63
Güdel, M. et al.2001a, A&A, 365, L336

138

BIBLIOGRAPHY

Güdel, M., Audard, M., Magee, H., Franciosini, E., Grosso, N., Cordova, F. A., Pallavicini, R., &Mewe, R. 2001b, A&A, 365, L344
Güdel, M. et al.2001c, astro-ph/0109267
Haisch, K. E., Jr., Barsony, M., Greene, T. P., & Ressler, M. E. 2002, AJ, 124, 2841
Hamaguchi, L., Terada, H., Bamba, A., & Koyama, K. 2000, ApJ, 532, 1111
Hamaguchi, K., ph.D. thesis, Kyoto University
Hamaguchi, K., Corcoran, M. F., & Imanishi, K. submitted to PASJ
Hartman, R. C. et al.1999, ApJS, 123, 79
Hayashi, C. 1966, ARA&A, 4, 171
Hayashi, M. R., Shibata, K., & Matsumoto, R. 1996, ApJ, 468, L37
Henyey, L. G., Lelevier, R., & Levée, R. D. 1955, PASP, 67, 154
Howard, E. M., Pipher, J. L., & Forrest, W. J. 1994, ApJ, 425, 707
Huenemoerder, D. P., Canizares, C. R., & Schulz, N. S. 2001, ApJ, 559, 1135
Imanishi, K., Koyama, K., & Tsuboi, Y. 2001a, ApJ, 557, 747
Imanishi, K., Tsujimoto, M., & Koyama, K. 2001b, ApJ, 563, 361
Imanishi, K., Tsujimoto, M., & Koyama, K. 2002, ApJ, 572, 300
Imanishi, K., Tsujimoto, M., & Koyama, K. 2002b, Proc. of The IAU 8th Asian-Pacific Regional
Meeting, 149.
Imanishi, K., Nakajima, H., Tsujimoto, M., Koyama, K., & Tsuboi, Y. PASJ, accepted
Inoue, H. 1985, Space Sci. Rev., 40, 317
Isobe, T., Feigelson, E. D., & Nelson, P. I. 1986, ApJ, 306, 490
Isobe, H., Yokoyama, T., Shimojo, M., Morimoto, T., Kozu, H., Eto, S., Narukage, N., & Shibata,K., 2002, ApJ, 566, 528
Jerius, D., Freeman, M., Gaetz, T., Hughes, J. P., & Podgorski, W. 1995, ASP Conf. Ser. 77:
Astronomical Data Analysis Software and Systems IV, 4, 357

BIBLIOGRAPHY

139

Johnstone, D., Wilson, C. D., Moriarty-Schieven, G., Joncas, G., Smith, G., Gregersen, E., & Fich,
M. 2000, ApJ, 545, 327
Kaastra, J.S. 1992, An X-Ray Spectral Code for Optically Thin Plasmas (Internal SRON-Leiden
Report, updated version 2.0)
Kamata, Y., Koyama, K., Tsuboi, Y., & Yamauchi, S. 1997, PASJ, 49, 461
Kamazaki, T., Saito, M., Hirano, N., & Kawabe, R. 2001, ApJ, 548, 278
Kastner, J. H., Huenemoerder, D. P., Schultz, N. S., Canizares, C. R., & Weintraub, D. A. 2002,
ApJ, 567, 434
Knude, J. & Høg, E. 1998, A&A, 338, 897
Kohno, M., Koyama, K., & Hamaguchi, K. 2002, ApJ, 567, 423
Koyama, K. 1987, PASJ, 39, 245
Koyama, K., Asaoka, I., Kuriyama, T., & Tawara, Y. 1992, PASJ, 44, L255
Koyama, K., Maeda, Y., Ozaki, M., Ueno, S., Kamata, Y., Tawara, Y., Skinner, S., & Yamauchi,
S. 1994, PASJ, 46, L125
Koyama, K., Hamaguchi, K., Ueno, S., Kobayashi, N., & Feigelson, E. D. 1996, PASJ, 48, L87
Lada, C. J., 1991, in The Physics of Star Formation and Early Stellar Evolution, NATO ASI, ed.
C. J. Lada & N. D. Kylafis, Kluwer, 329
Leous, J. A., Feigelson, E. D., André, P., & Montmerle, T. 1991, ApJ, 379, 683
Liedahl, D.A., Osterheld, A.L., & Goldstein, W.H. 1995, ApJ, 438, 115
Loren, R. B. 1989, ApJ, 338, 902
Loren, R. B., Wootten, A., & Wilking, B. A. 1990, ApJ, 365, 269
Lorenzani, A., & Palla, F. 2001, ASP Conf. Ser. 243, From Darkness to Light, Montmerle, T., &André, P. (San Francisco: ASP), 745
Luhman, K. L. & Rieke, G. H. 1999, ApJ, 525, 440
Mewe, R., Gronenschild, E. H. B. M., & van den Oord, G. H. J. 1985, A&AS, 62, 197
Montmerle, T., Koch-Miramond, L., Falgarone, E., & Grindlay, J. E. 1983, ApJ, 269, 182
Montmerle, T., Grosso, N., Tsuboi, Y., & Koyama, K. 2000, ApJ, 532, 1097

140

BIBLIOGRAPHY

Morrison, R. & McCammon, D. 1983, ApJ, 270, 119
Motte, F., André, P., & Neri, R. 1998, A&A, 336, 150
Mundy, L. G., Wootten, H. A., & Wilking, B. A. 1990, ApJ, 352, 159
Mushotzky, R. F., Cowie, L. L., Barger, A. J., & Arnaud, K. A. 2000, Nature, 404, 459
Nakajima, H., Imanishi, K., Takagi, S., Tsujimoto, M., & Koyama, K. 2003, submitted to PASJ
Neuhäuser, R. & Preibisch, T. 1997, A&A, 322, L37
Neuhäuser, R. & Comerón, F. 1998, Nature, 282, 83
Neuhäuser, R. et al. 1999, A&A, 343, 883
Nürnberger, D., Brandner, W., Yorke, H. W., Zinnecker, H. 1998, A&A, 330, 549
Ohashi, T. et al. 1996, PASJ, 48, 157
Ozawa, H., Nagase, F., Ueda, Y., Dotani, T., & Ishida, M. 1999, ApJ, 523, L81
Ozawa, H. 2000, ph.D. thesis, University of Tokyo
Parker, E. N. 1955, ApJ, 122, 293
Parker, E. N. 1957, J. Geophys. Res., 62, 509
Parker, E. N. 1970, ApJ, 162, 665
Petschek, H. E. 1964, in Proc. of AAS-NASA Symposium on the Physics of Solar Flares, ed. W.
N. Ness, NASA, SP-50, 425
Pravdo, S. H., Feigelson, E. D., Garmire, G., Maeda, Y., Tsuboi, Y., & Bally, J. 2001, Nature, 413,
708
Predehl, P., & Schmitt, J. H. M. M. 1995, A&A, 293, 889
Preibisch, T. 1997, A&A, 320, 525
Preibisch, T. & Zinnecker, H. 2002, AJ, 123, 1613
Press, W. H., Teukolsky, S. A.,Vetterling, W. T., & Flannery, B. P. 1992, Numerical Recipes in C
(2d ed.; Cambridge: Cambridge Univ. Press)
Reale, F., Betta, R., Peres, G., Serio, S., & McTiernan, J. 1997, A&A, 325, 782
Rodriguez, L. F., Myers, P. C., Cruz-Gonzalez, I., & Terebey, S. 1989, ApJ, 347, 461

BIBLIOGRAPHY

141

Rutledge, R. E., Basri, G., Martin, E. L., & Bildsten, L. 2000, ApJ, 538, L141
Schmelz, J. T. 1993, ApJ, 408, 373
Schmitt, J. H. M. M., & Favata, F. 1999, Nature, 401, 44
Sekimoto, Y., Tatematsu, K., Umemoto, T., Koyama, K., Tsuboi, Y., Hirano, N., & Yamamoto,S. 1997, ApJ, 489, L63
Serlemitsos, P. J. et al. 1995, PASJ, 47, 105
Shemi, A. 1991, MNRAS, 251, 221
Shevchenko, V. S., & Herbst, W. 1998, AJ, 116, 1419
Shibata, K. & Yokoyama, T. 1999, ApJ, 526, L49
Shibata, K. & Yokoyama, T. 2002, ApJ, 577, 742
Silva, A. V. R., White, S. M., Lin, R. P., de Pater, I., Shibasaki, K., Hudson, H. S., & Kundu, M.
R. 1996, ApJ, 458, L49
Simon, M. et al.1995, ApJ, 443, 625
Skinner, S. L. 2000, in IAU Symp. No.200, Birth and Evolution of Binary Stars, ed. B. Reipurth& H. Zinnecker (Boston: Kluwer Academic Publishers), 97
Smith, K., Güdel, M., & Benz, A. O. 1999, A&A, 349, 475
Spitzer, L. 1956, Physics of Fully Ionized Gases (New York: Interscience)Stelzer, B., Neuhäuser, R., & Hambaryan, V. 2000, A&A, 356, 949
Stine, P. C., Feigelson, E. D., André, P., & Montmerle, T. 1988, AJ, 96, 1394
Strom, K. M., Kepner, J., & Strom, S. E. 1995, ApJ, 438, 813
Struve, O. & Rudkjobing, M. 1949, ApJ, 109, 92
Swanenburg, B. N. et al.1981, ApJ, 243, 69
Sweet, P. A. 1958, in IAU Symp. No.6, Electromagnetic Phenomena in Cosmical Physics, ed. B.
Lenhnert (Cambridge Univ. Press), 123
Tachihara, K., Mizuno, A., & Fukui, Y. 2000, ApJ, 528, 817
Takagi, S., Murakami, H., & Koyama, K. 2002, ApJ, 573, 275

142

BIBLIOGRAPHY

Tanaka, Y., Inoue, H., & Holt, S. S. 1994, PASJ, 46, L37
Tsuboi, Y., Koyama, K., Murakami, H., Hayashi, M., Skinner, S., and Ueno, S. 1998, ApJ, 503,
894
Tsuboi, Y. 1999, PhD thesis, Kyoto University
Tsuboi, Y., Imanishi, K., Koyama, K., Grosso, N., & Montmerle, T. 2000, ApJ, 532, 1089
Tsuboi, Y., Koyama, K., Hamaguchi, K., Tatematsu, K., Sekimoto, Y., Bally, J., & Reipurth, B.
2001, ApJ, 554, 734
Tsujimoto, M., Koyama, K., Tsuboi, Y., Goto, M., & Kobayashi, N. 2001, ApJ, 566, 974
Tsujimoto, M., Koyama, K., Tsuboi, Y., Chartas, G., Goto, M., Kobayashi, N., Terada, H., &Tokunaga, A. T. 2002, ApJ, 573, 270
Tsujimoto, M. 2003, PhD thesis, Kyoto University
Tsuneta, S., Hara, H., Shimizu, T., Acton, L. W., Strong, K. T., Hudson, H. S., & Ogawaka, Y.
1992, PASJ, 44, L63
Vaiana, G. S. & Rosner, R. 1978, ARA&A, 16, 393
van den Oord, G. H. J., Mewe, R., & Brinkman, A. C. 1988, A&A, 205, 181
Vrba, F. J., Strom, K. M., Strom, S. E., & Grasdalen, G. L.
Waters, L. B. F. M. & Waelkens, C. 1998, ARA&A, 36, 233
Welty, A. D. 1995, AJ, 110, 776
Weisskopf, M. C., Brinkman, B., Canizares, C., Garmire, G., Murray, S., van Speybroeck, L. P.
2002, PASP, 114, 1
Wilking, B. A. & Lada, C. J. 1983, ApJ, 274, 698
Wilking, B. A. Schwartz, R. D., & Blackwell, J. H. 1987, AJ, 94, 106
Wilking, B. A., Lada, C. J., & Young, E. T. 1989, ApJ, 340, 823
Wilking, B. A., Greene, T. P., & Meyer, M. R. 1999, AJ, 117, 469
Xu, J. J., Wang, Z. R., & Qu, Q. Y. 1992, A&A, 256, 483
Yamauchi, S., Koyama, K., Sakano, M., Okada, K. 1996, PASJ, 48, 719

BIBLIOGRAPHY
Yokoyama, T. & Shibata, K. 1998, ApJ, 494, L113
Yokoyama, T., Akita, K., Morimoto, T., Inoue, K., & Newmark, J. 2001, ApJ, 546, L72
Yokoyama, T. & Shibata, K. 2001, ApJ, 549, 1160
Young, E. T., Lada, C. J., & Wilking, B. A. 1986, ApJ, 304, L45
Yu, T., & Chernin, L. M. 1997, ApJ, 479, L63
Zhang, Q., Wootten, A., & Ho, P. T. 1997, ApJ, 475, 713

143

Acknowledgment
I am deeply grateful to Prof. K. Koyama for his continuous guidance and support throughout thefive years of my graduate school period. I also thank my collaborators, Dr. Y. Tsuboi, Dr. K.
Hamaguchi, Dr. M. Tsujimoto, and Mr. H. Nakajima for their helpful and constructive discussions.
I would like to thank Prof. G. Garmire, Prof. E. D. Feigelson, Dr. L. Townsley, and Dr.
Y. Maeda for kind hospitality to the Chandra data analysis at the Pennsylvania State University.
Many helpful comments and discussions on X-ray observations of YSOs were given by Prof. T.
Montmerle, Prof. T. Preibisch, Prof. S. Yamauchi, Dr. N. Grosso, Dr. M. Güdel, and Dr. M.
Audard. Results of the solar X-ray observations, numerical simulations, and magnetic reconnection
theory were shown by Prof. K. Shibata, Dr. T. Yokoyama, Dr. T. Morimoto, and Mr. H. Isobe. I
also thank Dr. F. Motte, Dr. Y. Kamazaki, and Dr. K. Tachihara for their advice about the radio
observations on the ρ Ophiuchi cloud. Prof. T. Tanimori, Prof. T. Nakamura, Prof. T. Tanaka,and Dr. H. Matsumoto give me constructive referring and comments on this thesis.
I am grateful to all the members of the Cosmic-ray laboratory in Kyoto University for theircontinuous discussions, encouragements, and technical advice.
The Chandra data were obtained through the Chandra X-ray Observatory Science Center(CXC) operated for NASA by the Smithsonian Astrophysical Observatory. This publication makes
use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute ofTechnology. I was financially supported by Japan Society for the Promotion of Science and by the
Hayakawa Satio Foundation.
Finally, I would like to thank my family for supporting my graduate school life.

145