K YOTO U NIVERSITY
D OCTORAL T HESIS

Diagnostics for Physical Processesof X-ray Plasma in Supernova
Remnants

Author:Yuki A MANOA thesis submitted in fulfillment of the requirements
for the degree of Doctor of Philosophy
in theCosmic ray Group
Department of Physics

April 20, 2023

iii

KYOTO UNIVERSITY

Abstract
Faculty Name
Department of Physics
Doctor of Philosophy
Diagnostics for Physical Processes of X-ray Plasma in Supernova
Remnants
by Yuki A MANO

iv
X-ray imaging spectroscopy of Supernova remnants (SNRs) enables us to estimate the abundances and spatial distribution of heavy elements synthesized by the progenitors. This information provide us with many insights
into the evolution process and explosion mechanism of the progenitors. Improvement of the performance of X-ray detectors have revealed that the formation process, ionization state, and X-ray emission process of SNRs are
more complex than that expected by the conventional picture. For example,
recent studies point out the importance of charge exchange (CX) and resonance scattering (RS) in SNRs. Since signs of the RS or CX are expected to
be found in intensity ratios of multiplet lines (e.g., OVII Heα), which cannot
be resolved with widely-used detectors such as CCDs, it is difficult to obtain the observational evidence for CX or RS. If we can quantify the effectsof these processes, we can measure several important information such as
microturbulence velocities and 3D structure of SNRs.
We performed a high-resolution X-ray spectroscopy of N49, J0453.6−6829
with the RGS onboard XMM-Newton to obtain the observational evidence of
CX and RS. The RGS spectrum of N49 shows a high G-ratio of OVII Heαlines as well as OVIII Lyβ/α and FeXVII (3s–2p)/(3d–2p) ratios, which cannot be explained by the emission from an optically thin thermal plasma.
These line ratios can be well explained by the effect of RS. In the case of
the J0453.6−6829, we find a high f /r ratio of OVII Heα lines and OVIII
Lyβ/α ratio. The spectrum is fairly explained by taking into account a CX
emission in addition to the thermal component. Analyzing archival ATCA &Parkes radio data, we also reveal that H I cloud is possibly interacting with
J0453.6−6829. These results support the presence of CX in J0453.6−6829, as
the origin of the obtained high f /r ratio. Although a contribution of the RS
cannot be ruled out at this time, we conclude that CX seems more likely than
RS considering the relatively symmetric morphology of this remnant.
We present a method to constrain the 3D structure of SNRs by measuring
their line-of-sight length through the effects of RS. Additionally, we apply
this methodology to one of the brightest extra-Galactic SNRs, 1E 0102−72.9(hereafter E0102), whose 3D structure is currently a controversial. Previous
studies have proposed various structures for 1E 0102.2−7219, such as barrellike and cylindrical structure. These structures require strong scattering effect, which is inconsistent with our results. Therefore, we have proposed adouble-ring structure and a cylindrical structure with the expansion velocity as structures that are consistent with both our observations and previous
studies. This method can be applied to many celestial objects using the futuresatellite XRISM.
The RGS spectra of N49 and J0453.6−6829 show the excess around 16 Å,which cannot be explained by the effect of CX and RS. We attribute these discrepancies to uncertainties in the atomic data related to Fe-L emission lines.
To solve these problems, it is necessary to experimentally measure atomic
data. Ground experiments, such as EBIT are indispensable for the future development of X-ray astronomy.

v

Acknowledgements
I would like to express my gratitude to Prof. Takeshi Tsuru, Ass. Prof. Hiroyuki Uchida and Assoc. Prof. Takaaki Tanaka for guiding me throughout
the five years at Kyoto university. I am very grateful to my collaborators, Dr.
Liyi Gu, Dr. Hidetoshi Sano, Dr. Hiroya Yamaguchi and Mr. Yosuke Koshiba
I would like to thank Dr. Hidetoshi Sano for providing us with the ATCA
& Parke data. Mr. Yosuke Koshiba played the most important role in the
research of J0453. I appreciate the supports and encouragements of all thepeople in the Cosmic ray group and JAXA/ISAS. Finally, I thank my family
for their support.

viiContents
Abstractiv

Acknowledgementsv

1

Introduction

1

2

Supernova and Supernova Remnant
2.1 Stellar Evolution of the Progenitor . . . . . . . . . . . . . . . .
2.1.1 Stellar Evolution . . . . . . . . . . . . . . . . . . . . . .
2.1.2 Stellar Nucleosynthesis . . . . . . . . . . . . . . . . . .
2.1.3 Mass Loss of Massive Stars and Their Environment . .
Single Massive Stars . . . . . . . . . . . . . . . . . . . .
Binary Stars . . . . . . . . . . . . . . . . . . . . . . . . .
2.2 Supernova . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.2.1 Explosive Nucleosynthesis . . . . . . . . . . . . . . . .
2.2.2 Classification of Supernova . . . . . . . . . . . . . . . .
2.2.3 Type Ia SNe . . . . . . . . . . . . . . . . . . . . . . . . .
2.2.4 Core-collapse SNe . . . . . . . . . . . . . . . . . . . . .
2.3 Supernova remnant . . . . . . . . . . . . . . . . . . . . . . . . .
2.3.1 Basic Physics of Supernova Remnants . . . . . . . . . .
Shock Wave Heating . . . . . . . . . . . . . . . . . . . .
Evolution of Supernova Remnant . . . . . . . . . . . .
2.3.2 Observational View of Supernova Remnants . . . . . .

3
3
3
5
7
7
8
11
11
11
13
13
17
17
17
19
20

3

Plasma Processes and X-ray Radiation in Collisionally Ionized Plasma 27
3.1 Collisionally Ionized Plasma . . . . . . . . . . . . . . . . . . . . 27
3.1.1 Ionization and Recombination . . . . . . . . . . . . . . 27
Ionization and Recombination Processes . . . . . . . . 27
Ionization Distribution . . . . . . . . . . . . . . . . . . . 30
3.1.2 Atomic Structure and Radiation Processes . . . . . . . 33
Atomic Structure . . . . . . . . . . . . . . . . . . . . . . 33
Radiation Processes . . . . . . . . . . . . . . . . . . . . 36
Line Emission . . . . . . . . . . . . . . . . . . . . . . . . 36
Continuum Emission . . . . . . . . . . . . . . . . . . . . 39
3.2 Interaction between Collisionally Ionized Plasma and Neutral
Gas . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 40
3.2.1 Charge Exchange . . . . . . . . . . . . . . . . . . . . . . 40
3.3 Radiative Transfer . . . . . . . . . . . . . . . . . . . . . . . . . . 42
3.3.1 Optical Depth . . . . . . . . . . . . . . . . . . . . . . . . 42
3.3.2 Radiative Transfer . . . . . . . . . . . . . . . . . . . . . 42

viii

3.4

3.5

3.3.3 Resonance Scattering . . . . . . . . . . . . . . . . . . . .
Plasma Diagnostic . . . . . . . . . . . . . . . . . . . . . . . . . .
3.4.1 Electron Temperture & Charge Distribution (Ionization
timescale) . . . . . . . . . . . . . . . . . . . . . . . . . .
Electron Temperture . . . . . . . . . . . . . . . . . . . .
Charge Distribution (Ionization timescale) . . . . . . .
3.4.2 Diagnostic Using Heα Line ratio . . . . . . . . . . . . .
R-ratio . . . . . . . . . . . . . . . . . . . . . . . . . . . .
G-ratio . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Observational Evidence for Charge Exchange and Resonance
Scattering in Astrophysical Plasma . . . . . . . . . . . . . . . .
3.5.1 CX . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.5.2 RS . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

43
45
45
45
46
46
47
48
51
51
52

4

Instruments
4.1 XMM-Newton satelite . . . . . . . . . . . . . . . . . . . . . . .
4.2 X-ray Telescopes . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.1 Overview . . . . . . . . . . . . . . . . . . . . . . . . . .
4.2.2 Performance . . . . . . . . . . . . . . . . . . . . . . . . .
Imaging Performance . . . . . . . . . . . . . . . . . . .
Effective Area . . . . . . . . . . . . . . . . . . . . . . . .
4.3 X-ray CCD Camera (European Photon Imaging Camera: EPIC)
4.3.1 Overview and performance of MOS . . . . . . . . . . .
4.3.2 Background of the MOS . . . . . . . . . . . . . . . . . .
4.4 Reflection Grating Spectrometer (RGS) . . . . . . . . . . . . . .
4.4.1 Overview . . . . . . . . . . . . . . . . . . . . . . . . . .
RGA . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
RFC . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
4.4.2 Performance . . . . . . . . . . . . . . . . . . . . . . . . .
4.4.3 Background . . . . . . . . . . . . . . . . . . . . . . . . .
4.4.4 Response Function . . . . . . . . . . . . . . . . . . . . .

55
55
57
57
57
57
59
59
60
61
62
62
64
64
64
68
69

5

Evidence for RS and CX in SNRs
5.1 Our Objective and Target Selection . . . . . . . . . . . . . . . .
5.2 N49 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.2.1 Overview of N49 . . . . . . . . . . . . . . . . . . . . . .
5.2.2 Observation and Data Reduction . . . . . . . . . . . . .
5.2.3 Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . .
Spectral Analysis with CIP Model . . . . . . . . . . . .
Plasma Diagnostic Using O VII Heα Line Ratio . . . . .
5.2.4 Discussion . . . . . . . . . . . . . . . . . . . . . . . . . .
Physical Origin of High G-ratio . . . . . . . . . . . . . .
O VIII Lyβ/α ratio . . . . . . . . . . . . . . . . . . . . .
5.2.5 Conclusion of N49 . . . . . . . . . . . . . . . . . . . . .
5.3 J0453.6−6829 . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.3.1 Overview of J0453.6−6829 . . . . . . . . . . . . . . . . .
5.3.2 Observation and Data Reduction . . . . . . . . . . . . .
5.3.3 Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . .

71
71
72
72
72
73
75
75
76
76
83
85
86
86
87
87

ix

5.4

Spectral Analysis with CIP Model . . . . . . . . . . . . 87
Plasma Diagnostic Using O VII Heα Line Ratio . . . . . 90
Spectral Analysis with 2-NEI+CX Model . . . . . . . . 90
Spectral Analysis with 2-NEI−Gaus Model (RS Model) 92
5.3.4 Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . 92
X-ray Morphology of J0453.6−6829 . . . . . . . . . . . 94
Surrounding Environment of J0453.6−6829 . . . . . . . 95
5.3.5 Conclusion for J0453.6−6829 . . . . . . . . . . . . . . . 96
Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 100

6Constraint on 3D Structure of SNRs from X-ray Resonance Scattering
101
6.1 Method . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 101
6.2 Overview of E0102 . . . . . . . . . . . . . . . . . . . . . . . . . 103
6.3 Observations and Data Reduction . . . . . . . . . . . . . . . . 103
6.4 Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110
6.4.1 Spectral Modeling . . . . . . . . . . . . . . . . . . . . . 110
6.4.2 Plasma Diagnostics and Transmission Factor Measurement . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110
6.5 Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 116
6.5.1 Cylindrical Model . . . . . . . . . . . . . . . . . . . . . 116
6.5.2 Barrel-Like Structure . . . . . . . . . . . . . . . . . . . . 120
6.5.3 Structures That Explain Our Observation . . . . . . . . 124
Double-Ring Structure . . . . . . . . . . . . . . . . . . . 124
Cylindrical Model Considering Velocity Structure . . . 124
6.5.4 Progenitor of E0102 . . . . . . . . . . . . . . . . . . . . . 129
6.6 Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 129

7

Discussion
7.1 Summary of Our Results and Previous Studies . . . . . . . . .
7.1.1 Physical Conditions in which RS and/or CX Efficiently
Occur . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.1.2 Spectral Feature to Distinguish between CX and RS . .
7.2 Effect on Elemental Abundances Measurement and Progenitor
Estimation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
7.3 Constraint on 3D Structure of SNRs Using XRISM . . . . . . .
7.4 Origin of Excess around 16 Å (0.78 keV) . . . . . . . . . . . . .

8Conclusions

Bibliography

131
131
131
132
134
135
138
141
143

xi

List of Figures
1.1

History of energy resolution of X-ray spectrometers in X-ray
astronomy . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

2.1
2.2
2.3
2.4
2.5

Hertzsprung Russell diagram for massive stars in the SMC . .
p–p chain reaction . . . . . . . . . . . . . . . . . . . . . . . . . .
CNO cycle . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Onion-like structure in massive stars . . . . . . . . . . . . . . .
Evolution of the velocity and mass-loss rate of the wind fromthe star with time. . . . . . . . . . . . . . . . . . . . . . . . . . .
2.6 Density plots of the 2D evolution of the bubble during the WR
phase. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.7 Roche potential contours . . . . . . . . . . . . . . . . . . . . . .
2.8 Schematic view of mass transfer in binary stars . . . . . . . . .
2.9 Optical spectra of SN . . . . . . . . . . . . . . . . . . . . . . . .
2.10 Light curves of Type Ia, Ib, II-P, and II-L supernovae . . . . . .
2.11 Classification of SNs . . . . . . . . . . . . . . . . . . . . . . . .
2.12 SNSNR . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.13 Relationship between the progenitor mass and elemental abundances . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
2.14 Schematic image of an SNR . . . . . . . . . . . . . . . . . . . .
2.15 A schematic view around the shock front . . . . . . . . . . . .
2.16 Chandra X-ray (0.5–2.1 keV) images of the 24 SNRs . . . . . .
2.17 SNR morphologies . . . . . . . . . . . . . . . . . . . . . . . . .
2.18 X-ray and optical image of SN1987A and schematic view ofgeometrical model of SN1987A . . . . . . . . . . . . . . . . . .
2.19 Schematic image of the formation of the triple-ring nebula . .
2.20 3D simulation of SN and distribution of heavy elements in
Cassiopeia A . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.1
3.2

Binding energy of atomic subshells for neutral atoms . . . . .
Ionization and recombination processes in collisionally ionized plasma . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
3.3 Ionization distribution of CIE plasma . . . . . . . . . . . . . .
3.4 Ionization distribution of IP . . . . . . . . . . . . . . . . . . . .
3.5 Radiation processes in SNRs . . . . . . . . . . . . . . . . . . . .
3.6 A Diagram of CX . . . . . . . . . . . . . . . . . . . . . . . . . .
3.7 Lyβ/α ratio of O and Si as a function of electron temperature .
3.8 Temperature plot . . . . . . . . . . . . . . . . . . . . . . . . . .
3.9 Schematic Grotrian diagram for a He-like ion . . . . . . . . . .
3.10 R-ratio as a function of electron density . . . . . . . . . . . . .

2
4
5
6
7
8
9
10
11
12
12
13
14
16
17
18
21
22
23
24
25
28
28
32
34
37
41
45
46
47
48

xii
3.11 G-ratio as a function of electron temperature . . . . . . . . . .
3.12 Experimental He-like and H-like Fe spectra from collisional
excitation and charge exchange . . . . . . . . . . . . . . . . . .
3.13 Schematic image of the RS in SNRs . . . . . . . . . . . . . . . .
3.14 Contribution of X-ray emission from thermal plasma and CX
as functions of the impact parameter p . . . . . . . . . . . . . .
3.15 Hitomi SXS observation of Perseus cluster (Collaboration, 2017)3.16 RGS monochromatic images of DEM L 71 . . . . . . . . . . . .

52
53
54

4.1 Schematic image of the XMM-Newton satellite. . . . . . . . . .
4.2 Orbit of the XMM-Newton satellite . . . . . . . . . . . . . . . .
4.3 Light path of the X-ray tescope . . . . . . . . . . . . . . . . . .
4.4 Fractional encircled energy of the MOS 1 and 2 . . . . . . . . .
4.5 Effective areas of EPIC cameras and RGS spectrometers . . . .
4.6 Vignetting factor of the telescope of the XMM-Newton . . . .
4.7 Cross sectional view of CCDs . . . . . . . . . . . . . . . . . . .
4.8 Schematic view of MOS and pn . . . . . . . . . . . . . . . . . .
4.9 Energy resolution of MOS . . . . . . . . . . . . . . . . . . . . .
4.10 Background spectrum of MOS . . . . . . . . . . . . . . . . . . .
4.11 Schematic view of RGS . . . . . . . . . . . . . . . . . . . . . . .
4.12 Schematic view of RGA . . . . . . . . . . . . . . . . . . . . . .
4.13 Sketch of an RFC . . . . . . . . . . . . . . . . . . . . . . . . . .
4.14 Banana plot of RGS . . . . . . . . . . . . . . . . . . . . . . . . .
4.15 Effective area and energy resolution of the RGS . . . . . . . . .
4.16 Line spread function of RGS . . . . . . . . . . . . . . . . . . . .

56
57
58
58
59
60
61
62
63
63
65
66
66
67
68
69

5.1
5.2

49
50
51

X-ray image and radio continuum contour of N49 . . . . . . . 72
Radio continuum contours of N49 and emission from a molecular cloud (12 CO) . . . . . . . . . . . . . . . . . . . . . . . . . . 73
5.3 Light curve of the RGS 1 . . . . . . . . . . . . . . . . . . . . . . 74
5.4 MOS image of N49 . . . . . . . . . . . . . . . . . . . . . . . . . 74
5.5 RGS and MOS spectra of N49 and best-fit “NEI” model . . . . 76
5.6 Close-up views of the RGS spectrum of N49 around Fe XVII
Lα and O VIII Lyβ . . . . . . . . . . . . . . . . . . . . . . . . . . 77
5.7 Relation between kTe and the G-ratio of OVII Heα for CIE and
IP . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 78
5.8 Relation between kTe and the G-ratio of OVII Heα for RP . . . 80
5.9 The fitting result of O VIII Lyα and O VII Heα lines when the
ISM plasma of N49 is RP. . . . . . . . . . . . . . . . . . . . . . . 81
5.10 RGS and MOS spectra of N49 and best-fit “NEI + CX” model . 82
5.11 NEI - Gaussians . . . . . . . . . . . . . . . . . . . . . . . . . . . 82
5.12 Transmission factors of N49 . . . . . . . . . . . . . . . . . . . . 84
5.13 Fitting results of J0453.6−6829 (Haberl et al., 2012) . . . . . . . 86
5.14 Three-color image of SNR J0453.6−6829 . . . . . . . . . . . . . 88
5.15 RGS and MOS spectra of J0453.6−6829 . . . . . . . . . . . . . . 89
5.16 f /r ratio of OVII Heα as a function of electron temperature kTe 91
5.17 Fitting result of J0453 spectra with "2-NEI+CX" and "2-NEI−Gaus"
models . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 93

xiii
5.18 Transmission factors for each line of J0453.6−6829 . . . . . . . 94
5.19 Example schematic view of an X-ray emitting plasma of J0453
95
5.20 ATCA & Parkes HI channel maps . . . . . . . . . . . . . . . . . 97
5.21 Integrated intensity map of the ATCA & Parkes HI . . . . . . . 98
5.22 RGS1+2 spectra of the NE and SW regions of J0453.6−6829 . . 98
5.23 Spatially resolved spectroscopy of J0453.6−6829 with CX model 99
6.1
6.2
6.3

Sketch of an SNR . . . . . . . . . . . . . . . . . . . . . . . . . .
RGB image of E0102 . . . . . . . . . . . . . . . . . . . . . . . .
Radially-averaged X-ray surface brightness of E0102 and thebest-fit spatial model . . . . . . . . . . . . . . . . . . . . . . . .
6.4 Ne X Lyα image of E0102 and 3D structure proposed by Flanagan et al. (2004) . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.5 EPIC-MOS image of E0102 . . . . . . . . . . . . . . . . . . . . .
6.6 RGS spectra of E0102 . . . . . . . . . . . . . . . . . . . . . . . .
6.7 RGS spectra of E0102 and best-fit “NEI” model . . . . . . . . .
6.8 Line flux ratio map . . . . . . . . . . . . . . . . . . . . . . . . .
6.9 Cross sectional view of cylindrical model . . . . . . . . . . . .
6.10 Relation between the transmission factor and the electron column density of O resonance lines in cylindrical model . . . . .
6.11 Relation between the transmission factor and the electron column density of Ne resonance lines in cylindrical model . . . .
6.12 Cross sectional view of barrel-like model . . . . . . . . . . . .
6.13 Relation between the transmission factor and the electron column density of O resonance lines in barrel-like model . . . . .
6.14 Relation between the transmission factor and the electron column density of Ne resonance lines in barrel-like model . . . .
6.15 Cross sectional view of double ring structure . . . . . . . . . .
6.16 Cross sectional view of a cylindrical structure and the velocity
structure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
6.17 Line-of-sight distance dependence of RS cross section . . . . .
6.18 Relationship between cross section and line-of-sight distance .
7.1
7.2
7.3
7.4
7.5
7.6
7.7

102
104
105
105
106
108
111
113
117
118
119
121
122
123
124
126
127
128

Total cross sections of CX, as functions of collision velocity . . 133
Relationship between Fe/O abundance ratio and progenitor
mass of massive stars . . . . . . . . . . . . . . . . . . . . . . . . 135
X-ray image of W49B observed with Chandra . . . . . . . . . . 137
Simulated distribution of heavy elemant in jet-driven core-collapse
SN and SNR . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 137
X-Ray image and schematic view of W49B . . . . . . . . . . . . 138
Simulated XRISM spectrum of W49B . . . . . . . . . . . . . . . 139
Close-up views of the RGS spectrum around Fe XVII Lα andO VIII Lyβ obtained by our studies and previous studies . . . 140

xv

List of Tables
2.1

Explosive nucleosynthesis . . . . . . . . . . . . . . . . . . . . .

11

3.1

Basic information for O VII and Fe XXV Heα lines . . . . . . .

47

4.1
4.2

Design and performance of MOS . . . . . . . . . . . . . . . . .
Performance of the RGS . . . . . . . . . . . . . . . . . . . . . .

61
68

5.1
5.2
5.3

Best fit parameters of N49 . . . . . . . . . . . . . . . . . . . . .
Best fit parameters of J0453.6−6829 . . . . . . . . . . . . . . . .
f /r ratios or G-ratios for OVII Heα . . . . . . . . . . . . . . . .

79
90
92

6.1
6.2
6.3
6.4
6.5
6.6
6.7

Information on the XMM-Newton observations of E0102 . . .
Detected Emission Lines and their Identification . . . . . . . .
E0102 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
Observed line fluxes of E0102 . . . . . . . . . . . . . . . . . . .
Obtained kTe and ne t, and line emissivity ratios . . . . . . . .
Transmission factors of resonance lines . . . . . . . . . . . . . .
Optical depth of O VII Heα (r ) . . . . . . . . . . . . . . . . . .

107
109
114
114
115
115
125

7.1
7.2
7.3

Summary of previous and our RGS stidies . . . . . . . . . . . . 131
Oscillator Strength of Fe XVII Lα lines . . . . . . . . . . . . . . 134
Optical depth of W49B . . . . . . . . . . . . . . . . . . . . . . . 136

1

Chapter 1

Introduction
A supernova (SN) is an explosion at the end of a star’s life. Heavy elements synthesized in the interior of a star are released into interstellar spaceby an SN explosion. Ejected material from an SN collides with surrounding gases and forms a high-temperature plasma called supernova remnant(SNR). Since SNR plasmas are often optically thin in the X-ray band, X-ray
imaging spectroscopy enables us to directly estimate the abundances andspatial distribution of heavy elements. This information provide us withmany insights into the evolution process and explosion mechanism of theprogenitors (e.g., Sato et al., 2021).
Recent studies pointed out that the formation process and X-ray emission
process of SNRs are more complex than those expected in the conventional
picture assuming spherically symmetric explosions in uniform environments(e.g., Okon et al., 2020; Yamaguchi et al., 2018; Uchida et al., 2019). Such diversity in the physics of SNRs has been revealed by improvements in theperformance of X-ray detectors. Figure 1.1 show the improvement historyof energy resolution of X-ray instruments and important findings, which resulted from the spectral study. CCD detectors, which are widely used in
X-ray astronomy, can separate the major K-shell emission lines from heavy
elements with different ionization states. On the other hand, X-ray grating
spectrometers and future available X-ray calorimeters have an energy resolution one order of magnitude better than CCDs. Such energy resolution
enables us to measure the kinetic velocities of gases and the abundances ofrare elements, as well as to investigate physical processes such as charge exchange (CX) due to collisions between ions and neutral atoms, and scatteringof line photons by ions (resonance scattering: RS).
In this thesis, we perform diagnostics for physical processes such as CX
and RS in SNRs with the grating spectrometer. Although these processes are
predicted to occur in SNRs by previous calculations (e.g., Lallement, 2004;
Kaastra and Mewe, 1995), established observational evidence is still limited.
This is because it is difficult to apply the grating spectrometer to diffuse objects such as SNRs. Since the effect of these processes apparently reduce orenhance intensities of some lines, ignoring their contribution can sometimes
lead to, for example, biases in elemental abundance measurements. On theother hand, quantifying their contribution will allow us to obtain several important information such as microturbulence velocities, absolute abundancesand 3D structure of SNRs.

2

Chapter 1. Introduction

F IGURE 1.1: History of energy resolution of X-ray spectrometers in X-ray astronomy (Ezoe et al., 2021). The energy resolution is FWHM in energy for 6 keV X-rays.

In Chapter 2, we describe the reviews of SN and SNR, respectively. The
basic characteristics of the instruments aboard XMM-Newton which we usedin this study are summarized in Chapter 4. Chapter 5 provide the results of
the high resolution X-ray spectroscopy of SNRs located in Large Magellanic
Cloud (LMC): N49 and J0453−6829. In Chapter 6, we introduce a method toconstrain the 3D structure of SNR plasmas using the effect of RS. Combining all the works, we develop the discussions in Chapter 7, and we finallydescribe the summary and conclusion of this thesis in Chapter 8.

3Chapter 2

Supernova and Supernova
Remnant
A supernova (SN) is an explosion at the end of a star’s life. It is one of themost dramatic explosion phenomena in the universe, with a brightness comparable to that of a single galaxy and a released kinetic energy of ∼ 1051 erg.
SN explosions are the final stage of stellar evolution and the origin of compact objects such as black holes and neutron stars. In addition, heavy elements synthesized in the interior of stars during their evolution and explosions are released into interstellar space by SN explosions. This diversity
makes SNs one of the most interesting phenomena in the universe. In thischapter, we explain the evolutionary process leading up to the explosion ofa star, the classification and explosion mechanism of SN, and the supernova
remnant (SNR).

2.1

Stellar Evolution of the Progenitor

2.1.1

Stellar Evolution

Stars are classified according to the observed temperature (spectrum type)and luminosity. Figure 2.1 shows an HR diagram showing the types of stars.
Stars produce energy by nuclear fusion reactions in their cores, creating enoughpressure to support their self-gravity. Nuclear fusion reactions inside stars
begin with nuclear reactions of hydrogen, gradually synthesizing heavier elements. A star in the early stages of stellar evolution that is burning hydrogenis called a main sequence star (for example, the B-type star indicated by the
red circle in Figure 2.1). The temperature and brightness of main sequencestars depend primarily on their mass. When a star exhausts the hydrogen inits core, hydrogen nuclear reactions begin in a shell outside the core, then theradius of the star expands and the surface temperature drops. As the radius
expands and the surface temperature drops, the star moves to the right on
the HR diagram and becomes a giant star such as a red supergiant (RSG). A
Star with an initial mass of less than 25 M⊙ are considered to be an SN explosion in the giant star phase. Massive stars with initial mass exceeds 25 M⊙
move to the left again in the HR diagram due to their large amount of mass
loss. Such Massive stars considered to explode after becoming a Wolf-Rayetstar (WR star).

4

Chapter 2. Supernova and Supernova Remnant

F IGURE 2.1: Hertzsprung Russell diagram for massive stars
in the SMC (Millette, 2021) based on the observation of Ramachandran et al. (2019). The brown pentagons represent WRstars (encircled if in a binary system), yellow symbols represent
yellow supergiant (YSG) stars, blue triangles for blue supergiant stars (BSG), and red triangles for red supergiant (RSG).
Black tracks show standard stellar evolutionary paths, while
the blue tracks show the paths of WR stars.

2.1. Stellar Evolution of the Progenitor

5

F IGURE 2.2: p–p chain reaction (taken from astronomical dictionary)

2.1.2

Stellar Nucleosynthesis

The stars contract quasi-statically due to self-gravity while maintaining thebalance between the pressure gradient force and gravity. Part of the released
gravitational energy is released as photons and the rest becomes internal energy. Gradually, the interior of the star becomes hot and dense, and the nuclear fusion reaction of hydrogen begins.
Hydrogen Burning The phenomenon of synthesizing helium nuclei fromfour hydrogen nuclei (protons) is called hydrogen burning. Hydrogen burning occurs in the core of main-sequence stars with temperatures of ∼ 107 K.
Hydrogen burning includes the p–p chain reaction, which is dominant in
main-sequence stars of solar mass or less, and the CNO cycle, which is dominant in stars that are more than 1.3 times as massive as the Sun.
In the p–p chain reaction, two protons first become deuterium, and thedeuterium reacts with a proton to be 3 He. This reaction can be described as
p+p →
2
H+p →

H + e+ + ν
3
He + γ.
2

(2.1)
(2.2)

The p–p chain reaction is classified into ppI, ppII and ppIII depending on thesubsequent reactions. A summery of these reactions is shown in Figure 2.2.
In the CNO cycle, protons are successively captured by carbon, nitrogen,and oxygen nuclei and become helium nuclei. The series of reactions shownin the following Equations (2.3), (2.4), (2.5), (2.6), (2.7) and (2.8) is the main
reaction of the CNO cycle called the CN cycle, and CNO are used as catalysts
here.
12

C+p →
13

N →

13

N

(2.3)

13

C + e+ + γ

(2.4)

6

Chapter 2. Supernova and Supernova Remnant

F IGURE 2.3: CNO cycle (taken from astronomical dictionary)
13
14

C+p →

14

N+γ

(2.5)

N+p →

15

O+γ

(2.6)

15

+

(2.7)
(2.8)

15

15

O →
N+p →

12

N+e +ν
C + 4 He

The reaction of Equation (2.6) has the following branching reaction.
15

N+p →

15

O+γ

(2.9)

This reaction occurs at a rate of 10−4 . In the CNO cycle, there is also reactions
called the NO cycle that starts from this reaction. Figure 2.3 summarizes
these reactions.
Helium Burning and Fusion of Heavier elements When the central temperature reaches ∼ 108 K, helium is ignited via the triple-alpha process. In
the triple-alpha process, three helium nuclei (α particles) fuse to be a carbon
nucleus. This process occurs through the two-step reaction shown in Equations (2.10) and (2.11) below.
4

He + 4 He ⇀↽8
4
Be + He →

8

Be
C+γ

12

(2.10)(2.11)

Since 8 Be is unstable, it decays to 4 He in ∼ 10−16 seconds. Therefore, thereaction described as Equation (2.10) is in an equilibrium state, and 12 C is

2.1. Stellar Evolution of the Progenitor

7

F IGURE 2.4: Onion-like structure in massive stars. (The scale is
not exact.)

synthesized from the fusion of 4 He with a small amount of 8 Be. The 12 C
produced here reacts with helium to produce 16 O.
12

C + 4 He →

16

O+γ

(2.12)For stars with a mass of ≤ 8 M⊙ , the nuclear fusion reaction does not proceed
any further and eventually evolves into a white dwarf with a core of carbon
and oxygen.
Inside stars with a mass of ≥ 8 M⊙ , nuclear fusion reactions of heavier
elements take place. When the central temperature reaches ∼ 6 × 108 K,
carbon burning occurs, and the carbon nuclei become mainly oxygen, neon,and magnesium. For stars with a mass of ≥ 10 M⊙ , oxygen burning subsequently takes place at a central temperature above ∼ 2 × 109 K, producing
Si, S, and Ca. At temperatures above ∼ 3 × 109 K, high-energy photons
break silicon nuclei, and Fe is synthesized by silicon burning. As a result of
these reactions, the star has an onion-like structure as shown in Figure 2.4, at
the final stage of its evolution.

2.1.3

Mass Loss of Massive Stars and Their Environment

Massive stars lose a large amount of their mass before their explosion as SNe.
Mass loss from massive stars modifies the surrounding medium, forming a
dense and inhomogeneous circumstellar medium (CSM). The resulting CSM
subsequently affects the evolution of SNRs. Here, for single stars and binarystars, we describe the environmental structure resulting from the mass lossof massive stars during their evolutionary process.
Single Massive Stars
Since massive stars have high luminosity, their envelope is removed as a stellar wind due to high radiation pressure. The mass loss rate and the velocity ofa stellar wind differ depending on their evolutionary stage. Here, we explain
the evolutionary process of stars and the surrounding environment createdby stellar winds of each evolutionary stage. Main sequence stars evolve into
various stars, such as red supergiants or Wolf-Rayet stars, depending on theirinitial mass and metallicity (Figure 2.1), where metallicity is the abundance
ratio of heavy elements to hydrogen. Figure 2.5 shows the mass loss rate and8Chapter 2. Supernova and Supernova Remnant

F IGURE 2.5: Evolution of the velocity (left) and mass-loss rate(right) of the wind from the star with time (Dwarkadas et al.,2007).

stellar wind velocity at each evolutionary stage of a star with an initial massof 35 M⊙ . This star follows a sequence of evolution from a main sequence star
to a red supergiant and finally to a Wolf-Rayet star. Main sequence stars havehigh stellar wind velocity and low mass loss rate. The stellar wind of a main
sequence star expands while sweeping up the surrounding ISM, and forms
an ISM shell and main sequence wind bubble. At the RSG phase, contrary to
the main sequence star phase, stellar wind velocity is low and the mass loss
rate is high. Therefore, the RSG stellar wind forms a dense region inside themain-sequence wind bubble. On the other hand, a WR star has high stellar
wind velocity and a high mass loss rate. Consequently, the WR star wind creates a cavity wall inside the CSM produced by the RSG wind. As a result, thefinal stage of a single massive star’s evolution results in an ambient medium
with a radially non-uniform density profile, as demonstrated in Figure 2.6.
Binary Stars
Another important mass loss mechanism is "binary interaction". About halfof the stars are found in binary systems in which two or more stars revolvearound each other. When two stars evolve in a binary system that is close
to each other, they mutually change their masses and orbits by exchanging
mass and momentum. This process is called "binary interaction". Since it has
been reported that almost all O-type stars in the main sequence stage are inbinary systems (Sana et al., 2014), binary interaction is an important process
for massive stars.
In a binary star system consisting of two stars, the heavier star is called
the primary star, and the lighter one is called the companion star. Figure 2.7(a) shows the gravitational potential created by a binary star in a coordinate
system that revolves with the binary star. Figure 2.7 (b)–(d) show a crosssectional view of the potential when the axis passing through the two stars
is taken as the x-axis. As stars evolve, there are stages in which the stellar
radius expands significantly (such as a RSG). At this time, when the volume2.1. Stellar Evolution of the Progenitor

F IGURE 2.6: Density plots of the 2D evolution of the bubble
during the WR phase (Dwarkadas et al., 2007).

9

10

Chapter 2. Supernova and Supernova Remnant

F IGURE 2.7: left: (a) Dimensionless Roche potential contours
on the plane of orbit in a binary system with a mass ratio of
2. right: Roche potential and Roche lobe overflow (RLOF) diagrams with the x-axis passing through the two stars on theorbital plane. (b) Neither of the two stars exceeds their own
gravitational potential, and no mass transfer occurs. (c) Masstransfer from L1 to the companion star (stable RLOF) as theheavier star expands and exceeds its own gravitational potential. (d) Binary systems have a common envelope.

of the primary star expands beyond its own gravitational potential, its envelope is transferred to the companion star through the L1 point in Figure
2.7 (a). When the mass transfer is relatively slow and the envelope is stablytransferred to the companion star, it is called a stable Roche-Lobe overflow(left panel of Figure 2.8). On the other hand, when the mass transfer is fast,
the companion star also exceeds its own gravitational potential as shown in
Figure 2.7 (d). At this time, the companion star is incorporated into the primary star’s envelope, which is called the common envelope. Binaries thatform a common envelope undergo mass loss due to stellar winds in the envelope which has a structure extending in the equatorial direction, as shownin the right panel of Figure 2.8. A supernova explosion in such a structure
may form a peculiarly shaped SNR, such as that seen in 1987A (see Subsection 2.3.2).

2.2. Supernova

11

F IGURE 2.8: Schematic view of mass transfer in binary stars

2.2

Supernova

2.2.1

Explosive Nucleosynthesis

Neucleo fusion take place not only in a stellar evolution stage but also in an
SN due to explosive burning by the shock wave. The elements synthesizedby explosive nuclear fusion depend on the temperature and density. The
relation between dominant nucleosynthesis processes and temperature aresummarized in table 2.1.
TABLE 2.1: Explosive nucleosynthesis (Thielemann et al., 1996)

Burning Site
C/Ne burning
O burning
Incomplete Si burning
Complete Si burning

2.2.2

Main ProductsO, Mg, Si, Ne
O, Si, S, Ar, Ca
Si, S, Fe, Ar, Ca
Fe, HeIgnition temperature [109 K]
2.1
3.3
4
5

Classification of Supernova

SNs are classified into several types based on the optical spectra near maximum light. SNs with absorption lines of hydrogen Balmer series in theirspectra are classified as type II SNs and ones that do not are called type I
SNs. Type I SN is further divided based on the presence or absence of silicon(Si) and helium (He) in their spectra. Type Ia SN contain a Si absorption at
6150 Å, Type Ib has no Si but shows He as an emission signature, and TypeIc has neither. Type II SNs are classified into IIP and IIL types by the shapeof their light curves. The main feature of IIP SNs is their plateau-like light
curves, while the light curve of SNe IIL is a nearly linear decline. Figure 2.10
shows the optical spectra of type Ia, Ib, Ic, and type II SNs, and Figure 2.11
shows the SN classification. The above classification is not based on the explosion mechanism of SN. SNs are classified into Type Ia SN (Thermonuclear
supernova) and Core-Collapse SN based on their explosion mechanism.

12

Chapter 2. Supernova and Supernova Remnant

F IGURE 2.9: Optical spectra of SN (Filippenko, 1997).

F IGURE 2.10: Light curves of Type Ia, Ib, II-P, and II-L supernovae (Filippenko, 1997).

2.2. Supernova

13

F IGURE 2.11: Classification of SNs (Vink, 2012).

2.2.3

Type Ia SNe

Type Ia supernovae are characterized by no hydrogen absorption lines, andexponentially decaying light curves, and found in old galaxies such as elliptical galaxies. Based on these features, Type Ia SN is thought to be the resultof the explosion of a carbon-oxygen white dwarf (e.g., Nomoto, 1982). Whena white dwarf reaches Chandrasekhar mass (∼ 1.4 M⊙ ), it cannot support its
own gravity by the degeneracy pressure of electrons. As a result, the whitedwarf contracts, the density and temperature of the core increase, and thecarbon burning is ignited at the core. These fusion reactions proceed explosively and the white dwarf explodes with the energy of the nuclear reaction
released. Since the progenitor is completely distructed by the explosion, no
compact objects such as neutron stars or black holes are left at the center. Two
scenarios have been proposed for how white dwarfs reach Chandrasekhar
mass. One is the accretion of mass from a companion star to a white dwarf
in a binary system with a main sequence or a red giant (Single Degeneratemodel). The other is the merger of two white dwarfs in a binary star system (Double Degenerate model). Recent study suggests that both scenarios
are required to explain the elemental abundance of iron peak elements the
Perseus cluster (Collaboration, 2017).

2.2.4

Core-collapse SNe

Explosion Mechanism As described in Subsection 2.1.1, massive stars havea Fe core at the final stage of their evolution. Since Fe is a quite stable element, the core can no longer generate nuclear energy. Therefore, the density
and temperature increase due to the subsequent contraction. When the temperature of the core exceeds ∼ 3 × 109 K photo-disintegration of Fe nuclei

14

Chapter 2. Supernova and Supernova Remnant

F IGURE 2.12: (left) Radial positions of shock waves as afunction of time after bounce by one-dimentional simulation(Sumiyoshi et al., 2005). (right) Mean shock radii of 2D models(Burrows and Vartanyan, 2021). The different colors of the lines
represent the mass of the progenitor.

takes place:56

Fe → 134 He + 4n − 124.4 MeV
4He → 2p + 2n − 28.3 MeV.

(2.13)(2.14)

Since these processes are endothermic reaction, the pressure of the core decreases and collapses gravitationally. Electron capture takes place in thedense core, and a proto-neutron star is formed.
p+ + e− → n + νe

(2.15)

Falling material bounces at the surface of the proto-neutron star, forming an
outward shock wave. This shock wave is called "bounce shock". When the
bounce shock reaches the surface of the star, it is observed as an SN.
The bounce shock generally does not reach the surface of the star since
the shock stalls by the falling outer material. To revive bounce shocks, an additional mechanism to heat the shock wave is required. A promising process
is neutrino heating, but stars in most mas range do not explode in spherically symmetric (one-dimensional) simulation (the left panel of Figure 2.12).
Multi-dimensional effects are also the key of the explosion mechanism.
Neutrinos generated in the center of SN heat the inner part of the SN
efficiently. The resulting temperature gradient (more precisely, the entropy
gradient) creates convection inside the shock wave. Such convection may efficiently carry the energy due to neutrino heating to the location just behindthe shock, and the pressure from the vortices generated by the convection
also revives the shock wave (e.g., Kifonidis et al., 2003). Meanwhile, convection is not the only promising hydrodynamical instability to revive theshock wave. The standing accretion shock has been recognized to be generically unstable to non-radial deformation. In the stagnant shock front, theperturbation is amplified by the back and forth of sound waves and vortices

2.2. Supernova

15

between the protoneutron star and the shock wave. This causes the shock
wave to distort and vibrate. As a result, the shock wave becomes stronger
in the vibrating direction, and the shock wave revives in that direction (e.g.,
Hanke et al., 2013). Magnetic fields and rotation are also believed to be important asymmetric effects. The magnetic field of the PNS is amplified in thepolar direction by the fast rotation (e.g., Takiwaki and Kotake, 2011). Thiseffect causes a jet-like explosion (e.g., Khokhlov et al., 1999). It is promising
as a hypernova explosion mechanism because it reproduces high explosion
energy reaching 1052 erg.
Nucleosysthesis The maximum temperature due to shock wave heating
can be estimated as follows. We assume that the inside of the shock wave isfilled with uniform radiation and that its energy accounts for most of the explosion energy. At this time, the relationship between the temperature (Tmax )and the radius of the shock wave (rsh ) is described as:


4π 3
4
r
= Eexp ,(2.16)
aTmax
3 sh
where Eexp is the explosion energy, a = 7.57 × 10−15 ergs cm−3 K−4 . Therefore, Tmax is
10

Tmax = 1.33 × 10 K

Eexp
1051 erg

1/4 

rsh −3/4.
108 cm(2.17)

Therefore, the maximum temperature, which decides the synthesized elements, is lower for the outer layer. Outside the radius of 13000 km, the maximum temperature is below T ∼ 2 × 109 K, so explosive nucleosynthesis does
not occur.
The abundances of the ejected material from SN consists of the elements
synthesized by the explosive and stellar nucleosynthesis. The yields of theseelements depend sensitively on physical parameter of the progenitor and thedetails of the explosion mechanism, such as the progenitor mass. Overall theyields of core collapse supernovae are dominated by carbon, oxygen, neonand magnesium, which are products of the various stellar burning phases.
The yields of these elements reflect the physical parameter of the progenitor,
such as mass and metallicity (Figure 2.13).

16

Chapter 2. Supernova and Supernova Remnant

F IGURE 2.13: Relationship between the progenitor mass andabundance. Left: Abundance ratios (Fe/Si) as a function of the
CO core mass (Katsuda et al., 2018), based on the nucleosynthesis models by Sukhbold et al., 2016. The abundances are
relative to the solar values (Anders and Grevesse, 1989). Right:
Oxygen yield of core collapse supernovae as a function of main
sequence mass (Vink, 2012). The circles and squares are thepredictions of Woosley and Weaver, 1995, the triangles are predictions of Chieffi and Limongi, 2004, and the crosses of Thielemann et al., 1996. In general the oxygen yields obtained by thevarious groups are very similar, but above 30 M⊙ one sees thatcertain models predict a diminishing oxygen yield. The reason
is that above30 M⊙ stellar cores may collapse into black holes,
and part of the oxygen falls onto the black hole. The amount offall-back is governed by the explosion energy and the amountof pre-supernova mass loss, but it is also sensitive to the numerical treatment of the explosion.

2.3. Supernova remnant

17

F IGURE 2.14: Schematic image of an SNR.

2.3

Supernova remnant

A supernova remnant (SNR) is an object resulting from a supernova. The
SN ejecta expands in the interstellar medium (ISM). The expansion velocity(∼ 104 km s−1 ) is much larger than the sound speed in the ISM (∼ 10 km s−1 ),and hence shock wave (blast wave) is generated. The blast wave sweepups the surrounding ISM and ionizes them. As the mass of the swept-up
ISM increases, the decelerated outer-layer ejecta collides with the inner-layer
ejecta, forming an inward shock wave called the reverse shock. As a result,
the SNR have the structure shown in Figure 2.14. In this chapter, we explain
the overview of physics concerning SNRs: the heating process performed bytheir shock wave, the macroscopic evolution of SNRs.

2.3.1

Basic Physics of Supernova Remnants

Shock Wave Heating
The relationships between the properties of the upstream and downstream
fluids are derived by the conservations of some fluid quantities across theshock. We take a coordinate to be the rest frame of the shock, and define
upstream and downstream, as shown in Figure 2.15. The conservations of themass, momentum, and energy (per unit mass) across the shock are expressedas:ρu vu = ρd vd(2.18)

18

Chapter 2. Supernova and Supernova Remnant

F IGURE 2.15: Schematic view of the shock front in observerrest frame (left) and shock rest frame (right). ρ, v, and P are thedensity, velocity, and pressure, respectively.

ρu v2u + Pu = ρd v2d + Pu
1 2
γ Pu
1
γ Pd
vu +
= v2d +,
2
1 − γ ρu2
1 − γ ρd(2.19)(2.20)

where ρ, v, P, are the density, velocity, and pressure of upstream (subscriptu) and downstream (subscript d), γ is the specific heat ratio (γ = Cp /Cv ).
The Rankine-Hugoniot equations are obtained as:( γ + 1) ρ d − ( γ − 1) ρ uPu=Pd( γ + 1) ρ u − ( γ − 1) ρ dvuρ(γ − 1) Pu + (γ + 1) Pd
χ≡ d ==,
ρuvd(γ + 1) Pu − (γ − 1) Pd(2.21)(2.22)

where χ is the compression ratio. Here we introduce the Mach number of the
upstream fluid (Mu = vu /csu ), where csu is the sound speed in the upstream
fluid. Equations (2.21) and (2.22) can also be expressed as:Pu
2γMu2 − γ + 1
=
Pd
γ+1
ρ
vu
(γ + 1) Mu2
χ= d ==.
ρuvd(γ − 1) Mu2 + 2(2.23)(2.24)For monoatomic gas (γ = 5/3), assuming the strong shock (ρd /ρu ≫ 1),thecompression ratio χ is 4. Combining Equations (2.21) and (2.22) with theideal gas law, we obtain the downstream temperature to bekTd =2( γ − 1)
3µm p v2u =µm p v2u ,216( γ + 1)(2.25)

which corresponds to ∼ 1.2 keV (∼ 1.4 × 107 K) for a shock velocity of1000 km s−1 and the solar composition ISM (µ ≈ 0.6). We can see in Equation (2.25), the shock heating can be understood as a process that converts acertain fraction of the kinetic energy to internal (thermal) energy.

2.3. Supernova remnant

19

Evolution of Supernova Remnant
The evolution phase of SNRs can be classified into the following four stages
based on the relationship between the mass of the ejecta, Mejecta , and themass of the swept-up ISM, MISM .
1. Free expansion phase
2. Adiabatic (Sedov) phase
3. Radiative cooling phase
4. Disappearance phase
Free expansion phase At the early stage of SNRs evolution, MISM is muchsmaller than Mejecta . Therefore, the ejecta expands freely. Most of the kinetic
energy of the explosion (E) becomes the kinetic energy of the ejecta (Eejecta ).
The velocity (vs ) and the radius (Rs ) of the blast wave are described as:s1/2 

Mejecta 1/2
2EE
9[cm s−1 ] (2.26)= 1.0 × 10
vs =MejectaM⊙1051 erg
Rs = vs t,(2.27)

where t is the elapsed time since explosion. MISM become
MISM =

4
πR3s µm H n0 ,3(2.28)

where n0 is hydrogen number density of the ISM. From Equations (2.26),(2.27) and (2.28), the timescale for MISM ∼ Mejecta is2

t ∼ 1.9 × 10

−1/2 E
1051 erg

EejectaM⊙

5/6 

µ −1/3  n0 −1/3[yr(2.29)].
1.4
1 cm−3

Therefore, this phase typically continues for a few hundred years.
Adiabatic (Sedov) phase When MISM ≥ Mejecta , the eject is decelerated by
the swept ISM. The SNR expands adiabatically since the energy loss due toradiation is negligible. The Adiabatic stage is also called the Sedov stage,
since the exact solution for spherically symmetric shock waves is known as
the Sedov-Taylar solution. The shock wave radius (Rs ), velocity (vs ), andtemperature downstream of the shock wave (Ts ) are2/5 

1/5 

n0 −1/5[cm] (2.30)
1 cm−3104 yr1051 erg−3/5 
1/5 
tEn0 −1/5
7
vs = 5 × 10
[cm s−1 ] (2.31)
1 cm−3104 yr1051 erg
19

Rs = 4 × 10



t

E

20

Chapter 2. Supernova and Supernova Remnant

6

Ts = 3 × 10t
104 yr−6/5 E
1051 erg

2/5 

n0 −2/5[K]. (2.32)
1 cm−3The Sedov phase continues for tens of thousands of years and converts about
70% of the kinetic energy of the explosion into thermal energy (Chevalier,1974).
Radiative cooling phase At this stage, radiative cooling is non-negligible.
The dense shell cools faster, but the inner region is still hot and under highpressure. Therefore, the outer shell is pushed by the internal gas pressure
P ∝ V −γ where V is the volume of the internal gas. This phase is also called"pressure-driven snowplow" phase. Assuming a non-relativistic monatomicgas (γ = 5/3), the shock wave expands according to
Rs ∝ t2/7 .

(2.33)As the temperature decreases further, the internal pressure is also negligible and the shell expands while conserving momentum (Ms vs ). The shock
wave radius can be described as
Rs ∝ t1/4 .

(2.34)

This stage is also called "the momentum-conserving snowplow" stage.
Disappearance phase When the shock wave velocity equals the velocity ofthe proper motion of the ISM (10–20 km s−1 ), the SNR loses its boundaryand vanishes. This stage is reached on a timescale of ∼ 105 years from the
explosion.

2.3.2

Observational View of Supernova Remnants

In the previous section, we explained the evolution of SNRs by sphericallysymmetric explosions in a homogeneous environment. However, as shownin Figure 2.16, many supernova remnants have spherically asymmetrical morphologies. These morphologies may reflect the asymmetric effects of the explosion and the effects of non-uniform interactions with the surrounding environment, as explained in Section 2. Here we explain the information thatcan be obtained from such morphologies of SNRs.
Lopez et al., 2011 claims that it is possible to distinguish between typeIa SNR and CC SNR based on their X-ray morphologies. Figure 2.17 shows
P2 /P0 versus P3 /P0 plot for SNRs in the Milky Way and LMC where P2 /P0
and P3 /P0 are parameters that characterize the morphologies of SNRs. P2 /P0
is the quadrupole ratio. SNRs that have high P2 /P0 are those with elliptical/elongated morphologies or those with off-center centroids because one
side is substantially brighter than the other. P3 /P0 is the octupole ratio. Examples that have high P3 /P0 are those that have asymmetric or non-uniform

2.3. Supernova remnant

F IGURE 2.16: Chandra X-ray (0.5–2.1 keV) images of the 24
SNRs (Lopez et al., 2011). The cyan circles mark the full-band
centroids of each SNR determined by Lopez et al. (2011). Red
numbers denote Type Ia SNRs; light blue numbers denote CC
SNRs.

21

22

Chapter 2. Supernova and Supernova Remnant

F IGURE 2.17: Power ratios, the quadrupole ratio P2 /P0 vs. theoctupole ratio P3 /P0 , of the soft X-ray band (0.5–2.1 keV) for
SNRs in the Milky Way and LMC. Type Ia SNRs are in red,
CC SNRs are in blue, and 0548−70.4 is in purple because of its
anomalous ejecta properties that make its type uncertain. The
quadrupole ratio is a measure of ellipticity/elongation, and theoctupole ratio quantifies the mirror asymmetry of the emission.

surface brightness distributions. CC SNRs tend to have a non-spherical symmetry morphologies compared to Type Ia. This is because the progenitor of
the CC SNR is a massive star, and non-uniform CSM tends to form around it
as we explained in Subsection 2.1.3.
SN 1987A was first detected on February 23, 1987 in the Large Magellanic
Cloud. Pre-supernova image shows that the progenitor of SN 1987A is a blue
super giant (e.g., White and Malin, 1987). The X-ray emission of SN 1987A
has a ring-like structure. The optical observation reveals that there are tworings to the north and south of the X-ray ring. These optical rings are inapproximate alignment with the symmetry axis of the inner ring as shownin the right panel of Figure 2.18. The origin of this triple ring nebula is acontroversial, but as mentioned in Podsiadlowski, 2017, binary interactions
may play an important role. The progenitor of 1987A may be a binary starthat evolved in a common envelope. These stars merges to form a blue super
giant. The fast blue super giant wind sweep up all the slowly expanding
structures that had been ejected previously and produce two dense physicalrings at the end of two bipolar caps. This formation scenario for SN1987A issummarized in Figure 2.19.
Convection generated by neutrino heating forms high-entropy bubbles(the left panel of Figure 2.20). Some of the high-entropy bubbles can develop into Rayleigh-Taylor instabilities on a larger scale at late stages of the2.3. Supernova remnant

23

F IGURE 2.18: (right) X-ray and optical image of SN1987A observed with Chandra and HST (taken from Chandra photo album). The purple ring corresponds to X-rays. There are two
optical rings to the north and south of the X-ray ring. (left)
Schematic view of geometrical model for the triple-ring nebula (Podsiadlowski, 2017).

explosion and penetrate into the hydrogen envelope with large velocities creating iron-dominated plumes. The SNR Cassiopeia A has an Fe-protruding
structure that suggests this picture (the right panel of Figure 2.20). Based on
the Ti/Fe and Cr/Fe ratio, Sato et al. (2021) revealed that this structure wassynthesized in the center of the SN. This result supports the picture of theconvective SN explosion mechanism.
As these studies show, the morphology of the SNR has important information about the progenitor and its explosion mechanism. Our objective is toobtain observational evidence for CX and RS from SNRs with high resolution
X-ray spectroscopy. This allow us to accurate measure the abundances accurately, as we will discuss later, and also help us to estimate the 3D structureof the SNRs.

24

Chapter 2. Supernova and Supernova Remnant

F IGURE 2.19: Schematic image of the formation of the triplering nebula (Podsiadlowski, 2017).

2.3. Supernova remnant

F IGURE 2.20: (left) 3D simulation of SN (Burrows et al., 2012).
Convection is induced by neutrino heating. (right) Distribution of heavy elements in Cassiopeia A (Sato et al., 2021). The
iron ejecta are shown in red. The Chandra ratio map of the
Si/Mg band is shown in green. The 44 Ti observed by NuSTAR
is shown in blue.

25

27

Chapter 3

Plasma Processes and X-ray
Radiation in Collisionally Ionized
Plasma3.1

Collisionally Ionized Plasma

As explained in Chapter 2.1, hot gas (∼ 1–5 keV) is produced by the shock
wave heating. Free electrons in such hot gases have enough energy to remove electrons bound in inner- or outer-shell of heavy elements in the gases(see Figure 3.1). As a consequence, heavy elements in SNRs are ionized via
free electron collision. Such plasma is called Collisionally Ionized Plasma(CIP). The properties of CIP can be investigated through spectral diagnostics of thermal X-ray emission, which is characterized by atomic emission
lines and continuum emission components. In this section, we explain theionization and recombination processes that take place in CIP and introduce
the concept of ionization equilibrium. We also summarize X-ray radiation
processes of CIP.

3.1.1

Ionization and Recombination

Ionization and Recombination Processes
In the CIP, four distinct ionization processes occur, each of which has a corresponding recombination process (reverse reaction), as summarized below.
Collisional ionization
Collisional excitation auto-ionization
Photo ionization
Photo excitationauto-ionization

⇀
↽
⇀
↽
⇀
↽
⇀
↽

Three-body recombination
Dielectronic recombination collisional deexcitation
Radiative recombination
Dielectronic recombination

Figure 3.2 shows a schematic image of each process. When the plasma density is enough high, these ionization and recombination processes balanceeach other. However, SNR plasmas have low density (∼ 1 cm−3 ) and are basically optically thin. Therefore these individual reactions are not balanced.
While there are exceptions, SNR plasmas generally satisfy the following conditions.

28

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
PlasmaF IGURE 3.1: Binding energy of atomic subshells for neutral
atoms (Kaastra et al., 2008).

F IGURE 3.2: Ionization and recombination processes in CIP.
Horizontal lines represent energy levels, and above n = ∞ isthe free electron energy region. Ionization processes are shownin red, and recombination processes are shown in blue. The
dashed lines indicate processes that are important in plasmasof SNRs.

3.1. Collisionally Ionized Plasma29

1. The plasma is optically thin. Therefore, photons emitted by the plasma
are not attenuated by self-absorption and do not excite or ionize atoms.
2. Due to the low plasma density, the number of excited-state atoms is
negligible compared to that of ground-state atoms.
As a result, photoionization and photoexcitation autoionization are negligible in SNR plasma. Additionally, three-body reactions (such as three-bodyrecombination and Dielectronic recombination collisional deexcitation) arealso negligible due to the low density. Below is an overview of the important
processes that take place in SNRs.
Collisional ionization Collisional ionization is a process in which free electrons interact with ions through inelastic collisions. When a bombarding freeelectron has higher kinetic energy than the binding energy of the target electron bound in the ion (A+q in Equation (3.2)), the bound electron can be removed. This process is described as:
A+q + e− → A+(q+1) + 2e− .

(3.1)

The collisional ionization cross section is described by the equation publishedby Lotz, 1968:
σ = 4.5 × 10−20

Ne log( E/I )
cm2 ,
EI

(3.2)

where E is a kinetic energy of an electron, I is the binding energy of theatomic shell, and Ne is the total number of electrons bound to the shell. The
cross section mainly depends on the energy of the free electron. The crosssection is zero when the kinetic energy of the free electron is lower than thebinding energy, but increases as the kinetic energy exceeds the binding energy. It reaches a maximum value when the free electron energy is roughly
three to four times the binding energy (e.g., Gillaspy, 2001).
Collisional excitation auto-ionization In the case where E < I, ionization
can still occur through a collisional excitation auto-ionization. This process
can be described as follows. An electron bound to a shell is excited to ahigher state by a free electron collision. The electron is then removed fromthe ion through an Auger transition. The resulting vacancy is then filled byanother electron from the outermost shell. In this case, energy conservation
requires that the excited electron originated from an inner shell. This process
is described as:
A+q + e− → A+q∗ + e− → A+(q+1) + 2e− ,(3.3)

where A+q means the excited state of the ion.
Radiative recombination Radiative recombination is a process where a freeelectron is captured by an ion and a photon is emitted simultaneously. This

30

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
Plasmaprocess is the inverse process of photo ionization described as:
A+(q+1) + e− → A+q + hν

(3.4)

The energy of the emitted photon is defined as:
Eγ = Ekin − Ebind .

(3.5)

Radiative recombination is also important process as an X-ray radiation process.
Dielectronic recombination Another recombination process is dielectronic
recombination. When an electron with a certain kinetic energy is captured inan ion, part of the kinetic energy of the electron is non-radiatively transfersed
to a bound electron, which cause an excitation of the bound electron. This
process is described as:
A+(q+1) + e− → A+q∗∗ → A+q + hν,(3.6)

where A+q∗∗ means the doubly excited state of the ion. When the excited ion
relaxes to an energetically lower state, a photon is emitted. This process is aresonance process, which tends to take place when the kinetic energy of thefree electrons is equal to the energy required to excite a bound electron.
Ionization Distribution
When the electron temperature in the SNR plasma becomes sufficiently high,atoms are gradually ionized through several processes described in Subsubsection 3.1.1.1. Finally, SNR plasmas reach an ionization equilibrium state inwhich ionization and recombination are balanced. Here, we describe the time
evolution of the ionization distribution in CIP and explain the concepts ofcollisional ionization equilibrium (CIE) and ionization non-equilibrium ionization (NEI). Then, we explain the standard picture of plasma evolution in
SNRs and the origin of recombining plasmas, which is an ionized state thatchallenges this standard picture.
Time evolution of the charge number distribution of ions (ionization distribution) in CIP is given as
1 dNZ,z= NZ,z−1 SZ,z−1 ( Te ) − NZ,z (SZ,z ( Te ) + α Z,z ( Te )) + NZ,z+1 α Z,z+1 ( Te ),(3.7)
ne dt
where ne is the number density of electron, NZ,z is the number density ofions with atomic number Z and charge number +z, SZ,z and α Z,z are rate coefficient of ionization and recombination, both of which depend solely on theelectron temperature (Te ). Therefore, ionization distribution of CIP is detremined by the electron temperature (Te ), the number density of electron (ne ),and time (t) which corresponds to the elapsed time from shock wave heating in the case of SNRs. From condition 2 in Subsubsection 3.1.1.1, only ions3.1. Collisionally Ionized Plasma

31in the ground state are included when Equation (3.2) is calculated. Therefore, Equation (3.2) is the time-dependent equation of the (Z+1)-dimensional
vector.
Collisional Ionization Equilibrium (CIE) The equilibrium state of ionization and recombination in CIP is called Collisional Ionization Equilibrium(CIE). The ionization distribution of each element in CIE plasma can be calculated by giving the following two conditions to Equation (3.7).
1 dNZ,z
=0
ne dt

(3.8)

z= Z

NZ = ∑ NZ,z

(3.9)

z =0

NZ is the number density of atoms with atomic number Z. The ionization
distribution of each ion is
Z k −1

SZ,j ( Te )
)]
α
(
T
)
e
Z,j
+
1
k =1 j =0(3.10)

SZ,j ( Te )) ( k ≥ 1).
α Z,j+1 ( Te )j =0

(3.11)

NZ,0 = NZ [1 + ∑ ∏ (k −1

NZ,k = NZ,0 ∏ (According to Equation (3.10) and (3.11), the ionization distribution of eachion can be determined by the electron temperature (Te ). Figure 3.3 shows therelation between the ionization distribution and electron temperature. The
electron temperature of a typical SNR is about 0.1 – 2.0 keV. As shown in theupper panel of Figure 3.3, relatively light elements, such as O, exist stably inan ionized state of He-like (i.e., the state with two bound electrons) with aclosed K-shell structure in CIE plasma with an electron temperature of 0.1–
2.0 keV. On the other hand, heavy elements, such as Fe, remain stably in a
Ne-like state (i.e., the state with 10 electrons) with a closed L-shell structure.
Non-Equilibrium Ionization(NEI) The state in which the ionization rate
and the recombination rate are not balanced is called non-equilibrium ionization (NEI). Furthermore, the NEI plasma is classified into the following
two types of plasma.
• Ionizing Plasma (IP)A plasma that is in a lower ionization state (closer to neutral state) thanthe CIE plasma. The ionization rate is higher than the recombination
rate, and ionization progresses toward the CIE state.
• Recombining Plasma (RP)Contrary to the IP, plasmas in a higher ionization state than the CIE
plasma in which recombination process is dominant.

32

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
PlasmaF IGURE 3.3: Ionization distribution of CIE plasma. The horizontal axis corresponds to the electron temperature. The vertical axis corresponds to ion fraction (NZ,z /NZ ). The upper panel
shows that of oxygen ions (Vink, 2012). Solid lines represent ion
fractions of O+8 (blue), O+7 (brown), O+6 (yellow-green), O+5(purple), and O+4 (cyan), respectively. The bottom panel showsthat of iron ions (ASTRO-H White paper). Colored solid lines
represent ion fractions of Fe+26 (yellow-green), Fe+25 (blue),
Fe+24 (red), Fe+16 (magenta), and Fe+8 (orange).

3.1. Collisionally Ionized Plasma

33

From Equation (3.7), the ion fraction of IP is a function of ne and t. Masai,
1994 computed the timescale for the plasma to reach the CIE from the neutral
state as:ne t ∼ 1012 cm−3 s.

(3.12)

Therefore, the t is12

t ∼ 10

1 cm−3nes ∼ 3 × 10

4

1 cm−3ne
yr.

(3.13)

SNR plasma is often IP because this timescale is comparable to or longer than
SNR ages.
The standard evolution scenario of SNR plasma has been established basedon the above theoretical calculation and observational studies. IP is observedin young SNRs such as Cas A and Tycho (e.g., Tsunemi et al., 1986), and CIE
is observed in relatively old SNRs (e.g., Sato et al., 2016). Therefore, in thestandard scenario of SNR plasma evolution, the SNR plasma is consideredto evolve from the initial low ionization state IP to CIE. Therefore, the RP
has not been considered in the standard scenario of SNR plasma evolution.
However, a recent studies using an X-ray CCD detector revealed that the
SNR plasma can be the RP (e.g., Kawasaki et al., 2002). To date, RP have been
discovered in more than 10 SNRs, and two theories have been proposed as
the origin of RP. One is rapid cooling due to adiabatic expansion of the SNR(Itoh and Masai, 1989), and the other is cooling due to thermal conduction
with surrounding dense gas (Kawasaki et al., 2002). A recent study found aclear correlation between the ionization state and electron temperature from
W49B (Yamaguchi et al., 2018) and W44 (Okon et al., 2020), confirming thatboth adiabatic expansion and thermal conduction are the origins of RP. In
both cases, RP is caused by rapid plasma cooling due to interaction with thedense inhomogeneous surrounding environment.

3.1.2

Atomic Structure and Radiation Processes

Atomic Structure
As can be seen from Figure 3.3 and 3.4, SNRs include ions with multiple
electrons in their orbits. Line emissions from SNRs is due to the orbital transitions of the electrons bound to these ions. Here we describe the atomic
structure and the selection rules for electron transitions, that help us derive
plasma properties from line emissions.
Energy Level of Single-electron Atoms The energy level of single-electron
atoms can be characterized by solving the Schrödinger equation of a single
electron in the central field created by an atomic nucleus. The resulting energy level is defined by the principal quantum number (n), the azimuthal
quantum number (l), and the magnetic quantum number (m). Here, l and m

34

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
PlasmaF IGURE 3.4: Ionization distribution of IP. The horizontal axis
corresponds to the ionization timescale (ne t). The vertical axis
corresponds to ion fraction (NZ,z /NZ ). The upper panel showsthat of oxygen ions (Vink, 2012). Solid lines represent ion fractions of O+8 (blue), O+7 (brown), O+6 (yellow-green), O+5 (purple), and O+4 (cyan), respectively. The bottom panel shows thatof iron ions (ASTRO-H White paper). Colored solid lines represent ion fractions of Fe+26 (yellow-green), Fe+25 (blue), Fe+24(red), Fe+16 (magenta), and Fe+8 (orange).

3.1. Collisionally Ionized Plasma

35

are the following values.
l = 0, 1, 2, . . . , n − 1
m = −l, −l + 1, . . . , 0, . . . , l − 1, l

(3.14)(3.15)

Additionally, the total angular momentum quantum number, j, is introduced,
taking into account the interaction between the electron’s spin and orbital
motion (known as spin-orbit interaction). j is the following value dependingon the direction of spin and angular momentum.

|l − s| ≤ j ≤ l + s(3.16)

Without a strong external magnetic field, the energy level of the electron
is determined by n, l and m. The principal quantum number, n, corresponds
to the distance between the atomic nucleus and the electron. Orbitals withn = 1, 2, 3, . . . are referred to as the K-shell, L-shell, M-shell, and so on. The
photon emitted by the electron transition from n = 2 to n = 1 is called theKα line, while the transition from n = 3 to n = 1 is known as the Kβ line.
Photons emitted by electron transition from n > 3 to n = 1 are called Kγ, Kδ,. . ., respectively.
Energy Level of Multi-electron Atoms In the case of multi-electron atoms,
as with single-electron atoms, the energy levels of individual electrons boundby the central-field created by atomic nuclei are defined by the principal
quantum number (n), the orbital quantum number (l), and the total angular momentum quantum number (j) (independent particle approximation).
Furthermore, the electron-electron Coulomb interaction has to be considered
to determine the energy levels of multi-electron atoms. Therefore, we introduce the total angular momentum J, which is the sum of the spins and orbital
angular momentum of all electrons.
For atoms with small atomic numbers (Z < 30) observed in X-ray astronomy, the spin-orbit interaction is weaker than the Coulomb interaction
between electrons. The energy levels of such atoms can be well describedunder the assumption of LS coupling (Russell-Saunders coupling). In the
LS coupling scheme, we first consider the total orbital angular momentum
L and the total spin angular momentum S expressed as Equation (3.17) and(3.18). The L and S are the sum of orbital angular momentum (l) and spin
angular momentum (s) of all electrons, respectively.
L = ∑ li

(3.17)

S = ∑ si

(3.18)

i

i

Multi-electron atoms have energy levels classified by the L and S. The total
angular momentum J is the sum of the total orbital angular momentum L

36

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
Plasmaand the total spin angular momentum S expressed as:
J = L+S(3.19)The energy difference between the energy levels due to the difference in J issmaller than that due to L and S (see Figure 3.9).
In the LS coupling scheme, using the quantum number of J, L, and S, thestate of the atom is expressed as:
2S+1

LJ ,

(3.20)

where L = 0, 1, 2, 3, 4, 5, 6, . . . are written as S, P, D, F, G, H, I, . . ., respectively.
Selection Rule Under the assumption of LS coupling, selection rules ofelectric dipole transition are described as below.
∆S = 0∆L = 0, ±1∆J = 0, ±1 (except for J = 0 → 0)

(3.21)(3.22)(3.23)

Transitions that satisfy the selection rules are called allowed transitions, andthose that do not are called forbidden transitions. Allowed transitions havehigh transition probabilities, and line emission due to allowed transitions isgenerally observed prominently in X-ray spectra of astronomical objects.
Radiation Processes
The X-ray spectrum of the thermal plasma in supernova remnants consistsof continuum emissions due to scattering and recombination of free electrons, and line emissions due to state transitions of atoms. Photo ionization/excitation processes are negligible since the supernova remnant plasma
is optically thin. Due to their low density, atoms in an excited state de-excite
through spontaneous emission before the next collision. Therefore, the processes summarized in Figure 3.5 are the main X-ray emission processes in
SNRs. Here, we outline each process. In particular, for line emission, we introduce emissivity, which is an important physical quantity that determines
the observed flux of line emission.
Line Emissionlk (pthoton cm3 s−1 ) in a spectral line k → l emitted by ions
The emissivity JZ,zof Z +z in an optically thin plasma can be written as:
lk
DR
JZ,z
= [SzII−1 ηz−1 + (SzE/IE + αzDRS )ηz + (αzRR+1 + α z +1 ) η z +1 ] ,(3.24)

where SzII−1 and SzE/IE are the rate coefficients (cm3 s−1 ) of the innershell ionization (II) of ion Z +(z−1) and excitation (E) or inner shell excitation (IE) of3.1. Collisionally Ionized Plasma

37

F IGURE 3.5: X-ray Radiation processes in SNRs. Horizontal
lines represent energy levels, and above n = ∞ is the free electron energy region.
DR
ion Z +z (Mewe et al., 1981). Furthermore, αzRR
+1 and αz+1 are those of the radiative or dielectronic recombination (RR or DR), respectively. η is the ion
fraction described as:

η=

NZ,z
,
NZ

(3.25)

which is calculated from Equation 3.7. As described below, each rate coefficient is a function of electron temperature. Therefore, the emissivities ofline emissions are the function of the electron temperature (kTe ) and the ion
fraction (ne t).
Inner Shell Ionization In typical SNRs, the relatively light elements suchas O are dominated by the He-like and H-like ionization state (Figure 3.3 (a)).
On the other hand, high-Z elements, such as Fe, are often less ionized thanthe He-like state (Figure 3.3 (b)). Therefore, in SNR plasmas, inner K-shellionization is induced by collision interactions between high-energy free electrons and low-ionized ions. The vacancy left after the inner shell ionization
is filled another electron in the outermost shell. Then, Line photons are emitted by this orbital transition of the electron. Using the formula of Lotz, 1968,
Mewe et al., 1981 introduces the rate coefficient SzII−1 as:
I
SzII−1 = 6.49 × 10−4 Te−1/2 C I I (χIIz−1 )−1 × E1 (χzI −1 /kTe ),(3.26)

38

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
PlasmaI is the ionization energy, and C is described as:
where χzI −II
1CI I =wjAr,
Σwj A a + ΣAr(3.27)

where Σwj is the total statistical weight of the configuration to which thelevel j belongs and, A a and Ar are transition probabilities for decay by autoionization and radiation, where the summation is over all possible radiativetransitions from level j. C I I is a constant depending on the ion species andthe transition as calculated by Mewe et al., 1981.
Collisional Excitation Ions in the ground state are excited by free electron
collisions, and subsequent de-excitation induces line emission. The rate coefficient of collisional excitation for this process are calculated as follow. We
assume that all excited ions Z +z were in the ground state (of statical weightw g ). After excitation by electron impact to the upper level, the line transition
takes place with some branching ratio factor BR = Ar /( A a + Ar ) described
as:
!
E/IE
χ
,
(3.28)
SzE/IE = 8.62 × 10−6 (Ω/w g ) BR C1 Te−1/2 exp − z
kTe
where w g is the statistical weight of the ground state, χE/IE
is the photon enz
ergy, C1 is the correction factor that is a constant depending on the ion species
and the transition. Ω is the collision strength averaged over the Maxwellianelectron energy distribution (given in Mewe et al., 1978) and described as:√Ω/w g = (8π/ 3)(χH /χE/IE) f g,(3.29)zwhere f is the absorption oscillator strength of the considered line, g is theaverage of the Gaunt factor, which can be regarded as a constant dependingon the electron temperature.
As shown in Equations (3.28) and (3.29), the rate coefficient of collisional
excitation includes the absorption oscillator strength, which will be describedin detail in Subsection 3.3.3. Collisional excitation is the most general X-ray
line emission process in SNR plasmas. Therefore, the line emissions withlarge oscillator strength, such as a resonance line of He-like ion, have highintensity. Therefore, line emissions with large oscillator strength, such as aresonance line of He-like ion, are observed prominently in X-ray spectra of
SNRs.
Radiative Recombination When the electron is bound to an excited levelby radiative recombination, a line photon is emitted by spontaneous deexcitation (cascade decay: bottom left of Figure 3.5). The contribution to a

3.1. Collisionally Ionized Plasma

39certain line in ion Z +z due to radiative recombination of ion Z +(z+1) is written as:
−η

−11
αzRR
BR CRR (z + 1)2η +1 Te ,+1 = 10(3.30)

where η and CRR is a constant that differs for each ion species, and η has avalue of about 0.6 to 0.8 (see Table 1 in Aldrovandi et al., 1973).
Dielectronic Recombination The rate coefficient of dielectronic recombination of ion Z +(z+1) of a line emission by ion Z +z is given by


χs
DR
−16 −3/2
αz+1 = 2.07 × 10 Te
Bs exp −,(3.31)
kTe
where
Bs =

ws A a Ar.
w1 ( A a + ΣAr )(3.32)

χs is the energy difference between the upper level s of the satellite line (withstatistical weight ws ) in ion Z +z and the ground state (with statistical weightw1 ) of the recombining ion Z +(z+1) .
Continuum Emission
Bremsstrahlung A charged particle emits a photon when it is scattered byan electric field created by another charged particle (top left of Figure 3.5).
This process is called bremsstrahlung. Since a electron has a smaller mass
than an ion, bremsstrahlung emitted in plasma is mainly from electrons scattered by the electric field of ions.
f ( ve ) =

4πv2e



me
2πkTe

3/2

me v2e
exp −
2kTe



(3.33)

When the electrons follow the Maxwell-Boltzmann distribution described as
the equation (3.33), the intensity per unit time, unit volume, and unit frequency is


dW
hν
−1/2 2∝ Te
Z ne ni exp −gf f ,(3.34)
dVdtdvkTe
where ve , me , Te , ne are the velocity, mass, temperature, number density ofelectrons respectively, Z and ni are atomic number and number density ofthe scattering ion, k is the Boltzmann constant, h is the Planck constant. g f f
is velocity averaged Gaunt Factor. In X-ray emitting plasmas,

3kTe
gf f =.
(3.35)πhν40

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
PlasmaFrom Equations (3.34) and (3.35), the spectrum bremsstrahlung becomes acontinuum spectrum with a cutoff at hν ∼ kTe , so the electron temperature
can be determined by the cut-off energy.
Radiative Recombination Continuum Radiative Recombination Continuum (RRC) is electromagnetic radiation emitted when free electrons are captured by ions. As shown in the bottom left of Figure 3.5, a photon is emitted with an energy corresponding to the sum of the binding energy of theelectron and the kinetic energy of the free electron. The intensity per unit
time, unit volume, and unit frequency of the RRC of electrons following theMaxwell-Boltzmann distribution is E− E (dW
exp − kTedge( E ≥ Eedge )e∝(3.36)
dVdtdv
0( E < Eedge ),where Eedge is the binding energy of the electron. The spectrum of RRC has
an energy edge with an energy corresponding to the binding energy of electron.

3.2

Interaction between Collisionally Ionized Plasma
and Neutral Gas

When an ion collide with a neutral atom, various reactions occur such as
ionization and excitation of an ion or a neutral atom, and charge transferfrom a neutral atom to an ion (Charge eXchange: CX). Since SNRs is systems
in which neutral atoms and ions interact in the shock wave, these processes
cannot be ignored in some cases. CX is a particularly important process when
the collision velocity between ions and neutral atoms is lower than the orbital
velocity of electrons in the classical hydrogen atom ( ∼ 2000 km/s). CX
may be important because shock wave velocities of SNRs range from several
hundred to 2000 km/s. Here we describe the basic physics of CX. Detailed
spectral features of CX are summarized in Subsection 3.4.2.

3.2.1

Charge Exchange

When a neutral atom collides with an ion, the potential barrier between them
become lower, and an electron bound to the neutral atom can be transferedto the energy level of the ion (see Figure 3.6). This interaction is described as:
Xq+ + Y → X∗(q−1)+ + Y+ ,(3.37)

where X is the ion, q+ is its charge number, Y is the neutral atom, and thestar superscript denotes an excited state. If the neutral atom is not hydrogen, multiple electrons can be transferred, leading to a multiply excited ion.
However the cross section for multiple transfer is much smaller than that fora single transfer (Wargelin et al., 2008). Therefore, we consider CX between

3.2. Interaction between Collisionally Ionized Plasma and Neutral Gas

41

F IGURE 3.6: A Diagram of CX (Cravens, 2002). The vertical axis
corresponds to the potential depth. The solid cyan lines represent potential barriers and the solid red lines represent energy
levels.

ions and neutral hydrogen. Since the potential of hydrogen is shallower thanthat of ions, the electron first is captured the excited level. Then, CX producesa X-ray line emission due to subsequent de-excitation. The most probable energy level to which an electron is transferred from neutral hydrogen is given
by:
np ∼

q−1
1+ p
2q

!−1/2

∼ q3/4 ,(3.38)

where, q is the charge number of the ion (Janev and Winter, 1985). Therefore,
for example, in the case of the reaction described as Equation (3.39), electrons
are likely to be trapped to n ∼ 4.
O7 + + H → O ∗ 6 + + H +(3.39)

Besides this n-selectivity, the electron transfer process exhibits also selective properties with respect to the electron orbital angular momentum(l-selectivity). Contrary to the case of n-selectivity, l selective property no
underlying physical criteria have been establish. However, the l-selectivity
considered to be related to the collision velocity between the neutral hydrogen and an ion (e.g., Janev and Winter, 1985). In the low velocity regime, thepeak l shell is l = 1 or 2. On the other hands, in the high velocity regime, thel-selectivity follows the statistical weight (e.g., Gu et al., 2016).

42

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
Plasma3.3

Radiative Transfer

3.3.1

Optical Depth

When radiation passes through a material of optical depth τ, its intensity
becomes e−τ . The optical depth τ is defined asτ = αL = κρL = Nσ,(3.40)

where α (cm−1 ) is the absorption coefficient, L (cm) is the thickness of thematerial, κ (cm2 /g) is the mass absorption coefficient, N (cm−2 ) is the column density, σ (cm2 ) is the cross section. The column density is obtained byintegrating the number density n (cm−3 ) along the line of sight.
N=

Z L
0

n(s)ds(3.41)

In general, these values depend on photon energy.
The mean free path of photons is defined as the average distance a photon
travels through an absorbing material without being absorbed. The probability that a photon travels through the material with τ is e−τ . The mean opticaldepth traveled is equal to unity.

⟨τ ⟩ =Z ∞0τe−τ dτ = 1(3.42)The mean free path l in a homogeneous medium is determined by⟨τ ⟩ = αl = 1(3.43)

From Equations (3.43) and (3.40), the relation between the mean free path land the thickness of the material L (cm) is
l = L/τ.

(3.44)

A medium with a low optical depth (τ < 1) is described as optically thin,
whereas one with a high optical depth is optically thick. According to theequation (3.44), in an optically thick medium, the mean free path l is smaller
than the thickness of the material L, and most photons are absorbed. On theother hand, most photons pass through an optically thin medium without
being absorbed.

3.3.2

Radiative Transfer

The source function Sν is defined as the ratio of the emissivity jν [erg/s/cm3 /Hz/str]
to the absorption coefficient αν [cm−1 ]:
Sν ≡

jν
.
αν

(3.45)

3.3. Radiative Transfer

43

Using the source function, the radiative transfer equation can be written,dIν
= − Iν + Sν ,dτν(3.46)

where Iν is the specific intensity. Considering only absorption and radiation(without considering scattering), we can solve the radiative transfer equation
if Sν is independent of τ.
Iν (τ ) = Sν (1 − e−τ ) + Iν (0)e−τ(3.47)In the optically thin case (τ ≪ 1), the specific intensity described asIν (τ ) = Iν (0) + τν Sν .

(3.48)Iν (τ ) is proportional to the optical depth. Therefore, in three dimensions, thebrightness of an optically thin object is proportional to its volume. Emission
from the background of an astronomical object (Iν (0)) is also observed.
In the optically thick case (τ ≫ 1), the specific intensity described asIν (τ ) = Sν .

(3.49)The background emission (Iν (0)) cannot be observed. Iν (τ ) is independentof τ. Therefore, the brightness of an optically thick object is proportional toits surface area.
It is known that SNR plasmas are optically thin for continuum radiation.
For example, the optical depth of Thomson scattering (τT ) can be calculated
as:


L
ne 
−5
τT = 2.1 × 10,
(3.50)
10 pc
1 cm−3where the typical electron number density and the diameter of SNR are 1–
10 cm−3 and the diameter is 10 pc. Thus, the X-ray emission from SNRs is
optically thin in most energy bands, so self-absorption is negligible.

3.3.3

Resonance Scattering

Resonance Scattering (RS) is an apparent scattering phenomena due to anabsorption and re-emission of line photons by ions. When a photon with energy equal to the energy difference between an excited state and the groundstate of an ion is incident on the ion, the ion absorbs the photon and excited.
The ion then emits a photon of energy equal to the incident photon due to deexcitation. This makes the photon appear to be scattered without changing
its energy.
The optical depth τ of RS is expressed as:τ = Nion σ,(3.51)

44

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
Plasmawhere Nion is the column density of ions that cause RS, for example, the column density of He-like O when considering the RS of the O VII Heα line.
Using elemental abundance (nz /nH ) and ion fraction, τ can be described as:  
ninz
σ.
(3.52)
τ = NHnH
nz
The cross section of RS at a line center is√ 2  πe f v −1σ=
cm2 ,mc ν c(3.53)

where f is the oscillator strength, ν is the frequency, m is the electron mass, v
is the root-mean-square kinetic velocity of the ion described as:
v2 =

2kTi
+ v2tur .
mi(3.54)The kTi is the ion temperature in keV, mi is the ion mass, and vtur is the rootmean-square turbulent velocity due to motions other than thermal one. From
Equations (3.52), (3.53) and (3.54), Kaastra and Mewe, 1995 calculate the τ tobe:
  
nz
M
4.24 × 106 f NH nnzi
nH
TkeV
τ=


0.0522Mv2100 1/2
EeV 1 +
T

1/2,

(3.55)

keV

where f is the oscillator strength of the line, EeV is the line centroid energy ineV, NH is the hydrogen column density in cm−2 , ni is the number density of
the ion, nZ is the number density of the element, M is the atomic weight of
the ion, TkeV is the ion temperature in keV, and v100 is the micro-turbulencevelocity in units of 100 km s−1 . From Equation (3.55), it can be seen that an
optical depth of RS is proportional to an oscillator strength f . An oscillator
strength f lu for a transition from an energy level l to u is expressed as:
f lu =me chν B ,
4π 2 e2 lu lu(3.56)

where Blu is Einstein absorption coefficient. Therefore, it is a dimensionlessquantity that reflects the probability of electronic transitions between energy
levels, which are induced by the absorption of photons. Table 3.1 summarizes
the oscillator strengths of the O VII (He-like oxygen) Kα lines.
It can be seen that the oscillator strength differs by many orders of magnitude from line to line. Transitions that satisfy the selection rule introduced in
Subsubsection 3.1.2.1 have large oscillator strengths. Using Equation (3.55),
optical depths of Resonance line and Forbidden line listed in Table 3.1 arecalculated as:


L
ne 
τr = 1.06(3.57)
10 pc
1 cm−33.4. Plasma Diagnostic

45

F IGURE 3.7: Line intensity ratios expected for CIE or NEI (ionizing) plasma. (calculated with pyatomdb) (a) Lyβ/α ratio of O
as a function of electron temperature. (b) Same as panel (a), but
for Si.

τf = 2.95 × 10

−10



L
10 pc



ne 
1 cm−3(3.58)

where we assume the plasma is in the CIE state with an electron temperatureof 0.2 keV whose abundance of oxygen is solar value (Anders and Grevesse,
1989). Thus some resonance lines have an optical depth of approximately 1
in a typical SNR plasma, and the effect of RS may not be negligible.

3.4

Plasma Diagnostic

3.4.1

Electron Temperture & Charge Distribution (Ionization
timescale)

As mentioned in 3.1.1.2, the electron temperature and ionization timescale(ne t) are important parameters that characterize the physical condition of
SNR plasmas. In this subsection, we introduce how to constrain these parameters based on the line ratios of the X-ray spectrum.
Electron Temperture
To derive the electron temperature, it is useful to measure the line intensity
ratio between atomic transitions involving different quantum levels of identical ions. This is because collisional excitation to higher quantum levels requires a larger kinetic energy (i.e., higher temperature) of free electrons. Figure 3.7 shows Lyβ/Lyα ratios of H-like O and Si as a function of kTe . It canbe seen that this line ratios are sensitive to the electron temperature, while
the difference due to ne t is relatively small. The electron temperature canalso be measured by the shape of the continuum (bremsstrahlung and RRC).
However, when the continuum is dominated by non-thermal components,
or when multiple plasmas with different temperatures are observed, electron
temperature measurements based on line ratios are particularly useful.

46

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
PlasmaF IGURE 3.8: (a) Lyα/Heα ratios of O as a function of ne t (calculated with pyatomdb). (b) Same as panel (a), but for Si.

Charge Distribution (Ionization timescale)
By investigating the ion population together with the electron temperature,
we can determine whether the plasma has reached ionization equilibrium.
For example, the line ratio of Lyα and Heα reflects the population of H-likeand He-like ions, which can be used to constrain the ionization timescale ne t.
Figure 3.8 show the line ratio of Lyα and Heα of O as a function of ne t In
typical SNRs, the heavy elements lighter than Ca (Z ≤ 20) are dominated by
the He-like and H-like ions, and thus this diagnostic can be applied.

3.4.2

Diagnostic Using Heα Line ratio

Since He-like ions are closed-shell structures, they exist under a wide rangeof physical conditions (temperature, density, plasma ionization state, etc.).
Therefore, the line flux ratio of He-like ions is useful for probing a wide rangeof physical quantities. Here, we describe the basics of the line structure of
the Kα lines (Heα line) of He-like ions, and explain the diagnostic methodof physical quantities and physical processes of plasma using the Heα Lineratio.
Under the assumption of LS coupling, the energy levels of He-like ionscan be expressed as shown in the Figure 3.9. The Heα line emitted by thetransition indicated by w and z in the left panel of Figure 3.9 is called theresonance (r) and forbidden ( f ) line, respectively. The Heα lines indicated byboth x and y are called intercombination (i) line. Basic information for O VII
and Fe XXV Heα lines are summarized in Table 3.1. According to selection
rule (Subsubsection 3.1.2.1), the transition w is allowed transition. On theother hand, x, y, and z are forbidden transition. This is the reason why resonance line has several orders of magnitude larger oscillator strength thanthe other lines. The energy level shown in red in the right panel of Figure
3.9 is called a singlet where the two electrons have opposite spin directions.
The energy level in blue is called triplet with electrons with spins in the same
direction. x, y, and z are triplets, which have a lower oscillator strength thanw, but have a higher statistical weight. Therefore, electrons are more likely to
be trapped in the excited levels of x, y, and z due to recombination. Table 3.1

3.4. Plasma Diagnostic

47

F IGURE 3.9: (left) Schematic Grotrian diagram for a He-like ion(Porquet et al., 2001). Energy levels up to n = 2 are shown, considering the case where only one electron is excited. w (or r),x, y (or i), and z (or f): resonance, intercombination, and forbidden lines, respectively. Full upward arrows: collisional excitation transitions, broken arrows: radiative transitions, andthick skew arrows: recombination (radiative and dielectronic)
and cascade processes. (right) Same as left panel, but including
the energy level of n > 3 (Smith et al., 2014). g corresponds tostatistical weigh.

shows the centroid energy of Heα lines of O and Fe. These transitions can be
resolved only by high-resolution spectrometers, such as the Reflective Grating Spectrometer on board XMM–Newton and the X-ray microcalorimeter
Resolve on board XRISM.
TABLE 3.1: Basic information for O VII and Fe XXV Heα lines(taken from Atomdb)

Element
O

FeLine Name
Resonance Line (w)Intercombination Line (x)Intercombination Line (y)Forbidden Line (z)Resonance Line (w)Intercombination Line (x)Intercombination Line (y)Forbidden Line (z)Upper State
1P
1
3P
2
3P
1
3S
1
1P
1
3P
2
3P
1
3S
1

Oscillator Strength
7.20 × 10−1
1.14 × 10−7
8.19 × 10−5
2.00 × 10−10
7.20 × 10−1
1.14 × 10−7
8.19 × 10−5
2.00 × 10−10Centroid Energy
574 eV569 eV569 eV561 eV
6700 eV
6680 eV
6670 eV6640 eVR-ratio
The intensity ratio ( f /i) is called R-ratio. According to the selection rule,the transition between 3 S1 and 3 P2,1,0 is an allowed transition. In low density plasmas such as in SNRs, the transition of ions in the excited state to theground state by spontaneous emission takes place faster than the next collision. Therefore, the excitation from 3 S1 to 3 P2,1,0 is negligible. However, as

48

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
PlasmaF IGURE 3.10: R-ratio as a function of electron density.

the plasma density increases, the collisional excitation from 3 S1 to 3 P2,1,0 cannot be ignored, and the R-ratio increases. Figure 3.10 shows the R-ratio as afunction of density. Each element is sensitive to different density regions.
G-ratio
The G-ratio, defined as the ratio of ( f + i )/r, can be used to probe the electron
temperature. As mentioned above, the flux ratio between i and f depends
on the electron density of the plasma. Therefore, ( f + i )/r can be used toconstrain the electron temperature independent of electron density. Since
collisional excitation is the dominant radiation process of Heα lines in IP and
CIE, which are commonly seen in SNR, the value of G-ratio is generally small.
For example, the G-ratio of the O VII Heα lines is generally smaller than unity
in CIE and IP (Figure 3.11). On the other hand, RP, inner-shell ionization,
charge exchange, and resonance scattering, which are described below, arecharacterized by a high G-ratio. Therefore, as explained below, G-ratio canbe used to probe these physical processes and ionization states.
RP RP is the plasma in a higher ionization state than the CIE plasma inwhich recombination process is dominant. Therefore, its X-ray spectrum ischaracterized by relatively strong radiation components due to recombination, such as RRC and line emission from cascade decay appear. When freeelectrons are trapped by H-like ions, electrons tend to enter spin triplet energy levels (the energy levels shown in blue in the right panel of Figure 3.9)

3.4. Plasma Diagnostic

49

F IGURE 3.11: G-ratio of O VII Heα lines as a function of electron
temperature.

because the selectivity of angular momentum l and spin s of captured electrons follows statistical weights. Moreover, electrons captured into the spin
triplet energy level always goes through the n = 2 levels to the ground state
due to the selection rule. Therefore, the final de-excitation of the electrons
emits either f or i, resulting in a large G-ratio (Figure 3.11). In addition to
the high G-ratio, the Kβ and Kγ, ... lines due to the de-excitation of electrons
captured into n > 3 spin siglet energy levels to the ground state and the RRC
due to radiative recombination are also enhanced in the RP.
Inner-shell Ionization The electron configuration of Li-like ions in groundstate is 1s2 2s. There are two excited states, 1 S0 , and 3 S1 , induced by the innershell ionization of the electron in the 1s orbit of the Li-like ion. Therefore, thetwo-photon decay emission due to the 1s2s (1 S0 ) → 1s2 (1 S0 ) transition andthe forbidden line due to the 1s2s (3 S1 ) → 1s2 (1 S0 ) transition can be emitted
by inner-shell ionization, which enhance the observed G-ratio.
CX Since CX is a recombination process, the G-ratio is enhanced for thesame reason as RP. On the other hand, CX is a non-radiative recombination
process, RRC is not enhanced. The bottom panel of Figure 3.12 shows theexperimental CX spectra where the w line is relatively weaker than z line.
RS As explained in Subsection 2.3.3, a cross section of RS is proportional tothe oscillator strength. Table 3.1 summarize the oscillator strengths of Heαlines of O and Fe. Resonance lines have oscillator strengths several orders of

50

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
PlasmaF IGURE 3.12: Experimental He-like and H-like Fe spectra fromcollisional excitation by monochromatic electron as 15 keV (top)and from CX with N2 (bottom) (Wargelin et al., 2008). Bothspectra are observed with the X-ray microcalorimeter.

3.5. Observational Evidence for Charge Exchange and Resonance Scattering51
in Astrophysical Plasma

F IGURE 3.13: Schematic image of the RS in SNRs. (left) In thecase of spherical symmetric SNRs. The RS effect is cancelled
out and an enhancement of f /r is not observed 3.1. (right) In
the case of spherical asymmetric SNRs. The intensity reductionof resonance line would be observed.

magnitude larger than those of forbidden lines and intercombination lines,
and are particularly susceptible to the effects of RS. Therefore, the G-ratio
can be enhanced by the intensity reduction of resonance line due to RS.
The observed RS effect depends on the structure of the object. If an object
has an ideal spherical symmetric structure, the RS effect is cancelled out andan enhancement of f /r is not observed (the left panel in Figure ??). On theother hand, when observing an object with a long structure in the line-ofsight direction, only the resonance lines are scattered outside the line-of-sightdirection, resulting in a high G-ratio.

3.5

Observational Evidence for Charge Exchange
and Resonance Scattering in Astrophysical Plasma

3.5.1

CX

CX is a possible X-ray emission process in any system where ionized gas
interact with a neutral gas. CX X-ray emission from an astrophysical object
was first observed from the comet Hyakutake with ROSAT in 1996 (Lisseet al., 1996). Since then, CX X-ray emission has been detected from otherastronomical object, such as stars (e.g., Pollock, 2012), star forming regions(e.g., Liu et al., 2012), active galactic nuclei (e.g., Gu et al., 2022), and clustersof galaxies (e.g., Conselice et al., 2001).
Lallement, 2004 calculate the contribution of X-ray emission from CX inshock front of an SNR. Based on their estimation, the X-ray emitting region
would be the outer edge of the shock front (a few percent of the shock radius: Figure 3.14). The observational evidence is obtained from the Cygnus
Loop (Uchida et al., 2019), Puppis A (Katsuda et al., 2012), and G296.1−0.5(Tanaka et al., 2022). The elemental abundances in some of the outermost rim
regions of the Cygnus Loop are anomalously enhanced (∼ 1 solar), whereas

52

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
PlasmaF IGURE 3.14: (left) Contribution of X-ray emission from thermal plasma and CX as functions of the impact parameter p defined as showan in the right panel (Lallement, 2004). The CX
exceeds the thermal emission at p ∼ 0.99. (right) Schematic image of CX in the outer layer of SNR.

the abundances are typically ∼ 0.2 solar for most of the rim (e.g., Katsuda etal., 2008). Based on CCD observation, Katsuda et al., 2011 found the possible
CX X-ray emission from these abundance enhanced regions where a strongline emission is detected at ∼ 0.7 keV, which is likely cascade lines of Helike oxygen (O Heγ + δ + ϵ). They pointed out that the excess at ∼ 0.7 keV
due to CX emission causes an apparent abundance enhancement if an X-ray
spectrum is fitted with a CIP model without the CX component. RGS observation of the Southwestern knot region of the Cygnus Loop where elemental
abundances are enhanced show high O VII Heα F/R ratiom that cannot be reproduced by emission from CIP (Uchida et al., 2019), which provides strong
evidence for CX X-ray emission.

3.5.2

RS

Since RS has large cross sections, it may be observed even in objects that are
generally assumed to be optically thin, such as SNR and galaxy clusters. At
the center of the Perseus cluster, Hitomi Collaboration, 2018 observed a larger
F/R ratio than expected from an optically thin CIP, which is well explainedby the RS effect (Figure 3.15).
Kaastra and Mewe, 1995 predicted that RS of X-ray photons can occur ina plasma with a large depth along the line of sight such as a rim of SNRs.
Several observational signatures of RS have been reported such as a highforbidden-to-resonance ratio of O VII Heα obtained from grating spectra of
DEM L71 (van der Heyden et al., 2003) and N23 (Broersen et al., 2011). A
difference in surface brightness between forbidden and resonance lines also
supports the presence of RS (e.g., van der Heyden et al., 2003). Based ona observation of the Cygnus Loop with Suzaku satellite, Miyata et al., 2008
claimed that a depleted abundance of O may be partially explained by RS.
A recent grating observation of the Loop also hints at a possibility of RS

3.5. Observational Evidence for Charge Exchange and Resonance Scattering53
in Astrophysical Plasma

F IGURE 3.15: (left) The Hitomi SXS observation regions overlaid on the Chandra X- ray image of the Perseus Cluster in the
1.8–9.0 keV band (Collaboration, 2017). (top right) Flux suppression in the strongest line of He-like Fe XXV (w) in the centerregion of the Perseus Cluster. The red lines show the best-fitting
models. (bottom right) Comparisons of the observed and predicted ratios of the Fe Heα resonance (w), Heα forbidden (z)lines. Observations are shown as magenta crosses and the simulations with RS as black diamonds and the same without RS
as gray diamonds.

(Uchida et al., 2019). These studies suggest that the RS effect may potentially be significant in SNRs. However, no strong observational evidence has
been established so far.

54

Chapter 3. Plasma Processes and X-ray Radiation in Collisionally Ionized
PlasmaF IGURE 3.16: (top) Chandra X-ray image of DEM L 71 for theenergy range 0.7–1.0 keV (van der Heyden et al., 2003). The
RGS dispersion axis is also indicated by the black arrow. RGS
monochromatic images for O VII resonance (lower left) and O
VII forbidden (lower right).

55

Chapter 4

Instruments
As mentioned in Chapter 3, signs of the RS or CX are expected to be found inintensity ratios of multiplet lines (e.g., OVII Heα), which cannot be resolved
with widely-used detectors such as CCDs. In this study, we thus use a grating spectrometer, the RGS onboard the XMM-Newton satellite, which has
a much better energy resolution than the CCDs. In this chapter, we briefly
summarize the basic properties of the XMM-Newton satellite and its instruments.

4.1

XMM-Newton satelite

The X-ray Multi-Mirror Mission (XMM)-Newton satellite (Jansen et al., 2001)(hereafter XMM-Newton) was launched by the European Space Agency (ESA)in December 1999 and put into an elliptical orbit with a perigee of 5000 km
and an apogee of 115000 km. Figure 4.1 and Figure 4.2 show a schematic viewof the spacecraft and the orbit of XMM-Newton. XMM-Newton is equippedwith two types of telescopes: three X-ray reflecting telescopes with Wolter-I
type mirrors and an optical/ultraviolet telescope with an aperture of 30 cm.
As explained below, three CCDs and two grating spectrometer are mounted
on the focal plane of the telescopes. The European Photon Imaging Camera(EPIC) is an X-ray imager of CCDs for an energy range of 0.15–12 keV. EPIC
consists of two types of charge-coupled devices (CCDs), Metal Oxide Semiconductor (MOS) (Turner et al., 2001) and pn (Strüder et al., 2001). The Reflection Grating Spectrometer is suited for high spectral resolution X-ray spectroscopy in the energy range 5–38 Å(0.33-2.5 keV). The Optical/Ultraviolet
Monitor (OM) CCDs for optical/ultraviolet imaging and grism spectroscopy
is mounted on the focal plane of the optical/ultraviolet telescope telescopes.
In this study, we use the data obtained with MOS and RGS. Therefore, wesummarize the design and performance of the X-ray telescope, EPIC-MOS,
and RGS in the following section.

56

Chapter 4. Instruments

F IGURE 4.1: Sketch of the XMM-Newton satellite (XMM User’s
Handbook). The mirror modules, two of which are equippedwith Reflection Grating Arrays, are visible at the lower left.
The right side is the focal plane, where the CCDs are installed.
Green, purple, and light blue are radiators for MOS, pn, and
RGS, respectively.

4.2. X-ray Telescopes

57

F IGURE 4.2: Orbit of the XMM-Newton satellite (XMM User’s
Handbook).

4.2

X-ray Telescopes

4.2.1

Overview

Since the refractive index of materials for X-rays is approximately unity, we
commonly use a total refection or diffraction to focus X-rays. The X-ray telescopes on board XMM-Newton utilize the Wolter-I type optics which focus
the X-ray by total reflection. The Wolter-I type optics employ paraboloidal
and hyperboloidal surfaces as primary and secondary mirrors (see Figure4.3). The X-ray telescopes are comprised of 58 Wolter-I type optics with a 70
cm aperture, a focal length of 7500 mm and a field of view of 30 arcmin.

4.2.2

Performance

Here we describe a performance of the X-ray telescopes, which is mainly
characterized by the imaging capability and effective area.
Imaging Performance
The Point Spread Function (PSF) describes the distribution of observed Xrays in the focal plane for point sources. The fractional encircled energy (FEE)
is defined as the ratio of the number of X-ray photons obtained by integrating
the PSF in the radial direction to total X-ray photons. Figure 4.4 shows FEE
as a function of radius from the center of the PSF.

58

Chapter 4. Instruments

F IGURE 4.3: A schematic view of the Wolter-I type optics of the
XMM-Newton (XMM User’s Handbook).

F IGURE 4.4: Fractional encircled energy of the MOS1 (left) and
MOS2 (right)(XMM User’s Handbook).

4.3. X-ray CCD Camera (European Photon Imaging Camera: EPIC)

59

F IGURE 4.5: Effective areas of EPIC cameras and RGS spectrometers (XMM User’s Handbook).

Effective Area
The effective area describes the photon collection capability of X-ray telescopes, which depend on an actual aperture size of the telescope and an efficiency of the detector. Figure 4.5 shows the effective area on the opticalaxis considering the efficiency of each detector. The effective area for offaxis sources is smaller than that for on-axis sources (Figure 4.6). This effectis called "vignetting", and becomes significant in higher energy bands. We
corrected the vignetting effect in our image analysis and spectral analysis.

4.3

X-ray CCD Camera (European Photon Imaging
Camera: EPIC)

As mentioned above, the XMM-Newton satellite carries two types of CCD
detectors called MOS and pn. Among them, we only used the MOS for thefollowing imaging analysis and a wide-band spectroscopy. In this section,
we thus focus on a details of the design and performance of the MOS.

60

Chapter 4. Instruments

F IGURE 4.6: Vignetting factor defined as the ratio of the effective area against incident angle to that on on-axis (XMM User’s
Handbook).

4.3.1

Overview and performance of MOS

Photon-counting CCD detectors have a good imaging performance and energy resolution, and are widely used as focal plane detectors of X-ray astronomical satellite. A CCD is an image sensor that consists of two-dimensionallyarranged pixels. Photons incidenting on each pixel are absorbed in the depletion layer of the CCDs, and finally charges the number of which is proportional to the energy of the X-ray photons are generated. The number of X-ray
photons incident on one pixel is at most one per exposure when observing
a general X-ray astronomical object. Therefore, by measuring the number ofcharges generated at each pixel, the position and energy of the incident X-ray
can be determined.
Figure 4.7 represents an example of a schematic of the structure of CCDs
used for X-ray astronomy. CCDs are classified into two types, front-illuminated(FI) and back-illuminated (BI) by a direction of photons onto the detector.
Since X-rays enter the FI CCDs from the electrode side, some photons are absorbed by the electrodes. For this reason, FI CCDs has a disadvantage thatthe detection efficiency drops on the low energy side. In addition, since theelectrode of the FI CCDs is exposed to debris and micrometeorites, they are
more easily to be damaged than BI CCDs. On the other hand, BI CCDs aresuperior to FI CCDs in these respects, but they have a disadvantage that itis difficult to thicken the depletion layer because it is not easy to provide anelectrode on the back side. Figure 4.8 shows schematic views of the field ofview (FOV) of the telescope and MOS and pn. The MOS consists of seven FI
CCDs with an imaging area of 2.5 cm square.
The energy resolution of the MOS during the early in-orbit phase wasabout 10 % worse than the ground calibration value before launch. After that,

4.3. X-ray CCD Camera (European Photon Imaging Camera: EPIC)

61

F IGURE 4.7: Cross sectional view of a Front-illuminated CCD(left) and a Back-illuminated CCD (right).

additional cooling was performed to improve the energy resolution, and thecurrent energy resolution is almost stable at the pre-launch value (see Figure
4.9).
TABLE 4.1: design and performance of MOS (XMM User’s
Handbook)

Energy range
Pixel size
Pixel grid
Timing resolution
Energy resolution (FWHM)

4.3.2

0.15 − 12 keV
40 µm
600 × 600
1.5 ms∼ 70 eV (@1 keV)

Background of the MOS

The background of the MOS includes X-ray emission from celestial objects
and instrumental background. The instrumental background consists of detector noise which is important in the low energy band (< 0.3 keV) and theparticles-induced background which is important in the high energy band.
The particles-induced background is classified into a time-varying component and a non-time-varying component. The non-time-varying component
is observed steadily even outside the exposure time, and has a spectrum
consisting of a flat continuum and fluorescence lines as shown in Figure
4.10. These are high-energy charged particles directly incident on the detector from outside the field of view of the telescope, and X-ray fluorescence due
to interactions between the detector housing and the high-energy charged
particles. The time-varying component is detected within the telescope fieldof view on the CCD (the shaded area of Figure 4.8) and has a flat-shaped
spectrum. It is considered that this component is caused by the direct incidence on the CCD of low-energy protons.

62

Chapter 4. Instruments

F IGURE 4.8: Schematic view of MOS (left) and pn (right). The
shaded circle areas correspond to the field of view with 30 arcmin (XMM User’s Handbook).

4.4

Reflection Grating Spectrometer (RGS)

4.4.1

Overview

RGS consists of Reflection Grating Assemblies (RGA) and CCDs (RGS Focal
Cameras: RFC) that detect photons dispersed by the RGA. RGS is sensitive
to X-rays of 5-38 Å (0.35–2.5 keV) and has an energy resolution of 2.0 eV (@1 keV) for point sources. This makes it possible to separate and detect Kshell transition lines containing He-like triplet of C, N, O, Ne, Mg, and Si,
and L-shell transition lines of Fe and Ni. Figure 4.11 shows the layout of
the telescope, RGA and RFC. Two RGAs are placed between the telescopeand MOS1 and 2, respectively. About 44% of photons reflected by the telescope incident the MOS, 40% are dispersed, and the rest are absorbed by thegrating. RGS resolves the wavelength of the photon from the X-ray incident
position on the RFC by the following diffraction grating equation describedas:mλ = d (cos β − cos α),(4.1)

where m (= −1, −2, ...) is the spectral order, d is the groove spacing, λ is the
X-ray wavelength, α is the angle between the incoming ray and the grating
plane, and β is the angle between the outgoing ray and the grating plane. As
shown in Figure 4.11, RGA, RFC, and MOS are placed on the Rowland circle,
which reduces the aberration caused by arranging the gratings. 1st, 2nd,and higher order diffracted photons may be overlapped on the RFC. Since4.4. Reflection Grating Spectrometer (RGS)F IGURE 4.9: Energy resolution of MOS (XMM User’s Handbook). Solid lines are the results of the ground calibration test
and each data point is those of in-orbit calibration. The difference in the shapes of the data points corresponds to the difference in the date on which the data were obtained. Cooling to improve the energy resolution was performed between
Rev.#530 and Rev.#560.

F IGURE 4.10: Background spectrum of the MOS (XMM User’s
Handbook). The emission lines at 1.5 keV and 1.7 keV are theK-shell emission lines of Al and Si, respectively.

63

64

Chapter 4. Instruments

CCDs are used as the X-ray detector, two pieces of information regarding theincident photon energy are obtained: the signal charge amount of the CCD
and the incident position on the RFC. From these two pieces of information,
we can separate the spectral order as shown in Figure 4.14.
RGA
RGA consists of 182 blazed gratings. As shown in Figure 4.12, blazed gratings have a sawtooth shape, which makes it possible to cause interference
between diffracted photons of a specific wavelength and specific order andreflected photons by adjusting the incident angle (α) and the blaze angle. A
plane grating maximizes the intensity of the 0 th-order photon, but a blazed
grating can maximize the intensity of any order of diffracted photon due tothis structure. In the case of RGA, the blaze angle and incident angle (α) are
set so that the diffraction efficiency is maximized for the first-order photonsof 15 Å. Each grating is a SiC substrate about 10 cm × 20 cm in size coatedwith 2000 thick gold. The groove density varies by about 10 % to correctaberrations on the RFC, but is about 646 mm−1 .
RFC
The RFC consists of 9 CCDs with almost the same specifications as the EPICMOS detector, arranged along the dispersion axis as shown in Figure 4.13.
Each CCD is assigned an ID number from 1 to 9 in order from the longer
wavelength side. Figure 4.14 shows an X-ray image obtained with RFC and aplot with the dispersion angle on the horizontal axis and the pulse height output from the CCDs on the vertical axis. The first- and second-order spectrum
overlap on the RFC. The RFC has an energy resolution of 160 eV (FWHM)
for 2 keV X-rays. The spectral orders are resolved by the energy resolution of
the RFC (see the bottom panel of Figure 4.14).

4.4.2

Performance

We summarize the basic performance of RGS in Table 4.2. The centroid energies of the resonance and forbidden of OVII Heα lines that we focus on in thisstudy are ∼574 eV and ∼561 eV, respectively. The RGS has an energy resolution of ∼1 eV in the energy band of 500 eV, which allows us to separate the
OVII Heα lines.
Although there are X-ray detectors other than the RGS that have better energy resolution than that of CCDs, the RGS is particularly suited forspectroscopy of the O VII Heα lines. The grating spectrometers onboard
the Chandra satellite (High Energy Grating: HEG, Medium Energy Grating: MEG, Low-Energy Transmission Grating: LETG) and the calorimeter
onboard the Hitomi satellite (Soft X-ray Spectrometer: SXS) also have better
energy resolution than that of CCDs. Figure 4.15 shows a comparison of effective area and energy resolution between these detectors and RGS. RGS has
better performance than the SXS and the grating spectrometers onboard the4.4. Reflection Grating Spectrometer (RGS)F IGURE 4.11: Schematic view of RGS (taken from XMM Users
Handbook). The RGA is shown in blue on the dashed circle.
The unit of length is mm. The solid red line corresponds to theoptical path of X-rays.

65

66

Chapter 4. Instruments

F IGURE 4.12: (left) Schematic view of RGS (taken from XMM
Users Handbook). α, β, and σ represent the angle of incidence,
the angle of diffraction, and the tilt of the mirror surface, respectively. The unit of length is mm. (right) Schematic view ofa blazed grating.

F IGURE 4.13: Sketch of an RFC chip array (XMM User’s Handbook). The Z axis is the dispersion axis. The left side corresponds to the short wavelength side (high energy side and theright side corresponds to the long wavelength side (low energyside).

4.4. Reflection Grating Spectrometer (RGS)F IGURE 4.14: (Top) X-ray image obtained with RFC. The horizontal axis corresponds to the dispersion angle, and the verticalaxis corresponds to the off-axis angle along the vertical direction of the dispersion axis. (Bottom) The X-ray image called"Banana plot". The horizontal axis corresponds to the dispersion angle, and the vertical axis corresponds to the pulse height
output from the CCDs. The area enclosed by the solid red line
corresponds to the first-order on the bottom and the secondorder on the top (XMM Users Handbook).

67

68

Chapter 4. Instruments

F IGURE 4.15: Performance comparison between RGS and otherhigh-resolution X-ray spectrometers (Mitsuda et al., 2010). The
left and right figures show the effective area and energy resolution, respectively.

Chandra satellite in the energy band below 1 keV where the L-shell transitionlines of Fe and the K-shell transition lines of C, N, and O are detected.
TABLE 4.2: Performance of the RGS (XMM User’s Handbook)

Energy range

0.35–2.5 keV (first-order)
0.62–2.5 keV (second-order)
Pixel size
27 µm
Energy resolution (FWHM)∼ 1.0 eV (@0.5 keV)

4.4.3

Background

The RGS background consists of X-ray emissions from other celestial objects (cosmic X-ray background; CXB) and instrumental background, andthese backgrounds have to be removed for spectral analysis. Similar to MOS,
the instrumental background of the RGS includes detector noise and theparticles-induced background. The particles-induced background consistsof high-energy charged particles (minimum ionizing particles, protons, ionsof heavy elements) incident from outside the telescope field of view and soft
protons that incident from the field of view of the telescope. The high-energy
charged particles are uniformly incident on the RFC. The RGS has two pieces
of information about the energy of the detected event: the position (dispersion angle) and the pulse height of the RFC. The background induced byhigh-energy charged particles can be removed by extracting the event fromthe region where the energy obtained from the peak value of the CCD andthe energy obtained from the detection position are not inconsistent (region
enclosed by the solid red lines in Figure 4.14). Since the other background,which is caused by soft protons, is time-dependent, the flux level is estimatedfrom the light curve of detected photons and can be removed by setting anappropriate threshold.

4.4. Reflection Grating Spectrometer (RGS)69

F IGURE 4.16: Line spread function (LSF) of RGS (XMM User’s
Handbook): the mirror response (green), broadened after folding with the RGA response (red), and after applying the detector energy selections (blue). All components are normalised totheir maximum.

4.4.4

Response Function

The RGS response function for a point source consists of the Line Spread
Function (LSF) and the effective area. The LSF is the observed X-ray distribution along the dispersion axis when monochromatic X-rays are incident
from the optical axis. The LSF is a convolution of the response functions ofmirrors and the grating, followed by convolution with the detector response(see Figure 4.16).
Since RGS is not equipped with a slit, the spatial broadening of the sourceis convolved with the spectrum. In order to perform a high-resolution spectroscopy of a diffuse source with RGS, it is required to generate a response
function convolved with the brightness distribution along the dispersion axisof the source. The off-axis angle θ of the incident photon corresponds to thedeviation of the wavelength:∆λ = 0.138m−1 (θ/arcmin) ,(4.2)

where m is the spectral order. For a diffuse sources, the energy resolution of

70

Chapter 4. Instruments

the RGS is degraded. The centroid wavelengths of the resonance and forbidden of OVII Heα lines are ∼21.60 Å and ∼22.10 Å, respectively. A differenceof 0.5 Å between the centroid wavelengths of these emission lines corresponds to a source diameter of about 4 arcmin from Equation (4.1). Therefore, in this study, we observe objects with a diameter of less than 4 arcmin
for spectroscopy of the O VII Heα lines.

71

Chapter 5

Evidence for RS and CX in SNRs
5.1

Our Objective and Target Selection

As mentioned in Chapter 3, recent studies pointed out the importance of CX
and RS in SNRs. Both of these processes are predicted to occur in SNRs by
previous calculations (e.g., Lallement, 2004; Kaastra and Mewe, 1995). However, it is not clear under what physical conditions CX and/or RS efficiently
take place, since established observational evidence is limited.
Possible detections of the spectral feature of CX are reported from RGS
observations of the Cygnus Loop (Uchida et al., 2019) and Puppis A (Katsuda et al., 2012). If CX takes place in these SNRs, its emitting region would
be the outer edge of the shock front (a few percent of the shock radius) in contact with dense clouds (see Chapter 3). However no quantitative constraint
on the CX emitting region has been performed. In the case of RS, several
observational signatures have been reported such as a high forbidden-toresonance ratio of O VII Heα lines obtained from RGS spectra of DEM L71(van der Heyden et al., 2003) and N23 (Broersen et al., 2011). A difference insurface brightness between forbidden and resonance lines also supports thepresence of RS (e.g., van der Heyden et al., 2003). However there is no strong
observational evidence has been established so far.
We perform a high-resolution X-ray spectroscopy of N49, J0453.6−6829
with the RGS onboard XMM-Newton in order to obtain observational evidence for CX and RS from SNRs based on quantitative measurements. Thesetargets are chosen among the SNRs located in the LMC, since they are relatively compact (∼ 2′ in diameter) remnant, which makes it possible to perform high-resolution X-ray spectroscopy with the RGS. Furthermore, both
SNRs are located in dense gases. As previous studies pointed out, CX is considered to be induced by collisions between SNRs and surrounding dense
gases (e.g., Uchida et al., 2019). Additionally, N49 has a spherically asymmetric structure where the effect of RS is expected to be observed (Kaastra
and Mewe, 1995). Therefore, these SNRs are ideal targets for our study.

72

Chapter 5. Evidence for RS and CX in SNRs

F IGURE 5.1: X-ray image superimposed and the radio continuum contours of N49 (Yamane et al., 2018). The X-ray image
was observed with the Chandra satellite (Park et al., 2012). The
soft gamma-ray repeater and the "bullet" ejecta are indicated bywhite arrows.

5.2

N49

5.2.1

Overview of N49

N49 is a middle-aged (∼ 4800 yr; Park et al., 2012) SNR located in the LMC.
Its radio emission shows a clear shell structure, while the X-ray emission isbright in the southeastern part of the inner radio shell (Figure 5.1). The western part has a "bullet" like the structure of ejecta (Park et al., 2012). Although
the origin of N49 is somewhat controversial, most recent research supportsthat it originated from a core-collapse explosion based on a possible association with the soft gamma-ray repeater (SGR) 0526–66 (Cline et al., 1981) andthe presence of a dense interstellar medium (ISM; e.g., Banas et al., 1997; Yamaguchi et al., 2014). The thermal X-ray emission is explained by a mixtureof two components; metal-rich ejecta and a shock-heated ISM (e.g., Uchida
et al., 2015). It is also notable that N49 is in an overionized state (Uchida etal., 2015) and is interacting with dense molecular clouds on the eastern side
(Banas et al., 1997; Yamane et al., 2018). Due to the interaction with molecular clouds, the thermal X-ray emission of N49 is particularly bright on thesoutheastern rim. N49 is an attractive object since Kaastra and Mewe (1995)
pointed out that the effect of RS can show up if an SNR has such a sphericallyasymmetric structure.

5.2.2

Observation and Data Reduction

N49 was observed with the XMM-Newton satellite (Jansen et al., 2001) in
2001 (Obs.ID 0113000201) and 2007 (Obs.ID 0505310101). The observationin 2007 was performed with a roll angle which placed N49 and a nearby
SNR, N49B, along the dispersion direction of the RGS, making it difficult to5.2. N49

73

F IGURE 5.2: (left) 12 CO( J = 1 − 0) integrated intensity maps
(Yamane et al., 2018). White contours indicate the hard X-rays
observed with Chandra. (right) Distribution of the 12 CO( J =3 − 2)/12 CO( J = 1 − 0) intensity ratio (Yamane et al., 2018).
The contours indicate broadband X-rays.

extract the RGS spectra of N49. We thus analyzed only the data obtained in
2001. For our spectral analysis, we used the RGS (Herder et al., 2001) andthe EPIC MOS (Turner et al., 2001) data. We reduced the data using XMM
Science Analysis Software version 16.1.0. The RGS data were processed with
the RGS pipeline tool rgsproc. To discard periods of background flares, weapplied good time intervals based on the count rate in CCD9 (Figure 5.3),
which is the closest to the optical axis of the mirror and most affected by the
background flares. The resulting effective exposure time is 11 ks for both
RGS1 and RGS2. Since the second-order spectra are of low statistical quality,
we analyzed only the first-order spectra.

5.2.3

Analysis

We analyzed the spectra using version 3.04.0 of the SRON SPEX package(Kaastra et al., 1996) with the maximum likelihood C-statistic (Cash, 1979;
Kaastra, 2017). The RGS spectra were fitted simultaneously with those of
MOS1 and 2. To account for the spatial broadening of the source, we multiplied spectral models with the SPEX model lpro, to which we input the
MOS1 image of the source. Our models have two absorption models: one for
the Milky Way and the other for the LMC. The column density of the former
was fixed to 6 × 1020 cm−2 (Dickey and Lockman, 1990) whereas that of thelatter is left free. The elemental abundances for the LMC absorption were
fixed to values found in the literature (∼ 0.3 solar; Russell and Dopita, 1992;
Schenck et al., 2016).

Chapter 5. Evidence for RS and CX in SNRs

10.5
01.03115×108

1.0312×108

1.03125×108

1.0313×108
TIME s

1.03135×108

1.0314×108-65.95F IGURE 5.3: Light curve of the RGS 1. The areas indicated by
the arrows are the periods of background flares.

-66.05Declination-66.00N49B

N49-66.10RATE count/s

1.5

2

74

81.60

81.45

81.30

Right ascention

F IGURE 5.4: EPIC-MOS image (0.4–8.0 keV) of N49. The crossdispersion width of the RGS (5 arcmin) is in between the whitesolid line. The spectral extraction region is enclosed by thewhite dashed line. N49B is another SNR near N49.

5.2. N49

75

Spectral Analysis with CIP Model
Figure 5.5 shows MOS1 and combined RGS1+2 spectra of N49. Prominentlines are detected at ∼12 Å (Ne X Lyα), ∼13.5 Å (Ne IX Heα), ∼15 Å (Fe XVII
Lα; 3d–2p), ∼16 Å (O VIII Lyβ), ∼17 Å (Fe XVII Lα; 3s–2p), ∼20 Å (O VII
Heβ, O VIII Lyα), ∼22 Å (O VII Heα). We first applied a two-component
NEI model (neij; Kaastra and Jansen, 1993) absorbed by neutral gas to thespectra. The NEI model consists of emissions from the overionized hot ejecta
and an ionizing cool component originating from the swept-up ISM (Uchida
et al., 2015). Although SGR 0526−66 cannot be spatially resolved with the
RGS, according to the result of Uchida et al. (2015), the emission is negligible
compared to that from N49 in the energy band covered by the RGS (0.3–
2.0 keV). This model is described as:
Neutral Interstellar Abs. × [ NEI ( ISM) + NEI (ejecta)].

(5.1)

Free parameters of the NEI components include the electron temperature(kTe ), ionization time scale (ne t, where ne and t are the electron number density and the elapsed time since shock heating or rapid cooling, respectively),and emission measure (ne nH V). In addition to these parameters, the abundances of O (=N=C), Ne, Mg, Si, S, Ar, and Fe (=Ni) of the ejecta were set free.
The neij model has another parameter kTinit , which represents an initial temperature before a rapid cooling or a shock heating of the plasma. Uchida et al.
(2015) determined kTinit for the ejecta component mainly from the ionization
state of the Fe Kα emission at ∼6.6 keV. Since the line centroid is out of the
wavelength band of the RGS data, we fixed the kTinit of the ejecta component
at 11 keV based on Uchida et al. (2015). On the other hand, the kTinit of the
ISM component is fixed at ∼ 0 keV.
Figure 5.5 shows the result of the spectral fitting with the two-component
NEI model (hereafter, NEI model). The best-fit parameters are listed in Table 5.1. The model well reproduces the overall MOS spectrum; the best-fit
parameters of kTe and ne t of the ejecta component are consistent with thoseobtained by Uchida et al. (2015). Focusing on the RGS spectrum in Figure 5.6,
however, we found discrepancies between the model and the data especiallyaround the O VII triplet, O VIII Lyβ, and Fe XVII Lα series.
Plasma Diagnostic Using O VII Heα Line Ratio
We first focus on the O VII Heα line, more specifically on its G-ratio. As mentioned in Chapter 3, the G-ratio is strongly depends on kTe (see Figure 5.7).
We assume here that the emitting plasma is in an ionizing state since mostof the O VII Heα line emission of N49 is attributed to the shocked ISM (Figure 5.5 (a)). Fitting with four Gaussians, we estimated the G-ratio of the O
VII Heα line observed in N49, which is plotted in Figure 5.7. The observed
G-ratio obviously requires an unreasonably low kTe (< 0.15 keV) with any
reasonable range of ne t, indicating necessity of some physical process to enhance the G-ratio.

Chapter 5. Evidence for RS and CX in SNRs

(3s–2p)Fe

0.4

O

Si Heα
f

0.2

O

Heα
r

Fe

S Heα

VII Heβ

r

0.3

Lyβ

Mg Heα(3d–2p)Counts/s/Å

Lyα

Fe

ejectaO(3d–2p)0.5

HeαNe X Lyα

(a)

Ne

76

i

f

0.1

χ

N Lyα

0
4
2
0
−2
−4

5

10

15

20

25

Wavelength (Å)

F IGURE 5.5: RGS +2 (black) and MOS1 (gray) spectra of N49.
Overlaid is the best-fit “NEI” model.

5.2.4

Discussion

Physical Origin of High G-ratio
Our analysis revealed that the RGS spectrum of N49 shows a high G-ratio.
As described in Chapter 3, the G-ratio is enhanced by CX and RS, inner-shellionization and recombination plasma (RP). Here, we consider the physicalprocesses that explain the high G-ratio of N49.
Inner-Shell Ionization As mentioned in Chapter 3, X-ray fluorescence due
to inner-shell ionization of Li-like ions increase the G-ratio. The ion fractionof Li-like O is expected to increase in low-ionized plasma. Therefore, reducing the ne t is expected to increase the G-ratio (Figure 5.7). A low value ofne t < 3 × 109 cm−3 s is required to explain the G-ratio with an electron temperature of 0.2 keV. t corresponds to the elapsed time since the plasma wasshock heated. Therefore, in this case, t < 100 (1 cm−3 /ne ) yr for a typicalinterstellar space density of ne = 1. Because N49 is 5000 years old, it is unreasonable that only a few decades have passed since the ISM plasma of N49
was heated. Therefore, we rejected the possibility of the contribution of X-ray
fluorescences due to inner-shell ionization of Li-like O.
RP We next examine the possibility that the ISM plasma of N49 is in an
over-ionized state (RP). Figure 5.8 shows the temperature dependence of the
G-ratio of an over-ionized plasma. In this case, there seems to be a wide

5.2. N4977

0.2

O VII Heα

+

0.4

NEI

r

Gaussians
Fe XVII Lα(3s–2p)α(3d–2p)

0.15

f
O VIIILyβ+
Fe XVIII Lα(3s–2p)0.2

Counts/s/ÅCounts/s/Å

0.3

0

Model

NEI + CX

NEI

Data

Model

0

Data

i
0.05

0.1

0.1
0.05
0
−0.05
−0.1
0.1
0.05
0
−0.05
−0.1

0.1

NEI - Gaussian

NEI

15

16

Wavelength (Å)

17

0.1
0.05
0
−0.05
−0.1
0.1
0.05
0
−0.05
−0.1

NEI + CX

NEI

NEI - Gaussian

NEI

21.5

22

22.5

Wavelength (Å)

F IGURE 5.6: Close-up views of the RGS spectrum (Figure 5.5)
around Fe XVII Lα and O VIII Lyβ (left) and O VII Heα (right).
The dark green, cyan, and red solid lines in the top panels represent the best-fit “NEI”, “NEI+CX”, and “NEI−Gaussians”
models, respectively. The middle and bottom panels show
residuals from the models with the same color scheme as the
top panels.

78

Chapter 5. Evidence for RS and CX in SNRs

2

1.75

G-ratio

1.5

1.25

1
net = 3 × 109 cm-3 s
net = 1010 cm-3 s
net = 1011 cm-3 s
net = 1012 cm-3 s

0.75

0.5

0

0.05

0.1

0.15

0.2

0.25

electron temperature (keV)

F IGURE 5.7: Relation between kTe and the G-ratio of OVII Heα.
The solid curves represent the G-ratio expected for an ionizing plasma emission predicted by the neij model in SPEX. The
colors of each line indicate ne t assumed. The red hatched area
denotes the G-ratio derived from the observed line ratio and its
statistical error. The gray lines indicate kTe from the best-fit NEI
model.

5.2. N49

79
TABLE 5.1: Best fit parameters of N49

Component
Absorption
ISM

EjectaCXParameters (unit)
NH ( MW ) (1022 cm−2 )NH ( LMC ) (1022 cm−2 )kTe (keV)ne t (1011 cm−3 s)
EM (1056 cm−3 )kTe (keV)
kTinit (keV)
ne t (1011 cm−3 )O(= C = N)
Ne
Mg
Si
S
Ar
Fe
EM (1056 cm−3 )kTe (keV)abundance
vcollision
EM (1056 cm−3 )C-statistic/d.o.f.

NEI
0.6 (fixed)
3.1 ± 0.1
0.204 ± 0.003> 10
260 ± 200.61 ± 0.01
11 (fixed)
0.5
7.2+
−0.3
+0.16
0.710.06
0.09
0.96+
−0.07
0.75 ± 0.07
0.08
0.87+
−0.06
1.2 ± 0.1
1.8 ± 0.5
0.32 ± 0.03
58 ± 44700/3211

NEI + CX
0.6 (fixed)
2.8 ± 0.1
0.230 ± 0.006
> 10
20
170+
−10
0.63 ± 0.01
11 (fixed)
0.4
7.2+
−0.3
+0.20
0.67−0.14
0.09
0.94+
−0.06
+0.07
0.72−0.06
0.08
0.85+
−0.06
1.2 ± 0.1
0.4
1.8+
−0.3
+0.03
0.29−0.02
557+−2(= value of the ISM)(= abundances of the ISM)
270 ± 110
25
47+
−13
4660/3214

NEI - Gaussians
0.6 (fixed)
3.3 ± 0.1
0.205 ± 0.003
> 10
23
322+
−22
+0.08
0.56−0.01
11 (fixed)
1.0
6.9+
−0.4
1.3 ± 0.1
1.0 ± 0, 1
0.75(fixed)
0.87(fixed)
1.2(fixed)
1.8(fixed)
0.32 ± 0.01
53 ± 0.1

4611/3190

range of parameters (kTe and ne t) that reproduce the observed G-ratio. However, as described below, when we also focused on observational spectra inwavelength bands other than the O VII Heα emission line, we found that thepossibility that the ISM is the RP can be almost ruled out. As shown in Figure 5.9, the RP model that can explain the G-ratio overestimates the O VIII
Lyα intensity. The intensity ratio of O VIII Lyα and O VII Heα reflects the ion
fraction of H-like O and He-like O ions. The O VIII Lyα / O VII Heα ratio
in the RGS spectrum of N49 does not require the high H-like O / He-like O
ratio expected by the RP model. Therefore, we also rejected the possibility
that the ISM plasma of N49 is in an over-ionized state.
CX As suggested by Uchida et al. (2019), the CX emission can be enhanced
particularly in a region where a shock is interacting with dense gas. CX is
therefore promising in the case of N49 since it is interacting with surrounding
molecular clouds (Yamane et al., 2018). We thus added a CX model (Gu et al.,2016) to the NEI model. All the abundances and kTe for the CX component
were coupled to those of the ISM component. The collision velocity vcollision(i.e. shock velocity) was allowed to vary. We assumed a multiple collisioncase in which an ion continuously undergoes CX until it becomes neutral.
This model is described as:
Neutral Interstellar Abs. × [ NEI ( ISM ) + NEI (ejecta) + CX ].

(5.2)

The best-fit “NEI+CX” model is displayed in Figure 5.10, and the parameters
are listed in Table 5.1. Although the addition of the CX component improves
the residuals at the O VII triplet, the discrepancies still remain around O VIII

80

Chapter 5. Evidence for RS and CX in SNRs

2

G-ratio

1.75

1.5

1.25

1
net = 1010 cm-3 s
net = 1011 cm-3 s
net = 1012 cm-3 s

0.75

0.5

0

0.05

0.1

0.15

0.2

0.25

electron temperature (keV)
F IGURE 5.8: Same as 5.7 but the G-ratio for RP.

0.3

5.2. N49

81
1.25

O
Lyα

Counts/s/Å

1

0.75

0.5

0.25

0
18

VII
Heα
r

O
Heβ

f
i

19

20

21

22

23

Wavalength (Å)

F IGURE 5.9: The fitting result of O VIII Lyα and O VII Heα lines
when the ISM plasma of N49 is RP.

Lyβ and Fe XVII Lα series (Figure 5.6). Therefore, we conclude that the NEI+ CX model is insufficient to describe the spectrum of N49.
RS RS is other possible process that would be responsible for the enhanced
G-ratio. The observed high Fe XVII (3s–2p)/(3d–2p) and O VIII Lyβ/α ratios
can also be caused by RS, as pointed out by previous studies (e.g., Xu etal., 2002; Hitomi Collaboration, 2018). To quantify the contribution of thescattering effect, we added negative Gaussians at wavelengths where lines of
the ISM component are prominent in the best-fit NEI model: the O VII Heαintercombination, O VII Heα resonance, O VII Heβ, O VIII Lyα, O VIII Lyβ,
Fe XVII Lα, and Ne IX Heα resonance lines. Since the oscillator strengths of
the O VII and Ne IX forbidden lines are several orders of magnitude smaller
than those of the other lines, we assume that the scattering effect is negligible
for the forbidden lines. The free parameters of the NEI components are theelectron temperature (kTe ), ionization time scale (ne t), and emission measure(ne nH V). The abundances of O (=N=C), Ne, and Fe (=Ni) of the ejecta wereset free and the abundances of Mg, Si, S, and Ar were fixed to the values
obtained in the NEI model fit. This model is described as:
Neutral Interstellar Abs. × [ NEI ( ISM) + NEI (ejecta) − Gaussians]. (5.3)
The best-fit “NEI − Gaussians” model is displayed in Figure 5.11 and thebest-fit parameters are listed in Table 5.1.
We compare transmission factors estimated based on the result from the“NEI − Gaussians” model fit with those expected for RS. The transmission
factor p is defined as
A − ∆A
,
(5.4)
p=
A

Chapter 5. Evidence for RS and CX in SNRs

(3s–2p)O

Si HeαVII Heβ

r

0.3

Lyβ

Mg Heα(3d–2p)Counts/s/Å

Fe

CX

0.4

Lyα

Fe

ejectaO(3d–2p)

0.5

Heα

(b)

Ne X Lyα

0.6

Ne

82

f

0.2

Heα
r

Fe

S Heα

O
i

f

0.1

χ

N Lyα

0
4
2
0
−2
−4

5

10

15

20

25

Wavelength (Å)(3s–2p)Lyβ
O

0.3
Si Heα

0.2

Lyα

Fe

Heα

(3d–2p)Mg HeαO

VII HeβCounts/s/Å

0.4

Fe

negative gaussians(3d–2p)0.5

NeejectaFe

Ne X Lyα

F IGURE 5.10: Same as Figure 5.5 but with the “NEI + CX”
model. The red solid line represents the CX model.

O

Heα
r

S Heα

i f

0.1

N Lyα

0

χ

−0.1
4
2
0
−2
−4

5

10

15

20

Wavelength (Å)

F IGURE 5.11: Same as Figure 5.5 but with the “NEI −Gaussians” model. The magenta solid line represents the negative Gaussians.

25

5.2. N49

83

where A is the total number of photons emitted by the plasma, and ∆A is
the number of photons scattered out of our line of sight. We can derive p foreach line from the data, given that the best-fit normalizations of the negative
Gaussians and the line intensities of the ISM component correspond to ∆A
and A, respectively.
Referring to Kaastra and Mewe (1995), we calculate the transmission factors in a case where RS effectively occurs. Under the slab approximation as asimple geometrical model of the SNR rim, Kaastra and Mewe (1995) used asingle-scattering treatment where a photon completely escapes from the lineof sight at every scattering event. As shown by Park et al. (2003), the ISM
plasma of N49 has a particularly bright emission at the southeastern rim.
Assuming that the RS occurs dominantly at the southeastern rim, we adopt
the same assumption as Kaastra and Mewe (1995). Then, the transmission
factor is written as
1
,
(5.5)
p=
1 + 0.43τ
where τ is the optical depth of the ISM plasma (Kastner and Kastner, 1990).
The optical depth τ at the line centroid is given as Equation (3.55) in Chapter
3. We assumed a thermal equilibrium between all ions and electrons and neglected the micro-turbulence velocity. The oscillator strengths and ion fractions for each element were taken from SPEX. In our case, the absorption
column density (NH ) is the only free parameter.
In Figure 5.12, we compare the transmission factors estimated from thedata and those calculated with Equation (5.5). They are roughly consistent
if NH is (3.0–10) × 1019 cm−2 , which corresponds to a plasma depth of (10–34)×(nH /cm−3 ) pc. Since the plasma depth is comparable to the diameter ofN49, ∼ 20 pc, the result supports that RS occurs at the rim of N49.
O VIII Lyβ/α ratio
The effect of RS roughly explains the observed line ratios, but as seen in Figure 5.12, OVIII Lyα requires a higher column density than the other lines. A
possible explanation would be that RS occurs also in the Galactic Halo (GH),as proposed by Gu et al. (2016). According to Nakashima et al. (2018), GH
spectra are represented by a CIE plasma model with kTe ∼ 0.26 keV. Sincesuch a plasma has a larger optical depth for the OVIII Lyα line than that of the
OVII Heα resonance line, the GH should selectively reduce the intensity of
OVIII Lyα. We applied the foreground CIE absorption model, hot (Plaa et al.,
2004; Steenbrugge et al., 2005), to the NEI − Gaussians model. The electron
temperature and Fe abundance of the hot model are fixed to 0.26 keV and
0.56 solar, respectively (the other elemental abundances are fixed to the solar
values), by referring to Nakashima et al. (2018). We found that the discrepancy cannot be reduced with the model, where the transmission factors of
OVIII Lyα and OVII Heα are calculated to be ∼ 0.43 and ∼ 0.80, respectively.
The cause of the observed high OVIII Lyβ/α ratio (the measured value
is 0.18 as an upper limit) is not clear. Uncertainties in the model of the Fe-L
lines (e.g., Gu et al., 2019) might partially explain the result since OVIII Lyβ

84

Chapter 5. Evidence for RS and CX in SNRs

1transmission factor

0.8

0.6

0.4
absorption column density: plasma depth
0 /cm2: 0 pc
1.0 × 1019/cm2: 3 (nH/cm-3) pc
3.0 × 1019/cm2: 10 (nH/cm-3) pc
6.0 × 1019/cm2: 20 (nH/cm-3) pc (∼ diameter of N49)
10 × 1019/cm2: 32 (nH/cm-3) pc
30 × 1019/cm2: 97 (nH/cm-3) pc

0.2

0

10
Ne
He
(r)

FeL
@15.0

8
FeL
@15.3

O
Ly

6
FeL
@16.8

FeL
@17.1

4O
He

O
Ly

2O
He
(r)

O
He
(i)

F IGURE 5.12: Transmission factors for each line. The squares
and arrows are the values and lower limits derived from theobservational data whereas the circles are those expected for
RS. The colors indicate NH assumed in the calculation.

05.2. N49

85

overlaps with the FeXVIII Lα line. RS by the ejecta may selectively reduce
Lyα, which we did not take into account in our analysis. We will be able toevaluate the ejecta contribution by measuring the intensity ratio of the emission lines in Fe Kα, which originates only from the ejecta (e.g., Yamaguchi
et al., 2014; Uchida et al., 2015). Since Fe Kα is out of the wavelength bandof the RGS, further observations with X-ray microcalorimeters (e.g., Kelley
et al., 2016) are required to clarify this point.

5.2.5

Conclusion of N49

We analyzed a high-resolution X-ray grating spectrum of N49 obtained with
the RGS on board XMM-Newton. We found that the G-ratio of OVII Heα in
the RGS spectrum of N49 is higher than that expected for a CIP emission.
The ratios of FeXVII Lα (3s–2p)/(3d–2p) and OVIII Lyβ/α also show large
residuals from the CIP emission model. While an extra CX component well
reproduces the G-ratio of the OVII Heα triplet, the residuals around FeXVII
Lα and OVIII Lyα still remain. On the other hand, RS can reproduce the
RGS spectrum including the G-ratio of the OVII Heα triplet and the ratios
of FeXVII Lα (3s–2p)/(3d–2p). We estimated the optical depth for the RS
from the intensities of the scattered lines and found that the depth is roughly
consistent with the size of N49. Therefore, we conclude that RS is responsible
for the high G-ratio of the RGS spectra of N49.

86

Chapter 5. Evidence for RS and CX in SNRs

F IGURE 5.13: Fitting results for the EPIC and RGS of SNR
J0453.6−6829 (Haberl et al., 2012). The detectors for each color
are pn (black), MOS1 (blue), MOS2 (cyan), RGS1 (red), and
RGS2 (green). The bottom row shows residuals.

5.3

J0453.6−6829

5.3.1

Overview of J0453.6−6829

J0453.6−6829 is a relatively compact (∼ 2′ in diameter) remnant of a corecollapse explosion (Lopez et al., 2009; McEntaffer et al., 2012), containing apulsar wind nebula (PWN) at the center of the shell (Gaensler et al., 2003).
In addition to synchrotron radiation from the PWN, McEntaffer et al. (2012)indicated that the X-ray spectrum of J0453.6−6829 obtained with Chandra iswell explained by a shock-heated ISM. A similar conclusion was reached by
Haberl et al. (2012), who performed multi-frequency observations of J0453.6−6829,
including X-ray band with XMM-Newton. In their spectral fit, the forbiddenline of O VII is seemingly higher than a normally expected by CIP model (Figure 5.13). Furthermore, J0453.6−6829 is reported to have a high dust density
in the surrounding environment (Williams et al., 2006). We thus revisit the
RGS data with particular attention to the He-like lines, in conjunction withan HI observation around the remnant, in order to investigate the relation
between the f /r ratio and the surrounding environment of J0453.6−6829.

5.3. J0453.6−6829

5.3.2

87

Observation and Data Reduction

J0453.6−6829 was observed with XMM-Newton on 2001 March 29 (Obs. IDs
0062340101 and 0062340501). Since one of the data sets (Obs. ID 0062340501)
was affected by large soft-proton flares, we present results only from Obs.
ID 0062340101. The raw data were processed with the XMM Science Analysis Software (SAS) version 18.0.0 and the calibration data files released in2020 June. In the following spectral analysis, we combine RGS 1 and RGS 2
data with MOS spectra. After discarding periods of background flares, we
obtained MOS and RGS data with effective exposure times of ∼ 6 ks and∼ 20 ks, respectively. We do not analyze second-order spectra because of
their poor statistics.

5.3.3

Analysis

Spectral Analysis with CIP Model
Figure 5.14 shows a background-subtracted three-color image of J0453.6−6829
taken by EPIC (MOS and pn). We extracted RGS spectra by limiting the crossdispersion width so as to cover the whole of J0453.6−6829. MOS spectra wereobtained from the entire region of the remnant. Off-source regions in thefield of view (FOV) were used to extract background spectra for each instrument. We simultaneously fitted the unbinned RGS and MOS spectra using
SPEX version 3.06.01 (Kaastra et al., 1996), applying a maximum likelihood
method, W-stat (Wachter et al., 1979). Throughout this analysis, same as N49,
the hydrogen column density (NH ) of the Galactic absorption was fixed to
6 × 1020 cm−2 (Dickey and Lockman, 1990) and that of the LMC was left
free. We referred to Russell and Dopita (1992) for the elemental abundancesof the LMC.
Figure 5.15 presents the MOS and RGS spectra of J0453.6−6829. McEntaffer
et al. (2012) reported that the X-ray spectrum of J0453.6−6829 can be explained by a two-temperature NEI model (neij) with a power-law component for the PWN. We first applied this "2-NEI" model, in which photon index Γ and normalization of the power-law component were fixed to 2.0 and
3.5 × 1043 photons s−1 keV−1 (McEntaffer et al., 2012), respectively. Free parameters for the thermal components include the kTe , ne t, and emission measure (ne nH V. Abundances of C, N, O, Ne, Mg, Si, and Fe were set free andtied between the two components.
The best-fit model of the 2-NEI model is plotted in Figure 5.15, and its
parameters are summarized in Table 5.2. We found that whereas the model
can reproduce the MOS spectrum, significant residuals remain in the RGS
spectrum especially at the O VII Heα and O VIII Lyβ lines. The result implies
that some additional considerations are required to better explain the finestructures of the spectrum of J0453.6−6829.

28:00.0

Chapter 5. Evidence for RS and CX in SNRs

NE

-68:30:00.0

SW

32:00.0

Declination (J2000)

88

4:54:00.0

40.0

4:53:20.0

Right Ascension (J2000)

F IGURE 5.14: Three-color image of J0453.6−6829 obtained withEPIC (MOS and pn). Red, green, and blue correspond to theenergy bands of 0.3–1.2 keV (soft), 1.2–2.0 keV (medium), and
2.0–8.0 keV (hard), respectively. The magenta lines indicate thecross dispersion width of the RGS (5’). The spectra and background are extracted from the region sandwiched by the whitelines and the region above the source region sandwiched by themagenta and white lines, respectively. The dashed line shows
how the source region is divided for a spatially-resolved analysis.

5.3. J0453.6−6829

Fe XVII L (3s-2p)

high-temperature NEI(f)

C VI Lyα(f)low-temperature NEI

NⅦ Lyα(r)(r)

C VI Lyβ,$,+6

Ne Ⅹ Lyα

Mg ⅩⅠ Heα,$,=6

Si XIII Heα9(:02;<;<!

,$7

O Ⅶ HeαO Ⅶ HeγO Ⅶ Heβ

Ne IX Heα

,$7+6

,$,6

Fe ⅩⅦ L (3d-2p)
O VIII Lyβ + Fe ⅩⅦI L (3s-2p),$76

O Ⅷ Lyα89!"#$%&''('

,
*
+
,
)+
)*

6

7,

76

+,
-./"#"0123%4!5+6

8,

F IGURE 5.15: RGS1+2 (black) and MOS1 (gray) spectra of
J0453.6−6829. The colored solid curves indicate the contributions of the low-temperature (blue) and high-temperature NEI(orange) components. The bottom panel shows residuals fromthe best-fit model.

86

90

Chapter 5. Evidence for RS and CX in SNRs
TABLE 5.2: Best fit parameters of J0453.6−6829

Parameters (unit)
2-NEI
2-NEI+CX 2-NEI−Gaus (RS)
NH(Galactic) (1020 cm−2 )
6.0 (fixed)
6.0 (fixed)
6.0 (fixed)
2.8
+2.3
5.2
NH(LMC) (1020 cm−2 )
8.7+
7.2
6.9+
−2.7
−3.5
−2.4
44
−
1
−
1
Power law (PWN)
Normalization (10 photons s keV ) 0.35 (fixed) 0.35 (fixed)
0.35 (fixed)Γ
2.0 (fixed)
2.0 (fixed)
2.0 (fixed)
28
+13
55
Low-temperature NEI Emission Measure (1058 cm−3 )
81+
60
43+
−20
−25
−13
+0.01
+0.01
+0.05
kTe (keV)
0.15−0.01
0.18−0.02
0.15−0.03
ne t (1011 cm−3 s)
> 10
> 10
> 10
0.13
+0.15
+0.47
C
0.34+
0.40
0.44
−0.10
−0.1
−0.17
+0.05
+0.07
0.18
N
0.12−0.04
0.15−0.04
0.18+
−0.06
+0.05
+0.05
+0.39
O
0.26−0.04
0.23−0.03
0.53−0.10
0.07
+0.08
0.31
Ne
0.34+
0.33
0.38+
−0.05
−0.06
−0.07
+0.08
+0.13
+0.27
Mg
0.42−0.07
0.41−0.08
0.51−0.09
0.09
+0.11
0.20
Si
0.22+
0.19
0.32+
−0.08
−0.07
−0.11
0.04
0.06
0.06
Fe
0.25+
0.22+
0.25+
−0.03
−0.03
−0.04
+5
+7.9
+10
58
−
3
High-temperature NEI Emission Measure (10 cm )12−4
7.9−3.5
17−16
0.05
+0.09
+0.33
kTe (keV)
0.42+
0.48
0.35
−0.03
−0.08
−0.08
+0.5
+0.5
11
−
3
ne t (10 cm s)
1.6−0.3
1.4−0.3
> 2.2
100
CX
Emission Measure (1058 cm−3 )
···
18+
···
−9
−
1
vcol (km s )
···
< 286
···
W-statistic/d.o.f.
4124/3627 4107/3625
4085/3622
Component
Absorption

Plasma Diagnostic Using O VII Heα Line Ratio
We first focus on the O VII Heα line. In the case of J0453.6−6829, contrary
to N49, intercombination lines of O VII Heα could not be clearly separated.
Therefore we focus on the f /r ratio. We quantified the f /r intensity ratio of
O VII by adding four Gaussians instead of multiplet line components of OVIIHeα, (i.e., resonance, forbidden and intercombination lines) implemented in
the neij code: in this method, the other lines and continua were not changed
and were the same as those of the best-fit model. We compared the obtained
value with that expected from the NEI model as indicated in Figure 5.16 (a).
0.09
The resultant f /r ratio, 1.06+−0.10 requires kTe < 0.025 keV. On the other
hand, few O6+ (He-like) ions, which emit OVII, are present in such a lowtemperature plasma (panel (b) of Figure 5.16), being inconsistent with our
result. We thus conclude that any single or multiple NEI component(s) cannot reproduce the observed RGS spectrum. Another possible scenario to account for both f /r ratio and OVIII Lyβ line is an over-ionized plasma. This
scenario, however, would make a significant excess of radiative recombination continua, and thus contradicts the observed spectrum.
Spectral Analysis with 2-NEI+CX Model
We next added a CX component to the 2-NEI model (hereafter, 2-NEI+CX
model) to enhance the forbidden line intensity of OVII. The free parametersof the CX model are normalization (nH nnh V, where nnh is the density of theneutral materials) and shock velocity (zv ). The ionization temperature wastied to the kTe of the low-temperature NEI component. The best-fit result

5.3. J0453.6−6829

F IGURE 5.16: (a) f /r ratio of OVII Heα as a function of electron
temperature kTe in the case of the best-fit value of ne t of the
low-temperature NEI component. The horizontal orange andvertical blue hatched areas indicate the observed line ratio andkTe of the low-temperature NEI component, respectively. The
curve represents theoretically expected values calculated fromthe neij model in SPEX. (b) Oxygen ionization fraction as afunction of kTe .

91

92

Chapter 5. Evidence for RS and CX in SNRs

and parameters are shown in Figure 5.17 and Table 5.2, respectively. As aresult, we successfully fitted the RGS spectrum with the 2-NEI+CX model,
except that the discrepancy between the data and model is still seen at the
OVIII Lyβ line.
Spectral Analysis with 2-NEI−Gaus Model (RS Model)
We next applied negative Gaussians in addition to the 2-NEI model (hereafter, 2-NEI−Gaus model). The Gaussians were fixed at the centroid wavelengths of the lines whose oscillator strengths are relatively large: resonancelines of Ne IX and O VII, Fe XVII Lα (3d–2p), Fe XVII Lα (3s–2p), and O VIII
Lyα. Normalizations (photons s−1 ) of these five Gaussians were set free andthe other parameters are the same as the 2-NEI model. As shown in Figure
5.17, the 2-NEI−Gaus model globally reduces the residuals. The best-fit parameters (Table 5.2) are consistent with those expected for a typical middleaged SNR. We therefore claim that a presence of RS cannot be ruled out interms of the spectral fitting.

5.3.4

Discussion

As indicated in the previous section, the high-resolution X-ray spectrum of
J0453.6−6829 suggests the presence of CX or RS in the remnant. Similar cases
have often been discussed in the literature (e.g., Uchida et al., 2019; Amano etal., 2020; Suzuki et al., 2020). Although it is in general difficult to distinguishbetween these two possibilities with the available spectroscopies, the SNRmorphology and surrounding environment may provide a clue to the true
origin of the high f /r ratio. In Table 5.3, we summarize LMC/SMC SNRs forwhich the f /r (or G) ratios were measured so far using the RGS to compare
our results with those from other SNRs in the discussion below.

TABLE 5.3: f /r ratios or G-ratios for OVII Heα

Name
1E0102−7219
N132D
DEM L71
N23 (0506−68.0)
N49
J0453

Ref.
Age Ref. f /r or G ratios
Ref.
SurroundingType∗ (yr) Age∗
Ratios∗environments
Ib/c or IIL/b
1, 2 1000
3 0.55‡±0.03
4
no dataIb
1
2500
5 0.68†±0.02
6 CO and HI clouds
Ia
10
4400
11 0.65†12
no dataII
13 4600
13 0.99‡±0.06
14 star-forming region
II
16 6600
17 1.23‡±0.12
18 CO and HI clouds
0.09
II
20 13000
21 1.06†+
22
HI clouds−0.10Type of SNe

Ref. Surrounding
environments∗
7, 8, 915
7, 8, 19
22

∗ References. (1) Blair et al. (2000); (2) Chevalier (2005); (3) Hughes et al. (2000); (4) This work. (5) Vogt and Dopita (2011); (6) Suzuki
et al. (2020); (7) Banas et al. (1997); (8) Sano et al. (2017); (9) Sano et al. (2020); (10) Hughes et al. (1998); (11) Ghavamian et al. (2003);(12) van der Heyden et al. (2003); (13) Hughes et al. (2006); (14) Broersen et al. (2011); (15) Chu et al. (1988); (16) Park et al. (2012);(17) Park et al. (2003); (18) This work. (19) Yamane et al. (2018); (20) Lopez et al. (2009); (21) Gaensler et al. (2003); (22) This work.
† f /r ratio ‡ G-ratio ( f + i ) /r

5.3. J0453.6−6829

93

O Ⅷ Lyαhigh-temperature NEI

(f)

CX

C VI Lyα(f)low-temperature NEI(r)

C VI Lyβ

$9$#1

(r)

O Ⅶ HeαNⅦ Lyα$9$1

Fe XVII L (3s-2p)
O Ⅶ HeγO Ⅶ HeßMg ⅩⅠ Heα

$9$:1
Si XIII Heα456*,7878!

$92

Fe ⅩⅦ L(3d-2p)
O VIII Lyβ+ Fe ⅩⅦI L (3s-2p)(a)
Ne Ⅹ Lyα$92#1

Ne IX Heα

$921

"
2-NEI
#
$
!#
"
!"
2-NEI+CX#$!#!"Ne IX HeαNe Ⅹ Lyα

(r)
(f)

#1

3$

31

O Ⅶ Heαlow-temperature NEI(r)high-temperature NEI(f)negative Gaussians (RS)C VI Lyα$9$1

Mg ⅩⅠ Heα$92(b)
Si XIII Heα456*,7878!

$921

Fe ⅩⅦ L (3d-2p)
O VIII Lyβ+ Fe ⅩⅦI L (3s-2p)$9##$%&'()(*+,-./!0C VI Lyβ21

NⅦ Lyα2$

O Ⅷ Lyα1

Fe XVII L (3s-2p)
O Ⅶ HeγO Ⅶ HeßRel. Error$$

Rel. Error

!$9$1
"
2-NEI
#
$
!#
!"
"
2-NEI-Gaus
#
$
!#
!"

1

2$

21

#$
%&'()(*+,-./!0#1

3$

F IGURE 5.17: (a)Same as Figure 5.15 but with the 2-NEI+CX
model. The magenta curve represents the CX component. The
middle and lower panels show residuals from the 2-NEI and 2NEI+CX models, respectively. (b)Same as Figure 5.15 but withthe 2-NEI−Gaus model. The magenta curve indicates scatteredline intensities. The middle and lower panels show residuals
from the 2-NEI and 2-NEI−Gaus models, respectively.

31

94

Chapter 5. Evidence for RS and CX in SNRs

F IGURE 5.18: Transmission factors p for each line. The black
points are those estimated from the observed line intensities
and normalizations of the negative Gaussians. The coloredcurves represent expected transmission factors as a function of
NH .

X-ray Morphology of J0453.6−6829
For quantitative evaluation of the effect of RS, we calculated a transmission
factor p by applying the same method as N49. We compared theoretical values of p for several optical depths according to Equqtion (3.40) with thoseestimated from the best-fit normalizations of the negative Gaussians. Figure
5.18 shows the result. Although the model fails to explain the observed resonance line of Ne IX, the estimated p for J0453.6−6829 are roughly consistent
with those at NH = 1.0–5.0 × 1019 cm−2 , which corresponds to a plasma
depth of 3–18 pc under an assumption of the plasma density nH = 1.1 cm−3
(Williams et al., 2006). Since the diameter of J0453.6−6829 is estimated to be∼ 36 pc from the apparent angular size of the shell (∼ 2.5′ ), the line-of-sightplasma depth that contributes to RS should be 10–50 % of the diameter. If
this is really the case, the shell-type SNR would be required to have a highly
asymmetric morphology; for instance, a bright shell is prominent only on
one side of the remnant (see Figure 5.19).
If J0453.6−6829 has an ideal spherical symmetric structure, the RS effectwill be cancelled out and an enhancement of f /r will not occur. Note that
the RGS spectrum of J0453.6−6829 was obtained from the entire region. We
can thus postulate that asymmetry of an SNR is a key parameter to evaluate
the effect of RS. From a soft-band imaging analysis by Lopez et al. (2011),

5.3. J0453.6−6829

95

F IGURE 5.19: Example schematic view of an X-ray emitting
plasma that accounts for the result of the calculation of RS.

we found that J0453.6−6829 has a less asymmetric morphology among thesix SNRs listed in Table 5.3. Other core-collapse remnants are more “elliptical” (N23) or “non-uniform” (N132D), which are parametrized as P2 /P0
and P3 /P0 in their calculation. Although N49 was not analyzed by Lopezet al. (2011), it would also be categorized as a highly elliptical remnant due
to its morphology similar to that of N23. It is reasonable that N23, N49, andN132D show relatively high f /r ratios due to RS, as claimed by previous
studies (Broersen et al., 2011; Amano et al., 2020; Suzuki et al., 2020). On theother hand, in the case of J0453.6−6829, the effect of RS might be unlikely or
insufficient to satisfactorily explain the observed high f /r ratio.
Surrounding Environment of J0453.6−6829
CX is another possibility that causes the enhancement of f /r. An interaction
with a dense ambient medium is expected in this case, as in previous studies of Galactic SNRs with the RGS: Puppis A (Katsuda et al., 2012) and the
Cygnus Loop (Uchida et al., 2019). While McEntaffer et al. (2012) implieda presence of dense gas in the vicinity of J0453.6−6829 because of a spatial
correlation between the X-ray and infrared morphologies, the surrounding
environment of this remnant has still been unclear (Williams et al., 2006; Lakićević et al., 2015). As shown in Figure 5.20, we compared the ATCA & Parkes(Kim et al., 2003) HI velocity channel map around J0453.6−6829 with the Xray morphology and found HI clouds located along with the southwestern
half of J0453.6−6829. Figure 5.21 shows the integrated intensity maps of HI.
We also found the southwestern part of the remnant is increasingly covered
with an HI cloud. The position-velocity diagram suggests that the SNR shellis expanding into the dense gas (panel (b) of Figure 5.21).
If the SNR shell interacts with the HI cloud in the southwestern region,
relatively strong forbidden line emission would be detected there. We thus
divided the data into two in the cross-dispersion direction (namely, northeast; NE and southwest; SW) as indicated in Figure 5.14. As displayed in
Figure 5.22, the forbidden line intensity of OVII in SW is stronger than thatin NE. Applying the same method as the entire region, we obtained f /r ratios
0.18
+0.40
of 0.97+−0.14 and 1.41−0.29 for NE and SW, respectively. Although statistically96

Chapter 5. Evidence for RS and CX in SNRs

they are equal within the errors, the trend is consistent with the above expectation and thus strongly supports the presence of CX. The best-fit models of
2-NEI+CX for these regions are displayed in Figure 5.23. While the CX component is required both in NE and SW, its contribution is relatively dominant
in SW. We thus confirmed that the anomalous f /r ratios is due to the CX
emission, mainly caused by an interaction with the southwestern HI cloud.
According to the discussion above, we presume that the emitting regionof CX is the southwestern edge of J0453.6−6829, which is in contact with the
HI cloud. To quantitatively examine the possibility of CX, we estimate theemitting volume VCX using the volume emission measure of the CX component EMCX = nH nnh VCX , where nnh is the neutral material density of thesurrounding gas. Given that the HI gas has a typical density of nnh = 10 cm−3 ,we obtain the emitting volume VCX to be 2 × 1058 cm3 . Since the total volumeof J0453.6−6829 is estimated to be VSNR ∼ 1060 cm3 assuming a diameter of∼ 36 pc, we conclude that the CX occurs in ∼ 0.4 % of the SNR radius. The
result fits well with the calculation by Lallement (2004) (see Chapter 3) andthus supports the possibility that the observed anomalous f /r ratio is due
to CX. Note that the significant residuals seen at ∼ 16 Å (Section 5.3.3.3) isstill an open question; such discrepancies around the OVIII Lyβ line are often
pointed out by many RGS observations (e.g., N49: Amano et al., 2020), andmight be due to uncertainties in the atomic data (see also, Plaa et al., 2012).

5.3.5

Conclusion for J0453.6−6829

We performed a high-resolution spectroscopy of J0453.6−6829 with the RGS
onboard XMM-Newton and found that the intensity of the forbidden line of
OVII is significantly stronger than expected from a simple CIP (2-NEI) model.
0.09
To account for the obtained high f /r ratio (1.06+−0.10 ), we examined two possibilities: CX and RS. Both models are statistically acceptable, although small
residuals remain at ∼ 16 Å (around the OVIII Lyβ line) between the data andthe 2-NEI+CX model. Such discrepancies are often pointed out by many RGS
observations (e.g., Amano et al., 2020; Plaa et al., 2012) and are likely due touncertainties in the atomic data. From the best-fit result with the RS model,
we estimated a transmission factor τ; the result requires a significantly asymmetric shape along the line of sight. This may be inconsistent with the apparent morphology of J0453.6−6829, since a previous systematic X-ray study indicates that this remnant is one of the "least asymmetric" core-collapse SNRs
(Lopez et al., 2011). On the other hand, our estimate of the emitting volume for the CX component (∼ 0.4 % of the SNR radius) agrees well with atheoretical expectation (Lallement, 2004). We also found evidence of an interaction between J0453.6−6829 and the dense ambient gas in the ATCA &Parkes HI map, which supports the picture that the observed f /r ratio is dueto the CX emission at SNR shock fronts. In conclusion, the presence of CX
in J0453.6−6829 is favored in our study, while a slight or significant contribution of the RS effect also cannot be ruled out. Future spatially resolved
spectroscopies with high angular resolution missions like Athena will clarify
this point.

5.3. J0453.6−6829

mage: ATCA & Parkes HI
rays

F IGURE 5.20: ATCA & Parkes HI channel maps overlaid with
the Chandra X-ray intensity of J0453.6−6829 (white contours).
Each panel shows the HI intensity map integrated over the
2.1 km s−1 width evenly spaced in the 248.7 − 273.9 km s−1
range.

97

98

Chapter 5. Evidence for RS and CX in SNRs(a) HI intensity (VLSR.: 248.2 – 273.9 km s )-68°30’36"— -68°28'12")(b) p-v diagram (Dec.: −68.51– −68.47
deg.)−1

SNR shell size

26'

270

265
VLSR [km s−1]

Declination (J2000)

30 pc

30'

260

255
K km s−1−68°34'
h

0

600

m s
04 54 30

K degree
1200

HPBW
h

m s
04 53 30
Right Ascension (J2000)

250
h

0.00

m s
04 54 30

1.05

2.10
h m s
04 53 30
Right Ascension (J2000)

F IGURE 5.21: (a) Integrated intensity map of the ATCA &Parkes HI in VLSR of 248.7 − 273.9 km s−1 . (b) Position-velocity
diagram of H I image. The integration range in Dec. is from−68◦ 30’36" to −68◦ 28’12" (J2000.0). The white arrow indicatesthe position of J0453.6−6829.

F IGURE 5.22: RGS1+2 spectra of the NE (top) and SW (bottom)
regions of J0453.6−6829.

5.3. J0453.6−6829

,$7

99

low-temperature NEINE

high-temperature NEI

9(:02;<;<!

,$,>

CX

,$,=

,$,*

,$,+

!"#$%&''('

,
*
+
,
)+
)*

6

,$7

7,

76

+,
-./"#"0123%4!5+686

low-temperature NEI

SW

high-temperature NEI,$,>

9(:02;<;<!

8,

CX

,$,=

,$,*

,$,+

!"#$%&''('

,
*
+
,
)+
)*

6

7,

76

+,
-./"#"0123%4!5+6

8,

F IGURE 5.23: Same as the panel (a) of Figure 5.15 but for theNE (top) and SW (bottom) regions.

86

100

5.4

Chapter 5. Evidence for RS and CX in SNRs

Summary

We analyzed the RGS spectra of N49 and J0453.6−6829, which are surroundedby dense neutral gas. As a result, we obtained observational evidence for CX
and RS.
The RGS spectrum of N49 shows a high G-ratio of O VII Heα as well as
O VIII Lyβ/α and Fe XVII (3s–2p)/(3d–2p) ratios, which cannot be explainedby the emission from an optically thin CIP. We thus considered the possibilityof CX and RS. Although an additional CX emission well reproduces the Gratio of the O VII Heα lines, the O VIII Lyβ/α and Fe XVII (3s–2p)/(3d–2p)
ratios cannot be explained. The effect of RS can well explain the G-ratio of
the OVII Heα and FeXVII Lα lines. We estimated the optical depth for the RS
from the intensities of the scattered lines and found that the depth is roughly
consistent with the size of N49. Therefore, we conclude that RS occurs in therim of N49.
In the case of the J0453.6−6829, we found a high f /r ratio of OVII Heαlines and OVIII Lyβ/α ratio. Both 2-NEI+CX and NEI−Gaussians (RS) models well explain the f /r ratio of OVII Heα lines although small residuals
remain at ∼ 16 Å (around the OVIII Lyβ line) between the data and the 2NEI+CX model. This discrepancy are also observed in the RGS spectrumof N49 and are likely due to uncertainties in the atomic data (Plaa et al.,2012). We focused on the SNR morphology and surrounding environment,
since it is difficult to distinguish between the possibilities of CX and RS fromthe spectral features. Estimated transmission factors τ require a significantly
asymmetric shape along the line of sight. This may be inconsistent with theapparent morphology of J0453.6−6829 (Lopez et al., 2011). On the other
hand, our estimate of the emitting volume for the CX component (∼ 0.4% of the SNR radius) agrees well with a previous calculation (Lallement,
2004). We also found evidence of an interaction between J0453.6−6829 andthe dense ambient gas in the ATCA & Parkes HI map. Furthermore, relatively strong forbidden line emission was detected in the southwestern region of J0453.6−6829 where the SNR shell interacts with the HI cloud. In
conclusion, our results strongly support the presence of CX in J0453.6−6829.
Our results help us understand the physical conditions in which CX and/or
RS occur. A detailed discussion of the physical conditions is given in Chapter
7.

101

Chapter 6Constraint on 3D Structure of
SNRs from X-ray Resonance
Scattering
In Chapter 5, we showed that CX and RS can be observed in SNRs. These processes are promising means of constraining several important parameters. In
this section, we propose a method to constrain the 3D structure of SNRs fromthe effect of RS. This method would provide us many insight into progenitors and their explosion mechanism. Additionally, we report the result of
employing this methodology to one of the brightest extra-Galactic SNRs, 1E
0102−72.9 (hereafter E0102), whose 3D structure is currently a controversial
and still open question.

6.1

MethodIn our method, the 3D structure of SNRs is constrained by measurement of
its line-of-sight length. The column density (ne L) can be estimated by thetransmission factors of RS, as we performed in Chapter 5. On the other hand,the emission measure (∝ n2e L) can be derived from the X-ray brightness ofthe plasma. We can thus constrain ne and L independently from these two
observation values.
We describe the detailed method to measure the transmission factor andthe emission measure. For simplicity, we assume that the SNR has a cylindrical structure as displayed in Figure 6.1. For nearly optically thin emission
from an isothermal plasma with a uniform density, an observed flux (F) of aline emission can be expressed as:
F (/cm2 /s) = p J (kTe , ne t)

nZ
EM,
nH

(6.1)

where p is the transmission factor, J (kTe , ne t) is the emissivity introducedin Subsection 3.1.2, nZ /nH is the elemental abundance. EM is the emission
measure described as:
EM =

ne nH
SL
4πd2

(6.2)

102

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

F IGURE 6.1: Sketch of an SNR. For simplicity, we assume thatthe structure of the SNR is cylindrical column.

∼

0.8 n2e
SL,
4πd2

(6.3)

where ne is the number density of the electron, and nH is the number densityof the hydrogen, d is the distance to the object and SL is the volume of theobject (see Figure 6.1). Here, we used a relation ne ∼ 1.2nH , which approximately describes a plasma with a solar abundance.
Below, as an example, we describe the procedure to constrain the transmission factors of the O VII Heα (r ) line from the observed fluxes of the O
VII Heα (r ) (FHeα (r) ) and O VII Heα ( f ) (FHeα ( f ) ) lines. From Equation (6.1),
the observed fluxes are described as:
FHeα (r) = pHeα (r) JHeα (r) (kTe , ne t)
FHeα ( f ) = JHeα ( f ) (kTe , ne t)

nO
EM,
nH

nO
EM
nH

(6.4)
(6.5)

where we assume the transmission factor of O VII Heα ( f ) to be unity, since
the oscillator strengths of the O VII Heα ( f ) line is much smaller than that of
O VII Heα (r ). We can determine the emissivity of each line by constraining
the electron temperature (kTe ) and ionization timescale (ne t) as we describedin Chapter 3. Therefore, when kTe and ne t, and abundance ( nnHZ ) are determined, the EM and p can be calculated as:pHeα (r) =FHeα (r) /FHeα ( f )JHeα ( f ) /JHeα (r)nH FHeα ( f )EM =.
nO JHeα ( f )

(6.6)
(6.7)

From the transmission factor p, we can estimate ne L as we performed in
Chapter 5.

6.2. Overview of E0102

6.2

103

Overview of E0102

1E 0102−72.9 (hereafter E0102) is a young (∼ 2000 yr; e.g., Banovetz et al.,2021) SNR located in the Small Magellanic Cloud (SMC). Its thermal X-ray
emission is mainly explained by metal-rich ejecta in which the X-ray emission
is dominated by emission lines from highly ionized O, Ne, and Mg (e.g.,
Gaetz et al., 2000, Rasmussen et al., 2001). The progenitor of E0102 is still
unclear. E0102 is classified to be a Type IIb SNR due to its interaction with thedense wind of the progenitor (e.g., Chevalier, 2005; Seitenzahl et al., 2018).
Therefore, its progenitor is considered to be a single WR star or a binary star.
However, the progenitor mass is not well constrained, such as M < 15 M⊙(Katsuda et al., 2018), 25-40 M⊙ (Blair et al., 2000), or M > 50 M⊙ (Finkelstein
et al., 2006), and there is no firm evidence that it was binary.
Figure 6.2 presents an X-ray image of E0102 obtained with Chandra. Although the X-ray morphology of E0102 appears to be explained by a shelllike symmetry structure at first glance, there are various suggestions about
its 3D structure. Based on the X-ray radial profile of E0102, Hughes (1994)
suggested that it has a barrel-like partial spherical shell structure rather thana projected spherical shell (Figure 6.3). Flanagan et al. (2004) found that the
spectrum of E0102 can be explained by a superposition of red and blue shifting rings from Chandra HETG observations (the left panel of Figure 6.4).
Based on this observational result, they proposed that E0102 has a cylindrical structure expanding at a velocity of approximately ±900 km/s along theline of sight (the right panel of Figure 6.4).
We analyze RGS data, aiming to constrain the 3D structure of E0102 andprovide implications about its progenitor. E0102 is a relatively compact (∼ 1′
in diameter) SNR, which makes it possible to perform high-resolution X-ray
spectroscopy with the RGS. Furthermore, we can obtain high statistical data
since E0102 has been observed more than 50 times for the on-board calibration of XMM-Newton. These high-quality data enable us to perform robust
plasma diagnostics and measurements of the effect of RS.

6.3

Observations and Data Reduction

As a calibration source for XMM-Newton, E0102 has been observed several
times each year from 1999 to 2021. We use most of the RGS observations
except several data sets whose effective exposures are extremely short. The
basic information on the data we use are listed in Table 6.1. The following
reduction tasks were done with the Science Analysis System version 16.1.0.
The RGS data were processed with the pipeline tool rgsproc.
In this study, we analyze a spectrum in which all observations are stacked.
We used the publicly available script rgscombine in the Science Analysis System to combine all source and background spectra respectively. The total effective exposure time is ∼ 3 Ms for both the first- and second-order spectra.
To account for the spatial broadening of E0102, we convolved the response
files for point-like sources with the profiles of the surface brightness along
the dispersion direction using the ftools, ftrgsrmfsmooth. This process was

104

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

F IGURE 6.2: RGB image of E0102 obtained with Chandra. Red,
green and blue colors correspond to the energies 0.5–0.6, 1.0–1.4
and 4.0—7.0 keV.

performed for each observation respectively. To obtain the profiles of E0102,
we generated Chandra images in the six energy bands of 0.32–0.6 keV, 0.6–0.7
keV, 0.7–0.85 keV, 0.85–0.95 keV, 0.95–1.2 keV, 1.2–2.0 keV. After this convolution, we constructed the response matrix for the stacked grating spectrum by
averaging the response matrices of each observation weighted by the exposure time. We use version 12.10.0c of the XSPEC software for the following
spectral analysis. We fit the first- and second- order spectrum of RGS data
simultaneously.

6.3. Observations and Data Reduction

F IGURE 6.3: (left) Radially-averaged X-ray surface brightnessof E0102 obtained with ROSAT high resolution imager and thebest-fit spatial model (Hughes, 1994). The short-dash curve isthe barrel-like component, the long-dash curve is the spherical shell component, the dotted curve is the clumped component, and the background is the straight line. The bottom panel
shows the difference between the observed data and the bestfit model. (Right) Cross-sectional view of the geometric model.
The observer is in the direction of the vertical axis. The barrellike component is shown as the doubly hatched region, and thespherical shell component is shown singly hatched region. Circles near the bottom of the figure correspond to the clumped
components.

F IGURE 6.4: (Left) Ne X Lyα image of E0102 obtained by Chandra HETG (Flanagan et al., 2004). The red area corresponds tothe line-of-sight velocity of +900 km/s, and the green area corresponds to -900 km/s. (Right) 3D structure of E0102 proposedby Flanagan et al. (2004).

105

106

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

F IGURE 6.5: EPIC-MOS image (0.4–2.0 keV) of E0102. The
cross-dispersion width of the RGS (5 arcmin) is in between thewhite solid line. The spectral extraction region is enclosed by
the white dashed line.

6.3. Observations and Data Reduction

107

TABLE 6.1: Information on the XMM-Newton observations of
E0102 used in the analysis.
Obs. ID
0123110201
0123110301
0135720601
0135720801
0135720901
0135721001
0135721101
0135721301
0135721401
0135721501
0135721701
0135721901
0135722001
0135722101
0135722201
0135722301
0135722501
0135722701
0412980201
0412980301
0412980701
0412980801
0412980901
0412981001
0412981301
0412981401
0412981501
0412981701
0412981701
0412981701
0412982101
0412982201
0412982401
0412982501
0412983201
0412983301
0412983401
0412983501
0791580701
0791580801
0791580901
0791581001
0791581101
0791581201
0810880201
0810880301
0810880501

Obs. Start Date(yyyy-mm-dd)
2000-04-16
2000-04-17
2001-04-14
2001-12-25
2002-04-20
2002-05-18
2002-10-13
2002-12-14
2003-04-20
2003-10-27
2003-11-16
2004-04-28
2004-10-26
2004-11-07
2004-11-07
2004-11-07
2005-04-17
2006-04-20
2007-04-25
2007-10-26
2008-11-14
2009-04-13
2009-10-21
2010-04-21
2010-10-19
2011-04-20
2011-11-04
2012-12-06
2012-12-06
2012-12-07
2013-11-07
2014-10-20
2015-10-30
2015-10-28
2016-10-26
2016-12-03
2017-11-02
2017-11-04
2016-04-26
2016-04-26
2016-04-26
2016-04-26
2016-04-27
2016-04-27
2018-10-28
2019-11-18
2019-11-04

Exposure Time
(ks)
33.6
18.7
16.6
50.4
8.3
17.6
29.9
54.2
50.9
33.9
42.3
55.0
59.8
45.0
11.6
27.4
41.9
58.3
35.2
60.9
53.1
19.7
55.0
54.1
33.0
26.8
57.6
17.1
32.8
34.3
56.7
59.9
60.1
63.5
54.3
56.0
78.4
61.7
57.6
23.1
20.8
57.3
26.3
13.8
49.2
58.0
37.6

108

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

F IGURE 6.6: RGS spectra of E0102. Black and red are the firstorder and of second-order spectra of the RGS, respectively. The
detected lines are labeled and listed in Table 6.2 with their identification.

6.3. Observations and Data Reduction

109

TABLE 6.2: Detected Emission Lines and their Identification

Label
1
2
3
4
5
6
7
8
9
10
11
12
13
14
15
16
17
18
19
20
21
22
23
24
25
26

Energy
(keV)
0.329
0.341
0.347
0.368
0.381
0.407
0.420
0.426
0.431
0.436
0.443
0.451
0.459
0.470
0.500
0.511
0.561
0.569
0.574
0.654
0.666
0.698
0.713
0.725
0.739
0.775

Ion

TransitionLabel

S XIII
S XII
Ca XI
C VI
S XIV
S XIV
N VI
N VI
N VI
C VI
Ca XIIAr XIV
C VI
C VI
N VII
S XIV
O VII
O VII
O VII
O VIII
O VII
O VII
O VII
Fe XVIIFe XVIIO VIII

3s → 2p
3d → 2p
3s → 2p
2p → 1s
3d → 2p
3p → 2s
2s → 1s (f)2p → 1s (i)2p → 1s (r)3p → 1s
3d → 2p
3d → 2p
4p → 1s
5p → 1s
2p → 1s
4d → 2p
2s → 1s (f)2p → 1s (i)2p → 1s (r)2p → 1s
3p → 1s
4p → 1s
5p → 1s
3s → 2p
3s → 2p
3p → 1s

27
28
29
30
31
32
33
34
35
36
37
38
39
40
41
42
43
44
45
46
47
48
49
50
51

EnergyIon
Transition
(keV)
0.817
O VIII
4p → 1s
0.826
Fe XVII
3d → 2p
0.837
O VIII
5p → 1s
0.847
O VIII
6p → 1s0.873 Fe XVIII
3d → 2p
0.905
Ne IX
2s → 1s (f)
0.922
Ne IX
2p → 1s (r)
0.964
Fe XX
3d → 2p
1.022
Ne X
2p → 1s
1.056 Fe XXIII
3d → 2p
1.076
Ne IX
3p → 1s
1.127
Ne IX
4p → 1s
1.152
Ne IX
5p → 1s
1.168 Fe XXIV
3p → 2s
1.211
Ne X
3p → 1s
1.277
Ne X
4p → 1s
1.314
Fe XXI
4d → 2p
1.343
Mg XI
2s → 1s (f)1.352
Mg XI 2p → 1s (r)1.473
Mg XII
2p → 1s
1.579
Mg XI
3p → 1s
1.659
Mg XI
4p → 1s
1.696
Mg XI
5p → 1s
1.745
Mg XII
3p → 1s
1.865
Si XIII 2p → 1s (r)110

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

6.4

Analysis

The purpose of our analysis is to measure the transmission factors, according to the method described in Section 6.1. The analysis is carried out inthe following procedure. First, we measure the flux of several emission lines
accurately by using the "modified NEI" model described later. Next, we constrain the kTe and ne t by plasma diagnostics introduced in Subsection 3.4.1.
Finally, we measure the transmission factors of He α (r ) and Lyα lines of O
and Ne based on the He α (r ) / He α ( f ) and Ly α / Ly β ratios.

6.4.1

Spectral Modeling

Figure 6.6 shows the background-subtracted RGS spectra. K-shell emission
lines of He-like and H-like ions are clearly detected. In addition to theselines, we detect a dozen L-shell emissions from Fe and the intermediate-masselements (S, Ar, Ca) and K-shell emissions from highly excited states of O
and Ne. To measure the fluxes of individual lines accurately, contributionsof continuum emission and other weak lines cannot be ignored. Therefore,
we first tried to roughly reproduce the RGS spectrum using NEI models toestimate the contributions of these emission components. As indicated byprevious studies (e.g., Alan et al., 2019), the spectra of E0102 consist of emission from multiple-temperature plasmas. We thus applied two componentsof velocity-broadened, collisionally ionized plasma models, bvrnei in the
XSPEC that can reproduce both CIE and NEI plasmas. The free parameters
are the electron temperature (kTe ), ionization parameter (ne t), abundances of
C, N, O, Ne, Mg, Si, S, Ar, Ca, Fe (= Ni), velocity dispersion (σ), the redshift(or blueshift), and the emission measure. Our foreground absorption models consist of two tbnew models: one for the Milky Way and the other for
the SMC. The column density of the former was fixed to 5.36 × 1020 cm−2
(Dickey and Lockman, 1990) whereas that of the latter is left free. The elemental abundances for the SMC absorption were fixed to values found in theliterature (∼ 0.2 solar; Russell and Dopita, 1992). The best-fit model is shownin Figure 6.7, the best-fit parameters are listed in Table 6.3. We assume that
N abundances of both NEI models are 0 because that became small values(∼ 10−3 ) in our analysis. Although this model is not statistically acceptable,
the RGS spectra are roughly reproduced by this model.

6.4.2

Plasma Diagnostics and Transmission Factor Measurement

We measure the fluxes of several lines to constrain the transmission factors
by using the NEI model obtained above. The fluxes of the He α resonance,
forbidden, and Heβ of O VII and Ne IX, and Lyα and Lyβ lines of the O
VIII and Ne X are directly measured with the following procedure. A similar
procedure was performed by Suzuki et al. (2020). We remove all the above
emission lines from the NEI models by modifying one of the model data
files named apec_v3.0.9_nei_line.fits in the XSPEC package (hereafter6.4. Analysis

F IGURE 6.7: RGS 1+2 spectra of E0102. Overlaid is the bestfit NEI model. The colored solid curves indicate the contribution of the individual components, low-kTe (blue) and high-kTe(green) plasmas.

111

112

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering"modified NEI" model). The best-fit two-temperature NEI model obtained in
Subsection 6.4.1 is replaced with a model consisting of the modified NEI plus
ten Gaussians for the Heα resonance, forbidden, and Heβ and Lyα and Lyβlines of O and Ne. All the parameters (e.g., electron temperatures, elemental
abundances) of the modified NEI model components are fixed to the best-fit
values of the original NEI model. Then we fit the Gaussian components, andobtaining the line fluxes are listed in Table 6.4.
We estimate the kTe and ne t based on the line fluxes obtained above. We
utilize the fluxes of He α forbidden, Heβ, and Lyβ lines. Since these lines
have relatively small oscillator strengths, we can derive kTe and ne t without considering the effects of RS. Figure 6.8 (a-1) shows O VII Heβ/α ( f )flux ratio map calculated with AtonDB database and the observed flux ratio.
The observed O VII Heβ/α ( f ) ratio indicates that these lines are emitted
by the plasma with an electron temperature of 0.38–0.6 keV. Figure 6.8 (a-2)
shows O Lyβ/Heα ( f ) ratio map. The solid and dashed lines indicate theobserved Lyβ/Heα ( f ) and Heβ/α ( f ) flux ratios, respectively. The overlapof the contours indicates that oxygen lines are emitted by the plasma with
an electron temperature of 0.40 ± 0.01 keV and an ionization time scale of(7.1 ± 0.2) × 1010 cm−3 s. We applied the same procedure to Neon (Figure(b-1), (b-2)). The obtained kTe and ne t are summarized in Table 6.5.
Since the emissivity of each line has been constrained, we can derive thetransmission factors from the line flux ratios as described in Section 6.1. The
emissivity ratios calculated using the kTe and ne t obtained above are summarized in Table 6.5. From Equation (6.8), the transmission factor of Heα (r)and Lyα can be expressed as:pHeα (r) =pLyα =

FHeα (r) /FHeα ( f )JHeα ( f ) /JHeα (r)FLyα /FLyβ.
JLyβ /JLyα

(6.8)
(6.9)

By applying the emissivity ratios in Table 6.5 and the line fluxes in Table 6.4,
the transmission factors are calculated as summarized in Table 6.6.

6.4. Analysis

F IGURE 6.8: Theoretical predictions for line ratios of (a-1)
Heβ/Heα ( f ) of O, (a-2) Lyβ/Heα ( f ) of O, (b-1) Heβ/Heα ( f )of Ne, (a-2) Lyβ/Heα ( f ) of Ne, as a function of ionization timescale (ne t: horizontal axis) and electron temperature (kTe : vertical axis). The black solid lines indicate the regions constrainedby observed line ratios. The black dashed lines indicate the regions constrained by observed Heβ/Heα ( f ). We constrain theelectron temperature and ionization time scale of O and Ne by
the overlap of the contours.

113

114

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

TABLE 6.3: Best fit parameters of E0102

Component
Absorption

Parameters (unit)
NH ( MW ) (1020 cm−2 )NH (SMC ) (1020 cm−2 )NEI (low-kTe ) kTe (keV)
kTinit (keV)
ne t (1011 cm−3 )
C
N
O
Ne
Mg
Si
S
Ar
Ca
Fe
σ (103 km/s)
vredshift
EM (10−3 cm−3 )NEI (high-kTe ) kTe (keV)
kTinit (keV)
ne t (1011 cm−3 )
C
N
O
Ne
Mg
Si
S
Ar
Ca
Fe
σ (103 km/s)
vredshift
EM (1056 cm−3 )C-statistic/d.o.f.

Value
5.36 (fixed)
0.010
6.95+
−0.0040.375 ± 0.001
0.01 (fixed)
5.27 ± 0.01
0.62 ± −0.01
0 (fixed)
0.04
1.66+
−0.01
0.09
2.78+
−0.07
0.75 ± 0.07
0.035
0.679+−0.0360.028
0.724+−0.029+0.0390.805−0.041
0.040
0.260+−0.0410.201 ± 0.020
1.21 ± 0.01
1210 ± 10
0.02
4.30+
−0.06
1.15 ± 0.01
0.01 (fixed)
0.02
7.01+
−0.05
0.16
0.71+
−0.06
0 (fixed)
1.27 ± 0.01
0.09
2.67+
−0.07
1.60 ± 0.01
0.401 ± 0.013
1.2 ± 0.1
1.8 ± 0.5
1.8 ± 0.5
0.001
0.294+−0.002
1.17 ± 0.01
7
−997+
−10
0.04
4.30+
−0.02
36047/5772

TABLE 6.4: Observed line fluxes of E0102

O
Ne

Heα (r)(cm2 /s)(4.15 ± 0.01) 10−3
(1.39 ± 0.01) 10−3Heα ( f )(cm2 /s)(1.96 ± 0.01) 10−3
(6.23 ± 0.04 10−4Heβ(cm2 /s)(6.94 ± 0.04) 10−4
(2.66 ± 0.02) 10−4

Lyα
(cm2 /s)(5.62 ± 0.01) 10−3
(1.33 ± 0.01) 10−3

Lyβ(cm2 /s)(7.96 ± 0.03) 10−4
(1.82 ± 0.01) 10−4

6.4. Analysis

115

TABLE 6.5: Obtained kTe and ne t, and line emissivity ratiosONe

kTe (keV)
0.40 ± 0.01
0.74 ± 0.01

ne t (1010 cm−3 s)
7.1 ± 0.2
7.8 ± 0.2

Heα ( f )/Heα (r )0.037
0.446+−0.0200.361 ± 0.006

Lyβ/Lyα0.130 ± 0.008
0.134 ± 0.003

TABLE 6.6: Transmission factors

Line ID
O VII Heα (resonance)
O VIII Lyα
Ne IX Heα (resonance)
Ne X LyαTransmission Factor
0.077
0.943+−0.0430.054
0.915+−0.0580.806 ± 0.014
0.982 ± 0.018

116

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

6.5

Discussion

In this section, we examine the validity of 3D structures proposed in previous
studies by comparison between the transmission factors expected from the
proposed structure and those obtained from our analysis.

6.5.1

Cylindrical Model

We examine the possibility of a cylindrical structure proposed by Flanagan
et al. (2004). Based on their observations, the 3D structure of E0102 is determined as follows. Flanagan et al. (2004) claims that E0102 has the cylindricalstructure expanding at ±900 km/s along the line of sight. Since the age of
E0102 is estimated to be 1738 ± 175 yr (Banovetz et al., 2021), we assume
that E0102 is in free expansion at ±900 km/s for 1738 years. Therefore, theline-of-sight length of E0102 is about 3.2 pc (Figure 6.9). The inner and outer
diameter of E0102 are calculated as 4.2 pc and 5.4 pc from the X-ray image,assuming the distance to E0102 is 62 kpc (Graczyk et al., 2020). Since the 3D
structure of E0102 has been determined, we can estimate the electron column
density and the transmission factor from the observed line flux.
We describe the procedure to estimate the transmission factor of the O
VII Heα (r ) line. First, the electron column density is calculated from thevolume of E0102 and the observed line flux of the O VII Heα ( f ). By substituting the values of FHeα ( f ) in Table 6.4 and JHeα ( f ) constrained in subsection
6.4.2 into Equation (6.5), the emission measure ((0.8 n2e )/(4πd2 )V) is (3.32 ±0.01) × 1011 (cm−5 ), where we assume the oxygen abundance (nO /nH ) to be
2.7 based on the result of our spectral modeling obtained in Subsection 6.4.1.
The volume of E0102 is 3.40 × 1057 (cm3 ) from Figure 6.9, so the electron density is 7.30 ± 0.02 (cm−3 ), from the emission measure obtained above. Finally,
we obtain the electron column density to be 7.29 ± 0.01 (cm−2 ) from the electron density and the line-of-sight length of E0102.
Next, we estimate the transmission factor of O VII Heα (r ) in a similar
way as we did in Chapter 5. We adopt the same assumption as Kaastra andMewe (1995). Then, the transmission factor is written as Equation (5.5). The
optical depth τ at the line centroid is given as Equation (3.55) in Chapter 3.
We neglect the micro-turbulence velocity and assume a thermal equilibrium
between all ions and electrons, thus the ion temperature (TkeV ) is equal to theelectron temperature obtained in Subsection 6.4.2 (0.4 keV). The oscillator
strengths and ion fractions for each element are taken from Atomdb. In thecase of E0102 the electron column density is (7.29 ± 0.01) × 1020 (cm−2 ).
We also estimate the transmission factors of the Lyα line in a similar way
to that of the O VII Heα (r ) line. Figure 6.10 and 6.11 show the transmission factors of O and Ne lines as a function of the electron column density.
The range of column densities correspond to the blue shaded area, and thetransmission factors corresponding to these column densities are shown inthe cyan shaded area. The orange shaded areas indicate the transmission
factors obtained from our analysis summarized in Table 6.6. As can be seenfrom Figure 6.10 and 6.11, the cylindrical structure proposed by Flanagan

6.5. Discussion

117

F IGURE 6.9: Cross sectional view of E0102. Here E0102 has thecylindrical structure proposed by Flanagan et al. (2004).

et al. (2004) requires lower transmission factors for each line than those obtained in our analysis. Therefore, we conclude that the structure proposedby Flanagan et al. (2004) is inconsistent with our observational results.

118

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

F IGURE 6.10: Relation between the transmission factor and theelectron column density of (a) O VII Heα (r ), and (b) O VIII
Lyα. The black solid curves represent the transmission factor
calculated from Equation (5.5).

6.5. Discussion

F IGURE 6.11: Same as Figure 6.10 but for Ne lines.

119

120

6.5.2

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

Barrel-Like Structure

Another possible structure of E0102 is a partial spherical shell with a barrellike shape (Figure 6.12 (a)), which was suggested by Hughes (1994). Assuming the distance to E0102 is 62 kpc (Graczyk et al., 2020), their proposed structure has an inner diameter of 3.7 pc (1.14 × 1019 cm) and an outer diameter of
5.2 pc (1.60 × 1019 cm), as shown in Figure 6.12 (a). In this case, the volumeof E0102 is 5.5 × 1057 cm3 . As in Subsection 6.5.1, the electron density can becalculated using the volume of E0102 and the line fluxes.
To estimate the electron column density, we calculate the mean line-ofsight length (Lmean ) of E0102. Figure 6.12 (b) shows the area enclosed by the
red dashed line in Figure 6.12 (a). We define xi every 0.01 pc between xmin
and xmax , and calculate the line-of-sight length of E0102 (2 × li ) for each xi .
Using li , the mean line-of-sight length (Lmean ) of E0102 is expressed as:
Lmean =

1 n
li .
n i∑
=1

(6.10)The mean line-of-sight length (Lmean ) is 4.1 pc, which requires a higher electron column density than that of the cylindrical model calculated in Subsection 6.5.1. The results of calculating the transmission factors from the electron
column density obtained above are shown in Figures 6.13 and 6.14. Similar
to the cylindrical model, the barrel-like structure requires lower transmission factors than those of our observations for all emission lines. Therefore,
we conclude that the barrel-like structure proposed by Hughes (1994) is also
inconsistent with our observations.

6.5. Discussion

F IGURE 6.12: (a) Cross sectional view of E0102. Here E0102
has the barrel-like structure proposed by Hughes (1994). The
area enclosed by the red dashed line corresponds to the area
shown in the right panel of Figure 6.3. (b) Area enclosed by the
red dashed line in (a). We calculate the length li of E0102 along
the line-of-sight direction for each xi , and then determine theaverage value Lmean to obtain the electron column density.

121

122

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

F IGURE 6.13: Same as Figure 6.10, but for the barrel-like structure.

6.5. Discussion

123

Transmission Factor

1.0

0.8

0.6

0.4

(a)
Ne IX Heα (r)

0.2

1017

1018

1019

1020

1021

1018

1019

1020

1021

Transmission Factor

1.0

0.8

0.6

0.4

0.2

(b)
Ne X Lyα
1017

Electron column density (/cm2)
F IGURE 6.14: Same as Figure 6.11, but for the barrel-like structure.

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

124

(a)

(b)
-900 km/s

observer+900 km/s

observer1.2 pc1.2 pc1.2 pc1.2 pc1.2 pcF IGURE 6.15: (a) Cross sectional view of single ring structure.
(b) Cross sectional view of double ring structure.

6.5.3

Structures That Explain Our Observation

In this subsection, we propose possible structures that explain our observations. The structures proposed by previous studies require lower transmission factors, which is inconsistent with our observations. For the two structures described below, relatively high transmission factors is expected. We
consider that one of these structures is the plausible 3D structure for E0102.
Double-Ring Structure
According to the results in Subsection 6.5.1, a cylindrical structure with alarge line-of-sight length (∼ 3.2 pc) requires a stronger scattering effect thanthat we have observed in Subsection 6.4.2. In the case of a ring-like structure, as shown in Figure 6.15 (a), the thickness of E0102 in the line-of-sightdirection is equal to the thickness in the direction perpendicular to the lineof-sight direction. For such symmetrical structures, the RS effect is canceled
out, which could explain our observations. However, such a single ring-likestructure is inconsistent with the observation by Flanagan et al. (2004). We
have estimated that the line-of-sight length of E0102 is 3.2 pc in Subsection
6.5.1, assuming it is in free expansion at the current observed expansion velocity by Flanagan et al. (2004). Since it is unlikely that E0102 is in accelerating expansion, this value can be considered a lower limit for the line-of-sightlength of E0102. Therefore, we propose a double ring structure as shown in
Figure 6.15 (b) as a structure to explain both Flanagan et al. (2004) and our
results. If E0102 had such a structure, the observed scattering effect would
be slight and could explain the Doppler velocity (±900 km/s) observed by
Flanagan et al. (2004).
Cylindrical Model Considering Velocity Structure
The kinetic motion of ions can reduce the cross section and optical depth of
RS. In Subsection 6.5.1, we considered only thermal motion as the kinetic motion. However, we found that the cylindrical structure proposed by Flanagan

6.5. Discussion

125

et al. (2004) could also account for our observations by taking into accountthe bulk motion of ions due to the cylinder’s expansion in the line-of-sightdirection.
We estimate the effect of the line-of-sight expansion velocity on the optical depth of each emission line. Assuming a uniform density, the line-of-sightvelocity is proportional to the line-of-sight distance, as shown in Figure 6.16.
We use Figure 6.17 to consider the optical depth for photons emitted at arbitrary point A in E0102. For example, the cross section at point B, which is 0.3
pc away from point A, can be estimated as follows. In the point A stationary
system, the photon energy distribution of the O VII Heα emitted at A is represented by the solid black line in Figure 6.17 (b). Since the Doppler broadening due to the thermal motion of ions is significantly larger than the natural
width of O VII Heα (r ), the energy distribution of photons is expressed by aGaussian distribution characterized by the variance:
∆hν =



2kTi
m i c2

1/2
hν0 ,(6.11)

where hν0 is the rest frame photon energy, mi is the ion mass, and Ti is theion temperature. In this case, the cross section of the ion present at point B
is represented by the solid red line in Figure 6.17 (b). Of the photons emitted
at point A, those scattered at point B are included in σ2 in Figure 6.17 (b). We
thus define σ2 as the effective cross section at point B.
Figure 6.18 shows the relationship between the distance from where thephoton is emitted and the effective cross section at that location. At a distance of more than 0.6 pc away from where the photon is emitted (point A),
the effective cross section (σ2 ) is less than 0.01 times σ1 + σ2 . Since the optical depths of E0102 obtained in Subsection 6.5.1 are approximately 3 at most(see the middle column of Table 6.7), RS in the area where the cross section
becomes 0.01 times is negligible. Therefore, we adopt 0.6 pc as the upper
limit of effective line-of-sight length for E0102. The right column of Table 6.7
summarizes the optical depth τ of each emission line when the line-of-sightlength is 0.6 pc. The reduction in cross section due to the expansion significantly decreases the optical depth of each line, making all lines optically thin.
Therefore, such an effect may explain our observations.
TABLE 6.7: Optical depthLine ID
O VII Heα (resonance)
O VIII Lyα
Ne IX Heα (resonance)
Ne X Lyα

w/o velocity
2.92
3.01
1.01
0.603

w velocity
0.548
0.579
0.189
0.113

126

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering-900 km/s

900 km/s

observer3.2 pc

km/s

900

-1.6

1.6

-900

F IGURE 6.16: Cross sectional view of a cylindrical structure andthe velocity structure. The upper panel shows the cylindricalstructure proposed by Flanagan et al. (2004). The lower panel
shows the relationship between line-of-sight distance and velocity. The horizontal and vertical axis corresponds to the distance from the observer and the expansion velocity of E0102 inthe line-of-sight direction, respectively.

pc

6.5. Discussion

127

(a)
-900 km/s

900 km/s

observer3.2 pcB

A

0.3 pcNormalized Cross Section

(b)
3.0

2.0

σ1
1.0

σ2
573.5

574.0

574.5

Energy (eV)F IGURE 6.17: Line-of-sight distance dependence of RS crosssection. (a) Same as the upper panel of Figure 6.16. We consider the cross section at point B of a photon emitted at point
A. (b) Photon energy distribution ans RS cross section of the O
VII Heα(r ). The solid black line represents the photon energy
distribution of the O VII Heα emitted at A, in the point A stationary system. The solid red line represents the RS cross section at point B. σ1 + σ2 corresponds to the cross section without
considering the Doppler shift due to expansion. Of the photons
emitted at point A, the photons contained in σ2 are scattered.

128

Chapter 6. Constraint on 3D Structure of SNRs from X-ray Resonance
Scattering

1.0

σ2 / (σ1 + σ2)

0.8

0.6

0.4

0.2

0.01
0.0
0.05

0.1

0.5

1.0

Distance (pc)
F IGURE 6.18: Relationship between cross section and line-ofsight distance. The horizontal axis corresponds to the distance
between two points in E0102. The vertical axis corresponds to
the ratio of the cross section without considering the Doppler
shift due to expansion to the effective cross section.

6.6. Summary

6.5.4

129

Progenitor of E0102

We discuss the implications of the structures proposed in Subsection 6.5.3
for the progenitor of E0102. We proposed a double-ring structure or a cylindrical structure with the expansion velocity, both of which are asphericallysymmetric but have an axisymmetric structure. The key to the formation ofsuch structures of SNRs is binary interactions, which create disk-like CSMs
and bipolar cavities around SNRs as explained in Chapter 2. When an SN explodes in such a CSM structure, SNR can be elongated along the polar axis,
as demonstrated by Blondin et al. (1996). This may explain the cylindricalstructure of E0102. Multi-ring structures have been observed in other SNRssuch as N103B and SN1987A (e.g., Yamaguchi et al., 2021), whose progenitors are also considered to be binary stars, and their ring-like structures arethought to have been formed by CSMs structure resulted from binary interactions. Therefore, our results and previous studies suggest that the progenitorof E0102 may also be a binary star.

6.6

Summary

We presented a method to constrain the 3D structure of SNRs using the effects of RS and applied this method to the SNR E0102. The transmission
factors of the O and Ne lines were estimated to be larger than 0.79. The structures proposed by previous studies require lower transmission factors, whichis inconsistent with our results. Therefore, we have proposed a double-ring
structure or a cylindrical structure with the expansion velocity as structuresthat are consistent with both our observations and previous studies. Theseaxisymmetric structures imply that the progenitor of E0102 could be a binary
star. As we will explain in Chapter 7, this method can be applied to manycelestial objects using the future satellite XRISM.

131

Chapter 7

Discussion
7.1

Summary of Our Results and Previous Studies

In this section, based on our results and previous studies, we discuss thephysical conditions in which CX and/or RS efficiently take place, and the
key spectral signature to distinguish between the CX and RS. This information is useful for precise spectroscopy that properly takes into account thecontribution of the CX and RS.

7.1.1

Physical Conditions in which RS and/or CX Efficiently
Occur
TABLE 7.1: Summary of previous and our RGS stidies

Object name
E0102
N132D
Puppis A
DEM L71
N23
N49J0453.6−6829
Cygnus Loop
G296.1−0.5

SN typeIb/c or IIL/b
Ib
IIL/b
Ia
II
II
II
II
Ib/c

Ref.
1,2
1
2
14
17
20
23
26
30

age
1000–1600
2500
3700
4400
4600
5000
13000
14000
20000

Ref.
Environment
3, 4
no data6
CO and H I clouds
11
H I clouds
15
no data17
Star-forming region
20
CO and HI clouds
24
HI clouds
27 near by clouds interaction region
31
HI clouds

Ref.

RS or CX
RS
7, 8, 9
RS?
12
CX
RS?
18
RS?
7, 8, 21
RS
24
CX
28
CX and RS?
32
CX

Ref.
5
10
13
16
19
22
25
29
32

References. (1) Blair et al. (2000); (2) Chevalier (2005); (3) Hughes et al. (2000); (4) Banovetz et al. (2021); (5) This work (Chapter
6). (6) Vogt and Dopita (2011); (7) Banas et al. (1997); (8) Sano et al. (2017); (9) Sano et al. (2020); (10) Suzuki et al. (2020); (11)Winkler et al. (1988); (12) Reynoso et al. (1995); (13) Katsuda et al. (2012); (14) Hughes et al. (1998); (15) Ghavamian et al. (2003);(16) van der Heyden et al. (2003); (17) Hughes et al. (2006); (18) Chu et al. (1988); (19) Broersen et al. (2011); (20) Park et al. (2012);(21) Yamane et al. (2018); (22) This work (Chapter 5). (23) Lopez et al. (2009); (24) Gaensler et al. (2003); (25) This work (Chapter
5). (26) Miyata et al. (1998); (27) Levenson et al. (1998); (28) Patnaude and Fesen (2005); (29) Uchida et al. (2019); (30) Castro et al.
(2011); (31) Hwang et al. (1994); (32) Tanaka et al. (2022);

In Table 7.1, we summarize SNRs observed with the RGS to date to compare our results with other SNRs. Based on the result of N49, RS tends to beobserved in an asymmetric structure of SNRs with a large column density.
Therefore, for example, at the rim of SNRs, the effect of RS is expected to beobserved.
Our spatially-resolved spectral analysis of J0453.6−6829 has revealed asignificant enhancement of CX X-ray emission in the southwest region where

132

Chapter 7. Discussion

the SNR interact with the dense H I gas. Furthermore, we estimate the volume of the CX-emitting region to be roughly 0.4 % of the entire SNR volume.
CX X-ray emission is thus expected to be observed in the outer edge of theshock front in contact with dense clouds. In addition to J0453.6−6829, CX
emissions have been observed in other SNRs, such as Puppis A (Katsuda
et al., 2012), Cygnus Loop (Uchida et al., 2019), and G296.1−0.5 (Tanaka etal., 2022). All of these studies are based on observations of partially bright
structures in the galactic SNRs. These bright structures are considered to beformed as a result of collisions between SNRs and the surrounding dense
neutral gases (e.g., Hwang et al., 2005). Hence, our result is in line with theresults of these previous studies.
It is worth noting that most of the SNRs in which CX has been observed
are older than 10000 years. This tendency can be attributed to the dependence of the cross section of CX on the collision velocity between ions andneutral hydrogen. According to Equation (2.31) in Chapter 2, the shock wave
velocity of the young SNR is relatively high, and consequently, the collisionvelocity between ions and neutral atoms at the shock front is expected to belarge. Figure 7.1 shows the relationship between the collision velocity andthe total cross section of CX. The cross sections of CX are relatively large ina collision velocity below 2000 km/s and decrease sharply above 2000 km/s.
Therefore, it is important to consider the effect of CX in the spectroscopy of
evolved SNRs.

7.1.2

Spectral Feature to Distinguish between CX and RS

Both CX and RS have similar observational signatures such as a high f /r
ratio, making it difficult to distinguish between these two processes spectroscopically. However, we found that the (3s–2p)/(3d–2p) ratio of Fe XVII
Lα lines is a useful spectral feature for distinguishing between the two processes. Because of the following reasons, in typical SNR plasmas, the (3s–2p)/(3d–2p) ratio is sensitive to RS but rarely affected by CX. In the electron
temperature range (0.1–0.5 keV) of typical SNR plasmas, most Fe atoms arein the Ne-like ionization state and rarely in the Na-like ionization states (seeFigure 3.3). The Fe XVII (n = 3 → 2) lines are emitted via CX described as:
Fe17+ + H → Fe∗16+ + H+ → Fe16+ + hν + H+ .

(7.1)

Therefore, in typical SNR plasmas where most Fe elements are in a Ne-likeionization state (Fe16+ ), Fe XVII Lα lines are scarcely emitted by CX. On theother hand, since the Fe XVII Lα lines are scattered by Ne-like Fe, typical SNR
plasmas have large optical depths of RS for Fe XVII Lα lines. Since the (3d–2p) line has a larger oscillator strength than the (3s–2p) line (see Table 7.2), the(3s–2p)/(3d–2p) ratio of Fe XVII Lα lines is enhanced by RS. Therefore, a high(3s–2p)/(3d–2p) ratio of Fe XVII Lα lines is an exclusive spectral feature of
RS. In fact, a high (3s–2p)/(3d–2p) ratio was observed in the RGS spectrum of
N49, but not in the spectrum of J0453.6−6829, which provided an importantclue to distinguish between CX and RS.

7.1. Summary of Our Results and Previous Studies

133

Ar18+
Ne10+
O8+
B5+

F IGURE 7.1: Total cross sections of CX, as functions of collisionvelocity, for B5+ (black), O8+ (red), Ne10+ (green), and Ar18+(blue) interacting with a hydrogen atom (Gu et al., 2016).

134

Chapter 7. Discussion
TABLE 7.2: Oscillator Strength of Fe XVII Lα lines (taken from
Atomdb)Transition
2s2 2p5 3d (1 D1 ) → 2s2 2p6 (1 S0 ); (3d–2p)2s2 2p5 3s (1 P1 ) → 2s2 2p6 (1 S0 ); (3s–2p)

7.2

Centroid Energy (Wavelength) Oscillator Strength
0.826 keV (15.02 Å)
2.49
0.725 keV (17.05 Å)
0.126

Effect on Elemental Abundances Measurement
and Progenitor Estimation

Several previous studies pointed out that ignoring the effects of CX and RS
can lead to bias in parameter measurements (e.g., Kaastra and Mewe, 1995).
Based on the results of N49 and J0453.6−6829, we evaluate the impact of CX
and RS on parameter measurements and the progenitor estimation based onthese parameters. In our results, we found that the impact on the measurement of oxygen abundance is particularly large. By comparing the best-fit
parameters of the NEI and NEI-Gaussians models in Table 5.1, we find thatthe oxygen abundance of N49 is underestimated by a factor of ∼1.8, if we
do not take into account the RS effect. Katsuda et al. (2018) point out that
the Fe/O ratio is sensitive to the progenitor mass. We thus use the Fe/O ratio shown in Katsuda et al. (2018) to evaluate the impact on progenitor massestimation based on these parameters.
From the best-fit parameters of the NEI model in Table 5.1, without considering RS, the Fe/O ratio of N49 is measured to be about 0.45. On the other
hand, a model that takes into account the RS effect (NEI-Gaussians model)
measures the Fe/O ratio at 0.25. These Fe/O ratios correspond to CO coremasses (MCOcore ) of about 3.2 and 4.5, respectively (Figure 7.2 (a)). As argued
by Katsuda et al. (2018), there is a good linear relation between MCOcore andprogenitor mass (MZAMS : zero age main sequence mass) at a mass range of
MZAMS < 40 M⊙ (Figure 7.2 (b)). Therefore, the progenitor mass can be estimated using the CO core mass obtained above. As shown in Figure 7.2 (b),without considering RS, the progenitor mass of N49 is estimated to be about
16 M⊙ , and with RS, it is estimated to be about 20 M⊙ .
It is worth noting that N49 contains the magnetar candidate SGR 0526˘66(Cline et al., 1981). Although the origin of magnetars is still unclear, they may
have been created by very massive progenitors. For example, SGR 1806−20
and CXO J164710.2−455216 are considered to be produced by progenitorswith masses of M > 40 M⊙ (Muno et al., 2005; Figer et al., 2005). The progenitor mass of the AXP 1E 1841−045 and SGR 1900+14 are estimated to be∼ 20 M⊙ (Kumar et al., 2014) and ∼ 17 M⊙ (Davies et al., 2009), respectively.
The progenitor mass of ∼ 20 M⊙ obtained with the model considering RS isin line with the observational studies of progenitors of magnetars.

7.3. Constraint on 3D Structure of SNRs Using XRISM
(a)

135

(b)

without RS
with RSF IGURE 7.2: (a) Fe/O abundance ratios as a function of the COcore mass (Katsuda et al., 2018), based on the nucleosynthesis
models by Sukhbold et al. (2016). The abundances are relative
to the solar values Anders and Grevesse (1989). The Fe/O ratio
obtained by the NEI model (without considering RS) and thecorresponding CO core mass are shown in blue. The red arrow
indicates those obtained by the NEI-Gaussinas model. (b) COcore mass as a function of a progenitor mass, MZAMS (Sukhbold
et al., 2016). The red and blue arrows are the same as in (a).

7.3

Constraint on 3D Structure of SNRs Using XRISMIn Chapter 6, we proposed a method to constrain the 3D structure of SNR
plasmas using the effect of RS. This method can be applied to only compact
SNR with RGS. However, using this method with the XRISM satellite, we candetermine the detailed 3D structure of many objects including diffuse SNR in
the Milky Way. In this section, using the observation of W49B as an example,
we explain how to apply this method to galactic SNRs by XRISM.
The XRISM satellite is the 7th Japanese X-ray astronomical satellite, whichis to be launched in 2023. The XRISM satellite carries two types of X-ray detectors, an X-ray microcalorimeter array (Resolve) and an X-ray CCD camera(Xtend). The Resolve achieves high-resolution spectroscopy (FWHM: 5–7 eV)in the 0.3–12 keV bandpass. Since Resolve is a non-dispersive spectrometer,
it is possible to perform spatially resolved high-resolution spectroscopy ofthe diffuse SNR in the Milky Way.
Our target W49B is a core-collapse SNR located in the Milky Way. The
X-ray spectrum of W49B has strong emission lines from He-like and H-likeions, which indicate that the X-ray emission originates from the metal-richejecta. The morphology of W49B is elongated from east to west and is characterized by a bar-like distribution of iron surrounded by a relatively broad
distribution of silicon (Figure 7.3). Although the physical origin of the X-ray
morphology of W49B is still uncertain, two scenarios have been proposed:a jet-driven core-collapse supernova (hereafter "jet scenario"; Keohane et al.,2007; Lopez et al., 2013) or a spherically symmetric SN that expanded into an
inhomogeneous ISM (hereafter "shell scenario"; Zhou et al., 2011). If W49B is
the remnant of a jet-like explosion, it is a valuable sample of the remnants of

136

Chapter 7. Discussion

explosions associated with hypernovae and gamma-ray bursts.
To distinguish between the above two scenarios, we suggest a XRISM observation of the eastern edge of W49B. If the morphology of W49B is a resultof interaction with the inhomogeneous ISM, the eastern edge would correspond to the rim of the shell with a relatively long line-of-sight length asshown in the right panel of Figure 7.5. On the other hand, in the case of jet
scenarios, this eastern edge region may correspond to the jet-induced ringlike structures, which has been demonstrated by González-Casanova et al.
(2014). In the case of a simple structure shown in Figure 7.5, it is possible
to distinguish the two scenarios based on the RS effect. Figure 7.6 (a) represents the simulated spectrum of W49B where the spectral extraction regionis shown by the red square in Figure 7.5. The thermal emission was modeledaccording to the spatially resolved spectral analysis performed by HollandAshford et al. (2020). We investigate the effect of RS on the Si and Fe Kα linesof the simulated spectrum in both scenarios. Assuming the abundances ofheavy elements and ion temperatures (equal to electron temperature) as determined by Holland-Ashford et al. (2020), we calculate the optical depth of
Si Heα (r ), Si Lyα, and Fe Heα (r ) lines according to Equation (3.55). The lineof-sight length of the eastern edge of W49B is assumed to be 10 pc from the
X-ray image in the shell scenario and 1 pc from simulation results obtained
by González-Casanova et al. (2014) in the jet scenario. Table 7.3 summarizes
the optical depths of each emission line, assuming a typical electron densityof 10 cm−3 . In the jet scenario, the optical depths are roughly 0.1, whereas, inthe shell scenario, the depths are approximately 1, which implies that the RS
effect can be observed only in the shell scenario. Figure 7.6 (b) and (c) represent the spectra expected to be observed in the two scenarios. We utilize thesame assumptions as Kaastra and Mewe (1995) to calculate the transmission
factors. Significant scattering effects are expected to be observed from both
Si and Fe lines in the shell scenario. Especially, since Fe is the abundant element, the intensity of Fe XXV Heα (r ) is expected to be strongly reduced by
RS. Therefore, the Fe Heα line is particularly useful in separating these two
scenarios.
TABLE 7.3: Optical depths of resonance lines of Si and Fe

Line IDSi XIII Heα (resonance)
Si XIV Lyα
Fe XXV Heα (resonance)

jet scenario
0.07730.06250.143

shell scenario
0.773
0.625
1.43

7.3. Constraint on 3D Structure of SNRs Using XRISMF IGURE 7.3: X-ray image of W49B observed with Chandra
ACIS (Lopez et al., 2009). The left panel shows three color image in the soft band (0.3–1.50 keV; blue), medium band (1.5–2.5
keV; green), and hard band (2.5–8.0 keV; red). The middle andright panels represent X-ray images of Fe XXV and Si XIII, respectively.

F IGURE 7.4: (left) Density structure of SN simulated by Tominaga (2008). Color of the marks represents the abundance of themass element (H: yellow, He: cyan, O+C: green, O+Mg: blue,
Si: magenta, and Fe: red). A large amount of iron is synthesized in the region along the jet axis. (middle) Fe XXV X-ray
emission maps of W49B simulated by González-Casanova et al.
(2014). (right) Same as the middle panel but Si XIV map.

137

138

Chapter 7. Discussion

F IGURE 7.5: (right) X-ray full-band (0.5–8.0 keV) exposurecorrected image of W49B as observed for 220 ks by Chandra
ACIS (Lopez et al., 2013). The white dashed line corresponds tothe field of view of Resolve, and each pixel corresponds to one
pixel of Resolve. The spectral extraction region is shown by the
red square. (middle) Schematic image of W49B in the shell scenario. In this case, the spectral extraction region (the easternedge) corresponds to the rim of the SNR. (right) Schematic image of W49B in the jet scenario. The spectral extraction region
corresponds to the ring-like structure of the SNR.

7.4

Origin of Excess around 16 Å (0.78 keV)The RGS spectra of N49 and J0453.6−6829 show the excess around 16 Å(see(a-1) and (a-2) panels in Figure 7.7), which cannot be explained by the effect of
CX and RS. Similar discrepancies have been reported for other objects (Figure
7.7 (b-1) to (b-4)), and we attribute these discrepancies to uncertainties in theatomic data related to Fe-L emission lines, such as oscillator strength andrate coefficients of each atomic process (introduced in Chapter 3). The FeL emission from Fe XVI to Fe XXIV is observed in the wavelength rangeof 14-18 Å, which is emitted through a series of atomic processes, such as
collisional excitation, radiative recombination, dielectron recombination, andinner-shell ionization. Therefore, comprehensive atomic data covering theseprocesses are required to calculate the emissivities of Fe-L lines. Most ofthese atomic data have been obtained in theoretical calculations with only
a few benchmarks established with laboratory measurements (e.g., Gu etal., 2022). It has been pointed out that there is a discrepancy between theatomic data obtained by theoretical calculations and the data obtained byexperiments, which lead to uncertainty of emissivities of Fe-L lines (e.g.,
Kühn et al., 2022). Since there are Lα lines of Fe XVIII around 16 Å, thediscrepancy may be due to uncertainties in the Fe XVIII lines. Significant
excess around 16 Å is not observed in the spectrum of E0102, which contains
few Fe L-lines in its spectrum. This confirms our suggestion. To solve these
problems, it is necessary to experimentally measure atomic data. Groundexperiments, such as EBIT are indispensable for the future development of
X-ray astronomy.

7.4. Origin of Excess around 16 Å (0.78 keV)F IGURE 7.6: (a) Simulated Resolve spectrum of W49B. The exposure time is assumed to be 300 ks. (b) Close-up view of thesimulated Resolve spectrum around the Si Kα lines expected tobe observed in the two scenarios. Jet and shell scenarios areshown in black and red, respectively. The effect of RS reduces
the intensity of Heα(r ) and Lyα in the shell scenario. (c) Sameas (b) but Fe Heα lines.

139

140

Chapter 7. Discussion(a-1) SNR: N49(a-2) SNR: J0453-6829

0.4

Fe XVII(3d-2p)Fe XVII(3s-2p)

0,05

Fe XVIIIO VIII Lyβ

Fe XVIIIO VIII Lyβ

0,04

Counts/s/ÅCounts/s/Å

0.3

Fe XVII(3s-2p)Fe XVII(3d-2p)

0,06

0.2

0.1

0,03

0,02

0,01

0

0
15

16

Wavalength (Å)

17

18

−0,01

14,5

15

15,5

16

16,5

17

17,5

Wavelength (Å)(b-1) Elliptical Galaxy: NGC5813(b-2) SNR: N23

Fe XVII (3d-2p)

Fe XVII (3s-2p)

Fe XVIII(b-3) Stellar Binary: Capella(b-4) Elliptical Galaxy: NGC1404

Wavelength (Å)

F IGURE 7.7: (a): Close-up views of the RGS spectrum around
Fe XVII Lα and O VIII Lyβ obtained by our studies. There are
discrepancies between the data and the NEI model around the
OVIII Lyβ line (∼ 16 Å). (a-1) RGS spectrum of N49. The overlaid black solid line is the best-fit “NEI” model. The solid blueand green lines correspond to the ISM and ejecta components,
respectively. (a-2) RGS spectrum of J0453.6−6829. The coloredsolid curves indicate the contributions of the low-temperature(blue) and high-temperature NEI (orange) components. (b):Same as (a), but obtained by previous studies. Each panel
shows the RGS spectrum of (b-1) elliptical galaxy NGC5813(Plaa et al., 2012), (b-2) SNR N23 (Broersen et al., 2011), (b3) stellar binary Capella (Audard et al., 2001), elliptical galaxy
NGC1404 (Mernier et al., 2022), respectively.

18

141

Chapter 8

Conclusions
We performed a high-resolution X-ray spectroscopy of N49, J0453.6−6829
and E0102 with the RGS onboard XMM-Newton. The important results are
summarized as follows:• We obtained observational evidence for CX and RS from N49 and J0453.6−6829,
based on the quantitative evaluation of their effects.
• Based on our results and previous studies, RS seems to be observed in
an asymmetric structure of SNRs with a large column density, such as
rim of SNRs.
• On the other hand, CX X-ray emission is particularly enhanced at theouter edge of the shock front (a few percent of the shock radius) incontact with dense clouds. Moreover, most of the SNRs in which CX
has been observed are older than 10000 years. This tendency may be
attributed to the dependence of the cross section of CX on the collisionvelocity between ions and neutral hydrogen. Therefore, evolved SNRs
are suitable targets to search for CX X-ray emissions.
• We underestimate the O abundance by about a factor of 1.8 if we do nottake into account the RS effect, which leads to an underestimation of
the progenitor mass by about 75 % . This demonstrates the importanceof RS in measuring elemental abundances.
• We have presented a method to constrain the 3D structure of SNRs by
measuring their line-of-sight length through the effects of RS. We haveapplied this method to the SNR E0102 located in the SMC. The structures proposed by previous studies require strong scattering, which isinconsistent with our results. Therefore, we have proposed a doublering structure and a cylindrical structure with the expansion velocity as
structures that are consistent with both our observations and previous
studies. This method can be applied to many celestial objects using thefuture satellite XRISM.
• The RGS spectra of N49 and J0453.6−6829 show the excess around 16 Å,which cannot be explained by the effect of CX and RS. We attributethese discrepancies to uncertainties in the atomic data related to Fe-L
emission lines. To solve these problems, it is necessary to experimentally measure atomic data. Ground experiments, such as EBIT are indispensable for the future development of X-ray astronomy.

143

Bibliography
Alan et al. (2019). “A Detailed Archival CHANDRA Study of the Young
Core-collapse Supernova Remnant 1E 0102.2-7219 in the Small Magellanic
Cloud”. In: The Astrophysical Journal 873.1, p. 53.
Aldrovandi et al. (1973). “Radiative and dielectronic recombination coefficients for complex ions”. In: Astronomy and Astrophysics 25, p. 137.
Amano et al. (2020). “Evidence for Resonance Scattering in the X-Ray Grating
Spectrum of the Supernova Remnant N49”. In: The Astrophysical Journal897.1, p. 12.
Anders and Grevesse (1989). “Abundances of the elements: Meteoritic andsolar”. In: Geochimica et Cosmochimica acta 53.1, pp. 197–214.
Audard et al. (2001). “The XMM-Newton view of stellar coronae: High-resolution
X-ray spectroscopy of Capella”. In: Astronomy & Astrophysics 365.1, pp. L329–
L335.
Banas et al. (1997). “Supernova remnants associated with molecular clouds
in the large magellanic cloud”. In: The Astrophysical Journal 480.2, p. 607.
Banovetz et al. (2021). “The Center of Expansion and Age of the Oxygen-rich
Supernova Remnant 1E 0102.2-7219”. In: The Astrophysical Journal 912.1,
p. 33.
Blair et al. (2000). “Hubble Space Telescope observations of oxygen-rich supernova remnants in the magellanic clouds. II. Elemental abundances in
N132D and 1E 0102.2–7219”. In: The Astrophysical Journal 537.2, p. 667.
Blondin et al. (1996). “Axisymmetric circumstellar interaction in supernovae”.
In: The Astrophysical Journal 472.1, p. 257.
Broersen et al. (2011). “The high resolution X-ray spectrum of SNR 0506-68
using XMM-Newton”. In: Astronomy & Astrophysics 535, A11.
Burrows and Vartanyan (2021). “Core-collapse supernova explosion theory”.
In: Nature 589.7840, pp. 29–39.
Burrows et al. (2012). “An investigation into the character of pre-explosion
core-collapse supernova shock motion”. In: The Astrophysical Journal 759.1,
p. 5.
Cash (1979). “Parameter estimation in astronomy through application of thelikelihood ratio”. In: The Astrophysical Journal 228, pp. 939–947.
Castro et al. (2011). “AN XMM-NEWTON STUDY OF THE BRIGHT, NEARBYSUPERNOVA REMNANT G296. 1–0.5”. In: The Astrophysical Journal 734.2,
p. 86.
Chevalier (1974). “The evolution of supernova remnants. Spherically symmetric models”. In: The Astrophysical Journal 188, pp. 501–516.
— (2005). “Young core-collapse supernova remnants and their supernovae”.
In: The Astrophysical Journal 619.2, p. 839.

144

Bibliography

Chieffi and Limongi (2004). “Explosive Yields of Massive Stars from Z= 0 to
Z= Z⊙ ”. In: The Astrophysical Journal 608.1, p. 405.
Chu et al. (1988). “Environments and populations of supernova remnants in
the Large Magellanic Cloud”. In: The Astronomical Journal 96, pp. 1874–
1883.
Cline et al. (1981). “Precise source location of the anomalous 1979 March 5
gamma ray transient”. In.
Collaboration, Hitomi (2017). “Solar abundance ratios of the iron-peak elements in the Perseus cluster”. In: Nature 551.7681, pp. 478–480.
Conselice et al. (2001). “On the nature of the NGC 1275 system”. In: The Astronomical Journal 122.5, p. 2281.
Cravens (2002). “X-ray emission from comets”. In: Science 296.5570, pp. 1042–
1045.
Davies et al. (2009). “The progenitor mass of the magnetar SGR1900+ 14”. In:
The Astrophysical Journal 707.1, p. 844.
Dickey and Lockman (1990). “HI in the Galaxy”. In: Annual review of astronomy and astrophysics 28, pp. 215–261.
Dwarkadas et al. (2007). “The evolution of supernovae in circumstellar windbubbles. II. Case of a Wolf-Rayet star”. In: The Astrophysical Journal 667.1,
p. 226.
Ezoe et al. (2021). “High-resolution X-ray spectroscopy of astrophysical plasmas with X-ray microcalorimeters”. In: Reviews of Modern Plasma Physics
5, pp. 1–43.
Figer et al. (2005). “Massive stars in the SGR 1806–20 cluster”. In: The Astrophysical Journal 622.1, p. L49.
Filippenko (1997). “Optical spectra of supernovae”. In: Annual Review of Astronomy and Astrophysics 35.1, pp. 309–355.
Finkelstein et al. (2006). “Optical structure and proper-motion age of theoxygen-rich supernova remnant 1E 0102–7219 in the Small Magellanic
Cloud”. In: The Astrophysical Journal 641.2, p. 919.
Flanagan et al. (2004). “Chandra High-Resolution X-Ray Spectrum of Supernova Remnant 1E 0102.2–7219”. In: The Astrophysical Journal 605.1, p. 230.
Gaensler et al. (2003). “Discovery of a New Pulsar Wind Nebula in the Large
Magellanic Cloud”. In: The Astrophysical Journal 594.2, p. L111.
Gaetz et al. (2000). “Chandra X-Ray Observatory Arcsecond Imaging of theYoung, Oxygen-rich Supernova Remnant 1E 0102.2–7219”. In: The Astrophysical Journal 534.1, p. L47.
Ghavamian et al. (2003). “The physics of supernova blast waves. I. kinematicsof DEM L71 in the large Magellanic cloud”. In: The Astrophysical Journal590.2, p. 833.
Gillaspy (2001). “Highly charged ions”. In: Journal of Physics B: Atomic, Molecular and Optical Physics 34.19, R93.
González-Casanova et al. (2014). “The morphology and dynamics of jet-driven
supernova remnants: the case of W49B”. In: The Astrophysical Journal Letters 781.2, p. L26.

Bibliography

145

Graczyk et al. (2020). “A Distance determination to the Small Magellanic
Cloud with an accuracy of better than two percent based on late-type
eclipsing binary stars”. In: The Astrophysical Journal 904.1, p. 13.
Gu et al. (2016). “Plasma code for astrophysical charge exchange emission at
X-ray wavelengths”. In: Astronomy & Astrophysics 588, A52.
— (2019). “X-ray spectra of the Fe-L complex”. In: Astronomy & Astrophysics
627, A51.
— (2022). “X-ray spectra of the Fe-L complex-III. Systematic uncertainties inatomic data”. In: Astronomy & Astrophysics 664, A62.
Haberl et al. (2012). “Multi-frequency observations of SNR J0453–6829 in the
LMC-A composite supernova remnant with a pulsar wind nebula”. In:
Astronomy & Astrophysics 543, A154.
Hanke et al. (2013). “SASI activity in three-dimensional neutrino-hydrodynamics
simulations of supernova cores”. In: The Astrophysical Journal 770.1, p. 66.
Herder, Den et al. (2001). “The reflection grating spectrometer on board XMMNewton”. In: Astronomy & Astrophysics 365.1, pp. L7–L17.
Hitomi Collaboration (2018). “Measurements of resonant scattering in the
Perseus Cluster core with Hitomi SXS”. In: Publications of the Astronomical Society of Japan 70.2, p. 10.
Holland-Ashford et al. (2020). “Spatially Resolved Study of Recombining
Plasma in W49B Using XMM-Newton”. In: The Astrophysical Journal 903.2,
p. 108.
Hughes (Jan. 1994). “ROSAT HRI Observations of Magellanic Cloud Supernova Remnants (I)”. In: The Soft X-ray Cosmos. Ed. by Eric M. Schlegel andRobert Petre. Vol. 313. American Institute of Physics Conference Series,
p. 144. DOI: 10.1063/1.46727.
Hughes et al. (1998). “ASCA X-ray spectroscopy of large magellanic cloudsupernova remnants and the metal abundances of the large magellanic
cloud”. In: The Astrophysical Journal 505.2, p. 732.
— (2000). “Electron Heating and Cosmic Rays at a Supernova Shock from
Chandra X-Ray Observations of 1E 0102.2–7219”. In: The AstrophysicalJournal 543.1, p. L61.
— (2006). “The Chandra View of the Supernova Remnant 0506–68.0 in the
Large Magellanic Cloud”. In: The Astrophysical Journal 645.2, p. L117.
Hwang et al. (1994). “An X-ray study of five supernova remnants in the Carina spiral arm”. In: The Astrophysical Journal 431, pp. 819–825.
— (2005). “CHANDRA X-ray observation of a mature cloud-shock interaction in the bright eastern knot region of Puppis A”. In: The AstrophysicalJournal 635.1, p. 355.
Itoh and Masai (1989). “The effect of a circumstellar medium on the X-ray
emission of young remnants of Type II supernovae”. In: Monthly Noticesof the Royal Astronomical Society 236.4, pp. 885–899.
Janev and Winter (1985). “State-selective electron capture in atom-highly charged
ion collisions”. In: Physics Reports 117.5-6, pp. 265–387.
Jansen et al. (2001). “XMM-Newton observatory-I. The spacecraft and operations”. In: Astronomy & Astrophysics 365.1, pp. L1–L6.

146

Bibliography

Kaastra (2017). “On the use of C-stat in testing models for X-ray spectra”. In:
Astronomy & Astrophysics 605, A51.
Kaastra and Jansen (1993). “A spectral code for X-ray spectra of supernova
remnants”. In: Astronomy and Astrophysics Supplement Series 97, pp. 873–
885.
Kaastra and Mewe (1995). “Optical depth effects in the X-ray emission fromsupernova remnants”. In: Astronomy and Astrophysics 302, p. L13.
Kaastra et al. (Jan. 1996). “SPEX: a new code for spectral analysis of X & UV
spectra.” In: UV and X-ray Spectroscopy of Astrophysical and Laboratory Plasmas, pp. 411–414.
— (2008). “Thermal radiation processes”. In: Space Science Reviews 134, pp. 155–
190.
Kastner and Kastner (1990). “Opacity ratios for Doppler-broadened lines from
common upper levels”. In: Journal of Quantitative Spectroscopy and Radiative Transfer 44.2, pp. 275–288.
Katsuda et al. (2008). “Asymmetric ejecta distribution of the Cygnus Looprevealed with Suzaku”. In: Publications of the Astronomical Society of Japan
60.sp1, S107–S114.
— (2011). “Possible charge-exchange X-ray emission in the Cygnus Loop detected with Suzaku”. In: The Astrophysical Journal 730.1, p. 24.
— (2012). “High-resolution X-Ray Spectroscopy of the Galactic Supernova
Remnant Puppis A with XMM-Newton/RGS”. In: The Astrophysical Journal 756.1, p. 49.
— (2018). “Progenitor Mass Distribution of Core-collapse Supernova Remnants in Our Galaxy and Magellanic Clouds Based on Elemental Abundances”. In: The Astrophysical Journal 863.2, p. 127.
Kawasaki et al. (2002). “ASCA observations of the supernova remnant IC
443: thermal structure and detection of overionized plasma”. In: The Astrophysical Journal 572.2, p. 897.
Kelley et al. (2016). “The Astro-H high resolution soft x-ray spectrometer”. In:
Space Telescopes and Instrumentation 2016: Ultraviolet to Gamma Ray. Vol. 9905.
SPIE, pp. 229–245.
Keohane et al. (2007). “A near-infrared and X-ray study of W49 B: a windcavity explosion”. In: The Astrophysical Journal 654.2, p. 938.
Khokhlov et al. (1999). “Jet-induced explosions of core collapse supernovae”.
In: The Astrophysical Journal 524.2, p. L107.
Kifonidis et al. (2003). “Non-spherical core collapse supernovae-I. Neutrinodriven convection, Rayleigh-Taylor instabilities, and the formation andpropagation of metal clumps”. In: Astronomy & Astrophysics 408.2, pp. 621–
649.
Kim et al. (2003). “A neutral hydrogen survey of the Large Magellanic Cloud:aperture synthesis and multibeam data combined”. In: The AstrophysicalJournal Supplement Series 148.2, p. 473.
Kühn, Steffen et al. (2022). “New measurement resolves key astrophysicalFe XVII oscillator strength problem”. In: Physical Review Letters 129.24,
p. 245001.

Bibliography

147

Kumar et al. (2014). “CHANDRA AND XMM-NEWTON STUDY OF THE
SUPERNOVA REMNANT KES 73 HOSTING THE MAGNETAR 1E 1841–045”. In: The Astrophysical Journal 781.1, p. 41.
Lakićević et al. (2015). “The influence of supernova remnants on the interstellar medium in the Large Magellanic Cloud seen at 20–600 µm wavelengths”. In: The Astrophysical Journal 799.1, p. 50.
Lallement (2004). “On the contribution of charge-exchange induced X-ray
emission in the ISM and ICM”. In: Astronomy & Astrophysics 422.2, pp. 391–
400.
Levenson et al. (1998). “Panoramic views of the Cygnus Loop”. In: The Astrophysical Journal Supplement Series 118.2, p. 541.
Lisse et al. (1996). “Discovery of X-ray and extreme ultraviolet emission fromcomet C/Hyakutake 1996 B2”. In: Science 274.5285, pp. 205–209.
Liu et al. (2012). “Charge-exchange X-ray emission of nearby star-forming
galaxies”. In: Monthly Notices of the Royal Astronomical Society 420.4, pp. 3389–
3395.
Lopez et al. (2009). “Typing supernova remnants using X-ray line emission
morphologies”. In: The Astrophysical Journal 706.1, p. L106.
— (2011). “Using the X-ray morphology of young supernova remnants toconstrain explosion type, ejecta distribution, and chemical mixing”. In:
The Astrophysical Journal 732.2, p. 114.
— (2013). “The galactic supernova remnant W49B likely originates from ajet-driven, core-collapse explosion”. In: The Astrophysical Journal 764.1, p. 50.
Lotz (1968). “Electron-impact ionization cross-sections and ionization rate coefficients for atoms and ions from hydrogen to calcium”. In: Zeitschrift für
Physik 216.3, pp. 241–247.
Masai (1994). “Nonequilibria in thermal emission from supernova remnants”.
In: The Astrophysical Journal 437, pp. 770–780.
McEntaffer et al. (2012). “SNR 0453-68.5: An Asymmetric Remnant and its
Plerion in the Large Magellanic Cloud”. In: The Astrophysical Journal 756.1,
p. 17.
Mernier et al. (2022). “The cycle of metals in the infalling elliptical galaxy
NGC 1404”. In: Monthly Notices of the Royal Astronomical Society 511.3,
pp. 3159–3178.
Mewe et al. (1978). “Heliumlike ion line intensities. I-Stationary plasmas. IINon-stationary plasmas”. In: Astronomy and Astrophysics 65, pp. 99–121.
— (1981). “Calculated X-radiation from optically thin plasmas. IV-Atomic
data and rate coefficients for spectra in the range 1-270 A”. In: Astronomy
and Astrophysics Supplement Series 45, pp. 11–52.
Millette, Pierre A (2021). “Stellar Evolution of High Mass-Loss Stars”. In:
Progress in Physics 17.
Mitsuda et al. (2010). “The high-resolution x-ray microcalorimeter spectrometer system for the SXS on ASTRO-H”. In: Space Telescopes and Instrumentation 2010: Ultraviolet to Gamma Ray. Vol. 7732. SPIE, pp. 280–289.
Miyata et al. (1998). “Metal-rich plasma at the center portion of the Cygnus
Loop”. In: Publications of the Astronomical Society of Japan 50.2, pp. 257–269.

148

Bibliography

Miyata et al. (2008). “Evidence for resonance line scattering in the Suzaku
X-ray spectrum of the Cygnus Loop”. In: Publications of the Astronomical
Society of Japan 60.3, pp. 521–526.
Muno et al. (2005). “A neutron star with a massive progenitor in Westerlund
1”. In: The Astrophysical Journal 636.1, p. L41.
Nakashima et al. (2018). “Spatial distribution of the milky way hot gaseous
halo constrained by Suzaku X-ray observations”. In: The Astrophysical Journal 862.1, p. 34.
Nomoto (1982). “Accreting white dwarf models for type I supernovae. IPresupernova evolution and triggering mechanisms”. In: Astrophysical Journal, Part 1, vol. 253, Feb. 15, 1982, p. 798-810. 253, pp. 798–810.
Okon et al. (2020). “Deep XMM-Newton Observations Reveal the Origin of
Recombining Plasma in the Supernova Remnant W44”. In: The Astrophysical Journal 890.1, p. 62.
Park et al. (2003). “X-ray emission from multiphase shock in the large magellanic cloud supernova remnant N49”. In: The Astrophysical Journal 586.1,
p. 210.
— (2012). “An X-Ray Study of Supernova Remnant N49 and Soft GammaRay Repeater 0526–66 in the Large Magellanic Cloud”. In: The Astrophysical Journal 748.2, p. 117.
Patnaude and Fesen (2005). “Model simulations of a shock-cloud interactionin the Cygnus Loop”. In: The Astrophysical Journal 633.1, p. 240.
Plaa, de et al. (2004). “X-ray spectroscopy on Abell 478 with XMM-Newton”.
In: Astronomy & Astrophysics 423.1, pp. 49–56.
— (2012). “Estimating turbulent velocities in the elliptical galaxies NGC 5044
and NGC 5813”. In: Astronomy & Astrophysics 539, A34.
Podsiadlowski (2017). “Binary Evolution and the Progenitor of SN 1987A”.
In: arXiv preprint arXiv:1702.03973.
Pollock (2012). “Ion-ion charge exchange and X-rays from the winds of hotstars”. In: Astronomische Nachrichten 333.4, pp. 351–354.
Porquet et al. (2001). “Line ratios for helium-like ions: Applications to collisiondominated plasmas”. In: Astronomy & Astrophysics 376.3, pp. 1113–1122.
Ramachandran et al. (2019). “Testing massive star evolution, star formation
history, and feedback at low metallicity-Spectroscopic analysis of OB stars
in the SMC Wing”. In: Astronomy & Astrophysics 625, A104.
Rasmussen et al. (2001). “The X-ray spectrum of the supernova remnant 1E
0102.2-7219”. In: Astronomy & Astrophysics 365.1, pp. L231–L236.
Reynoso et al. (1995). “VLA observations of neutral hydrogen in the directionof Puppis A”. In: The Astronomical Journal 110, p. 318.
Russell and Dopita (1992). “Abundances of the heavy elements in the Magellanic Clouds. III-Interpretation of results”. In: The Astrophysical Journal384, pp. 508–522.
Sana et al. (2014). “Southern massive stars at high angular resolution: observational campaign and companion detection”. In: The Astrophysical Journal Supplement Series 215.1, p. 15.

Bibliography

149

Sano et al. (2017). “Interstellar gas toward the Magellanic supernova remnants”. In: AIP Conference Proceedings. Vol. 1792. 1. AIP Publishing LLC,
p. 040038.
— (2020). “ALMA CO Observations of Gamma-Ray Supernova Remnant
N132D in the Large Magellanic Cloud: Possible Evidence for Shocked
Molecular Clouds Illuminated by Cosmic-Ray Protons”. In: The Astrophysical Journal 902.1, p. 53.
Sato et al. (2016). “Suzaku spectra of a Type-II supernova remnant, Kes 79”.
In: Publications of the Astronomical Society of Japan 68.SP1.
— (2021). “High-entropy ejecta plumes in Cassiopeia A from neutrino-driven
convection”. In: Nature 592.7855, pp. 537–540.
Schenck et al. (2016). “A Chandra Study of the Interstellar Metallicity in the
Large Magellanic Cloud Using Supernova Remnants”. In: The Astronomical Journal 151.6, p. 161.
Seitenzahl et al. (2018). “Integral Field Spectroscopy of Supernova Remnant
1E0102–7219 Reveals Fast-moving Hydrogen and Sulfur-rich Ejecta”. In:
The Astrophysical Journal Letters 853.2, p. L32.
Smith et al. (2014). “ASTRO-H White Paper-New Spectral Features”. In: arXiv
preprint arXiv:1412.1172.
Steenbrugge et al. (2005). “Simultaneous X-ray and UV spectroscopy of the
Seyfert galaxy NGC 5548-II. Physical conditions in the X-ray absorber”.
In: Astronomy & Astrophysics 434.2, pp. 569–584.
Strüder et al. (2001). “The European photon imaging camera on XMM-Newton:
the pn-CCD camera”. In: Astronomy & Astrophysics 365.1, pp. L18–L26.
Sukhbold, Tuguldur et al. (2016). “Core-collapse supernovae from 9 to 120 solar masses based on neutrino-powered explosions”. In: The AstrophysicalJournal 821.1, p. 38.
Sumiyoshi et al. (2005). “Postbounce evolution of core-collapse supernovae:
long-term effects of the equation of state”. In: The Astrophysical Journal
629.2, p. 922.
Suzuki et al. (2020). “Plasma Diagnostics of the Supernova Remnant N132D
using Deep XMM–Newton Observations with the Reflection Grating Spectrometer”. In: The Astrophysical Journal 900.1, p. 39.
Takiwaki and Kotake (2011). “Gravitational wave signatures of magnetohydrodynamically driven core-collapse supernova explosions”. In: The Astrophysical Journal 743.1, p. 30.
Tanaka et al. (2022). “Charge Exchange X-Ray Emission Detected in Multiple
Shells of Supernova Remnant G296. 1–0.5”. In: The Astrophysical Journal933.1, p. 101.
Thielemann et al. (1996). “Core-collapse supernovae and their ejecta”. In: The
Astrophysical Journal 460, p. 408.
Tominaga (2008). “Aspherical properties of hydrodynamics and nucleosynthesis in jet-induced supernovae”. In: The Astrophysical Journal 690.1, p. 526.
Tsunemi et al. (1986). “X-ray spectra of the Cassiopeia A and TYCHO supernova remnants and their element abundances”. In: The AstrophysicalJournal 306, pp. 248–254.

150

Bibliography

Turner et al. (2001). “The European photon imaging camera on XMM-Newton:
the MOS cameras”. In: Astronomy & Astrophysics 365.1, pp. L27–L35.
Uchida et al. (2015). “N49: THE FIRST ROBUST DISCOVERY OF RECOMBINING PLASMA IN AN EXTRA GALACTIC SUPERNOVA REMNANT”.
In: The Astrophysical Journal 808.1, p. 77.
— (2019). “High Forbidden-to-resonance Line Ratio of O VII Discovered fromthe Cygnus Loop”. In: The Astrophysical Journal 871.2, p. 234.
van der Heyden et al. (2003). “High resolution spectroscopy and emission
line imaging of DEM L 71 with XMM-Newton”. In: Astronomy & Astrophysics 406.1, pp. 141–148.
Vink (2012). “Supernova remnants: the X-ray perspective”. In: The Astronomy
and Astrophysics Review 20, pp. 1–120.
Vogt and Dopita (2011). “The 3D structure of N132D in the LMC: a latestage young supernova remnant”. In: Astrophysics and Space Science 331,
pp. 521–535.
Wachter et al. (1979). “Parameter estimation in X-ray astronomy using maximum likelihood”. In: The Astrophysical Journal 230, pp. 274–287.
Wargelin et al. (Jan. 2008). “EBIT charge-exchange measurements and astrophysical applications”. In: Canadian Journal of Physics 86.1, pp. 151–169.
DOI : 10.1139/P07-125. arXiv: 0708.0233 [astro-ph].
White and Malin (1987). “Possible binary star progenitor for SN1987A”. In:
Nature 327.6117, pp. 36–38.
Williams et al. (2006). “Dust destruction in fast shocks of core-collapse supernova remnants in the large magellanic cloud”. In: The Astrophysical Journal652.1, p. L33.
Winkler et al. (1988). Supernova Remnants and the Interstellar Medium. Cambridge University Press Cambridge.
Woosley and Weaver (1995). The evolution and explosion of massive Stars II: Explosive hydrodynamics and nucleosynthesis. Tech. rep. Lawrence Livermore
National Lab.(LLNL), Livermore, CA (United States).
Xu et al. (2002). “High-resolution observations of the elliptical galaxy NGC
4636 with the reflection grating spectrometer on board XMM-Newton”.
In: The Astrophysical Journal 579.2, p. 600.
Yamaguchi et al. (2014). “Discriminating the progenitor type of supernova
remnants with iron K-shell emission”. In: The Astrophysical Journal Letters785.2, p. L27.
— (2018). “Evidence for Rapid Adiabatic Cooling as an Origin of the Recombining Plasma in the Supernova Remnant W49B Revealed by NuSTAR
Observations”. In: The Astrophysical Journal Letters 868.2, p. L35.
— (2021). “Discovery of Double-ring Structure in the Supernova Remnant
N103B: Evidence for Bipolar Winds from a Type Ia Supernova Progenitor”. In: The Astrophysical Journal Letters 910.2, p. L24.
Yamane et al. (2018). “ALMA Observations of Supernova Remnant N49 in
the LMC. I. Discovery of CO Clumps Associated with X-Ray and Radio
Continuum Shells”. In: The Astrophysical Journal 863.1, p. 55.

Bibliography151

Zhou et al. (2011). “Unveiling the spatial structure of the overionized plasma
in the supernova remnant W49B”. In: Monthly Notices of the Royal Astronomical Society 415.1, pp. 244–250.