Investigation of the Formation Process of Recombining Plasmas
in the Galactic Supernova Remnantswith the Suzaku X-Ray Observatory
Hideaki Matsumura

Department of Physics, Graduate School of Science, Kyoto University
Kitashirakawa Oiwake-cho, Sakyo, Kyoto, Kyoto, 606-8502, Japan
matumura@cr.scphys.kyoto-u.ac.jp

January 5, 2018

Abstract
In most of the supernova remnants (SNRs), plasmas are expected to have an ionization
degree below the equilibrium and to be in a ionization-dominant state. Recent observations with Japanese X-ray satellites, ASCA and Suzaku, revealed non-standard plasmas,
which have higher ionization degree than expected in a collisional ionization equilibrium.
The plasmas are thought to be in a recombination-dominant state, and thus called recombining plasmas (RPs). Such plasmas are not anticipated in a standard picture of the
SNR plasma evolution. Although some scenarios are proposed to explain the RPs, theformation process of the RP is not fully understood yet. In any scenario, a key seems to
be interactions between the SNR and ambient dense gas.
To understand the formation process of RPs, we perform spatially resolved spectroscopic analyses of X-ray emissions from the Galactic SNRs, G166.0+4.3, IC 443, W44, and
W49B with Suzaku. We compare spatial variations of physical parameters of the RPs with
distributions of the ambient gas. In all the remnants, the spectral analyses reveal goodcorrelations between locations of plasmas with lower electron temperatures and those of
surrounding dense molecular clouds. The lower electron temperatures of the RPs can be
explained well by the thermal conduction between an SNR plasma and a molecular cloud.
For a more quantitatively investigation, we estimated cooling timescales of the thermal
conduction in each remnant. In all the remnants, the conduction timescales are shorter
than a characteristic timescale for a plasma to reach collisional ionization equilibrium.
Therefore, we conclude that the RPs can be realized by the thermal conduction scenario.

Contents
1 Introduction

2

2 Overview of Supernova Remnants
2.1 Supernovae . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.1.1 Nucleosynthesis . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.1.2 Classification of Supernovae . . . . . . . . . . . . . . . . . . . . . .
2.2 Shock Wave Heating . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3 Evolution of SNRs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.1 Free expansion phase . . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.2 Adiabatic phase (Sedov phase) . . . . . . . . . . . . . . . . . . . . .
2.3.3 Radiative cooling phase . . . . . . . . . . . . . . . . . . . . . . . .
2.3.4 Disappearance phase . . . . . . . . . . . . . . . . . . . . . . . . . .
2.4 Classification of SNRs . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

4
4
4
6
9
11
11
11
13
14
14

3 Thermal X-ray Plasma in Supernova Remnants
3.1 Ionization State of Thermal Plasma . . . . . . . . . . . . . . . . . . . . . .
3.2 X-ray Emission from Thermal Plasmas . . . . . . . . . . . . . . . . . . . .
3.2.1 Continuum emission . . . . . . . . . . . . . . . . . . . . . . . . . .
3.2.2 Line emission . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3 Recombining Plasma in SNRs . . . . . . . . . . . . . . . . . . . . . . . . .
3.3.1 Discovery of recombining plasma . . . . . . . . . . . . . . . . . . .
3.3.2 Formation process of recombining plasma . . . . . . . . . . . . . . .
3.3.3 Toward the understanding of the origin of RPs . . . . . . . . . . . .

18
18
19
19
24
26
26
26
26

4 Instruments
4.1 Suzaku Satellite . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2 X-ray Telescope (XRT) . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.1 Mechanism and design . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.2 Performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.3 X-ray Imaging Spectrometer (XIS) . . . . . . . . . . . . . . . . . . . . . .
4.3.1 Overview of the XIS . . . . . . . . . . . . . . . . . . . . . . . . . .
4.3.2 Calibration; charge transfer ineﬃciency . . . . . . . . . . . . . . . .
4.3.3 Calibration; gain . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.3.4 Data reduction; flickering pixels . . . . . . . . . . . . . . . . . . . .
4.3.5 Non X-ray background (NXB) . . . . . . . . . . . . . . . . . . . . .

28
28
30
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31
31
32
34
36
39

ii

CONTENTS

iii

5 G166.0+4.3
5.1 Previous Works and Our Aims . . . . . . . . . . . . . . . . . . . . . . . . .
5.2 Suzaku Observation and Data Reduction . . . . . . . . . . . . . . . . . . .
5.3 Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3.1 Image . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3.2 X-ray background subtraction . . . . . . . . . . . . . . . . . . . . .
5.3.3 SNR spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.4 Discussion on G166.0+4.3 . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.4.1 Spatial structure . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.4.2 Ambient gas density and SNR age . . . . . . . . . . . . . . . . . . .
5.4.3 Origin of recombining plasma . . . . . . . . . . . . . . . . . . . . .

42
42
42
44
44
44
46
48
48
51
53

6 IC 443
6.1 Previous Works and Our Aims . . . . . . . . . . . . . . . . . . . . . . . . .
6.2 Suzaku Observation and Data Reduction . . . . . . . . . . . . . . . . . . .
6.3 Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.3.1 Image . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.3.2 SNR spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.4 Discussion on IC 443 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.4.1 Absorption by the molecular cloud . . . . . . . . . . . . . . . . . .
6.4.2 Origin of the CIE plasma . . . . . . . . . . . . . . . . . . . . . . . .
6.4.3 Origin of the recombining plasmas . . . . . . . . . . . . . . . . . . .

54
54
55
56
56
56
63
63
63
63

7 W44
7.1 Previous Works and Our Aims . . . . . . . . . . . . . . . . . . . . . . . . .
7.2 Suzaku Observations and Data Reduction . . . . . . . . . . . . . . . . . .
7.3 Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.3.1 Image . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.3.2 X-ray background subtraction . . . . . . . . . . . . . . . . . . . . .
7.3.3 SNR spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.3.4 Hard X-ray sources . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.4 Discussion on W44 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.4.1 Ambient gas and X-ray absorption . . . . . . . . . . . . . . . . . .
7.4.2 Origin of the CIE plasma . . . . . . . . . . . . . . . . . . . . . . . .
7.4.3 Origin of the recombining plasma . . . . . . . . . . . . . . . . . . .

66
66
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67
69
73
78
78
80
80

8 W49B
8.1 Previous Works and Our Aims . . . . . . . . . . . . . . . . . . . . . . . . .
8.2 Suzaku Observations and Data Reduction . . . . . . . . . . . . . . . . . .
8.3 Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
8.3.1 Image . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
8.3.2 SNR spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
8.4 Discussion on W49B . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

82
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84
85
85
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90

iv

CONTENTS

9 Discussion
9.1 Summary of the Results . . . . . . . . . . . . . . . . . . . . . . . . . . . .
9.1.1 Electron densities and recombination timescales . . . . . . . . . . .
9.1.2 Abundance patterns and progenitor masses . . . . . . . . . . . . . .
9.1.3 Spatial variation of kTe . . . . . . . . . . . . . . . . . . . . . . . . .
9.1.4 Spatial variation of ne t . . . . . . . . . . . . . . . . . . . . . . . . .
9.2 Rarefaction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
9.3 Thermal Conduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

92
92
92
93
94
96
97
98

10 Conclusions

104

A Spatially resolved spectroscopy of IC 443
106
A.1 Fitting results of the spectra in the regions A and B . . . . . . . . . . . . . 106
A.2 Fitting results of the spectra in the regions 1–20 . . . . . . . . . . . . . . . 107
B Spatially resolved spectroscopy of W44
113
B.1 Fitting results of the spectra in the regions 2–10 . . . . . . . . . . . . . . . 113

List of Figures
2.1 Scheme of onion-like structure . . . . . . . . . . . . . . . . . . . . . . . . .
2.2 Classification scheme of supernovae . . . . . . . . . . . . . . . . . . . . . .
2.3 Light curves for Type Ia, Ib, II-P, and II-L supernovae . . . . . . . . . . .
2.4 A schematic view of the shock wave . . . . . . . . . . . . . . . . . . . . . .
2.5 A schematic view of the SNR in the transition phase . . . . . . . . . . . .
2.6 Cooling coeﬃcient for an optically-thin thermal plasma . . . . . . . . . . .
2.7 X-ray and radio images of the shell-like SNR, Tycho . . . . . . . . . . . . .
2.8 X-ray and radio images of the plerionic SNR, 3C 58 . . . . . . . . . . . . .
2.9 X-ray and radio images of the composite SNR, G11.2-0.3 . . . . . . . . . .
2.10 X-ray and radio images of the mixd-morphology SNR, W44 . . . . . . . . .

6
8
8
9
12
13
14
15
16
16

3.1
3.2
3.3
3.4
3.5
3.6

Ion fractions as a function of the ionization temperatures . . . . . . . . . . 20
Ion fractions of Fe in non-ionization equilibrium states . . . . . . . . . . . 21
The emissivity of a pure silicon plasma in a recombined-dominant state . . 22
Emission mechanism for radiation recombination continua and cascade lines 23
Level scheme for a He-like ion . . . . . . . . . . . . . . . . . . . . . . . . . 24
Emission mechanism for the dielectronic recombination . . . . . . . . . . . 25

4.1
4.2
4.3
4.4
4.5
4.6
4.7
4.8
4.9

29
30
31
31
32
33
33
34

Schematic view of the Suzaku satellite . . . . . . . . . . . . . . . . . . . .
Schematic view of the XRT . . . . . . . . . . . . . . . . . . . . . . . . . .
Photograph of the XRT . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Eﬀective area of the XRT and XIS0 image . . . . . . . . . . . . . . . . . .
Vignetting of the XRT . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Picture and side view of the XIS . . . . . . . . . . . . . . . . . . . . . . . .
Schematic view of the XIS CCD . . . . . . . . . . . . . . . . . . . . . . . .
Schematic view of the XIS CCD . . . . . . . . . . . . . . . . . . . . . . . .
Trend of centroid energy and width of the Mn Kα lines for uncalibrated
data of the XIS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.10 Trend of centroid energy of the Mn Kα lines for calibrated data of the XIS
4.11 Trend of width of the Mn Kα lines for calibrated data of the XIS . . . . .
4.12 XIS0 Spectrum of the SNR E0102−72.3 . . . . . . . . . . . . . . . . . . .
4.13 Trend of centroid energies of lines in E0102−72.3 . . . . . . . . . . . . . .
4.14 Gain collection for W49B . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.15 Spectra of the SNR W49B before and after the gain correction . . . . . . .
4.16 Maps of the flickering pixels in each sensor . . . . . . . . . . . . . . . . . .
4.17 Spectra of the SNR G166.0+4.3 before and after the reduction of the flickering pixels . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
v

35
35
36
37
37
38
38
39
40

vi

LIST OF FIGURES
4.18 Spectra of the NXB in the XIS . . . . . . . . . . . . . . . . . . . . . . . . 41
4.19 Spectra of the NXB in the XIS . . . . . . . . . . . . . . . . . . . . . . . . 41
5.1
5.2
5.3
5.4
5.5
5.6
5.7
5.8
5.9

ROSAT image of G166.0+4.3 . . . . . . . . . . . . . . . . . . . . . . . . .
XIS image of G166.0+4.3 in the 0.5-2.0 keV band . . . . . . . . . . . . . .
XIS spectra of the IRAS 05262+4432 for a background spectra . . . . . . .
XIS spectra in the W and NE regions extracted from G166.0+4.3 . . . . .
NE spectrum and the plasma model with and without response matrixes
multiplied . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
XIS0+3 and XIS1spectra in the W and NE regions. . . . . . . . . . . . . .
XIS images of G166.0+4.3 in the Si-K and Fe-L bands . . . . . . . . . . .
Abundance ratios of Fe to Si . . . . . . . . . . . . . . . . . . . . . . . . . .
Fermi gamma-ray image in the 0.5–2 GeV band . . . . . . . . . . . . . . .

6.1
6.2
6.3
6.4
6.5
6.6

IC 443 image and spectrum taken from Yamaguchi et al. (2009) . . . . . . 55
Suzaku images of IC 443 in the energy band of 0.3–2.0 keV and 3.0–5.0 keV 57
XIS spectrum of the SE region of IC 443 . . . . . . . . . . . . . . . . . . . 58
Spectra and models in IC 443 . . . . . . . . . . . . . . . . . . . . . . . . . 60
Fitting results of spectral analysis in IC 443 . . . . . . . . . . . . . . . . . 62
Normalization ratio of the CIE/RPcold to the RPhot . . . . . . . . . . . . . 64

7.1
7.2
7.3
7.4
7.5
7.6
7.7
7.8
7.9

XIS image of W44 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
X-ray background spectra for the spectral analysis of W44 . . . . . . . . .
XIS0+3 spectra of the Whole region in the 0.6–10 keV band . . . . . . . .
XIS0+3 and XIS1 spectra of the Whole region with the best-fit models . .
XIS spectra and models in regions of hard X-ray sources in W44 . . . . . .
L–kTe relation for z < 0.2 . . . . . . . . . . . . . . . . . . . . . . . . . . .
Spectra in region 1 in the 5.8–7.3 keV band. . . . . . . . . . . . . . . . . .
Fitting results of spectral analysis in W44 . . . . . . . . . . . . . . . . . .
Distributions of the intensity ratio of the 12 CO(J = 2–1) to 12 CO(J = 1–0)
emissions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

68
70
71
72
74
77
78
79

8.1
8.2
8.3
8.4
8.5
8.6
8.7

Shocked H2 , FeII and X-ray color component image of W49B . . . . . . . .
13
CO(J = 1–0) integration intensity . . . . . . . . . . . . . . . . . . . . . .
Gamma-ray image of W49B with Fermi-LAT and H.E.S.S. . . . . . . . . .
XIS image of W49B in the 0.3–10 keV band . . . . . . . . . . . . . . . . .
XIS images of W49B in the 0.3–5 keV and 5–10 keV bands . . . . . . . . .
XIS spectrum of the Whole region of W49B . . . . . . . . . . . . . . . . .
Spectra and models in W49B . . . . . . . . . . . . . . . . . . . . . . . . .

83
83
84
85
86
87
88

9.1
9.2

Abundance ratios of G166.0+4.3, W44, and W49B . . . . . . . . . . . . . 95
Relationship between radii of the wind-blown bubbles and the stellar masses
for OB stars . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 97
Schema of the plasma structure in our calculation of the time developmentof kTe . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 99
Time evolution of the electron temperatures . . . . . . . . . . . . . . . . . 101
Conduction timescales as a function of the electron densities . . . . . . . . 102

9.3
9.4
9.5

43
44
45
46
47
49
51
52
52

81

LIST OF FIGURES
A.1
A.2
A.3
A.4
A.5

vii

Spectra in regions A and B of IC 443 . . . . . . . . . . . . . . . . . . . . . 106
Spectra in regions 1–4 of IC 443 . . . . . . . . . . . . . . . . . . . . . . . . 107
Spectra in regions 5–10 of IC 443 . . . . . . . . . . . . . . . . . . . . . . . 108
Spectra in regions 11–16 of IC 443 . . . . . . . . . . . . . . . . . . . . . . . 109
Spectra in regions 17 and 20 of IC 443 . . . . . . . . . . . . . . . . . . . . 110

B.1 Spectra in regions 1–4 of W44 . . . . . . . . . . . . . . . . . . . . . . . . . 113
B.2 Spectra in regions 5–10 of W44 . . . . . . . . . . . . . . . . . . . . . . . . 114

List of Tables
2.1

Classification of SNRs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 16

3.1
3.2
3.3

K-shell edge energies of He-RRC and H-RRC . . . . . . . . . . . . . . . . . 23
Centroid energies of K-shell lines of H-like and He-like ions. . . . . . . . . . 25
List of the SNRs where RPs are clearly or possibly detected with Suzaku
observations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 27

4.1

Overview of Suzaku capabilities . . . . . . . . . . . . . . . . . . . . . . . . 29

5.1
5.2
5.3

Observation log of G166.0+4.3 . . . . . . . . . . . . . . . . . . . . . . . . . 43
Best-fit model parameters of the X-ray background spectra for G166.0+4.3 45
Fit parameters of source regions of G166.0+4.3 . . . . . . . . . . . . . . . 50

6.1
6.2

Observation log of IC 443 . . . . . . . . . . . . . . . . . . . . . . . . . . . 55
Best-fit model parameters of the SE, NW, and NE spectra extracted fromIC 443 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 61

7.1
7.2
7.3
7.4

Observation log of W44 . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Best-fit model parameters of the background spectra for W44 . . . . . . .
Best-fit model parameters of the spectrum in the Whole region of W44 . .
Best-fit model parameters of the spectrum in the regions A, B, C and 1 . .

8.1
8.2

Observation log of W49B . . . . . . . . . . . . . . . . . . . . . . . . . . . . 84
Best-fit model parameters of the W49B spectrum . . . . . . . . . . . . . . 89

67
69
73
75

9.1 Summary of the results 1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 92
9.2 Summary of the results 2 . . . . . . . . . . . . . . . . . . . . . . . . . . . . 93
9.3 Initial parameters for the calculation of the time evolution of kTe . . . . . 99
9.4 List of timescales in each remnants . . . . . . . . . . . . . . . . . . . . . . 100
A.1 Best-fit model parameters of the spectra in regions A and B of IC 443 . . . 107
A.2 Best-fit model parameters of the spectra from regions 1–10 . . . . . . . . . 111
A.3 Best-fit model parameters of the spectra from regions 11–20 . . . . . . . . 112
B.1 Best-fit model parameters of the spectra from regions 2–10 . . . . . . . . . 115

viii

LIST OF TABLES

1

Chapter 1
Introduction
Supernova remnants (SNRs) are the objects resulting from an explosion of a star, i.e.,
supernova (SN). Elements up to Fe are created through stellar nucleosynthesis and SN
nucleosynthesis and are ejected into the interstellar space. A shock wave caused by an
SN explosion propagates in the interstellar space. The expanding shock wave supplies
thermal and non-thermal energy into the interstellar space. The shock heats the plasma
up to a temperature of ∼ 107 K. Non-thermal population of particles or cosmic rays
are widely believed to be generated through the diﬀusive shock acceleration mechanism
working at blast waves of SNRs. Thus, SNe and SNRs act as major sources of metal andenergy in galaxies.
X-ray observations provide a powerful probe to study physics regarding SNR plasmas,
which emit thermal electron bremsstrahlung and various characteristic emission lines fromions in the X-ray band. In an SNR plasma, ions are gradually ionized by collisions withelectrons. The plasma reaches a collisional ionization equilibrium after ∼ 105 yr. Since
most of the SNRs have an age of ≲ 104 yr, SNR plasmas are naturally expected to havean ionization degree below the equilibrium and to be in a ionization-dominant state.
It was found that many SNRs indeed have ionization-dominant plasmas through X-ray
observations.
Recent observations with Japanese X-ray satellites, ASCA and Suzaku, however, revealed peculiar plasmas in some SNRs, which have higher ionization degree than expected
in a collisional ionization equilibrium. The plasmas are thought to be in a recombinationdominant state, and thus called recombining plasmas (RPs). Such plasmas are not anticipated in the standard picture of the SNR plasma evolution described above. Althoughsome scenarios are proposed to explain the RPs, the formation process of the RP is not
fully understood yet. In any scenario, a key seems to be interactions between the SNR
and ambient dense gas.
In this thesis, aiming to understand the formation process of the RPs, we analyze
Suzaku data of four Galactic SNRs: G166.0+4.3, IC 443, W44, and W49B, and discuss
implications of the results. In particular, we carry out spatially resolved spectroscopy inorder to study the influence of interactions between the SNR shell and ambient gas on thecharacteristics of the SNR plasma. The overviews of SNRs and physics regarding SNR
plasmas are reviewed in Chapters 2 and 3, respectively. Chapter 4 gives a description of
the instruments used for the data analysis. The data analysis and results on G166.0+4.3,
IC 443, W44 and W49B are described in Chapters 5, 6, 7, and 8, respectively. In Chapter 9, the results are summarized and the origin of the RPs in the remnants are discussed.
2

3
The final conclusion of this thesis is given in Chapter 10.

Chapter 2
Overview of Supernova Remnants
Supernovae (SNe) are the energetic explosive events that occur at the last stage ofevolution of stars. An SN explosion releases kinetic energy of ∼ 1051 ergs, with an initial
velocity of ∼ 3 × 104 km s−1 . The expanding forward shock sweeps up the surrounding
interstellar medium (ISM) while a reverse shock propagates back into the SN ejecta.
The relics of the SNe, called Supernova remnants (SNRs), can be observed in various
energy bands such as radio, optical, X-ray, and gamma ray. SNRs provide us with information about the progenitors, explosion mechanisms, and evolution of the remnants. In
this chapter, we explain the basic characteristics of SNe and SNRs.

2.1

Supernovae

2.1.1

Nucleosynthesis

Stellar nucleosynthesis is a generic term of nuclear reactions that elements heavier than
H are generated in stars, but that at the time of supernova explosions is called supernova
nucleosynthesis and is distinguished. When a high-density part occurs in interstellar gas,
gravity in the part becomes strong. Surrounding gas are accumulated by the gravity andstellar temperature increase in inverse proportion to a radius of the star. A thermonuclear
fusion reaction will start when a distance between protons in the center is close to beyond
the Coulomb barrier which is the repulsion force between protons.
p-p chain process
The proton-proton chain (p-p chain) reaction starts at the phase of stellar evolution
when the core temperature reaches ∼ 107 K. The process is described as
H + 1H →2
H + 1H →3
He + 3 He →1

2

H + e+ + νe
3
He + γ4
He + 2 1 H.

The p-p chain is a main process in stars with masses < 2 M⊙ on their main sequence.
42.1. SUPERNOVAE

5

CNO cycle
When the core temperature reaches ∼ 2×107 K, the other reaction, the CNO (carbonnitrogen-oxygen) cycle becomes dominant;12

C + 1H
13
N
13
1
C+ H
14
N + 1H
15
O
15
1
N+ H

→
→
→
→
→
→

13

N+γ
C + e+ + νe
14
N+γ
15
O+γ
15
N + e+ + νe
12
C + 4 He.
13

He burning process
When H is depleted by the above nuclear fusion reactions, He burning begins in thestar’s core. Since 8 Be is an unstable nucleus, 12 C is finally synthesized by triple-alphaprocess;
He + 4 He →8
Be + 4 He →12
C + 4 He →48

Be + γC+γ16
O + γ.
12

The triple alpha process converts He to C and O.
Fusion of heavier elements
When the temperature exceeds ∼ 6 × 108 K in the cores of massive stars, carbonburning process takes place. The principal reactions are
C + 12 C →12
C + 12 C →1220

Ne + 4 He
23
Na + 1 H,

and when the temperature increases, Ne and O burning processes start. These processes
are described as
Ne + γ →20
Ne + 4 He →

16

→
→
→
→

32

20

O + 4 He
24
Mg + γ,and16

O + 16 O
16
O + 16 O
16
O + 16 O
16
O + 16 O

S+γ
S+n
31
P + 1H28
Si + 4 He,
31

respectively. Furthermore, 32 S, 36 Ar and 40 Ca are synthesized by the alpha capture of
28
Si.

6

CHAPTER 2. OVERVIEW OF SUPERNOVA REMNANTS

When the temperature exceeds ∼ 4 × 109 K, Si burning process starts. Si, S, Ar andCa capture 4 He and synthesize iron group elements such as Cr, Mn, Ti, Fe, Co, Ni and
Zn. In the complete Si burning process, most of Si burns out and a large amount of 56 Ni
is synthesized. 56 Fe and 56 Co are generated by the the beta decay of 56 Ni.
In a star whose mass is larger than ∼ 8 solar mass (M⊙ ), the C burning progresses
and produces O + Ne + Mg. If the mass exceeds ∼ 10 M⊙ , heavier elements from Si to
Fe are synthesized in order. Therefore, in a massive star, an onion-like structure formedat the final stage of its evolution (Figure 2.1).

H
He
C, O
O, Ne, Mg
SiFe

Figure 2.1: Scheme of onion-like structure in a massive star.

2.1.2

Classification of Supernovae

The fundamental classification scheme of SNe is traditionally based on their optical
spectra and light curves as shown in Figure 2.2. If an optical spectrum lacks absorption
lines of hydrogen Balmer series, the SN is classified as Type I. On the other hand, the SN
is classified as Type II when its spectrum contains the hydrogen lines (Minkowski, 1939,
1940). Type I SNe are sub-classificated according to the presence of Si and He absorption
feature (e.g., Elias et al., 1985; Wheeler et al., 1987). The SNe with strong silicon featureare classified as Type Ia. The SNe without strong silicon feature and with helium feature
are defined as Type Ib, whereas those without silicon and helium features are classified as
Type Ic. Subclasses of Type II SNe are distinguished by the shape of their light curves.
The SNe are defined as Type II-L if the shapes of the light curves are linear, whereas thelight caves are plateau in the Type II-P SNe (e.g., Barbon et al., 1973). The typical light
curves of each SN type are shown in Figure 2.3.
As can be seen from the definition, this classification dose not reflect explosion mechanisms of SNe. The progenitors of Type Ia SNe are presumed to be white dwarfs withnear-Chandrasekhar mass (∼ 1.4 M⊙ ), whereas Type II progenitors are considered to bemassive stars (M ≳ 10 M⊙ ). We describe details of the characteristics for Types Ia andII in the following. Although the spectra of Type Ib SNe lacks hydrogen absorption lines,
the property of their light curves is very similar to those of Type II SNe. Furthermore,Type Ib SNe are observed only in spiral galaxies and star forming regions. Therefore, the2.1. SUPERNOVAE

7progenitors of Type Ib SNe are presumed to be massive stars that have lost their hydrogenenvelopes. The progenitors of Type Ic SNe are also presumed to be massive stars, which
may be progenitors of gamma-ray bursts (GRBs).

Type Ia Supernovae
Type Ia SNe are observed in all types of galaxies including elliptical galaxies, where
no massive stars exist. The lack of hydrogen absorption lines in the Type Ia spectrasuggests that progenitors do not have outer hydrogen envelopes. As we mentioned, theprogenitor of Type Ia SNe is likely to be while dwarfs. When the mass of the white dwarf
reaches about the Chandrasekhar limit (1.4 M⊙ ), its gravity exceeds the electron degeneracy pressure, and the carbon burning is ignited as the core. Nuclear fusion advances
explosively, and the star is totally disrupted. This explosion does not leave a compact
object such as a neutron star or a black hole. Type Ia SNe mainly yield Fe-peak elements
and intermediate-mass elements Si to Ca. In the literature of nucleosynthesis models ofType Ia SNe, the heavy elements such as Si and Fe are selectively synthesized (Nomoto
et al., 1984; Iwamoto et al., 1999).
Two possible scenarios are considered for Type Ia SNe. One is the single degeneratescenario (Whelan & Iben, 1973), which assumes a white dwarf accreting from a nondegenerate companion. Webbink (1984) proposed the other scenario; the double degeneratescenario. They assumed the merger of two white dwarfs whose combined mass exceeds
the Chandrasekhar limit. One of the proofs to support the single degenerate scenario
is possible detection of Tyhco Brahe’s 1572 SN (Ruiz-Lapuente et al., 2004). They suggested that the SN was occurred by a binary system of the while dwarf and the red giant
star. On the other hand, the results of SN 2009dc, whose progenitor mass exceeds the
Chandrasekhar limit, supports the degenerate scenario (e.g., Yamanaka et al., 2009). The
origin of Type Ia SNe is still under debate.

Type II Supernovae
Type II SNe are the origin of massive stars with the main sequence masses of ≥ 10 M⊙ .
In these massive stars, the nucleosynthesis proceeds to produce Fe in the core. When thecentral temperature in the Fe core exceeds ∼ 3 × 109 K, photo-disintegration of Fe nuclei
takes place;
Fe → 13 4 He + 4 n − 124.4 MeV
4He → 2 1 H + 2 n − 28.3 MeV.

56

The star cannot support the gravity any more and collapses gravitationally because thethermal energy is absorbed in the reaction. The inflating matter released the gravitationalenergy which is ∼ 1053 ergs by the gravitational collapse in total. The gravitational energy
creates a shock wave which propagates outward. The kinetic energy of the explosion istypically ∼ 1051 ergs because neutrinos carried out 99% of the gravitational binding
energy. This explosion leaves a compact object such as a neutron star or a black hole.

8

CHAPTER 2. OVERVIEW OF SUPERNOVA REMNANTS

SN TypesnoType I
Type Ia

yes

Si ?

yes

Type IIfeature of its light curvenoType Ib

Type II-L

Type II-P

He ?
yes

thermonuclear

H?

othersnoType Ic
core collapse

Figure 2.2: Classification scheme of supernovae based on the features of the spectra and light
curves.

Figure 2.3: Light curves for Type Ia, Ib, II-P, and II-L supernovae (Wheeler & Harkness, 1990).

2.2. SHOCK WAVE HEATING

2.2

9

Shock Wave Heating

The supersonic SN explosion generates shock wave which compresses, heats up, andaccelerates the media. As shown in Figure 2.4, we take a coordinate to be the rest frameof the shock, and define upstream and downstream. The tangential component of the gas
velocity is assumed to be zero on both sides. Then, conservation of mass, momentum,
and energy across the shock front can be described asρ u vu = ρ d vd

(2.1)

ρu vu2 + pu = ρd vd2 + pd

(2.2)

1
1 2
vu + wu = vd2 + wd ,
(2.3)
2
2
respectively, where ρ, v, p, w are the mass density, velocity, pressure, and enthalpy of
upstream (subscript u) and downstream (subscript d), respectively. Considering idealgas, the enthalpy is given asγp
w = Cp T =,(2.4)(γ − 1)ρ
where Cp , γ are the heat capacity at constant pressure and the ratio of the heat capacity
at constant pressure to that at constant volume (= Cp /CV ), respectively.
From Equations 2.1, 2.2, 2.3 and 2.4, relations between the physical condition of the
upstream and the downstream, called Rankine-Hugoniot relations, are obtained asρuvd(γ + 1)pu + (γ − 1)pd
=
=
ρd
vu
(γ − 1)pu + (γ + 1)pd

vs ≠ 0
v2

SNR

(2.5)

vs` = 0v1 = 0ISM

Observer rest frame

vd

vu = vsρd , pd

ρu, puSNR(downstream)

ISM(upstream)

Shock rest frame

Figure 2.4: A schematic view of the shock wave in observer rest frame (left) and shock rest
frame (right). v, ρ, p, w, and T are the velocity, density, pressure, enthalpy, and temperature,
respectively. Subscripts s, u and d indicate shock, upstream and downstream, respectively.

10

CHAPTER 2. OVERVIEW OF SUPERNOVA REMNANTS

pu(γ + 1)ρu − (γ − 1)ρd=pd
(γ + 1)ρd − (γ − 1)ρu

(2.6)

Td
p d ρupd (γ + 1)pu + (γ − 1)pd==Tu
p u ρdpu (γ − 1)pu + (γ + 1)pd

(2.7)

1
{(γ − 1)pu + (γ + 1)pd }
2ρu

(2.8)

1 {(γ + 1)pu + (γ − 1)pd }2.
2ρu (γ − 1)pu + (γ + 1)pd

(2.9)

vu2 =
vd2 =

Assuming the strong shock (pd /pu ≫ 1), Equations 2.6, 2.7, 2.8 and 2.9 can be rewritten
as
ρu
vd
γ−1
=
=
(2.10)
ρdvu
γ+1Tu
pd γ − 1=Tdpu γ + 1(2.11)(γ + 1)pd
2ρu(2.12)(γ − 1)2 pd,
2(γ + 1)ρu

(2.13)

vu2 =
vd2 =

respectively. Since vu = vs (see Figure 2.4), we can derive the mean shock velocity at therest frame of the upstream using Equations 2.10 and 2.13 as
kTd = µmH2(γ − 1)pd=
µmH vs2 ,ρd(γ + 1)2(2.14)

where k, µ and mH are the Boltzmann constant, the mean atomic weight and the hydrogenmass, respectively. Assuming the non-relativistic monatomic gas (γ = 5/3), the densityand velocity ratios in the downstream to the upstream, and the post-shock temperature
are obtained asρd=4(2.15)
ρu
vd
1
=
vu
4
kTd =

3
µmH vs2 ,16(2.16)(2.17)

respectively. The post-shock temperature is determined only by the shock speed. In addition, the shock wave heats up all particles with diﬀerent masses to the similar velocities
determined by the shock speed. Therefore, the energy non-equipartition between protons
and electrons are expected. The shock wave heating can be understood as the energy
conversion from the kinetic energy to the internal (thermal) energy.

2.3. EVOLUTION OF SNRS

2.3

11

Evolution of SNRs

An SNR is generally composed of an expanding ejecta material (ejecta) and ISM swept
up by the shock wave. The evolution of SNRs can be divided into following several phases;
free expansion, adiabatic, radiative cooling, disappearance, and transition phases (see e.g.,
Cioﬃ et al., 1988; Truelove & McKee, 1999). These phases are mainly determined by theratio of a mass of the ejecta (Mej ) to that of the swept-up ISM (MISM ). For simplicity,
we assume that the density of the ambient gas is uniform.

2.3.1

Free expansion phase

After the SN explosion, the stellar materials are ejected around. The explosion generates a forward shock wave and compresses surrounding ISM. In the initial phase, theejecta expands keeping its kinetic energy and velocity constant because MISM is muchsmaller than Mej and is negligible. The expansion velocity (vs ) is ∼ 104 km s−1 , which
is much larger than the sound speed of ∼ 10 km s−1 in the surrounding ISM. Most of
the explosion energy (E) is released the kinetic energy while only 2–3% of its total isconverted into thermal energy. Therefore, vs can be described as√()1/2 (
)−1/2
2E
E
Mej
9
vs =
= 1.0 × 10
cm s−1 .
(2.18)
Mej
1051 erg
M⊙
The mass of the swept-up ISM is expressed by
4
MISM = πRs3 µmH n0 ,3

(2.19)

where Rs and n0 is the radius of the expanding shell and the hydrogen number density ofthe swept-up ISM, respectively. Since the free expansion phase continues until that MISM
is close to Mej , the timescale since the explosion (t) can be approximately estimated at(t ∼ 1.9 × 102E1051 erg)−1/2 (

Mej
M⊙

)5/6 (µ )−1/3 ( n0 )−1/3
yr
1.4
1 cm−3(2.20)

from an equation of Rs = vs t and Equations 2.18 and 2.19. The free expansion phasetypically continues for t ≲ 102−3 yr as shown in Equation 2.20.

2.3.2

Adiabatic phase (Sedov phase)

When MISM becomes much larger than Mej (typically ≳ 10 Mej ), the kinetic energy istransferred to the swept-up ISM. In a phase where energy loss by radiation cooling fromthe ISM can be negligible, the expansion is adiabatic, and therefore, this phase is called theadiabatic phase. In the adiabatic phase, we can be applied the Sedov solution to quantify
an estimate for the outcome of a point explosion; density, pressure, and temperature of
the gas, and the distribution of the shock velocity. Therefore, the phase is also called theSedov phase. The radius and velocity of the blast wave, and the mean temperature justbehind the shock front are described as

12

CHAPTER 2. OVERVIEW OF SUPERNOVA REMNANTS()2/5 ()1/5 (

n0 )−1/5cm
104 yr1051 erg
1 cm−3()−3/5 (
)1/5 (tdRs
En0 )−1/5
7
vs == 5 × 10cm s−1
4
51
−3
dt
10 yr
10 erg
1 cm()−6/5 (
)2/5 (tEn0 )−2/5
Ts = 3 × 106K,
104 yr1051 erg
1 cm−3
Rs = 4 × 10

19

t

E

(2.21)(2.22)(2.23)

respectively. As shown in those equations, the evolution during this phase is determined
only by t, E and n0 . Since the ISM density is much larger than that of the ejecta, X-ray
emission from the ISM is dominant at all the X-ray emission in the remnant. Therefore,
n0 and Ts can be estimated from the thermal X-ray spectrum of the SNR. The radius Rs is
also determined observationally if the distance to the remnant is known, and then, we canderive t and E from Equations 2.21 and 2.23. Until the end of the phase (t ∼ 103−4 yr),
∼ 70% of the initial explosion energy is converted into the thermal energy (Chevalier,1974).
In transition phase between the free expansion and the adiabatic phase, the outershock (forward shock) sweeps up the ISM and its speed decreases with time according
to Equation 2.22, whereas the inner ejecta keeps expanding at constant speed in Equation 2.18. The swept-up ISM pushes back on the ejecta and causes another shock wave
to propagate inward through the ejecta, called “reverse shock” (McKee, 1974). The forward shock heats surrounding ISM, whereas the ejecta is heated by the reverse shock.
Figure 2.5 shows a schematic view of the SNR in the transition phase.

blast
wave

forwardshock

reverseshock

shockedISM

unshocked
ejecta(invisible to X-ray
observations)shockedejectaFigure 2.5: A schematic view of the SNR in the transition phase between the free expansion
and the adiabatic phase.

2.3. EVOLUTION OF SNRS

2.3.3

13

Radiative cooling phase

As shown in Equations 2.22 and 2.22, vs and Ts decrease over time during the adiabatic
phase. When the radiative energy loss becomes large compared with the internal energyof the remnant, the evolution is no longer adiabatic. This phase is called the radiative
cooling phase. As shown in Figure 2.6, electrons recombine with the heavy elements, suchas C and O, with Ts ≲ 3 × 105 K, and a shell of the SNR becomes cooler and denser. The
shell expands conserving its momentum, and an equation of P V γ = const. is established,
where P and V are the mean pressure and the volume of the gas, respectively. Therefore,
assuming the non-relativistic monatomic gas (γ = 5/3), the time dependency of the shock
expansion is described as
Rs ∝ t2/7(2.24)vs ∝ t−5/7 .

(2.25)

McKee & Ostriker (1977) called this stage the pressure-driven snowplow phase.
When the temperature cools down further, the shock expands conserving its radial
momentum; Ms vs = const., where Ms is the mass of the shell because the pressure canbe ignored. In the stage called the momentum-conserving snowplow phase (Cioﬃ et al.,1988), the time dependency of the shock expansion is obtained as
Rs ∝ t1/4(2.26)vs ∝ t−3/4 .

(2.27)The radiative cooling phase continues for t ≳ 104 yr.

Figure 2.6: Cooling coeﬃcient for an optically-thin thermal plasma with solar abundances(Anders & Grevesse, 1989) as a function of plasma temperature (Sutherland & Dopita, 1993).

14

2.3.4

CHAPTER 2. OVERVIEW OF SUPERNOVA REMNANTS

Disappearance phase

After the shock velocity drops to the sound speed in the ISM, the SNR lose its boundary and goes out into the ISM. The age of SNRs in the disappearance phase is ∼ 105 –106 yr.

2.4

Classification of SNRs

In most of SNRs, their X-ray morphologies are almost the same as those in the radio
band. SNRs have been classified into broad categories of shell-like, plerionic (Crab-like)and composite, based on their radio morphologies. Shell-like SNRs are characterized bylimb-brightened morphology in radio and X-ray bands as shown in Tycho’s SNR (Figure 2.7). The radio emission is mainly produced by synchrotron radiation from electrons
which is accelerated up by a blast wave. The X-ray emission is composed of thermal
and/or non-thermal X-rays. The origin of the non-thermal X-ray emission is the same as
the radio emissions but the energy of the accelerated electrons is the TeV order. We describe the origin of the thermal X-rays in Capter 3. Most of the shell-like SNRs are young
SNRs in the free expansion or adiabatic phases such as Tycho, Kepler and Cassiopeia A.
VLA radio image (1.4 GHz)

Chandra X-ray image (1.6-2.0 keV)

Figure 2.7: X-ray and radio images of the shell-like SNR, Tyhco. (right) X-ray image in theenergy band of 1.6-2.0 keV with Chandra (Lu et al., 2015). (left) 1.4 GHz radio continuum
image from with the Very Large Array (VLA; Credit Radio: NSF/NRAO/VLA).

The plerionic SNRs have center-filled emissions in both the radio and X-ray bands(see emissions from 3C 58 in Figure 2.8). Crab Nebula is the most famous remnant of
this type, and therefore, the plerionic SNRs are also called Crab-like SNRs. The originof the central emission is a pulsar wind nebula (PWN) powered by a rapidly rotating
neutron star (pulsar). The energy loss of the pulsar causes relativistic electrons andpositrons which radiates a synchrotron radiation by interacting with the magnetic fieldof the pulsar. If the center-filled emission from a PWN coexists with a shell emission as
those of the shell-like SNRs, the remnants are categorized into composite SNRs. We show

2.4. CLASSIFICATION OF SNRS

15

the emissions from a typical composite SNR G11.2-0.3 in Figure 2.9. These remnantsassociating with a neutron star account for ∼ 15% of all X-ray-detected Galactic SNRs.
Chandra X-ray image (0.5-10 keV)VLA radio image (1.4 GHz)

Figure 2.8: X-ray and radio images of the plerionic SNR, 3C 58 (Slane et al., 2004). (upper) Xray image in the energy band of 0.5-10.0 keV with Chandra. (bottom) 1.4 GHz radio continuum
image from with VLA.

Rho & Petre (1998) claimed that some SNRs have a radio shell with a center-filled
X-ray morphology. They classified them into mixed-morphology SNRs. Figure 2.10 shows
X-ray and radio images of a typical mixd-morphology SNR, W44. Rho & Petre (1998)
also claimed that ∼ 25% of the X-ray-detected SNRs are categorized into the mixedmorphology class. While the formation mechanism of the mixed-morphology SNRs has
not been fully understood, two scenarios are mainly considered. One is the evaporating
cloudlet scenario (White & Long, 1991), and the other is the thermal conduction scenario(Cox et al., 1999; Shelton et al., 1999). White & Long (1991) suggested that many small
clouds in the center of the remnant are heated by the blast wave and evaporate, andthen, emit X-rays. Cox et al. (1999); Shelton et al. (1999) considered that outer X-ray
plasma are cooled by the thermal conduction between the plasma and molecular clouds
associated with the remnant, and the center-filled X-ray morphology are produced. In
either scenario, interactions between the blast wave and the clouds play an important rolein the formation of the X-ray morphology. Most mixed-morphology SNRs are actuallyassociated with molecular clouds as indicated by CO line emissions or OH (1720 MHz)masers (e.g., Lazendic & Slane, 2006).

16

CHAPTER 2. OVERVIEW OF SUPERNOVA REMNANTS

VLA radio image (8.4 GHz)

Chandra X-ray image (3.3-8.1 keV)

Figure 2.9: X-ray and radio images of the composite SNR, G11.2-0.3 (Borkowski et al., 2016).
(right) X-ray image in the energy band of 3.3-8.1 keV with Chandra. (left) 8.4 GHz radio
continuum image from with VLA.

Figure 2.10: X-ray and radio images of the mixd-morphology SNR, W44. Suzaku X-ray image
in the energy band of 0.3-10.0 keV with VLA 1.5 GHz radio contours.

Table 2.1: Classification of SNRs.

Shell-like SNRs
plerionic SNRs
Composite SNRs
Mixed-morphology SNRs

Radio
shell-like center-filled⃝
××⃝ (PWN)⃝⃝ (PWN)⃝×

X-rayshell-like center-filled⃝
××⃝ (PWN)⃝⃝ (PWN)×⃝

2.4. CLASSIFICATION OF SNRS

17

Chapter 3
Thermal X-ray Plasma in Supernova
Remnants
3.1

Ionization State of Thermal Plasma

Thermal plasma is classified into three states based on relation between its electron
temperature (kTe ) and ionization temperature (kTz ). The relation in each state is described as(1) kTe > kTz ; Ionization-dominant state,(2) kTe = kTz ; Collisional ionization equilibrium (CIE) state,(3) kTe < kTz ; Recombination-dominant state,
respectively. The ionization state of SNR plasmas changes during their evolution. The
shock wave heating converts the translational motion into the thermal random motion(§2.2).
The temperature of the particle behind the shock wave increases in direct proportion
to the product of its mass and the square of the shock velocity (kTi ∝ mi vs2 ). Since
an ion mass is much larger than the electron mass (me ), the ion temperature (kTi ) and
kTe is non-equilibrium (kTi ∼ 2000kTe ). Then, the temperatures become equilibrium by
the particle interaction with the timescale of the electron-electron and electron-protoninteraction described as()3/2 ()
ne )−1kT
ln Λ
τee ∼ 4.9 × 10
1 cm−3
1 keV
30.9
( n
)−1 ( kT )3/2 ( ln Λ )pτep ∼ 3.1 × 1011
1 cm−3
1 keV
30.9
8

(

s

(3.1)

s

(3.2)

where ln Λ is the Coulomb logarithm (Zeldovich & Raizer, 1966).
Following Masai (1994), the timescale to reach an ionization equilibrium is calculatedasτion =j=1∑
Z(ne Sj )−1 =(( n )−1ne )−1e−112(min(S
))
∼
10
i
−3
1 cm
1 cm−3
18

s

(3.3)

3.2. X-RAY EMISSION FROM THERMAL PLASMAS

19

where SZ is the ionization rate coeﬃcient from the ion of Z. From comparison of thesetimescales, it is interpreted that the protons heated the electrons, and then the electrons
ionize the ions.
We show ion fractions of O, Ne, Mg, Si, S and Fe in the CIE state as a function of kTz
in Figure 3.1. The fraction of highly ionized ions increases as the kTz rises, and therefore,a weighted centroid-energy of emission lines becomes high. As shown in Figure 3.2,
the fractions in non-equilibrium plasmas are significantly diﬀerent with that in the CIE
plasma.
In general, ages of old SNRs are comparable to the ionization timescale, so that theirplasmas are in a CIE state. The plasmas in young and middle age SNRs are in a ionizingdominant (underionized) state (called IP; ionizing plasma). However, in some SNRs,
recombination-dominant plasmas (called RP; recombining plasma) are discovered (e.g.,Kawasaki et al., 2002, 2005; Ozawa et al., 2009b; Yamaguchi et al., 2009). In order torealize the recombination-dominant state, special gas-environments around the SNR are
required. We describe the formation process of the RPs in §3.3.

3.2

X-ray Emission from Thermal Plasmas

Here, we show mechanisms of thermal X-ray emissions, composing of continuum andline components, from SNR plasmas.

3.2.1

Continuum emission

The continuum emission from the SNR plasma consists of three components; bremsstrahlung,radiative recombination continuum (RRC), and two-photon emission due to free-free, freebound, and bound-bound transitions of electrons, respectively.
Bremsstrahlung
When a free electron in the plasma is deflected by an ion, it decelerates and produces
electromagnetic radiation called bremsstrahlung. The emissivity of the bremsstrahlung
for a electron with the Maxwell-Boltzmann distribution is given by
ϵﬀ

(

)1/2

( −E )
Te−1/2 Zi2 ne ni exp
gﬀ
kTe( −E )= 6.8 × 10−38 Te−1/2 Zi2 ne ni exp
gﬀ ergs−1 cm−3 Hz−1 ,
kTe
32πe6
=
3me c3

2π
3kme

(3.4)
(3.5)

where e, c, k, me , Z, ne , ni , E, and gf f are the electron charge, light speed, Boltzmann
constant, electron mass, atomic number, densities of electron and ion, emitted photon
energy, and velocity averaged Gaunt factor (e.g., Rybicki & Lightman, 1979), respectively.
For a typical thermal plasma emitting X-rays, the Gaunt factor is described as
(
gﬀ =

3 kTe
π E

)1/2
,

(3.6)

CHAPTER 3. THERMAL X-RAY PLASMA IN SUPERNOVA REMNANTS

100

O

Ion concentration

Ion concentration

20

10

1
1

Ionization temperature (keV)

10

Mg

0.01

Ion concentration

Ion concentration

0.1

1

1

10

0.1

1

10

0.1

1

10

10

1
0.01

0.1

1

Ionization temperature (keV)

10

S

0.01

Ion concentration

Ion concentration

0.1

Ionization temperature (keV)

Si

100

10

100

10

1
0.01

100

Ne

100

Fe

100

10

1

Ionization temperature (keV)

10

1
0.01

0.1

1

Ionization temperature (keV)Bared
H-likeHe-like

Li-likeBe-likeB-like10

0.01

C-likeN-likeO-likeIonization temperature (keV)

F-likeNe-likeNa-likeMg-likeAl-likeLowersFigure 3.1: Ion populations as a function of the ionization temperatures for O, Ne, Mg, Si, S
from Arnaud & Rothenflug (1985) and Fe from Arnaud & Raymond (1992). The figure is takenfrom Sawada (2011).

(a)Ar-like

0.01

H-like

0.01

0.1

1

10

1010
1011
n e t (cm−3 s)1012

kTe (keV)(b)

Ionizing Plasma(c)Recombining Plasma

0.01

Fraction
0.1

1

Fraction
0.1

1

21

Bare

He-likeNe-likeFraction
0.1

13.2. X-RAY EMISSION FROM THERMAL PLASMAS

109

Figure 3.2: (a) Ion population of Fe in a CIE plasma as a function of the electron temperaturekTe . Orange, magenta, red, blue, and green lines represent Ar-like, Ne-like, He-like, H-like, andbare Fe, respectively. (b) Ion population of Fe in an kTe = 5 keV ionizing plasma as a functionof the ionization timescale ne t. (c) Same as (b) but for recombining plasma where the initialand current electron temperatures are assumed to be 30 keV and 0.3 keV, respectively. The
figure is taken from Smith et al. (2014).

and therefore, the spectral shape of the bremsstrahlung only depends on Te . Equation 3.4
can be approximated for each relation between E and kTe as 0(E ≪ kTe )
 E
−0.4
E
(E ≲ kTe )
(3.7)
ϵﬀ ∝

E
exp (− kTe ) (E > kTe ).
Thus, the bremsstrahlung spectrum has the energy cut-oﬀ near E ∼ kTe as shown in
Figure 3.3, and we can measure Te from the observed spectrum of hot plasmas.
Radiative recombination continuum
RRC is electromagnetic radiation emitted when a free electron recombines with an ion.
Figure 3.4 shows an emission process of the RRC. The energy of emitted photon is thesum of the kinetic energy of the free electron and the binding energy of the recombined
electron. The spectrum of the RRC has an energy edge depending on the binding energyof the recombined electron as shown in Figure 3.3. The emissivity of the RRC for electrons
with the Maxwell-Boltzmann distribution is given by
)√
( E−E
)
(E − E1edge
edge
σnrec (E − Eedge ) exp −, (3.8)ϵbf = 4ne nZ,j+1 E
kTe
2πme kTe
kTe
where nZ,j+1 , Eedge , and σnrec are the density of the ion with the atomic number Z and theionization state j+1, the binding energy of the recombined electron, and the recombination
cross section to the level n, respectively.

22

CHAPTER 3. THERMAL X-RAY PLASMA IN SUPERNOVA REMNANTS

Figure 3.3: The emissivity of a pure silicon plasma in a recombined-dominant state (kTe =1 keV, ne t = 5 × 1010 cm−3 s). Shown are the contributions of bremsstrahlung (red dotted line),radiative recombination continuum (red dashed line) and two-photon emission (red solid line).
The figure is taken from Vink (2012).
When the electron temperature is much lower than the binding energy (kTe ≪ Eedge ),Equation 3.8 can be approximated as
( E−E
)
edge
ϵbf ∝ exp −(E ≥ Eedge )kTe= 0(E < Eedge ).

(3.9)(3.10)In this dissertation, we denote the RRC emitted by the recombination of H-like ionsto the ground state of He-like ions as “He-RRC”, and call that of bare ions to the groundstate of H-like ions “H-RRC”. We summarize Eedge of He-RRC and H-RRC for heavy
elements in Table 3.1.

Two-photon emission
When an electron is excited from 1s2 1 S0 to 1s 2s 1 S0 states by a free electron, theelectron cannot transit directly from 1s 2s 1 S0 to 1s2 1 S0 because of a selection rule. Ina low-density environment, second collision does not happen easily, and therefore, theelectron transits while emitting two photons. Although the sum of the energy of the two
photons is conserved, the energy of the each photon is not uniquely determined. Therefore,
a spectrum of the two-photon emission is a continuum as shown in Figure 3.3.

3.2. X-RAY EMISSION FROM THERMAL PLASMAS

freeelectron-

Energy

-

-

n=∞
n=4
n=3

23

RRC
( n=2)

RRC( n=3)

RRC( n=4)

n=2
H-RRC( n=1)

Lyα(n=2 1)Lyβ(n=3 1)Lyγ(n=4 1)n=1
Figure 3.4: Schematic illustration of emission mechanism for radiation recombination continua(red arrows) and cascade lines (blue arrows).

Table 3.1: K-shell edge energies (Eedge ) of He-RRC and H-RRC. These values correspondto the K-shell binding potential of He-like and H-like ions, respectively.
Z

Element

12
14
16
18
20
26

Mg
Si
S
Ar
Ca
Fe

Eedge (keV)He-RRC H-RRC
1.763
1.958
2.439
2.666
3.225
3.482
4.121
4.406
5.128
5.440
8.830
9.194

24

CHAPTER 3. THERMAL X-RAY PLASMA IN SUPERNOVA REMNANTS

3.2.2

Line emission

The line emission is due to bound-bound transitions with diﬀerent energy levels forelectrons. For a hydrogen atom, the energy of a photon is described as()
1
1
E = Ry−,(3.11)
n2 n′2where Ry , n, and n′ are the Rydberg constant (13.6 eV), and the principal quantum
number of before and after the transition, respectively. For heavy H-like elements withthe atomic number of Z, that is given by
(
)
1
1
2
E ∼ Z Ry.
(3.12)−
n2 n′2The emission lines of H-like Kα (2p → 1s), H-like Kβ (3p → 1s), and H-like Kγ (4p → 1s)
are called Lyα, Lyβ, and Lyγ lines, respectively.
In the case of He-like ions, Kα emission are composed of resonance (1s2p 1 P1 →2 1
1s S0 ), forbidden (1s2s 3 S1 → 1s2 1 S0 ), and inter-combination (1s2p 3 P2,1 → 1s2 1 S0 )transition lines (Figure 3.5). The emission lines of He-like Kα (n = 2 → 1), He-like
Kβ (n = 3 → 1), and He-like Kγ (n = 4 → 1) are called Heα, Heβ, and Heγ lines,
respectively. We show centroid-energies of K-shell lines in Table 3.2.

1P
1
1

S0

7
6

3P
2
3P
3P1
0

5
4
3

2-photons

1

w

x
y

shell n=2
3S
1

2

z
1

S0

shell n=1 (ground)Figure 3.5: Level scheme for a He-like ion. Transitions of w, (x, y), and z indicate that ofresonance, inter-combination, and forbidden lines, respectively. (Porquet et al., 2001)
As radiation processes with the bound-bound transitions in SNRs, the photoelectriceﬀect and the collisional excitation are well known. However, in a recombination-dominantstate, the dielectronic recombination and the cascade decay are not negligible.

3.2. X-RAY EMISSION FROM THERMAL PLASMAS

25

Table 3.2: Centroid senergies (eV) of K-shell lines of H-like and He-like ions.
Lyα

Lyβ

Lyγ

Heα
Heβ
res.∗ frb.∗ int.∗
O
654
775
817
574
561
569
665
Ne 1022 1211 1277
922
905
915 1072
Mg 1472 1745 1840 1352 1330 1343 1579
Si
2006 2377 2506 1865 1840 1854 2183
S
2623 3107 3277 2461 2431 2447 2884
Ar 3323 3936 4151 3140 3104 3124 3685
Ca 4106 4864 5130 3908 3845 3892 4582
Fe 6966 8266 8732 6702 6641 6670 7798
Ni 8077 9590 10107 7806 7744 7766 9190
∗ res., frb., and int. indicate resonance, forbidden, andinter-combination lines, respectively.

Heγ
698
1127
1660
2294
3033
3875
4819
8217
9680

Dielectronic recombination
Dielectronic recombination is a radiative process when a free electron is captured
and in the end a photon is emitted. When a free electron recombines with an ion, a coreelectron is subsequently excited. The electron returns to a lower level and emits a satellite
line at a slightly lower energy than the normal transition. Figure 3.6 shows the emission
mechanism for the dielectronic recombination.
freeelectron-

Energy

n=∞
n=4
n=3

-

n=2

n=1

-

-

-

-

-

-

Figure 3.6: Schematic illustration of emission mechanism for the dielectronic recombination.

Cascade decay
Cascade decay is accompanied by the recombination of electrons into higher excitedlevels in the free-bound transitions (Figure 3.4). In plasmas with a recombinationdominant state, the intensity ratio of the cascade decay to the other normal transitions
is higher than that in plasmas with a state of collisional ionization equilibrium.

26

CHAPTER 3. THERMAL X-RAY PLASMA IN SUPERNOVA REMNANTS

3.3

Recombining Plasma in SNRs

3.3.1

Discovery of recombining plasma

Kawasaki et al. (2002) reported a first observation of RPs in SNRs. They analyzed
ASCA data of the mixed-morphology SNR IC 443 and measured kTz based on the H-liketo He-like Kα intensity ratios of Si and S, and compared them with kTe . They found thatkTz is significantly higher than kTe , and suggested that the plasma is overionized. In theobservations of IC 443 (Yamaguchi et al., 2009) and W49B (Ozawa et al., 2009b) with
Suzaku, they discovered RRCs, which provides clear evidence that the plasmas are inan extremely recombining state. Subsequent Suzaku observations revealed RPs in othermixed-morphology SNRs in the Galaxy as well as in the Large Magellanic Cloud. We summarized the SNRs where RPs are clearly or possibly detected with Suzaku observationsin Table 3.3.

3.3.2

Formation process of recombining plasma

The formation process of RPs has not been fully understood yet, but two scenarios aremainly discussed. One is the rarefaction scenario (Itoh & Masai, 1989). If a supernova
explodes in a dense circumstellar matter (CSM) around the massive progenitor, ejecta
and the CSM are shock-heated and quickly ionized because of the high density. When theshock breaks the CSM out to a lower density ISM, kTe is decreased by adiabatic cooling.
Kawasaki et al. (2002) proposed the other scenario, the thermal conduction scenario. If an
SNR shock is interacting with a molecular cloud, thermal conduction occurs between the
SNR plasma and the cloud, and, therefore, kTe drops. Since the recombination timescalesof ions are generally longer than the conduction timescales, kTz cannot follow the decreaseof kTe , and RPs can be realized.
Both of the scenarios require special gas-environments around the remnants; a dense
CSM and a molecular cloud for the rarefaction and thermal conduction scenarios, respectively. In order to investigate the origin of RPs, we need to obtain spatial distributions ofphysical parameters of the thermal plasmas (e.g., kTe , ne t) and to compare the parameters
and the gas environments.

3.3.3

Toward the understanding of the origin of RPs

In order to understand the origin of RPs, we used Suzaku data of the Galactic SNRs
G166.0+4.3, IC 443, W44 and W49B, and perform spatial resolved spectroscopy. These
SNRs are known to have unique ambient gas through observations in the radio, infrared, X-ray and/or gamma-ray bands. In some of the remnants, the previous observations
already reported the detection of RPs. The sizes of the SNRs allows us to divide theremnants into some regions by using the Suzaku telescope. The high sensitivity and thehigh energy resolution of Suzaku can detect RRCs of heavy elements and enable us toaccurately determine the X-ray spectral features of the thermal plasma.

W49B
IC 443 (1)IC 443 (2)′′G359.1-0.5
W28W44
G346.6-0.2
3C391 (1)3C391 (2)
CTB 37A
G290.1-0.8
N49
Kes 17
G166.0+4.3
3C400.2

SNR

kTinit(keV)––
5 (fixed)′′–
3 (fixed)
1.00 ± 0.05
5 (fixed)–
1.8 +1.6−0.6
5 (fixed)
1.7 +1.2−0.3
11 ± 15 (fixed)
3 (fixed)
3.15 +0.02
−0.01

kTz
(keV)
2.46 +0.08−0.07∼ 1.0(Si), ∼ 1.2(S)––0.77 +0.09
−0.08
–
–
–
–
–
–
–
–
–
–
–

kTe
(keV)
1.52 +0.01−0.02
0.61 +0.03−0.02
0.16–0.28 keV (RPcold )0.48–0.67 keV (RPhot )0.29 ± 0.02
0.40 +0.02−0.03
0.47 ± 0.01
0.30 +0.03−0.01
0.58 ± 0.01
0.18 +0.03−0.05
0.49 +0.09−0.06
0.45 +0.02−0.01
0.62 ± 0.01
0.73 ± 0.03
0.46 ± 0.03
0.71 +0.03
−0.02
–
–
3.8–15.1′′–
6.3
6.2 +0.9−0.4
4.8 +0.1−0.4–
14.0 +1.5−2.2
1.3 +0.3−0.1
12.2 +1.3−0.9
9.8 ± 1.1
1.6 +0.6−0.3
6.1 +0.5−0.4
1.8 +0.2
−0.5

ne t
11
(10 s cm−3 )
Yes
Yes
Yes
′′
Yes
Yes
Yes
Yes
Yes
Yes
Yes
Possible
Yes
Yes
Possible
Possible

Cloud Interaction

Ozawa et al. (2009b)
Yamaguchi et al. (2009)
Matsumura et al. (2017b)′′
Ohnishi et al. (2011)
Sawada & Koyama (2012)
Uchida et al. (2012)
Yamauchi et al. (2013)
Ergin et al. (2014)
Sato et al. (2014)
Yamauchi et al. (2014)
Kamitsukasa et al. (2015)
Uchida et al. (2015)
Washino et al. (2016)
Matsumura et al. (2017a)
Ergin et al. (2017)

Reference

Table 3.3: List of the SNRs where RPs are clearly or possibly detected with Suzaku observations.

3.3. RECOMBINING PLASMA IN SNRS
27

Chapter 4
Instruments
4.1

Suzaku Satellite

Suzaku is the fifth in a series of Japanese X-ray astronomy satellites devoted to observations of celestial X-ray sources (Mitsuda et al., 2007). It was launched by the Japan
Aerospace Exploration Agency (JAXA) aboard an M-V launch vehicle from JAXA’s Uchinoura Space Center on July 10, 2005. The spacecraft was put into a near-circular orbit
at a ∼ 570 km altitude with an inclination angle of ∼ 31◦ and an orbital period of ∼ 96
min. Figure 4.1a shows a schematic view of the Suzaku satellite in orbit. The length of
the spacecraft is 6.5 m after the deployment of the extensible optical bench (EOB) andthe total weight is 1706 kg. Suzaku had been operated for ∼ 10 yr, which exceeds its
planned lifetime of ∼ 2 yr. Because of problems on its batteries and attitude control, theoperation of Suzaku has been finished on June 1, 2015.
Suzaku carried two imaging detectors: the X-ray Imaging Spectrometer (XIS; Koyama
et al., 2007) and the X-ray Spectrometer (XRS; Kelley et al., 2007) combined with the Xray Telescopes (XRT; Serlemitsos et al., 2007) and a non-imaging hard X-ray detector: theHard X-ray Detector (HXD; Kokubun et al., 2007; Takahashi et al., 2007). Figure 4.1b andTable 4.1 present a side view and the specifications of the instruments and the telescopeson Suzaku, respectively. The XIS consists of four cameras with X-ray charge-coupleddevices (CCDs) named XIS0, 1, 2 and 3. The CCDs of XIS0, 2 and 3 are front-illuminated(FI) devices whereas that of XIS1 is a back-illuminated (BI) CCD. XIS2 and a part of
XIS0 have not been functioning since November 2006 and June 2009, respectively (Suzaku
XIS documents in 2007 and 2010). The HXD covers the bandpass of 10–600 keV. The
XRS is the first micro-calorimeter system for X-ray astronomy but it failed to work after
August 8, 2005 because of a liquid helium coolant loss. In this dissertation, we useddata of the XIS so that we describe the details of the XRT and the XIS in the following
sections.

28

4.1. SUZAKU SATELLITE

(a)

29

(b)

XRT-S

Sun shade
XRT-I(4 units)

5.4 m

EOB
Star
tracker
Side panel
Gyroscopes(8 panels in total)

XRS
6.5 m

HXD
Base
panel

XIS
(4 units)

Figure 4.1: Schematic view of the Suzaku satellite (Mitsuda et al., 2007).

Table 4.1: Overview of Suzaku capabilities (Mitsuda et al., 2007).
XRT

Focal length
Field of viewEﬀective area
Angular resolution
XIS
Field of view
Bandpass
Pixel grid / size
Energy resolution
Eﬀective area(including the XRT)HXD Field of view
Bandpass

4.75 m17′ at 1.5 keV, 13′ at 8 keV
440 cm2 at 1.5 keV, 250 cm2 at 8 keV
2′ (Half power diameter)
17′ .8 × 17′ .8
0.2–12 keV
1024 × 1024 / 24 µm × 24 µm∼ 130 eV FWHM at 6 keV
330 cm2 (FI), 370 cm2 (BI) at 1.5 keV
160 cm2 (FI), 110 cm2 (BI) at 8 keV
4◦ .5 × 4◦ .5 (≳ 100 keV), 34′ × 34′ (≲ 100 keV)
10–600 keV

30

CHAPTER 4. INSTRUMENTS

4.2

X-ray Telescope (XRT)

4.2.1

Mechanism and design

The refractive index of, for instance, glass for X-rays is close to unity as opposed toa much larger value for visible light. Therefore, we cannot use the same focusing method
as visible light to focus X-rays. The most common method to focus X-rays is using
total external reflection. X-rays are reflected when the grazing incidence angle is shallow
enough (≲ 1◦ ). In 1952, Wolter proposed the design composed of two-stage reflection
mirrors in order to realize a wide filed-of view (FOV). The Wolter type I telescope is themost widely used for X-ray astronomy. It is composed of primary and secondary mirrors
in parabolic and hyperbolic shapes, respectively (Figure 4.2).
Suzaku carried five Wolter type I telescopes. One is for the XRS (named XRT-S),
and the others are for the XIS (named XRT-I). We hereafter refer to XRT-I simply as
the XRT. The XRT contains in total 175 pairs of concentric thin-foil reflectors, whose
thickness is 175 µm (Figure 4.3). The diameter and weight of the XRT are 399 mm and19.3 kg, respectively. The total weight of the four modules of the XRT is only ≲ 80 kg
which is smaller than those of mirrors aboard other X-ray astronomy satellites such as
Chandra (∼ 950 kg) and XMM-Newton (∼ 1300 kg).

Figure 4.2: Schematic view of the XRT composed of nested reflectors: (a) front view; (b) crosssection including the optical axis (Mori et al., 2005).

4.2.2

Performance

In spite of its light weight, the XRT has a large eﬀective area, 450 cm2 at 1.5 keV
and 250 cm2 at 7 keV, which is comparable to those of the mirrors aboard XMM-Newton(Figure 4.4a). However, the XRT trades oﬀ the spatial resolution. In order to evaluate
the spatial resolution, we commonly use a half power diameter (HPD), which is defined
as a diameter within which 50% of focused X-rays from a point source are enclosed. The
HPDs of the XRT are 1′ .8–2′ .3 (see Figure 4.4b), which is about one order of magnitudelarger than those of Chandra and XMM-Newton.
The eﬀective area is a function of the oﬀ-axis angle, which tends to decrease toward theoﬀ-axis. This eﬀect is called “vignetting”. The vignetting eﬀect is more evident in higher

4.3. X-RAY IMAGING SPECTROMETER (XIS)

31

Figure 4.3: Photograph of the XRT (Serlemitsos et al., 2007).
energies because the critical angle for total external reflection is in inverse proportion to
the X-ray energy. Figure 4.5 shows the vignetting curves of the XRT in the 3–6 keV and
8–10 keV bands.

(b)

(a)

Figure 4.4: (a) Total eﬀective area of the four XRT-I modules with those of XMM-Newton andChandra. Transmissions of the thermal shield and the optical blocking filter, and the quantum
eﬃciency of the CCD are all taken into account. (b) XIS0 image of a bright point source,
SS Cygni. The figures are taken from Serlemitsos et al. (2007)

4.3

X-ray Imaging Spectrometer (XIS)

4.3.1

Overview of the XIS

Figure 4.6 shows a photograph and a side view of a XIS module. The pixel structureof the CCDs is classified as a MOS type which consists three layers; Metal, Oxide, andSemiconductor. The CCD chips have layers of poly-silicon, silicon-dioxide and p-type

32

CHAPTER 4. INSTRUMENTS

Figure 4.5: Vignetting curves of the XRT. The data points are obtained from observations of
the Crab nebula. The model curves are calculated with a ray-tracing simulator with spectral
parameters of the Crab nebula. The figures are taken from Serlemitsos et al. (2007).

silicon, which function as a MOS-junction. Each CCD has four segments (A, B, C, andD) consisting of an imaging area, a frame-store region and a readout node (Figure 4.7).
The imaging area has 1024 × 1024 pixels with a pixel size of 24 µm × 24 µm. Combinedwith the XRT, the imaging area covers the 17′ .8 × 17′ .8 region in the sky. When an X-ray
photon is absorbed in the depletion layer of the CCD, electron-hole pairs are produced inproportion to the X-ray energy. Generated charge is transferred from the imaging area
to the frame-store region after an exposure, and is read out from the node.
In the FI CCDs, incident X-rays are attenuated by the 0.7 µm thick gate structureso that quantum eﬃciency (QE) on the low-energy side is lower. On the other hand, theBI CCD gains high QE in the low-energy band because of its thin dead layer, ∼ 10 nm(Figure 4.8). In the high energy band above ∼ 4 keV, the QE of XIS0 is higher than thatof XIS1 because of the thicker depletion layer (∼ 76 µm for the FI CCDs and ∼ 42 µm for
the BI CCD). In the case of the BI CCD, X-ray photons are incident from the opposite
side of the electrode so that we need to fully deplete it. It is diﬃcult to process a thickfull-depletion layer.

4.3.2

Calibration; charge transfer ineﬃciency

In orbit, CCDs suﬀer from radiation damage such as crystallographic defects in thedepletion layer generated by cosmic rays. When signal charge is transferred to the readoutnodes, a part of the charge is trapped in the crystallographic defects and is lost, which
increases charge transfer ineﬃciency (CTI). The charge loss is a stochastic process so thatthe amount of the loss fluctuates and energy resolution also deteriorates (Nakajima et al.,2008; Ozawa et al., 2009a).
In order to reduce the CTI, we use a method of a spaced-raw charge injection (SCI)(Prigozhin et al., 2000; Bautz et al., 2004; LaMarr et al., 2004). Artificial charge is injected
into given pixels from serial registers attached to the top of each column and fills the traps
so that subsequent charge is transferred with less charge loss. Figure 4.9 shows a trendof centroid energy and width of the Mn Kα lines (5.895 keV) from the 55 Fe radioactive

4.3. X-RAY IMAGING SPECTROMETER (XIS)

(a)

33

(b)

Figure 4.6: (a) Photograph and (b) side view of the XIS (Koyama et al., 2007).

Chargeinjection
register
A

B

C

D

Illuminatedby 55Fe

Readout
nodes

Figure 4.7: Schematic view of the XIS CCD (Koyama et al., 2007).

34

CHAPTER 4. INSTRUMENTS

Quantum efficiency

1

FI (XIS0)BI (XIS1)

0.5

0

0

5
Energy (keV)

10

Figure 4.8: Quantum eﬃciency of the XIS0 (black) and the XIS1 (red) as a function of incident
photon energy (Koyama et al., 2007).

sources equipped on two corners of each XIS (see Figure 4.7). The centroid energy of eachsensor was decreasing from the launch to August 2006 and the width was increasing inthis period. The XIS team for the first time performed the SCI in August 2006, and then,the centroid energy and the width were recovered. At the start of the SCI operations, theamount of the artificial charge was equal to the amount produced by 6 keV and 2 keV
X-ray photons for the FI and BI CCDs, respectively. The CTI of XIS1 was more rapidly
increasing as shown in Figure 4.9, and, therefore, the injected charge amount was changed
from 2 keV to 6 keV in 2011.
The gain and the energy resolution of each sensor are improved by the SCI. However,the pulse height shows periodic variation because the SCI depends on the distance between
the pixel of interest and the preceding charge injection row. Therefore, all the XIS data are
calibrated with a method to correct the pulse height for the periodic variation (Uchiyama
et al., 2009). Figures 4.10 and 4.11 show the trend of centroid energy and width of thethe Mn Kα lines after the calibration, respectively. The uncertainty of the energy-scale
determination is ≲ 20 eV (0.3%) for each sensor.

4.3.3

Calibration; gain

For gain calibrations of the XIS in the high energy band, Mn Kα and Kβ lines of
the calibration source are used. For calibration in the low-energy band, the XIS teamuses emission lines of O Lyα (0.654 keV), Ne Heα (0.921 keV), Ne Lyα (1.021 keV) andMg Heα (1.352 keV) from the bright SNR E0102−72.3 (Figure 4.12). Figure 4.13 shows
trends of centroid energies of the emission lines of XIS0. The gains of all the XIS sensors
are calibrated ≲ 10 eV.
We describe below our own gain calibration method with an example of W49B. In
analyses of thermal emission of SNRs, centroid energies of emission lines provide us withinformation of electron temperatures and degrees of ionization of the plasmas. The un-4.3. X-RAY IMAGING SPECTROMETER (XIS)35

Figure 4.9: Trend of centroid energy (top panel) and width (bottom panel) of the Mn Kα lines
for uncalibrated data of the XIS (Suzaku Technical Description).

(b)
5.92

XIS0 (FI)5920

Mn K Line Center (keV)

Mn K Line Center (keV)

(a)
5900
5880
5920

XIS1 (BI)

5900
5880
0

200
400
Time since 2005-7-10 (day)

600

5.9

XIS0

5.88
5.92

5.9

XIS1

5.88
5.92

5.9

XIS3

5.88

2008

2010
Time (year)

2012

2014

Figure 4.10: Trend of centroid energy of the Mn Kα lines for calibrated data of the XIS (a)before and (b) after turning on the charge injection technique. The black and blue points inpanel (b) indicate the segment A and D data of each sensor, respectively (Suzaku Technical
Description).

36

(b)

250

250

XIS0 (FI)200

MnKa Line Width (eV)

MnKa Line Width (eV)(a)

CHAPTER 4. INSTRUMENTS

150
100
250

XIS1 (BI)

200
150
100

0

200
400
Time since 2005-7-10 (day)

600

200

XIS0

150
250
200

XIS1

150
250
200
150

XIS3

2008

2010
Time (year)

2012

2014

Figure 4.11: Trend of width of the Mn Kα lines for calibrated data of the XIS (a) before and (b)after turning on the charge injection technique. The black and blue points in panel (b) indicate
the segment A and D data of each sensor, respectively (Suzaku Technical Description).

certainty of ∼ 10 eV significantly aﬀects spectral fitting results. Therefore, we need amore accurate calibration than that performed by the XIS team.
For the gain collection, we used Lyα lines of Si (2.006 keV), S (2.628 keV), Ar(3.323 keV) and Ca (4.108 keV) from W49B and a Mn Kα line (5.895 keV) from thecalibration sources. Heα lines were not used because centroid energies of lower ionizedlines, such as Li-like and Be-like lines, are close to those of Heα lines, and intensity ratioof them depend on electron temperatures and ionization degrees of plasmas. We fitted
narrow-band spectra in each energy band around the lines with a Gaussian and a powerraw for the underlying continuum. We then obtained the relation between the true energies and the Gaussian centroids with linear regression (Figure 4.14). We corrected thegain using the relation. Figure 4.15 shows XIS0 spectra of the SNR W49B before andafter the gain correction. The emission lines are moved toward the true values.

4.3.4

Data reduction; flickering pixels

Before analyzing XIS data, we must discard hot pixels and flickering pixels. The
former always output pulse heights larger than the threshold for X-ray event detection,
whereas the latter output large pulse heights intermittently. The number of these pixels
increase in the orbit due to the radiation damage. Hot pixels are automatically identified
and removed in orbit. On the other hand, it is diﬃcult to remove flickering pixels in orbit,
and thus, they are filtered out in the on-ground data reduction process.
In order to remove the flickering pixels, the XIS team provides maps of the flickering
pixels called “noisy pixel maps” for each sensor (Figure 4.16). The XIS team defined theflickering pixels as those exceeding a threshold in a given period of time, and searched the
pixels from 19 month data in the non X-ray background database. Ratios of the flickering
pixels to the total are 1.5% (15594 pixels), 6.1% (64020 pixels), and < 0.1% (4139 pixels)
in XIS0, 1, and 3, respectively. We show spectra of the SNR G166.0+4.3 before and after
the filtering of the flickering pixels (Figure 4.17).

37

10

4.3. X-RAY IMAGING SPECTROMETER (XIS)Ne HeαNe LyαMg HeαMg Lyα

0.1

O Heα

10−3

0.01

Counts s −1 keV−1

1

O Lyα

0.5

1
Energy (keV)

2

5

Figure 4.12: XIS0 Spectrum of the SNR E0102−72.3.
O Lyα (654 eV)

680

935
Energy (eV)Energy (eV)

670
660
650

0

0

3000
1000
2000
Time from the launch (day)Mg Heα (1352 eV)

1375
1365
Energy (eV)1035
Energy (eV)

915

895

3000
1000
2000
Time from the launch (day)Ne Lyα (1021 eV)

1045

1025
1015

1355
1345
1335

1005
995

925

905

640
630

Ne Heα (921 eV)

945

0

3000
1000
2000
Time from the launch (day)

1325

0

3000
1000
2000
Time from the launch (day)Figure 4.13: Trends of centroid energies of O Lyα, Ne Heα, Ne Lyα and Mg Heα of E0102−72.3
obtained with XIS0. The black and red points are data taken in readout modes with no option
and with the 1/4 window option (for a bright source), respectively (Suzaku Technical Description).

38

CHAPTER 4. INSTRUMENTSCa
Lyα

ΔE = obtaind E - true E (eV)

20
Si
Lyα

10

S
Lyα

Ar
Lyα

Mn Kα

0
-10
-20
0

2000
4000
6000
8000
True centroid enegies of lines (eV)Figure 4.14: Comparison of centroid energies of emission lines from W49B before applying thecalibration described in §4.3.3. The red line indicates a function estimated by linear regression.

Si HeαCounts s -1 keV -1

1
10

-1

10

-2

10

-3

Si LyαS Lyα Ar Heα
Ca HeαAr LyαCa LyαFe HeαCr Heα1.6-2.8 keV band
1

Si LyαFe Fe HeβLyα + Ni Heα

7 eV

9 eV

Si Heα
0.1
1.6

10

S Heα

S Heα

S Lyα

2

2.8

-4

1

2

Energy (keV)

5

10

Figure 4.15: Green and Blue crosses show XIS0 spectra of the SNR W49B before and after thegain correction, respectively. Centroid energies of the Si and S Lyα lines shift 9 and 7 eV after
the correction, respectively.

4.3. X-RAY IMAGING SPECTROMETER (XIS)

400

600

800

1000
ActX

ActY
1000
800
600
400

0

200

400

600

800

1000
ActX

0

200

XIS3

200

400

600

800

ActY
1000

XIS1

200
0

0

0

200

400

600

800

ActY
1000

XIS0

39

0

200

400

600

800

1000
ActX

Figure 4.16: Maps of the flickering pixels provided in April 2016.

4.3.5

Non X-ray background (NXB)

In the orbit, high energy particles produce background signals, called non X-ray background (NXB). NXB spectra are obtained by observations pointing toward the Earth atnight. Figure 4.18 shows the NXB spectra of XIS0 (FI CCD) and XIS1 (BI CCD). The
NXB includes the fluorescence X-ray lines of Al Kα, Si Kα, Ni Kα and Kβ, and Au Lα,Lβ and Mα from materials of the spacecraft, and Mn Kα and Kβ from the 55 Fe calibration
sources. One of the origin of the continuum emissions is Compton scattering of cosmic
X-rays or X-rays produced by interactions between cosmic rays and the satellite structure.
The other is an ionization of cosmic-ray electrons or secondary electrons produced by theinteractions.
Figure 4.19 shows the NXB spectra obtained by CCDs on Chandra, XMM-Newton
and Suzaku in other to compare their background levels. The XIS has lower backgroundlevel in the energy band of 3–10 keV because of the lower altitude of the orbit of Suzaku.

40

CHAPTER 4. INSTRUMENTS

0.1
10−3

0.01

Counts s −1 keV−1

1

XIS0

0.5

1

2
Energy (keV)

5

0.1

1

XIS1

10−3

0.01

Counts s −1 keV−1

10

0.5

1

2
Energy (keV)

5

0.1

1

XIS3

10−3

0.01

Counts s −1 keV−1

10

0.5

1

2
Energy (keV)

5

10

Figure 4.17: Spectra of each sensor in the northwest region of the SNR G166.0+4.3. Black andred crosses indicate the spectra before and after the reduction of the flickering pixels, respectively.

4.3. X-RAY IMAGING SPECTROMETER (XIS)

41

Figure 4.18: Spectra of the NXB in the XIS0 (black) and the XIS1 (gray) (Tawa et al., 2008).

Figure 4.19: XIS background counting rate as a function of energy. The background rate wasnormalized with the eﬀective area and the field of view, which is a good measure of the sensitivity
determined by the background for spatially extended sources. The background rate of Chandra
and XMM-Newton adopted from Katayama et al. (2004) are shown for comparisons. (Mitsuda
et al., 2007)

Chapter 5
G166.0+4.3
5.1

Previous Works and Our Aims

G166.0+4.3 is a Galactic SNR whose radio shell has a large bipolar structure in thewest with a smaller semicircle shell in the east (Sharpless, 1959; Landecker et al., 1982;
Pineault et al., 1987). Landecker et al. (1982) estimated the distance to G166.0+4.3 as
5.0 ± 0.5 kpc by using the Σ-D-z (surface brightness-diameter-distance above the Galactic
plane) relationship. Pineault et al. (1987) suggested that the gas density in the east is
higher than that in the west because of the unusual morphology of the radio shell.
G166.0+4.3 was observed in the X-ray band with ROSAT (Burrows & Guo, 1994),
ASCA (Guo & Burrows, 1997) and XMM-Newton (Bocchino et al., 2009). Burrows& Guo (1994) found that the X-ray morphology is distinct from the radio morphology
and reproduced the spectrum with CIE models (Figure 5.1). They estimated the age ofthe remnant to be ∼ 2.4 × 104 yr by applying the evaporating cloudlet model (White& Long, 1991) and assuming the estimated distance of 5 kpc. Bocchino et al. (2009)classified G166.0+4.3 as an mixed-morphology SNR and reported that the spectra can be
reproduced by an IP model.
The gas environment of G166.0+4.3 is clearly diﬀerent between the east and thewest, according to the morphology of the radio shell. It is one of the best sources forstudying plasma evolution depending on the environment. In this Chapter, we reporton observational results of G166.0+4.3 with Suzaku, which has high sensitivity and highenergy resolution in the energy band of 0.3–12 keV. Throughout this chapter, errors arequoted at 90% confidence levels in the text and tables, and error bars in the figures indicate
1σ confidence intervals. We assume 5 kpc as the distance to the remnant. All spectral
fits are performed with XSPEC version 12.9.0n. The plasma models are calculated with
ATOMDB version 3.0.2. We used solar abundances given by Anders & Grevesse (1989).

5.2

Suzaku Observation and Data Reduction

We performed Suzaku observations of G166.0+4.3. Table 5.1 summarizes the observation log. In this analysis, we used data from the XIS installed on the focal planes of
the XRT. XIS2 and a part of XIS0 are not used in what follows (see §4.3.1).
We reduced the data using the HEADAS software version 6.19. We used the calibration
database released in 2015 October for data processing. The total exposure time of the five42

5.2. SUZAKU OBSERVATION AND DATA REDUCTION

43

observations is ∼ 226 ks after the standard data screening. In the processing of the XIS
data, we removed cumulative flickering pixels by referring to the noisy pixel maps provided
by the XIS team (§4.3.3). We also discarded pixels adjacent to the flickering pixels. NXBs
were estimated by xisnxbgen (Tawa et al., 2008). The redistribution matrix files and theancillary response files were produced by xisrmfgen and xissimarfgen (Ishisaki et al.,2007), respectively.

Figure 5.1: ROSAT image of G166.0+4.3 in the 0.52-2.02 keV band with 1.4 GHz radio contours(Burrows & Guo, 1994). The coordinates refer to the J2000.0 epoch. The red curve along thebottom of the image represents the X-ray intensity along the slice shown as a green line, which
passes through the center of the SNR near its plane of symmetry. The inset in the upper
left-hand corner is a color image of the radio continuum data, scaled from 0 to 1.67 K.

Table 5.1: Suzaku observation log of G166.0+4.3.
Target
G166.0+4.3 northeast
G166.0+4.3 northeast
G166.0+4.3 northwest
G166.0+4.3 southeast
G166.0+4.3 southeast
IRAS 05262+4432 (Background)Obs. ID
509022010
509022020
509023010
509024010
509024020
703019010

Obs. date2014-09-19
2014-09-22
2014-09-20
2014-09-21
2015-03-13
2008-09-14(R.A., Dec.)(5h 27m 07s .7, 42◦ 58′ 52′′ .2)(5h 27m 07s .7, 42◦ 58′ 52′′ .2)(5h 25m 46s .8, 42◦ 54′ 01′′ .7)(5h 26m 41s .2, 42◦ 39′ 06′′ .2)(5h 26m 41s .2, 42◦ 39′ 06′′ .2)(5h 29m 56s .0, 44◦ 34′ 39′′ .2)

Eﬀective Exposure

61 ks
61 ks
42 ks
34 ks
27 ks
82 ks

44

CHAPTER 5. G166.0+4.3

5.3

Analysis

5.3.1

Image

Figure 5.2 shows a mosaic image of G166.0+4.3 in the energy band of 0.5–2.0 keV.
We subtracted the NXB from the image and corrected it for the vignetting eﬀect of the
XRT. The contours are a 325 MHz radio image from the Westerbord Northern Sky Survey(WENSS).

5:28:00.0

5:26:00.0

5:24:00.0

NE

43:00:00.0

Declination

43:30:00.0

Right ascension

42:30:00.0

W

9.8e-43.9e-3

1.6e-26.3e-22.5e-1

1.0

Figure 5.2: XIS image of G166.0+4.3 in the 0.5-2.0 keV band after the NXB subtraction andcorrection of the vignetting eﬀect. The coordinates refer to the J2000.0 epoch. The X-ray countis normalized so that the peak becomes unity. We overlaid contours of the 325 MHz radio image
from the WENSS.

5.3.2

X-ray background subtraction

For X-ray background estimation, we used Suzaku observational data of IRAS 05262+4432,
which is located at (l, b) = (165◦ .1, 5◦ .7). The observation log is shown in Table 5.1. We
extracted a spectrum from a source-free region. In order to model the spectrum, we
adopted a model by Masui et al. (2009). They observed the direction (l, b) = (230◦ , 0◦ )with Suzaku in order to study the soft X-ray emission from the Galactic disk. Their model
consists of four components: the cosmic X-ray background (CXB), the local hot bubble(LHB), and two thermal components for the Galactic halo (GHcold and GHhot ).
We fit the X-ray background spectrum with the model after the NXB subtraction.
We fixed the electron temperatures of the LHB, GHcold and GHhot at the values given5.3. ANALYSIS

45

by Masui et al. (2009). The absorption column density (NH ) of the CXB was fixed at
3.8×1021 cm−2 , the Galactic value in the line of sight toward IRAS 05262+4432 (Dickey& Lockman, 1990). The normalizations of all the components were allowed to vary.
Figure 5.3 and Table 5.2 show the fitting result and the best-fit parameters, respectively.

Counts s −1 keV−1

0.1

0.01

CXB

10−3

GH cold

LHB

GH hot

10−4
4

χ

2
0
−2
−4
0.5

1

2
Energy (keV)

5

10

Figure 5.3: XIS0+3 (black crosses) and XIS1 (red crosses) spectra taken from the source-free
region of the IRAS 05262+4432 data with the best-fit model. The solid lines show the modelof XIS0+3. The blue, green, magenta and orange lines represent the CXB, LHB, GHcold andGHhot models, respectively. The bottom panel shows the data residuals from the best-fit model.

Table 5.2: Best-fit model parameters of the X-ray background spectra for
G166.0+4.3.
Component
CXB

Model function
Wabs
Power law

LHB

APEC

GHcold

APEC

GHhot

APEC

Parameter
NH (1021 cm−2 )Photon index
Normalization∗kTe (keV)
Normalization†kTe (keV)
Normalization†kTe (keV)
Normalization†
χ2ν (ν)Value
3.8 (fixed)
1.4 (fixed)
10.7 ± 0.4
0.105 (fixed)
13.4 ± 3.2
0.658 (fixed)
2.1 ± 0.2
1.50 (fixed)
3.1 ± 0.5
1.29 (234)∗ The unit is photons s−1 cm−2 keV−1 sr−1 at 1 keV.
†The emission measure integrated over the line of sight, i.e., (1 / 4π)
units of 1014 cm−5 sr−1 .

∫ne nH dl in

46

5.3.3

CHAPTER 5. G166.0+4.3

SNR spectra

Considering the morphology of the radio shell, we extracted spectra from the northeast
and west regions (named NE and W) indicated with the white solid lines in Figure 5.2.
Figure 5.4 shows XIS0+3 spectra of W and NE after the NXB subtraction. One canclearly see Fe-L and Ni-L line complexes, and Si-K and S-K lines. The Fe-L and Ni-L line
complexes are more prominent in W than those in NE.

0.1
0.01

S

Ne

Mg
Fe L
Ni L

10−4

10−3

Counts s −1 keV−1

1

Si

0.5

1

2
Energy (keV)

5

10

Figure 5.4: XIS0+3 spectra of the W (black) and NE (blue) regions. The vertical solid anddashed black lines show the center energies of the K lines from He-like and H-like ions, respectively. The arrows indicate the energy bands with Fe-L and Ni-L line complexes.

We fitted the spectra with the model consisting of the SNR component and the X-ray
background (§5.3.2). The overall normalization of the X-ray background component was
treated as a free parameter. The absorption column density of the CXB component was
fixed at 3.6×1021 cm−2 which is the Galactic value in the line of sight toward G166.0+4.3.
We applied a one-component IP model to the W spectrum using the VVRNEI modelin the XSPEC, which calculates the spectrum of a non-equilibrium ionization plasma after
a rapid transition of the electron temperature from kTinit to kTe . The initial temperaturekTinit is fixed at 0.01 keV whereas the present electron temperature kTe , ionization parameter ne t and normalization are free parameters. We used the Wisconsin absorption(Wabs) model (Morrison & McCammon, 1983) for the absorption, whose column density
is a free parameter. The abundances of O, Ne, Mg, Si, S, and Fe are also free parameters.
The Ar and Ca abundances are linked to S, and Ni is linked to Fe. The abundances ofthe other elements are fixed to the solar values. The energy band around the neutral Si
K-edge (1.73–1.78 keV) is ignored because of the known calibration uncertainty. The fit
left large residuals around 0.82 keV and 1.23 keV due to the lack of Fe-L lines in themodel (e.g., Yamaguchi et al., 2011). Therefore, we added two narrow lines at 0.82 keV
and 1.23 keV. The results are shown in Figure 5.6(w-i) and Table 5.3. We found that the
spectrum is well reproduced by a one-component IP model with the electron temperatureof 0.83 keV.

5.3. ANALYSIS

47

We applied the same IP model as W to the NE spectrum. Significant residuals arefound at ∼ 2.0 keV and ∼ 2.6 keV which correspond to Si Lyα (2.0 keV) and theedge of the RRC of He-like Si (2.67 keV) + S Lyα (2.63 keV), respectively (Figure 5.6(ne-i)). These spectral features, if confirmed, provide evidence that the plasma is in arecombination-dominant state. We tried a one-component RP model. However, a satisfactory fit was not obtained.
For a detailed investigation, we restricted the energy band to > 1.6 keV, where wefound the large residuals, and tried one-component IP or RP models. The absorption
column density is fixed to the best-fit value for W, NH = 0.8 × 1021 cm−2 , whereas kTe
and ne t are left free. The kTinit is a free parameter for the RP model, whereas it is fixed
at 0.01 keV for the IP model. The abundances of Si and S are free parameters, andthose of Ar and Ca are linked to S. The residuals were not improved with the IP model(Figure 5.6(ne-ii) and Table 5.3). On the other hand, we successfully fitted the spectrumwith the RP model (Figure 5.6(ne-iii) and Table 5.3).
We extrapolated the above RP model down to 0.5 keV. Figure 5.6(ne-iv) shows themodel and the full-band spectrum in the energy band of 0.5–10.0 keV. The spectral stru-

Photons cm-2 s −1 keV−1

Counts s −1 keV−1(a)Photons cm-2 s −1 keV−1

Counts s −1 keV−1

(b)

2.0
1.5

kTe = 0.40 keV

1.0
0.5
0
0.4
0.3
0.2
0.1
0

0.5

2.0
1.5

0.6

0.7

0.8

0.9

1.0

1.1

1.2

0.8

0.9

1.0

1.1

1.2

Energy (keV)

kTe = 0.85 keV

1.0
0.5
0
0.1
0.05
0

0.5

0.6

0.7

Energy (keV)

Figure 5.5: (a) NE spectrum of XIS0+3 (crosses) and the plasma model with (top) and without(bottom) response matrixes multiplied (solid lines) in the energy band of 0.5–1.2 keV. The plasma
model consists only of Fe and Ni components with NH = 1.7 × 1021 cm−2 , kTe = 0.40 keV,
kTinit = 3.0 keV and ne t = 6.0 × 1011 cm−3 s. (b) Same as (a) but with kTe = 0.85 keV.

48

CHAPTER 5. G166.0+4.3

cture in the 0.7–1.1 keV band cannot be reproduced. With the electron temperaturekTe = 0.37 keV, the model predicts a peak energy of the Fe and Ni L line complexes at∼ 0.8 keV, which is significantly lower than the observed value of ∼ 0.95 keV. Figure 5.5
shows the NE spectrum of XIS0+3 and the plasma model in the energy band of 0.5–1.2
keV. The model consists only of the Fe and Ni components with NH = 1.7 × 1021 cm−2 ,kTinit = 3.0 keV and ne t = 6.0 × 1011 cm−3 s. The electron temperatures are 0.40 keV and
0.85 keV in Figures 5.5(a) and 5.5(b), respectively. The dominant emission lines of the
Fe and Ni components are FeXVII 2p5 3s1→2p6 (0.725 keV and 0.739 keV) and FeXVII
2p5 3d1→2p6 (0.812 keV and 0.826 keV) with kTe = 0.40 keV, and FeXIX 2p3 3s1→2p4(0.822 keV), FeXVIII 2p4 3d1→2p5 (0.873 keV), FeXX 2p2 3d1→2p3 (0.965 keV) andFeXXI 2p1 3d1→2p2 (1.009 keV) with kTe = 0.85 keV. The comparison between the data
and the models indicates that the electron temperature of the Fe and Ni plasma shouldbe ∼ 0.85 keV in order to reproduce the observed structure. On the other hand, the Si
and S plasma demands lower kTe as we saw in the (ne-iii) fit.
The result indicates that the plasma has diﬀerent kTe between Fe (+Ni) and the otherelements. Kamitsukasa et al. (2016) decomposed the ejecta of SNR G272.2−3.2 into
components with diﬀerent elements, considering that each element has diﬀerent radial
distributions. Following their idea, we adopted a two-component VVRNEI model: one
for Fe and Ni, and the other for lighter elements. The spectra can be well fitted with themodel comprising a CIE for Fe and Ni and an RP for the other elements. Figure 5.6(ne-v)
shows the spectrum with the best-fit model. Although we tried IP and RP for the Fe andNi component, the ionization parameter ne t cannot be constrained well. Therefore, we
here assumed the CIE model for simplicity.

5.4

Discussion on G166.0+4.3

5.4.1

Spatial structure

The spectral analysis shows that the NE plasma contains two components; the lowtemperature RP component and the high temperature Fe-rich component. This suggeststhat the distributions of the two components are diﬀerent from each other. In order toinvestigate the distribution, we made XIS images in the Si-K band (1.75–2.10 keV) and the
Fe-L band (0.70–1.30 keV). We divided them by the corresponding underlying continuum
image for which we selected the line-free band of 1.50–1.70 keV. We show the images in
Figure 5.7, which indicates that Si distributes uniformly whereas the Fe distribution ismore compact. In some SNRs, regardless of the type of the explosion, distributions of Fe
are observed to be centered (e.g., Uchida et al., 2009; Hayato et al., 2010). The observed
distributions are consistent with the widely believed picture of supernova nucleosynthesisthat heavier elements tend to be produced closer to the center of the progenitor.
For a detailed analysis of the Fe distribution, we extracted spectra from the three
regions (i)–(iii) in NE indicated by the white lines in the bottom panel of Figure 5.7.
We fitted the spectra of each region with the same model as Figure 5.6(ne-v). We fixedkTe of the RP component to the best-fit value, 0.46 keV. Figure 5.8 shows the resultant
abundance ratios of Fe to Si. The ratio in the inner region is significantly higher andsuggests the center-filled distribution of Fe.

5.4. DISCUSSION ON G166.0+4.3

49

(w-i)

(ne-i)IP + 2gauss model(0.5-10.0 keV)Counts s −1 keV−1

0.1
0.01
10−3

2

0

0
−2
−4

1

(ne-ii)

2
Energy (keV)

5

10

0.5

1

(ne-iii)
IP model (1.6-10 keV)

1

2
Energy (keV)Counts s −1 keV−1

0.1
0.01
10−3

10

0.1
0.01
10−3

10−4
4

10−4
4

2

2

0

5

RP model (1.6-10 keV)1

χCounts s −1 keV−1

10−3

2

0.5

χ

0.01

10−4
4

−4

0
−2

−2

−4

−4
0.5

1

(ne-iv)

1.6 2
Energy (keV)

5

0.5

10

1

(ne-v)
RP + 2gauss model(0.5-10.0 keV)

1
0.1
0.01
10−3

1.6 2
Energy (keV)

5

10

CIE (only Fe & Ni) + RP (other elements) + 2gauss model

1
Counts s −1 keV−1

Counts s −1 keV−1

0.1

10−4
4

−2

(0.5-10 keV)

0.1
0.01
10−3
10−4
4

10−4
4

2

2
0

χ

χ

IP + 2gauss model(0.5-10.0 keV)

1

χ

χ

Counts s −1 keV−1

1

0
−2

−2

−4

−4
0.5

1

2
Energy (keV)

5

10

0.5

1

2
Energy (keV)

5

10

Figure 5.6: XIS0+3 (black crosses) and XIS1 (red crosses) spectra in the W and NE regions.
Each spectrum is fitted with the model consisting of the SNR and background components. The
solid and dotted lines show the SNR and background models of XIS0+3, respectively.

ParameterNH (1021 cm−2 )kTe (keV)
kTinit (keV)
ZH (solar)
ZHe (solar)
ZC (solar)
ZO (solar)
ZNe (solar)
ZMg (solar)
ZSi (solar)

Model function

Wabs
VVRNEI1(w-i)(ne-i)(ne-ii)(ne-iii)(ne-v)
0.8 ± 0.1
0.3 ± 0.1
0.8 (fixed)
0.8 (fixed)
1.7 ± 0.1
0.83 ± 0.01
0.90 +0.01
1.36 +0.51
0.37 ± 0.04
0.46 ± 0.03
−0.02
−0.27
0.01 (fixed) 0.01 (fixed) 0.01 (fixed)≥ 2.983.0 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)+0.2+0.2
0.4 −0.1
0.2 −0.1
1 (fixed)
1 (fixed)
0.02 +0.03
−0.01
+0.2
+0.2
0.6 −0.1
0.6 −0.1
1 (fixed)
1 (fixed)
0.4 ± 0.1
1.2 ± 0.2
0.6 ± 0.1
1 (fixed)
1 (fixed)
0.6 ± 0.1
+0.3
0.7 +0.1
0.6
±
0.1
1.4
2.3
±
0.3
1.0
± 0.1
−0.2
−0.1
+0.7
ZS = ZAr = ZCa (solar)
1.2 ± 0.3
1.2 ± 0.2
1.8 ± 0.4
3.3 −0.6
1.5 ± 0.2
ZFe = ZNi (solar)
1.2 ± 0.2
0.4 ± 0.1
1 (fixed)
1 (fixed)
0 (fixed)+1.7
ne t (1011 cm−3 s)
4.0 +0.7
3.1
±
0.5
1.6
6.2
±
0.6
6.1 +0.5
−0.5
−0.7
−0.4
+0.3
+1.6
+0.8
V EM (1057 cm−3 )†2.7 ± 0.2
2.9 +0.2
1.2
6.6
10.3
−0.3
−0.2
−1.3
−1.0
+0.02
VVRNEI2
kTe (keV)
————
————
————
————
0.87 −0.03(only Fe & Ni comp.)
kTinit (keV)
————
————
————
————
= kTe (VVRNEI2)
ZFe = ZNi (solar)
————
————
————
————
0.14 ± 0.01
Zothers (solar)
————
————
————
————
0 (fixed)
ne t (1011 cm−3 s)
————
————
————
————
100 (fixed)
V EM (1057 cm−3 )†
————
————
————
————
= V EM (VVRNEI1)
GaussianCentroid (keV)
0.82 (fixed) 0.82 (fixed)
————
————
0.82 (fixed)+0.8
Normalization‡
4.6 +1.3
2.1
————
————
6.7 +1.5
−0.7
−0.7
−1.3
GaussianCentroid (keV)
1.23 (fixed) 1.23 (fixed)
————
————
1.23 (fixed)
Normalization‡
3.4 +1.3
0.8 ± 0.1————————
0.7 ± 0.2
−0.2
2
1.37 (247)
1.77 (251)1.55 (150)
1.15 (150)
1.39 (250)χν (ν)∫†
Volume emission measure V EM = ne nH dV , where ne , nH , and V are the electron and hydrogen densities,
and the emitting volume, respectively.
‡
The unit is photons s−1 cm−2 sr−1 .

Table 5.3: Fit parameters of source regions.

50
CHAPTER 5. G166.0+4.3

5.4. DISCUSSION ON G166.0+4.3

51
Right ascension

Right ascension
5:26:00.0

5:28:00.0

5:24:00.0

Declination(i) (ii)
43:00:00.0

43:30:00.0

NE
43:00:00.0

Declination

5:26:00.0

5:24:00.0

43:30:00.0

5:28:00.0(iii)

42:30:00.0

42:30:00.0

W

0.03

0.06

0.13

0.25

0.50

1.0

0.03

0.06

0.13

0.25

0.50

1.0

Figure 5.7: XIS images of G166.0+4.3 in 1.75–2.10 keV band (left: Si-K band) and the 0.70–
1.30 keV band (right: Fe-L and Ni-L band) after subtraction of the NXB and correction of thevignetting eﬀect. The images are divided by the corresponding underlying continuum image forwhich we selected the line-free band of 1.50–1.70 keV. The scale is normalized so that the peak
becomes unity. The contours are a radio image of the WENSS at 325 MHz.

5.4.2

Ambient gas density and SNR age

In our spectral analysis, the volume emission measures are 2.7 × 1057 cm−3 and 1.0 ×
1058 cm−3 in W and NE, respectively. The volumes are obtained to be 4.1 × 1059 f cm3
and 1.8 × 1059 f cm3 for the W and NE plasmas, respectively, where f is the filling factor,
assuming the morphologies to be prisms whose bases are the W and NE regions and depths
are 40 pc and 20 pc for W and NE, respectively. With the volume emission measures,
we obtain the electron densities ne,W = 0.1 f −1/2 cm−3 and ne,NE = 0.3 f −1/2 cm−3 for
W and NE, respectively. They suggest that the ambient gas is denser in NE than that
in W. This is supported by the radio morphology, the smaller shell in the eastern side.
Another clue of denser gas in NE is the GeV gamma-ray emission detected in the easternpart of the remnant by Fermi LAT observations (Figure 5.9: Araya, 2013). If π 0 decays
are the main production process for the gamma rays, dense gas, which acts as targets foraccelerated protons for π 0 production, would be in NE.
In the spectral analysis, we obtained ne,W t = 4.0 × 1011 cm−3 s and ne,NE t = 6.1 ×
11
10 cm−3 s for the W and NE plasmas, respectively. With ne,W and ne,NE calculated above
based on the emission measures, the timescales are estimated to be t = 1.4 × 105 f 1/2 yr
for W and t = 7.1 × 104 f 1/2 yr for NE. The ionization timescale for the W plasma, which

52

CHAPTER 5. G166.0+4.3(ZFe / ZFe ) / (ZSi / ZSi )0.25

Region

(i)

(ii)

(iii)

0.2

0.15

0.1

0.05

0

6
12
18
24
Angular distance from the NE rim (arcmin)

Figure 5.8: Abundance ratios of Fe to Si in the regions (i)–(iii) as a function of angular distancesfrom the NE rim toward the center of the remnant.

NE

W

Figure 5.9: Signal-to-noise map in the G166.0+4.3 smoothed with a boxcar of length 0◦ .5 inthe energy band of 0.5–2 GeV with Fermi (Araya, 2013). The green contours show the WENSS
325 MHz radio image. The while lines indicate the NE and W regions. The Galactic coordinates
are referred.

5.4. DISCUSSION ON G166.0+4.3

53

is an IP, is expected to be almost the same as the age of the remnant, and, therefore, weestimate the age to be tage = 1 × 105 f 1/2 yr. With ne,NE t, we can estimate the elapsed
time since recombination started to dominate over ionization to be trec = 6 × 104 f 1/2 yr(< tage ).

5.4.3

Origin of recombining plasma

As shown in §3.3, the rarefaction (Itoh & Masai, 1989) and thermal conduction(Kawasaki et al., 2002) scenarios have been considered for the formation process of RPs.
In either case, the environment of the remnant is a key to the RP production process. We
discovered an RP only in a part of G166.0+4.3, the outer part of the NE region, where theambient gas is denser. In the W region, the spectrum is reproduced by the IP model. The
diﬀerent gas environments between the W and NE regions might actually be responsible
for the diﬀerence of the plasma states.
In the rarefaction scenario, the electron temperature is decreased by adiabatic coolingwhen the blast wave breaks out of dense CSM into rarefied ISM. The lower the ISM density
is, the more eﬀective the adiabatic cooling must be (Shimizu et al., 2012). Therefore, inthis scenario, we can naturally expect an RP in the W region whose ambient gas density
would be lower. Our result, where we found an RP only in the NE region, is not simply
explained by the rarefaction scenario.
In the thermal conduction scenario, an RP is anticipated in a part of the remnant
where blast waves are in contact with cool dense gas. In the case of G166.0+4.3, the RP
was indeed discovered in the denser NE region. In our results, the temperature of the inner
Fe-rich plasma is higher (kTe = 0.87 keV) than that of the outer RP (kTe = 0.46 keV). In
this scenario, the plasma is cooled from the outside layers. Therefore, it is possible thatthe plasma is cooled only in the outer part of the NE region.

Chapter 6
IC 443
6.1

Previous Works and Our Aims

IC 443 (G189.1+3.0) is a Galactic supernova remnant (SNR) at a distance of 1.5 kpc(Welsh & Sallmen, 2003). Its age is estimated to be in the range of 3–30 kyr (Petre etal., 1988; Olbert et al., 2001). Olbert et al. (2001) found a PWN in the southeastern partof the remnant, suggesting that IC 443 is a remnant of a core-collapse supernova. Seenin the radio and optical bands, IC 443 has two shells with diﬀerent radii (e.g., Lee et al.,2008). The remnant is categorized as a mixed-morphology SNR since the X-ray emission
has a center-filled morphology.
In X-rays, Kawasaki et al. (2002) measured the Lyα to Heα line intensity ratio of Si
and S with ASCA, and found that the plasma in the northeastern part of the remnant
is more ionized than expected in CIE. With Suzaku, Yamaguchi et al. (2009) discovered
RRCs of Si and S in the northeastern region, which provided unambiguous evidence for
an RP (Figure 6.1). Ohnishi et al. (2014) found RRCs of Ca and Fe in almost the same
region as Yamaguchi et al. (2009) and indicated that heavier elements are also overionized.
IC 443 is interacting with both molecular and atomic gas (e.g., Lee et al., 2008). Cornett et al. (1977) discovered a dense molecular cloud associated with IC 443. CO line
emissions in the southeastern part of the remnant suggest that the molecular cloud is interacting with IC 443 (e.g., Denoyer, 1979; Xu et al., 2011; Yoshiike, 2017). Measurementsof the CO line velocity by Yoshiike (2017) suggest that the most part of the interacting
molecular gas is located in front of IC 443. In the northeast, the SNR shock emits atomic
lines expected from post-shock recombining gas (Fesen & Kirshner, 1980), and therefore,
the shock is likely to be propagating into an atomic medium. Near-infrared observations
also support the atomic gas in the northeast (e.g., Kokusho et al., 2013).
In this chapter, we perform spatially resolved X-ray spectroscopy of the SNR IC 443
with Suzaku and compare the results with the distribution of interstellar gas interactingwith the remnant. We then discuss the implications particularly on the formation processof the RP. All spectral fits are performed with XSPEC version 12.9.0n. The plasma
models are calculated with ATOMDB version 3.0.8. We used solar abundances given by
Wilms et al. (2000).
54

55

Declination

37:12.0

22:45:00.0

52:48.0

6.2. SUZAKU OBSERVATION AND DATA REDUCTION

36.0

6:17:12.0
Right ascension

16:48.0

Figure 6.1: IC 443 image and spectrum taken from Yamaguchi et al. (2009). (left) Vignettingcorrected XIS image of the northern part of IC 443 in the 1.75–3.0 keV band. The optical
Digitized Sky Survey image is overlapped in contours. (right) XIS spectrum in the 1.75-6.0 keV.
Black and red represent FI and BI, respectively. Individual component of the best-fit model for
the FI data are shown with solid lines; blue, green, and magenta are the VAPEC, Gaussians (SiLyα, Si Lyβ, S Lyα and Ar Lyα), and RRC components of H-like Mg, Si and S, respectively.

6.2

Suzaku Observation and Data Reduction

Table 6.1 summarizes the Suzaku observation log of IC 443. We used the data fromthe XIS, which consists of four X-ray CCD cameras installed on the focal planes of the
XRT. We do not use XIS2 and the part of XIS0 for the data acquired after November
2006 and June 2009, respectively, because they have not been functioning since thesetimes (§4.3.1).
We reduced the data using the HEADAS software version 6.19. We used the calibration
database released in April 2016 for processing the data. We removed flickering pixels byreferring to the noisy pixel maps provided by the XIS team (§4.3.4). We also discarded
pixels adjacent to the flickering pixels. NXBs were estimated by xisnxbgen (Tawa et al.,2008). Redistribution matrix files and ancillary response files were produced by xisrmfgenand xissimarfgen (Ishisaki et al., 2007), respectively.
Analyzing the XIS0 and 3 data taken in 2007, we noticed that the centroid energies of
Lyα lines from IC 443, and Mn Kα and Kβ lines from the on-board calibration sources
are shifted by 2–8 eV from the true values. We corrected the gain by using the Lyα linesof O, Mg, Si, S from IC 443, and Mn Kα and Kβ lines from the calibration sources asmentioned in §4.3.3.
Table 6.1: Suzaku observation log of IC 443.
TargetIC 443 northeast
IC 443 southeast
IC 443 northwest
IRAS 05262+4432 (Background)Obs. ID
501006010
501006020
505001010
703019010

Obs. date
2007-03-06
2007-03-07
2010-09-17
2008-09-14(R.A., Dec.)(6h 17m 11s .4, 22◦ 46′ 32′′ .5)(6h 17m 11s .3, 22◦ 28′ 46′′ .9)(6h 15m 59s .4, 22◦ 45′ 18′′ .7)(5h 29m 56s .0, 44◦ 34′ 39′′ .2)

Exposure
42 ks
44 ks
83 ks
82 ks

56

CHAPTER 6. IC 443

6.3

Analysis

6.3.1

Image

Figure 6.2a shows an XIS3 image of IC 443 in the energy band of 0.3–2.0 keV. We
subtracted the NXB from the image and corrected for the vignetting eﬀect of the XRT.
The X-ray emissions in the east are center-filled as is previously known (e.g., Troja et al.,2006). In the northwest, the image shows a hint of a shell-like structure.
Figure 6.2b shows an XIS3 image in the 3.0–5.0 keV band. We found the PWN 1SAX
J0617.1+2221 (region A) and the point
source U061530.75+224910.6 (region B) (Bocchino & Bykov, 2001, 2003). We found two
additional point sources in the northwest and a source in the northeast which are alreadydetected in XMM-Newton data by Bocchino & Bykov (2003). We excluded the PWN andthe point sources for our spectral analysis of IC 443.

6.3.2

SNR spectraPrevious X-ray studies found that the spectra in the northeastern region can be reproduced with RP models (Yamaguchi et al., 2009; Ohnishi et al., 2014). In order to examine
ionization states of the SNR plasmas in the other regions, we extracted spectra from thesoutheast (SE), northwest (NW) and northeast (NE) regions shown in Figure 6.2a. We
used the spectrum of IRAS 05262+4432 (§5.3.2) for an X-ray background spectrum inthis spectral analysis.
Figure 6.3 shows the XIS0+3 spectrum of the SE region after NXB subtraction. The
characteristic structures of Si and S RRCs indicate that the plasma is in a recombinationdominant state. Following Yamaguchi et al. (2009), we first analyzed the spectrum inthe energy band above 1.6 keV in order to quantify the ionization states of Si and S.
We applied a model consisting of one RP component to the spectrum. We used the
VVRNEI model in XSPEC. The initial plasma temperature kTinit was fixed at 5 keV, inwhich most Si and S ions become bare nuclei. The present electron temperature kTe andionization parameter ne t were allowed to vary. The abundances of Si, S, Ar, Ca and Fe
were allowed to vary, whereas the Ni abundance was linked to Fe. The abundances ofthe other elements were fixed to solar. We used the Tuebingen-Boulder ISM absorption
model (TBabs; Wilms et al., 2000), whose column density was fixed at 7.0 × 1021 cm−2 asdetermined by Kawasaki et al. (2002). For the XIS0+3 spectrum, we ignored the energy
band of 1.78–1.92 keV because of the known calibration uncertainty around the neutral
Si K-edge. Although we excluded the PWN from the spectrum extraction region, there is
non-negligible leakage from the PWN due to the tail of the point spread function of the
XRT. We, therefore, added a power law to account for the PWN emission. We fixed thephoton index at 1.89 given by our spectral analysis of the PWN region (see Appendix A.1).
Figure 6.4a and Table 6.2 show the fitting results and the best-fit parameters, respectively. The spectrum is reproduced well by the model with kTe = 0.30 ± 0.01 keV andne t = (4.6 ± 0.1) × 1011 s cm−3 . The electron temperature is significantly smaller than
0.61+0.03−0.02 keV given by Yamaguchi et al. (2009). They fitted the spectrum with a phe-6.3. ANALYSIS

57

(a)

NE

NW

SE

19:00.0

0.010

0.013

0.025

(b)

6:17:00.0

0.26

1.0

15:00.0

50:00.0
22:40:00.0

9

30:00.0

15
17

0.013

6

7

8
5

4

15:00.0

16:00.0

12

10

11

20:00.0

Declination
0.010

16:00.0

Right ascension
18:00.0

19:00.0

0.072

17:00.0

18:00.0

3

2

B

18

19

20

1

13

A

0.025

14
16

0.072

0.26

1.0

Figure 6.2: XIS3 images of IC 443 in the energy band of (a) 0.3–2.0 keV and (b) 3.0–5.0 keV after
NXB subtraction and correction of the vignetting eﬀect. In panel (a), we overlaid green contoursof 1.4 GHz radio image which were obtained from the NRAO VLA Sky Survey (NVSS). The
cyan and magenta contours are radio images of 12 CO(J = 1 − 0) and 12 CO(J = 2 − 1) obtained
from the NANTEN2, respectively (Yoshiike, 2017). Sources in regions A and B in panel (b) are
the PWN 1SAX J0617.1+2221 and the point source U061530.75+224910.6, respectively. The
X-ray count rates are normalized so that the peaks become unity. The coordinate refers to theJ2000.0 epoch.

58

CHAPTER 6. IC 443

10

Ne

Mg

S

1

O

Ar

Fe L complex

0.1

Ca
Fe K

Si RRC
S RRC

10−3

0.01

Counts s −1 keV−1

Si

1

2
Energy (keV)

5

Figure 6.3: XIS0+3 spectrum of the SE region after NXB subtraction. The vertical solid anddashed black lines indicate the centroid-energies of Heα and Lyα lines, respectively. The arrows
indicate the energy bands of the Fe-L complex and the radiative recombining continua of Si andS.

nomenological model consisting of a CIE component, Gaussians for Si, S and Ar Lyα lines
and RRCs of Mg, Si and S. We here used a more physically oriented model that takes
into account ion fractions and ionization timescales of all elements in the recombinationdominant state. In this model, RRCs of Ne and of lighter elements, which are not includedin the model employed by Yamaguchi et al. (2009), extend up to the Si–S band. Thisdiﬀerence would be the reason why we obtained lower kTe .
We extrapolated the above model down to 0.6 keV, which results in large residuals
in the soft band. We added a CIE component to the model following Kawasaki et al.
(2002) and Bocchino et al. (2009), who modeled IC 443 spectra with two-component
models with a CIE component for the soft band. The result is shown in Figure 6.4b andTable 6.2. Since we found the abundances of the CIE component are close to solar, we
fixed them at solar values. The fit left residuals around 1.23 keV, most probably due tothe lack of Fe-L lines in the model (e.g., Yamaguchi et al., 2011), and therefore, we added
a Gaussian at 1.23 keV. The fit still left residuals at 0.87 keV, indicating a shortage ofthe emission line of FeXXVIII 1s2 2s2 2p4 3d1 → 1s2 2s2 2p5. To make the line stronger
with ne t ∼ 1011 s cm−3 , kTe is required to be higher than ∼ 0.5 keV.
We then fitted the spectrum with a model consisting of a power-law, a CIE, an RP
and an additional higher kTe components. We tried CIE, IP, and RP models as thehigh kTe component, and found that the CIE and IP models fail to fit the spectrum.
We applied a model consisting of a CIE and two RP components to the spectrum. The
spectrum is well reproduced by the model with the significantly diﬀerent kTe between
the two RP components whereas ne t and abundances of the two components are almost
the same. Therefore, we linked ne t and the abundances of the components to each other.

6.3. ANALYSIS

59

The best-fit model and parameters are shown in Figure 6.4c and Table 6.2, respectively.
The residuals at the FeXXVIII line was significantly improved. The electron temperaturesof the CIE and two RP (named RPcold and RPhot ) components are 0.22 ± 0.01 keV,
0.19 ± 0.01 keV and 0.54+0.03−0.01 keV, respectively. The ionization timescale of the RP+0.2
components is ne t = (4.2−0.1 ) × 1011 s cm−3 .
We performed spectral fitting of the NW and NE spectra in the same procedure as
we did for the SE spectrum. We found that the model consisting of a CIE and two-RP
components gives satisfactory fits to both spectra. The spectra with the best-fit models
are plotted in Figures 6.4d and 6.4e and the best-fit parameters are again summarized inTable 6.2. The ionization parameter ne t of the NW spectrum is close to 2 × 1012 s cm−3 ,which is a characteristic timescale for a plasma to reach CIE (Masai, 1994). Therefore, theRP in the NW is close to CIE. Since we cannot fit the spectrum in the energy band of the
Fe-L complex of the NW spectrum, we allowed the Fe abundance of the CIE component
to vary. The fit gave somewhat higher Fe abundance of 1.3 ± 0.1, but the abundance isstill close to solar. Although kTe and ne t of the NW and NE spectra are close to those of
the SE spectrum, normalization ratios of the RP components are significantly diﬀerent,
indicating that the RP components have spatial variation.
In order to conduct spatially resolved spectroscopy in greater detail, we divided theremnant into 20 regions as shown in Figure 6.2b and extracted spectra from each region.
We fitted the spectra with the above model, and found that it reproduces all the spectra
well. The details of the fits are described in Appendix A.2. Since the leakage from the
PWN and the point sources contaminate the spectra of regions 16, 17 and 20, we added
a power law to the model for these spectra. The photon indices were fixed at 1.89 and
2.22, which are determined by our analysis in Appendix A.1, for the PWN and the pointsource U061530.75+224910.6, respectively. The obtained ranges of kTe are 0.19–0.28 keV,
0.16–0.28 keV and 0.48–0.67 keV for the CIE, RPcold and RPhot components, respectively.
The electron temperature of the CIE component is uniform whereas those of the RPcold
and RPhot components decrease toward the southeast as shown in Figures 6.5a–6.5c.
Figures 6.5d and 6.5e show the map of NH and ne t, respectively. In the northwest, ne t islarger than those of the other regions.

60

CHAPTER 6. IC 443(a)Reg. SECounts s −1 keV−1

10
1
0.1

Counts s −1 keV−1(d) 10

0.01
10

10−4
4

2

2
χ

χ

−2

0
−2

−4

(b)

2
1.6
Energy (keV)

10

−4

5

1

(e)
Reg. SE1
0.1

Counts s −1 keV−1

1Counts s −1 keV−1

0.01

4
0

0.01

2
Energy (keV)

5Reg. NE

10
1
0.1
0.01

10−3

10−3
4

4
2

2
χ

0

χ

0.1

10−3

−3

−2

0
−2

−4
1

(c)

2
Energy (keV)

10Counts s −1 keV−1

Reg. NW

1

5

−4
1

2
Energy (keV)

5Reg. SE

1
0.1
0.01
10−3
4

χ

2
0
−2
−4
1

2
Energy (keV)

5

Figure 6.4: XIS0+3 (black crosses) and XIS1 (red crosses) spectra in the SE (a–c), NW (d), andNE (e) regions. Each spectrum is fitted with the model consisting of the SNR and backgroundcomponents. The solid and dotted lines show the SNR and background models of XIS0+3,
respectively. The SNR model consists of power-law (green lines), CIE (magenta lines), RPcold(cyan lines) and RPhot (blue lines) components, respectively.

ParameterNH (1022 cm−2 )kTe (keV)ZFe = ZNi (solar)
Zother (solar)
V EM (1057 cm−3 )†kTe (keV)
kTinit (keV)
ZO (solar)
ZNe (solar)
ZMg (solar)
ZSi (solar)
ZS (solar)
ZAr (solar)
ZCa (solar)
ZFe = ZNi (solar)
ne t (1011 s cm−3 )V EM (1057 cm−3 )†kTe (keV)
kTinit (keV)
Zall (solar)
ne t (1011 s cm−3 )V EM (1057 cm−3 )†Centroid (keV)
Normalization‡
Photon index
Normalization∗
χ2ν (ν)

SE
SE
NW
NE
0.6–7.5 keV
0.6–7.5 keV
0.6–7.5 keV
0.6–7.5 keV+0.01
0.90 +0.02
0.88 +0.01
0.74 +0.02
0.85 −0.01
−0.01
−0.02
−0.03
+0.01
0.28 +0.01
0.22 +0.01
0.26 +0.02
0.23 −0.01
−0.03
−0.01
−0.02
+0.1
1 (fixed)
1 (fixed)
1 (fixed)
1.3 −0.1
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)+2.2
8.9 +2.1
14.2 +1.6
3.5 +1.3
35.1 −3.9
−0.6
−1.9
−1.0
+0.01
+0.01
+0.06
+0.01
VVRNEI2
0.26 −0.01
0.19 −0.01
0.22 −0.01
0.24 −0.01(RPcold comp.)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)+0.6
1 (fixed)
1 (fixed)
1 (fixed)
2.4 −0.3
+0.2
+0.2
+0.2
+0.4
2.4 −0.1
3.2 −0.2
1.5 −0.2
3.8 −0.3
+0.1
+0.1
+0.3
1.5 +0.1
1.6
1.1
2.2
−0.1
−0.1
−0.1
−0.1
+0.1
+0.1
+0.1
+0.4
1.5 −0.1
1.8 −0.1
0.8 −0.1
3.9 −0.2
+0.1
+0.1
+0.4
0.9 +0.1
1.2
0.5
2.6
−0.1
−0.1
−0.1
−0.1
+0.1
+0.1
+0.2
0.8 −0.1
1.2 −0.1= ZS
2.4 −0.2+0.2+0.3
0.8 +0.2
1.2
=
Z
2.2
S
−0.1
−0.2
−0.2
+0.1
+0.1
+0.1
+0.1
0.1 −0.1
0.5 −0.1
0.2 −0.1
0.8 −0.1+0.1
4.1 +0.2
4.2 +0.2
9.7 +0.9
5.4 −0.1
−0.1
−0.1
−1.1
+0.5
+0.7
+0.8
+0.4
21.7 −0.9
12.6 −1.0
2.4 −1.4
7.0 −0.5+0.01
VVRNEI3————
0.54 +0.03
0.63 +0.02
0.61 −0.01
−0.01
−0.01
(RPhot comp.)————
5 (fixed)
5 (fixed)
5 (fixed)————= RPcold= RPcold= RPcold————= RPcold= RPcold= RPcold
+0.1
+0.4
————
4.1 +0.2
1.8
6.8 −0.9−0.4−0.2
Gaussian
1.23 (fixed)
1.23 (fixed)
1.23 (fixed)
1.23 (fixed)+8+16
132 +1556
49
142 +18
−15
−9
−6
−19
Power law
1.89 (fixed)
1.89 (fixed)
————
————
171 +22130 +12
————
————
−20
−14
1.71 (1126)
1.47 (1124)
1.23 (838)1.67 (1123)∫†
Volume emission measure at the distance of 1.5 kpc: V EM = ne nH dV , where ne , nH , and V are the electron and hydrogendensities, and the emitting volume, respectively.
‡
The unit is photons s−1 cm−2 sr−1 .
∗The unit is photons s−1 cm−2 keV−1 sr−1 at 1 keV.

Model function
TBabs
VVRNEI1(CIE comp.)

SE
1.6–7.5 keV
0.70 (fixed)
————
————
————
————
0.30 +0.01−0.01
5 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1.1 +0.1−0.1
0.6 +0.1−0.1
0.6 +0.1−0.1
0.6 +0.1−0.1
0.2 +0.1−0.1
4.6 +0.1−0.1
30.5 +1.4
−1.4
————
————
————
————
————
————
————
1.89 (fixed)
151 +23−23
1.55 (594)

Table 6.2: Best-fit model parameters of the SE, NW, and NE spectra.

6.3. ANALYSIS
61

62

CHAPTER 6. IC 443
(a)

(d)

0.00

0.06
0.11
0.17
0.22
kTe of the CIE comp. (keV)

0.28

(b)

5.0

2

6.0
7.0
8.0
9.0
Column density NH (× 1021 cm-2)

10.0

(e)

0.00

0.06
0.11
kTe of the RP

.5
3.51
0.17
0.22
comp.
(keV)
cold

3.53
0.28

0.30

0.37
0.45
0.52
0.60
kTe of the RP hot comp. (keV)

0.67

3.58
3.50

3.85
3.67
3.53
3.67 4.22
4.22 4.94
6.39 6.39
15.0
Ionization parameter net (× 1011 s cm-3)

9.26

(c)

6

Figure 6.5: (a–c) Maps of the electron temperatures kTe of the CIE, RPcold , RPhot components.

(d) Map of the column density NH with the 12 CO(J = 1 − 0) emissions (cyan contours; Yoshiike,
2017). (e) Map of the ionization parameters ne t.

6.4. DISCUSSION ON IC 443

6.4

Discussion on IC 443

6.4.1

Absorption by the molecular cloud63

Figure 6.5d shows the spatial distribution of the X-ray absorption column densities(hereafter NH,X ). Since our analysis covers almost the whole remnant, we can performa systematic comparison between NH,X and the ambient gas distribution. Our spectral
fits gave high NH,X in regions 1, 2, 5, 6, 13, 14, 16, 17 and 18, which coincide with thelocations of the molecular cloud traced by the 12 CO(J = 1–0) line (e.g., Xu et al., 2011;
Yoshiike, 2017). Our result indicates that the cloud is located in front of the remnant.
If we subtract the lowest NH,X (region 19) from the highest one (region 14), we obtain(3.1 ± 0.6) × 1021 cm−2 , which can be attributed to the absorption by the molecular cloud.
Yoshiike (2017) observed the 12 CO(J = 1–0) line with the NANTEN2 telescope, andconcluded that the molecular cloud is located in front of IC 443 based on line velocity
measurements. We note that Troja et al. (2006) also reached the same conclusion basedon a comparison between XMM-Newton data and CO data. Yoshiike (2017) estimated thecolumn density of the molecular cloud to be NH,CO = (6.1 ± 0.8) × 1021 cm−2 . Although
the estimate by Yoshiike (2017) is somewhat larger than ours, the two measurements canbe regarded to be consistent with each other, considering the fact that the estimate by
Yoshiike (2017) is obtained for the CO emission peak location and that our estimate isan average of a much larger region.

6.4.2

Origin of the CIE plasma

In the spectral analysis, we successfully reproduced the spectra from all the regionswith the model consisting of the CIE, RPcold and RPhot components. The abundances of
the CIE component are consistent with solar, suggesting that the component is of shocked
ISM origin. The spatial distribution of the component also supports this picture. Figure 6.6a-1 shows the normalization ratios of the CIE component to the RPhot component
plotted as a function of the angular distance from the center of the remnant which we
define as (R.A., Dec.) = (6h 17m 40s .0, 22◦ 38′ 00′′ .0). The map of the normalization ratio
is shown in Figure 6.6a-2. The normalization ratios tend to increase toward the outer regions, suggesting that the CIE plasma is concentrated in the rim as is generally expected
for a shocked ISM emission.

6.4.3

Origin of the recombining plasmas

In section 6.3.2, we found that the emission is well reproduced by a model which
includes two RP components with diﬀerent kTe . As shown in Figures 6.5a–6.5c, the electron temperatures of both components decrease toward the southeast. The southeasternregion coincides with the locations where the molecular cloud is known to be interactingwith the remnant based on the 12 CO(J = 2–1) to 12 CO(J = 1–0) intensity ratio (e.g., Xuet al., 2011; Yoshiike, 2017). The normalization ratio of the two RP components points
to a similar tendency. Figures 6.6b-1 and 6.6b-2 show a radial profile and a map of theRPcold to RPhot normalization ratio, respectively. As opposed to Figure 6.6a-1, no clear
correlation is seen in the radial profile. From the map, it is evident that the ratio is

64

CHAPTER 6. IC 443

(a-1)

(a-2)

Norm. (CIE) / Norm. (RP hot )

102

10

1019 cm at 1.5 kpc

1

0
1
2
3
19
cm)0.2
Distance from the center of the remnant
(×
10
05
0.1

(b-1)

Norm. (RP cold ) / Norm. (RP hot )

10

0.05

0.4

0.20

0.79

0.79

1.6

3.2

3.2

13

6.3

Norm. (CIE) / Norm. (RP hot )

50

13

(b-2)

1

1019 cm at 1.5 kpc

10-1
0
1
2
3
Distance from the center of the remnant(×1019 cm)0.04
01
0.02

0.01

0.079

0.04

0.16

0.16

0.32

0.63

0.63

2.5

1.3

Norm. (RP cold ) / Norm. (RP hot )

10

2.5

Figure 6.6: (a-1) Radial profile and (a-2) map of the CIE to RPhot normalization ratio withthe 1.4 GHz radio emissions (green contours). (b-1) Radial profile and (b-2) map of the RPcold
to RPhot normalization ratio with a map of the 12 CO(J = 2 − 1) to 12 CO(J = 1 − 0) intensity
ratio (red contours; Yoshiike, 2017). The yellow diamonds denote the center of the remnant(R.A., Dec.) = (6h 17m 40s .0, 22◦ 38′ 00′′ .0). The length of the white arrows correspond to 1019 cm
at the distance of 1.5 kpc.

6.4. DISCUSSION ON IC 443

65

larger in the southeastern region, where kTe is also lower than the other regions (see Figures 6.5a–6.5c). These results on kTe and the normalization ratio together indicate that
the SNR plasma is cooler in the region where the shock is in contact with the molecular
cloud.
One of the plausible mechanisms to explain the cooler plasma in the southeasternregion is cooling of the X-ray emitting plasma by the molecular cloud via thermal conduction. If the cooling proceeds faster than the recombination timescale of the plasma,
the recombining plasma can also be explained by the thermal conduction. The thermal
conduction scenario is proposed for the formation process of an RP in G166.0+4.3 (Chapter 5). We found kTe of a high-density region of the SNR is significantly lower than that ofa low-density region. The observational result is in fact similar to our findings on IC 443
reported in this chapter. It is interesting to point out that the plasma in the northwesternregion is close to CIE with ne t ∼ 1012 s cm−3 . A possible explanation is that the coolingof the plasma in this region is less eﬃcient since the region is far away from the molecular
cloud, which would be the major cooling source for IC 443. Lower ambient gas density,
which is suggested by the large shell radius in the west (Lee et al., 2008), would also makethe cooling less eﬃcient.

Chapter 7
W44
7.1

Previous Works and Our Aims

W44 (G34.7−0.4) is a middle-aged SNR with an estimated age of ∼ 20 kyr (Wolszczan
et al., 1991) and is located at a distance of ∼ 3 kpc (Claussen et al., 1997). W44 has asynchrotron radio shell (e.g., Kundu & Velusamy, 1972) and a center-filled X-ray emission(e.g., Rho et al., 1994), and thus it is classified into a mixed-morphology SNR. W44 is
known to be interacting with molecular clouds from CO line (e.g., Seta et al., 2004) and1720 MHz OH maser (e.g., Claussen et al., 1997) observations.
In the X-ray band, Rho et al. (1994) obtained spectra with ROSAT, EXOSAT and
Einstein, and interpreted them as an IP emission (White & Long, 1991). Kawasaki et al.
(2005) analyzed W44 spectra obtained with ASCA and measured ionization temperatures
based on the Si Lyα to Heα intensity ratios. They claimed that the plasma is nearly ina CIE state because kTz is almost equal to kTe . In a Suzaku observation of the central
region, Uchida et al. (2012) found RPs. The best-fit parameters indicate that the plasma
was cooled down from kTinit ∼ 1 keV to kTe ∼ 0.5 keV. They also found hard X-ray
emissions inside the SNR and interpreted them as a synchrotron enhancement in thevicinity of a molecular cloud.
After the work by Uchida et al. (2012), five more W44 observations with Suzaku werecarried out to cover the whole remnant. In this chapter, we analyze data from the six
Suzaku observations of W44 and perform spatially resolved X-ray spectroscopy to compare
the results with the distribution of the ambient gas interacting with the remnant. We
then discuss the implications particularly for the formation process of the RP. All spectral
fits are performed with XSPEC version 12.9.0n. The plasma models are calculated with
ATOMDB version 3.0.8. We used solar abundances given by Wilms et al. (2000).

7.2

Suzaku Observations and Data Reduction

We performed Suzaku deep observations of W44, whose log is summarized in Table 7.1.
The total exposure time is about 335 ks. The XIS FOVs of each observation are shownin Figure 7.1a. In this analysis, we used XIS data excluding those from XIS2 and a partof XIS0 because they were not functioning at the time of the observations (§4.3.1). We
reduced the data and estimated NXBs with the same method in Chapter 6.
667.3. ANALYSIS

67

Table 7.1: Suzaku observation log of W44.
TargetW44 center
W44 west
W44 south
W44 south
W44 east
W44 north
HESS J1857-A (Background)
HESS J1857-B (Background)

7.3

Analysis

7.3.1

Image

Obs. ID
505004010
508002010
508003010
508003020
508004010
508005010
506019010
506020010

Obs. date2010-04-10
2013-10-24
2013-10-22
2014-04-09
2013-10-18
2013-10-19
2011-09-18
2011-11-01(R.A., Dec.)

Eﬀective Exposure(18h 56m 08s .4, 1◦ 23′ 19′′ .0)

61 ks
61 ks
67 ks
32 ks
58 ks
56 ks
4 ks
4 ks

(18h 55m 45s .4, 1◦ 19′ 17′′ .8)(18h 56m 14s .7, 1◦ 12′ 36′′ .0)(18h 56m 13s .1, 1◦ 13′ 14′′ .9)(18h 56m 35s .0, 1◦ 17′ 56′′ .8)(18h 55m 50s .4, 1◦ 29′ 34′′ .1)(18h 56m 51s .0, 2◦ 45′ 06′′ .8)(18h 58m 11s .6, 2◦ 44′ 57′′ .8)

Figures 7.1a, 7.1b, 7.1c and 7.1d show XIS3 mosaic images of W44 in the energy bandsof 0.3–10 keV, 0.3–10 keV, 3–7 keV and 0.3–2 keV, respectively. We subtracted NXB fromthe images and corrected for the vignetting eﬀect of the XRT. As shown in Figure 7.1b,
W44 has a center-filled X-ray emission with a radio shell as is previously known (e.g.,
Rho et al., 1994). In the hard band image of Figure 7.1c, we found three hard sources inregions A, B and C. The hard X-ray source in region A is the one discovered by Uchida
et al. (2012). The hard source in region B is located at the position of the PWN (Petreet al., 2002).

7.3.2

X-ray background subtraction

In order to estimate X-ray background for the spectral analysis, we analyzed observation data of nearby targets (HESS J1857-A and HESS J1857-B; Table 7.1). We extracted
spectra from source-free regions and merged the spectra from the two observations. We
fitted the spectrum with a model by Uchiyama et al. (2013), who studied the X-ray emission from the Galactic plane. The model consists of three components; the Galactic ridge
X-ray emission (GRXE), the CXB and the foreground emission (FE).
The GRXE model consists of low-kTe thermal plasma (LP), high-kTe thermal plasma(HP), an emission from cold neutral matter (cold matter: CM) expressed by a powerlaw, and neutral Fe Kα (6.40 keV) and Kβ (7.06 keV) lines. We allowed the absorption
column density of the GRXE and normalizations of the LP, HP and CM components tovary. The other parameters were fixed at the values given by Uchiyama et al. (2013).
A normalization of the FE component was allowed to vary, whereas we fixed the otherparameters of the FE at the values by Uchiyama et al. (2013). Figure 7.2 and Table 7.2
show the fitting result and the best-fit parameters, respectively.

68

CHAPTER 7. W44

53:0

Right ascension(a)

18:57:00.0

55:00.0

56:00.0(b) -2.97e-08-3.22e-07

Whole

20:00.0

west

south

0.0

3–7 keV

1:10:00.0

Declination

centereast(c)

30:00.0

north

40:00.0

58:00.0×2

0.3–10 keV

0.2

0.4

0.6

2.59e-07

5.51e-07

8.41e-070.8

1.0

9.8×10-4

3.9×10-3(d) 2.00
1.13e-06
1.42e-06 2.36e-07
1.71e-06 5.54e-07
0.3–10e+00 keV
7.90e-08

1.6×10-2

1.18e-06

6.3×10-2

0.25

1.0

0.3–2
keV2.00e-05
9.95e-06

2.45e-06

4.95e-06

9
4
3

8

7

9.8×10-4

3.9×10-3

1.6×10-2

6.3×10-2

0.25

1.0

9.8×10-4

3.9×10-3

C

A

2

1

B

1.6×10-2

510

6

6.3×10-2

0.25

Figure 7.1: XIS3 images of W44 in the energy band of (a,b) 0.3–10 keV, (c) 3–7 keV and (d)0.3–2 keV after the NXB subtraction and correction for the vignetting eﬀect. The X–ray countrates are normalized so that the peaks become unity. In panel (a), the green squares are theFOVs of XIS3 in each observation. The coordinate refers to the J2000.0 epoch. In panel (b), thegreen contours are a radio image as seen in 20 cm wavelength by the Very Large Array (VLA)(Helfand et al., 2006). The magenta and cyan contours are the NANTEN2 12 CO(J = 2–1)
images in the velocity range of 40–50 km s−1 and 50–60 km s−1 , respectively (Uchida et al.,2012). The orange stars indicate the locations of the 1720 MHz OH maser emissions (Claussen
et al., 1997). The white lines in panels (c) and (d) indicate regions for the spectral extraction.

1.0

7.3. ANALYSIS

69

Table 7.2: Best-fit model parameters of the background spectra for W44.
ComponentLP

Model function
TBabs
VAPEC

HP

VAPEC

CM

Power law

Absorption GRXB

Gaussian1
Gaussian2
Absorption CXBCXBAbsorption FE

FE

TBabs
Power law
TBabs
APEC1

APEC2

Parameter
NH (1022 cm−2 )kTe (keV)
ZAr (solar)
Zother (solar)
Normalization†kTe (keV)
ZAr (solar)
Zother (solar)
Normalization†
Photon index
Normalization∗
Centroid (keV)
Normalization‡
Centroid (keV)
Normalization‡
NH (1022 cm−2 )Photon index
Normalization∗
NH (1022 cm−2 )kTe (keV)
Zall (solar)
Normalization†kTe (keV)
Zall (solar)
Normalization†
χ2ν (ν)Value
2.7 +0.2−0.3
1.33 (fixed)
1.07 (fixed)
0.81 (fixed)
126 ± 30
6.64 (fixed)
1.07 (fixed)
0.81 (fixed)
98 +7
−16
2.13 (fixed)< 9.0
6.4 (fixed)= Norm Power law × 1.103

7.06 (fixed)= Norm Gaussian1 × 0.125= Absorption GRXB × 2

1.4 (fixed)
9.9 (fixed)
0.56 (fixed)
0.59 (fixed)
0.05 (fixed)
68 ± 80.09 (fixed)
0.05 (fixed)= Norm APEC1 × 261.5∗ The unit is photons s−1 cm−2 keV−1 sr−1 at 1 keV.
†The emission measure integrated over the line of sight, i.e., (1 / 4π)
1014 cm−5 sr−1 .
‡The unit is photons s−1 cm−2 sr−1 .

7.3.3

1.37 (252)∫

ne nH dl in units of

SNR spectra

We extracted a spectrum from the Whole region indicated in Figure 7.1c. The regions A, B and C, where the hard X-ray emissions are located, are excluded from theWhole region. Figure 7.3 shows the XIS0+3 spectrum in the 0.6–5.0 keV band. Lineemissions of Ne, Mg, Si, S and Ar are clearly seen. We also found Si and S RRCs, which
provide clear evidence that the plasma is in a recombination-dominant state.
We fitted the spectrum with the model consisting of the SNR component and thebackground (§7.3.2). A previous X-ray study found that the W44 spectrum in the centeralregion can be reproduced by an RP model with kTe ∼ 0.5 keV, kTinit ∼ 1 keV andne t ∼ 6 × 1011 s cm−3 (Uchida et al., 2012). Following them, we first applied an RP modelto the Whole spectrum, by using the VVRNEI model in XSPEC. We allowed NH , kTe ,kTinit , ne t and the normalization of the RP component to vary. The abundances of Ne,
Mg, Si, S, Ar, and Fe were also allowed to vary, whereas the O, Ca, and Ni abundances

CHAPTER 7. W44

Counts s −1 keV−1

70

0.1
0.01
10−3
10−4

χ

5
0
−5
1

2

Energy (keV)

5

10

Figure 7.2: X-ray background spectra of XIS0+3 (black crosses) and XIS1 (red crosses) extracted from the source-free regions of HESS J1857-A and HESS J1857-B. The black and redsolid lines indicate the best-fit models of XIS0+3 and XIS1, respectively. The LP, HP, CM,
CXB, and FE components of XIS0+3 show the solid cyan, magenta, blue, green, and orange
lines, respectively. The bottom panel shows the residuals from the best-fit model.

were linked to Ne, Ar, and Fe, respectively. The abundances of the other elements were
fixed to solar. For the XIS0+3 spectrum, we ignored the 1.78-1.92 keV band because of the
known calibration uncertainty around the neutral Si K-edge. The fit left residuals around1.23 keV, most probably due to the lack of Fe-L lines in the model (e.g., Yamaguchi et al.,2011), and, therefore, we added a Gaussian at 1.23 keV. The result is shown in Figure 7.4aand Table 7.3. The spectrum in the high energy band above 1.7 keV is well reproduced by
the RP model, indicating that the plasma has a component in a recombination-dominantstate. On the other hand, the fit left residuals in the low energy band of 0.7–1.7 keV,suggesting that the plasma has an additional component in the soft band.
To account for the soft band residuals, we tried the CIE, IP and RP models. Abundances of Ne and Mg were allowed to vary, whereas that of O was linked to Ne. The otherabundances were fixed to solar. We allowed kTe and a normalization of the soft-band component to vary. The residuals in the soft band was significantly improved independentof the choice of the model. Nonetheless, we here chose to use the CIE model, following
the IC 443 analysis in §6.4.2. We present the fitting result and the best-fit parametersin Figure 7.4b and Table 7.3, respectively. We found that kTe of the CIE component is
0.22 ± 0.01 keV, which is almost the same as those of IC 443. The best-fit parametersof kTe , kTinit and ne t of the RP component are 0.37 ± 0.01 keV, 0.93 ± 0.01 keV and
11(3.5+0.2
s cm−3 , respectively. The obtained kTe is lower than that given by Uchida−0.3 ) × 10
et al. (2012), possibly indicating that the plasma in the Whole region on average is colder
than the plasma in the central region, where Uchida et al. (2012) analyzed. The obtained
ne t is smaller than that by Uchida et al. (2012), possibly indicating that the recombination timescale of the outer plasma is shorter than that of the central plasma. In order7.3. ANALYSIS

71

Ne

Mg

Si
S

1

Ar
Fe L complex

0.01

Si RRC
S RRC

0.1

Counts s −1 keV−1

10

to conduct spatially resolved spectroscopy, we divided the Whole region into 10 regions
as shown in Figure 7.1d and extracted spectra from each region. We fitted the spectrawith the above model but we fixed kTinit and the O, Ne, and Mg abundances of the CIE
component and the O and Ne abundances of the RP component to the best-fit values in
the Whole fit. In most of the regions, the spectra can be reproduced well by the model.
The fit for the region 1 spectrum, however, left residuals in the high energy band above
4 keV, indicating that the spectrum has a hard X-ray component. An additional fit for
the hard X-ray component is shown in the next subsection (§7.3.4). Each spectrum withthe best-fit models is plotted in Appendix B.1.

1

2
Energy (keV)

5

Figure 7.3: XIS0+3 spectra of the Whole region in the 0.6–10 keV band after NXB subtraction.
Emission lines of Ne, Mg, Si, S and Ar can be seen in the spectrum. The arrows indicate theenergy bands of the Fe L complex and the RRCs of Si and S.

72

CHAPTER 7. W44(a)Counts s −1 keV−1

10

Whole

1
0.1
0.01
10−3

5

χ

0
−5

1

(b)

2
Energy (keV)

5Counts s −1 keV−1

10

10

Whole

1
0.1
0.01
10−3

χ

5
0

−5

1

2
Energy (keV)

5

10

Figure 7.4: XIS0+3 (black crosses) and XIS1 (red crosses) spectra of the Whole region after
NXB subtraction. The magenta, blue and orange solid lines denotes the SNR models for XIS0+3,
consisting of the CIE, RP and Gaussian components, respectively. The dotted lines show thebackground model for XIS0+3.

7.3. ANALYSIS73

Table 7.3: Best-fit model parameters of the spectrum in the Whole region.

Model function

Whole

Whole

(RP)

(CIE+RP)

Parameter

NH (1022 cm−2 )kTe (keV)
ZO = ZNe (solar)
ZMg (solar)
Normalization†kTe (keV)
kTinit (keV)
ZO = ZNe (solar)
ZMg (solar)
ZSi (solar)
ZS (solar)
ZAr = ZCa (solar)
ZFe = ZNi (solar)
ne t (1011 s cm−3 )Normalization†
Centroid (keV)
Normalization‡
χ2ν (ν)1.92 ± 0.01
————
————
————
————
0.34 ± 0.01
0.97 ± 0.01
1.3 ± 0.1
1.6 ± 0.1
2.4 ± 0.1
2.7 ± 0.2
7.9 ± 0.9
0.31 ± 0.31
4.9 ± 0.1
3.14 ± 0.13
1.23 (fixed)
26 ± 22.31 (1227)

2.00 +0.02−0.01
0.22 ± 0.01
0.26 +0.14−0.12
0.94 +0.10−0.13
6.61+0.80−1.42
VVRNEI2
0.37 ± 0.01
(RP comp.)
0.93 ± 0.02
3.6 +0.3−0.2
2.8 ± 0.1
4.5 +0.2−0.1
4.3 +0.2−0.1
8.1 +1.2−1.1< 0.193.5 +0.2−0.3
1.23 +0.02−0.04
Gaussian
1.23 (fixed)
26 +2−3
1.92 (1223)∫†
The emission measure integrated over the line of sight, i.e., (1 / 4π) ne nH dl
in units of 1018 cm−5 sr−1 .
‡The unit is photons s−1 cm−2 sr−1 .
TBabs
VVRNEI1(CIE comp.)

7.3.4

Hard X-ray sources

In the imaging analysis, we found the hard X-ray sources in regions A, B and C. For
spectral analyses of the regions, we extracted spectra from each region and fitted themwith the model accounting for the sum of the emissions from the hard X-ray source, SNR
plasma, and X-ray background. Models for the hard X-ray sources are chosen based on
the characteristics of their emissions as described below. We applied the models obtainedin §7.3.2 for the X-ray background. For the SNR component, we used the best-fit model ofthe Whole spectrum in §7.3.3 but we allowed kTe , ne t, the Mg, Si, S, Ar and Fe abundancesof the RP component, and the normalization to vary. In the spectral analysis of region 1(§7.3.3), we found the residuals in the hard band so that we refitted the spectrum withthe model including a hard component.

74

CHAPTER 7. W44

(A-1)

(A-2)

0.1
0.01

0

0

χ

5

2
Energy (keV)

5

0.01

10−4

5

5

0

0

χ

10−4

(1-1)
10

5

0.01

0.01

10−4
5

0

0

χ

10−4
5

−5
5

10

10

0.1

10−3

2
Energy (keV)

5region 1

1

10−3

1

2
Energy (keV)

10

0.1

−5

1

(1-2)
region 1

1

−5

10

10

0.01
10−3

2
Energy (keV)

5region C

0.1

10−3

1

2
Energy (keV)

1

0.1

−5

1

(C)
region B

1

−5

10

Counts s −1 keV−1

Counts s −1 keV−1

0.01

5

10

χ

0.1

10−3

1

region A

1

10−3

(B)

Counts s −1 keV−1

Counts s −1 keV−1

1

−5

χ

10

region A

Counts s −1 keV−1χCounts s −1 keV−1

10

1

2
Energy (keV)

5

10

Figure 7.5: XIS0+3 (black crosses) and XIS1 (red crosses) spectra in the regions A, B, C
and 1 of W44. Each spectrum is fitted with the model consisting of the SNR and backgroundcomponents. The solid and dotted lines show the SNR and background models of XIS0+3,
respectively. The SNR model consists of CIE (magenta lines), RP (blue lines), power-law (green
lines), APEC (purple line), and Gaussian (orange lines) components, respectively.

ParameterNH (1022 cm−2 )kTe (keV)
Zall (solar)
Normalization†kTe (keV)
kTinit (keV)
ZO = ZNe (solar)
ZMg (solar)
ZSi (solar)
ZS (solar)
ZAr = ZCa (solar)
ZFe = ZNi (solar)
ne t (1011 s cm−3 )Normalization†
Centroid (keV)
Normalization‡
NH (1022 cm−2 )Photon index
Normalization∗
Centroid (keV)
Normalization‡kTe (keV)
Zall (solar)
Redshift (z)Normalization†
χ2ν (ν)1.92 ± 0.01
0.21 ± 0.01
= Zall,Whole
8.35 ± 1.04
0.44 ± 0.01
0.93 (fixed)
3.6 (fixed)
3.1 ± 0.1
5.9 ± 0.2
5.4 ± 0.2
7.4 ± 0.7
0.23 ± 0.03
3.4 ± 0.1
1.65 ± 0.03
1.23 (fixed)
58 ± 4= NH, SNR
2.48 ± 0.15
152 ± 37
————
————
————
————
————
————
2.01 (1116)

Region B(SNR+PL)(SNR)

Region C

SNR comp.

1.94 ± 0.02
1.83 ± 0.03
2.21 +0.05−0.04
0.21 ± 0.01
0.22 ± 0.02
0.22 +0.01−0.02= Zall,Whole= Zall,Whole= Zall,Whole
9.77 +1.77
5.92 +2.51
8.61 +2.82
−1.56
−2.40
−3.16
VVRNEI2
0.45 ± 0.01
0.28 ± 0.01
0.38 +0.01−0.02(RP comp.)
0.93 (fixed)
0.93 (fixed)
0.93 (fixed)
3.6 (fixed)
3.6 (fixed)
3.6 (fixed)
3.0 ± 0.1
2.8 +0.1
2.4 +0.3
−0.2
−0.2
+0.2
5.7−0.3
5.1 ± 0.3
6.6 +0.5−0.3
5.1 +0.4
5.3 +0.8
5.8 +0.6
−0.3
−0.6
−0.5
7.0 +1.1= ZS
10.0 +3.2
−1.0
−2.6
+0.05
+0.3
0.21 −0.06
1.0 −0.2< 0.193.3 ± 0.1
4.9 +0.3
2.4 +0.2
−0.2
−0.1
+0.03
+0.14
1.73 −0.04
1.73 −0.10
1.80 ± 0.09
Gaussian
1.23 (fixed)
1.23 (fixed)
1.23 (fixed)
57 ± 632 ± 723 +16−15
Hard X-ray
TBabs= NH, CXB
= NH, SNR
————
comp.
Power law————
2.34 ± 0.08
————
————
405 +52————−47
Gaussian
————
————
————
————
————
————
APEC
3.67 +0.67
————
————
−0.46
0.41 +0.10
————
————
−0.07
0.091 (fixed)
————
————
0.034 +0.006
————
————
−0.0071.93 (1115)
1.19 (1117)
1.26 (549)∫†The emission measure integrated over the line of sight, i.e., (1 / 4π) ne nH dl in units of 1018 cm−5 sr−1 .
‡The unit is photons s−1 cm−2 sr−1 .
∗The unit is photons s−1 cm−2 keV−1 sr−1 at 1 keV.

Model function
TBabs
VVRNEI1(CIE comp.)

Region A(SNR+APEC)

Region A(SNR+PL)

1.97 ± 0.02
0.23 +0.02−0.01= Zall,Whole
5.60 +1.97−1.51
0.35 +0.01−0.02
0.93 (fixed)
3.6 (fixed)
3.4 ± 0.1
6.1 +0.4−0.3
6.5 +0.6−0.8
14.6 +5.4−3.4< 0.06
4.0 +0.3−0.1
2.29 ± 0.10
1.23 (fixed)
39 ± 9= NH, SNR
1.69 +0.66−0.88
12 +18
−9
————
————
————
————
————
————
1.37 (1098)(SNR+PL)

Region 1

Table 7.4: Best-fit model parameters of the spectrum in the regions A, B, C and 1.
Region 1

1.97 ± 0.02
+0.02
0.23 −0.01= Zall,Whole
5.59 1.99
−1.53
+0.01
0.35 −0.02
0.93 (fixed)
3.6 (fixed)+0.1
3.4 −0.2+0.4
6.1 −0.3+0.6
6.5 −0.8+5.8
14.8 −3.4< 0.05+0.2
4.0 −0.1+0.15
2.29 −0.08
1.23 (fixed)
40 +8
−10
= NH, SNR
+0.79
1.74 −0.74
+31
11 −8
6.40 (fixed)+0.17
0.38 −0.18
————
————
————
————
1.36 (1097)(SNR+PL+Gauss)

7.3. ANALYSIS
75

76

CHAPTER 7. W44

Region A
We first used a power-law model for the hard X-ray component. The fitting results
are shown in Figure 7.5 (A-1) and the best-fit parameters are summarized in Table 7.4.
The model gives a satisfactory fit in the 0.6–10 keV band except for the line-like residuals
around 6.1 keV. We, therefore, added a Gaussian with its centroid energy treated asa free parameter. The fit with the Gaussian significantly improved the fitting and thecentroid energy was determined to be 6.14 ± 0.02 keV. There are major atomic lines atthis energy in the rest frame. The most probable explanation would be a redshifted Feline. A promising candidate of a redshifted Fe line emitter is an active galactic nucleus ora galaxy cluster. The hard X-ray source (Figure 7.1c) is clearly extended, and therefore,
most likely is a galaxy cluster. Considering of the typical temperature and ionization stateof a plasma in a galaxy cluster, the line should be Fe Heα at 6.70 keV. Then, the redshift
is estimated to be z = 0.091 which is consistent with the galaxy cluster explanation.
An X-ray emission from a galaxy cluster is generally reproduced by a CIE model.
Therefore, we adopted the APEC model as the hard component. The absorption column
density was fixed at the total Galactic absorption in the line of sight toward W44. The
model gave a satisfactory fit (Figure 7.5 (A-2) and Table 7.4). The abundance and kTe+0.67
of the galaxy cluster were determined to be 0.41+0.10−0.07 and 3.67−0.46 keV, respectively. The
flux is 1.6 × 10−12 ergs cm−2 s−1 at 0.5–2.0 keV. By systematically analyzing of ROSAT
data of a number of galaxy clusters, Fairley et al. (2000) reported relationship between
the bolometric luminosity (L) and electron temperature (kTe ). The obtained flux of thehard X-ray source can be translated into a luminosity of 5.6×1043 erg s−1 with the Hubble
constant of 50 km s−1 Mpc−1 assumed by Fairley et al. (2000). We converted the 0.5–
2.0 keV flux to the bolometric luminosity by multiplying a factor of 2–3 (Fairley et al.,2000), and obtained L = (1.1–1.7) × 1043 erg s−1 . The obtained L and kTe are consistent
with the L-kTe relationship by Fairley et al. (2000) as shown in Figure 7.6. We, therefore,
conclude that the hard X-ray source in region A is a background galaxy cluster.
Region B
We fitted the spectrum in region B with the model composed of the SNR, backgroundand power-law components. The spectrum can be reproduced well by the model as in
Figure 7.5 (B) and Table 7.4. The photon index and flux at 2–10 keV are 2.34 ± 0.08
and 4.0 × 10−12 ergs cm−2 s−1 , respectively. The photon index is consistent with thatof the PWN measured by Chandra (Petre et al., 2002), whereas the flux is an orderof magnitude higher than their measurement. The Suzaku image of the hard emission
spreads more than the PWN image with Chandra. Our result suggests that the PWN is
spatially extended more than previously known through the Chandra observation (Petreet al., 2002).
Region C
We performed a spectral analysis of region C and found that the spectrum can be well
reproduced by a model composed of the SNR and background components (Figure 7.5(C) and Table 7.4). Therefore, we conclude that no additional hard X-ray component isnecessary for this region.

7.3. ANALYSIS

77

Figure 7.6: L–kTe relation for z < 0.2 (Fairley et al., 2000). The black plots represent groupsof galaxy clusters from the Wide Angle ROSAT Pointed Survey (WARPS) bins, Helsdon &Ponman (2000), Mushotzky & Scharf (1997) and David et al. (1993). The line shown is thebest-fitting line to clusters with kTe > 1 keV, L ∝ kTe 3.15 . The red plot indicate the galaxy
cluster discovered by us.

Region 1
The spectrum of region 1 was refit with a model with a power-law component inaddition to the SNR and background components. Figure 7.5 (1-1) and Table 7.4 show
the fitting results. The spectra can be mostly reproduced by the model but residuals
around 6.4 keV are evident (see Figure 7.7). The residuals was improved by the Gaussian
whose centroid energy is 6.38±0.04 keV, which corresponds to the energy of the neutral or
low ionized Fe Kα line. However, a low-ionized Fe line is unlikely since Fe ions in middleaged SNRs such as W44 are expected to be highly ionized. The line is most probably
from neutral Fe. One of the plausible processes to emit the line is K-shell ionization of
Fe atoms in dense cloud by electrons or protons accelerated in the SNR. Such scenarios
are proposed to explain the neutral Fe Kα line from the Galactic ridge (Nobukawa etal., 2015) as well as from the SNR 3C391 (Sato et al., 2014) and Kes 79 (Sato et al.,2016). Following them, we fitted the spectrum with the model consisting of the SNR,
background, power-law components and a Gaussian whose centroid energy was fixed at
the line energy of the neutral Fe (6.40 keV). The fit improved the residuals around 6.4 keV
as shown in Table 7.4 and Figures 7.5 (1-2) and 7.7. The improvement is significant with
the F-test null probability of 1.9 × 10−3 .

CHAPTER 7. W44

Counts s −1 keV−1

78

0.015
0.010
0.005

Counts s −1 keV−1

0
0.015
0.010
0.005
0

6

6.5
Energy (keV)

7

Figure 7.7: XIS0+3 spectra from region 1 in the energy band of 5.8–7.3 keV with the modelsame as that in (Top) Figure 7.5 (1-1) and (Bottom) Figure 7.5 (1-2).

7.4

Discussion on W44

In Figure 7.8, we show the maps of NH , kTe of the CIE component, kTe and ne t of the
RP component, and normalization ratio of the CIE to the RP components in each region.
We also show the synchrotron and CO radio images and the locations of the 1720 MHz
OH maser emissions in the figure.

7.4.1

Ambient gas and X-ray absorption

As shown in Figure 7.8a, our spectral analysis has revealed the spatial distributionof the X-ray absorption column densities (NH,X ). We compare the NH,X distribution
with that of the ambient gas traced by radio line emissions. The location with higher
NH,X coincides with the 12 CO(J = 2–1) line in the 40–50 km s−1 range, indicating thatthe cloud is located in front of the remnant. The area-weighted average of NH,X ofregions 3, 8, and C, where the 12 CO(J = 2–1) line emission is detected, is calculated
to be 2.18 × 1022 cm−2 . On the other hand, the average across the other regions is
1.96 × 1022 cm−2 . The diﬀerence of the two should correspond to the column density of
the cloud responsible for the 12 CO(J = 2–1) line. The column density of the molecular
cloud was estimated by Seta et al. (2004) at ∼ 4 × 1021 cm−2 with the Nobeyama 45m
radio telescope. It is almost the same as the NH,X diﬀerence, supporting the suggestionthat the cloud is located in front of the remnant.

7.4. DISCUSSION ON W44

79

(a)

(b)

1.70
1.88

1.82

(c)

0.15
0.21

(e)

0.24

2.00
2.05

1.82
1.94
2.06
2.18
Column density NH (× 1022 cm-2)
1.94

2

2.06

15 2.12

0.21
0.27
0.33
0.39
kTe of the RP comp. (keV)
0.27

0.3

0.33

0

0.36

0.15

2.30
0.16 2.18

0.17 2.24

(d)

0.45
0.5 0.39

0.25

1.0
2.0
3.0
4.0
net of the RP comp. (× 1011 s cm-3)

5.0

0.18 2.

0.0
1

0.17
0.19
0.21
0.23
kTe of the CIE comp. (keV)

0.42

1.5 0.4

11

2

0.19

2

0.2

2.5

0.21

3

0.22

3.5

0.23

4

2.05
2.27
3.12
6.52
20.0
Norm. (CIE comp.) / Norm. (RP comp.)

2.13

2.27

2.55

3.12

4.25

6.52

Figure 7.8: Maps of (a) NH , (b) kTe of the CIE component, (c) kTe of the RP component,(d) ne t of of the RP component and (e) the normalization ratio of the CIE to RP components,
respectively. In panel (a–d), we overlaid the 12 CO(J = 2–1) images in the velocity range of 40–
50 km s−1 (magenta contours) and 50–60 km s−1 (cyan contours) obtained with the NANTEN2(Uchida et al., 2012). The orange stars indicates the locations of the 1720 MHz OH maser
emissions (Claussen et al., 1997). The green contours in panel (e) are a radio image as seen in
20 cm wavelength by the VLA (Helfand et al., 2006).

80

7.4.2

CHAPTER 7. W44

Origin of the CIE plasma

In the spectral analysis, we found that the SNR plasma emissions in the whole remnant
are well reproduced by the model consisting of the CIE and RP components. Most of
the abundances of the CIE component are close to solar (0.9–1.0), suggesting that theorigin is a shocked ISM. Figure 7.8e shows the distribution of the normalization ratios of
the CIE to RP components. At each azimuthal direction, the ratio is higher in the outerregions, indicating that the CIE plasma dominants in the rim. The results of W44 are
the same as those of IC 443 (Chapter 6), and further support the shocked ISM origin for
the CIE component.

7.4.3

Origin of the recombining plasma

In the spectral analysis, we found that the electron temperatures of the RPs in thecentral part of the remnant (regions A and 4) are higher than those in the other regions (Figure 7.8c). In the thermal conduction scenario, the SNR plasma is cooled bya molecular cloud as already discussed in Chapters 5 and 6. The intensity ratio of the
12
CO(J = 2–1) to 12 CO(J = 1–0) emissions (Figure 7.9: Yoshiike, 2017) is distributed inregions excluding these central regions, implying the interaction between the SNR shockand the molecular cloud only in the outer regions. The OH maser emissions (Claussen
et al., 1997) also support the picture. The lower kTe in the outer regions than those inregions A and 4 can be explained by the thermal conduction between the plasma and theinteracting molecular cloud if the plasmas in the central regions are not cooled completely
until the conduction.
The spectral analysis also revealed that ne t of the eastern part of the remnant is lower
than those of the other regions (Figure 7.8d). The regions with the lower ne t coincides withthe location of the molecular cloud indicated by the 12 CO(J = 2–1) emission. Assumingthat the electron density of the RP is uniform, we consider that the SNR shock recently
started to interact with the molecular cloud. In the southern part of W44, ne t is higher
than those of the plasma in the other regions (Figure 7.8d), suggesting that the shockcloud interaction started in the early phase of the SNR evolution. If this is the case, thelower kTe of the RP in the south can be explained as well. On the other hand, Uchida etal. (2012) claimed that a conduction timescale is much longer than the SNR age, and thus
the thermal conduction scenario is unlikely. We discuss the conduction timescale later in§9.3.
Let us try to explain the observed ne t distribution in the context of the alternative
scenario, the rarefaction scenario. Shimizu et al. (2012) carried out 3D-hydrodynamical
calculations for the interaction of expanding supernova ejecta with the CSM and the ISM
outside and concluded that the CSM radius needs to be smaller than 1 pc. Since theW44 radius of ∼ 10 pc is much larger than 1 pc, we suppose that the shock broke out of
the CSM in the early phase of the SNR evolution. Since the time before the break fromthe CSM is much shorter than that after the break, we can regard that the rarefaction
occurred almost at the same time independent of the location. It is not easy to explain
the spatial dependence of ne t in W44.

7.4. DISCUSSION ON W44

0.30

0.48

81

0.66

0.84

1.0

Intensity ratio of 12CO J=2–1/J=1–0

1.2

Figure 7.9: Spatial distribution of the intensity ratio of the 12 CO(J = 2–1) to 12 CO(J = 1–0)

emissions in the velocity range of 30–40 km s−1 with the NANTEN2 observations (Yoshiike,
2017). The solid green lines show the regions used for the spectral analysis.

Chapter 8
W49B
8.1

Previous Works and Our Aims

W49B (G43.3−0.2) is a mixed-morphology SNR with center-filled X-ray emissions anda radio shell (Pye et al., 1984) and is located at ∼ 8 kpc (Radhakrishnan et al., 1972;
Moﬀett & Reynolds, 1994). The age was estimated at ∼ 6000 yr (Smith et al., 1985),
which is one of the youngest mixed-morphology SNRs. Most of the mixed-morphology
SNRs are interacting with atomic or molecular gas. W49B is one of such remnants.
Based on infrared observations, Keohane et al. (2007) claimed that shocked molecularhydrogens are associated with eastern and southwestern parts of W49B (Figure 8.1). A
CO observation revealed a molecular cloud in the southwest (Figure 8.2: Simon et al.,2001). In the western part of the remnant, GeV gamma rays were found by Abdo et al.
(2010) and H. E. S. S. Collaboration et al. (2016) with Fermi-LAT as shown in Figure 8.3.
TeV gamma rays were also detected in W49B by H. E. S. S. Collaboration et al. (2016). If
the origin of the gamma-ray emissions is π 0 decays, the results also provide a hint of theshell-cloud interaction.
Kawasaki et al. (2005) observed W49B with ASCA and measured ionization temperatures. They found that kTz of ∼ 2.5 keV is higher than kTe of ∼ 1.8 keV, indicatingthat the W49B plasma is overionized. Ozawa et al. (2009b) analyzed Suzaku data andfound a strong Fe RRC and cascade lines which is robust evidence for the presence of RP
in W49B. Lopez et al. (2013b) performed a spatially resolved spectroscopy with Chandra
and measured intensity ratios of Si Lyα/Si Heα and Ar Lyα/Ar Heα in each region. They
claimed that the degree of ionization becomes higher from east to west, with more ionizedplasma toward the west of the remnant.
The previous X-ray studies applied phenomenological models consisting of a CIE component, Gaussians and RRCs for spectral fits (Kawasaki et al., 2005; Ozawa et al., 2009b;
Lopez et al., 2013b). We here adopted a more realistic model, VVRNEI, in order toextract more accurate physical parameters. All spectral fits are performed with XSPEC
version 12.9.0n and the plasma models are calculated with ATOMDB version 3.0.8. We
used solar abundances given by Wilms et al. (2000).

82

8.1. PREVIOUS WORKS AND OUR AIMS

83

Figure 8.1: X-ray emission (blue), 2.12 µm shocked molecular hydrogen (red) and 1.64 µm [FeII](green) images of W49B. The Ks -band image of foreground and background stars is overlaid inwhite. The figure is taken from Keohane et al. (2007).

Figure 8.2: 13 CO(J = 1–0) integration intensity from 0.5 to 9.7 K km s−1 with the Milky
Way Galactic Ring Survey, a Boston University and Five College Radio Astronomy Observatory
collaboration (Simon et al., 2001). The green contours indicate the radio intensity of W49B obtained at a wavelength of 20 cm with the Multi-Array Galactic Plane Imaging Survey (MAGPIS;
Helfand et al., 2006).

84

CHAPTER 8. W49B

Fermi-LATb (deg)

-0.15

-0.2

-0.25

43.34 43.32 43.3

43.28 43.26 43.24 43.22 43.2 43.18
l (deg)Figure 8.3: (Left) Fermi-LAT counts map in units of counts per pixel in the 6–30 GeV band
taken from Abdo et al. (2010). The LAT localization is represented by a black circle with aradius of 0◦ .024 (90% confidence level) centered at (α, δ) = (287◦ .756, 9◦ .096). The magenta
and green contours indicate W49B and W49A in the Spizer IRAC 5.8 µm, respectively. (Right)
Gamma-ray excess map of W49B with the H.E.S.S. taken from H. E. S. S. Collaboration et al.
(2016). The crosses show the best-fit centroid positions of gamma-ray signal with the H.E.S.S.
(≥ 600 GeV; black) and Fermi-LAT (≥ 1 GeV; blue). The size of the crosses represents the sizeof the 95% confidence level contours, including systematic uncertainties. The while contours
show the radio intensity obtained at a wavelength of 20 cm with the MAGPIS.

8.2

Suzaku Observations and Data Reduction

Table 8.1 summarizes the Suzaku observation log of W49B. We used data from the
XIS excluding XIS2 because of the malfunction described in §4.3.1. We reduced the dataand estimated NXBs with the same method in Chapter 6. As already explained in §4.3.3,
we corrected the gains of each sensor.
Table 8.1: Suzaku observation log of W49B.
Target
W49B
W49B

Obs. ID
503084010
504035010

Obs. date2009-03-29
2009-03-31(R.A., Dec.)(19h 11m 08s .3, 9◦ 06′ 56′′ .5)(19h 11m 08s .3, 9◦ 06′ 55′′ .8)

Exposure
52 ks
62 ks

8.3. ANALYSIS

85

8.3

Analysis

8.3.1

Image

Figure 8.4 shows an XIS3 image of W49B in the energy band of 0.3–10.0 keV. We
subtracted the NXB from the image and corrected for the vignetting eﬀect of the XRT.
The image shows a center-filled morphology, where the peak of the X-ray emission is oﬀset
toward the east.
As shown in Figure 8.5, we present XIS3 images of W49B in the soft band (0.3–5.0 keV)and the hard band (5.0–10.0 keV), which cover emission lines of Ne to Ca and Cr to Ni,
respectively. We found that the hard band image is more compact than the soft bandone. If the images reflect the distribution of the elements, the result suggests that the
distribution of the heavier element is center-filled whereas that of the lighter elements ismore extended. The distributions are consistent with the onion-like structure (§2.1) thatheavier elements tend to be produced closer to the center of the progenitor.

Right ascension
11:30.0

10:00.0

11:00.0

05:00.0

8:55:00.0

0.0

10:30.0

Whole

East

9:00:00.0

Declination

15:00.0

12:00.0West

Background

0.2

0.4

0.6

0.8

1.0

Figure 8.4: XIS3 image of W49B in the 0.3–10 keV band after NXB subtraction and correctionof the vignetting eﬀect. The yellow square indicates the FOV of the XIS. The magenta and cyandotted lines show the source and background regions for the spectral analysis, respectively. The
X-ray count rates are normalized so that the peaks become unity. The green contours are theradio intensity obtained at a wavelength of 20 cm with the MAGPIS. The coordinate refers to
the J2000.0 epoch.

86

CHAPTER 8. W49B

0.3–10.0 keV

0.0

0.2

0.4

0.6

0.3–5.0 keV

0.8

1.0

0.0

0.2

0.4

0.6

5.0–10.0 keV

0.8

1.0

0.0

0.2

0.4

0.6

0.8

1.0

Figure 8.5: Same as Figure 8.4 but in the energy bands of 0.3–5.0 keV (Left) and 5.0–10.0 keV(Right).

8.3.2

SNR spectra

For the spectral analysis, we extracted source and background spectra from the Wholeand Background regions shown in Figure 8.4. We subtracted the background spectrum
from the source spectrum. Figure 8.6 shows the XIS0+3 spectrum. One can clearly see
the emission lines of Mg, Si, S, Ar, Ca, Cr, Mn, Fe and Ni. The strong cascade lines and
RRC of Fe are also detected, suggesting that the plasma is in a recombination-dominantstate.
We first applied an RP model for the spectrum by using the VVRNEI model multipliedby the absorption model, TBabs. The abundances of Mg, Si, S, Ar, Ca, Cr, Mn, Fe, and Ni
were allowed to vary, whereas the others were fixed to solar. We also allowed kTe , kTinit ,ne t and the normalization to vary. For the XIS0+3 spectrum, we ignored the energy
band around Si K-edge because of the known calibration uncertainty. We used the XIS1
spectrum in the energy band below 6.8 keV because of the higher NXB level. Figure 8.7a
and Table 8.2 show the fitting results. The fit left residuals around the low-energy bandof ≲ 2 keV, indicating that the ionization degree of Si is significantly lower than the otherelements. Therefore, we tried a CIE component for the low-energy band as we did forthe spectral analyses of IC 443 and W44 in Chapters 6 and 7. For the CIE component,
the normalization and kTe were free parameters. We first fixed the abundances of all theelements to solar, but residuals around the Mg Kα line appeared. Thus, we allowed the
Mg abundance to vary. The best-fit model and parameters are presented in Figure 8.7b
and Table 8.2, respectively. The additional CIE component improved the residuals dueto the discrepancy of the ionization degrees between Si and the other elements. However,residuals in the hard band above 6 keV are still significant, which implies even higher
ionization degree for the plasma emitting the Fe lines and RRC.
We, therefore, divided the RP into two components (RPL and RPH ). For the RPL
model, we allowed the abundances of Mg, Si, S, Ar, and Ca to vary. The abundances of
the heavier elements (≥ Cr) were fixed at zero. The other abundances were fixed to solar.
We also allowed kTe , kTinit , ne t and the normalization to vary. For the RPH model, theabundances of Cr, Mn, Fe, and Ni were allowed to vary. The abundances of the lighter
elements were fixed at zero, and those of the other abundances were fixed to solar. The
normalization of RPH was linked to that of RPL , whereas kTe , kTinit and ne t were allowed8.3. ANALYSIS

87

10

to vary. We fitted the Whole spectrum with the model consisting of the CIE and two RP
components, and found that the spectrum can be reproduced by the model as shown in
Figure 8.7c and Table 8.2.
Considering of the shell-cloud interaction suggested by the infrared 2.12 µm line emission (Figure 8.1), we extracted spectra from the East and West regions shown in Figure 8.4
in order to examine possible influences of the interaction on the characteristics of the Xray emitting plasma. We adopted the same model as the Whole fit with kTinit of the RPL
component fixed at the best-fit value of the Whole fit. The model well fitted the spectra(Figures 8.7d and 8.7e; Table 8.2).

Si

S

Ar

Fe Kα

Ca

Ni +
Fe Kβ

0.1

Cr

Mn
0.01

Counts s −1 keV−1

1Mg

Fe cascade

10−3

Fe RRC
1

2

Energy (keV)

5

10

Figure 8.6: XIS0+3 spectrum of the Whole region after subtracting the background and NXB
spectra.

88

CHAPTER 8. W49B(d) 10

Whole

Counts s −1 keV−1

1
0.1

10−3
5

10−3
5

0

0
−5

1

2

Energy (keV)5(e) 10Counts s −1 keV−1

Whole

1

1

10

0.1

2

Energy (keV)

5

10

5

10

West

1
0.1

0.01

0.01

10−3
5

10−3
5

0

0

χ

Counts s −1 keV−1(b) 10

χ

0.1
0.01

−5

−5

−5
1

(c) 10Counts s −1 keV−1

East

1

0.01

χ

χ

Counts s −1 keV−1(a) 102

Energy (keV)

5

10

5

10

1

2

Energy (keV)

Whole

1
0.1
0.01

χ

10−3
5
0
−5
1

2

Energy (keV)

Figure 8.7: XIS0+3 (black crosses) and XIS1 (red crosses) spectra extracted from the Whole(a–c), East (d) and West (e) regions. The magenta, blue, cyan and purple lines indicate themodels of CIE, RP (all elements), RP (light elements) and RP (heavy elements), respectively.

ParameterNH (1022 cm−2 )kTe (keV)
ZMg (solar)
Zother (solar)
Normalization‡kTe (keV)
kTinit (keV)
ZMg (solar)
ZSi (solar)
ZS (solar)
ZAr (solar)
ZCa (solar)
ne t (1011 s cm−3 )Normalization†kTe (keV)
kTinit (keV)
ZCr (solar)
ZMn (solar)
ZFe (solar)
ZNi (solar)
ne t (1011 s cm−3 )Normalization†
χ2ν (ν)= ne t VVRNEI2= Norm. VVRNEI2

1.96 (2354)= ne t VVRNEI2

2.63 (2357)

40 ± 735 ± 723 ± 485 ± 18

20 ± 113 ± 27.7 ± 0.3
49 ± 3= Norm. VVRNEI2= kTe, VVRNEI2= kTinit, VVRNEI2= kTe, VVRNEI2

9.68 ± 0.07
0.174 ± 0.001
0.20 ± 0.07
1 (fixed)
1.1 ± 0.1
1.00 ± 0.01
6.2 ± 0.1
16 ± 421 ± 419 ± 320 ± 324 ± 45.3 ± 0.2
0.4 ± 0.1

= kTinit, VVRNEI2

7.34 ± 0.05
————
————
————
————
1.36 ± 0.02
8.3 ± 0.4
3.6 ± 0.4
5.4 ± 0.2
5.2 ± 0.2
6.4 ± 0.2
9.6 ± 0.3
7.7 ± 0.2
0.6 ± 0.1

(RP)

Whole

= Norm. VVRNEI2

9.22 +0.10−0.12
0.170 +0.003−0.0040.17 +0.09−0.08
1 (fixed)
1.1 +0.3−0.1
0.94 ± 0.02
5.9 +1.4−0.9
6.1 ± 1.4
7.0 +0.8−0.7
6.8 +0.7−0.5
7.5 +0.7−0.6
9.1 +0.6−0.8
4.4 ± 0.1
1.1 ± 0.1
1.46 +0.01−0.08
3.26 +0.01−0.10
7.7 ± 1.2
7.6 ± 1.4
3.0 ± 0.3
5.3 +1.2
−1.0
< 0.9

(CIE+2RP)

East= Norm. VVRNEI2

9.23 +0.20−0.15
0.175 +0.002−0.0070.17 ± 0.12
1 (fixed)
0.9 +0.2−0.1
1.05 +0.04−0.07
5.9 (fixed)
8.0 +3.1−2.4
7.7 +1.3−1.0
7.7 +1.1−0.8
8.3 +1.3−1.0
10.7 +1.3−1.1
5.6 +0.5−0.3
0.8 ± 0.1
1.66 +0.07−0.11
3.15 +0.02−0.12
10.7 +2.5−2.0
11.4 +3.6−1.6
4.2 +0.8−0.4
7.1 +1.8
−1.2
< 0.2

(CIE+2RP)‡

1.59 (2351)
1.33 (2028)∫
22
The emission measure integrated over the line of sight, i.e., (1 / 4π) ne nH dl in units of 10 cm−5 sr−1 .
∫†The emission measure integrated over the line of sight, i.e., (1 / 4π) ne nH dl in units of 1019 cm−5 sr−1 .

VVRNEI3(RPH comp.)

VVRNEI2(RPL comp.)

Model function
TBabs
VVRNEI1(CIE comp.)

Whole(CIE+RP)

Whole

Table 8.2: Best-fit model parameters of the W49B spectrum.
West

1.35 (2043)= Norm. VVRNEI2

9.25 +0.07−0.16
0.161 +0.007−0.0050.16 +0.10−0.09
1 (fixed)
1.7 ± 0.5
0.81 +0.02−0.01
5.9 (fixed)
4.6 ± 1.5
6.6 +1.0−0.9
6.3 ± 0.8
6.8 +0.9−0.8
8.1 +1.1−1.0
5.1 ± 0.1
1.4 ± 0.2
1.26 +0.02−0.09
3.37 ± 0.11
5.2 ± 1.1
4.6 +1.0−1.1
1.9 +0.3−0.2
3.6 +1.1
−1.0
< 0.4

(CIE+2RP)

8.3. ANALYSIS
89

90

8.4

CHAPTER 8. W49B

Discussion on W49B

We fitted the wide band spectra (1–10 keV band) and found the ejecta consisting of
the two components; the cold RPL component with kTe ∼ 0.9 keV and the hot RPH one
with kTe ∼ 1.5 keV. Ozawa et al. (2009b), who analyzed the same Suzaku data, did notreach this conclusion since they performed a spectral analysis only in the limited energy
band of 5–12 keV. With Chandra imaging analyses, the Fe band image is more compact
than those of the Si, S, Ar and Ca bands (Lopez et al., 2013a). The fact that the Suzaku
hard-band image is more compact than the soft band image (Figure 8.5) is consistent with
the Chandra result. The RPH component would be concentrated in the central region of
the remnant compared to RPL . If so, the plasma close to the rim would have lower kTe
than the plasma in the central region. This can be interpreted as the result of thermal
conduction between the SNR shock and the surrounding dense molecular gas.
In the spectral analysis of the East and West regions, we found that kTe of the western
plasma is lower than the eastern plasma, which is consistent with the results of thespatially resolved spectral analysis with Chandra by Lopez et al. (2013b). Lopez et al.
(2013b) claimed that the rarefaction scenario is likely to explain it. They claimed thatthe shock is interacting with a cloud only in the east, and that the ambient gas density
outside the western shell should be lower. They then concluded that the dominant coolingprocess is adiabatic expansion since the lower density plasma has lower temperature. This
interpretation, however, seems questionable to us. Since the 2.12 µm line was detectedboth on the eastern and western sides (Figures 8.1), the western shell should also beinteracting with a cloud. Gamma-ray data provide another clue of the interaction. Abdo
et al. (2010) and H. E. S. S. Collaboration et al. (2016) detected a GeV gamma-ray emission
on the western side of the remnant. According to H. E. S. S. Collaboration et al. (2016),
the gamma rays are most probably interpreted as π 0 -decay emission. Therefore, thegamma-ray results indicate that dense gas is present also on the eastern side otherwise
the bright gamma-ray emission would not be detected there.

8.4. DISCUSSION ON W49B

91

Chapter 9
Discussion
9.1

Summary of the Results

We analyzed the Suzaku data of four Galactic SNRs: G166.0+4.3, IC 443, W44, and
W49B. We obtained NH , kTe , kTinit and ne t of each component of the thermal plasma
in these remnants. Table 9.1 presents a summary of the parameters in the NE regionof G166.0+4.3 and the SE region of IC 443 which correspond to the locations where the
SNR shock is interacting with the molecular cloud, and the Whole regions of W44 andW49B. In the following, we determine electron densities, recombination timescales, andprogenitor masses which are key parameters to discuss the physical origin of the RPs.

Table 9.1: Summary of the fitting results.
SNR

RegionG166.0+4.3
IC 443
W44
W49B

NE
SE
Whole
Whole

ISM
kTe

kTe

Ejecta1
kTinit

(keV)

(keV)

(keV)

–
0.22 ± 0.01
0.22 ± 0.01
0.170 +0.003−0.0040.46 ± 0.03
0.54 +0.03−0.01
0.37 ± 0.01
0.94 ± 0.02

3 (fixed)
5 (fixed)
0.93 ± 0.02
5.9 +1.4−0.9

ne t(1011 s cm−3 )6.1 +0.5−0.4
4.2 +0.2−0.1
3.5 +0.2−0.3

5.3 ± 0.2

kTe
(keV)
0.87 +0.02−0.03

0.19 ± 0.01–
1.46 +0.10−0.08

Ejecta2
kTinit(keV)–
5 (fixed)–
3.26 +0.01
−0.10

ne t
(1011 s cm−3 )–= ne t (Ejecta1)–< 0.9

Ejecta1 and Ejecta2 indicate RP(H–Ca) and CIE(Fe&Ni) for G166.0+4.3, RPhot and RPcold for IC 443, and RPL and RPH
for W49B, respectively. For W44, Ejecta1 indicates the RP component.

9.1.1

Electron densities and recombination timescales

We estimate electron densities (ne ) and recombination timescales (trec ) of the RP components in each remnant. In §5.4.2, we estimated ne of the RP component in G166.0+4.3
from the volume emission measure. On the other hand, spectral fits directly give ne t ofeach remnant. By dividing the latter parameter by the former, we can obtain trec , which
we refer to as recombination timescale. In the same manner, we obtained ne and trec ofthe plasmas in IC 443, W44, and W49B as listed in Table 9.2. In all the remnants, trec isshorter than tage , which is consistent with a picture that the plasmas become to be in arecombination-dominant state at some point of the SNR evolution most probably either
by rarefaction or thermal conduction. We will discuss each possibility in §9.2 and §9.3.
92

9.1. SUMMARY OF THE RESULTS

93

Table 9.2: Summary of ne and trec .
Ejecta1
SNR
G166.0+4.3
IC 443
W44
W49B

Region
NE
SE
Whole
Whole

Ejecta2

ne

trec

ne

(cm−3 )
0.3f −1/2(kyr)
71f 1/2(cm−3 )
0.3f −1/2

0.9f −1/2
0.7f −1/2
2.7f −1/2

15f 1/2
16f 1/2
6.0f 1/2

1.6f −1/2–
2.7f −1/2

trec
(kyr)
–
8.3f 1/2–< 0.1f 1/2

Distance

Radius

(kpc)
5†

(pc)

1.5†
3†
8†

20
10
10
5

tage
(kyr)
140f 1/2
30‡
20‡
6‡

A filling factor is denoted by f .

† Estimated distances to G166.0+4.3 (White & Long, 1991), IC 443 (Welsh & Sallmen, 2003), W44 (Claussen et al., 1997),

and W49B (Moﬀett & Reynolds, 1994).
‡ Estimated ages of IC 443 (Olbert et al., 2001), W44 (Wolszczan et al., 1991), and W49B (Smith et al., 1985).

9.1.2

Abundance patterns and progenitor masses

X-ray spectral analyses of SNRs can be used to identify their SN types based onabundance patterns of their ejecta components. In our spectral analyses explained above,
we obtained the abundances of each element in each remnant (G166.0+4.3: Table 5.3;
IC 443: Tabale 6.2; W44: Table 7.3; W49B: Table 8.2). We here discuss SN types of
G166.0+4.3, IC 443, W44, and W49B by comparing the obtained abundances with thoseof theoretical predictions for core-collapse (Woosley & Weaver, 1995) and Type Ia (Maeda
et al., 2010) SNe.
Figure 9.1 shows abundance patterns normalized by Si in the NE and W regions of
G166.0+4.3, the NE, NW, and SE regions of IC 443, and the Whole regions of W44 andW49B. From our results of IC 443, W44, and W49B, we interpreted the CIE components
as the shocked ISM, whereas the RPs are attributed to the reverse-shock-heated ejecta.
We therefore compare the abundance patterns of the RP components to those of thetheoretical models. On the other hand, in the case of G166.0+4.3, we cannot distinguish
the ISM component from the ejecta component, and, therefore, the obtained abundances
may be a mixture of the ISM and ejecta. Even if this is the case, the contamination of
the ISM decreases the ratios of O/Si, Ne/Si, Mg/Si, Fe/Si, and Ni/Si in the NE region
but it does not aﬀect the determination of the SN types.
G166.0+4.3 The abundance pattern of the NE region roughly agrees with that of acore-collapse SNe model with a progenitor mass of 20–30 M⊙ . On the other hand, theabundance pattern in the W region favors a core-collapse SN with 11 M⊙ . In either case,these results indicate a core-collapse origin rather than that of Type Ia.
IC 443 The abundances pattern of IC 443 in all the regions are similar to each other.
They are consistent with those of core-collapse SNe with a progenitor mass of 20–30 M⊙ .
The core-collapse origin is supported by the presence of the associated pulsar, CXOU J061705.3+222127 (§A.1; Bocchino & Bykov, 2003; Swartz et al., 2015).
W44 The abundance pattern of W44 indicates a core-collapse SN origin since the Fegroup elements (Fe and Ni) are less abundant than the intermediate mass elements (O, Ne,
Mg, Si), which is typical for a nucleosynthesis model of core-collapse SNe. The estimatedabundances with O/Si, Ne/Si, and Mg/Si ∼ 1, and Fe/Si and Ni/Si < 1 exclude themodels with ≤ 11 M⊙ and ≥ 40 M⊙ . We therefore conclude that the progenitor mass islikely in the range of 20–30 M⊙ . The associated pulser PSR B1853+01 is detected in the

94

CHAPTER 9. DISCUSSION

central region of W44 (Chapter 7; Petre et al., 2002), which is supportive evidence forour conclusion.
W49B The Mg/Si ratio of ∼ 1 strongly favors the core-collapse origin. However, theType Ia origin cannot be ruled out by the Fe/Si and Ni/Si ratios of ∼ 1. The progenitor
type of W49B is still under debate: Lopez et al. (2013a) found an asymmetric distributionof emissions from FeXXV lines with the Chandra imaging analysis. They claimed that the
distribution suggests a Type Ib/c bipolar explosion. On the other hand, based on deep
Chandra observations of W49B, Zhou & Vink (2017) claimed the abundance pattern is
consistent with a Type Ia SN of a Chandrasekhar-mass white dwarf. They also reported
candidate point sources inside W49B, which may conflict with a picture of the Type Ia
SN if one of them is a neutron star.
We conclude that the origins of IC 443 and W44 are core-collapse SNe with progenitor
masses of 20–30 M⊙ . G166.0+4.3 and W49B are also likely core-collapse SNe but theirprogenitor masses have large uncertainties.

9.1.3

Spatial variation of kTe

In G166.0+4.3, we discovered an RP only in the outer part of the remnant in the dense
NE region whereas the plasma in the lower-density W region is an IP. The temperatureof the outer RP is lower than that of the inner Fe-rich plasma in the NE region andthat of the IP in the W region. In IC 443, kTe of both the RP components decrease
toward the southeast where the remnant is interacting with the molecular cloud. In
W44, kTe of the RP in the center of the remnant is higher than that of the outer RP,
suggesting that the SNR plasma is cooled sequentially from the rim to the center by thecool surrounding cloud. In W49B, the central plasma is hotter than the plasma in therim, which is similarly seen in G166.0+4.3 and W44. In the spectral analysis of W49B,
we also found that kTe of the western plasma is lower than that of the eastern plasma.
Since the gamma-ray emission is detected in the western side of the remnant, the westernshell should be interacting with a cloud.
All the results above point to the same tendency that the electron temperature kTe
is lower at the locations with shell-cloud interactions. This is qualitatively consistent
with the picture that the RPs are generated by rapid cooling of the plasmas throughthermal conduction with the molecular clouds. More quantitatively, however, the coolingmust proceed fast enough compared to the recombination of ions to realize RPs. We thus
estimate the cooling timescale in the following section (§9.3).

9.1. SUMMARY OF THE RESULTSO Ne Mg Si

S Ar CaFe Ni

O Ne Mg Si

G166.0+4.3 NE

10
1
0.1

0.01
10

15
20
Atomic Number
Element

S Ar Ca

0.1

25

30

5

10

15
20
Atomic Number

O Ne Mg Si

Fe Ni

Element

S Ar Ca

1
0.1

25

30

O Ne Mg Si

Fe Ni

IC 443 NE

10

Z X / Z Si

Z X / Z Si

10

1
0.1

10

15
20
Atomic Number

O Ne Mg Si

Element

S Ar Ca

25

5

30

Fe Ni

10

10

Z X / Z Si

0.1

0.01

Fe Ni

1
0.1

15
20
Atomic Number
Element

S Ar Ca

25

30

5

10

15
20
Atomic Number

25

Fe Ni

W49B Whole

10

1

S Ar Ca

0.01

O Ne Mg Si

W44 WholeElement

IC 443 NW

10

0.01

0.01

Z X / Z Si

1

IC 443 SE

5

Fe Ni

G166.0+4.3 W

0.01

O Ne Mg Si

5

S Ar CaZ X / Z Si

5Element

10

Z X / Z Si

Z X / Z Si

Element

95

1
0.1

0.01
10

15
20
Atomic Number

25

30

5

10

15
20
Atomic Number

25

30

Figure 9.1: Abundance ratios of each element to Si in the NE and W regions of G166.0+4.3, theNE, NW, and SE regions of IC 443, and the Whole regions of W44 and W49B. The red points
indicate several nucleosynthesis models for core-collapse SNe with diﬀerent progenitor massesof 11 M⊙ (red circle), 20 M⊙ (red square), 30 M⊙ (red traiangle), and 40 M⊙ (red star) from
Woosley & Weaver (1995). For Type Ia SNe, the blue points show W7 (blue circle) model from
Nomoto et al. (1984), C-DEF (blue square), and C-DDT (blue triangle) models from Maeda etal. (2010).

30

96

9.1.4

CHAPTER 9. DISCUSSION

Spatial variation of ne t

The spectral analyses of IC 443, W44, and W49B revealed the spatial variation ofne t of the RPs, which can be translated into the spatial variation of the recombination
timescales (trec ) once the electron density (ne ) is determined. For simplicity, we assume
that ne of the RP is uniform at the value listed in Table 9.2.
We note that, in G166.0+4.3, an RP is detected only in the NE region so that wecannot obtain the ne t distribution of this remnant. In the case of W44, ne t of the southernpart of the remnant is higher than those of the other regions. It suggests that the shockcloud interaction started in the early phase of the SNR evolution. We obtained similar
results for W49B as well. We found that ne t of the inner plasma is lower than that of
the outer part of the remnant (inner plasma: ne t < 0.9 × 1011 s cm−3 ; outer plasma:
ne t = 4.4 × 1011 s cm−3 ), indicating that the inner plasma started to be cooled later
than the outer plasma. These results on W44 and W49B are explained by the thermal
conduction scenario, in which molecular clouds first cool the plasmas in the rim regions,and then the inner plasmas via the outer plasmas.
In the case of IC 443, interestingly, a diﬀerent tendency was observed. The RP in
the southeast, where the shock-cloud interaction occurs, has a lower ionization timescale
ne t = 4.2 × 1011 s cm−3 than those of the northwestern parts of the remnant ne t ∼1012 s cm−3 . Nevertheless, several possibilities can be considered in the thermal conduction
scenario. One possibility would be that the dense molecular in the southeastern region
recently collided with the shell and cooled the plasma rapidly, and that the plasma in thenorthwestern region has been cooled by other lower density gas for a longer time since anearlier phase. The other possibility is that IC 443 has the same ne t variation only in theline-of-sight direction, which cannot be measured in our spectral analysis.

9.2. RAREFACTION

9.2

97

Rarefaction

As discussed in §9.3, the rarefaction scenario is unlikely to explain the kTe spatial
variation we observed. Nevertheless, we discuss the formation process of RPs in thecontext of the rarefaction scenario by using the progenitor masses given in §9.1.2. Itoh
& Masai (1989) proposed the rarefaction scenario, in which RPs are realized by adiabatic
cooling when the SNR shock breaks out of a dense CSM into a lower density ISM. Shimizu
et al. (2012) carried out 3D-hydrodynamical calculations of the interaction between the
expanding supernova ejecta and the CSM. They concluded that the CSM radius needs
to be a sub-pc scale (≲ 1 pc) in order for the adiabatic expansion significantly cools theplasma. Chen et al. (2013) found a linear relationship between the radius of a massive
star’s main-sequence bubble (Rb ) in a molecular environment and the star’s initial mass:
Rb ≈ 1.22M/M⊙ − 9.16 pc (Figure 9.2). In this case, the upper limit to the progenitor
mass is ∼ 8.3 M⊙ to realize the rarefaction if we accept the requirement of the CSM size
< 1 pc given by Shimizu et al. (2012) as mentioned above. As described in §9.1.2, theestimated progenitor masses of IC 443 and W44 are 20–30 M⊙ , which are much larger
than the upper limit; we therefore rule out the possibility of the rarefaction scenario.
On the other hand, since we cannot determine the progenitor masses of G166.0+4.3 andW49B, the rarefaction scenario is still possible for these remnants in terms of the CSM
size.
Our result is consistent with the conclusion by Moriya (2012). They studied the CSM
around several supernova progenitors and discussed which progenitors can have a CSM
dense enough to establish CIE soon after the explosion. They suggested that the CSM
around supergiants can establish CIE and can evolve to become recombining SNRs, but
massive progenitors of Wolf-Rayet stars are unlikely.

Figure 9.2: Relationship between radii of the wind-blown bubbles and the stellar masses for OB
stars (Chen et al., 2013). The plus and circle signs indicate the B and O stars, respectively. The
1/3
solid line represents a linear regression, p5 Rb = 1.22M/M⊙ − 9.16 pc, where Rb and p5 are themaximum size of the bubbles and normalized pressure p5 ≡ (p/k)(105 cm−3 K), respectively.

98

CHAPTER 9. DISCUSSION

9.3

Thermal Conduction

If the scale length of the temperature gradient lT ≡ Te /|∇Te | is much longer thanthe mean free path of electrons λe ≈ 1018 (Te /107 K)2 (ne /1 cm−3 )−1 cm, the heat flux is
calculated by the classical formula,
Qclass = −κ∇Te .

(9.1)

Spitzer (1956) gave the thermal conductivity as(
κ = 1.31ne λe k

kTe
me

)1/2

(
≈ 4.6 × 10

13

Te
108 K

)5/2 (

ln Λ
40

)−1

ergs−1 cm−1 K−1 ,(9.2)

where the plasma parameter is described as
4
Λ= π
3

(

ϵ0 kTene e 2

)3/2
ne .

(9.3)

If lT is comparable to or less than λe , the electrons diﬀuse at a speed greater than theiraverage thermal speed, and the conduction is in a “saturation” regime. The heat flux
approaches a limited value given by Cowie & McKee (1977),(
Qsat = 0.4ne kTe

2kTe
πme

)1/2
.

(9.4)

We divide the SNR plasma into 10 layers with the same thickness (d), and calculate heat flows between each layer at each time step (see a schematic picture shown in
Figure 9.3). The heat flow into and out of the i-th layer (∆Econd ) can be calculated as∆Econd ≃ (Qi,i−1 Si,i−1 − Qi+1,i Si+1,i )∆t2∆Econdk∆Te =,
3ne Vi

(9.5)
(9.6)

where Qi,j , Si,j , and Vi are the heat flux and the contact area between the layers i and j,and the volume of the layer i, respectively. We adopt the heat flux of(lT /λe > 10)
 Qclass−1 −1)(1 ≤ lT /λe ≤ 10)
+
Q
(Q−1
(9.7)
Qi,j =
sat
class

Qsat
(lT /λe < 1).
In the intermediate regime of 1 ≤ lT /λe ≤ 10, we smoothly connect them between theclassical and saturated conduction regime by following Dalton & Balbus (1993).
Table 9.3 and Figure 9.4 show the initial parameters and the results in each remnant.
We assumed kTinit of the RP component determined by the spectral fitting as the initial
temperature of the plasma. We assumed that the temperature of the outer cool gas isconstant at zero. The thickness of the layers are one tenth of the SNR radius. As theinitial electron density, we adopted ne obtained by the spectral analyses when the filling
factors were assumed to be 1 (see Table 9.2). We also assumed that heat flows only
between adjacent layers. For simplicity, we assumed that a volume of the plasma and9.3. THERMAL CONDUCTION99

ne are constant at their initial values. This assumption is justified as explained below.
According to McKee & Cowie (1975), a ratio of the expansion velocity propagating into
the molecular cloud to that before the shock-cloud interaction is described as,
(
)−1/2
nMC
vMC
,
(9.8)
=k
vISM
nISM
where k is a factor which depends weakly on the density ratio of ISM to a molecular cloud(nMC /nISM ), ranging k = 1–1.5. If we adopt the densities of ISM and a molecular cloudas ∼ 1 cm−3 and 102 –103 cm−3 , respectively, vMC /vISM would be ≲ 0.15.

layer9layer8outer coldgas
layer1

Qout,9 Q9,8 Q8,7
kTe̲ouerQ2,1 Q1,0 layer0
kTe̲1

kTe̲9 kTe̲8d= 0 keV(ﬁxed)

d

d

kTe̲0

inner hotplasmad

Figure 9.3: Schema of the plasma structure in our calculation of the time development of kTe
with the thermal conduction between the inner hot plasma and the outer cold gas.

Table 9.3: Initial parameters for the calculation of the time evolution of kTe .
SNR
G166.0+4.3
IC 443
W44
W49B

kTouter
(keV)
0
0
0
0

kTinner
(keV)
3.0
5.0
0.93
5.9

d
(pc)
2
1
1
0.5

ne
(cm−3 )
0.3
0.9
0.7
2.7

From Figure 9.4, we confirm that the cooling timescale of kTe of the inner plasma ismuch longer than that of the outer plasma. The results reproduce the observed distributions of kTe in G166.0+4.3 and W49B. Here we define a conduction timescale tcond as atime when the calculations reach the observed temperatures kTe of the RP components in
the innermost layer (see Table 9.4). From Figure 9.4, we estimated tcond of G166.0+4.3,
IC 443, W44 and W49B at ∼ 40 kyr, ∼ 20 kyr, ∼ 35 kyr and ∼ 4 kyr, respectively.

100

CHAPTER 9. DISCUSSION

Kawasaki et al. (2002) estimated a conduction timescale of the plasma in IC 443 witha diﬀerent assumption from us. They found that the plasma is composed of the outer
cold CIE with kTe ∼ 0.2 keV and the inner hot RP with kTe ∼ 1.0 keV with ASCA, andcalculated the timescale to equilibrate the temperatures of the two components. They
compared the conduction timescale of ∼ 6 kyr with tCIE ∼ 63 (ne /1 cm−3 )−1 kyr which isa characteristic timescale for a plasma to reach CIE (Masai, 1994). They concluded thatthe thermal conduction can realize RPs because the conduction timescale is shorter thantCIE .
To investigate whether the RPs can be realized with the conduction timescales estimated by us under the diﬀerent assumption from Kawasaki et al. (2002), we compare tcond
with tCIE . Table 9.4 lists tcond and tCIE of each remnant. We confirmed that the conduction timescales are shorter than tCIE for all the remnants, indicating that the plasmas canbecome RPs with the thermal conduction between the plasma and the molecular cloud.
The result contradicts with conclusions by other authors that the thermal conduction
timescale is too long to realize RPs in some SNRs. Uchida et al. (2012) analyzed theSuzaku data of W44 and interpreted that the temperatures of the inner and outer plasmas are 0.48 keV and 0.40 keV, respectively. They estimated the conduction timescaleof ∼ 10000 kyr with the same method as presented by Kawasaki et al. (2002). They
claimed that the thermal conduction scenario is unlikely because of the longer conduction
timescale than tCIE . We reached a diﬀerent conclusion from theirs because of the diﬀerent
definitions of the timescale as mentioned above.
If the thermal conduction is a cause of the RPs, it is also required that tcond should
be shorter than the SNR age, tage . Table 9.4 lists tcond and tage of each remnant. For
G166.0+4.3, we estimated the age at ∼ 130 (ne /0.3 cm−3 )−1 kyr (§5.4.2). We adopted
30 kyr, 20 kyr and 6 kyr given by Olbert et al. (2001), Wolszczan et al. (1991), and Smithet al. (1985) as the ages of IC 443, W44 and W49B, respectively. The requirement of
tcond ≤ tage is satisfied for G166.0+4.3, IC 443, and W49B, whereas tcond of W44 is longer
than its age. The calculations of tcond , however, highly depend on ne which has large
uncertainties with the assumed volumes of the plasma and the distance to the remnant.
We then calculated tcond as a function of electron density ne as plotted in Figure 9.5.
Under the condition of tcond < tage , the electron densities are limited to ne ≲ 0.55 cm−3
for G166.0+4.3, ne ≲ 1.4 cm−3 for IC 443, ne ≲ 0.40 cm−3 for W44, and ne ≲ 4.5 cm−3
for W49B. If the electron densities of the SNR plasmas satisfy these requirements, theRPs can be explained by the thermal conduction scenario.
Table 9.4: List of timescales in each remnants.
SNR
ne
kTe
tcond tCIE
tage
(cm−3 ) (keV) (kyr) (kyr) (kyr)G166.0+4.3
0.3
0.46
40
200
130
IC 443
0.9
0.54
20
70
30
W44
0.7
0.37
35
90
20
W49B
2.7
0.94
4
20
6

9.3. THERMAL CONDUCTION

G166.0+4.3

3n e = 0.3 cm-3
kT init, RP = 3.0 keV
kT e, RP = 0.46 keV
0 kyr

2

1 kyr
2 kyr
5 kyr
10 kyr
20 kyr1

kTe (keV)

kTe (keV)

101

5

IC 4434n e = 0.9 cm-3
kT init, RP = 5.0 keV
kT e, RP = 0.54 keV
0 kyr

3

1 kyr
2 kyr

2

5 kyr
10 kyr
20 kyr1

50 kyr

50 kyr100 kyr

100 kyr
0

2

4

6

W44

1.0

n e = 0.7 cm
kT init, RP = 0.94 keV
kT e, RP = 0.37 keV

0

1

2

3

4

5 6 7 8R (pc)

9 10

6

W49B

5

n e = 2.7 cm-3
kT init, RP = 5.9 keV
kT e, RP = 0.94 keV-30.8
kTe (keV)0

8 10 12 14 16 18 20
R (pc)

0 kyr
1 kyr

0.6

2 kyr
5 kyr
10 kyr
20 kyr

0.4
0.2

50 kyr4kTe (keV)

00 kyr
0.1 kyr3

0.2 kyr2

0.5 kyr
1 kyr
2 kyr
5 kyr

1

10 kyr100 kyr
0

0

1

2

3

4

5 6 7 8R (pc)

9 10

0

0

1

2

3
R (pc)

4

5

100 kyr

Figure 9.4: Time evolution of the electron temperatures of the RPs in G166.0+4.3 (top left),
IC 443 (top right), W44 (bottom left), and W49B (bottom right) as a function of distances fromthe center of the remnant. The dotted lines indicate kTe of the RPs obtained by the spectral
analyses.

102

CHAPTER 9. DISCUSSION

60

G166.0+4.3

t cond

50

100

t CIE

Time (kyr)Time (kyr)IC 443t cond

t age

50

40

t age

30

t CIE

20
10

0

ne,VEM

ne,VEM
0

0.2

0.4
0.6
0.8
Electron density (cm-3)

0

1

100

1
1.5
2
2.5
Electron density (cm-3)

3

W49B

t cond

80

t CIE
10

60

Time (kyr)Time (kyr)

0.5

15

W44

t CIE

40

t cond
t age
5

t age

20
0

0

ne,VEM
00.5
1
1.5
Electron density (cm-3)

20ne,VEM
02

4
6-3
Electron density (cm )8Figure 9.5: Conduction timescales of G166.0+4.3 (top left), IC 443 (top right), W44 (bottom
left), and W49B (bottom right) as a function of the electron densities denoted by the black solidlines. The red solid lines indicate the characteristic timescale of 63 (ne /1 cm−3 )−1 kyr for aplasma to reach CIE (Masai, 1994). The blue and green solid lines show the age of G166.0+4.3,
IC 443, W44, and W49B estimated by us (§5.4.2), Olbert et al. (2001), Wolszczan et al. (1991),
and Smith et al. (1985), respectively. The electron density obtained from a volume emission
measure (ne,VEM ) is denoted by the black dotted lines.

9.3. THERMAL CONDUCTION

103

Chapter 10
Conclusions
We studied the RPs in the four Galactic SNRs, G166.0+4.3, IC 443, W44, and W49B
with Suzaku. The results are summarized as follows:
1. The spectrum of the W region in G166.0+4.3 is well fitted with a one-component
IP model, whereas the plasma in the dense NE region consists of the Fe-rich plasma
and RP components. The Fe-rich plasma is concentrated near the center of theremnant whereas the other elements are distributed in the outer region. The electron
temperature of the RP is lower than those of the inner Fe-rich plasma and the IP
in the W region. The results of the spectral analysis are simply explained by thethermal conduction with the cool dense gas in the northeastern part of the remnant.
2. We performed spatially resolved spectroscopy of X-ray emissions from IC 443. The
spectra extracted from each region of the remnant are all fitted well with a model
consisting of CIE and two RP components. The electron temperature of one of the
RP component ranges from 0.16 keV to 0.28 keV whereas that of the other RP
component is in the range of 0.48–0.67 keV. The electron temperatures of both RPcomponents decrease toward the southeast, where the remnant is interacting withthe molecular cloud. Also, the normalization ratio of the lower-kTe RP to higher-kTe
RP components increases toward the southeast. The two findings may be a result ofcooling of the X-ray emitting plasma by the interacting molecular cloud via thermal
conduction.
3. The spectra extracted from 13 regions in W44 are well reproduced by a model
consisting of CIE and RP components. We found that the electron temperaturesof the RPs in the central part of the remnant are higher than those in the otherregions. The intensity ratios of the 12 CO(J = 2–1) to 12 CO(J = 1–0) emissions arehigher in the outer regions, suggesting the interaction between the SNR shock andthe molecular cloud there. The lower kTe in the outer regions can be explained by
the thermal conduction between the plasma and the interacting molecular cloud.
4. In W49B, we found that the spectra are well fitted with a model consisting of CIE
and two RPs (RPL for the lighter elements and RPH for the heavier elements). RPH ,which is in the central region as indicated by our imaging analysis, is hotter than the
RPL in the rim, which is similarly seen in G166.0+4.3 and W44. We also found thatkTe of the western plasma is lower than that of the eastern plasma. The western
104

105
shell should be interacting with a cloud from the gamma-ray emission, suggestingcooling by the thermal conduction there.
In all the remnants, the spectral analyses revealed good correlations between theelectron temperatures and locations of the surrounding molecular clouds. For a morequantitatively investigation, we estimated the cooling timescale by the thermal conductions in each remnant. The conduction timescales tcond of G166.0+4.3, IC 443, W44 and
W49B are estimated at ∼ 40 kyr, ∼ 20 kyr, ∼ 35 kyr and ∼ 4 kyr, respectively. These
are shorter than a characteristic timescale for a plasma to reach CIE, indicating thatthat the plasmas can become RPs with the thermal conduction between the shock andthe molecular cloud. On the other hand, the requirement of tcond ≤ tage is also satisfied
for all the remnants but W44. The calculations of tcond , however, highly depend on anelectron density ne which has large uncertainties depending on the assumptions about thevolumes of the plasma and the distance to the remnant. We, therefore, calculated tcond
as a function of electron density ne . We found that the electron densities are limited tone ≲ 0.55 cm−3 for G166.0+4.3, ne ≲ 1.4 cm−3 for IC 443, ne ≲ 0.40 cm−3 for W44,
and ne ≲ 4.5 cm−3 for W49B. If the electron densities of the SNR plasmas satisfy theserequirements, the RPs can be explained by the thermal conduction scenario.

Appendix A
Spatially resolved spectroscopy of
IC 443
A.1

Fitting results of the spectra in the regions A
and B

Figure A.1 shows spectra of the PWN 1SAX J0617.1+2221 in region A and the point
source U061530.75+224910.6 in region B in the 1.6-10.0 keV band after NXB subtraction.
We fitted the spectra with a model consisting of a power law and an RP component
for the sources and the SNR plasma, respectively. The X-ray absorption was fixed at
7.0 × 1021 cm−2 as determined by Kawasaki et al. (2002). We fixed kTinit at 5 keV,
whereas kTe and ne t were allowed to vary. The abundances of Si and S were allowed tovary, whereas those of Ar and Ca were linked to S. The abundances of the other elements
were fixed to solar. The spectra can be reproduced well by the model (Table A.1). The
photon indices of PWN 1SAX J0617.1+2221 and U061530.75+224910.6 are 1.89+0.08−0.07 and
+0.15
2.22−0.16 , respectively, which we used in the SNR emission analysis (see §6.3.2). The
photon indices and the fluxes of the power-law components are consistent with the results
by Bocchino & Bykov (2003), who analyzed XMM-Newton data.
0.1

0.1

Counts s −1 keV−1

Region A

0.01
10−3

10−3
10−4
10−5

10−5
4

10−6
4

2

2

0
−2

Region B

0.01

10−4

χ

χ

Counts s −1 keV−1

1

0
−2

−4

−4
2

Energy (keV)

5

10

2

Energy (keV)

Figure A.1: Same as Figure 6.4 but for regions A and B.
106

5

10

A.2. FITTING RESULTS OF THE SPECTRA IN THE REGIONS 1–20

107

Table A.1: Best-fit model parameters of the spectra in regions A and B.
1SAX J0617.1+2221 U061530.75+224910.6
region A
region B
0.7 (fixed)
0.7 (fixed)
1.89 +0.08
2.22 +0.15
−0.07
−0.16
+114
801 −94
254 +30−50VVRNEI
0.34 +0.05
0.31 +0.49
−0.04
−0.20
(RP comp.)
5 (fixed)
5 (fixed)
1 (fixed)
2.4 +0.7−0.5
1.3 +0.4
1 (fixed)−0.3
4.2 +0.6
> 100
−0.4
+5.7
1.2 +0.2
1.0
−0.4
−1.0
1.16 (241)1.05 (149)∫†
Volume emission measure at the distance of 1.5 kpc: V EM = ne nH dV , where ne , nH ,and V are the electron and hydrogen densities, and the emitting volume, respectively.
∗
The unit is photons s−1 cm−2 keV−1 sr−1 at 1 keV.

Model function
TBabs
Power law

A.2

Parameter
NH (1022 cm−2 )Photon index
Normalization∗kTe (keV)
kTinit (keV)
ZSi (solar)
ZS = ZAr = ZCa (solar)
ne t (1011 s cm−3 )V EM (1057 cm−3 )†χ2ν (ν)Fitting results of the spectra in the regions 1–20
10

Region 1

1
0.1
0.01

Counts s −1 keV−1

Counts s −1 keV−1

10

10

10−4

10−4
4

χ

χ

2

0
−2

0
−2

−4
1

2
Energy (keV)

−4

5

1

10

2
Energy (keV)

5

Region 3

1
0.1
0.01

Counts s −1 keV−1

10

0.1
0.01
10−3

10−4

10−4

4

4

2

2

0
−2

Region 4

1

10−3

χ

Counts s −1 keV−1

0.01
10−3

2

χ

0.1

−3

4

Region 2

1

0
−2

−4
1

2
Energy (keV)

5

−4
1

2
Energy (keV)

Figure A.2: Same as Figure 6.4 but for regions 1–4.

5

108

APPENDIX A. SPATIALLY RESOLVED SPECTROSCOPY OF IC 443

10

0.1
0.01

Counts s −1 keV−1

Region 5

1

0.1
0.01
10−3

10−4

10−4

4

4

2

2
χ

0

−4

−4
1

2
Energy (keV)152Energy (keV)

5

Region 7

0.1
0.01

Counts s −1 keV−1

10

1

0.1
0.01
10−3

10−4

10−4

4

4

2

2
χ

0
−2

Region 8

1

10−3

χ

Counts s −1 keV−1

10

0
−2

−4
1

2
Energy (keV)

−4

5

1

10

Region 9

1
0.1
0.01

2
Energy (keV)

5

10Counts s −1 keV−1

Counts s −1 keV−1

0
−2

−2

10

Region 10

1
0.1
0.01

−3

10−3

10−4

10−4

4

4

2

2

0

χ

χ

Region 6

1

10−3

χ

Counts s −1 keV−1

10

−2

0
−2

−4
1

2
Energy (keV)

5

−4
1

2
Energy (keV)

Figure A.3: Same as Figure 6.4 but for regions 5–10.

5A.2. FITTING RESULTS OF THE SPECTRA IN THE REGIONS 1–20

0.1
0.01

Counts s −1 keV−1

10Region 11

1

10−4

4

4

2

2
χ

−4
1

2
Energy (keV)

512Energy (keV)

5

10

Region 13

1
0.1
0.01

Counts s −1 keV−1

Counts s −1 keV−1

1010

Region 14

1
0.1
0.01
10−3

−3

10

10−4

−4

4

4

2

2

0

χ

χ

0
−2

−4

−2

0
−2

−4
1

2
Energy (keV)

−4

5

1

10

2
Energy (keV)

5

10

Region 15

1
0.1
0.01

Counts s −1 keV−1

Counts s −1 keV−1

0.01

10−4

−2

Region 16

1
0.1
0.01

10−3

10−3

10−4

10−4

4

4

2

2

0

χ

χ

0.1

10−3

0

Region 12

1

10−3

χ

Counts s −1 keV−1

10

109

−2

0
−2

−4
1

2
Energy (keV)

5

−4
1

2
Energy (keV)

Figure A.4: Same as Figure 6.4 but for regions 11–16.

5

110

APPENDIX A. SPATIALLY RESOLVED SPECTROSCOPY OF IC 443

Region 17

1
0.1
0.01

Counts s −1 keV−1

10

10−4

4

4

2

2
0
−2

−4
1

2
Energy (keV)

−4

5

1

10

2
Energy (keV)

5

10

Region 19

1
0.1
0.01

Counts s −1 keV−1

Counts s −1 keV−1

0.01

10−4

−2

Region 20

1
0.1
0.01

10−3

10−3

10−4

10−4

4

4

2

2

0

χ

χ

0.1

10−3

0

Region 18

1

10−3

χ

χ

Counts s −1 keV−1

10

−2

0
−2

−4
1

2
Energy (keV)

5

−4
1

2
Energy (keV)

Figure A.5: Same as Figure 6.4 but for regions 17–20.

5

∗

‡

The unit is photons s−1 cm−2 sr−1 .
The unit is photons s−1 cm−2 keV−1 sr−1 at 1 keV.

TBabs
VVRNEI1(CIE comp.)

reg.2
reg.3
reg.4
reg.5
reg.6
reg.7
reg.8
reg.9
reg.100.95 +0.04
0.82 +0.03
0.79 +0.03
0.88 +0.05
0.89 +0.02
0.79 +0.02
0.70 +0.03
0.84 +0.04
0.74 +0.02
−0.04
−0.04
−0.03
−0.06
−0.02
−0.02
−0.02
−0.05
−0.03
0.26 +0.01
0.28 +0.05
0.25 +0.04
0.22 +0.05
0.21 +0.01
0.24 +0.01
0.26 +0.01
0.20 +0.06
0.23 +0.01
−0.02
−0.01
−0.01
−0.03
−0.01
−0.01
−0.01
−0.06
−0.01
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1.7 +0.6
1.1 +0.5
0.8 +0.4
0.8 +0.7
7.5 +1.1
4.3 +1.7
2.7 +0.6
1.1 +1.1
4.5 +1.0
−0.5
−0.4
−0.4
−0.3
−0.9
−0.4
−0.5
−1.0
−0.7
+0.05
+0.05
+0.02
+0.02
+0.07
+0.08
+0.01
+0.04
VVRNEI2
0.28 −0.07
0.25 −0.05
0.27 −0.02
0.25 −0.01
0.27 −0.09
0.16 −0.07
0.26 −0.01
0.25 −0.05
0.22 +0.07−0.03(RPcold comp.)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
+0.3
+0.2
+0.2
+0.3
+1.7
+0.3
+0.1
+0.2
+0.5
1.9 −0.4
2.2 −0.3
1.9 −0.1
2.3 −0.3
5.1 −1.5
2.2 −0.2
1.8 −0.1
2.0 −0.3
3.4 −0.5
1.2 +0.2
1.5 +0.1
1.4 +0.1
1.3 +0.3
2.2 +0.8
1.3 +0.2
1.4 +0.1
1.4 +0.2
1.7 +0.3
−0.2
−0.2
−0.1
−0.1
−0.6
−0.1
−0.1
−0.1
−0.3
+0.2
+0.2
+0.2
+1.6
+0.1
+0.1
+0.2
1.7 +0.2
2.7
2.5
2.4
4.9
3.4
2.2
1.5
4.2 +0.6
−0.3
−0.2
−0.1
−0.2
−1.1
−0.3
−0.2
−0.1
−0.5
1.0 +0.2
1.7 +0.2
1.6 +0.1
1.4 +0.2
3.6 +0.9
2.4 +0.3
1.6 +0.1
0.9 +0.1
3.3 +0.4
−0.2
−0.1
−0.2
−0.1
−0.9
−0.2
−0.2
−0.1
−0.4
0.9 +0.5
1.6 +0.3
1.4 +0.2
1.3 +0.3
2.4 +1.2
2.5 +0.4
1.9 +0.3
1.0 +0.3
3.1 +0.8
−0.4
−0.2
−0.3
−0.3
−1.2
−0.4
−0.2
−0.2
−0.6
0.4 +0.1
0.5 +0.1
0.5 +0.1
0.8 +0.1
1.1 +0.5
0.6 +0.1
0.5 +0.1
0.8 +0.2
0.8 +0.2
−0.1
−0.1
−0.1
−0.2
−0.3
−0.1
−0.1
−0.1
−0.1
+0.6
+0.4
+0.2
+0.2
+0.5
+0.1
+0.2
+0.4
5.6 −0.6
5.4 −0.2
5.1 −0.2
4.4 −0.3
5.6 −0.4
5.4 −0.3
5.7 −0.2
5.4 −0.3
5.6 +0.2−0.3
1.1 +0.7
1.2 +0.4
1.8 +0.3
2.1 +0.3
0.1 +0.1
0.3 +0.2
1.8 +0.6
3.2 +0.5
0.3 +0.1
−0.5
−0.3
−0.4
−0.3
−0.1
−0.2
−0.5
−0.5
−0.1
VVRNEI3
0.62 +0.07
0.62 +0.02
0.59 +0.03
0.67 +0.06
0.64 +0.06
0.55 +0.03
0.57 +0.02
0.61 +0.04
0.61 +0.03
−0.07
−0.03
−0.02
−0.05
−0.04
−0.02
−0.01
−0.06
−0.03
(RPhot comp.)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold
0.5 +0.3
1.4 +0.2
1.4 +0.2
0.4 +0.1
0.3 +0.1
1.5 +0.2
2.2 +0.2
0.6 +0.1
0.6 +0.1
−0.2
−0.2
−0.2
−0.1
−0.1
−0.1
−0.2
−0.1
−0.1
Gaussian
1.23 (fixed)
1.23 (fixed) 1.23 (fixed) 1.23 (fixed) 1.23 (fixed) 1.23 (fixed) 1.23 (fixed) 1.23 (fixed) 1.23 (fixed)
19 +11
47 +1046 +1324 +133 +9
42 +1130 +8
22 +830 +10
−11
−13
−12
−6
−3
−12
−8
−8
−12
Power law
————
————
————
————
————
————
————
————
————
————
————
————
————
————
————
————
————
————
1.15 (477)
1.12 (485)1.32 (477)
1.17 (477)
1.24 (365)1.39 (477)
1.30 (477)1.04 (365)1.16 (477)∫†
Volume emission measure at the distance of 1.5 kpc: V EM = ne nH dV , where ne , nH , and V are the electron and hydrogen densities, and the emitting volume, respectively.

Parameter

NH (1022 cm−2 )kTe (keV)
Zall (solar)
V EM (1057 cm−3 )†kTe (keV)
kTinit (keV)
ZNe (solar)
ZMg (solar)
ZSi (solar)
ZS (solar)
ZAr = ZCa (solar)
ZFe = ZNi (solar)
ne t (1011 s cm−3 )V EM (1057 cm−3 )†kTe (keV)
kTinit (keV)
Zall (solar)
ne t (1011 s cm−3 )V EM (1057 cm−3 )†Centroid (keV)
Normalization‡
Photon index
Normalization∗
χ2ν (ν)Model function

reg.1
0.94 +0.04−0.04
0.24 +0.02−0.01
1 (fixed)
2.9 +1.3−0.9
0.22 +0.01−0.01
5 (fixed)+0.4
2.5 −0.3
1.5 +0.1−0.2
2.0 +0.2−0.2
1.2 +0.1−0.2
1.3 +0.3−0.2
0.7 +0.1−0.1
4.1 +0.3−0.2
3.3 +0.5−0.5
0.56 +0.04−0.05
5 (fixed)= RPcold= RPcold
1.4 +0.2−0.3
1.23 (fixed)
14 +7
−7
————
————
1.13 (477)

Table A.2: Best-fit model parameters of the spectra from regions 1–10.

A.2. FITTING RESULTS OF THE SPECTRA IN THE REGIONS 1–20
111

∗

‡

reg.110.74 +0.04−0.03
0.24 +0.01−0.01
1 (fixed)
2.7 +0.8−0.4
0.24 +0.05−0.07
5 (fixed)+0.3
2.0 −0.3
1.5 +0.3−0.2
1.7 +0.4−0.2
1.1 +0.2−0.2
1.8 +0.6−0.5
0.5 +0.2−0.1
5.8 +0.5−0.4
0.5 +0.4−0.4
0.54 +0.03−0.08
5 (fixed)= RPcold= RPcold
0.5 +0.3−0.2
1.23 (fixed)
28 +8
−8
————
————
1.18 (365)

Table A.3: Best-fit model parameters of the spectra from regions 11–20.

The unit is photons s−1 cm−2 sr−1 .
The unit is photons s−1 cm−2 keV−1 sr−1 at 1 keV.

NH

Parameter(1022 cm−2 )reg.12
reg.13
reg.14
reg.15
reg.16
reg.17
reg.18reg.19reg.200.81 +0.02
0.88 +0.03
0.98 +0.04
0.77 +0.03
0.91 +0.14
0.96 +0.03
0.85 +0.04
0.67 +0.02
0.77 +0.05
−0.03
−0.03
−0.05
−0.03
−0.09
−0.03
−0.03
−0.03
−0.02
kTe (keV)
0.22 +0.01
0.24 +0.02
0.27 +0.02
0.20 +0.02
0.24 +0.05
0.19 +0.01
0.25 +0.05
0.27 +0.01
0.24 +0.01
−0.01
−0.02
−0.03
−0.01
−0.05
−0.01
−0.06
−0.01
−0.01
Zall (solar)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
1 (fixed)
V EM (1057 cm−3 )†4.5 +0.6
2.1 +0.6
0.9 +0.3
5.5 +1.5
0.7 +2.0
5.3 +1.3
1.0 +0.3
1.3 +0.2
1.7 +0.8
−0.7
−0.5
−0.3
−1.5
−0.4
−0.7
−0.4
−0.2
−0.6
+0.02
+0.02
+0.01
+0.01
+0.04
+0.02
+0.04
+0.09
VVRNEI1
kTe (keV)
0.26 +0.09
0.21
0.19
0.18
0.19
0.19
0.22
0.24
0.22
−0.08
−0.11
−0.01
−0.01
−0.02
−0.03
−0.02
−0.05
−0.08
(RPcold comp.)
kTinit (keV)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
+5.4
+0.3
+0.5
+0.7
+2.0
+0.4
+0.3
+0.1
+1.1
ZNe (solar)
6.8 −2.9
2.3 −0.4
2.4 −0.4
4.4 −0.6
3.9 −0.9
5.0 −0.5
1.9 −0.3
1.5 −0.2
2.6 −0.7
+0.1
+0.2
+0.3
+0.8
+0.2
+0.1
+0.1
+0.4
ZMg (solar)
2.9 +4.1
1.3
1.4
1.9
2.0
2.4
1.1
1.1
1.5
−1.3
−0.3
−0.3
−0.1
−0.5
−0.2
−0.2
−0.1
−0.2
+5.3
+0.1
+0.2
+0.3
+0.6
+0.2
+0.1
+0.1
+0.4
ZSi (solar)
9.5 −3.5
1.6 −0.2
1.2 −0.2
2.3 −0.2
1.9 −0.3
2.5 −0.2
0.7 −0.2
0.7 −0.1
1.4 −0.2
ZS (solar)
6.7 +14.3
0.9 +0.1
0.8 +0.1
1.4 +0.1
1.1 +0.4
1.8 +0.2
0.3 +0.1
0.4 +0.1
0.8 +0.3
−2.5
−0.1
−0.1
−0.2
−0.2
−0.2
−0.1
−0.1
−0.2
+0.2
+0.2
+0.3
+0.6
+0.5
ZAr = ZCa (solar)
5.2 +5.4
0.8
0.8
1.4
1.2
2.0
=
Z
=
Z
= ZS
S
S
−3.1
−0.2
−0.2
−0.2
−0.5
−0.5
ZFe = ZNi (solar)
0.8 +1.2
0.3 +0.1
0.4 +0.2
0.8 +0.2
0.7 +0.5
0.6 +0.1
0.3 +0.2
0.2 +0.1
0.4 +0.1
−0.5
−0.1
−0.2
−0.2
−0.2
−0.2
−0.1
−0.1
−0.1
+0.7
+0.3
+0.4
+0.2
+0.4
+0.2
+0.7
+2.3
11
−3
ne t (10 s cm )5.7 −0.5
3.8 −0.3
3.8 −0.4
4.3 −0.1
4.7 −0.5
4.8 −0.2
5.8 −0.7
12.3 −1.7
15.1 +8.3−4.1
V EM (1057 cm−3 )†≤ 0.1
3.7 +0.5
2.3 +0.3
3.2 +0.6
1.1 +0.3
1.7 +0.6
1.5 +0.6≤ 0.2
0.2 +0.4
−0.4
−0.5
−0.5
−0.4
−0.8
−0.4
−0.2
VVRNEI3
kTe (keV)
0.63 +0.10
0.48 +0.07
0.59 +0.07
0.54 +0.04
0.64 +0.06
0.55 +0.04
0.63 +0.07
0.62 +0.02
0.62 +0.03
−0.08
−0.19
−0.14
−0.03
−0.11
−0.04
−0.04
−0.01
−0.04
(RPhot comp.)
kTinit (keV)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
5 (fixed)
Zall (solar)= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcoldne t (1011 s cm−3 )= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold= RPcold
V EM (1057 cm−3 )†0.1 +0.1
1.2 +1.1
0.5 +0.4
0.9 +0.2
0.3 +0.1
0.7 +0.2
0.3 +0.1
0.8 +0.1
0.3 +0.1
−0.1
−1.0
−0.1
−0.1
−0.1
−0.1
−0.1
−0.1
−0.1
GaussianCentroid (keV)
1.23 (fixed)
1.23 (fixed) 1.23 (fixed) 1.23 (fixed) 1.23 (fixed) 1.23 (fixed) 1.23 (fixed) 1.23 (fixed) 1.23 (fixed)
Normalization‡
5 +4
12 +7
15 +8
9 +5
11 +7
4 +4
4 +3
5 +1
6 +2
−5
−7
−4
−5
−8
−4
−2
−1
−2
Power law
Photon index
————
————
————
————
1.89 (fixed) 1.89 (fixed)
————
————
2.22 (fixed)
Normalization∗
————
————
————
————
45 +9
97 +7
————
————
18 +4
−9
−7
−5
χ2ν (ν)1.35 (250)
1.25 (477)
1.22 (365)1.53 (477)
1.13 (364)
1.36 (476)1.19 (273)
1.27 (257)
1.18 (256)∫†
Volume emission measure at the distance of 1.5 kpc: V EM = ne nH dV , where ne , nH , and V are the electron and hydrogen densities, and the emitting volume, respectively.

TBabs
VVRNEI(CIE comp.)

Model function

112
APPENDIX A. SPATIALLY RESOLVED SPECTROSCOPY OF IC 443

Appendix B
Spatially resolved spectroscopy of
W44
B.1

Fitting results of the spectra in the regions 2–10
10

Region 1

1

Counts s −1 keV−1

Counts s −1 keV−1

10

0.1
0.01
10−3

0.1
0.01
10−3

0

0

1

2
Energy (keV)

5Counts s −1 keV−1

1

−5

10

Region 3

0.1
0.01
10−3

2
Energy (keV)

5

10

Region 4

0.1
0.01
10−3

10

−4

10−4

5

5

χ

0
−5

1

1
Counts s −1 keV−1

−5

χ

5

χ

10−4
5

χ

Region 2

1

1

2
Energy (keV)

5

10

0
−5

1

2
Energy (keV)

Figure B.1: Same as Figure 7.5 but for regions 1–4.

113

5

10

114

APPENDIX B. SPATIALLY RESOLVED SPECTROSCOPY OF W44

1

Region 5

0.1

Counts s −1 keV−1

Counts s −1 keV−1

1

0.01
10−3
10−4
10

0

χ

0

1

2
Energy (keV)

5Counts s −1 keV−1

0.1
0.01
10−3

5

0

0
−5

10

Region 8

1

2
Energy (keV)

5

10

1

1

Region 9

0.1
0.01
10−3

Region 10

0.1
0.01
10−3
10−4

10−4

10−5
5

0

0

χ

5

−5

10

10−3

5

5

5

0.01

10−4

2
Energy (keV)

2
Energy (keV)

0.1

10−4

1

1

1

Region 7

χCounts s −1 keV−1
χ

−5

10

Counts s −1 keV−1

χ

5

−5

Counts s −1 keV−1

10−3

5

1

χ

0.01

10−4

−5

−5

Region 6

0.1

12
Energy (keV)

5

10

−5

1

2
Energy (keV)

Figure B.2: Same as Figure 7.5 but for regions 5–10.

510

2.19 +0.07−0.05
0.23 +0.01−0.02= Zall,Whole+2.35
5.13 −1.58
0.35 ± 0.01
0.93 (fixed)
3.6 (fixed)
2.9 ± 0.3
4.8 +0.3−0.5
4.6 +0.5−0.6
7.3 +3.2−1.8< 0.362.2 ± 0.3
1.10 ± 0.07
1.23 (fixed)
17 ± 81.37 (499)reg.3

The unit is photons s−1 cm−2 sr−1 .

The emission measure integrated over the line of sight, i.e., (1 / 4π)

Gaussian

2.03 +0.01−0.02
0.22 ± 0.01
= Zall,Whole
12.47 +1.64−3.10
0.42 +0.01−0.02
0.93 (fixed)
3.6 (fixed)
2.5 +0.2−0.1
4.1 ± 0.2
3.5 +0.3−0.2
5.4 +1.3−1.2< 0.094.1 ± 0.1
2.00 +0.07−0.06
1.23 (fixed)
38 ± 81.39 (1100)

reg.2∗ The unit is photons s−1 cm−2 keV−1 sr−1 at 1 keV.

‡†

NH (1022 cm−2 )kTe (keV)
Zall (solar)
Normalization†kTe (keV)
kTinit (keV)
ZO = ZNe (solar)
ZMg (solar)
ZSi (solar)
ZS (solar)
ZAr = ZCa (solar)
ZFe = ZNi (solar)
ne t (1011 s cm−3 )Normalization†
Centroid (keV)
Normalization‡
χ2ν (ν)TBabs
VVRNEI1(CIE comp.)

VVRNEI2(RP comp.)

Parameter

Model function∫

1.87 ± 0.06
0.20 +0.03−0.01= Zall,Whole+3.56
6.35 −2.96
0.38 +0.05−0.04
0.93 (fixed)
3.6 (fixed)
2.7 +0.6−0.4
3.4 +0.8−0.6
3.0 +1.4−0.8
3.7 +8.8−3.5< 0.76
3.7 +0.3−0.5
0.77 +0.12−0.09
1.23 (fixed)
31 +12−11
1.08 (265)

reg.5

reg.6
1.93 ± 0.06
0.20 +0.03−0.02= Zall,Whole+1.74
1.82 −1.17
0.26 ± 0.01
0.93 (fixed)
3.6 (fixed)
3.0 ± 0.3
7.9 ± 0.7
7.9 +1.5−1.3= ZS< 0.494.0 ± 0.4
0.63 +0.05−0.06
1.23 (fixed)
14 ± 51.09 (395)

ne nH dl in units of 1018 cm−5 sr−1 .

1.96 ± 0.03
0.20 +0.02−0.01= Zall,Whole
11.78 +3.17−3.06
0.44 ± 0.02
0.93 (fixed)
3.6 (fixed)
2.8 ± 0.2
3.1 +0.2−0.3
2.6 ± 0.4
3.4 +1.1−1.5
0.31 +0.11−0.10
3.1 ± 0.2
1.90 +0.08−0.10
1.23 (fixed)
52 +12−13
1.41 (498)

reg.4
2.03 +0.07−0.06
0.21 +0.01−0.02= Zall,Whole
4.72 +2.79−1.94
0.28 ± 0.02
0.93 (fixed)
3.6 (fixed)
2.1 ± 0.5
3.5 +0.7−0.6
2.3 +1.3
−0.9
= ZS
< 1.7
4.9 1.1
1.0
0.70 +0.16−0.11
1.23 (fixed)
18 +6−5
1.02 (395)

reg.7

reg.8
2.16 +0.06−0.07
0.21 +0.01−0.02= Zall,Whole
3.13 +1.60−1.18
0.37 +0.05−0.03
0.93 (fixed)
3.6 (fixed)
2.6 +0.5−0.4
3.6 +0.8−0.7
3.2 +0.9−0.7
1.9 +4.2−1.7< 0.492.0 +0.6−0.4
0.44 +0.05−0.04
1.23 (fixed)
9 +6−5
1.19 (394)

Table B.1: Best-fit model parameters of the spectra from regions 2–10.
2.14 ± 0.05
0.18 ± 0.01
= Zall,Whole
17.55 +8.78−5.84
0.42 ± 0.04
0.93 (fixed)
3.6 (fixed)
2.3 ± 0.3
3.0 +0.5−0.6
2.6 ± 0.7
5.4 +4.4−3.4
0.43 +0.25−0.19
3.7 +0.5−0.4
1.15 +0.12−0.13
1.23 (fixed)
22 ± 14
1.33 (328)

reg.9

1.90 +0.08−0.09
0.18 +0.03−0.01= Zall,Whole
4.50 +3.69−2.72
0.37 +0.13−0.05
0.93 (fixed)
3.6 (fixed)
2.8 +0.8−0.6
4.9 +1.4−1.6
4.2 +2.6
−2.3
= ZS
< 1.0
3.9 +0.8−0.7
0.31 +0.08−0.12
1.23 (fixed)
16 +7−6
1.26 (266)

reg.10B.1. FITTING RESULTS OF THE SPECTRA IN THE REGIONS 2–10
115

Acknowledgement
I am deeply grateful to Dr. Takaaki Tanaka, Dr. Hiroyuki Uchida, and Prof. Takeshi
Go Tsuru for their helpful advice and support. I am very grateful to my collaborators, Dr. Masatoshi Nobukawa, Dr. Kumiko Kawabata Nobukawa, Mr. Makoto Itou, andMr. Hiromichi Okon. I would like to thank Dr. Satoshi Yoshiike for providing us with
the NANTEN2 data. I also grateful to the Suzaku team members for their constanteﬀorts. This work is supported by Research Fellowships of the Japan Society for the
Promotion of Science for Young Scientists. I am indebted to my friends and family for
their understanding and kind supports.

116

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