Study on Energetics and Particle Acceleration inSuperbubbles and Development of X-ray CCD Cameras
onboard the Suzaku Satellite
Hiroshi Nakajima
Department of Physics, Faculty of Science
Graduate School of Kyoto University
January 5, 2007

1

Abstract
This thesis has two distinct components. One concerns the ground- and inorbit calibration of the X-ray Imaging Spectrometer (XIS), the Charge CoupledDevice (CCD) onboard the Suzaku X-ray astronomical satellite. The other involves
the X-ray observations and analyses of the superbubbles (SBs), which have been
considered to be formed by the composite contributions from the strong stellar
winds from massive stars in the shell and successive supernova (SN) explosions,
utilizing Suzaku and other satellites.
Successive SN explosions in SBs may re-heat the cavity and its wall leaving
sometimes stellar remnants such as neutron stars and black holes inside. Both of theremnants may emit either diffuse or point like X-rays. Young massive stars in an OB
association are also known to be moderately bright X-ray sources. Therefore, X-ray
observations can reveal valuable information on the overall structure and evolutionof SBs. The SBs would also be the key cites of particle acceleration. However,no observational evidence for particle acceleration and/or limited information ofhigh temperature plasma from SBs have only been obtained. We found strong nonthermal X-ray emission from 30 Dor C, an SB in the Large Magellanic Cloud (LMC).
Assuming that this emission is due to the synchrotron X-rays and that the thermalpressure inside the SB is equal to that of the magnetic field, we showed that theelectrons are accelerated up to ∼90 TeV in the shell of 30 Dor C. This is the firstclear evidence of the long-term particle acceleration inside an SB. On the other hand,N44, another SB in the LMC does not show non-thermal X-ray emission. If we sum
up the thermal energies of hot gases including the breakout regions around the mainshell of N44 and kinetic energy, it reaches up to the expected injected energy by thestellar winds and SNe. The lack of the non-thermal X-ray emission is seen in othergalactic SB candidates Sgr C and CTB 37A, which makes 30 Dor C the conspicuous
object.
The instrumental portion of this thesis especially concerns the verification andperformance investigation of the Charge Injection (CI) capability of the XIS. The
existing X-ray CCDs working in orbit suffer from the radiation damage due tothe charged particles and subsequent degradation of the charge transfer inefficiency(CTI). Then we have tried to run the CI for the first time among the X-ray CCDs
in orbit. In principle, charge packets are artificially injected into each transfer
channel and are subsequently read after the transfer in the same manner as X-ray
events. This method allows us to measure and/or mitigate the amount of charge
loss on a column-to-column basis, which in turn, can potentially be a powerfultool for the calibration of the CTI. The column-to-column CTI correction for thecalibration source spectra significantly reduces the line widths compared to thosewith a column-averaged CTI correction (from 193 eV to 173 eV in FWHM on anaverage). In addition, this method significantly reduces the low energy tail in theline profile of the calibration source spectrum.

Contents
1 Introduction

8

2 Observatories
2.1 Suzaku . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.1.1 X-ray Telescope . . . . . . . . . . . . . . . . . . . . . . . . . .
2.1.2 X-ray Imaging Spectrometer . . . . . . . . . . . . . . . . . . .
2.1.3 Hard X-ray Detector . . . . . . . . . . . . . . . . . . . . . . .
2.2 XM M -N ewton . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2.1 X-ray Telescope . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2.2 European Photon Imaging Camera . . . . . . . . . . . . . . .
2.3 Chandra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.1 X-ray Telescope . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.2 Advanced CCD Imaging Spectrometer . . . . . . . . . . . . .

11
11
11
13
16
17
19
20
20
21
21

3 XIS
3.1 Basic Structure of the XIS . . . . . . . . . . . . . . . . . . . . . . . .
3.2 Specifications of X-ray CCDs . . . . . . . . . . . . . . . . . . . . . .
3.2.1 Transfer Mechanism . . . . . . . . . . . . . . . . . . . . . . .
3.2.2 Event Detection . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.3 Energy Resolution . . . . . . . . . . . . . . . . . . . . . . . .
3.2.4 Radiation Environment . . . . . . . . . . . . . . . . . . . . . .
3.2.5 Charge Injection Capability . . . . . . . . . . . . . . . . . . .
3.3 Charge Transfer Inefficiency . . . . . . . . . . . . . . . . . . . . . . .

23
23
25
26
26
27
28
28
32

4 Ground Calibration of XIS
4.1 Goal for Calibration Accuracy . . . . . . . . . . . . . . . . . . . . . .
4.2 Calibration System . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.3 Readout Noise and Split Threshold . . . . . . . . . . . . . . . . . . .
4.4 Gain and Energy resolution . . . . . . . . . . . . . . . . . . . . . . .
4.5 Quantum Efficiency . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.6 Performance of Charge Injection Capability . . . . . . . . . . . . . .
4.6.1 Stability of the Amount of Injected Charges . . . . . . . . . .
4.7 Energy Redistribution Function . . . . . . . . . . . . . . . . . . . . .

40
40
40
42
44
47
48
48
50

2

CONTENTS

3

5 Onboard Calibration of XIS
5.1 Performance Degradation in Orbit . . . . . . . . . . . . . . . . . . . .
5.2 CTI correction with CI Capability . . . . . . . . . . . . . . . . . . . .
5.2.1 CTI Measurement with Checker Flag CI . . . . . . . . . . . .
5.2.2 Compensating for the Charge Loss . . . . . . . . . . . . . . .
5.2.3 Energy Resolution Improvement . . . . . . . . . . . . . . . . .
5.3 Gain and Charge Transfer Efficiency . . . . . . . . . . . . . . . . . .

52
52
52
52
55
57
61

6 Review of Multiple Supernovae and Superbubble
6.1 Supernova Explosion . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.1.1 Shock Waves . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.1.2 Nuclear Synthesis . . . . . . . . . . . . . . . . . . . . . . . . .
6.1.3 Evolution of Supernova Remnants . . . . . . . . . . . . . . . .
6.2 Superbubble . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.2.1 Formation of Superbubble . . . . . . . . . . . . . . . . . . . .
6.2.2 Evolution of Superbubble . . . . . . . . . . . . . . . . . . . .
6.2.3 Cosmic Ray Acceleration in Superbubble . . . . . . . . . . . .
6.2.4 X-ray emission from SBs . . . . . . . . . . . . . . . . . . . . .
6.2.5 Past Observations . . . . . . . . . . . . . . . . . . . . . . . . .

64
64
64
66
67
71
71
72
74
77
79

7 Observations and Results
7.1 30 Dor C . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.2 N44 (DEM L 152) . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.3 Sgr C . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.4 CTB 37A . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

83
83
90
94
97

8 Discussion
103
8.1 Cosmic Ray Acceleration in SB . . . . . . . . . . . . . . . . . . . . . 103
8.2 Energetics of Superbubble . . . . . . . . . . . . . . . . . . . . . . . . 105
9 Conclusion
107
9.1 Conclusion for XIS calibration . . . . . . . . . . . . . . . . . . . . . . 107
9.2 Conclusion for SBs study . . . . . . . . . . . . . . . . . . . . . . . . . 107
A Abbreviations

118

List of Figures
1.1

The extinction cross section by interstellar medium. . . . . . . . . . .

9Schematic picture of the Suzaku satellite. . . . . . . . . . . . . . . .
Schematic view of the XRT onboard Suzaku. . . . . . . . . . . . . .
Walter type I optics. . . . . . . . . . . . . . . . . . . . . . . . . . . .
An outside view of the XIS sensor component. . . . . . . . . . . . . .
Schematic view of the back surface of the BI CCD chip treated withchemisorption process. . . . . . . . . . . . . . . . . . . . . . . . . . .
2.6 Quantum efficiency of the FI and BI CCD calculated from a CCD
model with designed thickness of each layer. . . . . . . . . . . . . . .
2.7 Schematic view of the CI structure and how to inject charges at the
input gate. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.8 Internal background spectra of X-ray satellites. . . . . . . . . . . . .
2.9 Schematic view of the HXD sensor component. . . . . . . . . . . . . .
2.10 Predicted sensitivity of Suzaku/HXD vs. IN T EGRAL calculated
on ground. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.11 Schematic view of the XM M -N ewton telescope. . . . . . . . . . . . .
2.12 EPIC FWHM as a function of energy. . . . . . . . . . . . . . . . . . .
2.13 Schematic view of the Chandra satellite. . . . . . . . . . . . . . . . .
2.14 Schematic view of the ACIS focal plane. . . . . . . . . . . . . . . . .

12
12
13
14

2.1
2.2
2.3
2.4
2.5

3.1
3.2
3.3
3.4
3.5

Schematic view of the XIS sensor component. . . . . . . . . . . . . .
The CCD and heat sink assembly of the BI CCD flight model. . . . .
Schematic view of the XIS CCD chip. . . . . . . . . . . . . . . . . . .
Grade definition and the rule to add PHA. . . . . . . . . . . . . . . .
The predicted proton spectrum inside a simple model of the XIS
shielding. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.6 The specific energy loss for protons in silicon. . . . . . . . . . . . . .
3.7 Schematic view of how to measure δQ using the test and reference
charges. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.8 The PHA of injected charges as a function of the Y coordinate. . . .
3.9 Frame image of the XIS during the CI experiment. . . . . . . . . . .
3.10 The detrapping times for traps of various energies as a function oftemperature. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

4

15
15
16
17
18
18
19
21
22
22
24
24
25
27
28
29
31
31
32
33

LIST OF FIGURES

5

3.11 Schematic view of the case where only the preceding charge packet
lose electrons to a given trap. . . . . . . . . . . . . . . . . . . . . . . 34
3.12 Energy spectra of 55 Fe events of XIS2 collected in orbit. . . . . . . . . 35
3.13 The line center of the beneficial CPs in figure 3.12 subtracted by thatof the sacrificial CPs. . . . . . . . . . . . . . . . . . . . . . . . . . . . 39
A spectrum of the XIS FI sensor of the 55 Fe source collected on ground. 41
Schematic view and photograph of the X-ray generation system at
Kyoto University. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 41
4.3 Frame image of the proton-damaged FI sensor used in the groundexperiment. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 42
4.4 Spectra of 55 Fe extracted from the proton undamaged and damaged
regions. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 43
4.5 Percentage of the Grade 0 and 7 events as a function of the energy ofincident X-rays. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 44
4.6 Energy resolution (FWHM) and detected number of events as a function of the split threshold. . . . . . . . . . . . . . . . . . . . . . . . . 45
4.7 The optimized split threshold levels as a function of the energy of the
X-rays. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 46
4.8 The XIS energy scale linearity of the XIS. . . . . . . . . . . . . . . . 46
4.9 The spectral resolution as a function of energy of the XIS. . . . . . . 47
4.10 The QE as a function of energy for XIS FI. . . . . . . . . . . . . . . . 48
4.11 Stability of the amount of the injected charges. . . . . . . . . . . . . 49
4.12 Ratio of Q0test and Q0ref for the CI events and that of Q0upper and Q0lower
for the 55 Fe events. . . . . . . . . . . . . . . . . . . . . . . . . . . . . 50
4.13 The spectra of 55 Fe collected the response function we obtained withground experiments. . . . . . . . . . . . . . . . . . . . . . . . . . . . 51
4.14 Fit results of 55 Fe spectrum with the response function around theKα and Kβ line. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51
4.1
4.2

5.1
5.2
5.3
5.4
5.5
5.6
5.7
5.8

The line center energy of the each calibration source as a function oftime without any CTI correction. . . . . . . . . . . . . . . . . . . . .
The line widths of the each calibration source as a function of time
without any CTI correction. . . . . . . . . . . . . . . . . . . . . . . .
PHA distribution of the test and reference events as a function of the
X coordinate. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Distribution of the index in the relation between δQ and Q. . . . . .
The Y coordinate dependence of the center energy of 6.4 keV line
emission from the Sgr C region. . . . . . . . . . . . . . . . . . . . . .
Energy spectra of the calibration source after charge compensation
with CI capability and with the calibration source itself. . . . . . . .
FWHM of the calibration source spectra simultaneously obtainedwith the CI experiments. . . . . . . . . . . . . . . . . . . . . . . . . .
Spectra around the He-like Si Kα emission line of the west part ofTycho’s SNR. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

53
54
55
56
57
58
58
59

LIST OF FIGURES
FWHM of the calibration source spectra simultaneously obtainedwith the CI experiment. . . . . . . . . . . . . . . . . . . . . . . . . .
5.10 FWHM of the CI test events spectra before and after charge compensation, and that of the reference events. . . . . . . . . . . . . . . . . .
5.11 The line center energy of the each calibration source as a function oftime after CTI correction. . . . . . . . . . . . . . . . . . . . . . . . .
5.12 The line widths of the each calibration source as a function of timeafter CTI correction. . . . . . . . . . . . . . . . . . . . . . . . . . . .

6

5.9

6.1
6.2
6.3
6.4
6.5
7.1
7.2
7.3

A schematic view around the shock front. . . . . . . . . . . . . . . . .
The predicted relative abundances of synthesized heavy elements. . .
Schematic view of the evolution of SNRs. . . . . . . . . . . . . . . . .
The CR spectrum for all nuclei, as measured with various facilities,
and the expected behavior of the proton and iron spectra. . . . . . .
SB X-ray luminosity vs. bright star count for the best-fit NH X-ray
spectral fits and fixed NH X-ray spectral fits. . . . . . . . . . . . . . .

60
60
62
63
64
67
70
75
81

Chandra ACIS images around 30 Dor C. . . . . . . . . . . . . . . . . 85
Background-subtracted spectra for shells A - D of 30 Dor C. . . . . . 87
Hα image of N44, showing the nebular components and the OB associations. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 91
7.4 X-ray image of the N44 region. . . . . . . . . . . . . . . . . . . . . . 92
7.5 Background-subtracted spectrum extracted from entire main shell region. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 93
7.6 X-ray spectra extracted from southeast and northwest part of the N44
main shell region. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 93
7.7 Background-subtracted spectrum extracted from vicinity of OB-typestars association LH 47. . . . . . . . . . . . . . . . . . . . . . . . . . 95
7.8 Radio continuum contour map around Sgr C region. . . . . . . . . . . 96
7.9 X-ray images of XIS (6.4 keV, 2.45 keV, 6.7 keV) of Sgr C region. . . 100
7.10 Background-subtracted spectrum extracted from G359.4–0.1 and “blowout”
region. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 101
7.11 Radio continuum contour map around CTB 37A region. . . . . . . . 101
7.12 Chandra three-color X-ray image of CTB 37A region. . . . . . . . . . 102
7.13 Background-subtracted spectrum extracted from G348.5–0.0 and knot
region. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 102

List of Tables
2.1
2.2

The performance description of XIS front-illuminated CCD chips. . . 14
The on-axis performance description of XRT onboard XM M -N ewton
at 1.5keV. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 20

3.1

The notation list of parameters for the CI. . . . . . . . . . . . . . . . 30

5.1

The log of the CI experiment on orbit. . . . . . . . . . . . . . . . . . 55

6.1

Best-Fit LMC SB X-Ray Spectral Fits. . . . . . . . . . . . . . . . . . 80

7.1
7.2
7.3
7.4
7.5
7.6
7.7
7.8

Observation Log around 30 Dor C . . . . . . . . . . . . . . . . . . . .
Point Source Data around 30 Dor C . . . . . . . . . . . . . . . . . . .
Best-Fit Parameters of Shell A . . . . . . . . . . . . . . . . . . . . . .
Best-Fit Parameters of Shell B - D . . . . . . . . . . . . . . . . . . .
Fitting parameters for each region. . . . . . . . . . . . . . . . . . . .
Best-Fit Parameters of G359.4–0.1 and “blowout” region . . . . . . .
Observation Log around CTB 37A . . . . . . . . . . . . . . . . . . .
Best-Fit Parameters of G348.5–0.0 and knot . . . . . . . . . . . . . .

7

84
86
88
89
94
97
98
99

Chapter 1
Introduction
Since the daybreak of the astronomy, its progress has been facilitated by observations
with novel instruments at each time. Especially with respect to the observationsof the photons from celestial objects, many astronomers have taken their efforts toobtain the precise and accurate values of the wavelength (i.e. energy), flux intensity,
direction, and polarity with wide range. The sophisticated spectra, light curves, andimages enable one to reveal the consciousness and/or the emission mechanisms of
the objects.
Extending the effective energy range to the X-ray band, we have been able to see
and investigate energetic phenomena hidden behind dense absorption columns suchas shocked thermal plasma in supernova remnants (SNRs), accretion disks aroundblack holes, hot intracluster gases in clusters of galaxies, and so on. In fact, in thecase of the Galactic Center (GC), the absorption column density is 6 × 1022 cm−2 ,which is consistent with the suggested distance of 8.0 kpc to the GC (McNamara
et al. 2000) and the average number density of the hydrogen atoms along the galactic
plane of 1 cm−3 .
The first imaging in the hard X-ray band (2 - 10 keV) was performed by ASCA
satellite. Since then, Charge Coupled Device (CCD) has been the primary instrument because of its high spatial resolution, moderate energy resolution and temporal
resolution. Through the successful operations and results by the Chandra/ACIS andXM M -N ewton/EPIC, the X-ray CCD has also been the most reliable and stabledetector in the X-ray astronomy.
At the time of writing this thesis, the X-ray Imaging Spectrometer (XIS) onboard
Suzaku, the fifth Japanese X-ray astronomical satellite, is the brand-new, superior
X-ray CCDs in orbit. Its low and stable particle background and good energy
resolution are the complementary characteristics against those of the other X-ray
CCDs in orbit: high spatial resolution of the Chandra/ACIS, large effective area ofthe XM M -N ewton/EPIC.
We, the XIS team, have made experiments to improve the above performanceand make correct energy redistribution function, which defines a mapping from thesource photon energy to the detected pulse height amplitude (PHA).
Another advantage of the XIS is that it equips a register for the injection of
artificial charge packets. Injection and subsequent readout of the charge packets
8CHAPTER 1. INTRODUCTION

9

Figure 1.1: The extinction cross section by interstellar medium from far-infrared
to the X-rays. Both of the absorption and scattering are included. This figure isadopted from Ryter (1996).
enable us to measure charge losses and/or mitigate charge traps in a transfer channel.
Then, compensating the lost charges for each event leads to the precise estimationof the PHA, that is, the better gain and energy resolution. Because the degradationof the gain and energy resolution of X-ray CCDs in orbit is inevitable due to theradiation damage by charged particle, the charge injection (CI) capability can be
the most powerful tool to have the advantage of other satellites.
Hot ionized gas is an essential component of the interstellar medium (ISM) andsupernovae (SNe) and their remnants as well as the stellar winds from massive
stars are the main site for the formation of the hot gas with a temperature fromseveral 106 K to over 107 K. They interact with ISM through ionizing, accelerating,and heating up. The heavy elements are spread out by the SN explosions into
the interstellar space, thus the X-ray spectra of young SNRs strongly reflect theabundance of ejected matter. As the SNRs get older, the amount of the ambient
matter gathered up by their shock fronts. Hence the middle-aged SNRs show the Xray spectra which characterize the abundance of the ambient ISM. More older SNRs
are then considered to gradually lose their identities by merging into the ambient
ISM. Thus, the ISM can be considered as the wreck of SNRs.
However, in most cases, SNRs occur spatially and temporally concentrated andnear the massive stars. The enormous mechanical energy released by these stars
through stellar winds and SNe creates a large pressure in the surrounding medium.
The dense neutral ISM is displaced by low-density coronal gas, producing a super-CHAPTER 1. INTRODUCTION

10

bubble (SB) surrounded by a dense ringlike supershell. Through the SB evolution,
the front shocks of each SNR and strong stellar winds interact with each other producing the turbulence of plasma with a hundred to thousand parsec scale. The
association of massive stars and successive SNRs potentially distribute around thegalactic center, spiral arms, and other dense molecular clouds. Although detailed
calculations of the evolution of SBs may differ depending on various assumptions
about energy deposition and the interstellar environment, this qualitative picture
generally holds. Because of its enormous energy and mass, SBs contribute largepart of the energy input and chemical evolution of a galaxy. Therefore, we can trace
the course of the formation of hot ISM in the galaxy by systematically investigating
SBs and the site of SB formation as well as the hot ISM. Furthermore, the multiple
front shocks in the hot and tenuous gas and their interaction inside SB lead theacceleration of cosmic rays (CR). Compared with the isolated SNRs, the durationof the acceleration is long and the particle injection rate to the acceleration site ismore efficient, the SBs are one of the most possible candidates of the most powerfulacceleration site of Galactic CRs.
Scope of this thesis is:
1. Bring out the superior performances of the XIS through ground and in-orbit
calibrations.
2. Improve the in-orbit performance, especially the gain and energy resolution
with the execution of the CI capability for the first time among the X-ray
CCDs used for satellite.
3. Study the energetics and particle acceleration in the SBs by investigating somesamples of SBs and multiple SNRs sites.
We will introduce the observatories which are used in this thesis including Suzaku
in chapter 2 at first. The hardware portion of this thesis consists of chapter 3 which
describes the detailed specification of XIS. The following chapter 4 and 5 are dedicated to explain the results of the ground- and in-orbit calibrations, respectively.
The other part concerning SB consists of the review of SNR and SB formation andtheir evolution in chapter 6, our observation results in chapter 7, and the discussionin chapter 8. Finally the conclusion is shown in chapter 9.

Chapter 2
Observatories
2.1

SuzakuSuzaku, the fifth Japanese X-ray astronomical satellite, was put into orbit on 10th
July, 2005 as the recovery mission of Astro-E, the launch of which was unsuccessful in 2000. It was developed at the Institute of Space and Astronautical Science of Japan Aerospace Exploration Agency (ISAS/JAXA) in collaboration with
NASA Goddard Space Flight Center (GSFC), Massachusetts Institute of Technology(MIT), and other Japanese institutions. The spacecraft is in a near-circular orbit
with an apogee of 568 km, inclination angle of 31.9◦ , and the orbital period of approximately 96 minutes. Its mass and length are about 1300 kg and 6.5 m along thetelescope axis after the deployment of the extensible optical bench (EOB). Schematic
views are shown in figure 2.1.
Suzaku is equipped with three scientific instruments: the Hard X-ray detector(HXD), the X-ray Imaging spectrometer (XIS), and the X-ray spectrometer (XRS).
Although the basic specifications of these instruments are the same as those of
Astro-E, some improvements have been performed. In spite of the loss of the XRS,
the main instrument of Suzaku, just before its first light, Suzaku still has the wideenergy range of 0.2 - 600 keV, high energy resolution at low energy band, and an
extremely low background level compared with other major X-ray satellites.

2.1.1

X-ray Telescope

Suzaku has five light-weight thin-foil X-Ray Telescopes (XRTs). One of them isdedicated for the XRS and is called XRT-S, although it is no more functional. The
others are equipped for the XIS and are called XRT-I. They have been developed
jointly by NASA/GSFC, Nagoya University, Tokyo Metropolitan University, andISAS/JAXA. These are Walter type I grazing-incidence reflective optics consistingof compactly nested, thin conical elements as shown in figure 2.3. Because of thereflectors’ small thickness, they permit high density nesting and thus provide large
collecting efficiency with a moderate imaging capability in the energy range of 0.2 12 keV, all accomplished in telescope units under 20 kg each.
The XRTs are arranged on the EOB on the spacecraft in the manner shown in
11CHAPTER 2. OBSERVATORIES

12

Figure 2.1: (Left panel:) Schematic picture of the bottom of the Suzaku satellite.
(Middle panel:) Side view of the instruments and telescopes on Suzaku. (Rightpanel:) Top view of the instruments on Suzaku.

Figure 2.2: Schematic view of the XRTs mounted on the top plate of the EOB. The
courtesy of K. Abe, NIPPI corporation.
figure 2.2. The external dimensions of the 4 XRT-Is, however, are the same. The
angular resolutions of the XRTs range from 1.80 to 2.30 , expressed in terms of halfpower diameter, which is the diameter within which half of the focused X-ray is
enclosed. The angular resolution does not significantly depend on the energy of theincident X-ray in the energy range of Suzaku, 0.2 - 12 keV. The effective areas are
typically 440 cm2 at 1.5 keV and 250 cm2 at 8 keV. The focal lengths are 4.75 m for
the XRT-I. Individual XRT quadrants have their component focal lengths deviatedfrom the design values by a few cm. The optical axes of the quadrants of each XRT
are aligned within 20 from the mechanical axis. The field of view (FOV) for XRT-Is
is about 170 at 1.5 keV and 130 at 8 keV.

13

CHAPTER 2. OBSERVATORIES

parabolahyperbola

focal plane
Figure 2.3: Walter type I optics. This is accomplished by two successive grazingincidence reflections on a parabolic and a hyperbolic shaped metal surface.

2.1.2

X-ray Imaging Spectrometer

The XIS consists of the sensor, analog electronics (AE), and digital electronics (DE).
The picture of the sensor is shown in figure 2.4. The readout analog signals from thesensor are transferred to the AE and converted into the digital signals. The digitized
pixel data are transferred into the DE and then each X-ray event is extracted. The
X-ray event information is formatted into telemeter data and transferred to the dataprocessor.
All sensors of XIS were developed through the collaboration of the Center forSpace Research (CSR) at MIT, ISAS/JAXA, Osaka Univ., Kyoto Univ., Rikkyo
Univ., Kogakuin Univ., and Ehime Univ. The AE and CCD chips were developed
at CSR and the Lincoln Laboratory at MIT, respectively. The upper part of thesensors (hereafter referred to as bonnet) and DE were developed by ISAS/JAXA,
Osaka Univ., and Kyoto Univ.. The CCD calibrations have been done by MIT,
Osaka Univ. and Kyoto Univ. Gain, energy resolution, and quantum efficiency(QE) have been investigated. After simple calibrations at MIT, the Osaka Univ.
group has investigated the low energy (≤Si K-edge ' 1.8 keV) X-ray response of
the sensors. At Kyoto Univ., the high energy (≥Al Kα line ' 1.5 keV) response has
been intensively investigated.
Each sensor consists of an X-ray CCD camera combined with the single XRT.
Three sets of XIS utilize front-illuminated (FI) chips, while the other is equippedwith a back-illuminated (BI) chip (MIT Lincoln Laboratory model CCID41). All
sensors including spare sensors utilize the three phase CCD, and the FI CCDs havebasically the same chip structure as the ACIS-I chips onboard Chandra (CCID17;
Bautz et al. 1996). Table 2.1 summarizes the specification of the XIS CCD chip.
There are several improvements against the XIS onboard Astro-E (Tsuru et al.
2000). First, the BI chip is used. The BI chip has the same basic specifications as
the FI chips, except that it has a larger QE in the soft energy band as shown in figure

14

CHAPTER 2. OBSERVATORIES

Figure 2.4: An outside view of the XIS sensor component.
Table 2.1: The performance description of XIS front-illuminated CCD chips.
Pixel Size

24µm×24µm

Pixel Number

1024×1024

Chip Size

25mm×25mm

Time Resolution

8 msec (Timing Mode), 8 sec (Normal Mode)

2.6. Furthermore, the BI CCD for XIS is treated with the chemisorption processdeveloped by Lesser et al. (1998). Coating the thin (∼ 1 nm) silver on the backside
surface in a fabrication process, the silver catalyzes dissociation of the molecularoxygens leaving fixed, negatively charged oxygen atoms on the surface as shown infigure 2.5. Because of the potential due to these oxygen ions, the photoelectrons
created around the backside edge of the depletion layer can be successfully drifted to
the front surface. Hence, the BI CCD of XIS has almost the same energy resolution
as FI CCDs in the soft energy band, while other BI CCDs such as ACIS-S onboard
Chandra and pn-CCD onboard XM M -N ewton have energy resolution worse thanFI devices roughly by a factor of two.
Second, a charge injection (CI) capability is added (Prigozhin et al. 2004; Bautzet al. 2004; LaMarr et al. 2004). Injection of charge packets (CPs) into a transfer
channel of the CCD and a subsequent readout process allow the charge transfer
inefficiency (CTI) in each transfer channel to be measured. I will describe theperformance of the CI capability in chapter 5 in detail. Here I describe the basic
specification of the CI capability.
In order to maintain the performance of CCDs in orbit, the CTI must be frequently measured and corrected. Furthermore, the ASCA/SIS raised an important
issue: The CTI varies significantly from column to column. Therefore some capability is needed such that the CTI can be measured for each column. Then, theCI structure was introduced to the Astro-E2/XIS (figure 2.7). The structure of CI

CHAPTER 2. OBSERVATORIES

15

Figure 2.5: Schematic view of the back surface of the BI CCD chip treated withchemisorption process. This figure is adopted from Burke et al. (2004).

Figure 2.6: Quantum efficiency of the FI and BI CCD calculated from a CCD modelwith designed thickness of the electrode layer, insulation layer, and depletion layer.

CHAPTER 2. OBSERVATORIES

16

Figure 2.7: (Left panel:) Schematic view of the CI structure. The artificial charges
are injected to the serial register from the input gate located at the edge of theregister. After positioning the charges over the register, they are injected to theimaging area simultaneously by the vertical clock. (Right panel:) Schematic viewof how to inject charges at the input gate. The amount of injected charges canbe controlled by adjusting the offset voltage between the input gate and the S3electrode.
register and how to inject charges are shown in figure 2.7. A serial register (hereafter
CI register) equipped at the edge of the imaging area (IA) and an artificial charge
packet can be injected from the input gate equipped at the edge of the CI register.
Pulling down the potential for electrons at the input gate and the next electrode ( S3in figure 2.7), “fill” the bucket formed at the electrode. The amount of the injectedcharges can be controlled by the offset voltage between the input gate and the S3electrode. Then “spill” the charges pulling up the potential. We repeat this “fill andspill” cycle with 1/40960 sec ' 24 µsec in the normal manner. After positioning thepackets in the register with a desired pattern, they are simultaneously transferred
into the IA by vertical clock.
Another characteristics of XIS is the low background level. While Chandra and
XM M -N ewton suffer from the charged particle background, primarily low energy
protons (∼ 10 - 100 keV) in the van Allen belts (Plucinsky et al. 2000; Read &Ponman 2003), that of Suzaku is low and stable compared to other satellites dueto the low earth orbit. Figure 2.8 shows the comparison of the non X-ray backgroundbetween other satellites (ASCA, Chandra, and XM M -N ewton).

2.1.3

Hard X-ray Detector

One of the main advantage of Suzaku is a spectroscopy in the wide energy band of
0.2 - 600 keV. The HXD takes charge of high energy band of 12 - 600 keV. The HXD
consists of two detectors: 2 mm thick silicon PIN diodes sensitive over 10 - 70 keV,
and GSO crystal scintillators placed behind the PIN diodes covering 40 - 600 keV.
Figure 2.9 shows a schematic view of the HXD sensor. The HXD is composed of 16

CHAPTER 2. OBSERVATORIES

17

Figure 2.8:
Internal background spectra of ASCA/SIS, Suzaku/XIS-FI,
Suzaku/XIS-BI, XM M -N ewton/MOS, XM M -N ewton/pn. Each spectrum is normalized by the CCD area.
(4×4) identical detector units. The four sides of the detector units are surroundedby 20 thick anti-counters made of BGO scintillators for background shield. The
anti-counters are also expected to detect transient high energy phenomena such as
the gamma-ray bursts.
Inside of each well, passive fine collimators made of phosphor bronze sheet areplaced to limit the FOV for soft X-rays (< 100 keV), so that the FOV becomescomparable to those of the XRS and XIS. The fine collimators are also expected toreduce the cosmic diffuse X-ray background, which otherwise might be a significant
background source for the PIN diodes. With this collimator, the HXD achieves
the lowest background level in this energy band ever as shown in figure 2.10. The
detailed description of the detector and the results of the flight calibration of the
HXD can be seen in Takahashi et al. (2006) and Kokubun et al. (2006).

2.2

XM M -N ewton

The European Space Agency’s X − rayM ultiM irrorM ission − N ewton(XM M −N ewton) was launched by Ariane 504 on December 10, 1999. XM M -N ewton carries
three XRT with the largest effective area of a focusing telescope ever: The total
mirror geometric effective area at 1.5 keV energy is 1550 cm2 for each telescope, i.e.,
4650 cm2 in total (Jansen et al. 2001 for details). The main instrument, European
Photon Imaging Camera (EPIC), is the CCD cameras, which have moderate spectral

CHAPTER 2. OBSERVATORIES

18

Figure 2.9: Schematic view of the HXD sensor component. The green and blue partis the PIN diodes and BGO crystal, respectively.

Figure 2.10: Predicted sensitivity of Suzaku/HXD vs.
IBIS/ISGLI and
IBIS/PlCsIT onboard IN T EGRAL calculated on ground. 100 ks observation andthe energy resolution of ∆E = 0.5E are assumed.

CHAPTER 2. OBSERVATORIES

19

Figure 2.11: Schematic view of the XM M -N ewton telescope. X-rays irradiate fromthe bottom-left XRT to the upper-right instruments.
resolution (with a resolving power, E/∆E, of ∼ 20 - 50). The Reflection Grating
Spectrometers offer much higher spectral resolution, with a resolving power in therange of 200 - 800. A highly elliptical orbit offers continuous target visibility of up
to about 40 hours, with a minimum height for science observations of 46,000 km.
This is very favourable for studies of source variability and also in order to achievea high overall observatory efficiency.

2.2.1

X-ray Telescope

XM M -N ewton’s high sensitivity is achieved by using 58 thin nested mirror shells
in each XRT. The XRT onboard XM M -N ewton consists of three Wolter-I type
X-ray telescopes that are co-aligned with a relative astrometry between the three
EPIC cameras with the focal length of 7.5 m. They are calibrated to better than 1- 200 across the full FOV. The system is illustrated in figure 2.11. Two of them havegrating assemblies in their light paths, diffracting a part of the incoming radiation
onto their secondary focus. Values for the full width at half maximum (FWHM)
and half energy width (HEW) of the PSFs and effective area of each XRT are listed
in table 2.2.

20

CHAPTER 2. OBSERVATORIES

Table 2.2: The on-axis performance description of XRT onboard XM M -N ewton at
1.5keV.

2.2.2

Instr. chainpn

FWHM (00 )MOS-1+RGS-1 MOS-2+RGS-2

<12.5

4.3

4.4

00

HEW ( )

15.2

13.8

13.0

Eff. Area (cm2 )

1380

520

520

European Photon Imaging Camera

Two of the XM M -N ewton XRT are equipped with EPIC MOS (Metal Oxide Semiconductor) CCD arrays, the third carries a different CCD camera called EPIC pn.
The EPIC cameras offer the possibility to perform sensitive imaging observations
over an FOV of 300 and an energy range from 0.15 to 15 keV, with moderate spectral
resolution ( E/∆E ∼ 20 − 50).
Each of the MOS CCD arrays consists of seven FI CCDs. Each chip is a threephase frame transfer device and has an imaging area of 600 pixels × 600 pixels withthe pixel size of 40 µm square, so that the FOV of a mosaic of seven chips is 280 .4.
The mean depletion of the MOS CCDs is between 35 and 40 µm. Each camera isequipped with the calibration source of 55 Fe with Al-targets for the calibration of
the absolute gain and it provides the Al Kα, Mn Kα, and Mn Kβ line. Holland et al.
(1996) and Sembay et al. (2004) describe the specification and in-orbit performanceof the MOS detectors in detail.
The heart of the EPIC pn camera is an array of 12 monolithically integrated
pn-CCDs. The pixel size is 150 µm square, arranged in 200 rows by 64 channels
forming a sensitive area of 6 cm × 6 cm. A 300 µm silicon wafer is fully depletedwith a reverse biased pn-diode and reverse biased strip diodes at the opposite wafer
surface. The large area diode acts as a homogeneous radiation entrance window.
The pn-strip diodes (shift registers) are used for electron storage and transfer tothe anode, similar to the MOS gates of the other CCDs. The detector is kept at atemperature of –90 ◦ C.
The in-orbit FWHM of EPIC cameras of the Al Kα and Mn Kα on-board calibration lines is plotted in figure 2.12. The CTI correction is already applied. Thisshows a degradation of approximately 13% in the energy resolution of the MOS since
the launch. In this point the EPIC MOS is a formidable competitor for the XIS.

2.3

Chandra

The Chandra X-ray Observatory was deployed by the Space Shuttle Columbia on
the 23 of July 1999. The Chandra spacecraft carries a high resolution mirror, two
imaging detectors, and two sets of transmission gratings. Important Chandra features are spatial resolution of < 100 and good sensitivity from 0.1 to 10 keV for point

CHAPTER 2. OBSERVATORIES

21

Figure 2.12: (Left panel:) EPIC MOS FWHM as a function of energy. The solid
curve is a best-fit E0.5 function to ground calibration data between 0.1 - 12.0 keV.
(Right panel:) EPIC pn FWHM as a function of energy as present in the response
matrices. Curves are given for single- and double-pixel events at the focus position
as well as at a position 10 pixels away from the readout node.
sources. Chandra is in an elliptical high-earth orbit allowing observing intervalsof more than 48 hours in length. The details for the satellite and instruments aredescribed by Weisskopf et al. (2002).

2.3.1

X-ray Telescope

The mirror consists of four pairs of nested reflecting surfaces, arranged in the usual
Wolter-I type geometry. The high energy response is achieved by use of relatively
small reflection angles and by coating the mirrors with iridium. Improvements inmirror technology since Einstein include advances in grinding, polishing, alignment,
and testing. Mirrors with resolution of 0.500 have been achieved.

2.3.2

Advanced CCD Imaging Spectrometer

The Advanced CCD Imaging Spectrometer (ACIS) offers the capability to simultaneously acquire high-resolution images and moderate resolution spectra. The instrument also is used in conjunction with the High Energy Transmission Grating(HETG) or Low Energy Transmission Grating (LETG) to obtain higher resolution
spectra. ACIS contains 10 planar, 1024 pixels × 1024 pixels CCDs; four arranged in
a 2×2 array (ACIS-I) used for imaging, and six arranged in a 1×6 array (ACIS-S)
used either for imaging or as a grating readout. Any combination of up to 6 CCDs
may be operated simultaneously. Two CCDs are BI chips and eight are FI chips.
The response of the BI devices extends to energies below that accessible to the FI
chips. The chip-averaged energy resolution of the BI devices is better than that of
the FI devices in a high energy band.

22

CHAPTER 2. OBSERVATORIES

Figure 2.13: Schematic view of the Chandra spacecraft with an exploded view ofthe science instruments.

ACIS FLIGHT FOCAL PLANE
not constant with Z~22 pixels ~11"22 pixels
= 11"

I0

I1

w203c4r
0

w193c2
1
x

I2

}

I3

w158c4r

w215c2r

2

3

ACIS-I(aimpoint on I3 = (949, 978))330 pixels = 163"
S0

S1

S2

w168c4r
4

w140c4r
5

w182c4r
6

S3+
w134c4r
7

S4S5

w457c4w201c3r
98}

ACIS-S(aimpoint on S3 = (235, 497))+∆Y
18 pixels = 8".8

Image Region

Bottom.

column

three

two

one

Top

BI chip indicator

node zero

+Z

.

row

TargetOffsetCoordinates+Z
Pointing+∆Z

Coordinates-Z

Pixel (0,0)+Y

Sim Motion

Frame StoreCCD Key

NodeDefinitions

Row/Column
Definition

Coordinate
Orientations

Figure 2.14: Schematic view of the ACIS focal plane. The nominal aimpoints for“imaging” and “grating” mode are marked with cross on S3 and I3 chip, respectively.

Chapter 3
XIS
Here I describe the detail of the XIS and its components.

3.1

Basic Structure of the XISXIS camera body consists of a hood, bonnet, and base as shown in figure 3.1. The
hood has five baffles to block stray light. The bonnet includes a vacuum valve,an optical blocking filter (OBF), calibration sources, a pressure sensor, a paraffin
actuator, and a door. The base contains an X-ray CCD mounted on an alumina(Al2 O3 ) substrate and is attached to a heat sink assembly made of copper (Cu).
Over the frame-store region (FS) of the CCD, an aluminum (Al) shield with a gold(Au)-plated nickel (Ni) surface treatment is placed to block X-ray irradiation. The
surfaces of the substrate, the heat sink assembly, and the shield are also plated with
Au. The bonnet and the hood are made of Al with black surface finish. The inside of
the base is Ni-plated, except for the feed-through plate which is Au electro-plated.
The XIS sensor is covered by multi-layer insulators to reduce heat-flow from thesatellite body.
The XIS CCD are sensitive to optical and ultraviolet light as well as X-rays,and hence have an OBF 20 mm above the IA. The OBF is made of a polyimide(C22 H10 N2 O4 ) film with vapor-deposited Al on both sides, in order to avoid lightleakage from pinholes. The thicknesses of the Al on the two sides of the filter are
different from each other. This reduces the light transmission by the interference
effect (Born & Wolf 1999). The thickness of the polyimide film and Al layers
are ∼ 1400Å and ∼ 800Å + 400Å, respectively. This realizes a low transmission
coefficient for optical light (≤ 5 × 10−5 ), but is transparent to X-rays (≥ 80% above
0.7 keV). The measured transmissions for optical light and soft X-rays are reported
in Kitamoto et al. (2004). Since the OBF is very fragile and vibrations during launch
may destroy it; at the time of the launch, the vicinity of the chip and the OBF was
kept in an air-tight vacuum by closing a door equipped with the bonnet. The door
was opened by the paraffin actuator after 18 days from the launch.
The X-ray CCDs generally suffer from dark current due to the thermal excitationof electrons and the small energy band gap of Si semiconductor crystals (1.12 eV).
Then, a three-stage thermo-electric cooler (TEC; Peltier sensor) is used to cool the23

24

CHAPTER 3. XIS

Baﬄe

Hood

230

Electric
Valve

Door

Door Calibration
Source
CCD

Wall Calibration
Source70

Optical Blocking
Filter (OBF)
Cover Shield

Thermo-Electric
Cooler (TEC)

Heat Sink

Flexible Print

Figure 3.1: Schematic view of the XIS sensor component. The single-hatched andcross-hatched regions are bonnet and base, respectively.

Figure 3.2: The CCD and heat sink assembly of the BI CCD flight model. The
cover shield and base are removed in this picture. The bottom stage of the TEC
can be seen under the CCD.
CCD to the nominal operating temperature of −90◦ C. The cold-end of the TEC isdirectly connected to the substrate of the CCD, which is mechanically supportedby three Torlon posts (polyamide-imide plastic) attached to the head sink. The
heat is transferred through a heat pipe to a radiator panel on the satellite surface,
and is radiated away to space. The radiator and heat pipe are designed to cool thebase below −40◦ C under the nominal TEC operating conditions. Figure 3.2 is aphotograph of the BI CCD flight model on the heat sink assembly with the framestore cover shield removed. The flex cable put at the bottom of the chip is seen at
the left side. The bottom stage of the TEC can be seen under the FS region of the
CCD.
Each XIS sensor has three 55 Fe calibration sources. One is attached to the door,
and illuminates the whole IA almost uniformly. This calibration source was usedon the ground and for initial in-orbit calibration before the door was opened. For

CHAPTER 3. XIS

25

Figure 3.3: Schematic view of the CCD chip. A serial register is equipped at the
top of the imaging area for CI capability. The chip has four read out nodes. Signals
are readout simultaneously from these nodes.
normal observations, after the door is open, this source is out of the FOV of the
CCD. The other two calibration sources are located on a side wall of the bonnet
and illuminate the two far-end corners of the IA from the readout node. The fluxesof the these calibration sources were selected so that the gain can be determined at
5.9 keV with a statistical error of 0.1% within a single orbit early in the mission.

3.2

Specifications of X-ray CCDs

In order to see the function of the CCD, we give the schematic view of the XIS
FI chip in figure 3.3. Each CCD chip has four segments (from A to D), and eachsegment has one readout node. A simplified model of the dead layers in the FI
devices, in which the gate structure is approximated as a stack of nearly uniform
slabs, has been fit to the ground calibration data. Then the dead layer thicknessesof ∼0.28 µm (Si) and ∼0.44 µm (SiO2 ) are inferred. The surface dead layers of the
BI CCD is very thin, consisting of 5 nm of HfO2 , 1 nm Ag, and 3 nm SiO2 . As forthe depletion layer thickness, the ground calibrations verified that the thickness ofthe depletion layer is ∼ 65 µm for the FI chips, and ∼ 42 µm for the BI chip. This
leads the QEs for the FI and BI CCDs as a function of energy as shown in figure
2.6.
Position in the IA is defined by the detector-fixed coordinates (ActX, ActY),
where the origin (0, 0) is taken to be the first pixel readout of segment A in normal

CHAPTER 3. XIS

26clocking mode. The positive directions are shown by arrows in figure 3.3. In the IA,
ActX runs from 0 to 1023 (from segment A to D), while ActY runs from 0 to 1023
(from the readout node to the CI register).
In this thesis, we converge to the description and experiment results of the normal clocking mode and 5×5/3×3 edit mode because most of the observations isconducted with the above configurations.

3.2.1

Transfer Mechanism

In the normal clocking mode of frame transfer type CCDs, after the exposure is over,the charge in the IA is transferred at a high speed into the X-ray shielded FS region(FS transfer). Once the FS transfer is completed, the IA begins its next exposure
while the FS is slowly readout. The readout is accomplished by transferring the
FS region row by row into the serial register (parallel transfer), where each row isclocked out serially into an output node (serial transfer).

3.2.2

Event Detection

When a CCD pixel absorbs an X-ray photon, the X-ray is converted to an electriccharge, which in turn produces a voltage at the analog output of the CCD. This
voltage (“pulse-height amplitude”; PHA) is proportional to the energy of the incident X-ray. In order to determine the true PHA corresponding to the input X-ray
energy, it is necessary to subtract the Dark Levels and correct possible optical Light
Leaks. Dark Levels are non-zero pixel PHA caused by leakage currents in the CCD.
After each passage of South Atlantic Anomaly (SAA), the bias-level is determinedfor each pixel by taking average of the PHA during an exposure of about ten frames.
Another components of the bias-level is the light leak level. Optical and UV lightmight enter the CCD due to imperfect shielding (“light leak”) and produce PHA
that is not related to X-rays. The Light Leaks are calculated on board with the
PHA data after the subtraction of the Dark Levels. Analysis of the ASCA SIS data
showed that light leaks are considered to be rather uniform over the CCD, hence anaverage is calculated for 256 × 114 pixels (this size was 64 × 64 before January 18,2006) in every exposure and its running average produces the Light Leak. In spiteof the name of Light Leak, Light Leak does not represent in reality the optical/UV
light leak to the CCD. It mostly represents fluctuation of the CCD output correlated
to the variations of the satellite bus voltage. XIS has little optical/UV light leak,
which is negligible unless the bright earth comes close to the XIS field of view.
The main purpose of the on-board processing of the CCD data is to reduce
the total amount transmitted to ground. For this purpose, the DE searches for acharacteristic pattern of charge distribution (called an event) in the pre-processed(post- Dark Levels and Light Leaks subtraction) frame data. When an X-ray photon
is absorbed in the depletion layer with photo-absorption process, a primary electron
cloud with a size of ≤ 1µm is produced. The primary cloud grows in size to ≤ 5µm
as it drifts toward the electrode. Because the sizes of the clouds at the electrode aresmaller than the pixel size, most of the clouds are captured by four neighboring pixels

27

CHAPTER 3. XIS[Definition][Examples]

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

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

Figure 3.4: Grade definition and the rule to add PHA. The events out of thesecriteria are defined as Grade 7 to 11.
at most. An event is recognized when a valid PHA (one between the Event Lower
and Upper Thresholds) is found that exceeds the PHA in the eight adjacent pixels.
Then, the pulse height patterns around the local maximum pulse height pixels areclassified into 12 kinds of so-called grades on ground. At the time of writing, theevents classified as Grade 0, 2, 3, 4, and 6 (figure 3.4) are only regarded as X-ray
events. On the other hand, charged particles events produce a larger electron cloudthan X-ray events, hence the events are regarded as grade 7 or larger.
Determining the optimum split threshold level (figure 3.4) is important to measure the pulse height correctly. In the section 4.3, I will describe the results on the
ground calibration to determine the optimum value.

3.2.3

Energy Resolution

Although the number of electron-hole pairs created is a function of the energy of
the X-ray, it is a statistical process and there is a normal distribution around theactual X-ray energy. The energy resolution of the Si crystal detector is followed by
FWHM(eV ) =√s

8 ln 2 × W N 2 +

E0 F
,
W

(3.1)

where W is an energy required for one electron-hole pair creation (3.65 eV/pair), N
is a rms readout noise in electrons, E0 is an incident photon energy, and F is theFano factor, which is the ratio of the observed statistical dispersion of the electron
number to E0 /W , that is, the deviation from the Poisson statistics in the electronhole creation process. Actually W and the Fano factor for Si semiconductor has

CHAPTER 3. XIS

28

Figure 3.5: The predicted proton spectrum inside a simple model of the XIS shielding. The shielding is assumed to be spherically symmetric with a thickness of 10 g
cm−2 .
the energy and temperature dependance. According to a Monte Carlo simulation,
the former varies from ∼ 4.0 to ∼ 3.65 in the soft energy band of 0.1 - 2.0 keV,
while nearly constant (∼ 3.65) in the high energy band of 2.0 - 10.0 keV. The
latter ranges from 0.08 to 0.16 according to various experiments (Knoll (1989) andreferences therein).

3.2.4

Radiation Environment

Suzaku operates in a low earth orbit with several daily passes through the SAA.
The initial orbit is 568 km at an apogee, which is very close to that of ASCA (anapogee of 613 km). Results from a NASA charged particle environment study ofthe predicted ASCA orbit (Stassinopoulos & Barth 1991) indicated that geomagnetically trapped protons are the greatest concern for radiation damage. Figure 3.5
shows the predicted proton spectrum incident on the CCDs based on this study
assuming a shielding depth of 10 g cm−2 . Figure 3.6 shows the specific energy loss
for protons in silicon. Integrating the product of figure 3.5 and 3.6 yield the daily
dose in “rads” of silicon expected to be about 0.5 rads per day (A rad is a unit ofenergy loss equal to 100 ergs per gram).

3.2.5

Charge Injection Capability

One drawback to an X-ray CCD in-orbit is the degradation of the gain and energy
resolution due to an increase of the CTI and dark current. When a proton dislodges

CHAPTER 3. XIS

29Figure 3.6: The specific energy loss for protons in silicon.
a silicon atom in the lattice of the CCD, the vacancy produced is fairly mobile.It
will move through the lattice and possible reach the edge of the lattice where it is
no longer a vacancy. Sometimes, the vacancy will come into the vicinity of anothervacancy or an impurity atom also in the lattice. Often, the pair will form a complex
which have a minimum energy and become physically stuck at a localized position in
the lattice. The complex modifies the silicon band gap structure of the local lattice
by introducing a mid band gap energy level, which is regarded as a charge trap (CT).
The exact energy of this CT depends on the components of the trap complex. This
defect is especially severe for low-energy protons because they deposit more energy
than high-energy protons in the CCD transfer channel. The main origin of the CTI
and consequent gain degradation is the increase of CTs. In fact, soon after thelaunch of the Chandra, all of the FI CCDs of the ACIS suffered the degradation of
the CTI. Devices similar in type to ACIS CCDs have been tested by the high-energy
protons (10 - 40 MeV) but not by the low-energy protons before the launch. The
low-energy protons with an energy of ∼150 keV release a major part of their energy
at the transfer channel of the ACIS CCDs, which is located roughly 1 µm below
the electrodes. Although a thick shielding around the CCD camera can significantly
reduce the proton flux on the CCD, the radiation damage cannot be ignored over a
mission lifetime of several years.
Recently a CI technique has been developed (Prigozhin et al. 2004; Bautz et al.
2004; LaMarr et al. 2004; Smith et al. 2004; Meidinger et al. 2000). In principle,
CPs are artificially injected through a CI gate (Tompsett et al. 1975) into eachtransfer channel (hereafter, we call this vertical transfer channel a column, and thehorizontal channel a row) and are subsequently read after the charge transfer in thesame manner as the X-ray events. This method allows the amount of charge loss

CHAPTER 3. XIS30(δQ) on a column-to-column basis to be measured and/or mitigated, which in turn,can potentially be a powerful tool for CTI calibration. By referring to figures 3.3
and 2.7, here I explain the essential function of CI. For the brevity to describe the
CI technique and its results, notations of parameters, which will be frequently usedin this paper are listed in table 3.1.
Table 3.1: The notation list of parameters for the CI.
Parameters
Notation
Injected charge (for each column)
Q (QCOL )Readout charge (for each column)
Q0 (Q0COL )Charge loss in the transfer (for one column)δQ (δQCOL )Charge Transfer Inefficiency (for one column)
CTI (CTICOL )column-dependent CTI obtained with CI
CTICI
column-averaged CTI obtained with the cal. source CTICAL
Because a fraction of the charge (δQ) is lost due to the CTs in the transfer
channel, the readout charge (Q0 ) is Q0 =Q-δQ. On the other hand, we need theinformation of δQ in order to estimate the CTI. We hence need to adopt the injection
pattern with which we can know both of of Q0 and Q simultaneously with normal
readout. Then we employ the following CI. After injecting a test CP with charge
Q in one row, we inject the same charge (Q) in five subsequent rows: the preceding
four rows are called the sacrificial CPs and the last one is the reference CP. Figure
3.7 shows a schematic view. The test charge is separated from trains of sacrificial
charges at the request of the event detection algorithm.
The test charge Q may suffer from traps in the transfer channel (column), andtherefore, the readout charge (Q0test ) should be Q-δQ. On the other hand, thepreceding sacrificial charges may fill the CTs, and hence, the subsequent reference
charge Q may not be trapped if the clocking time is shorter than the de-trapping
time scale (Gendreau et al. 1993), which is the time until the trapped charge isemitted to the conduction band.
We put the four CPs as sacrificial charges because if the loss in the first packet
is large, the following packet may also lose some charge because as the first packet
loses electrons, it occupies smaller volume and does not fill some of the traps inside
the volume of the following larger packet. This effect can be seen in figure 3.8, whichshows the PHA of the injected CPs before the dark-level subtraction as a functionof Y coordinate averaged among columns in an arbitrary segment of the proton
damaged chip (the ground experiment is described in chapter 4). This confirmsthat the technique of the sacrificial charges works well and four pixel is sufficientas the sacrificial CPs satisfy the technique. Figure 3.8 also shows the emission oftrapped electrons back into the empty potential wells as a tail after the last injected
packet.
Thus, the readout charge (Q0ref ) from the reference CP should be approximately
equal to Q. Figure 3.9 shows a frame image taken during our experiment. The
positions of the CP trains are periodically shifted by one column and the sacrificial

CHAPTER 3. XIS

31

Figure 3.7: Schematic view of how to measure δQ using the test and reference
charges. The amount of lost charges (δQ) can be estimated by comparing the PHAof test charge, which suffer from CTI, to that of reference charge, which is notaffected by the traps.

Figure 3.8: The PHA of injected charges averaged among columns in an arbitrary
segment (seg.C) of the proton-damaged chip as a function of the Y coordinate. The
vertical axis is chopped for the brevity.

32

CHAPTER 3. XIS

Figure 3.9: Frame image of the XIS during the CI experiment. The gray scale shows
the pixel level, not the number of events. There are two events per each column:
The lower one is the test charge and the upper one is the reference charge. Notethat the sacrificial charges are not displayed at the request of the event detection
algorithm.
CPs are not shown to allow the proper event detection algorithm. δQ during thetransfer can be measured by the ratio of mean pulse height amplitude (PHA) of thetest CP to that of the reference CP. By selecting different Qs, we can investigate therelation between Q and δQ. The results of the ground and in-orbit experiments on
the CI capability will be described in section 4.6 and 5.2.

3.3

Charge Transfer Inefficiency

Shockley-Read-Hall theory (Shockley & Read 1952) describes the dependance ofminority carrier capture and release time scales on temperature and carrier density.
Traps capture electrons from passing charge packets on a capture time scale (tc ) is.
1
,
(3.2)
vt nσ
and the traps then release the captured electrons on an escape time scale (te ) of
tc =

te =

1
−(Ec −Et )vt Nc σe kT
The resulting trap occupancy rate is given by

.

(3.3)

dnt
nt N t − n t
=− +
,
(3.4)
dt
te
tc
where σ is the electron capture cross section, vt is the thermal velocity of an electron,n is the number density of free electrons, Nc is the density of states in the conduction
band, Ec is the band gap, Et is the trap energy, k is Boltzmann’s constant, T is thetemperature, nt is the number density of trapped electrons, and Nt is the number
density of traps.

CHAPTER 3. XIS

33

Figure 3.10: The detrapping times for traps of various energies as a function oftemperature. The P-V center trap has an energy level of 0.3, 0.4, and 0.5 eV. The
trap cross section is 3.5×10−15 cm−2 for each curve.
The phosphorus-vacancy (P-V) complex is a displacement induced trap whichis the most significant in the phosphorus doped buried channels of CCD such asdevices used in ACIS and XIS. The proton induced lattice vacancies combine withthe phosphorus atoms and make P-V complexes. The P-V trap is characterized witha cross section of 3.5×10−15 cm−2 and a trap energy of ∼0.4 eV. With such a trap anda typical CP of 1000 electrons, the capture time is virtually instantaneous (less than
10−7 s) compared to transfer rates used in our devices (24 µs per shift). However,as shown in equation 3.2, trapping time is a function of the density. For small CPs(0 - 10 electrons for dark current and 50 - 3000 for X-ray signals), the density ofelectrons is difficult to calculate. However, the density should be within an order ofmagnitude of the phosphorus dopant density (2×1016 cm−3 ). The detrapping times,
however, are on the order of tens of milliseconds.
Figure 3.10 shows the detrapping times for traps of various energies as a functionof temperature, using equation 3.3.
The P-V center trapping model explains the amount of charges, proton fluence,clock speed, and temperature dependance of our X-ray CCD performance. Whena CP from an incident X-ray encounters an empty trap during a CCD readout, if
the trapping time is sufficiently shorter than the time the CP is shifted away, it isvery possible that the trap is left full and the CTs will contribute little to the gainand energy resolution. Each trap has a capacity to hold a single electron. Given an
empty trap, the probability that a passing CP will lose an electron is nearly unity.
Once a CP has suffered a loss to a trap, if the detrapping time is long comparedto the time till the next X-ray induced CP arrives, the trap will remain filled andCHAPTER 3. XIS

34

Figure 3.11: The case where only the preceding charge packet lose electrons to agiven trap. The detrapping time is greater than the time between passing charge
packets.
will not affect the second CP; thus minimizing the spectral degradation. If thedetrapping time is short compared to the time till the next X-ray induced CP arrives,
the trap will probably have emptied and would have a chance to degrade the second
CP, and hence maximizes the damage to the spectral response. The former situationis shown in figure 3.11.
In this section, we call the preceding CP in figure 3.11 a sacrificial CP andfollowing CP in figure 3.11 a beneficial CP. Figure 3.12 shows the spectra of 55 Fe
data collected from the data obtained in 2005 (upper panels) and 2006 (bottom
panels) in-orbit. We chose the 55 Fe data of XIS2 from the statistical point of view.
As the distance between a sacrificial CP and beneficial CP increase (from (a) to(c) in figure 3.12), the line center of sacrificial CPs decrease against beneficial CPs.
In other words, charges are lost with increasing distance between two CPs because
increasing distance corresponds to an increasing time and an increasing probability
that the traps filled by the sacrificial CP have trapped.
Because the XIS utilizes the frame transfer type CCD, even if the distance between the sacrificial CP and beneficial CP is fixed, the amount of lost charge isdifferent among ActY. This is because the number of slow parallel transfer in the
FS depends on ActY. Furthermore, the difference of the pixel sizes between the IA
and FS makes more difficult to investigate the charge loss at each region. The total
number of traps seen by CPs as a function of the distance to the sacrificial CP canbe formulated as followings (Gendreau 1995) with an assumption that a trap cancontain at most one electron.
Consider two CPs in one column of the IA. Both the IA and FS have N pixels.

CHAPTER 3. XIS

35

Figure 3.12: Energy spectra of 55 Fe events of XIS2 collected in orbit. For each figure,
upper and bottom panels is the data collected in 2005 (from August to December)and 2006 (from January to June), respectively. The distance between a sacrificial
CP and beneficial CP is shorter than 50 pixels (upperleft panel), from 50 to 100
pixels (upperright panel), and larger than 100 pixels (bottomleft panel). The blackand red spectra show those of sacrificial CPs and beneficial CPs, respectively.

36

CHAPTER 3. XIS

The sacrificial CP, Qs is located at pixel y and the beneficial CP Qb is located a
distance d behind Qs . The following constraint holds on y, d, and N :y + d ≤ N.

(3.5)

Thus, for a given distance to the sacrificial CP, the maximum row number that thesacrificial CP can come from isymax = N − d.

(3.6)

The number of traps per pixel in the IA and FS are respectively given as
Ni = N L pi , Nf = N L pf ,(3.7)

where NL is the number of traps per micron along the transfer direction, pi and pf
are the pixel length along the transfer direction in the IA and FS, respectively.
Then the number seen by the sacrificial CP is
Ns = N Nf + yNi ,(3.8)

and the number seen by the beneficial CP is
Nb = N Nf + yNi + dNi .

(3.9)

For the simplicity, here is another assumption that the CPs are large enough sothat the capture times are very short compared to the pixel to pixel transfer times
so that empty traps have unity probability of capturing electrons from the passing
packets. Then, the number of empty traps (Ne ) seen by the beneficial CP is

Ne (y, d) = Ni d + Ni y(1 − exp(

−dcif
−dcif
)) + Nf (N − y − d)(1 − exp(
))
τe
τe

−dcf s))τedX−dcif + j(cif − cf s )+Nf(1 − exp(
)),
τe
j=1
+Nf y(1 − exp(

(3.10)

where τe is the trap emission time constant, cif and cf s are the time it takes totransfer from one pixel to the next during the FS transfer and parallel transfer in
the FS, respectively.
The first three terms in equation 3.10 describe traps encountered exclusively
during the FS transfer. The first term corresponds to the traps which could not be−dctouched by the sacrificial CP. The expression (1 − exp( τe if )) is the availability oftraps which increases with a time constant τe or equivalently with a position constantτe. The fourth term corresponds to traps encountered by both packets exclusivelycifduring the slow parallel transfer in the FS. The availability of these traps have aposition constant cτfes . For the XIS CCD, cf s /cif = 277. The fifth term is due totraps filled by the sacrificial CP during the FS transfer, but seen by the beneficial

37

CHAPTER 3. XIS

CP during the slow parallel transfer in the FS. For these traps, the passing time
between the sacrificial CP and beneficial CP is a function of position, hence eachterm in the summation corresponds to a different pixel. The summation term isgiven by

Nf

d
X

j=1

(1 − exp(= Nf d −−dcif + j(cif − cf s )))τeexp(cif − dcif − dcf s ){exp(dcf s /τe ) − exp(dcif /τe )}.
exp(cf s /τe ) − exp(dcif /τe )(3.11)

Hence we get Ne as
−dcif
−dcif
)) − Nf {d(1 − exp())τeτe
exp(cif − dcif − dcf s ){exp(dcf s /τe ) − exp(dcif /τe )}}−d +exp(cf s /τe ) − exp(dcif /τe )−dcif−dcf s+y{(Ni − Nf )(1 − exp()) + Nf (1 − exp())}. (3.12)τeτe

Ne (y, d) = Ni d + Nf N (1 − exp(

Suppose that the flux of X-rays per row is given by the function g(y), then we cancalculate the average number of empty traps, Navg (d) seen by CPs a distance d awayfrom the sacrificial CP. This would be the average loss of charge due to trappingand is given by
Navg (d) =yXmaxf (y, d)Ne(y, d),(3.13)

y=0

where
g(y)
f (y, d) = Pymax.
y=0 g(y)(3.14)In order to estimate the trap density for the XIS CCDs using the in-orbit data, theflux per row, g(y), must be determined for a given observation by looking at theimage. In general, the resulting f (y, d) may be very complicated. In the case of
uniform illumination,f (y, d) =1
ymax

Then the equation 3.13 becomes as

Navg (d) =

1

yX
max

ymax y=0

=

1
.
N −d

Ne (y, d)= (Ni + Nf )d + Nf N (1 − exp(

−dcif
))
τe

(3.15)

CHAPTER 3. XIS

38

−dcif
))
τe
exp(cif − dcif − dcf s ){exp(dcf s /τe ) − exp(dcif /τe )}+}exp(cf s /τe ) − exp(dcif /τe )−dcif−dcf s(ymax + 1){(Ni − Nf )(1 − exp()) + Nf (1 − exp(
))}
+
2
τe
τe
−dcif
))
(3.16)= NL {(pf + pi )d + pf N (1 − exp(τe−dcif
−pf {d(1 − exp())
τe
exp(cif − dcif − dcf s ){exp(dcf s /τe ) − exp(dcif /τe )}+}exp(cf s /τe ) − exp(dcif /τe )−dcf s−dcif(ymax + 1))) + pf (1 − exp())}}{(pi − pf )(1 − exp(
+
2
τe
τe
−Nf {d(1 − exp(

In equation 3.16, NL , pi , and pf are introduced to get rid of Ni and Nf . The only
unknown parameters are the linear density of traps, NL , and the detrapping time
constant τe . In the case of the XIS CCDs, the following values hold: pf is 18µm, pi
is 24µm, cif is 24 µsec, and cf s is 6.8 msec.
Although we cannot estimate NL and τe precisely only with the 55 Fe data becauseof the poor statistics. Nevertheless in figure 3.13, we can see the tendency thatthe difference of the line center energy between the sacrificial and beneficial CPs
decrease as the distance of the two CPs increase. This means that increasing distancecorresponds to an increasing probability that the traps filled by the sacrificial CP
have detrapped and that the beneficial CPs encounters the empty traps (equation
3.16).

CHAPTER 3. XIS

39Figure 3.13: The line center of the beneficial CPs in figure 3.12 subtracted by thatof the sacrificial CPs. The black and red data indicates the data of 2005 and 2006,
respectively. The negative correlation means that increasing distance corresponds toan increasing probability that the traps filled by the sacrificial CP have detrapped.

Chapter 4
Ground Calibration of XIS
4.1

Goal for Calibration Accuracy

Figure 4.1 shows a spectrum of 55 Fe obtained with the XIS Flight Model (FM) sensor.
We can see not only the main peaks of Mn Kα and Kβ, but also the escape peaks,
fluorescent peak, tail, and constant component. The ratios between the intensitiesof these components are sensitive to the X-ray energy, gate structure, and thicknessof the depletion layer. In order to interpret X-ray spectra obtained in orbit, we
aim to construct a precise energy redistribution function (response function) for
monochromatic X-rays.
The goals of our calibration are as follows:
1. The line center energy should be correctly measured with a systematic errorof <0.1% of the incident X-ray energy.
2. The energy resolution should be determined within an accuracy of 1%.
3. The QE should be measured within an accuracy of 10%.
4. The spectral shape of monochromatic X-rays, including the Si fluorescent peak,escape peaks, tail, and constant component should be measured with errors of
<3%.

4.2

Calibration System

The calibration was conducted with the fluorescent X-ray generation system at Kyoto University (Hamaguchi et al. 2000). The X-ray generation system involves continuous X-rays from an X-ray generator which shines an selected target within achamber. The targets (Al, Cl, Ti, Fe, Zn, Se, and Polyethylene) are on a rotatingstage. Secondary fluorescent X-rays are incident on the XIS in the other chamber.
In the case that XIS is moved out of the beam line, the intensity and spectrum of
the beam can be taken using a backward windowless solid-state detector (SSD).

40

41

CHAPTER 4. GROUND CALIBRATION OF XIS

Figure 4.1: A spectrum of the XIS FI sensor of the 55 Fe source. There are several
components in this spectrum such as the main peak, sub peak, constant component, escape peak, and Si Kα peak, all of which are labeled. The final purpose ofcalibration is to determine the flux of each component to within an error of 3%.

5

3
1

2

4

Figure 4.2: Schematic view and photograph of the X-ray generation system at Kyoto
University. The primary X-rays from the Kevex X-ray generator irradiate the second
targets and the secondary fluorescent X-rays irradiate the XIS sensor placed in themain chamber or the detector (SSD in this figure) behind the XIS.

CHAPTER 4. GROUND CALIBRATION OF XIS

42Figure 4.3: Frame image of the proton-damaged FI sensor used in the ground experiment. The gray scale is the pixel level, not the number of events. Gaps between
segments are horizontal over-clock regions. Proton irradiation causes the circularregion in segments A to C to have systematically different pixel levels due to thedark current.
We also conducted the experiments with the proton-damaged CCD chip at MIT
for the CI issues. The CCD chip was damaged by protons utilizing the cyclotron
at the Northeast Proton Therapy Center(NTPC) at Boston. The 160 MeV primarybeam of protons was degraded with lexan blocks to a spectrum peaked at 40 MeV
with an FWHM of about 4 MeV. The beam was irradiated onto the circular region
shown in figure 4.3. The total fluence was 2.0 × 109 cm−2 , which is approximately
the same as that the XIS receives during several years in orbit. For the ground test,
we used the secondary fluorescent X-rays and a 55 Fe radioisotope at MIT. Figure
4.4 compares PHA spectra from 55 Fe source extracted from proton-damaged andundamaged region in segment B, respectively.
During the both calibrations, the sensors were in environments with pressuresof ∼ 10−6 Torr and CCD temperatures of –90 ◦ C. In the following, we describe theresults of our calibration.

4.3

Readout Noise and Split Threshold

First, we investigated the readout noise of each sensor segment. All sensors showednoise levels of ∼2 - 3 ADU (RMS). From these values, we determined the split
threshold for event extraction for the FI chips. In order to measure the pulse heightcorrectly, it is essential to determine the value of the split threshold considering thereadout noise. Then, ∼5 σ of the readout noise, that is, ∼ 15 ADU, is preferable as
the split threshold. We then set the threshold to be 20 ADU considering the margin.
In the case of the BI chip, we found that the energy-dependent threshold levelshould be introduced in order to optimize the QE and energy resolution (Yamaguchi
et al. 2006a). Contrary to the case of the FI chip, the electron clouds produced by

CHAPTER 4. GROUND CALIBRATION OF XIS

43Figure 4.4: Spectra of 55 Fe extracted from the proton undamaged region (upper
panel) and damaged region (bottom panel). The line center and width are bothdegraded in the bottom panel compared with that of upper panel.
the low energy X-ray photons drift in the depletion layer for a longer distance thanthose produced by the high energy X-ray photons. Then, the FWHMs of the spatial
distribution of the electron clouds by the low energy X-rays are larger than thoseby the high energy X-rays, although the amount of the split charges is smaller thanthose by the high energy X-rays. As a result, if one set the low split thresholdconsidering the large FWHM of low energy events, the significant part of the highenergy events is regarded as grade 7 as long as one adopt the energy-independent
split threshold. On the other hand, one fails to sum up the pulse height in the split
pixel for the low energy events with the high split threshold optimized for the highenergy events. In fact the grade distribution with the split threshold of 7 ch is asshown in figure 4.5.
Then, we made the spectra with various split threshold levels for each fluorescentline prior to calibrations of other issues and investigated the FWHM of the spectrum
peak and the number of events as shown in figure 4.6. While the FWHM improvesas the split threshold decrease, the event detection efficiency (i.e. QE) degrades. We
determined the optimum threshold level from these plots so that the increment of
the QE becomes less than 3% when one increments the split threshold level by one
channel.
Then we determined the optimized split threshold levels as a function of energy
assplitth. = 2.21 × log 10E0 + 10.36,(4.1)

as shown in figure 4.7, where E0 is the energy of X-rays. This improves the QE 13,
36, and 40% at the energy of 2.6, 6.4, 8.6 keV compared to the values with constantsplit threshold of 7 ch, respectively.

CHAPTER 4. GROUND CALIBRATION OF XIS

44Figure 4.5: Percentage of the Grade 0 and 7 events as a function of the energy
of incident X-rays. The significant percentage of the X-ray events are regarded as
grade 7 in the high energy band and hence virtually diminish the QE of the BI chip.

4.4

Gain and Energy resolution

Adopting the split threshold determined in the above section, we extracted the Xray events from each data set and made a spectrum. For the line spectrum of eachfluorescent X-ray, we fitted each peak (Kα and Kβ) with two gaussians (main and
sub peak component) considering the response distribution function (figure 4.1).
Figure 4.8 shows the gain of XIS FI1 segment A. The grade of the extracted events
are 0, 2, 3, 4, and 6. The energy-pulse height relation is quite linear in the energy
band of 1 - 12 keV. The error of the gain is less than 0.1% for all segments of all
sensors. The difference in gain among the segments is less than 10%.
Figure 4.9 shows the spectral resolution of XIS FI1 segment A. The grades of theextracted events are the same as in the previous section. We fitted the data withFF
the equation 3.1. The free parameters are N and W. The best fit value of N and W−2
are 2.6±1.2 electrons and 4.34±0.06×10 , which corresponds to the Fano factor of∼ 0.16±0.02 when we assume that W = 3.65. Considering the gain obtained in theprevious section, the values for the noise and Fano factor are consistent with theresults of Section 4.3 and with the literature (Knoll 1989). We achieve an energy
resolution of ∼ 130eV@5.9 keV for all sensors. The error in the resolution is lessthan 1% for all segments of all sensors.

CHAPTER 4. GROUND CALIBRATION OF XIS

45

Figure 4.6: Energy resolution (FWHM) and detected number of events as a functionof the split threshold. The latter is normalized to that in the case with a split
threshold of 7 ch.

CHAPTER 4. GROUND CALIBRATION OF XIS

46

Figure 4.7: The optimized split threshold levels as a function of the energy of the
X-rays. A solid line is the best-fit result of the data (equation 4.1).

Figure 4.8: The XIS energy scale linearity of XIS1 segment C (left panel) and XIS2
segment C (right panel). We obtained a gain of 265.1 ch/keV.

47

CHAPTER 4. GROUND CALIBRATION OF XIS

Figure 4.9: The spectral resolution as a function of energy for the XIS1 segment C(left panel) and XIS2 segment C (right panel). We obtained a resolution of 131 eV
in FWHM at 5.89 keV.

4.5

Quantum Efficiency

Here we converge on the QE measurement in the high energy band executed at
Kyoto University. Figure 4.10 shows the QE of XIS FI1 segment A. The solid
curves represent the simulated function with a simple layer model described in thefollowing.
The probability that the X-ray with energy of E0 is absorbed within the depletion
layer isPabsorp = exp(−µSi dm nm ) exp(−µSiO2 do mo )(1 − exp(−µSi dd nd )),(4.2)

where dm , do , dd , nm ,no , and nd are thicknesses and density of metal, oxide, depletion
layers, respectively. µSi and µO are the mass absorption coefficients of silicon andoxygen, respectively. We assume that
µSi = 104.3−2.78 log10 (E0 ) (E0 ≤ 1.839keV)
103.19−2.71 log10 (E0 )+0.0755 log10 (E0 )2.0 +0.873 log10 (E0 )3(4.3)(E0 > 1.839keV), (4.4)µO = 103.65−2.74 log10 (E0 ) (E0 > 0.543keV).

(4.5)

The central curve assumes nominal values for the thickness of the electrode layer,
insulation layer, and depletion layer, that is, 0.33 µm, 0.27 µm, and 74 µm respectively. The uppermost and lowermost curves represent the cases when the thicknessof each layer deviates from the nominal value by ∼ ± 10%. The QE value in thesoft energy band is subject to the thicknesses of the ineffective layers, that is, theelectrode and insulation layer. In the high energy band, the QE depends on the
thickness of the depletion layer. As shown in figure 4.10, the QE of XIS at ∼4.5 keV

CHAPTER 4. GROUND CALIBRATION OF XIS

48

Figure 4.10: The QE as a function of energy for XIS FI segment A. The solid curves
show simulated QE, where the input parameters are the thickness of the electrodelayer, insulation layer, and depletion layer. We determined the nominal value to be
0.33 (0.25–0.45) µm, 0.27 (0.20–0.35) µm, and 74 (60–80) µm, respectively. Having
confirmed that the QE at 4.5 keV is almost constant (96±0.7%), we normalized the
QE to 96%@4.5 keV.
(Ti Kα line) is 96±0.7% as far as the thicknesses of each layer deviates 10% fromits nominal values. Hence the QE value at 4.5 keV can be a good reference value
considering the requested calibration accuracy. Hereafter we assume the QE at 4.5
keV is 96%.
After the data acquisition with the XIS, we shift the XIS and collect data with awindowless SSD, which is located behind XIS. Since the QE of the SSD is ∼ 100%at the energy band of > 1 keV, we measured the normalized QE of XIS as a functionof the intrinsic energy, and ∼52%@8.6keV is achieved. The error in the normalized
QE is less than 3% for all segments of all sensors.

4.6

Performance of Charge Injection Capability

4.6.1

Stability of the Amount of Injected Charges

If the amount of injected charge (Q) can be more accurately controlled than that ofcharge dispersion of X-ray events, then CI capability offers obvious advantages overconventional CTI measurements using X-ray calibration sources. We hence checked
the stability of Q0 with the non-damaged chip, for which Q0 most precisely reflect Q
because of the small number of CTs. Throughout this experiment, a constant offset
voltage was applied at the input gate.
Figure 4.11 shows the spectra of fluorescent X-rays (Ti Kα) and the CI test
events, which have approximately the same equivalent energy as X-rays. The datais collected with a non-damaged chip and a duration of 350 frames (2.8 ksec). Eventsare extracted from a specified column for all data sets. The FWHM of the CI data
(80±5 eV) is significantly better than that of the X-ray spectrum (130±10 eV). The
former is due to the stability of the amount of injected charges and the noise on

CHAPTER 4. GROUND CALIBRATION OF XIS

49

Figure 4.11: (Upper panel:) Spectra of Ti Kα fluorescent X-rays from an exposure of
350 frames taken in the ground experiments. (Middle panel:) Spectra of CI events
collected on ground from an exposure of 400 frames. (Bottom panel:) Same as
middle panel but collected in orbit from an exposure of 800 frames. The PHA levelis slightly different from that in middle panel because the offset voltage is varied
from that adopted in ground experiments. For all panels, the events are extracted
from an arbitrary column (X=2). Horizontal axis is the dark-level-subtracted PHA.
the charges (Prigozhin et al. 2004) while the latter is primarily due to the Poisson
statistics and the Fano factor. Thus, we verify that Q in a specified column issufficiently stable to estimate δQ on a short timescale of several ksec.
Next, we investigate whether the ratio of Q0test to Q0ref is proportional to the CTI,
which can be reflected by 55 Fe events, for each column with the damaged chip at
MIT and non-damaged chip at Kyoto Univ. Due to the stability of Q over several
hundreds of frames, the histograms of the PHA of Q0ref and Q0test are successfully
fitted with a single Gaussian for each column. For the 55 Fe events, we extracted
the events from upper and lower edges of the imaging area (Q0upper and Q0lower ). The
extracting regions were 100 rows in order to satisfy the statistics requirement forthe spectral fitting. Figure 4.12 shows the correlations between Q0test /Q0reference of
the CI events and Q0upper /Q0lower of the 55 Fe events. For the non-damaged chip (theleft panel), because δQ during the parallel transfer is 0 or 1 electron for 55 Fe events,
the correlation can be hardly seen. In the case of the damaged chip, the CTI is
significant and the most of the charge loss occur in the circular region as shown infigure 4.3. Then we can see a clear positive correlation between the 55 Fe and CI
events for damaged chip (especially in segment A, B, and C). The best-fit slope is∼1.05 and the correlation coefficient is 0.94 (d.o.f.=976). Hence, the ratio of thetest and reference charges properly represents the CTI of each column.

50

CHAPTER 4. GROUND CALIBRATION OF XIS
1Typical error

Typical error

1.004

1.002

Qupper/Qlower for 55Fe

Qupper/Qlower for 55Fe

0.98

1

0.998

0.96

0.996

0.996

0.998

1
Qtest/Qref. for CI

1.002

1.004

0.96

0.98

1

Qtest/Qref. for CI

Figure 4.12: Ratio of Q0test and Q0ref for the CI events (horizontal axis) and thatof Q0upper and Q0lower for the 55 Fe events (vertical axis). Open triangles, circles,
crosses, and squares represent the columns in segment A, B, C, and D, respectively.
Left:Plot for non-damaged chip. Note that plots for segments B, C, and D have thesame distribution as segment A. Hence, segments B, C, and D are ignored from thefigure for simplicity. Right:Plot for a proton-damaged chip.

4.7

Energy Redistribution Function

The branching ratio for each component such as in figure 4.1 is estimated. Wethen finally made the response function for FI and BI chips. determined with theground experiments. Figure 4.13 shows the spectrum of 55 Fe collected with FI and
BI chip. Although the response function reproduce the spectrum well over the entire
energy band, some residual is seen between the Kα and Kβ line. This indicates the
additional component in the spectra. The Radiative Auger Emission (RAE;Åberg
1971; Richard et al. 1975; Jamison et al. 1976) is the appropriate candidate to explain
this residual as shown in figure 4.14.

CHAPTER 4. GROUND CALIBRATION OF XIS

51

Figure 4.13: The spectra of 55 Fe collected the response function we obtained withground experiments. The fitted model is the response function we determined withthe ground experiments. Left and right panel show the spectra of the BI and FI
chip, respectively.

Figure 4.14: Fit results of 55 Fe spectrum with the response function around theKα and Kβ line. (Left panel:) The XIS1 spectrum is fitted with a model of two
monochromatic lines with energies of 5.9 and 6.5 keV. (Right panel:)The XIS1 spectrum is fitted with a model of two monochromatic line and a broad line emission,
which approximate the RAE component.

Chapter 5
Onboard Calibration of XIS
5.1

Performance Degradation in Orbit

As described in 3.2.5, the radiation damage onto the CCD has been integrated by
the protons in SAA. Figure 5.1 and 5.2 show the degradation of the line center
energy and line width of 55 Fe Kα line without any CTI correction as a function oftime from the first light to the June 2006. Here the line width means the value when
one fits the main peak with single gaussian using the response function made fromthe ground experiments (see 4.7). The both parameters gradually decrease at a rateof ∼ 120 eV per year and ∼ 50 eV per year, respectively. The goal of the in-orbit
calibration for the gain is the same as that in the ground calibration.

5.2

CTI correction with CI Capability

Most of the major X-ray missions are provided with one or more calibration sources
to measure the CTI and gain. However, the limited flux of calibration X-rays impedes a detailed and frequent investigation of the spatial structure of the CTI overthe chip, which makes it difficult to perform a column-to-column charge compensation. Hence we have been periodically executing CI to determine the CTI parameters.

5.2.1

CTI Measurement with Checker Flag CI

In order to apply the CI function to the observational data, the on-orbit CI experiment was conducted during the calibration observations of the SNR 1E0102–72.3 inthe Large Magellanic Cloud. Table 5.1 summarizes the experimental logs. All thedata was acquired with the normal clock and with the 3×3 or 5×5 mode (Koyama
et al. 2006a). Table 5.1 also lists the equivalent energy of the CPs. Figure 3.9 showsa frame image taken during our experiments. The positions of the CP trains are
periodically shifted by one column to allow the proper event detection algorithm.
The sacrificial CPs are not read because of the same reason. The value δQ after thetransfer can be measured by the difference of the mean PHA of the test events to52

CHAPTER 5. ONBOARD CALIBRATION OF XIS

53

Figure 5.1: The line center energy of the each calibration source as a function oftime without any CTI correction. The dashed lines indicate the intrinsic value ofthe Kα line center, 5.89507 keV (Bearden 1967; Krause & Oliver 1979). Each plotcorresponds to a continuous exposure toward a celestial object. The horizontal axis
AETIME is the time in seconds since 0000 UT Jan 1st, 2000.

CHAPTER 5. ONBOARD CALIBRATION OF XIS

54

Figure 5.2: The line widths of the each calibration source as a function of time without any CTI correction when one fits the Kα main peak with single gaussian using
the response function made from the ground experiments. Each plot corresponds toa continuous exposure toward any celestial object. The horizontal axis is the same
as figure 5.1.

55

CHAPTER 5. ONBOARD CALIBRATION OF XIS
150

seg.A

seg.C

seg.A
seg.C

140
line center PHA

line center PHA

1200

1150

130

seg.D
seg.B

seg.B

seg.D

120

1100

0

200

400

600

800

1000

0

200

X coordinate

400

600

800

1000

X coordinate

Figure 5.3: PHA distribution of the test and reference events as a function of the
X coordinate of XIS3 for two different amounts of injected charges. The equivalentenergies of the CPs are shown in table 5.1. Note that because each segment has its
own readout transistor and analog-to-digital converter, the gain and zero level varyfrom segment to segment, and hence, the PHA level varies. Anomalous columns,
including a hot or flickering pixel are eliminated.
that of the reference events. By selecting different Qs, we can also investigate therelation between Q and δQ.
Table 5.1: The log of the CI experiment on orbit.
XIS0 & XIS2
XIS1 & XIS3
Date

2006/7/17

2006/6/26-27

Time

06:06:50-21:39:46

02:47:07-02:37:55

The equivalent energy (keV)

0.6/4.2/8.0 for XIS0

0.3/7.3 for XIS1of the injected CPs

0.6/3.9/7.8 for XIS2

0.5/4.6 for XIS3

6.0

5.2

Effective Exposure (ksec)

5.2.2

Compensating for the Charge Loss

Given the results of the ground verification experiments (section 4.6.1), our next
aim is to apply the CI techniques to the onboard data.
Figure 5.3 shows the PHA distribution for Q0ref (open circle) and Q0test (cross)as a function of X-coordinate (column). The δQCOL is clearly observed in orbit for
the first time. In order to estimate the δQ dependence on Q, we selected two or
three different Q values for the in-orbit CI experiment (table 5.1). Assuming thesingle power law function of δQ ∝ Qα (e.g. Grant et al. 2004), we derive α for eachcolumn. The results are given in figure 5.4. The weighted mean values of α range
0.62 - 1.00 among sensors. These values are roughly consistent with another groundexperiment (Prigozhin et al. 2004).
The CI data provides only the information on the δQ at the edge of the IA andhence we need to know the Y coordinate dependence of δQ from the data of celestial

56

CHAPTER 5. ONBOARD CALIBRATION OF XIS

3

Typical error

α

2

1

0
0

200

400

600

800

1000

X coordinate

Figure 5.4: Distribution of the index in the relation between δQ and Q (XIS3).
Some anomalous columns are eliminated.
objects that extend over the field of view of the XIS. Figure 5.5 shows the center
energy of the 6.4 keV line as a function of the Y coordinate for the diffuse X-rays
from the Sgr C region (Obs. Sequence = 500018010, Obs. Date = 2006-02-20). Theline center at the lower edge of the image (Y=0) deviates from 6.40 keV due to the
CTI during the quick transfer of all the data in the IA to the frame-store (FS) region(hereafter we call this FS-transfer). However, the origin of the deviations at the otherimage regions is complicated because the charges suffer from three kinds of the CTI:
the CTI in the IA due to the FS-transfer, that in the FS region due to the both of the
FS-transfer and subsequent slow vertical transfer. δQ during the horizontal transfer
is ignored in this work. The CTI depends on the number density of CTs and thetransfer speed (Hardy et al. 1998). The shielding depth and pixel size are different
in the IA and FS. The CTIs therefore may be largely different between in theseareas. However, we cannot estimate each CTI component separately from the total
CTI seen in figure 5.5. Due to this limitation, we phenomenologically compensate
charges utilizing figure 5.5 and CI data as followings. At first we estimate thecolumn-averaged δQ as a function of Y coordinate at the energy of CI events fromthe Q dependance of δQ and figure 5.5. Then, we determine the ratio of the columnaveraged δQ to the δQCOL at Y=0. Finally we estimate δQCOL as a function of Y
coordinate so that the ratio of the column-averaged δQ to the δQCOL at Y=0 is
maintained at any other Y coordinate. Hence we can compensate charges correctlyin the entire region of the IA.

57

CHAPTER 5. ONBOARD CALIBRATION OF XISLine Center Energy (keV)

6.4

6.35

6.3

0

200

400

600

800

1000

Y coordinate

Figure 5.5: The Y coordinate dependence of the center energy of 6.4 keV line emission from the Sgr C region. Photons are each extracted from 1024 × 256 region.
Open triangles, circles, crosses, and squares represent the plots of XIS 0, 1, 3, and3, respectively.

5.2.3

Energy Resolution Improvement

Using the δQCOL determined with the CI experiment, we make the new spectra forthe calibration sources in the observations given in table 5.1. Figure 5.6 shows thecalibration source spectra after the CTICI (upper) and CTICAL (lower) correction.
The tail component seen after the CTICI correction is significantly reduced comparedto that after the CTICAL correction. This strongly indicates that the origin of thetail component is the dispersion of the CTICOL among columns. Figure 5.7 shows
the FWHM of the calibration source spectra after the CTICI and CTICAL correction.
On average, the FWHM is significantly improved from 193 eV to 173 eV. These are
the first in-orbit results for the CI function.
Our final purpose is to demonstrate that δQ parameters are effective for celestial
objects. We apply the δQ parameters to the Tycho’s SNR data (Obs. Sequence= 500024010, Obs. Date = 2006-06-27). Figure 5.8 shows the spectra around the
He-like Si Kα emission line in the west part of Tycho’s SNR for the XIS3 after the
CTICI correction and CTICAL correction. We see the same benefit as shown in figure5.6. Note that the latter is multiplied by 0.8 to avoid confusion. Radiation damage
continuously increases while in orbit, and hence, the benefit of the CI technique will
become more apparent over time.
Because the CTICI correction parameters are time dependent, δQCOL must be
periodically measured. We make two sets of CTICI using δQCOL derived from the CI
experiment in May and July 2006, and apply to the calibration source data taken in
May 2006 for the XIS0 and XIS2. The results are shown in figure 5.9. The CTI in
July is normalized to that in May, while the relative variation in CTICI is preserved.
If the observation time of the CI and 55 Fe data differs by two months, the FWHM
is significantly degraded only for the brightest calibration source. This confirms
that, in practical sense, the interval of two months between each CTI measurement

58

CHAPTER 5. ONBOARD CALIBRATION OF XISCounts/sec/keV

1

0.1

0.01

Counts/sec/keV

10−3
1

0.1

0.01

10−3
5

5.5

6

channel energy (keV)

Figure 5.6: Energy spectra of the calibration source in the segment D of XIS2, after
charge compensation with CI capability (upper panel) and with the calibration
source itself (lower panel). Data is simultaneously obtained with the CI experiment.

FWHM (eV)

200

150

seg.A

seg.D

XIS0

XIS1

XIS2

XIS3

Figure 5.7: FWHM of the calibration source spectra simultaneously obtained with
the CI experiments. Open triangles and circles represent the data after the charge
compensation with the CI technique and with calibration source itself, respectively.

59

CHAPTER 5. ONBOARD CALIBRATION OF XIS

2counts/sec/keV

1

0.5

0.2

1.7

1.8

1.9

2

channel energy (keV)

Figure 5.8: Spectra around the He-like Si Kα emission line of the west part ofTycho’s SNR (XIS3). Data represented by open circles is charge-compensated using
a calibration source, while the other (open triangles) is compensated using the CI
technique. Note that the former is multiplied by 0.8 to avoid a complex expression,and the data around the Si K-edge (1.839 keV) is ignored due to an incomplete
calibration.
is sufficient because there are few celestial objects which has emission line brighter
than this calibration source.
Although the CI technique significantly improves the energy resolution of thecalibration source spectra, the FWHMs are still larger than those before the launch(130 eV) as shown in figure 5.7. Because the charge trapping is a probability process,the number of trapped electrons (δQ) may also have a probability deviation, which
would increase the line width after the transfer. Hence the CTI correction with CI
technique cannot completely restore the line broadening. We confirm this effect infigure 5.10, which shows the FWHM of the CI reference events and those of thetest events before and after the charge compensation for the XIS3 in-orbit data.
In addition, the test events measured on ground (before the radiation damage) are
shown. The smaller FWHMs in all segments of the reference events are due to thefact that the reference events may not loss charges because the CTs are alreadyfilled by the sacrificial events. In fact, the FWHMs of the reference events are
consistent with those of the test events collected on ground although segment D
show anomalous trend.
Another characteristic of the figure 5.10 is that the FWHMs increase along with
X coordinate for the reference events and test events collected on ground. This is dueto the CTI in the CI register. This influences the accuracy of the mean PHA of thereference events and hence the accuracy of δQ. However, the CTI in the CI register
is rather lower than that in the IA and FS because the CPs are injected with aninterval of 3 pixels and hence they can work as sacrificial charges (Gendreau 1995).
In fact, as shown in figure 5.7, there is not significant difference between segment A
and D for the improvement of the FWHM of the calibration source spectra.
These results lead us to use the CI capability to fill the traps in the transfer

60

CHAPTER 5. ONBOARD CALIBRATION OF XIS

180

FWHM (eV)

170

160

seg.A

seg.D

seg.A

seg.D

150

XIS0

XIS2

Figure 5.9: FWHM of the calibration source spectra simultaneously obtained with
the CI experiment in May 2006. The δQ parameters derived from the CI dataobtained in May (open triangles) and July 2006 (open circles) are also applied.

120

FWHM (eV)

110

100

90

seg.A

seg.B

seg.C

seg.D

80

Figure 5.10: FWHM of the CI test events spectra before (cross) and after (opentriangle) charge compensation, and that of the reference events (open circle) for
each segment of XIS3. The test events collected on ground are also shown (opensquare). Even after the compensating, the FWHM of the test events spectra arelarger than those of the reference events due to the probability process of the charge
trap.

CHAPTER 5. ONBOARD CALIBRATION OF XIS

61channel by periodically injection of Q. A ground experiment shows that charges
injected into every 54th row improve the energy resolution (Bautz et al. 2004). We
are trying to utilize this CI technique for the onboard observations. The results will
be presented in a separate paper.

5.3

Gain and Charge Transfer Efficiency

After the CTI correction, we determined the gain parameters as a function of timefor each segment. With these CTI and gain parameters, the systematic error of Kαline center energy of 55 Fe became ≤ 2% ≈ 12eV as shown in figure 5.11. The line
width is also improved after the CTI correction (figure 5.12).

CHAPTER 5. ONBOARD CALIBRATION OF XIS

Figure 5.11: The same as figure 5.1, but after the CTICI correction.

62

CHAPTER 5. ONBOARD CALIBRATION OF XIS

Figure 5.12: The same as figure 5.2, but after the CTICI correction.

63

Chapter 6
Review of Multiple Supernovae
and Superbubble
6.1

Supernova Explosion

6.1.1

Shock Waves

Shock waves play the essential roles in the dynamics in SNRs. It is a transition
region between upstream and downstream states of a fluid through which there is a
net mass flux and across which mass, momentum, and energy flux are conserved but
entropy per unit mass and mass density increase. Here we describe fluid equations
and relations between some physical parameters.

(upstream)Puρu(downstream)

Tu

Pd

ρd

Td

shock front
Figure 6.1: A schematic view around the shock front.
In a coordinate at the shock front, the conservation of mass, energy, and momentum are described as∂ρ ∂(ρv)
+
=0
∂t
∂x
64

(6.1)

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE65∂(ρv) ∂(ρv 2 + P )
+
=0
(6.2)
∂t
∂x
∂
1
∂
1
{ρ( v 2 + E)} +{P v + ρ( v 2 + E)v} = 0
(6.3)
∂t
2
∂x
2
where ρ, v, P , and E are density, velocity, pressure, and energy per unit mass,
respectively. These continuity equations can be rewritten as follows in a Lagrangian
coordinate for the situation in figure 6.1.
ρu vu = ρ d vd
2
ρu vu + Pu = ρd vd2 + Pd

(6.4)

(6.5)
1 2
1 2
vu {Pu + ρu ( vu + Eu )} = vd {Pd + ρd ( vd + Ed )}
(6.6)
2
2
where subscripts u and d denote upstream and downstream, respectively. The equation 6.6 for the ideal gas is changed to the formula, together with the equation 6.4
1 2
γ Pu
1
γ Pd
vu +
= vd2 +
2γ − 1 ρu2
γ − 1 ρd(6.7)

where γ is specific-heat ratio. From the equations 6.4, 6.5, 6.7, and the equation ofstate for ideal gas, we obtainρdvu(γ − 1)Pu + (γ + 1)Pd
=
=
ρu
vd
(γ + 1)Pu + (γ − 1)PdTu
Pu (γ − 1)Pu + (γ + 1)Pd=·.
Td
Pd (γ + 1)Pu + (γ − 1)Pd

(6.8)
(6.9)

These formulae holds between thermodynamic quantities on the two sides of the surface of discontinuity and is known as Rankine-Hugoniot relations. For the velocitiesof the gas across the shock wave, we have1{(γ − 1)Pu + (γ + 1)Pd }2ρu
1 {(γ + 1)Pu + (γ − 1)Pd }2=2ρu {(γ − 1)Pu + (γ + 1)Pd }vu2 =(6.10)

vd2(6.11)

Here we supply some useful equations for the ratio of densities, pressures, and temperatures in terms of the Mach number Mu = vu /Cu , where Cu is the sound velocity.
ρd(γ + 1)Mu2
4Mu2
=
= 2
ρu
(γ − 1)Mu2 + 2
Mu + 322γMu
γ−1
5Mu2 − 1
pd
=
−
=
pu
γ+1
γ+1
4
2
Td
{2γMu − (γ − 1)}{(γ − 1)Mu2 + 2}=
Tu(γ + 1)2 Mu2(5Mu2 − 1)(Mu2 + 3)=
16Mu2

(6.12)(6.13)(6.14)

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE66
The Mach number Md is also given in terms of Mu by
Md2 =Mu2 + 32 + (γ − 1)Mu2=
2γMu2 − (γ − 1)
5Mu2 − 1(6.15)In the case of Pu  Pd (strong shock),

ρd
vu
γ+1
=
=
ρu
vd
γ−1
s
1
Pd
vu =
(γ + 1)
2
ρu
vd =

v
u
u 1 (γ − 1)2 Pdt

2 (γ + 1)ρu

(6.16)(6.17)(6.18)

vu and vd are defined in a coordinate where the shock front has null velocity. In theobserver’s coordinate, vu equals the shock velocity (≡ vs ) . Using the equation ofstate, the post-shock velocity vp ≡ vu − vd , post-shock pressure Pp ≡ Pd , and thepost-shock temperature Tp ≡ Td are obtained as
2vs
γ+1
2ρu vs2=γ+12(γ − 1) 2µµv
= Pp =
ρd(γ + 1)2 s

vp =(6.19)

Pp(6.20)

kTp(6.21)

where µ is the mean atomic weight, and k is the Boltzmann constant. In the caseof γ = 53 (monoatomic gas), we obtain The post-shock temperature is determinedby the shock speed only.
3 2
µv
kTp =
16 s
ρd
vu
=
=4
ρu
vd

6.1.2

(6.22)(6.23)

Nuclear Synthesis

Just after the Big Bang, the universe contained only hydrogen and helium with tiny
traces of other light elements like lithium. Heavier elements were, and continue tobe, “synthesized” deep in the interiors of stars by step-by-step process of nuclear
fusion, so-called nucleosynthesis. These newly formed chemical species were spewed
out into the interstellar space by the supernova explosions, together with still heavier
atomic nuclei that were fused during the explosion.
How far the process of nucleosynthesis proceeds in a star depends on its mass.
Stars like Sun normally produce elements as heavy as carbon and oxygen, but their

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE67
densities and temperatures are too low to fuse these into heavier elements. Moremassive stars do support the required conditions to synthesize elements such as
oxygen, magnesium, silicon, and iron, as well as some heavier ones. Therefore,
different types of SNe have different progenitors, thus producing different heavy
elements on different time-scales during the chemical evolution of galaxies.
Tsujimoto et al. (1995) predicted nucleosynthesis products of Type Ia and TypeII SNe, assuming that the stars initially more massive than 10 M explode as TypeII SNe if they are single stars, or as Type Ib/Ic SNe if they belong to close binary
systems. Figure 6.2 shows the predicted relative abundances of synthesized heavy
elements by Tsujimoto et al. (1995), which are normalized to the corresponding
cosmic abundances of Tokunaga (2000). The abundance of Mg is normalized to 1.

Figure 6.2: The predicted relative abundances of synthesized heavy elements (Tsujimoto et al. 1995). The predicted relative abundances of synthesized heavy elements.
The ratios to the cosmic abundances of Tokunaga (2000) are plotted for SN Ia andSN II of 13, 18, 25, and 70M stars.

6.1.3

Evolution of Supernova Remnants

At the time of SN explosion, most of the mass of the progenitor is ejected outward
with high velocity. The evolution of the SNR depends on the ejected mass (M0 ),swept up mass (Mb ), and radiated energy. Here we describe the brief summary of
SNR evolution.
The almost all of the initial explosion energy is converted to the kinetic energyof the ejecta. Just after the explosion, the ejecta flies supersonically with respect tothe background and thus establishes the main shock that sweeps the surroundinggas. At first, the ejecta is undecelerated and the remnant is in its “free expansion”
phase (M0 > Mb ). The velocity of the ejecta is as high as v0 ∼(5 - 10)×103 km s−1 .
Soon a reverse shock, built and supported by the high pressure of the increasing
amount of swept up matter, begins to thermalize it.

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE68
When the thermalized ejecta mass and that of the swept up ISM become comparable, the radius and age of the remnant are(3M0 )(1/3)=
Rs =(4πµmH nb )(1/3)
t ≈ Rs v0−1 ≈ 190 yr.

M0
1M!1/3 

µ1.36

−1/3 n0
1 cm−3

−1/3

pc (6.24)(6.25)

Here we assumed that the mean ambient hydrogen density of n0 = 1 cm−3 , and with
M0 = 1M , v0 = 104 km s−1 , and µ = 1.36 (mean atomic weight of cosmic material
per H atom; Tokunaga 2000). The thermalized ejecta radiates X-ray, hence the
X-ray spectroscopy of this phase of SNRs can be a powerful tool to investigate thechemical composition of heavy elements synthesized by SN events.
The free expansion lasts until the reverse shock runs through the whole ejecta.
By then, the remnant would have enter the so-called “adiabatic” or “Sedov” phaseof its evolution. Shklovskii (1962) proved that the SN explosion within the ISM canbe linked to a powerful point explosion in a gas with constant heat capacity. He also
proved that the self-similar (Sedov 1959) solution was applicable to this phenomena.
However, as shown by Gull (1973), a mass comparable to ten times that of theejected matter must be swept up before the outer portions of the remnant begin toresemble the standard (Sedov 1959) solutions, and some 50 times the ejected mass
should be swept up before full thermalization of the ejecta is completed. Thus, free
expansion, thermalization and the beginning of the Sedov phase overlap in time.
Ostriker & McKee (1988) introduced the equations which describes the motion
during Sedov phase as follows:
0.2 0.4
Rs = 0.31 n−0.2
E51
t pc
0

(6.26)

dRs
2 Rs
=
.
(6.27)dt
5 t
Rs , Vs , n0 , E51 , and t are the radius, velocity of the shock front, ambient density,
initial explosion energy in the unit of 1051 ergs, and the age in the unit of year,
respectively.
With the equation (6.22), the shock temperature is given asVs ≡3
0.4 −1.2µ Vs2 = 1.7 × 104 n−0.4
E51
t
keV
(6.28)
a16
During the Sedov phase, remnants conserve energy, both thermal (Eth = 0.72E0 )and kinetic (Eth = 0.28E0 ) energy (Chevalier 1974), and grow with time as t2/5
while the leading shock velocity decays as t−3/5 .
The decay in shock speed leads eventually to a cooling time comparable to the
dynamical time, and this marks the beginning of the radiative phase which promotes the formation of a dense shell at the edge of the remnant (Cox 1972). At
this moment, the shock velocity is decelerated down to ∼ 200 km s−1 and the gas
temperature is about kT ∼ (4 − 5) × 10−2 keV. By this time, about half of thethermal energy has been radiated away. In the early radiative stages, the interior ofkTs =CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE69
the remnant is still filled with the hot ejecta and the hot ISM, shocked earlier in theevolution, and both do work on the outer shock front for as long as the pressure iscontinuous across the remnant and with the same value across the dense shell. According to Falle (1981) and Falle (1988), the age tcool , radius Rcool , and expansion
velocity Vcool of SNR in the radiative cooling phase are expressed as,
0.24 −0.52tcool = 2.7 × 104 E51
n0yr,
0.295 −0.409
Rcool = 20E51 n0pc,
0.0554 0.111
Vcool = 277E51 n0
km s−1 .

(6.29)(6.30)(6.31)The remnant is then in the “pressure-driven snowplow” phase (Oort 1951). The
law of motion governing the shell at this stage can be derived from the adiabatic
expansion condition imposed on the interior hot gas,dRs
dET
= −4πRs P
,
dt
dt

4π 3
R P = (γ − 1)ET
3 s
and the equation for the mass and momentum of the shell
4π 3
R ρ0 ,
3 s

M =(6.32)(6.33)(6.34)d(M Vs )= 4πRs2 P,(6.35)dt
dRs= Vs .
(6.36)
dt
Here P is the pressure of the hot gas, the thickness of the shell is assumed to besmall compared to its radius, and the mass of the hot gas is assumed to be much lessthan that of the shell. With equation(6.33) and equation(6.36), McKee & Ostriker(1977) and Blinnikov et al. (1982) got a solution as follows,−5/21

Rs = 38(E51 )5/21 n0(t
105 yr)2/7 pc,(6.37)

2 Rs,(6.38)
7 t
where  = ET /E0 = 0.2 - 0.35.
Once the interior pressure becomes negligible, the remnant evolves as a pure“momentum-driven snowplow”, its radius then grow as t1/4 , and the shock speed
decays as t−3/4 , even faster than in the Sedov phase. When the expansion velocity is
as small as ∼ 10 - 20 km s−1 , which is about the same as the velocity of proper motionof the ambient interstellar matter, the remnant begins to melt into the ISM. If theremnant structure does not interact with other ones, it will eventually disappear
by fragmentation of the shell and by replacement of the cooling low-density gas byexpanding cooler material from the ISM (Cox & Smith 1974).
Vs =CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE70

Figure 6.3: The evolution of SNRs. Schematic representation of the main structuralchanges adopted by the shocked gases as a result of a sudden release of energy(E ∼ 1051 ergs) into a uniform density medium (n ∼ 1 cm−3 ). This figure is adopted
from Tenorio-Tagle (1996).

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE71

6.2

Superbubble

6.2.1

Formation of Superbubble

Young massive stars which have strong winds and explode as type Ib/Ic or II SN
are known to be spatially and temporally correlated with OB associations. Massive
stars are formed in groups by the collapse of molecular clouds. Because of theirshort lifetimes of 3 - 20 Myr (Schaller et al. 1992) and small dispersion velocities of
4 - 6 km s−1 (Blaauw 1991; Mel’nik & Efremov 1995), SN progenitors remain andexplode close to their birth place. Although identifying the membership of a given
OB association is not an easy task from the observational point of view, it is reliably
estimated that between 60% (Garmany 1994) and 95% (Higdon et al. 1998) of all
OB stars belong to such associations, which contains up to several OB stars withinregions of radius ROB ∼ 35 pc (Garmany 1994; Bresolin et al. 1999; Pietrzyński
et al. 2001, and references therein).
For evenly distributed stars, the mean distance between two closest neighbours
3
can be evaluated as Ddis ' (4πROB/3NOB )1/3 , so that each star can be consideredas occupying an individual spherical volume of radius Rdis ' Ddis /2, with typicalvalue:
Rdis ' 6 pc

ROB
35pc

!

NOB
100

!1/3

.

(6.39)

An important characteristic of such stars is that they experience strong winds during
most of their lifetime. The mass-loss rate and the wind velocity, and thus the wind
power, are not constant during stellar evolution (Meynet et al. 1994), but the total
wind energy, integrated over a massive star’s lifetime, amounts typically to 1051 erg
and is therefore comparable to the final SN explosion energy itself. When considering
the energy output of OB stars in the Galaxy, one thus has to include the contributionof the winds, which can roughly double the energy imparted to cosmic rays if thewind energy can somehow be used to accelerate particles.
After a few (≤5) explosions and/or strong stellar winds, once the remnant has
grown to a large radius, all subsequent explosions lead to blast waves which, by
evolving into the hot gas produced by former SNe decay into sound waves (TenorioTagle et al. 1987; Mac Low & McCray 1988; De Young & Gallagher 1990; TenorioTagle 1996). In this way, the kinetic energy deposited by the stellar cluster isfully thermalized without disrupting the pattern shown in figure 6.3. Following SN
explosions may occur in the cavity formed by preceding ones and/or strong stellar
winds from OB-type stars. As a result, the correlated core-collapsed SNe explosions
create giant cavities of hot, tenuous plasma, which is called SBs, rather than manysmaller isolated remnant bubbles (McCray & Snow 1979).
Higdon & Lingenfelter (2005) determined quantitatively the fraction of Galactic
core-collapse SN that occur in SBs. The fraction of core-collapse SN (Type II andIb/c) occurring in SBs is high, ranging from ∼80% (solely temporal correlations) to∼90% (only spatial sorrelations). Core-collapse SNe constitute 85% of Galacic SNe(van den Bergh & McClure 1994); only a small fraction of the remaining SNe, TypeCHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE72
Ia, occur in SBs. Thus, ∼75% of all Galactic SNe are expected to occur within SBs.
The deficiency of bright Galactic radio SNR shells is well-known issue (Clark &Caswell 1976). This can be most simply seen from the fact that there are only a few
SNR shells (Green 2004) with radio surface brightness >1% of that of the brightest
remnant, the ∼300 yr old Cas A. From various models (Berezhko & Völk 2004) ofradio emission in the Sedov phase of SN expansion, the radio surface brightness of
Cas A would be expected to remain above 1% of its present value for another 2000
- 4000 yr, so with a Galactic SN rate of one every 40 yr, we could expect to findbetween 50 and 100 radio SNR at least that bright. This is more than 10 times
that observed. Higdon & Lingenfelter (1980) suggest that this discrepancy resultsfrom the fact that most SN occur in the hot, tenuous (∼3×10−3 cm−3 ) phase of
the ISM, where their remnants have to expand to much larger volumes than in thedenser phases of the ISM to sweep up enough matter to enter the Sedov phase. As
a result, the peak surface brightness of remnants is the hot, tenuous phase is much
lower than that of those in the denser phases (Berezhko & Völk 2004). Koo & Kang
(2004) have recently suggested that the same effect may account for the deficiencyof observed HI SNRs in the inner Galaxy.

6.2.2

Evolution of Superbubble

The theories of SB evolution are based on the standard wind bubble theory (Weaver
et al. 1977). In order to grasp the typical values of basic physical quantities inside
the SB, here we follow Mac Low & McCray (1988) and assume that the SB isexpanding in a homogeneous medium of density n0 , powered by the activity of an
OB association providing a constant mechanical luminosity of LOB . Although theenergy release is not continuous and experience strong peaks when an OB star enters
the Wolf-Rayet stage or when a new SN explodes, the variation of the driving power
are smoothed out for sufficiently evolved SBs because the shells of individual SN
shocks become subsonic before they reach the shell. In addition, star formation is asequential process in giant molecular clouds, and massive stars have a whole range of
lifetimes. Hence as a first approximation, it can be assumed that the energy releaseis indeed roughly constant inside the SB, which allow us to treat the whole SB as avery large wind bubble.
With the above assumption, Mac Low & McCray (1988) follow Weaver et al.
(1977) and Castor et al. (1975) to find the temperature and density inside the SB:
8/35

Tr = (3.1 × 106 K) L37= Tc (1 − x)2/5

n0

2/35

t6

−6/35

× (1 − x)2/5(6.40)

6/35

n0

19/35

t6

−22/35(6.41)

nr = (1.1 × 10−2 cm−3 ) L37= Tc (1 − x)−2/5

× (1 − x)−2/5

where the n0 is the external ISM density in cm−3 , t6 is the age of the SB in units of
106 yr, x = r/R is the relative distance from the SB center, and (1−x)2/5 is a function
giving the temperature and density gradient inside the SB, assuming that the energy

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE73
is injected at x = 0. The internal pressure, i.e. the product PSB = µnSB kB TSB , which
independent of (1 − x)2/5 as the sound crossing time is lower than the SB dynamical
time. Here we assume a particle multiplicity µ ' 2.3, taking into account thecontribution of the electrons to the pressure (the SB interior is here assumed to be
fully ionized, with solar abundances).
The SB internal pressure thus reads:
2/5

3/5

PSB = (4.3 × 10−12 dyne cm−2 ) L37 n0

−4/5

t7

(6.42)

Given the PSB , one can calculate the radius of a stellar wind terminal shock,around a given OB stars with a wind power (Lw ) and a wind velocity (Vw ) byconsidering the condition of PSB is equal to the wind ram pressure (Pram ). Then one
obtain1/2−1/2

Rterm = (20 pc) Lw,37 Vw,3

−1/5

−3/10

LOB,38 n02/5

t7(6.43)

This radius can be larger than R, hence the direct wind-wind interaction should
occurs inside SBs. Due to the low pressure inside the SB, the region containing
unshocked wind material extends far enough around the star so that it may enter
directly into contact with the unshocked wind material of another star, that is, thewinds actually terminate each other, instead of being terminated by the surrounding
medium. Then, the wind energy is efficiently converted into strong turbulence,and since the material is fully ionized, plasma waves should also rapidly develop
and produce a magnetic turbulence with values of the magnetic field close to theequipartition value.
We now discuss in what differs the revolution of SNRs in SBs from those which
explode in the average ISM, that is, isolated SNe.
A major difference is the fact that the density and temperature of the internal
cavity are modified by subsequent explosions, and soon achieve values typical of a“coronal medium” (n = 102 - 103 cm−3 , T > 106 K). Thus, all SN exploding after
the first one at the center of the evolving remnant encounter an increasingly hotter
and rarefied medium. The lower density results in a quicker growth of the SNR shell,
and a higher temperature in a higher sound speed. The transition to a Sedov-like
expansion occurs when a mass of roughly 1.6 Meje is swept-up by the ejecta (McKee& Truelove 1995). In a medium of density 4 × 10−3 cm−3 , which corresponds tothe case of typical SB after 107 yr of evolution, this occurs when the SNR reaches
5/6 0.21 −1/2
5/6
4
after the explosion
a radius R0 ' 30 pcM10 t0.21
7 , i.e. ∼ 1.3 × 10 yrsM10 t7 E51(Truelove & McKee 1999).
In the Sedov-like phase, the SNR expands almost selfsimilarly, as from a point
explosion, according to:
2/5

22/175RSNR = (38 pc) tSNR,4 tSB,7(6.44)

and−3/5

22/175VSNR ' (1470 km s−1 ) tSNR,4 tSB,7(6.45)

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE74
where we have replaced the ambient gas density by that of the SB interior given byequation 6.41 and used the values corresponding to typical SB, and where tSB,7 isthe age of the SB in units of 107 yr and tSNR,4 the age of the SNR in units of 104 yr.
The above equations allow one to calculate the time when the shock becomessubsonic. Replacing the qtemperatureq TSB , from equation 6.40, in the expression of
the sound velocity, cs = γp/ρ ' γkT /(1.4mp ) ' 99km s−1 (T /106 K)1/2 for γ=
3/5, one finds that VSNR > cs until
37/105tsub ' (3.1 × 105 yr) tSB,7 .

(6.46)

By that time, the SNR has reached a radius
4/15

Rsub ' (150 pc) tSB,7 .

(6.47)

which can be compared to the radius of the SB itself:
1/5

−1/5

RSB ' (267 pc) LOB,38 n0

3/5

t7 .

(6.48)

Rsub < RSB is satisfied even in the very early phase in the evolution of a SB (i.e. for−1/3
the very first SN), and that Rsub /RSB still decrease as 0.56t7
as the SB evolves.
This justified the statement above that the discrete energy release inside the SB are
actually smoothed out and the growth of the SB can be worked out by assuming acontinuous driving power.
Another characteristics is that the shell of a SN exploding inside a SB becomessubsonic before becoming radiative phase. Indeed, evaluating the cooling time ofthe shocked gas compared to the age of the SNR, Blondin et al. (1998) obtained thetimescale for the end of the Sedov-like phase and the formation of a radiative shell
4/17 −9/17. Replacing n0 by the SB density and dividing by
as trad = 2.9 × 104 yrE51 n0
the time corresponding to the sonic transition, one find that
−1/51

trad /tsub ' 1.7 × t7

(6.49)is always larger than 1, indicating that the SNR will never become radiative
inside the SB. One may therefore expect that SN shock waves remain in the Sedovlike phase (and thus keep its initial energy ) until they die well inside the SB.
The radii of SBs formed as mentioned above range from those of ordinary SNR(<100 pc) to more than 1000 pc, and kinetic energies of expansion are up to 1054
ergs.

6.2.3

Cosmic Ray Acceleration in Superbubble

The question over the origin of cosmic rays (CRs) has persisted ever since they
were first detected in 1912 by Victor Hess using a balloon-borne electroscope device(Hess 1912). Recently the spectra of CRs up to energies of 105 TeV were taken inballoon measurements by Apanasenko et al. (2001). These are power-law spectra
without any features in this energy spectrum at energies E = 102 - 105 GeV per
nucleon is approximately constant and close to γ ≈ 2.75. The spectrum of heavy

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE75

Figure 6.4: The CR spectrum for all nuclei, as measured with various facilities,
and the expected behavior of the proton and iron spectra (Kampert et al. 2001).
dj/dE0 is the differential CR intensity related to the distribution function F (E0 ) by
dj/dE0 = cF (E0 ), where c is the speed of light.
nuclei in the same energy range is flatter, and its index appears to be close to γ ≈2.65. These results were obtained both from observations of extended air showers
and by direct measurements. At energies E = 105 - 107 GeV per nucleus, there isa feature (apparently a kink or a knee) in the CR energy spectrum summed over
all nuclei. The shape of the feature has not been studied adequately. After theknee, the spectrum is again close to a power law with γ ≈ 3.1 - 3.2 at energies
E ≈ 107 - 109.5 GeV per nucleus (Bird et al. 1993). The spectrum breaks at energy
E ≈ 109.5 GeV per nucleus, and γ ≈ 2.7 at E ≈ 109.5 - 1010.6 GeV per nucleus (Birdet al. 1993).
In the energy band of E ≤ 1010.6 GeV, the gyro radius due to the Galacticmagnetic field is significantly smaller than the thickness of the Galactic disk, andCR is blocked up in the Galactic disk changing its direction. In the case E > Eknee ,it is hard to bound CR inside the disk and they leak outside the disk, although
the firm confirmation has not been obtained yet. In the very-high energy band ofE ≥ 1010 GeV the gyro radius is almost the same or larger than the Galactic disk,hence it is hard to bound the CR in the disk. Because the CR of E ≥ 1010 GeV show
almost uniform distribution, the origin of them seems to be extragalactic objects.
There is now broad consensus that the shocks in the expanding blast waves of SNRs
accelerate the bulk of the CR ions up to energies of ∼300 TeV nucleon−1 .
In fact, the discovery of power-law (non-thermal) X-rays from SN 1006 (Koyama
et al. 1995) indicates the existence of extreme high energy charged particles by the
first order Fermi process. This lead a successive discovery of new class of shell-likeCHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE76
SNRs such as RCW 86 (Bamba et al. 2000; Borkowski et al. 2001b), G266.2–1.2(Slane et al. 2001), and G28.6–0.1 (Ueno et al. 2003). Some of the Galactic SNRs
are also observed with γ-rays. RX J1713.7–3946 has proven to be a source of nonthermal X-rays, GeV γ-rays, and TeV γ-rays. The CANGAROO collaboration
detected TeV γ-ray emission from the northwest rim, and the spectrum taken in the
γ-ray band is consistent with π 0 -decay γ-rays (Enomoto et al. 2002), which indicates
SNR molecular cloud interaction. A molecular cloud has indeed been found at thenorthern perimeter by CO observation (Butt et al. 2001). Another SNR, γ Cygni,
is associated with one of the brightest EGRET sources, and is also interacting witha dense cloud (Uchiyama et al. 2002)
However, according to the most optimistic estimates, the acceleration by single
shock waves from SN explosions cannot account for the origin of the particles above
the knee in the energy spectrum. Hence we take notice of SBs as potential candidatesof the accelerator above the knee.
Although the diffusive shock acceleration mechanism is not fundamentally modified inside a SB, there are some important differences compared to the isolated
SNRs.
First, the presence of a turbulent and magnetized medium in front of the shockexits. Recent study of diffusive shock acceleration has shown that the efficiency ofparticle acceleration around strong shocks and the maximum energy (Emax ) cruciallydepend on the level of turbulence and the value of the magnetic field (Lucke & Hodge
1970; Berezhko et al. 2003; Ptuskin & Zirakashvili 2003). Since the size of SNRs
and the acceleration duration is limited, large values of Emax require low diffusion
coefficient. A lower limit to the diffusion coefficient along magnetic field lines isprovided by the so-called Bohm scaling, where DB = 31 vrg and rg = γmv/qB is
the gyro radius of the particle of mass m, charge q = Ze and Lorentz factor γ in−1
a field of strength B. These give DB ' 3.1 × 1022 γβ 2 Z −1 BµG
cm2 s−1 . To lower
this value, and thus increase Emax , one needs larger magnetic fields. Although
direct measurements of the magnetic fields inside SBs are not available, one canestimate that it is indeed larger than in the average ISM, due to the mechanisms
such as turbulence generation through direct wind-wind interactions. Assuming
the equipartition of the mechanical energy released by the massive stars between
thermal pressure, turbulence and magnetic fields, one can obtain magnetic fields ofthe order of 10 - 20 µG.
Furthermore, the turbulence flow inside the SB should also make the acceleration
efficiency higher than an isolated SNR. In the isolated SNR, particle injection in theacceleration process is provided by the tail of the thermal distribution (Völk et al.
2003), which limits the fraction of particles flowing through the shock front to beeventually accelerated to about 10−4 or 10−3 at most. In the case of an SNR ina SB, all the pre-existing energetic particles passing through the shock will see thediscontinuity and be able to gain energy by diffusing back and forth across the shock
front.
Assuming an amplification factor αB , one can roughly estimate the maximum
energy obtained from standard diffusive acceleration inside SBs by following Ptuskin& Zirakashvili (2003) as

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE77

Emax ' (1.7 × 10

17αBeV) × Z
20

!

!

BSB
.
10µG

(6.50)We see that values of Emax of the order of Z × 1017 eV (as would be required in orderto reach the ankle of the CR energy distribution), require values of αB of the orderof 10 - 20. This corresponds to enhanced values of the magnetic field at the shockof the order of 100 - 200 µG, which is unreasonable compared to what is usuallyassumed in isolated SNRs.
According to Bykov & Fleishman (1992), CR acceleration time is estimated to
be about 3 × 105 yr and the maximum energy should exceed 3 × 1014 eV.
Because the lifetime of an SN blast wave is of the same order as the time scale
between two explosions (∼ 3 × 105 yr), the number of coexisting primary SN shocks
inside an SB must be small (≤2). Then multiple strong spherical SN shocks simultaneously in one SB would require an extreme SN rate in SB. However, repeatedshock acceleration is possible and must actually occur for relatively low-energy particle. If the time required for them to leave the SB is larger than the typical time
between two SN explosions, they may be overcome by a subsequent shock, and thus
be injected into a new Diffusive Sshock Acceleration (DSA) site.
The typical escape time is given by τesc ∼ R2 /2D, where D(E) is the average
diffusion coefficient in the SB. Casse et al. (2002) showed that variation law of the
diffusion coefficient with rigidity depends on the ratio of the particle gyro radius to
the principal length scale of the turbulence. The diffusion coefficient is estimatedfor a turbulent length scale of the order of the typical distance between massive
stars. For the characteristics of typical OB association, one obtains DE ' (1.0 ×
1027 cm2 s−1 .
Then maximum energy acquired by the repeated shock acceleration can be estimated to be
Emax ' 11 TeV × ηT

BSB
10µG

!

λmax
6pc

!−2

,

(6.51)

where λmax is the principal length of the turbulence.

6.2.4

X-ray emission from SBs

To find LX , we integrate the X-ray emissivity ΛX in the layer of shocked stellar wind:
LX =Z

n(r)2 ΛX [T (r)] d3 r

(6.52)For the temperature range of 106 −107 K, then emissivity within the soft energy bandof <2 keV is virtually independent of temperature, and ΛX (T ) can be approximatedby a constant, 3×10−23 ergscm3 s−1 for a solar abundance (Raymond & Smith 1977).
On the context of equations 6.40 and 6.41, the radius of the bubble is
1/5 −1/5 3/5
t6 .

R = (42 pc) L37 n0(6.53)

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE78
The X-ray luminosity can be driven by integrating equation 6.52 from the center
to a radius Rmax , while the temperature drops to the minimum that will emit in
the soft-energy band, Tmin . Defining a dimensionless temperature τ = Tmin /Tc , the
X-ray luminosity can be expressed as33/35 17/35 19/35
n0 t 6 ,LX = (1.1 × 1035 ergs s−1 ) ξI(τ )L37

(6.54)

where the dimensionless integral is I(τ ) = (125/33) − 5τ 1/2 + (5/3)τ 3 − (5/11)τ 11/2 ,and the metallicity relative to the Galactic (or solar) value is ξ. In scaled units−8/35 −2/35 6/35
here, τ = 0.16 L37 n0
t6 , so I(τ ) ∼ 2. For the practical use, the X-ray
emission should be expressed in physical parameters that are observable, such assize, expansion velocity, and density. The time derivative of equation 6.53 gives
expansion velocity Vexp ,Vexp = (0.59 km s−1 ) Rpc /t6 ,(6.55)

where the Rpc is the radius in units of pc. Combining equations 6.53, 6.54, and 6.55,
10/7

LX = (8.2 × 1027 ergs s−1 ) ξI(τ )n0

16/7

17/7
Rpc
V5

,

(6.56)

where V5 is the expansion velocity in units of km s−1 . The ambient density n0 cannot
be measured directly, but it can be determined indirectly from the electron density
in the shell, ne . For a pressure-driven SB, the ambient density is related to the shelldensity by2n0 = (9/7) ne kTe /(µa Vexp),

(6.57)

where Te is the electron temperature in the shell and µa = (14/11)mH (Weaver et al.
1977; Chu & Mac Low 1990). The shell electron density ne can be approximatedby the rms electron density ne (rms) determined from the Hα surface brightness.
For a shell of thickness ∆R, measured in pc,
ne = hn2e i =EM,
2Rpc [2(∆R/Rpc ) − (∆R/Rpc )2 ]1/2

(6.58)

where EM is the peak emission measure in units of cm−6 pc. It may be further
assumed that the interstellar matter in the volume within radius R has been swept
into a shell of thickness ∆R, hence
∆R
1 n0
=
.
(6.59)
R
3 ne
If the expansion velocity, electron temperature, and peak emission measure of a
SB are known, shell thickness and ambient density may be derived from equations 6.57, 6.58, and 6.59. equation 6.54 indicates that LX grows as a bubble ages.
However, detectability depends on the surface brightness, which is proportional to19/35 31/35 −23/35
L37 n0 t6. Hence, the detectability actually diminishes as a bubble ages.

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE79

6.2.5

Past Observations

The observations of SBs have been concentrated on the Large Magellanic Cloud(LMC) because the LMC is an ideal laboratory for the observation of SBs. Because
we know the distance to the LMC as 50 kpc with small uncertainty (Feast 1999), weare able to derive diameter of SBs (1 pc = 4 arcsec). There are many extended HII
regions of 75 - 150 pc in size connected with OB associations in the LMC (McGeeet al. 1978; Mathewson et al. 1985), which have the radio-bright emission. A weak
linear polarization has been detected in the radio emission from 30 Dor nebula in the
LMC (Haynes et al. 1990). Hence it can be said that the radio fluxes are producedby the synchrotron emission of non-thermal electrons. Calculated synchrotron flux
fits the observed one if Ne ∼ 10−11 cm−3 (Bykov & Fleishman 1992). Here theenergy density of the non-thermal electrons appears to be equal to a few percent of
the kinetic energy density of the system.
An excess of diffuse X-ray emission in SBs indicates the presence of interior supernova remnants shocking the inner walls of the SB shell (Chu & Mac Low 1990;
Wang & Helfand 1991; Dunne et al. 2001). In Chu & Mac Low (1990), Einsteinobservations were used to show that seven LMC SBs are diffuse X-ray sources withluminosities much higher than those expected by the wind-blown, pressure-driven
bubble models of Weaver et al. (1977). Off-center SNRs are proposed to be responsible for the excess X-ray emission. ROSAT observations of the SB N44 confirmed its
diffuse X-ray emission and provided the first useful X-ray spectra of N44 for determinations of plasma temperatures (Chu et al. 1993). To illustrate that excess X-ray
emission from SBs is caused by an intermittent process, Chu et al. (1995) analyzed
the ROSAT observations of four X-ray–dim SBs and showed that these SBs do nothave excess X-ray emission. However, comparing the observed X-ray luminosities
to the predicted luminosities, they found that the observed values are surprisingly
similar to the theoretical predictions. For high-resolution spectral analysis, ASCA
observations of N44 were made. The ASCA data showed that the hot gas in the
breakout region is slightly cooler than that in the SB interior (Magnier et al. 1996).
ROSAT observations of the HII complex N11 were analyzed to study the interaction
between OB associations, HII regions, and SBs (?). In the latest study, Dunne et al.
(2001) analyzed ROSAT observations of eleven Hα-identified SBs in the LMC.
In this study, the derived X-ray luminosities of the SB were compared with thepressure-driven bubble model of Weaver et al. (1977). Assuming that the shell thickness is small compared to the radius of the bubble, that the electron temperaturein the shell is Te ∼ 104 K, and that the mean atomic mass of the ambient mediumµa = (14/11) mH , the X-ray luminosity can be derived from the equations 6.56 6.59 as follows:
12/7 1/7
LX ' (6.7 × 1029 ergs s−1 ) EM5/7 ξIRpc
vexp ,(6.60)

where the vexp is the expansion velocity in km s−1 . They used the vexp fromthe Hα broad line spectra, EM from the continuum-subtracted Hα images. The
predicted X-ray luminosities range from 1034.3 to 1035.1 ergs s−1 . These luminosities
range from ∼3 to ∼50 times lower than the X-ray luminosities determined fromCHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE80

Table 6.1: Best-Fit LMC SB X-Ray Spectral Fits.
log NHkT

log LX †√Ne f(cm−2 )

(keV)

(ergs s−1 )(10−2 cm−3 )N11 shell 1...

21.9

0.13

37.17

31.

N44...

20.6

0.71

35.73

3.8

N51 DEM L 192...

20.4

0.28

35.18

1.5

N51 DEM L 205...

20.5

0.25

34.86

3.7

N57...

20.6

0.33

35.32

1.8

N103B...

21.3

0.72

34.85

0.8

N105...

21.6

0.31

35.40

4.6

N144...

20.5

0.32

34.74

1.3

N154...

20.7

0.29

35.52

1.8

N158...

20.7

0.60

35.85

3.0

N160...

20.9

3.40

35.80

1.9

N206...

20.7

0.31

35.09

2.2

30 Dor C East...

21.1

1.19

35.49

4.0

30 Dor C est...

22.0

1.22

35.84

7.4

Object†: In the energy band of 0.5 - 2.4 keV.

the ROSAT data. This suggests that the majority of the X-rays are producedby mechanisms different than those described in the pressure-driven bubble model,
confirming that the SBs in their sample are X-ray bright. They also compared the
SB volume, Hα luminosity, expansion velocity, and and bright star count (based on
the OB association star counts in Lucke & Hodge (1970)) with the X-ray luminosity.
The most significant positive correlation was found between the X-ray luminosityand bright star count, although the moderate positive correlation could be seen also
in other parameters. This correlation demonstrate that the X-ray luminosity of a
SB is affected by the richness and age of OB associations within its shell walls.
Recently the some observations of SBs with Chandra and XM M -N ewton havebeen done. The X-ray bright SBs in the LMC, N51D, and N11 were observed with
XM M -N ewton (Nazé et al. 2004). In addition to the X-ray emission from massive
stars, possible hidden SNRs, colliding-wind binaries, numerous pre-main sequencestars of some dense clusters inside the SB, they found that the SB may be leaking
some hot gas in the ISM. Cooper et al. (2004) observed the N51D, which is oneof the X-ray brightest SBs in the LMC and from which the emission line featuresin the X-ray spectra suggest an overabundance in oxygen and neon which is usedargue for the occurrence of SN explosion (Bomans et al. 2003). The X-ray spectrumof the diffuse emission from N51D requires a power-law component to explain theCHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE81

Figure 6.5: SB X-ray luminosity vs. bright star count for the best-fit NH X-ray
spectral fits (top) and fixed NH X-ray spectral fits (bottom). Luminosity is plottedon a logarithmic scale. Bright star counts are based on the star counts in local OB
associations (Lucke & Hodge 1970).

CHAPTER 6. REVIEW OF MULTIPLE SUPERNOVAE AND SUPERBUBBLE82
featureless emission at 1.0 - 3.0 keV. The origin of this power-law component isunclear, but it may be responsible for the discrepancy between the stellar energy
input and the observed interstellar energy in N51D.

Chapter 7
Observations and Results
Here we describe the results on the analyses of SBs and/or SB formation sites.

7.1

30 Dor C

First we describe the X-rays observation of the SB 30 Dor C in the LMC. As we mentioned in section 6.2.5, the LMC provides a sample of numerous SBs at a common
distance that is resolvable by modern X-ray detectors. For this study, we singled
out 30 Dor C as being the best target for our research. A radio source was discovered by Le Marne (1968) southwest of 30 Dor and was named 30 Dor C, which
is now categorized as an SB. Mills et al. (1984) found a shell-like structure in the
843 MHz band observation with a radius of about 30 (40 pc on 50 kpc distance).
Along the radio shell, complex Hα, Hβ, and [S II] emissions were found, which arebright in the southeast but dim in the other side (Mathewson et al. 1985). In the
X-ray band, Einstein Observatory detected the eastern shell of the 30 Dor region
for the first time (Long et al. 1981; Chu & Mac Low 1990). Dunne et al. (2001)
reported that the ROSAT spectrum required a thermal model with a rather hightemperature (1 keV). Itoh et al. (2001) found non-thermal X-rays in the ASCA data
but could not spatially resolve the thermal and non-thermal components in the SB.
Recently, Dennerl et al. (2001) found a complete ring in soft and hard X-rays witha diameter of 60 with XM M -N ewton. The hard X-rays from the shell resemble thesynchrotron X-rays in the shock front of SNRs, which is a site of cosmic-ray acceleration. Therefore, 30 Dor C would be the first and best candidate of cosmic-ray
accelerating SBs.
Chandra observed SN 1987A several times with the ACIS-S array. Some of the
SN 1987A observations covered the 30 Dor C region by chance, which is locatedat about 50 northeast from SN 1987A. In order to study the diffuse structure of
30 Dor C, we selected two observations in which the grating instrument is removed
from the X-ray mirror axis (observation IDs [Obs. IDs] 1044 and 1967; hereafter Obs.
1 and 2). The observed dates and targeted positions are given in table 7.1. Dataacquisition from the ACIS was made in the timed-exposure faint mode with a readout
time of 3.24 s in both observations. The data reductions and analyses were madeusing the Chandra Interactive Analysis of Observations (CIAO) software, version
83

84

CHAPTER 7. OBSERVATIONS AND RESULTS

2.3. For level-1 processed events provided by a pipeline processing at the Chandra
X-Ray Center, we made a charge transfer inefficiency correction and selected ASCA
grades 0, 2, 3, 4, and 6 as the X-ray events. High-energy events due to charged
particles and hot and flickering pixels were removed. There are some ”streaks” in
the CCD chip S4, which are caused by a random deposition of a significant amountof charges in pixels along the row during the serial readout process. These ”streak”
events were removed by using the program destreak in CIAO. The total available
times for each observation, after the screening, are listed in table 7.1.
Table 7.1: Observation Log around 30 Dor C

Obs.

Satellite

Obs. ID

R.A.

Decl.

(J2000.0)(J2000.0)

Date

Exposure
(ks)

1...

Chandra1044

05 32 22.0 -69 16 33.0

2001 Apr 25

18

2...

Chandra1967

05 32 22.0 -69 16 33.0

2000 Dec 07

99

3...

XM M -N ewton

0104660301 05 35 28.0 -69 16 11.0 2000 Nov 25

21

4...

XM M -N ewton

0113020201 05 37 47.6 -69 10 20.0 2001 Nov 19

16

Note. Units of right ascension are hours, minutes, and seconds, and units of declination aredegrees, arcminutes, and arcseconds.

XM M -N ewton also observed SN 1987A and the 30 Dor region several times.
We selected two observations (Obs. IDs 0104660301 and 0113020201, hereafterObs. 3 and 4), which cover 30 Dor C and are relatively free from high background
flares because of low-energy protons. The observed dates and targeted positions areshown in table 7.1. In both observations, only the MOS CCDs were operated in the
full-frame mode with the medium filter (Stephan et al. 1996; Villa et al. 1998) forblocking ultraviolet photons. The data reductions and analyses were made using the
XM M -N ewton Standard Analysis System (SAS; Watson et al. 2001) version 5.4.1;
we performed the basic pipeline process following the SAS guide. The backgroundlevel was changed significantly, particularly in Obs. 4. Hence, we removed the datawith a high background level (>0.6 counts s−1 in the 10.0 - 15.0 keV band). The
exposure times in each observation after the screenings are listed in table 7.1.
Figure 7.1 shows the soft (0.7 - 2.0 keV) and hard (2.0 - 7.0 keV) band Chandra
images around 30 Dor C, in which the two observations (Obs. 1 and 2) are combinedwith a correction of the exposure times. A clear shell-like structure with a radius of∼17000 (R ∼ 40 pc radius at the 50 kpc distance) is seen in both bands. In detail,however, the morphologies differ from each other; the entire shell is seen in the soft
band, whereas the hard X-rays are visible only at the western part. Cataloged SNRs,
the Honeycomb nebula (SNR 0536–69.3), and SN 1987A are also seen mainly in thesoft band (see figure 7.1). The X-ray features of these objects have been reportedwith XM M -N ewton (Dennerl et al. 2001) and Chandra (Burrows et al. 2000; Parket al. 2002; Michael et al. 2002) monitoring observations.

CHAPTER 7. OBSERVATIONS AND RESULTS

85

Figure 7.1: Chandra ACIS images around 30 Dor C in the (a) 0.7 - 2.0 keV bandand (b) 2.0 - 7.0 keV band, with J2000.0 coordinates. The scales are logarithmic,
with units of ×10−5 counts s−1 arcmin−2 cm−2 , as shown in the left bar for each
image. The source regions used for the spectral analyses are shown by the solidlines in the left panel (figure 7.2, tables 7.3 and 7.4). The plus signs represent theposition of the detected point sources shown in table 7.2. The position of SN 1987A(Burrows et al. 2000; Park et al. 2002; Michael et al. 2002) and the Honeycomb
nebula (Dennerl et al. 2001) are also shown.
In addition to the diffuse structure, some pointlike sources are found inside
30 Dor C. The exposure times of the XM M -N ewton observations are short, most of
the data suffer from a high background, and the spatial resolution is not sufficient.
Hence, we concentrated on the Chandra data for the point-source search and analysis. At first, point sources were searched for with the wavdetect software in the 0.5- 8.0 keV band images, then manually inspected for any spurious pointlike structuredue mainly to a part of the diffuse emission. We thus found six point sources witha significance level of greater than 7.0 σ, as shown in figure 7.1 and table 7.2. For
these six point sources, we searched for optical, infrared, and radio counterparts andfound that three (sources 1, 3, and 4) coincide at the positions of the brightest starclusters: α, β, and γ (Lortet & Testor 1984). We therefore checked further for any
X-ray emission from the other clusters (δ, , and ζ), which are also members of the
OB association LH 90 (Haberl & Pietsch 1999; Sasaki et al. 2000) but found no
candidate.
In the Chandra observations, we can see that the diffuse structure consists ofseveral shell fragments (see figure 7.1), and the whole structure is widely spreadover the two observed regions with different configurations of the CCD types (BIand FI). Therefore, a spectral analysis on all of the diffuse structure is technically andscientifically complicated. For this reason, we divided the diffuse structure into four
regions (hereafter shells A - D), as shown in figure 7.1), and performed the spectral
analysis separately for each shell. We excluded all of the detected point sources (see7.1) from each shell. The background regions were also selected separately for each

86

CHAPTER 7. OBSERVATIONS AND RESULTS

Table 7.2: Point Source Data around 30 Dor CNum.

CXO J05

SMa

Γ/kTb/(keV)

1...

3542.4-691152

MK

2.1 (1.4 - 3.6)

Nb,c

Fluxd

χ2 /d.o.f.

(×1022 Hcm−2 ) (ergs cm−2 s−1 )3.5 (2.5 - 5.1)

2...

3542.9-691206

PL

2.1 (1.4 - 3.2)

7.7 (4.4 - 13)3...

3559.9-691150

PL

1.7 (1.3 - 2.6)

0.13 (<1.3)

4...

3606.6-691147

MK

1.0 (0.19 - 1.3)

0 (<1.9)5...

3620.7-691303

PL

1.9 (1.5 - 2.4)

0 (<0.22)

6...

3633.3-691140

PL

1.8 (1.5 - 2.3)

0.5 (0.08 - 1.1)

1.8 ×10−14

9.53/12−14

8.25/13

6.5 ×10−15

4.03/7−15

19.8/20

2.0 ×10−14

21.4/18

2.0 ×10

2.8 ×10−15
8.9 ×10Notes. Sources 1, 3, and 4 have the optical counterparts, the stellar clusters α, β, and γ, respectively (Lortet & Testor 1984).
a: Spectral model. PL: power law, MK: thin-thermal plasma in collisional equilibrium (Mewe et al.
1985; Kaastra 1992), with abundance of 0.3 solar.
b: Parentheses indicate single-parameter 90% confidence regions.
c: Absorption column in the LMC. The abundance is assumed to be the average LMC values(Russel & Dopita 1992; Hughes et al. 1998).
d: Absorbed flux in the 0.5 - 9.0 keV band. (Feast 1999).

shell from the source-free regions in the same CCD and observations as those of theshells.
The XM M -N ewton fields covered the whole diffuse structure. However, forconsistency with the Chandra analysis, we divided the diffuse structure into shells
A - D, the same as in the Chandra case. The background regions were selectedfrom the source-free regions in the same observation. The spectra of each shelltaken from all the available observations with Chandra and XM M -N ewton wereanalyzed simultaneously. However, for brevity, we show only the Chandra spectra(background-subtracted) and the fitting results in figure 7.2.
For spectral analyses, we used XSP EC, version 11.00 (Arnaud 1996). The
spectrum of shell A shows many line-like structures with the center energies at 0.58,
0.68, 0.92, and 1.35 keV, which correspond to the emission lines of He-like O Kα andKβ, He-like Ne, and He-like Mg, respectively. We therefore fitted the spectrum witha thin thermal plasma model in non-equilibrium ionization (NEI; Borkowski et al.
2001a) with the mean LMC abundances (Russel & Dopita 1992; Hughes et al. 1998).
The interstellar absorption in our Galaxy and LMC were treated separately. The
Galactic absorption column was estimated using the H I data by Dickey & Lockman(1990) as N = 6.35 × 1020 cm−2 . Arabadjis & Bregman (1999) reported that thevalue of NH , measured in the X-ray band, is twice that of N in the case of |b| > 25◦
and N > 5 × 1020 cm−2 . Therefore, we fixed the galactic absorption column to be
NH = 1.27×1021 cm−2 ; we used the cross sections of Morrison & McCammon (1983)
and the solar abundances of Anders & Grevesse (1989). The absorption column in

8.64/9

CHAPTER 7. OBSERVATIONS AND RESULTS

87

Figure 7.2: (Upper panels:) Background-subtracted spectra for shells A - D (plussigns). Although the spectral fittings were made with both the Chandra and XM M N ewton data simultaneously, only the Chandra data and results are shown forbrevity. The best-fit models are shown with solid (power-law) and dotted (NEI)
lines. (Lower panels:) Data residuals from the best-fit models.

88

CHAPTER 7. OBSERVATIONS AND RESULTS

the LMC was, on the other hand, treated as a free parameter with the mean LMC
abundance (Russel & Dopita 1992; Hughes et al. 1998).
This thin-thermal plasma model was rejected with a χ2 /degrees of freedom(d.o.f.) = 331.8/211, even if we allow the abundances to be free, leaving a systematic
data residual at the high-energy band. We hence added a power-law component on
the thin-thermal model. Since the two-component model still leaves large residuals
at about 0.8 and 1.3 keV, we allowed the abundance of Fe and Mg in the thermal plasma (”NEI” component) to be free. The fitting was then greatly improved
with a χ2 /d.o.f = 264.6/209. Although this two-component model is still rejected
in a statistical point of view, further fine-tuning on the model is beyond the scopeof this paper. Figure 7.2 and table 7.3 show the best-fit models and parameters,
respectively.
Table 7.3: Best-Fit Parameters of Shell A
Parameters

NEI

kT /Γ (keV/M )...
ne tap (1012 cm−3 s)...
b

EM (1058−3cm )...

[Mg/H]c ...
c[Fe/H] ...
Fluxd (ergs cm−2 s−1 )...
NHLMC e (×1022 cm−2 )...

Power Law

0.21 (0.19 - 0.23) 2.9 (2.7 - 3.1)
9.9 (>1.4)...

2.6 (2.1 - 2.8)...

3.3 (2.5 - 4.0)...

0.10 (0.06 - 0.14)...

6.8 × 10−14

8.4 × 10−14

0.24 (0.21 - 0.28)...fNotes. Parentheses indicate single-parameter 90% confidence regions.
a: Ionization timescale, where ne and tp are the electron density and age of theplasma.
R
b: Emission measure EM = ne np dV ' n2e V where np and V are the proton density
and the plasma volume, respectively. The distance to LMC is assumed to be 50 kpc.
c: Abundance ratio relative to the solar value (Anders & Grevesse 1989).
d: Absorbed flux in the 0.5 - 9.0 keV band.
e: Absorption column in the LMC. The abundances are assumed to be the average
LMC values (Russel & Dopita 1992; Hughes et al. 1998).
f: Fixed to the same value as that for the NEI component.

Unlike shell A, the X-ray spectra of shells B - D are hard and featureless, suggesting a non-thermal origin. In fact, a thin-thermal model fitting requires an unrealistically high temperature (>2 keV) and low abundances (z < 0.3). We therefore
fitted the spectra with a power-law model with absorption, which was calculated in
the same way as that for shell A, and found acceptable fits for all of the spectra. The
best-fit models and parameters are shown in figure 7.2 and table 7.4, respectively.
It is conceivable that the spectra of shells B - D may include a small fraction of
the thin-thermal component found in shell A. We therefore added the same thinthermal spectrum as that for shell A and fitted them with this composite model(thin-thermal plus power-law). The free parameters are normalizations of the two

89

CHAPTER 7. OBSERVATIONS AND RESULTS

Table 7.4: Best-Fit Parameters of Shell B - D

Parameter

Shell
B

Flux (ergs s−1

DPower Law
2.7 (2.5 - 2.9)
2.3 (2.1 - 2.4)

Photon index
a

C

−2

cm )

NHLMC b (1022 cm−2 )2χ /d.o.f.

4.7 × 10−130.10 (0.06 - 0.19)

4.5 × 10−13

1.1 (0.94 - 1.4)

νrolloff (1016 Hz)

193.0/138
68.3/84
SRCUT (α = 0.5)
6.2 (4.1 - 7.9)
24 (15 - 47)

Flux density at 1 GHz

1.9 (1.5 - 3.0)

2.5 (2.3 - 2.7)
1.5 × 10−13

1.0 (0.85 - 1.3)
121.5/110
12 (6.4 - 23)

0.71 (0.67 - 0.75) 0.42 (0.25 - 0.83)(10−2 Jy)
NHLMC (1022 cm−2 )χ2 /d.o.f.

(<0.04)1.0 (0.90 - 1.1)

0.89 (0.75 - 1.0)

νrolloff (1016 Hz)

195.5/138
67.0/84
SRCUT (α = 0.6)
8.0 (5.3 - 11)
36 (20 - 75)

16 (8.3 - 35)

Flux density at 1 GHz

11 (8.2 - 16)

4.1 (3.9 - 4.3)

2.5 (1.6 - 4.6)

NHLMC (1022 cm−2 )(<0.01)1.0 (0.91 - 1.2)

1.1 (1.0 - 1.3)

χ2 /d.o.f.

195.2/138

67.1/84

121.8/110(10−2120.2/110

Jy)

Note. Parentheses indicate single-parameter 90% confidence regions.
a: Flux in the 0.5 - 9.0 keV band.
b: Absorption column in the LMC. The abundances are assumed to be the average LMC values(Russel & Dopita 1992; Hughes et al. 1998).

components: the power-law index and NH value. However, no significant fraction of
the thin-thermal component is found from shells B - D.
For all of the point sources, the X-ray photons are collected from an ellipse
with the radii of the PSF, as listed in table 7.2. We note that all of the sources
are located far from the on-axis position of the X-ray mirror, and hence the PSFs
are larger than the best value of the Chandra PSF (∼ 0.500 on the aim point).
The background regions were selected from source-free regions in the same way
as the diffuse emissions. We first fitted the spectra with a thin-thermal plasma
model in collisional equilibrium (MEKAL; Mewe et al. 1985; Kaastra 1992) with
an absorption calculated in the same way as diffuse emission. The abundances are
fixed to be 0.3 solar, the average value of the interstellar medium in the LMC.
The fittings are acceptable for two sources (1 and 4) with reasonable temperature(2.1 and 1.0 keV), but for the spectra of the other four sources, the models are

CHAPTER 7. OBSERVATIONS AND RESULTS

90

either rejected or require an unreasonably high temperature. We therefore fitted
the spectra of these sources with a power-law model. The best-fit parameters andthe reduced χ2 are listed in table 7.2.
The absorption columns of the northeastern shells (shells A and B) are similar to
those of most sources in the LMC (∼ 1021 cm−2 ), whereas those of the other shells(C and D) are significantly larger (∼ 1022 cm−2 ) than the typical LMC absorption.
A similar trend was found for the point sources; those in the western half, sources
1 and 2, have a larger NH (∼ 1022 cm−2 ) than those in the eastern half, sources 3- 6 (≤ 1021 cm−2 ). Since Dunne et al. (2001) have already reported this tendency,
we have thus confirmed the results with the better spatial and spectral capabilityof Chandra. This systematic increase of absorption toward the western region of
30 Dor C may be due to extra absorption of a molecular cloud located in front of
the western half of 30 Dor C. To verify our conjecture, we searched for the molecular
cloud in the CO map (Yamaguchi et al. 2001, uppermost panel of figure 7.2) andfound a candidate with intensity of I(CO) ∼ 3.6 K km s−1 . With a conversion factorof N (H2 )/I(CO) ∼ 9 × 1020 cm−2 (K km s−1 )−1 (Fukui et al. 1999), the estimatedabsorption column due to the molecular cloud is NHMC ∼ 6.5 × 1021 cm−2 , which isconsistent with our result.

7.2

N44 (DEM L 152)

N44 is a bright H II complex containing an 44 pc×67 pc SB cataloged as DEM L 152
around the OB association LH 47. The Hα shell is well detected with well-defined
shell walls as shown in figure 7.3. A variety of Hα-bright structures are seen in N44,
and identifications of the components, both bright knots and large filamentary arcs,
have been made by Henize (1956) with labels N44 A - N Davies et al. (1976) andwith labels DEM L 140 - 170. The radio emission from the ring of N44 is thermal(Mathewson & Clarke 1973), and the expansion velocity within the ring is 30 km s−1(Georgelin et al. 1983).
At the southwest corner, coinciding with a bright shell-like H II region in which
high-velocity material moving at –120 km s−1 with respect to the main component
has been detected (Goudis & Meaburn 1984). The expansion patterns of the mainshell and the large shell to the east are studied by Meaburn & Laspias (1991). The
main shell was modeled as SB using energy input implied by the observed massive
stellar component. It is shown that the expansion velocity of the main shell is muchhigher than expected (Oey & Massey 1995). This discrepancy may due to thesignificant energy/pressure lost by the blowout shown at the southwest of the SB.
Another important result they reported is that the stellar component interior to
the shell is generally older than the population immediately exterior to the shell (∼10 Myr vs. <5 Myr), suggestive of sequential star formation.
Because the shell is one of the brightest SB in the LMC in the soft X-ray band(see table 6.1) and then has been observed with many instruments. Diffuse X-ray
emission from various portions of the N44 region was detected by the Einstein
Observatory (Chu & Mac Low 1990; Wang & Helfand 1991). Chu et al. (1993)

CHAPTER 7. OBSERVATIONS AND RESULTS

91

Figure 7.3: Hα image of N44, showing the nebular components defined by Henize(1956) and the OB associations cataloged by Lucke & Hodge (1970). This figure isadopted from Nazé et al. (2002).
reported a breakout region on the southern edge of the SB and an X-ray blisteron the eastern edge of the shell are detected in the ROSAT image as well. The
high temperature plasma indicated by the ROSAT results stimulates us to search
non-thermal X-ray emission as well as 30 Dor C.
N44 was observed with EPIC onboard XM M -N ewton on the 2004 December 4- 5 (Obs. ID 0205740101). The original exposure time was ∼50 ks. The telescope
optical axis position on the EPIC was R.A. = 05h 22m 20.s 00, Decl.= −67◦ 560 40.00 0(J2000). Thanks to the large FOV (300 of diameter), the whole region of N44
including nearby SNR was covered. We used the Observation Data Files (ODF)
data provided by the XM M -N ewton Science Operations Centre. For the data
reduction, we used emchain and epchain in the analysis software SAS version 6.5.0
with the bad pixel finding algorithm on. Then we cut the events in the duration of
background flare with the following way. First, we made light curves of the Pattern
0 (single pixel) events in the energy band of 10 - 15 keV with the 100 sec binningand made histogram of the count rate. Then we fit the distribution of count ratewith gauss function and cut the noisy duration in which the count rate is higher
with 2 σ significance than the gaussian peak. Second, we did the same cutting but
using the events detected in the region of >120 from the center of FOV in the energy
band of 1.0 - 5.0 keV and Pattern <12 for MOS and <4 for pn chip. Then, theeffective exposure times were ∼23 ks for MOS and ∼17 ks for pn chip. Hereafter,
we used those events that has Pattern < 12 for MOS, <4 for pn chip in order toeliminate the events due to charged particles and hot and flickering pixels.
Figure 7.4 shows soft (0.2 - 2.0 keV) and hard (2.0 - 7.0 keV) band X-ray image of92

CHAPTER 7. OBSERVATIONS AND RESULTS

N44. As shown in the past X-ray observations, the distribution of the X-ray emission
has clear correlation with that of Hα (Magnier et al. 1996). In the soft band, there
is an almost uniform X-ray emission from entire region of main shell, some blowout
regions, and the nearby SNR 0523–67.9. However in the hard band, there seems tobe almost no significant diffuse emission in the main shell and blowout regions. Thisfeature is obviously different from that was seen in 30 Dor C (Bamba et al. 2004).

-67:50:00-67:50:00-67:54:00-67:54:00-67:58:00-67:58:00-68:02:00-68:02:00

5:23:00

5:22:00

5:21:00

5:23:00

5:22:00

5:21:00

Figure 7.4: Soft (left panel; 0.2 - 2.0 keV) and hard (right panel; 2.0 - 7.0 keV)band X-ray image of the N44 region. The regions from which we extracted photons
are shown in the soft band image. The main shell are divided into northwest andsoutheast regions. Hereafter we concentrate our study on the main shell. In thehard band image we show the point sources detected in the band.
We conducted the point source search with eboxdetect and emldetect commandin SAS 6.5.0 in the soft (0.2 - 2.0 keV) and hard (2.0 - 7.0 keV) energy band. The
detection minimum likelihood was 6 for eboxdetect and 10 for emldetect. While we
did not detect any plausible point source in the soft band, we detected 13 sourcesfrom the surrounding region of the main shell (7.4) in the hard band. However, we
could not detect any point source from the OB association in the vicinity of mainshell (LH 47 and LH48). Since the HEW of the XRT is ∼ 1500 (see section 2.2.1),
the mean distance between the closest neighbours in the typical OB association(Rdis ' 6 pc ' 2700 at the distance of 50 kpc) can be resolved. Our result indicatesthat the luminosities of the members in LH 47 and LH 48 is not so high that they
are detected as point sources.
Figure 7.5 shows the spectrum of MOS 1 and 2 chip extracted from the entire
main shell region (the left panel in figure 7.4). For the brevity to see the spectra, weonly show the spectra of EPIC MOS detectors and hide that of pn detector, although
we used all of them in the fitting. As in the case of 30 Dor C, some emission lines from
He-like O Kα and Kβ, He-like Ne, and He-like Mg. The instrumental backgroundcomponent (Neutral Al line; 1486 eV) is also seen. Hence we tested thin-thermal
plasma model in NEI with the mean LMC abundance and composite (thermal andnon-thermal) plasma model in the spectral fit. We fixed the galactic absorption
column to be NH = 1.27 × 1021 cm−2 (see section 7.1). In the case of the latter

CHAPTER 7. OBSERVATIONS AND RESULTS

93

model, fit result was statistically rejected while the thin-thermal model was acceptedwith a χ2 /d.o.f. = 188.24/153. The fitting parameters are shown in table 7.5.

Figure 7.5: Background-subtracted spectrum extracted from entire main shell region. We fit this spectrum with thin thermal plasma model with two temperature.
Al fluorescence line is also considered.
We then further divide main shell into three regions: northwest, southeast, andthe region around LH 47 (figure 7.4). The energy spectra from formers are shown infigure 7.6 and their fitting parameters are presented in table 7.5. Northwest region
shows the high temperature feature seen in figure 7.5 and significantly low ionization
parameter than southeast region.
The plasma temperature seen in the northwest region is exceedingly high compared with other SBs (Nazé et al. 2004; Cooper et al. 2004). This feature suggeststhat N44 is the most active, that is, experienced the latest SN among these SBs.

Figure 7.6: Background-subtracted spectrum extracted from northwest and southeast region of main shell (left and right panel, respectively). Al fluorescence line isalso considered.
Oey & Massey (1995) and Will et al. (1997) have observed OB associations
LH 47, which is located inside the main shell and its surrounding regions with

94

CHAPTER 7. OBSERVATIONS AND RESULTS

Table 7.5: Fitting parameters for each region.

Regionfluxa

χ2 /d.o.f.

2.5 (1.9–3.8)

3.51

188/153

2.2(1.5–2.5)

5.3 (4.6–7.0)––

1.8 (1.7–1.9)

1.4 (1.1–1.7)

1.30

180/139

91 (68–240)

0.35

72/95

820 (48–5000)

0.12

22/20

NH

kT

τ

(×1021 Hcm−2 )(keV)(×1010 s cm−3 )(kTlow component)

2.92 (2.42–3.46)

0.6(0.4–0.8)(kThigh component)–

North-West

2.40 (1.51–3.03)

South-East

1.00 (0.73–1.19) 0.40 (0.39–0.43)

Entire Main Shell

Vicinity of OB-stars

1.81 (<2.20)

0.40 (0.32–0.64)

Note. Parentheses indicate single-parameter 90% confidence regions.
a: Absorption corrected flux in the 0.5 - 7.0 keV band (×10−12 ergs cm−2 s−1 ).

optical light. Oey & Massey (1995) investigated spectroscopic data of the 62 bluest
stars (> 10 M )in the LH 47/48, and derived the spectral types of 54 stars in LH
47. Will et al. (1997) also took spectra of V band brightest 14 stars and got almost
consistent spectral types of them. On the whole, OB-type stars are concentrated at
center of main shell.
Hence we defined the regions from which we extract the photons in figure 7.4
and got the spectrum shown in figure 7.7 and fitting parameters presented in table
7.5. Furthermore in the case of this region, we set the abundance of O, Ne, and Mg
to be free. Then we got the abundance of 0.46 (>0.10), 1.15 (>0.26), 0.91 (>0.21),
respectively.

7.3

Sgr C

Sgr C region, the complex of the massive molecular clouds and hence the promising candidate for the clustering of massive stars and SNe. It is located at ∼ 340
southwestward from the GC, near the base of one of the two large radio continuumarcs and many other bright non-thermal filaments (NTFs) as shown in figure 7.8.
The neutral and ionized gas distribute hard against not only the H II region but thebright portion of Sgr C’s NTF (Liszt & Spiker 1995).
The spectra of J = 1 − 0 13 CO emission near Sgr C show the terrifically complex distribution around Sgr C. Hence it is very easy to make false physical associations based on random spatial coincidence. However, most possibly associatedcomponent which we should pay attention is at the recombination-line velocity of
Sgr C (∼65.5 km s−1 ). The sum of all the emission integrated over the range(−84 km s−1 < v < −24 km s−1 ) shows a cavity or cleft. The Sgr C shell-like H II
region sits within the cavity of the associated molecular gas. The origin of this cavity

CHAPTER 7. OBSERVATIONS AND RESULTS

95

Figure 7.7: Background-subtracted spectrum extracted from vicinity of OB-typestars association LH 47. Al fluorescence line is also considered.
must be the energetics of the H II region, which is dominated by two influences. One
is the thermal pressure maintained by photoionization, and the other is the strong
stellar winds expected from early O-type stars (Abbott 1982) which was inferred
by Odenwald & Fazio (1984). The mass of the main body of this cavity gas below
Sgr C (M359.5–0.15 in Liszt & Spiker 1995), calculated by assuming that 13 CO is
20% as bright as 12 CO, taking a CO-H2 conversion of 3×1020 H2 cm−2 (K km s−1 )−1
is 6.1 × 105 M .
In the X-ray band, Murakami et al. (2001) discovered the strong 6.4 keV line
emission, which is the key characteristic of X-ray Reflection Nebula (XRN; Koyama
et al. 1996; Sunyaev & Churazov 1998). The 6.4 keV band image shows an X-ray
peak near the molecular cloud (MC) CO 359.4+0.0 (Oka et al. 1998), which is oneof the giant MC in the Sgr C complex. The spectrum from X-ray peak show a large
absorption in the low energy band, which was also a good evidence that Sgr C isa new XRN. Although the absorption of Sgr C was smaller than that of Sgr B2,
the well-established XRN, by an order of magnitude, this is reasonable because themass of the Sgr C cloud is about one-seventh that of the Sgr B2 cloud (Oka et al.
1998).
Suzaku observed Sgr C complex region from 20th to 23rd February 2006 during
the Performance Verification stage. After the data reduction, the effective time is
115 ks for each XIS. The observation mode is the 3×3 mode for almost all (∼97%)of the data.
Figure 7.9 shows the X-ray intensity maps at various energies. 6.4 keV line
(upperleft panel) emission consists of primarily neutral Fe atoms of which a K-shellelectron is released. 2.45 keV and 6.7 keV lines are from He-like S and He-like Fe,
respectively. Thanks to the Suzaku’s superior energy resolution, we could reliably

CHAPTER 7. OBSERVATIONS AND RESULTS

96

Figure 7.8: Contour map of radio continuum emission around Sgr C region taken at1616 MHz with Very Large Array. The beam width was 1300 .2 by 2300 (α × δ).
make narrow-band images. We can see some new structures. First, new 6.4 keVclump around (l,b) = (359.5, –0.2), which is consistent with the position of M359.5–
0.15 in Liszt & Spiker (1995) can be seen. Second, a strong emission around thecenter of FOV in the 2.45 keV map (hereafter we call this emission G359.4–0.1)
and somewhat “blowout” structure which extends upperward in the figure from
G359.4–0.1.
G359.4–0.1 and “blowout” structure indicate that there is thin-thermal plasmacomponent superposed on the GC hot plasma and 6.4 keV clumps. In order to
see the characteristics of this component, we extract spectra from these regions.
Figure 7.10 shows the XIS spectra from G359.4–0.1 and “blowout” region. Since
the three components shown in figure 7.9 overlap each other, the contamination by
the other two components is inevitable. So we add the 6.4 keV (Kα) and 7.1 keV(Kβ) line and power-law components to the fitting model in the case of G359.4–0.1.
the interval and intensity ratio of the Kα and Kβ was fixed to the value obtainedwith ground experiments. For the “blowout” spectrum, we also add the GC hot
plasma component, of which all the parameter such as absorption, temperature andabundance is fixed to the value in Koyama et al. (2006b) and only the normalization
is free. The derived parameter for the component we interested in is shown in table
7.6. Both emissions are well fitted with thin-thermal plasma model with strong
absorption and temperature of ∼0.7 keV.

97

CHAPTER 7. OBSERVATIONS AND RESULTS

Table 7.6: Best-Fit Parameters of G359.4–0.1 and “blowout” region
Parameters

G359.4–0.1“blowout”

NH (10 cm )...

12.0 (11.0 - 15.1)

9.6 (8.9 - 9.9)

kT (keV)...

0.71 (0.56 - 0.94) 0.71 (0.56 - 0.75)

22

−2

a

[S/H] ...

1.6 (1.1 - 2.0)

1.3 (1.0 - 1.7)a2.0 (1.0 - 2.8)

2.7 (1.5 - 3.7)

2.2 (0.6 - 3.8)

1.4 (< 4.1)

[Ar/H] ...
[Ca/H]a ...
b

Flux (ergs cm
χ2 /d.o.f. ...

−2−1s )...

3.6 × 10−1354.0/82

2.9 × 10−1379.0/84

Note. Parentheses indicate single-parameter 90% confidence regions.
a: Abundance ratio relative to the solar value (Anders & Grevesse 1989).
b: Absorbed flux in the 2.0 - 10.0 keV band.

7.4

CTB 37A

CTB 37A is one of the radio-bright SNRs in the CTB 37 complex. Figure 7.11 showsa contour map from the radio (90 cm) observation. A bright radio shell with a radiusof ∼ 30 can be clearly seen, overlapping G348.5–0.0 in the east side. CTB 37A isactually two SNRs overlapped in projection (Kassim et al. 1991). One of them,G348.5+0.1 (to the west in figure 7.11), is a partial shell with a faint extension out
the open end of the shell (“breakout” morphology), while the other one, G348.5–
0.0 (to the east in figure 7.11), is a partial shell. CTB 37B, another radio SNRs
complex locates near from CTB 37A (∼ 200 ), hence this region can be categorized
to be the SB formation site. Caswell et al. (1975) concluded by H I observations
that CTB 37A is beyond the tangential point, and Frail et al. (1996) obtained theconsistent result of 11 pc by OH maser (1720 MHz) observations, which correspondswith a shell radius of ∼ 10 pc. An inspection of the CO emission toward CTB 37A
detected the MC that were probably accompanied by OH masers in the northernpart of the shell (Reynoso & Mangum 2000).
In the X-ray band, CTB 37A was detected by the ASCA Galactic plane survey
data. However, due to the stray light from a bright southward source and the limited
statistics, we could not determine whether the X-ray emission from CTB 37A isthermal or non-thermal. In order to clarify the X-ray emission mechanism andto obtain the X-ray emission map of the inside the SNRs, we observed CT 37A
with Chandra/ACIS-I and XM M -N ewton/EPIC. The observed dates and targeted
positions are given in table 7.7. Data acquisition from the ACIS was made in
the timed-exposure very faint mode in order to significantly reduce the particle
background. The data reductions and analyses were made using the CIAO software,version 3.0. For the XM M -N ewton data, EPIC CCDs were operated in the fullframe mode with the medium filter. We conducted the background flare cut withthe same way as section 7.2. Then, we obtain the effective exposure as shown in

98

CHAPTER 7. OBSERVATIONS AND RESULTS
table 7.7, while the intrinsic exposure was 20 ksec.
Table 7.7: Observation Log around CTB 37A

Obs.

Satellite

1...

Chandra2...

XMM-Newton

Obs. ID
6721

R.A.

Decl.

(J2000.0)(J2000.0)

17 14 35.2 -38 31 38.4

Date

Exposure(ks)

2006 Oct 7

200306510101 17 14 31.2 -38 29 54.2 2006 May 1

13

Note. Units of right ascension are hours, minutes, and seconds, and units of declination aredegrees, arcminutes, and arcseconds.

Figure 7.12 shows the three color X-ray image obtained with Chandra. Theenergy bands are 0.5 - 2.0 keV for red, 2.0 - 7.0 keV for green, and 7.0 - 10.0 keV
for blue image, respectively. The new features discovered by our observations arethat the northeast part of the G348.5–0.0 show the strong soft X-ray emission while
G348.5+0.1 does not show significant emission and that there is a knot-like structure
which emit strong hard X-rays at the wast edge of G348.5–0.0.
We then extracted the X-ray spectra from the primarily X-ray emitting regionin G348.5–0.0 except for the knot. Figure 7.13 shows two spectra, one of whichis extracted from the northeast part of G348.5–0.0 except for knot region, and theother is from knot region. While the former shows some emission lines from He-like
Si and S, the latter shows the smooth continuum emission. Hence the X-rays from
G348.5–0.0 can be fitted with the thin-thermal plasma model and those from theknot can be reproduced by power-law emission. The best fit parameters are shownin table 7.8. Because the absorption of the latter is about three times stronger than
the former, the origin of the emission from the knot structure, which we expected
to be the non-thermal X-rays from this SNR in the ASCA spectrum, is possibly anobject which locates on the line of sight along G348.5–0.0 accidentally.

99

CHAPTER 7. OBSERVATIONS AND RESULTS

Table 7.8: Best-Fit Parameters of G348.5–0.0 and knot
Parameters

G348.5–0.0

knot

kT1 /Γ (keV/M )...

0.38 (0.32 - 0.51) 1.8 (1.3 - 2.3)

kT2 (keV )...

1.55 (1.41 - 1.95)...

c[Mg/H] ...

1.2 (0.86 - 1.7)...

[Si/H]c ...

1.3 (1.0 - 1.6)...

1.6 (1.2 - 2.0)...

[S/H]c ...
dFlux (ergs cm−2−1s )...

NHLMC e (×1022 cm−2 )...

5.6 × 10−12

3.7 (3.3 - 4.0)

3.7 × 10−12

10 (7.9 - 12)

Notes. Parentheses indicate single-parameter 90% confidence regions.
a: Ionization timescale, where ne and tp are the electron density and age of theplasma.
R
b: Emission measure EM = ne np dV ' n2e V where np and V are the proton density
and the plasma volume, respectively. The distance to LMC is assumed to be 50 kpc.
c: Abundance ratio relative to the solar value (Anders & Grevesse 1989).
d: Absorbed flux in the 0.5 - 9.0 keV band.
e: Absorption column in the LMC. The abundances are assumed to be the average
LMC values (Russel & Dopita 1992; Hughes et al. 1998).
f: Fixed to the same value as that for the NEI component.

CHAPTER 7. OBSERVATIONS AND RESULTS

100

Figure 7.9: X-ray images of Sgr C region. The energy bands are 6.4 keV, 2.45 keV(He-like S Kα line), 6.7 keV (He-like Fe Kα) for upperleft, upperright, bottomleft
panel, respectively. Particle background component is subtracted and the effective
exposure is corrected. Red and green polygonal regions indicate the source andbackground regions adopted in the following spectral analysis, respectively.

CHAPTER 7. OBSERVATIONS AND RESULTS

101

Figure 7.10: (Left panel:)Background-subtracted spectrum extracted from G359.4–
0.1 region. The power-law (green) component and 6.4 keV line emission in the fitted
model function are the contaminants by the emission distributed over the FOV.
(Right panel:) Background-subtracted spectrum extracted from “blowout” region.
The hot plasma component of GC (red), power-law component (sky blue), and6.4 keV line emission (blue) are also in the fitted model function as the contaminants.

Figure 7.11: A radio continuum image of CTB 37A region at 1.4 GHz from Kassim
et al. (1991). The positions of OH (1720 MHz) masers are superimposed as crosseson the image. The beam size is 3200 by 1800 (α × δ).

CHAPTER 7. OBSERVATIONS AND RESULTS

102

Figure 7.12: The Chandra three-color image of CTB 37A region (red = 0.5 - 2.0 keV,green = 2.0 - 7.0 keV, blue = 7.0 - 10.0 keV). Red, green, and blue regions indicate
the source regions adopted in the following spectral analysis.

Figure 7.13: (Left panel:)Background-subtracted spectra of ACIS-I and EPIC extracted from the northeast part of G348.5–0.0 although the “knot” region is eliminated. (Right panel:) The same as left panel but for knot region.

Chapter 8
Discussion
8.1

Cosmic Ray Acceleration in SB

30 Dor C is an SB made by a strong stellar wind and/or successive SN explosionsof massive stars located in the OB star association LH 90. The age of this starassociation, or that of 30 Dor C, is on the order of a few to 10 Myr (Lucke &Hodge 1970). The thermal emissions are enhanced in the southeastern side of
30 Dor C (around shell A). The position of this component coincides with thatof the Hα emission (Dunne et al. 2001). The plasma temperature of 30 Dor C israther high compared to those of other LMC SBs (Dunne et al. 2001), although
it becomes significantly lower than previous results. Perhaps because of the poor
spectral resolution of ROSAT , previous observations could not resolve the powerlaw component (hard spectrum) from the thermal emission.
The X-ray luminosity of 30 Dor C is significantly lower than that of the otherSBs (Dunne et al. 2001). With the assumption that the plasma in shell A distributes
uniformly in the ellipsoid with radii of 5800×3400×3400 (total volume V = 1.2×1059 cm3 ),the mean density (ne ), thermal energy (E), and the age of the plasma (tp ) were
calculated as follows using the emission measure EM = n2e V and ionization timescale(see table 7.3)
ne = 4.7(4.2 − 4.8) × 10−1 cm−3 ,(8.1)

E ' 3ne kT V = 5.7(4.6 − 6.4) × 1049 ergs,(8.2)

tp = 6.7(> 0.5) × 105 yr.

(8.3)

Together with the Hα emission around shell A (Dunne et al. 2001), we inferthat the shell of 30 Dor C collides with dense matter and temporally emits thermal
X-rays with a relatively high temperature. The overabundance of a light element(Mg) relative to a heavier element (Fe) may indicate that a Type II SN occurred(Tsujimoto et al. 1995). Thus, the progenitor is a massive star that is a member of
4.4).
cluster LH 90 (Lucke & Hodge 1970) near the center of 30 Dor C (see103

CHAPTER 8. DISCUSSION

104

In the 843 MHz band, Mills et al. (1984) found a clear radio shell that is brightest
on the southwestern side (around shells C and D) and dim on the eastern side. The
nonthermal X-ray emissions are enhanced at the radio bright shell. This fact implies
that the nonthermal X-rays are emitted by the same mechanism as the radio band.
Therefore, the nonthermal X-rays are likely to be synchrotron radiation from theaccelerated electrons, like SN 1006 (Koyama et al. 1995) and other SNRs, which
accelerate particles up to TeV. The X-ray photon index of 2.1 - 2.9 (see table 7.4)
is, in fact, typical to synchrotron emissions.
To verify the synchrotron origin, we fitted the X-ray spectra with an SRCUT
model, which represents the synchrotron emission from electrons with an energy
distribution of a power law plus exponential cutoff (Reynolds 1998, 1999). Since
the radio index data are not accurate as a result of the large background and contamination of thermal emission (Mathewson et al. 1985), we fixed the spectral index(α) at 1 GHz to be 0.5, which is expected from the first-order Fermi accelerationand similar to that of SN 1006 (α = 0.57; Allen et al. 2001).
The fittings were statistically acceptable; the best-fit parameters are listed intable 7.4. The best-fit cutoff frequency is also similar to that of SN 1006, although
the age of 30 Dor C may be on the same order as that of LH 90 (10 Myr; Lucke &Hodge 1970), which is far larger than SN 1006. This implies that, depending on
the environment, the acceleration of high-energy electrons up to the knee energy cancontinue for a far longer time than the previous consensus for the SNR case (∼ 103 yr;
Reynolds 1999). The electron acceleration time may have been extended, because
successive SN explosions in 30 Dor C, possibly over the course of a few Myr, may
more-or-less continuously produce high-energy electrons. The total luminosity of thenonthermal component (∼ 5.3 × 1035 ergs s−1 ) is about 10 times larger than that of
SN 1006 (Koyama et al. 1995). This large nonthermal flux in 30 Dor C would be
due to the large energy supply from multiple successive SN explosions.
The expected flux density at 1 GHz is 3.0 (2.4 - 4.6) ×10−2 Jy with α = 0.5,
which is significantly smaller than the observed value (1.0 Jy; Mills et al. 1984). This
”inconsistency” is not relaxed, even if we assume a larger radio index of α = 0.6[1.8 (1.4 - 2.5) ×10−1 Jy, as shown in table 7.4]. We infer that the larger observed
radio flux than that expected from X-rays would be due to either the contaminationof thermal radio flux or a large uncertainty of the background level (see Mills et al.
1984).
The Hα emission (Dunne et al. 2001) is anticorrelated with the nonthermal
components. Similar features have been observed in some of other SNRs with synchrotron X-rays: SN 1006 (Winkler et al. 1997) and RCW 86 (Smith 1997). Because
the Hα region has higher density, it may have a higher magnetic field. Therefore, themaximum electron energy is limited by the quick synchrotron energy loss, leading
to reduced nonthermal X-rays.
We have identified three pointlike X-ray sources (sources 1, 3, and 4) to the tightclusters of massive stars α, β, and γ, respectively. The spectrum of source 1 (α) is
fitted with a thin-thermal plasma model of ∼2.1 keV, which is consistent with stellar
X-rays from massive stars. The cluster α is the brightest of the three, with X-ray
luminosity of ∼ 1034 ergs s−1 . Optical spectroscopy of this cluster revealed that it

CHAPTER 8. DISCUSSION

105

includes one red giant and one Wolf-Rayet (W-R) star (MG 41 and Brey 58; Lortet& Testor 1984). Hence, the latter is a possible counterpart of the 2 keV source. The
X-ray luminosity is near to the upper end of a massive star, or its binary (Maeda
et al. 1999). Although the X-ray luminosity of source 3, a counterpart of the starcluster β, is also consistent with a massive young star, the best-fit photon index,∼2 (see table 7.2), is rather typical of a rotation-powered neutron star, probably astellar remnant of an SN explosion in the active star cluster. The spectrum of source
3 (γ) is soft and well fitted, with a kT = 1.0 keV thermal plasma model, similar toa cluster α. The star cluster γ also includes one OB star (Sk -69◦ .212; Sanduleak1970), and the X-ray luminosity is consistent with that of a young massive star.
The other three sources (2, 5, and 6) have no counterpart in optical, in theinfrared band (Breysacher et al. 1981; Lortet & Testor 1984), nor in the SIMBAD
database. Their spectra are relatively hard and are consistent with being background
active galactic nuclei (AGNs) or stellar remnants (black hole or neutron star) bysuccessive SN explosions.
In order to constrain the nature of the pointlike X-ray sources with a power-law
spectrum, we further examined the time variability with the Kolmogorov-Smirnov
test (Press et al. 1992). However, no significant time variability was found evenbetween the two observations. For high-resolution timing, we examined the Chandra
HRC data (ObsID = 738), but no pulsation was found from any of these point
sources.
As shown in chapter 7, significant non-thermal X-rays were not detected fromother SBs or SB formation sites. This may have relation to the breakout, which is
obviously seen in N44. As described in section 6.2.3, hot and tenuous environment
inside SBs are essential condition for the effective particle acceleration. If the gas
inside the SB leak into the outer ISM, the strength and turbulance of the magnetic
field become weak and hence results as the less-effective acceleration.

8.2

Energetics of Superbubble

Utilizing the thermal energy derived in section 7.2, we investigate the energy balancein the case of N44.
For the amount of energy injection to SB, the combined action from SN andstellar winds from OB-type stars must be considered. Oey & Massey (1995) derived
the approximate age of LH 47 from the initial mass function of members to be∼10 Myr, then the integral of mechanical luminosity of stellar wind from OB-typestars and SN explosion are ∼3×1051 ergs and ∼1–4×1051 ergs respectively.
On the other hand, the amount of energy ejection can be estimated by measuring
the thermal energy of hot gas in the SB and kinetic energy of H I and H II shell.
Kim et al. (1998) estimated the latters to be ∼2.4×1050 ergs and ∼5.0×1049 ergs
respectively, which is negligible compared with SN explosion energy. Thermalenergy√inside the main shell, which can be calculated from flux, is 2.62×1051  ergs (wherethe  is filling factor). Furthermore we also derived the flux and thermal energy fromblowout region and western another shell (Figure 7.4). The sum of the thermalCHAPTER 8. DISCUSSION

106

√
√
energy is 2.15×1051  ergs. Then the total thermal energy is 4.77×1051  ergs,
which reaches to the total injected energy.
This feature shows a clear contrast against N51D, another SB in the LMC. In
the case of N51D, approximately two third of the injected energy is missing (Cooper
et al. 2004) which means a significant fraction of the stellar mechanical energy musthave been converted into other forms of energy. The X-ray spectrum of the diffuse
emission from N51D requires a power-law component to explain the featureless emission at 1.0 - 3.0 keV. The origin of this power-law component is unclear, but it may
be responsible for the discrepancy between the stellar energy input and the observed
interstellar energy in N51D. In the case of N44, no significant diffuse non-thermal
emission possibly corresponds to the negligible amount of missing energy. The casedependency of the amount of the missing energy have to be investigated more detailcomparing our X-ray dataset with data of other bands.
In the typical case of X-ray emission from OB-type stars, the temperature of
hot plasma is <1.0 keV and the solar abundance. The hot plasma in the vicinity ofOB-type stars shows quite a similar feature; the abundance of Ne and Mg is similar
to the solar value, while other region inside the main shell show consistent resultswith the mean LMC value of 0.3. According to the catalog in the Oey & Massey(1995), there are at least 4 O-type stars and 5 B-type stars in this region. On
the other hand the absorption corrected flux in the energy band of 0.2–2.0 keV is
1.49×10−13 ergs cm−2 s−1 which corresponds to the luminosity of 4.50×1034 ergs s−1 .
Considering the luminosity of galactic O-type stars (Babel & Montmerle 1997), thisvalue is not too large because there must be some contamination from diffuse hot
gas which fills entire shell. Hence we possibly first detect the contribution from
OB-type stars to the hot gas in an SB separately from that of SNRs.

Chapter 9
Conclusion
9.1

Conclusion for XIS calibration

The results of ground and in-orbit experiments of the CI capability of the XIS areas follows:
1. The amount of injected charges (Q) are sufficiently stable (dispersion of 91 eV
in FWHM), which should be compared to the X-ray energy resolution (FWHM
of 113 eV) with the same amount of charges.
2. The CTI depends on PHA of the charge. The charge loss can be explained asδQ ∝ Q0.62−1.00
3. With the CTICI correction , the energy resolution (FWHM) of 55 Fe is improved
from 193 eV to 173 eV approximately one year after the launch and the residual
around the tail component in the line profile is also significantly reduced.
4. The improved charge compensation method is applied to the Tycho’s SNR
data, which results in the same benefit as the calibration source data.
5. We confirm that the energy resolution can be largely improved by filling thecharge trap. Hence, we are currently trying another in-orbit CI capability
experiment to actively fill the CTs in the transfer channel by the CI technique.

9.2

Conclusion for SBs study

The results of study of the energetics and particle acceleration in the SBs are as
follows:
1. In the shell of the superbubble 30 Dor C, we resolved non-thermal and thermal
X-rays using Chandra and XM M -N ewton data. The nonthermal X-rays arelocated in the northern and western parts of the SB. The power-law model iswell fitted to the spectra with Γ = 2.1 - 2.9, similar to that of the cosmic-ray
accelerating SNR, SN 1006. This is the first discovery of nonthermal X-rays
107

CHAPTER 9. CONCLUSION

108

from the shells of SBs. The total luminosity is 10 times larger than that of
SN 1006. Only thermal emission could be seen in the case of N44, CTB 37A,
and Sgr C.
2. If we consider the breakout region around the main shell of N44, the sum of thethermal energy of hot gas and kinetic energy is consistent with the expected
injected energy by the stellar wind and supernovae.
3. The contrast between 30 Dor C and other SBs may have relation to the breakout, which is obviously seen in N44. While, hot and tenuous environment
inside SBs are essential condition for the effective particle acceleration, thestrength and turbulance of the magnetic field become weak and hence results
as the less-effective acceleration due to the inner gas leakage into the outerISM.

Bibliography
Abbott, D. C., 1982, ApJ, 263, 723Åberg, T., 1971, Phys. Rev. A, 4, 1735
Allen, G. E., Petre, R., & Gotthelf, E. V., 2001, ApJ, 558, 739
Anders, E., & Grevesse, N., 1989, Geochim. Cosmochim. Acta, 53, 197
Apanasenko, A. V. et al., 2001, Izv. Akad. Nauk, Ser. Fiz., 65, 433
Arabadjis, J. S., & Bregman, J. N., 1999, ApJ, 510, 806
Arnaud, K. A. 1996, in ASP Conf. Ser. 101, Astronomical Data Analysis Softwareand Systems V, ed. G. H. Jacoby & J. Barnes (San Francisco: ASP), 17
Babel, J., & Montmerle, T., 1997, ApJ, 485, L29
Bamba, A., Koyama, K., & Tomida, H., 2000, PASJ, 52, 1157
Bamba, A., Ueno, M., Nakajima, H., & Koyama, K., 2004, ApJ, 602, 257
Bautz, M. W., Kissel, S. E., Prigozhin, G. Y., Jones, S. E., Isobe, T., Manning, H. L.,
Pivovaroff, M., Ricker, G. R., & Woo, J., 1996, Proc. SPIE, 2808, 170
Bautz, M. W., Kissel, S. M., Prigozhin, G., LaMarr, B., Burke, B. E., Gregory, J. A.,& The XIS Team, 2004, Proc. SPIE, 5501, 111
Bearden, J. A., 1967, Rev. Mod. Phys., 39, 78
Bell, A. R., & Lucek, S. G., 2001, MNRAS, 321, 433
Berezhko, E. G., & Völk, H. J., 2004, A&A, 427, 525
Berezhko, E. G., Ksenofontov, L. T., & Völk, H. J., 2003, A&A, 412, L11
Bird, D. J. et al., 1993, Phys. Rev. Lett., 71, 3403
Blaauw, A. 1991, in the Physics of Star Formation and Early Stellar Evolution,
NATO Advanced Science Institutes (ASI) Ser. C, ed. C. J. Lada & N. D. Kylafis(Kluwer: Dordrecht), 342, 125
Blinnikov, S. I., Imshennik, V. S., & Utrobin, V. P., 1982, SvAL, 8, 361
109

BIBLIOGRAPHY

110

Blondin, J. M.. Wright, E. B., Borkowski, K. J., & Reynolds, S. P., 1998, ApJ, 500,
342
Bomans, D. J., Rossa, J., Weis, K., & Dennerl, K., 2003, in IAU Symposium, No.
212, ed. K. A. van der Hucht, A. Herrero, & C. Esteban (San Francisco: ASP),
637
Borkowski, K. J., Lyerly, W. J., & Reynolds, S. P., 2001, ApJ, 548, 820
Borkowski, K. J., Rho, J., Reynolds, S. P., & Dyer, K. K., 2001, ApJ, 550, 334
Born, M., & Wolf, E., 1999, Principles of optics, 7th, ed. (Cambridge: Cambridge
University Press), 67
Bresolin, F., Kennicutt, R. C., Jr., & Garnett, D. R., 1999, ApJ, 510, 104
Breysacher, J., 1981, A&AS, 43, 203
Burke, B. E., Gregory, J. A., Loomis, A. H., Lesser, M., Bautz, M. W., Kissel, S. E.,
Rathman, D. D., Osgood, R. M., III, & Cooper, M. J., 2004, IEEE Trans. Nucl.
Sci., 51, 2322
Burrows, D. N. et al., 2000, ApJ, 543, L149
Butt, Y. M., Torres, D. F., Combi, J. A., Dame, T., & Romero, G. E., 2001, ApJ,
562, L167
Bykov, A. M., 2001, Space Sci. Rev., 99, 317
Bykov, A. M., & Fleishman, G. D., 1992, MNRAS, 255, 269
Bykov, A. M., & Toptygin, I. N., 2001, Ast. L, 27, 625
Casse, F., Lemoine, M., & Pelletier, G., 2002, Phys. Rev. D, 65, 3002
Castor, J., McCray, R., & Weaver, R., 1975, ApJ, 200, L107
Caswell, J. L., Murray, J. D., Roger, R. S., Cole, D. J., & Cooke, D. J., 1975, A&A,
45, 239
Chevalier, R. A., 1974, ApJ, 188, 501
Chevalier, R. A., 1976, ApJ, 207, 872
Chu, Y. -H., & Mac Low, M. -M., 1990, ApJ, 365, 510
Chu, Y. -H., Mac Low, M. -M., Garcia-Segura, G., Wakker, B., & Kennicutt, R. C., Jr., 1993, ApJ, 414, 213
Chu, Y. -H., Chang, H. -W., Su, Y. -L., & Mac Low, M. -M., 1995, ApJ, 450, 157
Clark, D.H., & Caswell, J.L., 1976, MNRAS, 174, 267

BIBLIOGRAPHY

111

Cooper, R. L., Guerrero, M. A., Chu, Y. -H., Chen, C. -H. R., & Dunne, B. C.,
2004, ApJ, 605, 751
Cox, D. P., 1972, ApJ, 178, 159
Cox, D. P., & Smith, B. W., 1974, ApJL, 189, 105
Davies, R. D., Elliott, K. H., &Meaburn, J., 1976, MNRAS, 81, 89
De Young, D. S. & Gallagher, J. S., III, 1990, ApJL, 356, 15
Dennerl, K. et al., 2001, A&A, 365, L202
Dickey, J. M., & Lockman, F. J., 1990, ARAA, 28, 215
Dunne, B. C., Points, S. D., & Chu, Y. -H., 2001, ApJS, 136, 119
Enomoto, R. et. al., 2002, Nature, 416, 823
Falle, S. A. E. G., 1981, MNRAS, 195, 1011
Falle, S. A. E. G., 1988, Proceedings of IAU Colloq. 101, 419
Feast, M., 1999, PASP, 111, 775
Frail, D. A., Goss, W. M., Reynoso, E. M., Giacani, E. B., Green, A. J., &Otrupcek, R., 1996, AJ, 111, 1651
Fukui, Y. et al., 1999, PASJ, 51, 745
Garmany, C. D., 1994, PASP, 106, 25
Gendreau, K. C. 1995, PhD, Massachusetts Institute of Technology
Gendreau, K., Bautz, M. W., & Ricker, G., 1993, Proc. SPIE, 335, 318
Georgelin, Y. M., Georgelin, Y. P., Laval, A., Monnet, G., & Rosado, M., 1983,
A&AS, 54, 459
Goudis, C., & Meaburn, J., 1984, A&A, 137, 152
Grant, C.E., Bautz, M.W., Kissel, S.M., & LaMarr, B., 2004, Nucl. Instr. and Meth.
A, 5501, 177
Green, D. A., 2004, BASI, 32, 335
Gull, S. F., 1973, MNRAS, 161, 47
Haberl, F., & Pietsch, W., 1999, A&AS, 139, 277
Hamaguchi, K., Maeda, Y., Matsumoto, H., Nishiuchi, M., Tomida, H., Koyama, K.,
Awaki, H., & Tsuru, G. T., 2000, Nucl. Instr. and Meth. A, 450, 360

BIBLIOGRAPHY

112

Hardy, T., Murowinski, R., & Deen, M. J., 1998, IEEE Trans. Nucl. Sci., 45, 154
Haynes, R. F. et al., 1990, In “Galactic and Intergalactic Magnetic fields” IAU
Symp., No.140, 205, eds. Beck, R., Kronberg, P., & Dordrecht, R. W. K.
Henize, K. G., 1956, ApJS, 2, 315
Hess, V. F., 1912, Phys. Z., 13, 1084
Higdon, J. C., & Lingenfelter, R. E., 2005, ApJ, 239, 867
Higdon, J. C., & Lingenfelter, R. E., 2005, ApJ, 628, 738
Higdon, J. C., Lingenfelter, R. E., & Ramaty, R., 1998, ApJL, 509, 33
Holland, A., Abbey, A., Lumb, D., & McCarthy, K., 1990, Proc. SPIE, 1344, 378
Holland, A. D., Turner, M. J., Abbey, A. F., & Pool, P. J., 1996, Proc. SPIE, 2808,
414
Hughes, J. P., Hayashi, I., & Koyama, K., 1998, ApJ, 505, 732
Itoh, M., Tanaka, R., & Komatsu, H. 2001, in ASP Conf. Ser. 251, New Century of
X-Ray Astronomy, ed. H. Inoue & H. Kuneida (San Francisco: ASP), 250
Jamison, K. A., Hall, J. M., Oltjen, J., Woods, C. W., Kauffman, R. L., Gray, T. J.,& Richard, P., 1976, Phys. Rev. A, 14, 937
Janesick, J. R., Soli, G., Elliott, T. S., & Collins, S. A., 1991, Proc. SPIE, 1447, 87
Jansen, F. et al., 2001, A&A, 365, L1
Kaastra, J. S. 1992, An X-Ray Spectral Code for Optically Thin Plasmas (Internal
SRON-Leiden Rep., updated ver. 2.0; Leiden: SRON)
Kampert, K. -H. et al., 2001, in 27th International Cosmic Ray Conference, ed.
R. Schlickeiser (Copernicus Gesellschaft), 240
Kassim, N. E.. Weiler, K. W., & Baum, S. A., 1991, ApJ, 374, 212
Kim, S., Chu, Y. -H., Staveley-Smith, L., & Smith, R. C., 1998, ApJ, 503, 729
Kitamoto, S. et al., 2004, Proc. SPIE, 5168, 376
Knoll, G. F., 1989, Radiation Detection and Measurement, John Wiley & Sons, New
York, pp.405
Kokubun, M. et al., 2006, PASJ, accepted (astro-ph 0611233)
Koo, B. -C., & Kang, J., 2001, MNRAS, 349, 983
Koyama, K., Petre, R., Gotthelf, E. V., Hwang, U., Matsura, M., Ozaki, M., &Holt, S. S., 1995, Nature, 378, 255

BIBLIOGRAPHY

113

Koyama, K., Maeda, Y., Sonobe, T., Takeshima, T., Tanaka, Y., & Yamauchi, S.,
1996, PASJ, 48, 249
Koyama, K. et al., 2006, PASJ, accepted
Koyama, K. et al., 2006, PASJ, accepted (astro-ph 0609215)
Krause, M. O., & Oliver, J. H., 1979, J. Phys. Chem. Ref. Data, 8, 329
LaMarr, B., Bautz, M. W., Kissel, S. M., Prigozhin, G., Hayashida, K., Tsuru, T.,& Matsumoto, H., 2004, Proc. SPIE, 5501, 385
Le Marne, A. E., 1968, MNRAS, 139, 461
Lesser, M., & Iyer, V., 1998, Proc. SPIE, 3355, 446
Liszt, H. S., & Spiker, R. W., 1995, ApJS, 98, 259
Long, K. S., Helfand, D. J., & Grabelsky, D. A., 1981, ApJ, 248, 925
Lortet, M. C., & Testor, G. 1984, A&A, 139, 330
Lucke, P. B., & Hodge, P. W., 1970, AJ, 75, 171
Lumb, D. H., & Nousek, J. A., 1990, Proc. SPIE, 1736, 138
Mac Low, M. -M., & McCray, R., 1988, ApJ, 324, 776
Maeda, Y., Koyama, K., Yokogawa, J., & Skinner, S., 1999, ApJ, 510, 967
Magnier, E. A., Chu, Y. -H., Points, S. D., Hwang, U., & Smith, R. C., 1996, ApJ,
464, 829
Mathewson, D. S., & Clarke, J. N., 1973, ApJ, 180, 725
Mathewson, D. S., Ford, V. L., Tuohy, I. R., Mills, B. Y., Turtle, A. J., &Helfand, D. J., 1985, ApJS, 58, 197
Mathewson, D. S., Ford, V. L., Tuohy, I. R., Mills, B. Y., Turtle, A. J., &Helfand, D. J., 1985, ApJL, 509, 33
Matsumoto, H. et. al., 2005, Nucl. Instr. and Meth. A, 541, 357
Matsumoto, H. et. al., 2006, Proc. SPIE, 6266, 120
McCray, R., & Snow, T. P. 1979, ARAA, 17, 213
McGee, R. X., Newton, L. M., & Butler, P. W., 1978, MNRAS, 183, 799
McKee, C. F., & Ostriker, J. P., 1977, ApJ, 218, 148
McKee, C. F., & Truelove, J. K., 1995, Phys. Rep., 256, 157

BIBLIOGRAPHY

114

McNamara, D. H., Madsen, J. B., Barnes, J., & Ericksen, B. F., 2000, PASP, 112,202
Meaburn, J., & Laspias, V. N., 1991, A&A, 245, 635
Meidinger, N., Schmalhofer, B., & Strüder, L., 2000, Nucl. Instr. and Meth. A, 439,
319
Mel’nik, A. M., & Efremov, Y. N. 1995, Astron. Lett., 21, 10
Mewe, R., Gronenschild, E. H. B. M., & van den Oord, G. H. J. 1985, A&AS, 62,197
Meynet, G., Maeder, A., Schaller, G., Schaerer, D., & Charbonnel, C., 1994, ApJS,
103, 97
Michael, E. et al., 2002, ApJ, 574, 166
Mills, B. Y., Turtle, A. J., Little, A. G., & Durdin, J. M., 1984, Australian J. Phys.,
37, 321
Mitsuda, K. et al., 2006, PASJ, acceptedMorrison, R., & McCammon, D., 1983, ApJ, 270, 119
Murakami, H., Koyama, K., Tsujimoto, M., Maeda, Y., & Sakano, M., 2001, ApJ,
550, 297
Nakajima, H. et. al., 2005, Nucl. Instr. and Meth. A, 541, 365
Nazé, Y., Chu, Y. -H., Guerrero, M. A., Oey, M. S., Gruendl, R. A., & Smith, R. C.,
2002, ApJ, 124, 3325
Nazé, Y., Antokhin, I. I., Rauw, G., Chu, Y. -H., Gosset, E., & Vreux, J. -M., 2004,
A&A., 418, 841
Odenwald, S. F., & Fazio, G. G., 1984, ApJ, 283, 601
Oey, M. S. & Massey, P., 1995, ApJ, 452, 216
Oka, T., Hasegawa, T., Hayashi, M., Handa, T., & Sakamoto, S., 1998, ApJ, 493,
730
Oort, J. H., Problems of Cosmic Aerodynamics (Central Air Document Office Dayton, OH)
Ostriker, J. P., & McKee, C. F., 1988, Reviews of M odern P hysics, 60, 1
Park, S., Burrows, D. N., Garmire, G. P., Nousek, J. A., McCray, R., Michael, E.,& Zhekov, S., 2002, ApJ, 567, 314
Pietrzyński, G., Gieren, W., Fouqué, P., & Pont, F., 2001, A&A, 371, 497

BIBLIOGRAPHY

115

Plucinsky, P. P. & Virani, S.N., 2000, Proc. SPIE, 4012, 681
Press, W. H., Teukolsky, S. A., Vetterling, W. T., & Flannery, B. P., 1992, Numerical
Recipes in C (2d ed.; Cambridge: Cambridge Univ. Press)
Prigozhin, G., Burke, B., Bautz, M. W., Kissel, S. M., LaMarr, B., & Freytsis, M.,
2004, Proc. SPIE, 5501, 357
Ptuskin, V. S., & Zirakashvili, V. N., 2003, A&A, 403, 1
Raymond, J. C., & Smith, B. W., 1977, ApJS, 35, 419
Read, A. M., & Ponman, T. J., 2003, A&A, 409, 395
Reynolds, S. P., 1998, ApJ, 493, 375
Reynolds, S. P., & Keohane, J. W., 1999, ApJ, 525, 368
Reynoso, E. M., & Mangum, J. G., 2000, ApJ, 545, 874
Richard, P., Oltjen, J., Jamison, K. A., Kauffman, R. L., Woods, C. W., &Hall, J. M., 1975, Phys. Letters A, 54, 169Różyczka, M., & Tenorio-Tagle, G., 1985, A&A, 147, 202
Russel, S. C., & Dopita, M. A., 1992, ApJ, 384, 508
Ryter, C. E., 1996, Ap&SS, 236, 285
Sasaki, M., Haberl, F., & Pietsch, W., 2000, A&AS, 143, 391
Schaller, G., Schaerer, D., Meynet, G., & Maeder, A., 1992, A&APS, 96, 269
Sedov, L. I., 1959, Similarity and Dimensional Methods in Mechanics (Academic
Press, New York)
Sembay, S., Abbey, A., Altieri, B., Ambrosi, R., Baskill, D, Ferrando, P., Mukerjee, K., Read, A.M., & Turner, M.J.L., 2004. Proc. SPIE, 5488, 264
Shklovskii, I. S., 1962, SvA, 6, 162
Shockley, E. G., & Read, J. M., 1952, Phys. Rev., 87, 835
Slane, P., Hughes, J. P., Edgar, R. J., Plucinsky, P. P., Miyata, E/, Tsunemi, H., &Aschenbach, B., 2001, ApJ, 548, 814
Smith, R. C., 1997, AJ, 114, 2664
Smith, D. R., Holland, A. D., Hutchinson, I., Abbey, A. F., Pool, P., Burt, D., &Morris, D., 2004, Proc. SPIE, 5501, 189
Stassinopoulos, E. G., & Barth, J. M., 1991, “Astro-D Charged Particle Environment
Study”, NASA report X-900-91-07

BIBLIOGRAPHY

116

Stephan, K., Reppin, C., Hirschinger, M., Maier, H. J., Frischke, D. Fuchs, D.,
Mueller, P., & Guertler, P., 1996, Proc. SPIE, 2808, 421
Sunyaev, R., & Churazov, E., 1998, MNRAS, 297, 1279
Takahashi, T., et al., 2006, PASJ, accepted (astro-ph 0611232)
Tenorio-Tagle, G., 1996, AJ, 111, 1641
Tenorio-Tagle, G., Bodenheimer, P., & Rozyczka, M., 1987, A&A, 182, 120
Tokunaga, A. T., 2000, Allen’s Astrophysical Quantities, 4th ed., ed. A. N.Cox (NewYork: Springer)
Tompsett, M. F. et al., 1975, IEEE Trans. ED, 22, 305
Truelove, J. K., & McKee, C. F., 1999, ApJS, 120, 299
Tsujimoto, T., Nomoto, K., Yoshii, Y., Hashimoto, M., Yanagida, S., & Thielemann, F. -K., 1995, MNRAS, 277, 945
Tsuru, T., Nishiuchi, M., Awaki, H., Hamaguchi, K., Koyama, K., Tomida, H.,
Sakano, M., ASTRO-E XIS Team, 2000, AdSpR, 25, 885
Uchiyama, Y., Takahashi, T., & Aharonian, F. A., 2002, ApJ, 571, 866
Ueno, M., Bamba, A., Koyama, K., & Ebisawa, K., 2003, ApJ, 588, 338
van den Bergh, S., & McClure, R. D., 1994, ApJ, 425, 205
Villa, G. E., et al., 1998, IEEE Trans. Nucl. Sci., 45, 921
Völk, H. J., Berezhko, E. G., & Ksenofontov, L. T., 2003, A&A, 409, 563
Wang, Q., & Helfand, D. J., 1991, ApJ, 373, 497
Watson, M. G. et al., 2001, A&A, 365, L51
Weaver, G. et al., 1997, ApJ, 218, 377
Weisskopf, M. C., Brinkman, B., Canizares, C., Garmire, G., Murray, S., & Van
Speybroeck, L. P. 2002, PASP, 114, 1
Will J. -M., Bomans D. J., & Dieball A., 1997, A&AS, 123, 455
Winkler, P. F., & Long, K. S., 1997, ApJ, 491, 829
Yamaguchi, R., Mizuno, N., Onishi, T., Mizuno, A., & Fukui, Y., 2001, PASJ, 53,
959

BIBLIOGRAPHY

117

Yamaguchi, H., Nakajima, H., Koyama, K., Tsuru, T. G., Matsumoto, H.,
Tsunemi, H., Hayashida, K., Torii, K., Namiki, M., Katsuda, S., Matuura, D.,& Miyauchi, T., 2006, Proceedings of the The X-ray Universe 2005 Ed. A. Wilson, pp.949
Yamaguchi, H. et al., 2006, Proc. SPIE, 6266, 121

Appendix A
Abbreviations
CP Charge Packet
CTI Charge Transfer Ineeficiency
CTICOL CTI measured for each column
CTIAVE CTI measured averaged among columns
CT Charge Trap
GC Galactic Center
LMC Large Magellanic Cloud
SB Superbubble
SNR Supernova Remnant
SN Supernova

118

Acknowledgment
First of all, I am deeply grateful to my supervisor, Prof. K. Koyama,
Assoc. Prof. T. Tsuru, and Res. Assoc. H. Matsumoto for their continual guidance, advice and support throughout my graduate career.
Their patience and sometimes strict guidance led me to complete thiswork.
I express my thanks to Dr. A. Bamba and M. Ueno, who carried out
part of the study of SB samples, the results of them are so amazing
and wonderful. Dr. A. Bamba, especially, kindly gave me advices andsuggestions on the LMC SNRs, which were so useful and helpful tome.
I am also grateful to my collaborators, Prof. H. Tsunemi, Assoc.
Prof. K. Hayashida, Res. Assoc. K. Torii, Mr. S. Katsuda, M. Shoji,
D. Matsuura, and T. Miyauchi for their great helps in conducting the
ground experiments at Osaka University.
I also thank all the members of the Cosmic-ray laboratory in Kyoto
University for the daily discussion and insightful comments, as well as
for the kind and heartful attitude to me.
I express my gratitude to all the members of the Suzaku XIS teamfor their support of this study. All the observational data were obtained by the members of the Suzaku team.
I also thank to Dr. M. Bautz, S. Kissel, and B. LaMarr XIS team.
for their hospitality during my visit to MIT in U.S.A..
I was financially supported by the Japan Society for the Promotionof Science.
Finally I express supreme thanks to my family, especially for myparents. They have provided me not only financial but also unstinted
emotional support.

119