Study of Recombination X-rays
from Supernova Remnants with Suzaku
Midori Ozawa
Department of Physics, Graduate School of Science, Kyoto University
Kitashirakawa Oiwake-cho, Sakyo-ku, Kyoto, 606-8502, Japan

midori@cr.scphys.kyoto-u.ac.jp

This thesis was submitted to the Department of Physics,
Graduate School of Science, Kyoto University on January 2010in partial fulfillment of the requirements
for the degree of Doctor of Philosophy in physics.

Abstract
A gigantic explosion of a star forms a supernova remnant (SNR). Huge kinetic energy ofsupernova explosion is transduced into the thermal energy of electrons and atoms by the shock
wave heating, and high-temperature (∼107 K) plasma is formed. Because the plasma in an SNR is
so thin (∼1 cm−3 ), it takes 104−5 years to reach the ionization equilibrium state between electrons
and atoms. Therefore, X-ray investigation of ionization states of the plasma in SNRs provides agood indication of the evolutionary stage of the remnants.
It is the traditional common knowledge that all the plasma in SNRs is in the ionizing stateand end up in ionization equilibrium. In order to check this validity, we performed deep Suzaku
observations of SNRs, W49B, IC 443, and G359.1-0.5, whose X-ray spectra are known to exhibit
highly-ionized states of atoms.
The W49B spectrum exhibits an unusual structure of a saw-edged bump above 8 keV. We
ﬁrmly conclude that this bump is caused by the strong radiative recombination continuum of iron,detected for the ﬁrst time in SNRs. We also detected the strong radiative recombination structuresof silicon and sulfur from IC 443 and G359.1-0.5. For all these SNRs, the electron temperatures
determined from the bremsstrahlung and recombination continuum shapes are consistent witheach other. Furthermore, volume emission measures determined independently from the ﬂuxes ofthe bremsstrahlung and recombination components are also consistent. The results indicate thecommon origin of these components.
We further investigate the spatial distribution of the recombining plasma in W49B and IC 443.
Both SNRs exhibit signiﬁcant anisotropy. Combined with the spatial distribution, lifetime of highlyionized atoms, metallicity of the recombining plasma, and surrounding environments, we discuss theformation process of this anomalous coexisting state of highly-ionized atoms and low-temperatureelectrons. We propose the plausible origin is supernova explosion in a dense circumstellar matter.

Contents
1 Introduction

1

2 Supernovae and Supernova Remnants

5

2.1

Supernovae . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

5

2.2Supernova Remnants . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

8

2.2.1

Shock Wave Heating . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

8

2.2.2

Evolution . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

9

2.2.3

Classiﬁcation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 11

3 Instruments

15

3.1

Suzaku . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15

3.2

XRT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 17

3.3

XIS

. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19

3.3.1

Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19

3.3.2

On-orbit Performance . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 22

4 Thin Thermal Plasma

27

4.1

Classiﬁcation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 27

4.2

X-ray Emission . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 29

4.3

4.2.1

Continua . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 30

4.2.2

Lines . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 35

Temperature Measurement . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 41
4.3.1

Electron Temperature (kTe ) . . . . . . . . . . . . . . . . . . . . . . . . . . . . 41

4.3.2

Ionization Temperature (kTz ) . . . . . . . . . . . . . . . . . . . . . . . . . . . 42

3

CONTENTS

4
4.4

Volume Emission Measure Measurement . . . . . . . . . . . . . . . . . . . . . . . . . 44
4.4.1

Bremsstrahlung Component . . . . . . . . . . . . . . . . . . . . . . . . . . . . 44

4.4.2

Recombination Component . . . . . . . . . . . . . . . . . . . . . . . . . . . . 45

5 W49B47

5.1

Previous Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 47

5.2

Suzaku Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 49

5.3

Results for Global Region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 50
5.3.1

Overall Structure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 50

5.3.2

Spectral Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 50

5.3.3

Contribution of Recombination Lines . . . . . . . . . . . . . . . . . . . . . . . 53

5.3.4

Electron and Ionization Temperatures . . . . . . . . . . . . . . . . . . . . . . 54

5.3.5

Abundance Correction for Overionization . . . . . . . . . . . . . . . . . . . . 55

5.3.6

Volume Emission Measure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 55

5.4

Results for Segmented Region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 56

5.5

Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 62

6 IC 443

63

6.1

Previous Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 63

6.2

Suzaku Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 65

6.3

Results for Global Region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 66
6.3.1

Overall Structure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 66

6.3.2

Spectral Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 67

6.3.3

Contribution of Recombination Lines . . . . . . . . . . . . . . . . . . . . . . . 70

6.3.4

Electron and Ionization Temperatures . . . . . . . . . . . . . . . . . . . . . . 71

6.3.5

Abundance Correction for Overionization . . . . . . . . . . . . . . . . . . . . 72

6.3.6

Volume Emission Measure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 72

6.4

Results for Segmented Region . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 73

6.5

Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 78

7 G359.1-0.5
7.1

79Previous Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 79

CONTENTS

5

7.2

Suzaku Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 80

7.3

Analysis and Results for Global Region . . . . . . . . . . . . . . . . . . . . . . . . . 82

7.4

7.3.1

Overall Structure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 82

7.3.2

Spectral Analysis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83

7.3.3

Electron and Ionization Temperatures . . . . . . . . . . . . . . . . . . . . . . 85

7.3.4

Abundance Correction for Overionization . . . . . . . . . . . . . . . . . . . . 86

7.3.5

Volume Emission Measure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 86

Summary . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 86

8 Discussion on Recombination-dominant SNRs

87

8.1

Spectral Features . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 87

8.2

Characteristics . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 97

8.3

Formation Mechanisms . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 99

8.4

Future Works . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 108

9 Conclusion

109

A Suzaku Observation of Tycho

111

B Abbreviation

113

Chapter 1

Introduction
A supernova is a dynamic explosion of a star. The huge kinetic energy is released in the
explosion. This kinetic energy of the ejecta is converted into the thermal energy of atoms andelectrons via the shock wave heating and the hot plasma is formed. Hence, the supernova remnant(SNR) is composed of high-temperature (∼107 K) and low-density (∼1 cm−3 ) plasma.
The physical states of high temperature plasma in an SNR reﬂect the energy transduction
condition. The traditional picture of the time evolution for the SNR plasma ionization states is
described as follows. Soon after the supernova explosion, a blast shock wave goes through theprogenitor star and ejecta atoms are shock-heated and fully ionized. Because of the high electron
and atom density in this stage, however, highly ionized atoms rapidly recombine with electrons
within one day. Almost all the ejecta atoms are in the ﬁrst- or second-ionized states. This phaseis well-observed in in the visible light. The ionization state in this phase last until the passage ofa reverse shock which is formed when a swept-up interstellar matter dominate the ejecta mass.
When the reverse shock wave passes through the plasma, a particle’s kinetic bulk motion is
gradually randomized and transduced into the thermal motion. Among the constituents of plasma,massive ions with large electrical charges are ﬁrstly randomized and heated (ion temperature; kTi ).
Subsequently, electrons are heated (electron temperature; kTe ) by the thermal transport from theions. Finally, heated electrons gradually ionize the ions (ionization temperature; kTz ). Figure 1.1
shows the scheme of the plasma temperature time evolution. Because the plasma density in SNRs
is so low (∼ 1 cm−3 ), it takes long time (∼104 –105 years) to reach the equilibrium. In this manner,
young SNRs (<103 years) obey the following relation:
kTi > kTe > kTz .

1

(1.1)

CHAPTER 1. INTRODUCTION

2

Fig. 1.1.—: Time evolution of temperature for the normal SNR plasma. n and t are plasma density andtime after the reverse shock passage, respectively. Kinetic bulk energy of the shock wave is ﬁrst tranduced
into atoms (ions), then transferred into electrons. Heated electrons then ionize the atoms.

As for the ion temperature, there is no observational results so far. In principle, it can be
determined by the Doppler broadening of line emission from ions. However, it requires detectors
with very high energy resolution. Because it is diﬃcult to determine the ion temperature in ouractual condition, we concentrate on the study of electron and ionization temperatures 1 .
The picture kTe > kTz is supported by many observational results of ASCA, which is theprevious Japanese X-ray satellite. As we can see in ﬁgure 1.2, almost all the SNRs are in theunderionized state (kTe > kTz ). Even old SNRs (> 104 years), such as Cygnus Loop, are still in theunderionized state.

The energy resolution of the CCD (current standard X-ray detector) is ∼130 eV (at 6 keV, FWHM). X-ray microcalorimeter on board Astro-H (next Japanese X-ray satellite) has a very high energy resolution of 7 eV (FWHM)
between 0.3–12keV. Because the expected Doppler broadening is the order of ∼10 eV, we can detect this eﬀect using
the micro-calorimeter. The determination of the ion temperature is an issue for the next-generation X-ray satellites
and hence we will not discuss in this thesis.
13Overionized

W49B ?

Io

niz
ati
on

eq
uil
ibr
ium

z

IC 443 ?

UnderionizedeFig. 1.2.—: Plasma states of the major supernova remnants obtained with the ASCA satellite (Kawasaki etal. 2005). A horizontal axis means electron temperature (kTe ), which reﬂects the kinetic energy of electrons.
A vertical axis is ionization temperature (kTz ), which indicates an ionization degree of heavy elements. The
diagonal line represents kTe = kTz and thus ionization equilibrium condition, and kTz < kTe (lower right)
and kTz > kTe (upper right) regimes correspond to overionized and underionized conditions, respectively.
Note that all remnants, except for W49B and IC 443, are in the underionized state.

Only two exceptions are W49B and IC 443. Using the ASCA data, Kawasaki et al. (2002,
2005) claimed that the plasma is in the overionized state (kTe < kTz ) for these two SNRs (ﬁgure 1.2).
They derived kTe from the spectral continuum shape, assuming it is completely originates fromthe bremsstrahlung emission. However, this method is not appropriate when there is a signiﬁcant
contribution of the non-thermal and/or other continuous spectral components. In order to checkthis validity, we examine these SNRs. We also examine G359.1-0.5, whose spectrum exhibits asignal of highly-ionized sulfur atoms (Bamba et al. 2000).

4CHAPTER 1. INTRODUCTION

In principle, ionization dominates recombination in the underionized state. On the contrary,recombination should dominate ionization in the overionized plasma. Although Kawasaki et al.
(2002, 2005) insist the existence of overionized plasma in two SNRs, they cannot ﬁnd the recombination structure. If there are some structures originate from the recombination of atoms, theoverionized plasma would become conclusive.
The emission from high temperature plasma dominates in X-ray band, hence X-ray data aresuitable to investigate ionization states. Since the plasma states of SNRs can be known by measuring
the intensities of the emission lines of heavy elements, the detector sensitivity for the line spectrum
is essential to the accurate measurements. Suzaku, an actively working Japanese X-ray satellite, has
good energy resolution, high eﬃciency, and low background level for diﬀuse sources. In particular,
it is superior to the other X-ray missions, in the energy band corresponds to the K-shell emission
structures of iron, which is one of the major products of supernova explosions.
In this thesis, we will focus on the study of X-ray emission from SNRs with highly ionizedatoms utilizing the excellent performances of Suzaku. In chapter 2, we describe the reviews ofsupernovae and SNRs. The basic characteristics of the instruments onboard Suzaku are summarizedin chapter 3. In chapter 4, we describe the classiﬁcation of the thin thermal plasma and its detailed
emission mechanisms. The following chapters provide the ﬁrst results on the analyses of individual
Galactic SNRs with Suzaku: W49B in chapter 5, IC 443 in chapter 6, and G359.1-0.5 in chapter 7.
Combining all the works, we develop the discussions in chapter 8, and we ﬁnally summarize our
ﬁndings in chapter 9.

Chapter 2

Supernovae and Supernova Remnants
Supernovae (SNe) are one of the most energetic events in the universe. Its typical kinetic energy
is ∼ 1051 ergs1 . This is as large as the total radiation energy that the sun radiates in its whole
life time. The remnants created by the supernovae is called supernova remnants (SNRs). SNRs
can yield information about the progenitors, explosion mechanism, and evolution of the remnants.
This chapter is devoted to explain the basic characteristics of SNe and SNRs.

2.1

Supernovae

Around 300–500 SNe are discovered in a year, recently. In our Galaxy, however, SNe have notbeen discovered for ∼400 years 2 . SNe play quite important roles in several aspects. They scatterthe heavy elements into the interstellar medium synthesized in the stars and the explosion via
explosive nucleosynthesis. Furthermore, the shock wave generated by the SN explosions produce
density ﬂuctuation of an interstellar matter (ISM), which promote new star formation.
Figure 2.1 shows the fundamental classiﬁcation scheme of SNe. It is traditionally based on the
optical spectra. If a spectrum lacks absorption lines of hydrogen Balmer series, the SN is classiﬁed
as Type I, otherwise as Type II. Type I SNe are sub-classiﬁed due to the presence of Si and He.
The SNe with strong Si line are deﬁned as Type Ia. The SNe without Si line and with He line areclassiﬁed as Type Ib. Those without Si and He lines are Type Ic.
As we see from the deﬁnition, this classiﬁcation does not reﬂect explosion mechanisms of SNe.
The progenitors of Type Ia SNe are believed to be white dwarfs with nearly Chandrasekhar mass(∼ 1.4M ), while those of Type II to be massive stars (M >∼ 10M ). Details of the characteristics
and origins of classiﬁed SNe are described in the following.

This can be roughly estimated by a typical ejecta mass (M ∼ 1 M ) and a typical expansion velocity (v ∼ 10,000
km s−1 ) as 1/2M v 2 ∼ 1/2×(2 × 1033 g) ×(109 cm/s)2 ∼ 1051 ergs.
2
There are some Galactic supernova remnants whose ages are estimated to be younger than 400 years. However,these SNe were not observed by human.
1

5

6

CHAPTER 2. SUPERNOVAE AND SUPERNOVA REMNANTS

*!
PQ
5K!

[GU

PQ
[GU

*G!
[GU

6[RG+C

6[RG+D

6JGTOQPWENGCTPQ
6[RG+E

6[RG++%QTG%QNNCRUG

Fig. 2.1.—: Classiﬁcation scheme for supernovae.

Type Ia
Type Ia SNe are observed in all types of galaxies. They are even observed in the ellipticalgalaxies where the star formation has ﬁnished and no massive stars exist. The lack of hydrogen
revealed by optical spectra indicates that the progenitors have no hydrogen envelope. These facts
indicate it is plausible that Type Ia explosions originate from thermonuclear explosions of massaccreting white dwarfs in binary systems.
When the mass of the white dwarf reaches about the Chandrasekhar limit (∼ 1.4M ), thegravity exceeds the electron degeneracy pressure. Then, the carbon core burning occurs. Nuclear
fusion advances explosively, and the star is totally disrupted. This explosion leaves no compact
star nor black hole. Nucleosynthesis models of Type Ia SNe predict that the elements such as Si
and Fe are selectively synthesized, while those as O, Ne, and Mg are hardly synthesized (Nomoto
et al. 1984, Iwamoto et al. 1999).

Type II
Because Type II SNe are concentrated in the spiral galaxy’s arm, it is plausible that theirprogenitors are massive stars. A high mass star (≥10M ) ﬁnally has a Fe core. When the internal
temperature exceeds ∼ 3 × 109 K, photo-disintegration of Fe nuclei takes place.
56
4

Fe → 13 4 He + 4n − 124.4 MeV

He → 2p + 2n − 28.3 MeV ,Because this is an endoergic reaction, the star cannot support its gravity and collapses. The
gravitational energy released by the infalling matter creates a shock wave and the shock propagates
outward. The total energy generated by the gravitational collapse is ∼ 1053 ergs. Because most2.1. SUPERNOVAE

7of the explosion energy is carried out by neutrinos, the kinetic energy of the shock wave becomes∼ 1051 ergs. This explosion leaves a neutron star or a black hole. Nucleosynthesis models of TypeII SNe predict that O, Ne, and Mg, as well as Si and Fe are synthesized during this explosion(Thielemann et al. 1996).

Type Ib
Although the optical spectra of Type Ib SNe lack hydrogen absorption lines, the propertyof light curves is very similar to Type II. Furthermore, Type Ib SNe are observed only in spiralgalaxies and star forming regions. These facts indicate that the progenitors of Type Ib SNe are thesame as Type II (massive stars which have lost their hydrogen envelope). The plausible candidate
is a Wolf-Rayet star, because they lose an outer hydrogen envelope by the strong stellar wind.

Type Ic
It is believed that the progenitors of Type Ic are also massive stars. The plausible candidate
is a Wolf-Rayet star which has lost an outer hydrogen and helium envelope. A few fraction of
the Type Ic SNe is said to be accompanied with γ-ray bursts (GRBs). For example, the optical
spectrum of GRB980425 is very similar to Type Ic. The broad absorption feature indicate thevelocity of elements range from 10-20% of the light speed. From these facts, it is concluded thatthe progenitor of this GRB is a high mass star with ∼40M , and associated with a hypernova
phenomenon. Similar phenomena are also detected from GRB030329 and GRB031203.

CHAPTER 2. SUPERNOVAE AND SUPERNOVA REMNANTS

8

2.2

Supernova Remnants

274 SNRs are identiﬁed in our Galaxy (Green 2009), and about a half are observed in X-rays.
In this chapter, we explain the shock heating process, and evolution and classiﬁcation of SNRs.

2.2.1

Shock Wave Heating

Because the shell of an SNR expands supersonically, a shock wave is formed. As the shock
wave passes, the gas is compressed and heated. Here we consider a coordinate system in ﬁgure 2.2(right). Conservation of mass, momentum, and energy across the shock front can be expressed asρ u vu = ρ d vd

(2.1)

ρu vu2 + pu = ρd vd2 + pd

(2.2)

1
1 2
vu + wu = vd2 + wd ,
(2.3)
2
2
where ρ, v, p, and w are the density, velocity, pressure, and enthalpy of upstream (subscript u) anddownstream (subscript d). Assuming the ﬂuid is ideal gas, the enthalpy is given as
w = Cp T =γp,(γ − 1)ρ

(2.4)

where γ is the speciﬁc heat ratio (= Cp /CV ).
From equations 2.1, 2.2, 2.3, and 2.4, and assuming the strong shock (pd /pu  1), we obtain
ρu
vd
γ −1
=
=
ρd
vu γ + 1(2.5)

pd γ − 1Tu=Tdpu γ + 1(2.6)(γ + 1)pd
2ρu

(2.7)

(γ − 1)2 pd.
2(γ + 1)ρu

(2.8)

vu2 =
vd2 =

Because vu = vs (see ﬁgure 2.2), we can derive the mean post-shock temperature using equations 2.5
and 2.8 as:
pd 2(γ − 1)μmH vs2 ,(2.9)
kTd = μmH =ρd(γ + 1)2where k, μ, and mH are the Boltzmann constant, the mean atomic weight, and the hydrogen mass,
respectively. Assuming the non-relativistic single atomic gas (γ = 5/3), we obtain the relations ofρd /ρu = 4 and vd /vu = 1/4. The post-shock temperature is described askTd =3
μmH vs2 .
16(2.10)

As we can see from this equation, the shock wave heating can be understood as the energy conversionof the bulk kinetic energy into the randomized thermal energy.

2.2. SUPERNOVA REMNANTS

91DUGTXGTTGUVHTCOG
504

5JQEMTGUVHTCOG

+5/

XU

504

+5/

XF

XW XU

XFF
YF6F

XWW
YW6W

FQYPUVTGCO WRUVTGCOUJQEMHTQPVFig. 2.2.—: A schematic view of the shock wave in observer rest frame (left) and shock rest frame (right).
v, ρ, p, w, and T are the velocity, density, pressure, enthalpy, and temperature, respectively.

2.2.2

Evolution

Basically, an SNR is composed of an expanding ejected material (ejecta) and a swept-up ISM.
The evolution phase of an SNR is mainly determined by the relation between the mass of the ejecta(MEj ) and swept-up ISM (MISM ). We here assume that the density of the ambient medium is
uniform 3 .
1) MISM  MEj : Free Expansion Phase
The ejecta expands without deceleration. The expansion velocity (∼ 104 km s−1 as we see
below) is much larger than the sound speed in the ambient gas (∼ 10 km s−1 ), and hence shock
wave (blast wave) precede the ejecta. Since most of the explosion energy (E) is released kinetically
4 , the expansion velocity (v ) and the radius (R ) of the shock front can be expressed by
s
s


vs =



2E
E
= 1.0 × 109
51
MEj
10 erg1/2 

Rs = vs t ,

MEj
M

−1/2

[cm s−1 ] ,(2.11)(2.12)

where t is the time since the explosion. The mass of the ISM swept-up by the blast wave is describedas
4
MISM = πRs3 μmH n0 ,(2.13)
3
where n0 is hydrogen number density of the ISM. Since the free expansion phase continues until
MISM becomes comparable to MEj , its timescale can be approximately estimated from equation 2.11,
2.12, and 2.13 to be
3

The detailed descriptions are shown in Truelove & McKee (1999) for young stages (from free expansion phasethrough Sedov phase) and Cioﬃ et al. (1988) for evolved stages (radiation phase).
4
The energy converted into thermal energy is only 2–3% in this phase.

CHAPTER 2. SUPERNOVAE AND SUPERNOVA REMNANTS

10t ∼ 1.9 × 10

−1/2 E

21051 ergMEjM5/6 

μ1.4

−1/3 n0
1 cm−3

−1/3[yr] .

(2.14)

Therefore, this phase typically continues for a few hundred years.
2) MISM ∼ MEj : Transition Phase
In this phase, the interaction with the ISM cannot be ignored, although it is not yet dominant.
The swept-up ISM decelerates the ejecta and causes another shock wave (reverse shock). The ejecta
is compressed and heated by the reverse shock. The boundary between the ISM and the ejecta
is called the “contact discontinuity”. Note that reverse shock initially propagate outward. WhenMISM becomes greater than MEj , the reverse shock propagates back to the center. Most of young
SNRs, such as Tycho, are probably in this stage.
3) MISM ≥ MEj : Adiabatic (Sedov) Phase
The blast wave is decelerated by the swept-up ISM, but the energy loss via radiation is still
negligible. Therefore, the SNR expands adiabatically. In this stage, the shock evolution can beapproximated by a point explosion in a uniform medium 5 . The gas ﬂow can be characterized by
two parameters, E and n0 . The radius (Rs ) and velocity (vs ) of the blast wave, and the mean
temperature just behind the shock front (Ts ) are written asRs = 4 × 10

104 yrvs =2/5 

t

19

dRs
t
= 5 × 107dt
104 yr

Ts = 3 × 106t
104 yrE

1/5 

1051 erg
−3/5 

−6/5 

1/5 

E
1051 erg
E1051 ergn01 cm−3

2/5 

−1/5

n0
1 cm−3n0
1 cm−3

−1/5

[cm]

(2.15)[cm s−1 ](2.16)[K] .

(2.17)

−2/5

In this phase, the density of the swept-up ISM is much higher than that of ejecta, and hence
X-ray emission is dominated by the ISM. The Sedov phase lasts for a few ten thousand years. Until
the end of this phase, ∼ 70% of the initial explosion energy (bulk kinetic energy) is transduced into
thermal energy of the swept-up ISM (Chevalier 1974).
Radiative Cooling Phase
In the Sedov phase, the temperature behind the shock decreases with time (equation 2.17).
The cooling rate gradually increases, and the evolution becomes no longer adiabatic. When thetemperature cools down to ∼ 3×105 K, electrons eﬃciently recombine with abundant carbon andoxygen, and radiate ultraviolet emission lines. The shell is driven by the pressure of the internal
gas obeying P V γ = const., where P and V are the mean pressure and the volume of the internal
5This is explained by a self-similar solution by Sedov (1959). Therefore, this phase is often called as “Sedov
phase”.

2.2. SUPERNOVA REMNANTS

11gas, respectively. If we assume the non-relativistic single atomic gas (γ = 5/3), the shock expands
with the time asRs ∝ t2/7 .
(2.18)
When the temperature further cools down, the pressure can be ignored and the cool shell
expands with the conserved radial momentum (Ms vs = const., where Ms is the mass of the shell).
The shock expands as(2.19)
Rs ∝ t1/4 .
When the shock speed becomes comparable to the proper motion of the ISM (10–20 km s−1 ),the SNR lose its boundary. This phase is called as “Disappearance phase”. The age of SNRs inthis phase is larger than ∼ 106 years.

2.2.3

Classification

Based on the radio morphology, SNRs are classiﬁed mainly into three classes; shell-type,plerion-type, and composite. These classiﬁcation is summarized in table 2.1. Among 274 Galactic
SNRs, ∼65% , ∼4% , and ∼10% are classiﬁed as shell-type, plerion-type, and composite, respectively. Remains are not ﬁrmly classiﬁed. Addition to these, Rho & Petre (1998) claimed a new
category of SNRs, called ”mixed-morphology” SNRs. They have shell-like morphology in radio
band, and center-ﬁlled one in X-rays. Here we mention each type.
Shell-type
This type has limb brightened morphology in both X-ray and radio bands. Figure 2.3a shows
one example. X-ray emission is thermal or non-thermal. Examples for the former are Cassiopeia A,
Tycho, and Kepler. Almost all the observed SNRs in this category are considered to be between
free expansion and adiabatic phases. Examples for the latter are SN1006 and RX J1713.7-3946.
The origin of the non-thermal X-rays is synchrotron radiation from relativistic electrons accelerated
to the energy order of TeV by a blast wave. On the other hand, the origin of the non-thermal radio
photons is due to electrons accelerated up to the energy order of GeV.
Plerion-type
This type shows non-thermal emission with center-ﬁlled morphology. Figure 2.3b shows one
example. The origin of X-rays is synchrotron radiation from a neutron star or its nebula. Examples
are, Crab Nebula, 3C 58, and G21.5-0.9.
Composite
We can see both shell-like and center-ﬁlled morphology from this category. In ﬁgure 2.3c, weshow one example. The central object is associated with a pulsar, while the emission from limb is

CHAPTER 2. SUPERNOVAE AND SUPERNOVA REMNANTS

12

Table 2.1—: Classiﬁcation of SNRs.

Shell-type SNR
Plerion-type SNR
Composite SNR
Mixed-morphology SNR(Radio)(X-ray)Shell-like ?
yes
no
yes
yes
no

Center-ﬁlled ?
no
yes
yes
no
yes

associated with the blast wave. Thus, this is a mixture of the shell-type and plerion-type. Examples
are, CTB 109, Vela SNR, and G11.2-0.3.
Mixed-morphology
This is another group of SNRs which exhibit diﬀerent morphology in X-ray and radio bands.
One example is shown in ﬁgure 2.3. The bulk of X-ray emissions is thermal, mainly originatesfrom the swept-up ISM. Examples are, W28, W44, and 3C 391. The population represents ∼10%of all Galactic SNRs and as much as ∼25 % of all X-ray detected Galactic SNRs. Rho & Petre(1998) proposed two formation mechanisms to explain this category. One is a cloud evaporation
model. In this scenario, many small interior clouds evaporate and emit X-rays. The other is a
fossil conduction model (Cox et al 1999, Shelton et al. 1999a). This predicts the hot interior
plasma gradually became uniform by thermal conduction and creates centrally brightened X-ray
morphology. X-ray emission from the shell, on the other hand, became undetectable because thetemperature became too low (∼ 106 K) and absorbed by the ISM.
In this thesis, we analyze W49B, IC 443, and G359.1-0.5. All exhibit center-ﬁlled morphology
in X-rays, and a shell-like feature in the radio band. So we can say, all these SNRs are classiﬁed
into the mixed-morphology type.

2.2. SUPERNOVA REMNANTS

13(a) Shell-type(b) Plerion-type(c) Composite(d) Mixed-morphology

Fig. 2.3.—: Examples of classiﬁed SNRs. Gray scale and contours indicate X-ray and radio maps, respectively. (Credit: Chandra Supernova Remnant Catalog. http://hea-www.cfa.harvard.edu/ChandraSNR/) (a)Shell-type SNR; Tycho. Gray scale: 0.3-10 keV by Chandra. Contour: 1.375 GHz by VLA. (b) Plerion-type
SNR; 3C 58. Gray scale: 0.3-10 keV by Chandra. Contour: 21 cm by VLA. (c) Composite SNR; G11.2-0.3.
Gray scale: 0.3-10 keV by Chandra. Contour: 21 cm by VLA. (d) Mixed-morphology SNR; W44. Gray
scale: 0.7-10 keV by ASCA. Contour:1.4 GHz by VLA. These ﬁgures are taken from Kawasaki (ph.D.).

Chapter 3

Instruments
Suzaku is the Japan’s ﬁfth X-ray astronomy satellite. Throughout this thesis, we use thedata obtained with Suzaku. The X-ray CCD (XIS) onboard Suzaku has good energy resolution,high eﬃciency, and low and stable background for diﬀuse sources. All these characteristics are
essential for the accurate analyses of SNRs. In this chapter, we review the Suzaku satellite, onboard
instruments, and onboard calibration issues.

3.1

SuzakuSuzaku (Mitsuda et al. 2007) was launched on July 10, 2005, and put into a near circularorbit at a ∼ 570 km altitude with an inclination angle of ∼ 31◦ and orbital period of ∼ 96 minutes.
Figure 3.1 shows a schematic view of Suzaku. The spacecraft length is 6.5 m along the telescope
axis and the total weight is ∼ 1700 kg.
Suzaku carries three scientiﬁc detector systems: the X-ray Imaging Spectrometer (XIS:Koyama et al. 2007a), the Hard X-ray Detector (HXD: Takahashi et al. 2007; Kokubun et al.
2007), and the X-ray Spectrometer (XRS: Kelley et al. 2007). The XRS and four XIS cameras arelocated at the focal plane of the corresponding X-ray Telescope modules (XRT: Serlemitsos et al.
2007). The capability of these instruments is summarized in table 3.1.
The XRS is the ﬁrst orbiting X-ray microcalorimeter, and supposed to be one of the main
instruments of Suzaku. However, it became nonoperational due to the evaporation of the liquid
helium. The HXD is a non-imaging, collimated detector, which covers band pass of 10–600 keV.
In this thesis, we do not use the XRS and HXD.

15

CHAPTER 3. INSTRUMENTS

16

Table 3.1—: Overview of Suzaku capabilities (Mitsuda et al. 2007).

XRT

XIS

HXD

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

4.75 m20 (@ 1 keV) / 14 (@ 7 keV)
440 cm2 (@ 1.5 keV) / 250 cm2 (@ 8 keV)
2 (Half power diameter)17. 8 × 17. 8
0.2–12 keV
1024 × 1024 / 24 μm × 24 μm∼130 eV (@ 6 keV: FWHM)
330 (FI) & 370 (BI) cm2 (@ 1.5 keV)
160 (FI) & 110 (BI) cm2 (@ 8 keV)4.◦ 5 × 4.◦ 5 (≥ 100 keV) / 34 × 34  (≤ 100 keV)
PIN:10–70 keV / GSO: 40–600 keV

(b)

(a)

Sun shade
XRT
EOB

31

Star
trackerEquator570 kmThe EarthGyroscopes

Side panel(8 panels)

XRS

Orbital path

Base
panel

HXD
XIS (4 units)

Fig. 3.1.—: (a) Schematic view of Suzaku and its orbital path. (b) Side view of Suzaku with the internal
structure after the EOB extension (Mitsuda et al. 2007).

3.2. XRT

3.2

17

XRT

Fig. 3.2.—: Picture of the XRT (Serlemitsos et al. 2007).

There are ﬁve XRT modules (ﬁgure 3.2). One is equipped for the XRS, and the others are for
the XIS. All XRTs are co-aligned to image the same region of the sky. The XRT is conﬁgured toform a Wolter I grazing-incidence reﬂective optics (ﬁgure 3.3). Incident X-rays are reﬂected twiceby the parabolic and hyperbolic surfaces, and are condensed to the focal plane, when their incident
angle is less than ∼ 1◦ . In order to increase the eﬀective area, each XRT module has 175 nested
thin foil layers. Because it is diﬃcult to construct a precise Wolter geometry with thin foils, boththe primary and secondary reﬂectors are approximated by simple cones.
Figure 3.4a shows the energy dependence of the total on-axis eﬀective area of the four XRTs,
including the detector eﬃciency. It is compared with those of other X-ray missions. The XRTprovide a large eﬃciency in the energy range of 0.2–12 keV. The eﬀective area of the four XRT
modules is larger than Chandra in most energy band, and comparable to XMM-Newton above∼5 keV. As shown in ﬁgure 3.4b, the eﬀective area is a function of the oﬀ-axis angle. This eﬀect iscalled “vignetting”. The eﬀect is calculated by the ray-tracing simulator, and conﬁrmed by the observational data of Crab Nebula. The response of the XRT is described by Ancillary Response Files(ARFs). In all spectral analyses in following chapters, we use ARFs generated by xissimarfgen
software (Ishisaki et al. 2007).

CHAPTER 3. INSTRUMENTS

18

Mirrors

X-raysNestedParaboloids

NestedHyperboloids

Focal Point
Fig. 3.3.—: Sectional view of Wolter I reﬂecting optics. Incident X-rays are reﬂected twice at very shallow
angles and concentrated at the focal point.

(a)

(b)

SuzakuFig. 3.4.—: (a) Total eﬀective area of the XRT (combined with the XIS). Those of other X-ray missions(XMM-Newton and Chandra) are also shown for comparison. (b) Oﬀ-axis angle dependence of the eﬀective
area normalized at the optical axes. Data points are the Crab nebula data, and solid lines are the simulated
model (Serlemitsos et al. 2007).

3.3. XIS

19

XIS

3.3.1

Overview(a)(b)
BaffleElectric
valve

Hood
105 mm3.3

Wall calibration
sourceOptical blocking
filter (OBF)
Cover shield
Thermo-electric
CCD
cooler (TEC)
Heat sink

Base
70 mm

Door calibration
sourceBonnet
125 mmDoor

Flexible print

Fig. 3.5.—: (a) Picture and (b) side view of the XIS (Koyama et al. 2007a).

The XISs (ﬁgure 3.5) have four sensors (XIS 0, 1, 2, and 3), and each sensor has one charge
coupled device (CCD) chip. Each CCD consists of four segments (A, B, C, and D) with a dedicatedread-out node. Three sensors (XIS 0, 2, and 3) contain front-illuminated (FI) CCDs, and the other(XIS 1) has a back-illuminated (BI) CCD. Combined with the XRT, the ﬁeld of view is 17 .8×17 .8.
There are 55 Fe calibration sources located on a side wall of the bonnet (ﬁgure 3.5b). 55 Fe decays
to leave 55 Mn via electron capture mainly from the K-shell. Subsequently, a ﬂuorescent X-ray isemitted (Kα: 5.895 keV) accompanied by an electron transition. Calibration sources illuminate
the two far-end corners from the read-out node of the imaging area (ﬁgure 3.6).
The depletion layer of the BI CCD (∼ 42μm) is thinner than that of the FI CCD (∼ 65μm).
The quantum eﬃciency at high energy of the BI is, therefore, lower than FI (ﬁgure 3.7). The clocksof the XIS are generated by the sequence in the analog electronics according to the micro-code. The”normal” clocking mode has imaging capability with time resolution of 8 second. In all observations
described in this thesis, the XIS were operated in the normal mode.

CHAPTER 3. INSTRUMENTS

20

Chargeinjection
register
A

B

C

D

Illuminatedby 55Fe

Read-out
nodes

Quantum efficiency

Fig. 3.6.—: Schematic view of the XIS CCD (Koyama et al. 2007a).

1
FI (XIS 0)
BI (XIS 1)

0.5

0

0

5
10
Energy (keV)

Fig. 3.7.—: Quantum eﬃciency of the XIS as a function of incident photon energy (Koyama et al. 2007a).

3.3. XIS

21

Even if incident X-rays are monochrome, the spectral shape is complicated (ﬁgure 3.8). This
spectral response is called “response function”, which is examined by ground experiments. The response function is included in the response matrix ﬁles (RMFs), which can be made using xisrmfgen
software (Ishisaki et al. 2007).

Fig. 3.8.—: XIS response to monochromatic X-rays. There are (1) main peak, (2) sub peak, (3) triangle
component, (4) Si escape peak, (5) Si ﬂuorescent peak, and (6) constant component. (Matsumoto et al.
2006)

Counts s −1 keV −1 pixel −1

Figure 3.9 shows non X-ray background (NXB) spectra obtained while Suzaku was pointed
toward the night earth. Since the orbit altitude of Suzaku is much lower than that of XMM-Newton
or Chandra, Suzaku is less inﬂuenced by solar ﬂares and keeps the very stable NXB level. Hence,
the XIS is sensitive to diﬀuse sources.

Au Lα
−7

10

Mn Kα Ni Kα

Al Kα
Si Kα
Au Mα

Au Lβ
−8

10

−9

10

0

XIS0 XIS2
Ni Kβ
XIS1 XIS3 Mn Kβ

5

10Energy (keV)

Fig. 3.9.—: XIS non X-ray background spectra (Tawa et al. 2008). Note that the count rate of the BI(XIS 1) is higher than that of the FI (XIS 0, 2, and 3) especially above ∼7 keV.

CHAPTER 3. INSTRUMENTS

22

3.3.2

On-orbit Performance

Since the launch, the energy gain and resolution of the XIS is gradually degrading due to theradiation damage (ﬁgure 3.10). The damage increases a charge-transfer ineﬃciency (CTI).
In order to mitigate this eﬀect, we (= XIS team members) have carried out two methods
utilizing the charge injection technique. Using this technique, we can inject an arbitrary amountof charge in arbitrary positions (Prigozhin et al. 2004, Bautz et al. 2004, LaMarr et al. 2004,
Prigozhin et al. 2008). One method we have carried out is ”a checker-ﬂag charge injection (CFCI)”
1 , and the other is ”a spaced-row charge injection (SCI)”. All the results using these two methods
have been implemented to the Suzaku calibration database and applied to all the observational
data. In this section, we summarize the CFCI and SCI, and mention the observing condition inthis thesis.

(a)

(b)

Fig. 3.10.—: Time history of the (a) energy gain and (b) energy resolution of the Mn Kα line from the55
Fe calibration sources (Koyama et al. 2007a).

1

We can obtain a checker ﬂag pattern by these injected packets in the X-ray image (ﬁgure 3.11b) because of theonboard event-detection algorithm. Therefore, we call this technique a “checker ﬂag” charge injection.

3.3. XIS

23

Checker-flag Charge Injection (CFCI)
We summarize CFCI on-orbit experiments. The details are summarized in Ozawa et al. (2008,2009b). We brieﬂy explain the CFCI method (ﬁgure 3.11). First, we inject a “test” charge packet
into the top CCD row. Then, after the gap of a few rows, ﬁve continuous packets are injectedwith the same amount of charge of the test packet. The former four are called “sacriﬁce”, while
the last one is called “reference”. ”Test” loses its charge by the charge traps. On the other hand,”reference” does not lose the charge because the traps are already ﬁlled by ”sacriﬁce”. Thus, we
can measure the charge loss by comparing the pulse-height of ”reference” and ”test”.
We have conducted CFCI experiments six times in orbit. Eﬀective exposure time ranges froma few to ∼20 ks, and the equivalent X-ray energy of the injected charge packets ranges from 0.3 to
8 keV.
Using this CFCI capability, we can precisely measure the CTI of each CCD column. By
applying the column-to-column CTI correction, the energy scale have been greatly reduced, andthe over-all energy resolution have been also improved.

(a)

(b)

Fig. 3.11.—: (a) Schematic view of the charge injection. (b) Frame image of the XIS during the chargeinjection experiment. The gray scale shows the pixel level. Note that the sacriﬁcial packets are not displayed
due to the event detection algorithm (Nakajima et al. 2008).

CHAPTER 3. INSTRUMENTS

24

Here we show some results after the CTI correction. Figure 3.12a shows center energies of the
Mn I Kα line as a function of time. The time-averaged uncertainty of the absolute energy is greatly
improved to ∼ ± 0.1 % . Figure 3.12b shows the center energies of the O VIII Kα line from thesupernova remnant 1E 0102.2 − 7219. The uncertainty of the absolute energy is ∼ ± 0.2%.
Figure 3.13 shows the energy resolution of the Mn I Kα line after our new CTI correction.
We also plot the results of the traditional correction method. The energy resolution has greatly
improved by the column-to-column CTI correction using the results of the CFCI experiments.

(b)

(a)

$SUOVIII Kα peak [eV]

7\SLFDOHUURU

MnI Kα peak [eV]6HS

Time since 2005-7-10 [day]Time since 2005-7-10 [day]Fig. 3.12.—: Center energy after the CTI correction for (a) the Mn I Kα line from onboard calibration
sources and (b) the O VIII Kα line from 1E 0102.2 − 7219. The theoretical centroids are shown with dottedlines (Ozawa et al. 2009b).

3.3. XIS

25(b) XIS 1

MnI Kα FWHM [eV](a) XIS 0FROXPQWRFROXPQ$SU
7\SLFDOHUURU

7\SLFDOHUURU$XJ

Time since 2005-7-10 [day]Time since 2005-7-10 [day]Fig. 3.13.—: Energy resolution of Mn I Kα for (a) XIS 0 and (b) XIS 1. Arrows indicate the days
we performed the CFCI experiments. We compare the results of a new column-to-column (circle) andtraditional column-averaged (triangle) CTI corrections. Note our new methods has greatly improved theenergy resolution (Ozawa et al. 2009b).

Spaced-row Charge Injection (SCI)
The new CTI correction using the CFCI have greatly improved the energy resolution. The
degradation, however, cannot be fully restored using this method. The charge injection can be usedin another way to mitigate the eﬀect of the radiation damage. The SCI technique can reduce the
CTI actively and improve the energy resolution. In the SCI, a charge is injected into CCD rows
periodically. The injected charge ﬁlls the charge traps as a “sacriﬁce”.
The XIS performed the SCI in orbit for the ﬁrst time in August 2006, and the energy resolution
was improved from ∼200 eV to ∼140 eV (@5.9 keV). The SCI has been a normal observation mode
since October 2006. The details of the SCI are in Uchiyama et al. (2009).

CHAPTER 3. INSTRUMENTS

26
Observing Condition

Since the Suzaku launch, the XIS has been functioning well. However, in November 2006,
XIS 2 suddenly showed an anomaly. About 2/3 of the image was unusable due to a large amount ofcharge leakage. This event was possibly caused by a micro-meteoroid impact. Since this anomaly,
XIS 2 has been dysfunctional. Furthermore, in June 2009, a part of XIS 0 suddenly showed ananomaly in Segment A. The history of these anomaly and the observing condition of the data inthis thesis are summarized in ﬁgure 3.14.

Number of
XIS cameras

Time2005 Aug.

SCI

Observation

Observation start

2006 Sep.

SCI mode start

2006 Nov.

Anomaly in XIS2

4

Off2007 Mar. IC 443

3

On2008 Mar.* HESSJ1745-303
2008 May. **Lockman Hole2008 Aug. Tycho
2008 Sep. G359.1-0.5
2009 Mar. W49B

2009 Jul. Anomaly in XIS0 SegA

Fig. 3.14.—: History of the observing condition of the XIS. *Used as background data of G359.1-0.5.
**Used as background data of Tycho.

Chapter 4Thin Thermal Plasma
In this chapter, we classify the thin thermal plasma based on the ionization balance, andintroduce ionizing, equilibrium, and recombining plasmas. Because these plasmas exhibit diﬀerent
spectra, we can classify and diagnose the plasma states from an observed X-ray spectrum. For
example, electron and ionization temperatures, and volume emission measures are determined fromthe spectra in these states.

4.1

Classification

In a thin thermal plasma, we assume that the electron distribution is given by the MaxwellBoltzmann distribution law
f (v) = 4πv 2 mv 2 m 3/2
exp −
,
2πkTe
2kTe

(4.1)

where f (v), k, m, and v are the velocity probability density, Boltzmann constant, mass and velocityof an electron, respectively. The key parameter is the electron temperature designated as kTe . We
also designate the ionization temperature as kTz , which determines the ion fractions (ionization
degree of ions). With these deﬁnitions, the plasma can be classiﬁed into three cases as below.
1. kTz < kTe : underionized
2. kTz = kTe : ionization equilibrium
3. kTz > kTe : overionized
These three cases are determined by the competition between ionization and recombination
processes. The ionization processes are due to electron impact (collisional ionization), photo electric
ionization, and autoionization. As the inverse of these processes, the recombination processes are
due to 3-body recombination, radiative recombination, and dielectronic recombination.

27

CHAPTER 4. THIN THERMAL PLASMA

28

ᵲᶃᴾᵜᴾᵲᶘ

ᵳᶌᶂᶃᶐᶇᶍᶌᶇᶘᶃᶂ

ᵧᶍᶌᶇᶘᶇᶌᶅ
ᵲᶃᴾᵛᴾᵲᶘ

ᵧᶍᶌᶇᶘᵿᶒᶇᶍᶌ
ᵣᶏᶓᶇᶊᶇᶀᶐᶇᶓᶋ

ᵷᶍᶓᶌᶅᴾᵱᵬᵰᶍᶀᶑᶃᶐᶔᶃᶂ
ᶇᶌᴾᵱᵬᵰᶑ
ᵭᶊᶂᴾᵱᵬᵰ

ᵰᶃᶁᶍᶋᶀᶇᶌᶇᶌᶅ
ᵲᶃᴾᵚᴾᵲᶘ

ᵭᶔᶃᶐᶇᶍᶌᶇᶘᶃᶂ

Fig. 4.1.—: Schematic view of plasma states. Note that almost all the observed thin thermal X-ray spectraof SNRs indicate underionized or ionization equilibrium states.

1) Underionized: Thermal energy of electrons exceeds the ionization energy. In the underionized case1 , ionization dominates recombination. We hence describe the plasma in this state”ionizing plasma”. Ionizing plasma has been observed in almost all the young SNRs, such as
Cassiopeia A, SN 1006, and Tycho.
2) Ionization equilibrium: This is the ﬁnal state for a thermal plasma, in which the balance between the ionization and recombination rates is achieved. Because ionization is due toelectron collisions in most cases, we describe the plasma in this state as ”collisional ionization
equilibrium (CIE) plasma”. Such a CIE plasma has been observed in old SNRs.
3) Overionized: This is the opposite of case (1). Recombination dominates ionization. We
designate the plasma in this state as ”recombining plasma”. No one has observed the recombining
plasma in SNRs.
A schematic view of these plasma states is summarized in ﬁgure 4.1.

1

We here only assume the case with no external factor, such as a strong X-ray irradiation source, electric or
magnetic ﬁeld, or rapid expansion or contraction of the plasma.

4.2. X-RAY EMISSION

4.2

29

X-ray Emission

The appropriate plasma emission model depends not only on the temperature, but also on thedensity. Focusing on low density (∼ 1/cm3 ) and high temperature (106 ∼ 108 K) plasma condition
which is typical of SNRs, we assume some approximations below.
• Low density and collisionally ionized plasma (Coronal approximation).
In this approximation, most ions are in the ground state.
• Electrons are Maxwellian distributed.
• External electric and magnetic ﬁeld eﬀects are ignored.
Figure 4.2 shows the predicted spectrum from a thin thermal plasma in CIE state with above
assumptions. We can see complicated structures of many line and continuum emissions. In thissection, we summarize these emission mechanisms.

kTe = 1.5 keV
ZFe=1, ZNi=1

Fig. 4.2.—: Predicted thin thermal spectrum for electron temperature of 1.5 keV and abundances ofsolar values. We show the 5–12 keV band, where emission lines mainly originate from Fe. Most of thecontinuum emission is due to the bremsstrahlung radiation. A very small continuum jump at 8.83 keV(arrow) corresponds to the edge energy of a radiative recombination continuum of H-like Fe. This ﬁgure isderived from the APEC code assuming the CIE state.

CHAPTER 4. THIN THERMAL PLASMA

30

4.2.1

Continua

There are three radiative processes for continuum emission; thermal bremsstrahlung, radiative
recombination, and 2-photon radiation. These emissions are due to free-free, free-bound, andbound-bound transitions of electrons, respectively. The most intense continuum emission is thethermal bremsstrahlung in a CIE plasma of kT of the order of keV.
Free-free Emission: Bremsstrahlung
Electrons emit photons by the bremsstrahlung process in the Coulomb ﬁeld made by ions. The
emission power in an isotropic thermal distribution (thermal equilibrium) is described asdW
dV dtdν

32πe6 2π 1/2 −1/2 2=Te
A ne ni e−hν/kTe gf f
3mc3 3km= 6.8 × 10−38 A2 ne ni Te−1/2 e−hν/kTe gf f [ergs s−1 cm−3 Hz−1 ],

(4.2)
(4.3)

where e, c, A, ne , ni , and gf f are the electron charge, light speed, atomic number, electron and ion
densities, and velocity averaged Gaunt factor.
Gaunt factor is a function of the electron energy and the emission frequency 2 . For a typical
X-ray emitting thermal plasma,
3 kTe 1/2
gf f =.
(4.4)
π hν
As we see in the above formula, we can investigate electron and ion densities and electron temperature from the intensity and shape of the bremsstrahlung continuum spectrum. In ﬁgure 4.3a,
we show one example of bremsstrahlung emission.

2

This is reviewed in Bressaard & van de Hulst (1962), and Karzas and Latter (1961).

4.2. X-RAY EMISSION31

kTe = 1.5keV(a)kTe = 1.5keV

(b)




Fig. 4.3.—: Predicted continuum spectra in the 5–12 keV band for electron temperature of 1.5 keV. (a)Bremsstrahlung continuum due to free-free transitions of electrons. (b) Radiative recombination continuum
due to free-bound transitions of electrons. We show the continuum structure made by H-like Fe ions.

ᶙ

ᵋ

ᵋ

ᵋ

ᵋ

ᵰᵰᵡᵆḷᶌᵛᵐᵇ ᵰᵰᵡᵆḷᶌᵛᵑᵇ ᵰᵰᵡᵆḷᶌᵛᵒᵇ⁝⁛⁠⁗⁦⁛⁕‒⁗⁠⁗⁤⁙⁫‒⁡⁘‒⁗⁞⁗⁕⁦⁤⁡⁠

ᵌᵌᵌ

ᵌᵌᵌ

ᵌᵌᵌ

ᵌᵌᵌ

Q ҄
Q 

(QHUJ\

Q 
Q 

ᵦᵋᵰᵰᵡ
ᵆḷᶌᵛᵏᵇᵪᶗᶽᶐᶃᶁ

ᵪᶗᶾᶐᶃᶁ

ᵪᶗᶿᶐᶃᶁ

Q 
Fig. 4.4.—: Schematic illustration of emission mechanisms for radiative recombination continua andrecombination lines.

CHAPTER 4. THIN THERMAL PLASMA

32

Free-bound Emission: Radiative Recombination Continuum
When an electron recombines with an ion, a photon is emitted. The energy of this photon
equals the kinetic energy of the electron plus the binding energy of the newly-recombined electron.
Since the kinetic energy of the electron is not unique, this forms a continuous spectrum with asharp edge at the binding energy of the levels. This emission is called a radiative recombination
continuum (RRC). We show the schematic illustration of the RRC emission mechanism in ﬁgure 4.4.
The power emitted by this process is described in Tucker & Gould (1996) as:dE
dWf (v)dv== ne nA,j+1 vE · σnrec (Ee ) ·,dV dtdω
dE
dE

(4.5)

where nA,j+1 , E, and σnrec (Ee ) are the density of the ion of the atomic number A and the ionization
state j + 1, the energy of the emitted photon, and the recombination cross section to the level n at
the electron energy Ee , respectively. In our assumption, f (v) is the Maxwell-Boltzman distribution
described in equation 4.1.
This equation can be simpliﬁed considerably if we deﬁne Eedge to be the electron binding
energy3 . Substituting Ee = mv 2 /2 and the Maxwell-Boltzmann equation for f (v), we can rewrite
equation 4.5 as below.
E − EdWedge= 4ne nA,j+1 E
dE
kTe



 E −E1edge· σnrec (E − Eedge ) · exp −,2πmkTekTe(4.6)

This equation completely describes the spectrum of the RRC4 .
We can approximate this formula when the electron temperature is much lower than the K-edgeenergy (kTe  Eedge ) as below:

E − EedgedW∝ exp −,dE
kTe

for E ≥ Eedge .

(4.7)

Note that dW/dE = 0 for E < Eedge . In the spectral analyses of the following chapters, we adopt
this formula. In ﬁgure 4.3b, we show the RRC spectrum from H-like Fe ions assuming this equation.
In this thesis, we describe the RRC accompanied by the recombination of H-like ions to theground state of He-like ions as ”He-RRC”. In the same way, we designate that of fully ionized ionsto the ground state of H-like ions as ”H-RRC”. In table 4.1, we summarize the edge energies of
He-RRC and H-RRC for major elements.

3
4

Eedge corresponds to the spectral edge energy of the recombination continuum.
More details are described in http://cxc.harvard.edu/atomdb/physics/plasma/plasma.html.

4.2. X-RAY EMISSION

33

Table 4.1—: K-shell edge energies (Eedge ) of He-RRC and H-RRC. These values correspond to the K-shellbinding potential of He-like and H-like ions, respectively.

A

Element

12
14
16
18
20
26

Mg
Si
S
Ar
Ca
Fe

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

Bound-bound Emission: 2-Photon Decay
If the direct transition between two bound states is forbidden by a selection rule, two-photon
process occurs. The simpliﬁed level scheme is shown in ﬁgure 4.5a. Because only the sum of theemitted two photons are conserved, the energy of each photon is not uniquely determined. For thisreason, bound-bound transition associated with two photon radiation forms a continuous spectrum.

CHAPTER 4. THIN THERMAL PLASMA

34

D



6

Z

E
] \[

Fig. 4.5.—: (a) Simpliﬁed level scheme for He-like ions. w, x and y, and z represent resonance, intercombination, and forbidden lines, respectively. Solid upward arrows indicate the collisional excitation transitions. Broken arrows are radiative transitions (including photo-excitation from 3 S1 to 3 P0,1,2 levels, and2-photon continuum from 1 S0 to the ground level). Thick skew arrows are recombination (radiative and dielecrtonic) plus cascade processes. This ﬁgure is taken from Porquet et al. (2001). (b) Prospective spectrum
for the electron temperature of 1.5 keV in CIE. We show the energy band around the Kα energy of a He-like
Fe ion. We can see the resonance (w), inter-combination (x and y), and forbidden (z) lines. Other lines
mainly originate from the dielectronic recombination processes of He, Li, and Be-like Fe ions. This ﬁgure isderived from the APEC code.

4.2. X-RAY EMISSION

4.2.2

35

Lines

The bound-bound transitions between two discrete quantum levels accompany photon emission
called line emission. In the cases of a hydrogen atom or other H-like ion5 , a fairly complete discussionof transition rates can be given by quantum theory. For a hydrogen atom, the energy of a photon
absorbed or emitted in a transition between two discrete levels (n and n ) is given by

E = Ry

1
1
− 2
2
n
n



,

(4.8)

where Ry is Rydberg constant (13.6 eV). For heavy elements, E roughly becomesE ∼ A2 Ry

11−n2 n2



.

(4.9)

The emission lines are referred to as Lyα (H-like Kα; 2p → 1s), Lyβ (H-like Kβ; 3p → 1s), Lyγ(H-like Kγ; 4p → 1s), and so on. Mean azimuthal (orbital angular momentum) quantum numbers(l) are 0 for s, and 1 for p. The line energies of Lyman series of major elements are listed intable 4.2.
The emission mechanisms for He-like ions6 are summarized in ﬁgure 4.5. He-like ions emit
three intense lines: resonance (1s2p 1 P1 → 1s2 1 S0 ), forbidden (1s2s 3 S1 → 1s2 1 S0 ), and intercombination (1s2p 3 P2,1 → 1s2 1 S0 ) transition lines. In this thesis, the sum of Kα lines is denotedby ”Heα” (n = 2 → 1), and He-like Kβ by ”Heβ” (3 → 1), He-like Kγ by ”Heγ” (4 → 1), and so on.
Table 4.2—: K-shell line energies of H-like and He-like ions.

Lyα

Heα
Heβ
Heγ
res.∗ forb.∗ int.∗
Mg 1472 1745 1840
1352 1330
1343
1579
1660
Si
2006 2377 2506
1865 1840
1854
2183
2294
S
2623 3107 3277
2461 2431
2447
2884
3033
Ar 3323 3936 4151
3140 3104
3124
3685
3875
Ca 4106 4864 5130
3908 3845
3892
4582
4819
Fe 6966 8266 8732
6702 6641
6670
7798
8217
∗ res., forb., and int. indicate resonance, forbidden, and inter-combination lines, respectively.

5
6

Lyβ

Lyγ

An ion which contains only one electron is called a H-like ion.
An ion which contains two electrons is called a He-like ion.

36

CHAPTER 4. THIN THERMAL PLASMA

The ionization states of ions depend on the plasma temperature. Figure 4.6 shows the ion
fractions of major elements in the CIE state as a function of temperature. When the plasma
temperature becomes low, the fraction of low-ionized ions increases. For this reason, a weighted
line center energy of emission lines becomes low. We show this eﬀect in ﬁgure 4.7.

Fig. 4.6.—: Ion fractions of heavy elements as a function of plasma (ionization) temperature. Red, green,
blue, aqua, purple, brown and orange colors represent fully-ionized, H-, He-, Li-, Be- B- and C-like ions,
respectively.

4.2. X-RAY EMISSION37Si HeαAr HeαSi HeβAr HeβFe HeαFe HeβS HeαCa HeαFe HeγS HeβCa HeβFe Heδ

Fig. 4.7.—: Weighted line center energies for He-like and lower ionized ions as a function of plasma
temperature. Plasma condition is assumed be in a CIE state. Low center energies at low temperatures aremainly due to dielectronic recombination processes. These ﬁgures are derived from the APEC code.

CHAPTER 4. THIN THERMAL PLASMA

38

Here we mention the radiative processes of line emission. When an electron collides with anion, it may excite the ion (electron impact excitation), or it may ionize the ion (electron
impact ionization) ﬂicking the bound electron at the inner-shell7 . If an electron in an upper level
decays to a lower bound state subsequently, a monochromatic photon is emitted.
Dielectronic Recombination is other important process. An electron recombines with an
ion to an excited level and subsequently another electron is excited to a higher level. Thus theion is left in a highly excited state, which may then autoionize or radiatively decay. In the latter
case, one of the excited electrons decay to a lower level, creates a satellite line at a slightly lower
energy than the normal transition. ”Satellite” is named for the energy in the vicinity of the maintransition line, because another electron is still at the upper level.

ᵆᶀᵇᴾᵦᶃᵋᶊᶇᶉᶃ

ᵆᵿᵇᴾᵦᵋᶊᶇᶉᶃ
Q!

ᵋ

ᵐ

ᵮ

ᵋ

Q 

(QHUJ\

LyβrecLyγrec

ᵐ

ᵮ

Q!

Q 

Lyαrec

ᵋ
ᵋ

ᵏ
ᵑ

ᵮ
ᵱᵊᴾᵑᵮ

ᵏ

ᵋ

ᵮ
ᵱᵊᴾᵑᵮ

ᵑ

Heβrec Heαrec Heαrec Heαrec
Heγrec
ᶃᶒᶁᵌᵌ

ᶃᶒᶁᵌᵌ

Q 

ᵋ

ᵐ

ᵱ

Q 

ᵏ

ᵱ

Fig. 4.8.—: Schematic illustrations of emission mechanisms for cascade lines. (a) Simpliﬁed level schemefor H-like ions. Most free-bound electrons captured into excited levels are straightforwardly transit to theground state and produce radiate recombination (cascade) lines. (b) Same for He-like ions. Most electrons
captured into excited ”singlet” levels are straightforwardly transit to the ground state. On the contrary,most electrons into ”triplet” levels are once transit to triplet levels of n = 2 (dotted arrow), and then transit
to the ground state because of the selection rules for radiative transitions.

7

It is called the inner-shell ionization.

4.2. X-RAY EMISSION

39

Cascade Decay is an important process in recombination-dominant plasmas. Cascade decay
is accompanied by the recombination of electrons into higher excited levels in the free–boundtransitions. Figure 4.8 shows schematic illustrations of cascade line emission mechanisms for H-like(ﬁgure 4.8a) and He-like (ﬁgure 4.8b) ions. Throughout this thesis, we designate the cascade linesof H-like ions as Lyαrec (n = 2→1), Lyβrec (3→1), Lyγrec (4→1), and so on. We also designate thecascade lines of He-like ions as Heαrec (n = 2→1), Heβrec (3→1), Heγrec (4→1), and so on.
In the case of H-like ions, most electrons captured into upper levels are straightforwardly decay
to the ground state. In the case of He-like ions, the transition processes are very complicated. In
the ﬁrst-order approximation, however, most electrons captured into upper ”singlet” levels are
straightforwardly decay to the ground state. On the other hand, most electrons captured into
”triplet” levels transit via triplet levels of n = 2 to the ground state. This is because of theselection rules for radiative transitions.
In recombination-dominant plasmas, as we can expect from ﬁgure 4.8b, these cascade processes
increase the intensity ratio of the forbidden and inter-combination lines to the resonance line compared to the ratio in collisional-dominant plasmas. Hence, the Heα line center of a recombinationdominant plasma is signiﬁcantly lower than that of the collisional-dominant case. In fact, theexperimental weighted line center for Fe Heα in the recombination-dominant condition is 6.666(±0.005) keV (Wargelin et al. 2005) 8 . On the other hand, the predicted weighted center for Fe
Heα in collisional-dominant plasma at suﬃciently high temperatures (kT >4 keV) is 6.685 keV (ﬁgrue 4.7). There is a ∼20 eV line center diﬀerence between collisional-dominant and recombinationdominant plasmas. This diﬀerence (∼20 eV) is signiﬁcantly larger than the calibration uncertaintyof the Suzaku XIS at ∼6 keV (∼12 eV; Ozawa et al. 2009b). Thus, we can in principle derive thediﬀerence with very good conditions 9 of high statics in the nearly CIE state plasma 10 .

8Strictly speaking, this is the result of the charge exchange process. In this process, a H-like Fe ion rip an electronfrom the neutral atom or molecular. In perspective of the recombination-dominant condition, however, the essential
mechanisms are the same.
9
This method is applied to the Galactic center diﬀuse X-rays (Koyama et al. 2007b). Because the observed line
center of Fe Heα is 6.680 (±0.001) keV, they concluded the X-ray emission mainly originate from the excitation ofelectrons.
10
Note that the weighted line center energy also becomes low in the underionized case because there are many
low-ionized ions.

CHAPTER 4. THIN THERMAL PLASMA

40

In ﬁgure 4.9, we illustrate the expected spectrum for a recombination-dominant plasma. We
ignore the dielectronic recombination lines here. The recombination structures are composed of
RRC into each principle number (n), and cascade lines below the K-edge energy of the RRC(Eedge ).
The emission ﬂux of the RRC is most intense in the case of RRC[→1] 11 . We can conﬁrm thisfrom the approximated recombination cross section for H-like ions into a level of n (Nakayama &Masai 2001).

1 3 kTe1 −1σn ∝ 3+ 2(4.10)
n 2 Eedge nNote that in the case of He-like ions, this formula is still a good approximation but σ1 should be
reduced by half because one electron exists already in the ground state. Using this equation, we
can estimate the ﬂux ratios between RRC[→>1] 12 and RRC[→1]. For example, in the case of Si
H-RRC at kTe = 0.5 keV, the ﬂux ratio of RRC[→2]/RRC[→1] is derived from σ2 /σ1 as ∼0.29.
In the same way, the ratio RRC[→3]/RRC[→1] is derived from σ3 /σ1 as ∼0.12. In the following
analysis chapters, we mention only the most intense RRC[→1].

ᶙᶁᵿᶑᶁᵿᶂᶃᴾᶊᶇᶌᶃᶑ ᵆḷᶌᵛᵏᵇ
ᵦᶃᶽᶐᶃᶁ

ᶁᵿᶑᶁᵿᶂᶃᴾ ᶁᵿᶑᶁᵿᶂᶃᴾᶊᶇᶌᶃᶑ
ᴾᶊᶇᶌᶃᶑ
ᵆḷᶌᵛᵐᵇᵆḷᶌᵛᵑᵇᶙᵦᶃᶾᶐᶃᶁ

ᶙ
ᵰᵰᵡᴾᵆḷᶌᵛᵐᵇᵌᵌᵌ

ᵦᶃᶿᶐᶃᶁᵌᵌᵌ ᶉᵲᶃᵌᵌᵌ ᵰᵰᵡᴾ ᴾᵆḷᶌᵛᵑᵇ(QHUJ\ NH9ᵦᶃᵋᵰᵰᵡ
ᵆḷᶌᵛᵏᵇFig. 4.9.—: Spectrum expected for a recombination-dominant plasma. We show the structures accompanied by the electron recombination with H-like Fe ions to He-like Fe ions. Here we ignore the dielectronic
recombination lines. In our spectral analysis in the W49B chapter (in the energy band of 5–12 keV), major
recombination structures are due to the transition of electrons into the ground state (n = 1).

1112This notation means the RRC into the ground state.
This means the RRC into n >1.

4.3. TEMPERATURE MEASUREMENT

4.3

41

Temperature Measurement

We can derive the electron and ionization temperatures from an observed X-ray spectrum inseveral independent methods. We explain how to obtain these temperatures.

4.3.1

Electron Temperature (kTe )The electron temperature is measured by three independent methods. As we see from equation 4.2, the Bremsstrahlung Continuum Shape is a simple function of kTe , hence kTe isdetermined. This has been almost only way to determine kTe before the Suzaku era.
In addition to this method, we have developed two independent methods. One way is to
use the Line Flux Ratios. The probability ratio of collisional excitation between n = 1 → 2
and n = 1 → 3 only depends on the energy of a colliding electron (i.e. electron temperature). If
collisional excitation is a major process in the plasma, which is true for most SNRs, the Kβ/Kα
ﬂux ratio directly reﬂects the transition probability ratio. In the extreme recombining plasma,
there is cascade decay contribution for line emission in some measure. However, this contribution
is negligible for most cases 13 . In ﬁgure 4.10, we show the ﬂux ratios of Heβ/Heα (a) and Lyβ/Lyα(b) of Si, S, Ar, Ca, Fe in the CIE state. We also show Heγ/Heα and Heδ/Heα of Fe (a). All of
the ratios are monotonically increasing functions of the plasma temperature.
The other way is to use the Width of Recombination Continuum. From equation 4.7, we
see that kTe reﬂect the continuum shape. Thus, we can derive kTe from the width of the RRC.

E Lyβ/Lyα

D Heβ/Heαࠉ

Fe Heβ/HeαFe Heγ/HeαFe Heδ/HeαeeFig. 4.10.—: (a) Predicted emissivity ratios of Heβ/Heα for Si (black solid), S (black dashed), Ar (graysolid), and Ca (gray dashed) as a function of plasma temperature. We also show ratios of Heβ/Heα (solid),Heγ/Heα (dashed) and Heδ/Heα (dotted) of Fe. (b) Predicted emissivity ratios of Lyβ/Lyα as a functionof plasma temperature. The meanings of colors and lines are same as (a). These ﬁgures are derived fromthe APEC code, assuming the CIE state.

13

Note in the photoionized plasma, the contribution becomes dominant.

CHAPTER 4. THIN THERMAL PLASMA

42

4.3.2

Ionization Temperature (kTz )The ionization temperature directly reﬂects the ion fractions by deﬁnition. Hence, we canderive the ionization temperature using the ﬂux ratios of ions in diﬀerent ionization states.
The most direct method is to use the Line Flux Ratios. One example is Lyα/Heα. In
ﬁgure 4.11, we show ﬂux ratios of Lyα/Heα in the (a) CIE, (b) overionized, and (c) underionizedcases. All cases exhibit monotonically increasing functions of the plasma temperature.

(a) CIE

Si

S

Ar
Ca
Fe

(b) overionized(kTe=0.4 keV)SiS(c) underionized(kTe=2.0 keV)

Fe

Fig. 4.11.—: Predicted emissivity ratios of Lyα/Heα for several elements as a function of ionization
temperature. (a) Ratios for the CIE state. (b) Same for the overionized state at kTe = 0.4 keV (c) Samefor the underionized state at kTe = 2 keV. These ﬁgures are derived from the plasma code of APEC (a) and
Masai (b and c).

4.3. TEMPERATURE MEASUREMENT

43

The other new method we have developed is to use the fluxes of He-RRC and H-RRC.
Since the ﬂuxes of He-RRC and H-RRC directly reﬂect the amount of H-like and fully ionized atoms,
respectively, their ﬂux ratio or to Heα is sensitive to the ionization temperature. In ﬁgure 4.12, weshow some examples of (a) He-RRC/Heα, (b) H-RRC/He-RRC, and (c) H-RRC/Heα for several
kTe cases. All ratios are monotonically increasing functions of kTz , and hence these ratios can beused to derive kTz .

(b)H-RRC/He-RRC(a)He-RRC/Heα

Fe

Fe

14
15
16

14
15
16

(c)H-RRC/HeαFig. 4.12.—: Predicted emissivity ratios using the RRC as a function of ionization temperature for several
electron temperatures. (a) He-RRC/Heα of Fe. (b) H-RRC/He-RRC of Fe. (c) H-RRC/Heα of Si (left) andS (right). These ﬁgures are derived from the plasma radiation code of Masai (1994).

CHAPTER 4. THIN THERMAL PLASMA

44

4.4

Volume Emission Measure Measurement
A volume emission measure (VEM) is one of the physical parameters. It is given by


VEM ≡

ne nH dV,
4πD2

(4.11)

where ne , nH , V , and D are the electron and hydrogen densities (cm−3 ), emitting volume (cm3 ),and distance to the source (cm), respectively. A VEM is proportional to the observed ﬂux andused to estimate the total amount of the plasma. In the following chapters of SNR analyses, wederive VEMs from the emission of electrons’ bremsstrahlung and recombination components. In
this section, we summarize how to derive VEMs of these components from an observed X-ray
spectrum.

4.4.1

Bremsstrahlung Component

The continuum spectrum of bremsstrahlung is given in equation 4.2. We obtain the VEM
from the normalization parameter of the ﬁtting model which reproduce the observed spectrum.
For the VAPEC model which we use in our analyses14 , the normalization is given in unit of
10−14 ne nH dV /(4πD2 ), and hence it is directly converted to the VEM.
Table 4.3—: Solar number densities of major elements normalized by hydrogen (Anders & Grevesse 1989).

A
1
2
6
7
8
10
12
14
16
18
20
24
25
26
28

14

Element
H
He
C
N
O
Ne
Mg
Si
S
Ar
Ca
Cr
Mn
Fe
Ni

nX /nH |solar
1.00×10+0
9.77×10−2
3.63×10−4
1.12×10−4
8.51×10−4
1.23×10−4
3.80×10−5
3.55×10−5
1.62×10−5
3.63×10−6
2.29×10−6
4.68×10−7
2.45×10−7
4.68×10−5
1.78×10−6

This is a thin thermal plasma model in the CIE state (Smith et al. 2001).

4.4. VOLUME EMISSION MEASURE MEASUREMENT

4.4.2

45

Recombination Component

We here explain how to derive the VEM from the H-RRC of an element X. Note that theprocedure is essentially the same for the He-RRC.
The ﬂux of RRC (FRRC ) accompanied by the capture of electrons into the ground state of
H-like X is given as,(4.12)FRRC = α1 (kTe ) · YX · VEMRRC ,where α1 (kTe ), YX , and VEMRRC are a recombination rate coeﬃcient into the ground state of
H-like X at electron temperature kTe (cm3 s−1 ), a particle number of H-like X, and a VEM for therecombination component (cm−5 ).
The value α1 (kTe ) is given by
α1 (kTe ) = α(kTe ) ·σ1,(σ1 + σ2 + σ3 + ...)

(4.13)

where α(kTe ) is the total K-shell recombination rate coeﬃcient (cm3 s−1 ). σn (cm2 ) means transition cross section to the excited level n. For α(kTe ), we use the results of Badnell (2006), who
calculated the collisional-radiative population equations. The value σ1 /(σ1 + σ2 + σ3 + ...) can bederived from equation 4.10.
The value YX is given by
YX = ZX ·nX · κ(kTz ),nH solar(4.14)

where ZX , nX /nH |solar , and κ(kTz ) are an abundance of X in unit of the solar value, number density
ratios between X and hydrogen in the solar photosphere, and an ion fraction of fully ionized X at
ionization temperature kTz . In Anders & Grevesse (1989), number densities (nX ) of major elements
are listed up. In table 4.3, we summarize these values. We should note that ZX is the abundance
for the ”recombining” component. Since we have no available radiative code for the recombining
plasma, we substitute ZX for the abundance of the ”bremsstrahlung” (APEC) component, unlessotherwise mentioned. From equations 4.12, 4.13 and 4.14 , we can derive VEMRRC .

Chapter 5

W49B
W49B is one of the X-ray brightest Galactic SNRs. Using ASCA data, Kawasaki et al. (2005)
claimed that the plasma of W49B is overionized. Hence, this is one of the best candidates ofrecombining plasma in SNRs. In fact, we can see strong Lyα lines of Si, S, Ar, and Ca in the ASCA
1 spectrum, which are the indicators of highly ionized plasma. The spectrum of W49B is notable
for a very strong Kα line of Fe. This SNR has a jet-like X-ray morphology, which is unusual for
SNRs, and an incomplete shell-like structure in the radio band. In this chapter, we summarize theresults of a Suzaku observation of W49B and discuss the state of the plasma.

5.1

Previous Results

W49B (G43.3−0.2) exhibits centrally ﬁlled X-rays inside a bright radio shell with a radius of
100 arcsec (Pye et al. 1984). The distance to W49B remains uncertain. It was estimated to lie at
a distance of ∼8 kpc (Radhakrishnan et al. 1972; Moﬀett & Reynolds 1994). Using the spectral
distribution of H I absorption, however, Brogan & Troland (2001) found no clear evidence that
W49B is closer than W49A, which is located at a distance of ∼11.4 kpc (Gwinn et al. 1992). This
may indicate a possible association of W49B with the star-forming region W49A (see ﬁgure 5.1).
The near-infrared narrowband images indicate a barrel-shaped structure with coaxial rings, which
is suggestive of bipolar wind structures surrounding massive stars (Keohane et al. 2007; ﬁgure 5.2).
They also showed an X-ray image from Chandra, which has a jet-like structure along the axis of thebarrel. They interpreted these ﬁndings as evidence that W49B had exploded inside a wind-blown
bubble in a dense molecular cloud.
Using ASCA data, Hwang et al. (2000) showed that broadband modeling of the remnant’s
spectrum required two thermal components (0.2 keV and 2 keV) and signiﬁcant overabundances of
Si, S, Ar, Ca, Fe, and Ni. They conﬁrmed that most of the X-ray emitting plasma was nearly incollisional ionization equilibrium (CIE). They also found evidence for Cr and Mn Kα emission.
Kawasaki et al. (2005) claimed the presence of “overionized” plasma in W49B through theanalysis of ASCA 2.75–6.0 keV spectrum. They measured intensity ratios of the H-like Kα (Lyα)
1ASCA is the previous Japanese X-ray satellite.

47

CHAPTER 5. W49B

48

to He-like Kα (Heα) lines of Ar and Ca to obtain the ionization temperature (kTz ), and foundthat kTz (∼2.5 keV) is signiﬁcantly higher than the electron temperature (kTe ) determined fromthe bremsstrahlung continuum shape (∼1.8 keV). Miceli et al. (2006) adopted the same analysis
procedure for the XMM-Newton spectrum of the central region but found no evidence for theoverionized state.

Fig. 5.1.—: 21 cm continuum image of the W49 complex with the H II region W49A to the west and the
SNR W49B to the east. (Brogan & Troland 2001).

CTEOKP
Fig. 5.2.—: Multiwavelength view of W49B (Keohane et al. 2007). (White) Ks band (2.2 μm). (Red)Infrared (2.12 μm) H2 by Palomar. (Green) Infrared (1.64 μm) [Fe II] by Palomar. (Blue) X-ray by Chandra.
We can see a barrel-shaped nebula consisting of bright infrared rings around a glowing bar of intense X-ray
radiation along the axis. Red circle is a source region in our analysis.

5.2. SUZAKU OBSERVATION

5.2

49

Suzaku Observation
Table 5.1—: Observation log of W49B.

Start Date

Obs.ID

2009-03-29
2009-03-31

503084010
504035010

R. A.
Dec.
(J2000.0)
19h 11m 08s 09◦ 06 57
19h 11m 08s 09◦ 06 55Exposure
(ks)
52.2
61.4

10:00
05:00
09:00:00

Declination

15:00

We observed W49B twice with Suzaku. The observation log is given in table 5.1. Because
the pointing positions and rotation angles do not diﬀer signiﬁcantly, we merged these two data
sets. Three of the XISs were operated in the normal full-frame clocking mode with a SCI techniqueduring our observations. Two are FI CCDs and the other is a BI CCD. Data were cleaned using
processing version 2.2.11.24. We used HEASOFT version 6.5.1 for data reduction and XSPEC
version 11.3.2 for spectral analysis. After screening with the standard criteria, 2 the net integration
time was ∼113 ks.

19:11:30

11:00

10:30

Right ascension
Fig. 5.3.—: Vignetting-corrected XIS image of W49B in the 1.5–7 keV band shown on a linear intensity
scale. Data from the three active XISs are combined. Gray contours indicate every 10% intensity level relative
to the peak surface brightness. The black circle and gray annulus indicate the source and background regions,
respectively. The XIS ﬁeld of view is shown by the black square. (Ozawa et al. 2009a)2

http://heasarc.nasa.gov/docs/suzaku/processing/criteria xis.html

CHAPTER 5. W49B

50

5.3

Results for Global Region

5.3.1

Overall Structure

Figure 5.3 shows the vignetting-corrected XIS image in 1.5–7 keV, the energy band including
the major lines of K-shell emissions from Si, S, Ar, Ca, and Fe. We extracted the source spectrum
from a circular region with a radius of 2.5 arcmin and the background spectrum from an annulus
surrounding the source with inner and outer radii of 5 and 8 arcmin, respectively. The source spectrum was made by subtracting the background spectrum. The spectra of the two FI sensors weremerged to improve the statistics because their response functions were almost identical. Figure 5.4
shows the merged FI spectrum. We can see several prominent lines of K-shell emission from Heand H-like ions.

5.3.2

Spectral Analysis

As shown in ﬁgure 5.4, the wide-band spectrum includes too many emission lines and possibly
other complicated structures. For clean analyses and discussion, we focus on the 5–12 keV energy
band, where emission lines and other structures mainly originate from Fe. We used only the FI data
because the non-X-ray background (NXB) count rate of the BI data is much larger than the source
rate in this energy band (see ﬁgure 3.9). We ﬁxed the interstellar extinction to a hydrogen column
density of 5 × 1022 cm−2 with the solar elemental abundances, following Hwang et al. (2000). To
ﬁne tune the energy scale, an oﬀset was added as a free parameter and found to be –4.7 eV, which
was within the allowable range of the calibration uncertainties. We added this oﬀset for further
spectral ﬁtting.

Si
Mg

S Ar Ca Cr
Mn
Fe
Ni+
FeKβ

Fig. 5.4.—: Background-subtracted XIS FI spectrum. The energies of the prominent emission lines from
speciﬁc elements are labeled. (Ozawa et al. 2009a)

5.3. RESULTS FOR GLOBAL REGION

51

(a)

CrKα MnKα(b)Heαrec Ly
He
Lyα Heβrec Heγ-∞recHe-RRC

Fig. 5.5.—: (a) XIS FI spectrum in the 5–12 keV band. The best-ﬁt VAPEC model and additional Kα linesof Cr and Mn (black arrows) with Gaussian functions are shown by solid lines. The lower panel shows theresidual from the best-ﬁt model. We can see large bumps at ∼7 keV (blue arrows) and around ∼9 keV (redarrows). (b) Same spectrum as (a), but with the radiative recombination continuum (red), recombinationlines (orange), and a Lyα line (blue) of Fe. (Ozawa et al. 2009a)

CIE Model Fitting
We ﬁrst ﬁtted the spectrum with a VAPEC (CIE plasma) model (Smith et al. 2001). The
Fe and Ni abundances (hereafter ZFe and ZNi ) normalized by the number fraction of the solar
photosphere (Anders & Grevesse 1989) were free parameters. Since the Cr and Mn Kα emission
lines are not included in the VAPEC model, we added these lines with Gaussian functions.
Figure 5.5a shows the ﬁtting result of the one-VAPEC model, with a temperature of 1.64 keV.
The model exhibited signiﬁcant excess around the Fe Lyα line and above 8 keV, and hence wasrejected with a large χ2 /degree of freedom (dof) of 1051/138.
We then tried a two-VAPEC model, assuming equal abundances between the two components.
Fe Lyα was successfully reproduced, but the temperature of one component was unreasonably high(∼70 keV), and a large residual above 8 keV remained with an unacceptable χ2 /dof of 562/136.
Adding a further plasma or power-law component did not improve the ﬁtting any more. In every
case, the large bump above 8 keV was present.
To examine whether this bump is real or artiﬁcial, we checked the light curves of the sourceand background regions in the 8.5–10 keV band. They showed almost constant ﬂuxes, indicating
no ﬂare-like event had occurred during the observations. In addition, we detected the bump in both
FI spectra (XIS 0 and XIS 3) and even in the BI spectrum. Thus, the bump is a real structure.

CHAPTER 5. W49B

52
Recombination Structures

The edge energy of the bump (∼9 keV) corresponds to the electron binding energy of Fe. Thissuggests that the saw-edged bump is likely due to a radiative recombination continuum (RRC).
Both the line-like excess around Fe Lyα and this bump are suggestive of the overionized statebecause there should be a greater fraction of H-like ions in the overionized state than in the CIE.
This excessive number of H-like ions causes strong He-RRC and Fe Lyα emissions. However, nocurrent plasma code can be applied to the overionized plasma. We, therefore, introduce the Fe Lyαline and recombination structures consisting of the RRC and several emission lines below the Kedge energy of the RRC (Eedge ), in addition to the one-VAPEC model. We consider both He-RRC
and H-RRC3 for a consistency check.
We assume that the RRC is expressed as equation 4.7. This formula gives a good approximation
when the electron temperature is much lower than the K-edge energy (kTe  Eedge ). In the caseof W49B, kTe is ∼1.5 keV, as we conﬁrm later, while Eedge of Fe is ∼9 keV. Thus, we can safely
adopt this formula.
For the recombination lines which originate from cascade decays of free-bound electrons, weconsider Heαrec (n = 2→1), Heβrec (3→1), and Heγ-∞rec (n ≥ 4→1; we combine them into one
broad Gaussian function) 4 . The line widths and center energies of Heαrec and Heβrec are ﬁxed atzero and the experimental values for a charge exchange process, respectively (Wargelin et al. 2005).
Those of Heγ-∞rec are allowed to vary freely. The ﬂuxes of all these lines are free parameters.
The resultant best-ﬁt parameters and models are given in table 5.2 and ﬁgure 5.5b, respectively.
With this model, χ2 /dof is greatly improved to 193/128, although this is still unacceptable in apurely statistical sense. We can see signiﬁcant data excesses at the lower energy sides of the Heαlines of Fe (∼6.5 keV) and Ni (∼7.6 keV). The former is likely due to the incomplete response
function. We should note that Fe Heα statics are superior to those for any other objects observedwith Suzaku, and no signiﬁcant residual has appeared so far. If we observe carefully, we ﬁnd
a similar feature in the Mn-Kα line from the onboard calibration source. The latter would bepartially due to the same reason given above, but is mainly due to lack of satellite lines of theLi-like and lower ionization states5 of Ni in the VAPEC model. If we ignore the 6.4–6.6 keV and7.6–7.7 keV band to escape these eﬀects, an acceptable χ2 /dof of 118/107 is obtained.

3

See chapter 4 for these notations.
See chapter 4 for details of cascade lines.
5
As for Fe, such satellite lines are included in the VAPEC model.
45.3. RESULTS FOR GLOBAL REGION

53

Table 5.2—: Best-ﬁt spectral parameters of W49B.

CIE (VAPEC)

kTe [keV]
1.52 (1.50–1.53)

Line
Cr Kα
Mn Kα
Fe Heαrec d
Fe LyαFe Heβrec d
Fe Heγ-∞rec d
RRC
Fe He-RRC
Fe H-RRC

ZFe [solar]
4.44 (4.36–4.53)b
Center energy [keV]
5.655 (5.646–5.663)
6.162 (6.142–6.183)
6.666 (ﬁxede)
6.961 (6.946–6.971)
7.880 (ﬁxede)
8.538 (8.492–8.585)
Edge energy [keV]
8.830 (ﬁxed)
9.194 (ﬁxed)

ZNi [solar]
10.9 (9.14–12.7)bLine width [keV]
0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)0.200 (0.163–0.246)
kTe [keV]
1.43 (1.30–1.59)f
1.43 (1.30–1.59)fNormalizationa
1.61 (1.60–1.62)
Normalizationc
2.85 (2.59–3.11)
0.957 (0.750–1.16)
5.30 (4.20–6.39)
1.46 (1.27–1.65)
0.593 (0.338–0.847)
2.12 (1.81–2.43)
Normalizationc
11.2 (10.4–11.9)
0.257 (0.00–1.04)

The uncertainties in the parentheses are the 90% conﬁdence range. a The unit is 10−13 ne nH dV /(4πD2 ) (cm−5 ),where ne , nH , V , and D are the electron and hydrogen densities (cm−3 ), emitting volume (cm3 ), and distance tothe source (cm), respectively. b These values should be modiﬁed in the case of the overionized plasma. See discussion
for W49B for details. c The unit is 10−5 cm−2 s−1 . d Lines emitted by a recombination process. e Fixed at theexperimental values (Wargelin et al. 2005). f We assumed the same kTe .

5.3.3

Contribution of Recombination Lines

We have found the strong He-RRC of Fe from W49B. We also discovered RRC-accompanied
recombination lines, which may provide good diagnostics for the overionized plasma. We ﬁrst examine the validity of the best-ﬁt ﬂuxes of the RRC and recombination lines in table 5.2. The
recombination cross section of the H-like ions into a level of n is approximately given in equation 4.10. We apply this approximation for the He-like ions, but σ1 is reduced by half. In principle,
the recombination line ﬂux can be estimated by the branching ratio to various levels, but theseprocesses are very complicated. We, therefore, base our discussion only on a simple picture.
We compare the predicted capture and observed transition rates normalized with the n = 1
value. We can estimate the capture rates using equation 4.10 as σ2 /σ1 = 0.62, σ3 /σ1 = 0.25, and(σ4 + σ5 + · · ·)/σ1 = 0.34. On the other hand, the observed ﬂux rates are Heαrec /He-RRC = 0.47(±0.10), Heβrec /He-RRC = 0.053 (±0.023), and Heγ-∞rec /He-RRC = 0.19 (±0.03).6 By taking
the fractions between the observed and predicted rates, we obtain 76 (±16)%, 21 (±9)%, and 56(±9)% for the Heαrec , Heβrec , and Heγ-∞rec lines, respectively.
These fractions for the Heβrec and Heγ-∞rec lines may be conceivable, because one electron at
n = 1 suppresses direct transitions from excited levels (n ≥ 3→1) for He-like ions (ﬁgure 4.4). The
fraction of Heαrec is slightly smaller than expected because there should be a signiﬁcant contributionof the cascade decay electrons (n ≥ 3→2→1). The real Heαrec ﬂux may be somewhat larger, but
the Heαrec ﬂux ratio relative to total Heα is only <∼10% and does not aﬀect results and following
discussion.
6

Throughout this thesis, the errors in the parentheses are at 90% conﬁdence level.

CHAPTER 5. W49B

54

5.3.4

Electron and Ionization Temperatures

The electron temperatures determined by the bremsstrahlung continuum shape and the RRC
width are 1.52 (1.50–1.53) keV and 1.43 (1.30–1.59) keV, respectively. These consistent results
indicate a common origin of these emissions.
The ionization temperatures are determined from several ﬂux ratios. From the best-ﬁt modelin table 5.2, the ﬂux ratios of Lyα/Heα, He-RRC/Heα, and H-RRC/He-RRC are given as 0.016(0.014–0.018), 0.12 (0.11–0.13), and 0.023 (≤0.10). In ﬁgure 5.6, we compare these values withthe modeled emissivity ratios derived from the radiation code of Masai (1994) for a plasma of kTe= 1.5 keV. We obtain kTz = 2.58 (2.46–2.68) keV, 2.68 (2.63–2.73) keV, and 2.55 (≤3.65) keV,
respectively, for the above ratios. The ionization temperatures (∼2.7 keV) are signiﬁcantly higher
than the electron temperatures (∼1.5 keV), indicating that the plasma is in a highly overionizedstate.
The ﬁrst hint of overionized plasma in W49B was found by Kawasaki et al. (2005). Although
they analyzed diﬀerent elements (Ar and Ca) in diﬀerent energy bands (2.75–6.0 keV), and derived
kTz from the Lyα/Heα ratio using the CIE plasma code, their results (kTz ∼2.5 keV and kTe∼1.8 keV) are nearly consistent with ours. Our claim is more essential because it is based on clear
detection of the recombination structures.

Lyα/HeαHe-RRC/HeαH-RRC/He-RRC

Fig. 5.6.—: Predicted emissivity ratios of Lyα/Heα (black), He-RRC/Heα (dark gray), and H-RRC/HeRRC (light gray) of Fe as a function of the ionization temperature (kTz ) for an electron temperature of
1.5 keV (Masai 1994). The horizontal dashed lines represent 90% errors of the observed values. (Ozawa et
al. 2009a)

5.3. RESULTS FOR GLOBAL REGION

5.3.5

55

Abundance Correction for Overionization

The abundances listed in table 5.2 are valid only for the CIE state and should be modiﬁed in
the overionized case. Since no plasma code can be applied to the overionized plasma currently, we
make possible modiﬁcations using an available APEC code.
The Heα intensity is proportional to Z × (kTe ,kTz ), where Z and (kTe ,kTz ) are the abundanceof the element (solar) and the total emissivity for the Heα for kTe and kTz (cm3 s−1 ), respectively.
The emissivities of the He-, Li-, and Be-like ions for Fe and the He-like ions for Ni are modiﬁed bymultiplying the ion-fraction ratio between kTz = 2.7 keV and 1.5 keV (Mazzotta et al. 1998). The
total emissivity is given by adding those in individual ionization states. Multiplying by (1.5 keV,
1.5 keV)/ (1.5 keV, 2.7 keV), we obtain the real abundances in the overionized state as ZFe ∼3.0 solar and ZNi ∼2.8 solar. Both the elements are highly over abundant, indicating an ejecta origin of
the plasma.

5.3.6

Volume Emission Measure

To check the consistency of the common origin of the bremsstrahlung and RRC emissions, we
compare the volume emission measure (VEM) 7 .
The VEM of the VAPEC component (VEMVAPEC ) is derived from table 5.2 as VEMVAPEC= 1.61 (1.60–1.62) × 1013 cm−5 . On the other hand, the VEM of the RRC plasma (VEMRRC )is calculated from equation 4.12. According to Badnell (2006), the total radiative recombination
rate of He-like Fe at kTe = 1.5 keV is given as ∼3.9 ×10−12 cm3 s−1 . The recombination rate
into the ground state is given using equation 4.10 as σ1 /(σ1 + σ2 + · · ·) ∼ 0.45. The value ofnFe /nH is calculated using ZFe in previous subsection 5.3.5 and the number density of the solar
photosphere. Using the observed He-RRC ﬂux, we obtain8 VEMRRC = 1.3 (1.0–2.6) ×1013 cm−5 .
The two independent estimations of VEM give consistent results, supporting the same origin of theoverionized plasma.

7
8

See Section 4.4 for detailed derivation methods.
Here, we consider that the kTz error is 2.4–2.8 keV. The large error of VEMRRC is due to this eﬀect.

CHAPTER 5. W49B

56

5.4

Results for Segmented Region

In this section, we investigate the spatial distribution of the recombining plasma. We ﬁrst madecount maps to examine the plasma state. Figure 5.7 shows the NXB-subtracted and vignettingcorrected XIS band maps of Fe (a) Heα, (b) Heβ, and (c) He-RRC. All maps exhibit a center-ﬁlled
structure.
In ﬁgure 5.8, we represent the count ratio maps. The ratio Heβ/Heα (a) is an indicator of theelectron temperature, while the ratio He-RRC/Heα (b) shows the recombination degree. We cansee the He-RRC/Heα ratio is signiﬁcantly high in the west region.
D )H+HȘ NH9F )H+H55& NH9E )H+Hș NH9FRXQWV





 
FRXQWV
FRXQWV

Fig. 5.7.—: NXB subtracted and vignetting corrected XIS FI count maps shown on a linear intensity scale.
Each sell in these maps includes 32×32 CCD pixels. The red circle with a radius of 2.5 arcmin corresponds
to the source region shown in ﬁgure 5.3.

E +H55&+HȘ

D +Hș+HȘ







 







Fig. 5.8.—: XIS FI count ratio maps derived by taking ratios of ﬁgure 5.7. (a) Ratio between Fe Heβ(7.7–7.9 keV) and Heα (6.5–6.8 keV). (b) Ratio between Fe He-RRC (8.5–10 keV) and Heα (6.5–6.8 keV).
We only show the results of statistically-signiﬁcant regions inside the red circle with a radius of 2.5 arcmin.

5.4. RESULTS FOR SEGMENTED REGION

10:00

'

$ $
'
$ $
% %
% %
&
'
&
& &
'

05:00
09:00:00

Declination

15:00

57

'CUV

19:11:30

9GUV

11:00

10:30

Right ascension
Fig. 5.9.—: Regions used in the segmented analyses. We named Regions A to D and Segment 1 to 4 fora 16-fraction analysis (boxes). Ellipses are East and West regions we used in an east-west analysis. Gray
annulus represents the background region. We show the source region used in the global analysis with ablack circle as a reference.

16-fraction Analysis
The count ratio map showed a spatial diﬀerence of the RRC intensity. In order to quantitatively
conﬁrm this implication, we carry out the spectral analysis for the 16 segmented regions (ﬁgure 5.9).
We subtract the background spectrum taken from the orange annulus in ﬁgure 5.9 from the eachsource spectrum. Then we ﬁt the each spectrum with the same model and energy band with thosein table 5.2. Because of the limited photon ﬂux in 16-fraction analysis, however, we take some
restrictions into account as below: 1) Fix all the line centers and widths to those in table 5.2. 2)
Fix the ZNi /ZFe ratio to that in table 5.2. 3) Link electron temperatures of bremsstrahlung andrecombination components 9 . 4) Ignore a Fe H-RRC. 5) Fix ﬂux ratios between Heαrec , Heβrec ,Heγ-∞rec , and He-RRC to those in table 5.2 10 .
Figure 5.10 shows the ﬁtted spectra. We can see the ﬂux ratios between recombination (red)and bremsstrahlung (black) components are high in the west and low in the east. Figure 5.11 is
the ﬁtting results. Electron temperatures are high in east (Segment 1 and 2), while low in west
(Segment 3 and 4). In contrast, He-RRC/Heα ratios are high in west. Ionization temperatures
are also high in west reﬂecting the high He-RRC/Heα ratios. In ﬁgure 5.12, we show color maps,
which represent the best-ﬁt values in ﬁgure 5.11.

9This procedure is natural because these components are likely to be the same origin. See section 5.3.4 for details.
Strictly speaking, these ratios depend on the electron temperature. The inﬂuence, however, is small in the caseof kTe = 1.4-1.6 keV (we conﬁrm these values later). For example, the ratio Heαrec /He-RRC diﬀers only by 3% in
the above kTe range if we estimate from equation 4.10.
10

CHAPTER 5. W49B

58

$

$

$

$

%

%

%

%

&

&

&

&

'

'

'

'

Fig. 5.10.—: Spectral ﬁttings for the 16-fraction analysis. We modeled each spectrum with a CIE
component, a Fe Lyα line, He-RRC, and recombination lines.

5.4. RESULTS FOR SEGMENTED REGION

59

D

F

E

G

$ $
$ $
% '
%
% %
'
& &
& &
'

'

6HJPHQW6HJPHQW6HJPHQW6HJPHQW(photons cm-2 s-1 arcmin-2)'&%
5HJLRQ

Surface brightness$

H

Fig. 5.11.—: Fitting parameters for the 16-fraction analysis. (a) Electron temperatures. (b) Fe abundances. Note that these values should be modiﬁed in the overionized case. (c) He-RRC/Heα ﬂux ratios ofFe. (d) Ionization temperatures derived from (d) and a plasma radiation code of Masai (1994). (e) Surface
brightness in the 5.0-6.4 keV band. We excluded the Cr and Mn Kα contributions.

CHAPTER 5. W49B

60

F 55&+HȘIOX[UDWLR

D N7HE )H$EXQGDQFH

G N7]H 6XUIDFHEULJKWQHVV#NH9
D N7H E )H F IOX[ G N7] H 6%



㹼

㹼

㹼











㹼

㹼



㹼



㹼



㹼



㹼



㹼

㹼


㹼


㹼


㹼


㹼

㹼

SKFP
VDUFPLQ

㹼

NH9

㹼

UDWLR

㹼

VRODU

㹼

NH9

Fig. 5.12.—: Color maps representing the best-ﬁt values of physical parameters which are derived fromthe 16-fraction analysis. Red, orange, yellow, and green colors are in the order of highest to lowest for thephysical parameters. The meanings of (a) to (e) are the same as those in ﬁgure 5.11.

5.4. RESULTS FOR SEGMENTED REGION

61

East-West Asymmetry Analysis
In the count map and 16-fraction analyses, we can see the anisotropy of physical parameters in
the east and west regions of W49B. Hence, we directly investigate the diﬀerence between these tworegions. We took the spectra from the two ellipse regions in ﬁgure 5.9. Hereafter we call ”East”
and ”West” for the east and west ellipse regions, respectively. We took the background spectrum
from the orange annulus in ﬁgure 5.9 and subtracted it from each source spectrum.
We ﬁtted the spectra with the same model and restrictions with those of 16-fraction analysis.
The ﬁtted spectra and best-ﬁt parameters are summarized in ﬁgure 5.13 and table 5.3. The line
intensity of Fe Heα (∼6.7 keV) and Heβ (∼7.8 keV) is apparently more intense in East than West.
On the contrary, the bump around ∼9 keV is evidently intense in West.
In ﬁgure 5.14, we show the electron and ionization temperatures of the East and West regions.
We derived the ionization temperatures from the He-RRC/Heα ﬂux ratios of Fe and ﬁgure 4.12. The
electron temperature of West (∼1.4 keV) is lower than that of East (∼1.7 keV). On the other hand,
ionization temperature of West (∼3.0 keV) is signiﬁcantly higher than that of East (∼2.6 keV).
The recombination and overionization degree is explicit in the West region, where the ionization
temperature is more than twice as high as the electron temperature.
Table 5.3—: Best-ﬁt spectral parameters in the East and West regions.

CIE (VAPEC)

Cr Kα
Mn Kα
Fe Heαrec f
Fe LyαFe Heβrec f
Fe Heγ-∞rec f
Fe He-RRC

kTe [keV]ZFe [solar]
ZNi [solar]
Normalizationc
Center energyd
5.655 (ﬁxedg)
6.162 (ﬁxedg)
6.666 (ﬁxedh)
6.961 (ﬁxedg)
7.880 (ﬁxedh)
8.538 (ﬁxedg)
Edge energyd
8.830 (ﬁxed)

Line widthd
0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)0.20 (ﬁxedg)
kTe [keV]
linkediEast
West
1.64 (1.60–1.68) 1.42 (1.41–1.45)
5.18 (5.00–5.37)a 3.32 (3.18–3.43)a
2.45×[ZFe ]b
0.42 (0.42–0.44) 0.75 (0.72–0.76)
Normalizatione
1.14 (0.97–1.31) 1.10 (0.81–1.20)
0.45 (0.32–0.59) 0.20 (0.06–0.35)
0.47×[RRC Norm]b
0.63 (0.45–0.77) 0.47 (0.34–0.61)
0.053×[RRC Norm]b
0.19×[RRC Norm]b
Normalizatione
3.01 (2.58–3.41) 4.51 (3.98–4.93)
Should be modiﬁed in the overionized case. b Linked following table 5.2. c The unit is 10−13 ne nH dV /(4πD2 )(cm−5 ), d (keV), and e 10−5 (cm−2 s−1 ). f Lines emitted by a recombination process. g Fixed following table
5.2. h Fixed at the experimental values (Wargelin et al. 2005). i Linked to kTe of the CIE component (VAPEC).
a

CHAPTER 5. W49B

62

D (DVWE :HVW

Fig. 5.13.—: Spectral ﬁttings for the (a) East and (b) West regions. We modeled spectra with a CIE
component, a Fe Lyα line, He-RRC, and recombination lines.

Fig. 5.14.—: Electron and ionization temperatures in the East and West regions of W49B.

5.5

Summary

We have presented an XIS spectrum of unprecedented quality of the Galactic SNR W49B
obtained with Suzaku. The results and interpretations are summarized as follows:• The spectrum exhibits an unusual structure consisting of a saw-edged bump above 8 keV. We
ﬁrmly conclude that this bump is caused by the strong radiative recombination continuum(RRC) of iron, detected for the ﬁrst time in an SNR.
• The electron temperature derived from the bremsstrahlung and RRC spectral shape is∼1.5 keV. On the other hand, the ionization temperature derived from the observed intensity ratios between the RRC and lines of iron is ∼2.7 keV. Thus, the plasma is in a highly
overionized state.
• Volume emission measures independently determined from the ﬂuxes of the bremsstrahlung
and recombination components are consistent, suggesting the same emission origin.
• There is signiﬁcant spatial anisotropy of physical parameters between the west and east area of
W49B. The strong plasma recombining region corresponds to the X-ray dim and low electron
temperature area.

Chapter 6

IC 443
IC 443 is the ﬁrst SNR whose plasma was claimed to be in the overionized state by the ASCA
data (Kawasaki et al. 2002). The X-ray spectrum exhibits strong Lyα lines of Si and S, suggesting
the plasma is highly ionized. Thus, this SNR is the best candidate which possesses a recombining
plasma. In this chapter, we summarize the Suzaku results and discuss the plasma state 1 .

6.1

Previous Results

IC 443 (G189.1+3.0, 3C157) is a Galactic mixed-morphology SNR. Figure 6.1 shows the multiwavelength view of IC 443. It has a center ﬁlled morphology in the X-ray band and a shell-likestructure in the radio and optical bands (Rho & Petre 1998; Duin & van der Laan 1975; Braun& Strom 1986; Leahy 2004). It is located at a distance of 1.5 kpc (Welsh & Sallmen 2003) near
the Gem OB1 association and a dense giant molecular cloud (Cornett et al. 1977) with OH maser
emission (Claussen et al. 1997). A comprehensive X-ray study of IC 443 was ﬁrst made with the
Einstein and HEAO-A2 satellites (Petre et al. 1988). They estimated the SNR age to be ∼3000 yr.
Recently, Troja et al. (2008) derived the age of ∼4000 yr from the morphologies of the shocked
ejecta and interstellar medium (ISM) using XMM-Newton data. Thus, IC 443 is a middle-aged
SNR. In ﬁgure 6.2, we show the XMM-Newton MOS images of IC 443. We can see the ring-shaped
structure which is considered to be ejecta recently heated by the reverse shock (ﬁgure 6.2b).
The northeast shock front has been decelerated by the encounter with a H I cloud (Denoyer
1978; Dickel et al. 1989; Rho et al. 2001). It is strongly interacting with a nearby giant molecular
cloud at several positions along the southern rim (Cornett et al. 1977; Giovanelli & Haynes 1979;
Troja et al. 2006).

1

Details are also in Yamaguchi et al. (2009).

63

CHAPTER 6. IC 443

64

DUFPLQ

Fig. 6.1.—: Multiwavelength view of IC 443 (Gaensler et al. 2006). (Red) 670 nm emission [Second
Palomar Observatory Sky Survey]. The bright star to the west is η Gem. (Green) 1.4 GHz radio data
[DRAO Synthesis Telescope]. (Blue) 0.1–2.4 keV X-ray data [ROSAT PSPC]. (White contour) 8.5 GHz
data, indicating a pulser wind nebula [Very Large Array]. We superimpose two yellow squares to show
Suzaku XIS FOVs. Because IC 443 is extended with a diameter of ∼45 arcmin, two FOVs cannot cover thewhole remnant.

D

E

Fig. 6.2.—: XMM-Newton MOS image of IC 443. (a) Count map in the 1.4–5.0 keV energy band. Contour
levels correspond to 0.4, 1.0, 2.7, and 5.8 × 10 −2 cts s−1 arcmin−2 . Black crosses mark the position of the
12 hard X-ray point sources identiﬁed by Bocchino & Bikov (2003). The black arrow indicates the directionof PWN’s motion (Gaensler et al. 2006). (b) Equivalent width map of S lines. We can see the ring-shaped
structure. This is considered to be the ejecta recently heated by a reverse shock. Figures are taken from
Troja et al. (2008). We superimposed yellow squares to show the XIS FOVs.

6.2. SUZAKU OBSERVATION

65

The bulk of its X-ray emission is thermal, fairly well described with two CIE components with
∼0.2 and ∼1.0 keV (Petre et al. 1988; Troja et al. 2006). Using Chandra data, Olbert et al. (2001)
determined that the hard X-ray source at the southern edge of radio shell is a synchrotron pulsar
wind nebula (PWN) powered by a compact source, which is possibly associated with the remnant.
Hence, IC 443 is regarded as a core collapse SNR. The PWN can be identiﬁed by radio data withwhite contours in ﬁgure 6.1.
A previous ASCA observation revealed that the Lyα ﬂuxes of Si and S are signiﬁcantly higher
than those expected from the electron temperature (Kawasaki et al. 2002). They then derived an
ionization temperature of ∼1.5 keV, signiﬁcantly above the electron temperature of ∼1.0 keV which
was determined by its continuum shape of the spectrum. Thus, it was argued that the thermal
plasma is in the overionized state (Kawasaki et al. 2002). On the other hand, XMM-Newton foundthat the plasma is in a CIE, or the overionization is only marginal (Troja et al. 2008).

6.2

Suzaku Observation
Table 6.1—: Observation log of IC 443.

TargetStart Date

Obs.IDIC 443 North
IC 443 South

2007-03-06
2007-03-07

501006010
501006020

R. A.
Dec.
(J2000.0)hm
06 17 11s 22◦ 46 43
06h 17m 11s 22◦ 29 02Exposure
(ks)
42.0
44.0

Suzaku observed IC 443 with two pointings. The observation log is summarized in table 6.1.
These two ﬁeld of views (FOVs) are shown with yellow squares in ﬁgure 6.1 and 6.2. Note that
IC 443 is extended ∼45 arcmin that two FOVs cannot cover the whole remnant.
Three of the XISs were operated in the normal full-frame clocking mode with a SCI techniqueduring our observations. Two are FI CCDs and the other is a BI CCD. We used HEASOFT version
6.5.1 for data reduction and XSPEC version 11.3.2 for spectral analysis. After screening with thestandard criteria, the net integration time is ∼40 ks for both observations.

CHAPTER 6. IC 443

66

6.3

Results for Global Region

6.3.1

Overall Structure

Figure 6.3 shows the vignetting-corrected XIS image in 0.7–3.0 keV, the energy band including
the major K-shell emission lines from Ne, Mg, Si, and S, and L-shell emission blends from Fe. We
can see the X-ray emission covers most of the FOVs, and is bright in the north FOV. We also
show the hard (3.0-5.5 keV) emission with orange contours. The bright hard compact source at thesouthern edge corresponds to the PWN.
We extracted the source spectrum from an X-ray bright region with a 7.5 arcmin radius
circle. Because there is no nearby background observation for IC 443, we subtracted the non-X-ray
background (NXB) spectrum constructed with the xisnxbgen software from the source spectrum.
The spectra of the two FI sensors were merged to improve the statistics. Figure 6.4 shows themerged FI spectrum. We can see prominent lines of K-shell emission from He- and H-like ions ofO, Ne, Mg, Si, S, Ar, Ca, and Fe as well as L shell emission blends of Fe. We successfully resolved
the Fe Kα line for the ﬁrst time.
23°00’

Declination

22°50’

22°40’

22°30’

PWN
22°20’
6h18m

6h17m

6h16m

Right ascension
Fig. 6.3.—: Vignetting-corrected XIS image of IC 443 in the soft (0.7–3.0 keV; gray scale) and hard
(3.0–5.5 keV; contour) bands shown on a logarithmic intensity scale. Data from the three active XISs are
combined. The white circle indicates the source region. The black circle and ellipse are the source andbackground regions of the pulser wind nebula used in the following segmented analysis. The XIS ﬁeld ofviews are shown by black squares.

6.3. RESULTS FOR GLOBAL REGION67

In ﬁgure 6.4, we also show the cosmic X-ray background (CXB) level, which is still includedin the NXB-subtracted spectrum. which is modeled following Kushino et al. (2002). This CXB
spectrum is given by a power-law with an index of 1.412, and a ﬂux in the 2–10 keV band with 6.38
× 10−8 erg cm−2 s−1 sr−1 . Throughout the spectral ﬁtting, we ﬁxed the CXB to the above value.
Note that the contribution of the Galactic ridge emission is negligible because IC 443 is located at
the Galactic anti-center direction (l=189.1).

6.3.2

Spectral Analysis

As shown in ﬁgure 6.4, the wide-band spectrum includes too many emission lines and possibly
other complicated structures. In order to precisely examine the ionization states of Si and S, wehereafter focus on the spectrum in the 1.75–6.5 keV energy band. During the spectral analysis,
we ﬁxed the interstellar extinction to a hydrogen column density of 7 × 1021 cm−2 with the solar
elemental abundances, following Kawasaki et al. (2002) and Troja et al. (2008).
The centroids of the Lyα lines of Si and S for the FI spectrum are measured with a power-law
and Gaussian line model, and compared with the canonical values of the Astrophysical Plasma
Emission Database (APED: Smith et al. 2001). The averaged center energy diﬀerence was +4 eV.
Therefore, we added a 4 eV oﬀset in the FI spectrum. In the same way, an oﬀset of −8 eV wasadded to correct the energy scale for the BI spectrum.

Fe L

Si
S
Ca
Ne

O

Fe

Mg
Ar

level

Fig. 6.4.—: NXB-subtracted XIS FI spectrum extracted from the red circle region in ﬁgure 6.3. The
contribution of the cosmic X-ray background which is still included in the spectrum is shown by a solid line.

CHAPTER 6. IC 443

68
CIE Model Fitting

We ﬁrst ﬁtted the spectrum with a VAPEC model. The Si, S, and Ar abundances (hereafter
ZSi , ZS , and ZAr ) were free parameters, while the Ca abundance (ZCa ) was tied to that for Ar
because the Ca Kα line is not intense enough to independently determine the value. In the initialﬁt, we found a signiﬁcant inconsistency between the FI and BI spectra around 1.84 keV, whichcorresponds to the neutral Si K-edge energy. This is due to the well-known calibration issue of the
XIS. Since the calibration for the FI is currently better than that of the BI, we decided to ignore
below 1.9 keV for the BI spectrum.
Figure 6.5a shows the ﬁtting result of the one-VAPEC model, with a temperature of 1.1 keV.
The model exhibited signiﬁcant excess around the Lyα lines of S and Ar, and around 2.7 keV and3.5 keV, and hence was rejected with a large χ2 /dof of 1466/242.
We then tried a two-VAPEC model, assuming equal abundances between the two components.
However, it did not improve the ﬁt with the temperatures of ∼0.2 keV and ∼1.1 keV, and theresults were essentially the same as ﬁgure 6.5a with a large χ2 /dof of 1431/240. Adding a further
plasma model component did not improve the ﬁtting any more. In every case, the bumps around2.7 keV and 3.5 keV were present.
We checked the light curves of the source region in the 2.5–4.0 keV band. They showed almost
constant ﬂuxes, indicating no ﬂare-like event had occurred during the observations. We detected
the bumps in all XIS sensors. Thus, the bumps are real structures.

+ FI
+ BI

E

D

CXB level
S Lyα Ar Lyα

Si
Lyα

Si
H-RRC
S H-RRCSiLyβ SLyα Ar Lyα

Fig. 6.5.—: (a) XIS spectra in the 1.75–6.5 keV band. The best-ﬁt VPAEC model is shown by a solidline. The lower panel shows the residuals from the best-ﬁt model. There are bumps at the energies of Lyαlines of S and Ar (down-pointing arrows), and around 2.7 keV and 3.5 keV (up-pointing arrows). (b) Sameas (a), but with radiative recombination continua of Si and S and Lyα and Lyβ lines of Si, S, and Ar. The
residuals seen in (a) have disappeared.

6.3. RESULTS FOR GLOBAL REGION

69

Recombination Structures
The energies of the bumps at ∼2.6 keV and ∼3.3 keV correspond to the Kα potential energiesof H-like S and Ar ions, respectively. On the other hand, around the energies of the bumps for∼2.7 keV and ∼3.5 keV, there is no emission line candidate for an abundant element. Furthermore,these bumps are too broad to be accounted for by a single emission line. The lower edges ofunidentiﬁable bumps correspond to the electron K-shell binding potentials of H-like Si (2.666 keV)and S (3.482 keV). This suggests that these bumps are likely due to RRC.
We hence introduced the RRC and Lyα lines of abundant elements (Si, S, and Ar). Lyβ linesof these elements are also introduced, but they are not signiﬁcant except for Si. The widths andcenter energies of these lines are ﬁxed at zero and their theoretical values, respectively. The ﬂuxesof all lines are free parameters. We assume that the RRC is expressed following equation 4.7. This
formula gives a good approximation when kTe  Eedge . In the case of IC 443, kTe is ∼0.7 keV, as
we conﬁrm later, while Eedge of Si and S are ∼2.7 keV and ∼3.5 keV, respectively. Thus, we cansafely adopt this formula. During the ﬁtting in this subsection, we ignore the cascade lines. We
consider them in subsection 6.3.3.
The resultant best-ﬁt parameters and models are given in table 6.2 and ﬁgure 6.5b, respectively.
With this model, χ2 /dof is greatly improved to 331/235, although this is still unacceptable in apurely statistical sense. This would be partially due to the calibration uncertainty. We can see
the wave-shaped residuals in ﬁgure 6.5b around ∼2 keV. The XIS calibration around ∼2 keV isdiﬃcult and uncertainty of the gain and response is relatively high because there is a very limited
number of calibration sources which can be used in this energy band. If we ignore below 2.1 keV, a
χ2 /dof of 251/208 with null hypothesis probability of 0.021 is obtained. In this chapter, we discuss
the implications based on the result in table 6.2.
Table 6.2—: Best-ﬁt spectral parameters of IC 443.

CIE (VAPEC)

kTe [keV]ZSi [solar]
ZS [solar]
ZAr =ZCa [solar]
Normalizationa

Line
Si LyαSi LyβS LyαAr LyαRRCSi H-RRCS H-RRC

0.74 (0.71–0.78)
0.82 (0.79–0.89)b
1.40 (1.34–1.44)b
1.96 (1.76–2.10)b
1.66 (1.56–1.75)
Center energy [keV]
2.006 (ﬁxed)
2.377 (ﬁxed)
2.623 (ﬁxed)
3.323 (ﬁxed)
Edge energy [keV]
2.666 (ﬁxed)
3.482 (ﬁxed)

Width [keV]
Normalizationc
0 (ﬁxed)76 (69–78)
0 (ﬁxed)5.5 (3.8–7.0)
0 (ﬁxed)25 (23–26)
0 (ﬁxed)3.7 (3.1–4.3)
kTe [keV]
Normalizationc
0.58 (0.53–0.62)d 51 (45–55)
0.58 (0.53–0.62)d 13 (12–15)

The unit is 10−13 ne nH dV /(4πD2 ) (cm−5 ), where ne , nH , V , and D are the electron and hydrogen densities(cm−3 ), emitting volume (cm3 ), and distance to the source (cm), respectively. b Should be modiﬁed in theoverionization case (subsection 6.3.5). c The unit is 10−5 cm−2 s−1 . d We assumed the same kTe .
aCHAPTER 6. IC 443

70

6.3.3

Contribution of Recombination Lines

We have found the strong H-RRC of Si and S from IC 443 for the ﬁrst time, hence there shouldbe accompanying recombination (cascade) lines. We ﬁrst examine the contribution of cascade lines.
In the spectral analysis, we required the Lyα lines of Si, S, and Ar, and Lyβ of Si. These lines are
partially due to a greater number of H-like ions in the overionized state than in the CIE, and theremains originate from the recombination process. In the previous analysis, however, we did notconsider the cascade lines from upper levels (n≥4→1). Here we estimate this eﬀect.
We ﬁx the contribution of the cascade lines of H-like Si and S using equation 4.10. We assume
the electrons trapped in the upper levels straightforwardly fall into the ground state in the caseof H-like ions. Based on this assumption, the prospective ﬂux ratios of Lyαrec /H-RRC, Lyβrec /HRRC, and Lyγ-∞rec /H-RRC are derived from σ2 /σ1 , σ3 /σ1 , and (σ4 +σ5 +σ6 +...)/σ1 , respectively.
The resultant ratios are 0.27, 0.10, and 0.13 for Si, and 0.29, 0.12, and 0.15 for S, respectively. In
the following ﬁtting, we ﬁx these ratios to above values.
The widths and center energies of Lyαrec and Lyβrec are ﬁxed to 0 and theoretical values,
respectively. Those of Lyγ-∞rec are derived using the results of W49B as follows. The center
energy is determined assuming the equation
E(Heγ -∞rec ) E(Lyγ -∞rec )=,E(Heαrec )E(Lyαrec )(6.1)

where E(Heγ-∞rec ), E(Heαrec ), E(Lyγ-∞rec ), and E(Lyαrec ) are the observed center energy of Fe
Heγ-∞rec for W49B (8.538 keV), the ﬁxed one of Fe Heαrec for W49B (6.666 keV), the one we wantto know of Si (or S) Lyγ-∞rec for IC 443, and the theoretical one of Si (or S) Lyαrec (2.006 keVfor Si and 2.623 keV for S), respectively. In the same way, the width is determined assuming theequation
W (Heγ -∞rec ) W (Lyγ -∞rec )
=
,
(6.2)
E(Heγ -∞rec )E(Lyγ -∞rec )where W (Heγ-∞rec ) and W (Lyγ-∞rec ) are the observed line width of Fe Heγ-∞rec for W49B
(0.200 keV), and the one we want to know of Si (or S) Lyγ-∞rec for IC 443, respectively. The
resultant center energies and widths of Lyγ-∞rec are 2.57 keV and 0.060 keV for Si, and 3.36 keV
and 0.078 keV for S, respectively.
We ﬁx center energies and widths of cascade lines and ﬁt the spectra (ﬁgure 6.6). The resultant parameters including the temperatures of the bremsstrahlung and recombining components(∼0.74 keV and ∼0.57 keV, respectively), and the RRC ﬂuxes (∼5.0 × 10−4 cm−2 s−1 for Si H-RRC
and ∼1.4 × 10−4 cm−2 s−1 for S H-RRC) do not signiﬁcantly diﬀer from those in table 6.2 within
the margin of statistical error. Thus we can safely say that the inclusion of cascade lines does notchange the result.

6.3. RESULTS FOR GLOBAL REGION

71

D

E

Fig. 6.6.—: (a) Same as ﬁgure 6.5b, but cascade lines (Lyαrec , Lyβrec , and Lyγ-∞rec ) of H-like Si and S
are added. (b) Same as ﬁgure 6.5b, but He-RRC of Si and S are added.

6.3.4

Electron and Ionization Temperatures

The electron temperatures in table 6.2 determined by the bremsstrahlung and RRC shapes
are 0.74 (0.71–0.78) keV and 0.58 (0.53–0.62) keV, respectively. The kTe derived from thebremsstrahlung is slightly higher than that from the RRC. However, if we introduce the He-RRC
of Si and S (ﬁgure 6.6b), kTe become 0.68 (0.65–0.72) keV and 0.62 (0.57–0.65) keV for each component. This procedure is natural, because in the case of kTz ∼1.0 keV for Si and kTz ∼1.2 keVfor S as we conﬁrm later, fully-ionized/H-like ion fraction ratios are ∼1/2 for Si and ∼1/5 for S.
Considering that the recombination cross section into the ground state for fully-ionized ions is about
twice as large as H-like ions, the resultant He-RRC ﬂux is reasonable. Although we cannot strongly
insist the existence of the He-RRC, it is plausible to think a common origin of the bremsstrahlung
and recombining components with kTe ∼0.7 keV.
The ionization temperatures are derived from H-RRC/Heα ﬂux ratios of Si and S using ﬁgure 4.12. From the best-ﬁt model in table 6.2, these ratios are given as 0.11 (±0.01) and 0.075(±0.01), respectively. Assuming kTe = 0.7 keV, we obtain kTz = 0.96 (±0.04) keV for Si and1.18 (±0.05) keV for S. The ionization temperatures (0.9–1.2 keV) are signiﬁcantly higher than theelectron temperatures (0.5–0.8 keV), conﬁrming that the plasma is in a highly overionized state.
Kawasaki et al. (2002) found the ﬁrst hint of overionized plasma in IC 443. Because they
derived kTe without considering the contribution of the RRC, they obtained relatively high valueof ∼1.0 keV. They derived kTz ∼ 1.5 keV from the Lyα/Heα ratio of S using the CIE plasma code.
Our claim of overionization is more essential because it is based on clear detection of the RRC.

CHAPTER 6. IC 443

72

6.3.5

Abundance Correction for Overionization

The abundances listed in table 6.2 are valid only for the CIE state and should be modiﬁed in
the overionized case. We make modiﬁcations using an available APEC code.
The Heα intensity is proportional to Z × (kTe ,kTz ), where Z and (kTe ,kTz ) are the abundanceof the element (solar) and the total emissivity for the Heα for kTe and kTz (cm3 s−1 ), respectively.
In the kT range of 0.7–1.2 keV, most Kα emissions of Si or S originate from He-like ions. Hence,the emissivity of the He-like ions is modiﬁed by multiplying the ion-fraction ratio between kTz =1.0 keV (for Si) or 1.2 keV (for S) and 0.7 keV. The ion fraction of Si in kTz = 0.7 keV and 1.0 keV
are 0.72 and 0.35, respectively. Thus, we can obtain the real abundance of Si is 0.82 × (0.72/0.35)∼1.7 solar. In the same way, the ion fraction of S in kTz = 0.7 keV and 1.2 keV are 0.90 and 0.59,
respectively. Thus, we obtain the real abundance of S is 1.4 × (0.90/0.59) ∼2.1 solar.

6.3.6

Volume Emission Measure

We compare VEM to check the consistency of the common origin of the bremsstrahlung and
RRC emissions. The VEM of the VAPEC component is derived from table 6.2 as ∼1.7 × 1013 cm−5 .
On the other hand, the VEM of the RRC plasma (VEMRRC ) is calculated using equation 4.12.
According to Badnell (2006), the total radiative recombination rate at kTe = 0.7 keV is given as∼2.0 × 10−12 cm3 s−1 for H-like Si. The recombination rate into the ground state is given using
equation 4.10 as σ1 /(σ1 +σ2 +···) ∼ 0.67. The value of nSi /nH is calculated using ZSi in the previous
subsection. Using the observed H-RRC ﬂux of Si and κfully−ionized(Si) (1.0 keV) = 0.19, we obtain
VEMRRC(Si) ∼3.3 × 1013 cm−5 . In the same way for S, using the total radiative recombination rate
∼2.8 × 10−12 cm3 s−1 , σ1 /(σ1 + σ2 + · · ·) ∼ 0.64, and κfully−ionized(S) (1.2 keV) = 0.06, we obtain
VEMRRC(S) ∼3.7 × 1013 cm−5 . These values are almost consistent with that of the bremsstrahlungcomponent. Two independent estimations of VEM give nearly consistent results, supporting thesame origin of the bremsstrahlung and recombination components in IC 443.

6.4. RESULTS FOR SEGMENTED REGION

6.4

73

Results for Segmented Region

In this section, we investigate the spatial distribution of the recombining plasma. We divided
each FOV into 16 regions as is shown in ﬁgure 6.7. We name Region A to H from north to south,and Segment 1 to 4 from east to west. We used only FI data for simplicity in this section. We
subtract the NXB spectrum from each source spectrum. We modeled the CXB with a power-law
model.
The PWN detected nearly at the edge of the south FOV is reported to be extended ∼2 arcmin(Gaensler et al. 2006). We therefore modeled the PWN spectrum and ﬁxed it during the analysis.
We extracted the PWN spectrum from the blue circle region in ﬁgue 6.3 with a radius of 2 arcmin
and background from the surrounding ellipse region. The background subtracted spectrum is wellapproximated by a power-law model with a photon index of Γ = 1.78 and interstellar extinction of
NH = 4.4 × 10 21 cm−2 (ﬁgure 6.8). During the ﬁtting in Region G and H in ﬁgure 6.7, we ﬁxedthe ”shape” of the PWN spectrum with this model. The normalization of this component is a freeparameter.
23°00’$ $ $ $

Declination

22°50’% % % %

22°40’& & & &' ' ' '

22°30’( ( ( () ) ) )
* * * *

22°20’
6h18m+ + + +
6h17m

Right ascension
Fig. 6.7.—: Segmented regions used in our analysis. We named Region A to H from north to south, andSegment 1 to 4 from east to west. Contours indicate hard (3.0-5.5 keV) X-ray intensity on a logarithmic
scale.

74

CHAPTER 6. IC 443

Fig. 6.8.—: Background subtracted spectrum of the pulser wind nebula. The best-ﬁt power-law model isshown by a solid line.

We ﬁtted each spectrum with the same model and energy band with those in table 6.2. Becauseof the limited photon statics, however, we ignored the Si Lyβ line. We further linked the temperatures of the CIE and RRC components because it is plausible that these components are the same
origin. Figure 6.9 shows the ﬁtted spectra. We excluded regions A1 and A4 because of poor statics.
We can see the ﬂux ratio between the RRC (red) and CIE (black) components is high around theborder area of the two FOVs, while low in the X-ray bright north area.
We can see the hard X-ray component representing the PWN spectrum (shown in green curves
in ﬁgure 6.9) signiﬁcantly contaminate G1–G3 and H1–H3 regions. We should note that at theboundary regions of G2 and G3, and G1 and H1, hard X-ray point sources are detected with
Chandra (Bocchino & Bicov 2003; the black crosses in ﬁgure 6.2). The ﬂuxes of these two sourcesin the 2.0–10.0 keV band are both around one tenth of the PWN. Hence, the power-law spectra inthese regions would include the emission from these point sources. Because spectra of these point
sources are unclear, we should care that the systematic uncertainty of the parameters derived from
Region G and H is high.
Figure 6.10 is the ﬁtting results. We show the parameters of electron temperatures (a), ZSi (b)and ZS (c) without correction of the overionization eﬀect, H-RRC/Heα ratios of Si (d), ionization
temperatures derived from the Si H-RRC/Heα ratios and ﬁgure 4.12 (e), and surface brightness in
the 1.75–6.5 keV band excluding PWN and CXB contributions (f).
Electron temperatures are high in the X-ray bright north region (Region A, B and C, especiallyin Segment 2 and 3), while low in the dim region (Region D, E, F, and G). The H-RRC/Heα ratioof Si, which is a good indicator for the recombination degree, is low in the X-ray bright region,
while high in the dim regions. In ﬁgure 6.11, we show color maps, which represent the best-ﬁtvalues in ﬁgure 6.10.

6.4. RESULTS FOR SEGMENTED REGIONCounts s−1keV−1
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0.1

10−3

0.1

10−3

2

5
Energy (keV)

2

5Energy (keV)

2

5Energy (keV)

52Energy (keV)

Fig. 6.9.—: Spectral ﬁttings for the segmented regions. We modeled each spectrum with a CIE component,
Lyα lines of Si, S, and Ar, and H-RRC of Si and S. The dotted curves represents the contribution of thepulser wind nebula and possibly other point sources.

CHAPTER 6. IC 443

76

D

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E

Fig. 6.10.—: Fitting parameters for the segmented analysis. (a) Electron temperatures. Si (b) and S (c)abundances. Note that these values should be modiﬁed in the overionized case. (d) H-RRC/Heα ﬂux ratiosof Si. (e) Ionization temperatures derived from (d) and a plasma radiation code of Masai (1994). (f) Surface
brightness in the 1.75–6.5 keV band. We excluded the PWN and CXB contribution.

6.4. RESULTS FOR SEGMENTED REGION77

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Fig. 6.11.—: Color maps derived from the segmented analysis. Red, orange, yellow, and green colors arein the order of highest to lowest for the physical parameters. The meanings of (a) to (f) are the same as
ﬁgure 6.10.

CHAPTER 6. IC 443

78

6.5

Summary

We have presented the Suzaku XIS spectroscopic study of IC 443 in the X-ray bright northregion. We have also examined the spatial distribution of the physical parameters. The results andinterpretations are summarized as follows:• Many emission lines from highly ionized heavy elements are clearly detected. Fe XXV Kα is
resolved for the ﬁrst time.
• The X-ray spectrum in the 1.75–6.0 keV band is described by a thin-thermal plasma with theelectron temperature of ∼ 0.7 keV and several additional Lyman lines. In addition, we ﬁrmly
detect strong RRC of H-like Si and S for the ﬁrst time.
• The ionization temperatures of Si and S determined from the intensity ratio between the
RRC and He-like Kα line are ∼ 1.0 keV and ∼ 1.2 keV, respectively. We thus ﬁnd evidence
for strongly recombining and highly overionized plasma.
• Volume emission measures independently determined from the ﬂuxes of the bremsstrahlung
and recombination components are almost consistent, indicating the same plasma origin.
• Spatial analysis indicates the signiﬁcant anisotropy of physical parameters in IC 443. Strongly
recombining region corresponds to the high gradient area of the electron temperature andsurface brightness.

Chapter 7

G359.1-0.5
Bamba et al. (2000) claimed that the ionization degree of Si and S in the Galactic SNR
G359.1-0.5 is high using the ASCA data. Thus, this SNR is a candidate of the overionized andrecombining plasma. We present a Suzaku spectroscopic analysis of G359.1-0.5.

7.1

Previous Results

G359.1-0.5 was identiﬁed as an SNR by radio observations (Downes et al. 1979). Uchida,
Morris, and Yusef-Zadeh (1992) found the structure of the radio continuum emission surroundedby the 12 CO ring. Comparing the 21 cm absorption feature of the 12 CO ring with the Galacticrotation curve, they concluded the location of this SNR to be near to the Galactic center. We cansee the shell-like radio morphology with a radius of 12 arcmin in ﬁgure 7.1a (gray scale). Unlike
the radio, X-ray emission shows the center-ﬁlled morphology. Thus this SNR is classiﬁed as amixed-morphology SNR.
Bamba et al. (2000) analyzed the X-ray spectrum of G359.1-0.5 taken with ASCA (ﬁgure 7.1).
It shows prominent Kα lines of Si and S. They argued that the line center of Si Kα line wasconsistent with that from He-like ions, whereas that of S was consistent from H-like ions. However,this is unusual since the heavier elements are generally less ionized compared to the lighter ones in
the other typical SNRs. In the ASCA GIS spectrum, however, the H-like and He-like lines could
not be separately detected because of its limited energy resolution. It is important to separate
these lines and examine the ionization states.

79

CHAPTER 7. G359.1-0.580Ȯ*

10-2

(b)

10-3

(a)

6

-2 0 2

Counts s-1 keV-1

6L

1

5

2

Energy (keV)
Fig. 7.1.—: (a) ASCA GIS contour map in the 1.6–2.1 keV band, superposed on the radio map (grayscale) by Gray (1994). The contour levels are linearly spaced and are saturated at A 1742-294 and 1E 1740.72942. (b) Background-subtracted ASCA GIS spectrum of G359.1-0.5. The plasma required two temperaturecomponents to describe prominent Kα lines of Si and S. These ﬁgures are taken from Bamba et al. (2000).

7.2

Suzaku Observation

We observed G359.1-0.5 with Suzaku (table 7.1). We used HESS J1745-303 data, a TeV
unidentiﬁed source, as background. The X-ray emission from the HESS source has not been detected, and Bamba et al. (2009) derived a tight upper limit on the 2–10 keV emission as 2.1 ×
10−13 erg s−1 cm−2 . Hence we can safely use the HESS observation for background.
Three of the XISs (two for FI and one for BI) were operated in the normal full-frame clocking
mode with a SCI technique during all the observations. Data were cleaned using processing version
2.2.11.22. We used HEASOFT version 6.5.1 for data reduction and XSPEC version 11.3.2 forspectral analysis. After screening with the standard criteria, the net integration time is ∼60 ks forthe source data and ∼70 ks for the background data.
Table 7.1—: Observation log of G359.1-0.5 and HESS J1745-303.

∗

TargetStart Date

Obs.ID

G359.1-0.5 ∗HESS J1745-303 ∗∗2008-09-14
2008-03-02

503012010
502016010

Used as source data.∗∗ Used as background data.

R. A.
Dec.
(J2000.0)hm
17 45 12s -29◦ 45 12
17h 44m 54s -30◦ 06 35Exposure
(ks)
57.7
70.5

7.2. SUZAKU OBSERVATION

81

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ᵥᵿᶊᵿᶁᶒᶇᶁᴾᶊᶍᶌᶅᶇᶒᶓᶂᶃᴾᵆᶂᶃᶅᶐᶃᶃᵇᵑᵓᵖᵌᵔ

Fig. 7.2.—: Vignetting-corrected XIS image of G359.1-0.5 shown on a logarithmic intensity scale. Data
from the three active XISs are combined. The yellow ellipse and white pentagon indicate the source andbackground regions, respectively.

CHAPTER 7. G359.1-0.5

82

7.3

Analysis and Results for Global Region

7.3.1

Overall Structureᵤᶃᴾᵶᵶᵴᵧ

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ᵱᶇᴾᵶᵧᵴ
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ᵤᶃᴾᵶᵶᵴ

Figure 7.2 shows the vignetting-corrected XIS image in several energy bands. The 1.76–
1.96 keV band (ﬁgure 7.2d) corresponds to Kα emission of He-like Si. A strong Si Kα line ischaracteristic of SNR plasma. We can see a bright emission inside the yellow ellipse of ﬁgure 7.2d.
We extracted a source spectrum from this region with radii of the major and minor axes of 4.9
and 7.9 arcmin, respectively. We excluded a background spectrum from the pentagonal region inﬁgure 7.2. NXB spectra constructed with the xisnxbgen software were subtracted from both thesource and background data. After correcting the vignetting and accumulation area, the source
spectrum was made by subtracting the background spectrum. The spectra of the two FI sensors
were merged to improve the statistics. In ﬁgure 7.3, we show the merged FI spectra. We can see
several prominent lines of K-shell emission from He- and H-like ions of Mg, Si, and S from thesource spectrum. Because G359.1-0.5 is located near the Galactic center, the emission is heavily
contaminated by the Galactic center diﬀuse X-ray emission which is roughly expressed by the thinthermal plasma with kT ∼ 6 keV and the neutral Fe line at ∼6.4 keV.

VRXUFHEDFNJURXQG
VRXUFHEDFNJURXQG

Fig. 7.3.—: XIS FI spectra of G359.1-0.5. We show the NXB-subtracted spectrum in the source region(black), NXB-subtracted and vignetting and accumulation corrected background spectrum (light gray), andbackground-subtracted source spectrum (dark gray). Note that, in the energy band below 1 keV and above
5 keV, the source emission is heavily contaminated by Galactic center diﬀuse X-ray emission. The energiesof the prominent emission lines from speciﬁc elements are labeled.

7.3. ANALYSIS AND RESULTS FOR GLOBAL REGION

7.3.2

83

Spectral Analysis

CIE Model Fitting
We here focus on 1-5 keV band, because the spectrum is heavily absorbed by the interstellar
extinction below 1 keV and contaminated by Galactic center diﬀuse emission above 5 keV. We use
only the FI data because the calibration for the FI is currently better around ∼2 keV.
We ﬁrst ﬁtted the spectrum with one VAPEC model. The Si and S abundances (ZSi and ZS )were free parameters, while the Mg abundance (ZMg ) was ﬁxed to the solar value. Figure 7.4a shows
the ﬁtting result with kTe ∼0.8 keV. The model exhibited signiﬁcant discrepancy around 2.7 keV,and was rejected with a large χ2 /dof of 382/96. We then tried a two-VAPEC model, assuming
equal abundances between two components. Figure 7.4b shows the result with kTe ∼0.5 keV and∼2.6 keV. The large residual still remained with an unacceptable χ2 /dof of 224/94. Adding afurther plasma component did not improve the ﬁtting any more.
We checked the light curves of the source and background regions in the 2.5–3.0 keV band.
They showed almost constant ﬂuxes, hence no ﬂare-like event had occurred during the observations.
We detected the bump in all XIS sensors. Thus, the bump is a real structure.

D

F

E

G

Fig. 7.4.—: XIS FI spectrum of G359.1-0.5 in the 1–5 keV band. (a) One- and (b) Two-VAPEC modelﬁtting. The best-ﬁt model is shown by solid lines. (c) Same as (a), but with He-RRC and Lyα lines of Si
and S. (d) Same as (c), but the temperatures of the VAPEC and RRC components are linked.

CHAPTER 7. G359.1-0.5

84
Recombination Structures

There are no emission lines of abundant elements in the energy band of the bump. Furthermore,
the bump is too broad to be explained by a single emission line. Instead, the edge of the bump(∼2.4 keV) corresponds to the electron binding energy of H-like Si, suggesting that this bump islikely due to RRC. Assuming this is suggestive of an overionized state, we introduced Si He-RRC andLyα lines of Si and S in addition to the one-VAPEC model. A S He-RRC is also considered because
a marginal bump is seen around 3.5 keV in ﬁgure 7.4a. During the analysis in this chapter, we
ignore the contribution of the cascade lines because they are not signiﬁcant due to poor statics. We
assume that the RRC is expressed following equation 4.7. This formula gives a good approximation
when kTe  Eedge . In the case of G359.1-0.5, kTe is ∼0.4 keV, as we conﬁrm later, while Eedge ofSi and S are ∼2–3 keV. Thus, we can safely adopt this formula.
The ﬁtting result is given in ﬁgure 7.4c. With this model, χ2 /dof was greatly improved
to 115/91. The temperature of the CIE component (0.55–0.63 keV) is higher than that of therecombination component (0.27–0.37 keV). As we see in the analysis of IC 443, however, thetemperature of the CIE component becomes lower if we include the RRC of other lighter elements(such as Ne and Mg) or other ionization states (such as H-RRC). However, we cannot estimate
these components with the statics of G359.1-0.5. Instead, we linked the temperatures of the CIE
and RRC components. The resultant best-ﬁt parameters and models are given in table 7.2 andﬁgure 7.4d, respectively. With this model, χ2 /dof is 120/92, which is almost same with ﬁgure 7.4c.
Because this model is acceptable with null hypothesis probability of 0.026, we use the result oftable 7.2 in the following discussion.
Table 7.2—: Best-ﬁt spectral parameters of G359.1-0.5.

CIE (VAPEC)
Line

RRC

kTe [keV]ZSi [solar]
0.38 (0.28–0.41)b 5.8 (4.9–10.8)cCenter energy [keV]
Si Lyα2.006 (ﬁxed)
S Lyα
2.623 (ﬁxed)
Edge energy [keV]
Si He-RRC
2.439 (ﬁxed)
S He-RRC
3.225 (ﬁxed)

ZS [solar]
11.0 (7.7–22.4)cLine width [keV]
0 (ﬁxed)0 (ﬁxed)kTe [keV]
0.38 (0.28–0.41)b0.38 (0.28–0.41)b
Normalizationa
2.17 (1.78–2.58)
Normalizationd
7.94 (5.19–10.4)
1.36 (0.00–2.32)
Normalizationd
20.2 (18.0–24.7)
2.15 (0.29–3.11)

The uncertainties in the parentheses are the 90% conﬁdence range. a The unit is 10−12 ne nH dV /(4πD2 ) (cm−5 ),where ne , nH , V , and D are the electron and hydrogen densities (cm−3 ), emitting volume (cm3 ), and distance tothe source (cm), respectively. b We assumed the same kTe . c These values should be modiﬁed in the overionizationcase (section 7.3.4). d The unit is 10−5 cm−2 s−1 .

7.3. ANALYSIS AND RESULTS FOR GLOBAL REGION

7.3.3

85

Electron and Ionization Temperatures

We successfully detected He-RRC of Si and S from G359.1-0.5, for the ﬁrst time. Due to thelimited energy resolution and photon statistics of ASCA, the Si He-RRC was misunderstood as Hlike S Kα line and the existence of high-temperature (∼4 keV) S-rich plasma was claimed (Bambaet al. 2000). Here, we mention the plasma temperature.
The electron temperature determined by the bremsstrahlung continuum shape and the RRC
width is ∼0.4 keV. Using this value, we derive the ionization temperature. From the best-ﬁt modelin table 7.2, the ﬂux ratios of Lyα/Heα and He-RRC/Heα of Si are given as 0.15 (0.10–0.20) and
0.38 (0.34–0.46), while those of S are 0.13 (≤0.23) and 0.20 (0.03–0.30), respectively. In ﬁgure 7.5,
we compare these values with the modeled emissivity ratios at kTe = 0.4 keV (Masai 1994) . We
obtain kTz = 0.79 (0.68–0.86) keV, 1.13 (≥1.00) keV for Si, and 1.3 (≤1.6) keV and 0.77 (0.51–
0.92) keV for S using the above ratios, respectively. All these ionization temperatures are consistent
with ∼0.8 keV 1 . The result that the kTz ∼0.8 (0.7-0.9) keV is signiﬁcantly higher than kTe ∼0.4
(0.3-0.4) keV indicates the plasma is in a highly overionized state.

D

E

Fig. 7.5.—: Predicted emissivity ratios of (a) Lyα/Heα and (b) He-RRC/Heα for Si and S as a functionof the ionization temperature for electron temperature of 0.4 keV (Masai 1994). The horizontal dashed lines
represent 90% errors of the observed values.

1

Only Si He-RRC/Heα indicates high kTz (≥1.0 keV). This would be due to the signiﬁcant contribution of Si
H-RRC. In the case of kTz = 0.8 keV for the Si plasma, an ion fraction of fully-ionized ions is about 1/6 of that of
H-like ones (ﬁgure 4.6). Because the recombination cross section for the fully-ionized ions into the ground state isabout twice as large as that for H-like ions, we estimate the reasonable ﬂux portion of H-RRC is about ∼1/3 of the
He-RRC. If we assume that the real Si He-RRC ﬂux is ∼2/3 of that in ﬁgure 7.5b, it is consistent with kTz ∼0.8 keV.

CHAPTER 7. G359.1-0.5

86

7.3.4

Abundance Correction for Overionization

We make possible modiﬁcations for the abundance in the case of overionization using anavailable APEC code. In the kT range of 0.4–0.8 keV, most Heα emission originates from theHe-like ions. Hence, the emissivities of the He-like ions for Si and S are modiﬁed by multiplying
the ion-fraction ratio between kTz = 0.8 keV and 0.4 keV. The Heα intensity is proportional
to Z × (kTe , kTz ), where Z and (kTe , kTz ) are the abundance of the element (solar) and thetotal emissivity for the Heα for kTe and kTz (cm3 s−1 ), respectively. Multiplying by (0.4 keV,
0.4 keV)/ (0.4 keV, 0.8 keV), we obtain the real abundances in the overionized state as ZSi ∼5.8×(0.94/0.60)∼9.1 solar and ZS ∼ 11.0 × (0.85/0.86) ∼11.0 solar. We should note that thesevalues become lower if we consider the cascade lines which occupy signiﬁcant fraction of the Heα
emission2 .

7.3.5

Volume Emission Measure

The VEM of the VAPEC component is derived from table 7.2 as ∼2.2 ×1012 cm−5 . On
the other hand, the VEM of the recombination component is derived from equation 4.7. The
total radiative recombination rate of He-like Si at kTe = 0.4 keV is ∼2.1 ×10−12 cm3 s−1 . The
recombination rate into the ground state of the H-like ions is derived from equation 4.10 for the
He-like Si as σ1 /(σ1 + σ2 + · · ·) ∼ 0.45. The value of nSi /nH is calculated using ZSi in the previous
subsection. Using the observed He-RRC ﬂux in table 7.2 and κH−like (0.8 keV) ∼ 0.33, we obtain
VEMRRC(Si) ∼ 2.0 ×1012 cm−5 . In the same way for S, using α1 ∼3.1 ×10−12 cm3 s−1 , σ1 /(σ1 + σ2 +· · ·) ∼ 0.44, and κH−like (0.8 keV) ∼ 0.09, we obtain VEMRRC(S) ∼ 1.0 ×1012 cm−5 . These consistent
results support the same origin of the bremsstrahlung and recombining components.

7.4

Summary• The Suzaku spectrum exhibits unusual bump-like residuals around 2.7 keV when ﬁtted by
the CIE plasma model. This bump is well explained by the RRC of Si. This is the third
examples of the recombination-dominant SNR following after W49B and IC 443.
• The electron temperature derived from the bremsstrahlung and RRC shapes is ∼0.4 keV.
On the other hand, the ionization temperatures derived from the observed intensity ratios of
RRC/Heα and Lyα/Heα for Si and S are consistent with ∼0.8 keV. This indicated the theplasma of G359.1-0.5 is in a highly overionized state.
• VEMs independently determined from the ﬂuxes of the bremsstrahlung and recombination
components are consistent with each other. This supports the common origin of these components.

2

If we estimate from the observed He-RRC ﬂux and equation 4.10, about half of the total Heα emission for Si andS originate from the cascade process.

Chapter 8

Discussion on
Recombination-dominant SNRs
In the previous chapters, we have analyzed X-ray data of three SNRs (W49B, IC 443, andG359.1-0.5) obtained with Suzaku. The superior sensitivity and the energy resolution of Suzaku
XIS enable us to discover the brand-new spectral structure; the strong RRC in SNRs. We havefound the Fe He-RRC from W49B, Si and S H-RRC from IC 443, and Si and S He-RRC from
G359.1-0.5. In particular, we have further discovered cascade line structures from W49B.
All these SNRs exhibit the equal electron temperatures and volume emission measures forthe bremsstrahlung and RRC components. This fact robustly indicates that the X-ray emission isdue to a single overionized plasma, or a recombining plasma. Thus, we ﬁnd a new aspect of the
SNR evolution, presence of a strong recombining plasma, for the ﬁrst time. In this chapter, weinvestigate the characteristics of the recombination-dominant (hereafter RD) SNRs by comparing
with a well-known ionization-dominant (ID) SNR and discuss the origin of the strong recombining
plasma.

8.1

Spectral Features

In order to clarify the characteristics of the RD SNR 1 , we compare them with a typical ID
SNR (Tycho’s SNR 2 , hereafter Tycho for short). The Suzaku observation of Tycho is summarized
in appendix A.

1

Here we focus on W49B, which exhibit the most prominent structure of recombination X-rays.
In AD 1572, the astronomer Tycho Brahe observed a bright supernova in detail. It is now ﬁrmly classiﬁed as a
Type Ia SNR based on the optical spectrum from a scattered-light echo (Krause et al. 2008).
2

87

88

CHAPTER 8. DISCUSSION ON RECOMBINATION-DOMINANT SNRS

Fig. 8.1.—: Chandra three-color composite image of Tycho (Warren et al. 2005). (Blue) 4.1-6.1 keV band.
Mainly originates from a non-thermal component. (Red) 0.95-1.26 keV. Mainly emitted by Fe ions. (Green)
1.63-2.26 keV. Mainly emitted by Si ions.

Fe He-RRC Edge

He-like KαH-like Kα
He-like Kβ

Tychoo

S

{

Ar Ca

{

Si

{

{
{

W49B
B
Fe

Fig. 8.2.—: Background subtracted Suzaku XIS FI spectra of Tycho (black) and W49B (red) taken fromthe whole remnant region. The energies of the prominent Kα emission lines are labeled. We show He-likeKα lines in green, He-like Kβ in purple, and H-like Kα in blue. We also show the edge energy of the Fe
He-RRC.

8.1. SPECTRAL FEATURES

89

Tycho is one of the youngest (∼440 years) and X-ray brightest SNRs in our Galaxy. As
is seen in ﬁgure 8.1, it exhibits nearly spherically-symmetric shape. The X-ray spectrum shows
intense emission lines which were studied in detail with the previous satellites. Using Tenma3 data,
Tsunemi et al. (1986) showed that Kα center energies of Si, S, Ar, Ca, and Fe are consistent withthose of underionized states, which is revealed for the ﬁrst time in SNRs. Thus, this source is atypical ID SNR in a strongly underionized state.
In ﬁgure 8.2, we show the Suzaku spectra of Tycho and W49B. Both exhibit intense line
emission, but there are signiﬁcant diﬀerences. For the quantitative comparison of the spectral
structures, we construct a semi-phenomenological model, which is composed of a bremsstrahlung,
Gaussian function lines, and an RRC. We focus on the W49B and Tycho spectra in the energy
band of 5–12 keV, where emission structures mainly originate from Fe. We simply assume that therelevant RRC is a Fe He-RRC.
For the emission lines of Tycho, we treat Fe Kα, Kβ, Cr Kα, and Ni Kα. Line centroids arefree parameters 4 . We also include Fe Lyα with ﬁxed line center in order to derive the ﬂux upper
limit. Because heavy elements in Tycho are Doppler-shifted (Furuzawa et al. 2009), we treat theline widths as free parameters. All these widths are linked each other for simplicity.
As for the emission lines of W49B, Fe Heα, Heβ, Heγ-∞, Lyα, and Kα lines of Cr, Mn, andNi are included. We ﬁx the center energies of Cr Kα, Mn Kα, and Fe Lyα to those in table 5.1,
and Heβ to the expected value for kTz = 2.7 keV 5 . The center energies of the other lines are freeparameters. Since Fe Heα and Heγ-∞ include many emission lines from diﬀerent ionization statesand energy levels, we treat the widths of these lines as free parameters. Widths of the other lines
are ﬁxed to 0.
Using this semi-phenomenological model, we can independently determine the electron temperatures with three independent methods; the bremsstrahlung shape, line ﬂux ratios, and thewidth of RRC 6 . We, however, link the temperatures of bremsstrahlung and RRC components for
Tycho because we see no signiﬁcant excess of the RRC structure.
We ﬁrst ﬁt the spectra with a model of bremsstrahlung plus Gaussian lines. Figure 8.3a shows
the best-ﬁt results. We can see clear saw-edged bump structure around 9 keV for W49B with
χ2 /dof = 850/133, while no sign for Tycho with χ2 /dof = 144/58 7 .
We then add an Fe He-RRC. The best-ﬁt results are summarized in ﬁgure 8.3b and table 8.1.
For W49B, χ2 /dof is greatly improved to 270/131, while χ2 /dof = 89/57 for Tycho is not so large
but signiﬁcant. The signiﬁcant improvement of the Tycho spectrum would be due to artiﬁcial
compensation of the non-thermal contribution by the RRC.

3This is the Japanese second X-ray satellite, launched in 1983.
We only ﬁx the centroid of Ni Kα to the predicted value of neutral line because it is weak.
5
This ionization temperature is derived in chapter 5.
6
See section 4.3 for details.
7
We can see a rising slope above 7.5 keV in the lower panel (residual) of Tycho (right) in ﬁgure 8.3a. This would
be due to the non-thermal contribution as we brieﬂy discuss later.
4

CHAPTER 8. DISCUSSION ON RECOMBINATION-DOMINANT SNRS

90

In the spectrum of W49B, a large residual around Fe Heα are found. This is partially due tothe response uncertainty as we noted in chapter 5. However, the main reason is due to the single
Gaussian line approximation for the composite of many lines in diﬀerent ionization states. In fact,
if we ignore 6.4–6.9 keV, we obtain an acceptable χ2 /dof = 113/98.
The electron temperatures of W49B derived in the three diﬀerent methods are; ∼1.7 (1.6–1.7)
keV from the bremsstrahlung shape, ∼1.8 (1.6–1.9) keV from the ﬂux ratios of Fe Heβ/Heα 8 , and∼1.3 (1.2–1.6) keV from the RRC width. These results are consistent with each other, and hence
no additional component such as non-thermal emission is required. In the previous APEC model
analysis (chapter 5), we implicitly assumed that the temperatures of the bremsstrahlung and thatfrom the Kβ/Kα ﬂux ratio are the same. Here we conﬁrm the validity of this assumption.
The Fe Kβ/Kα ﬂux ratio of Tycho can be converted to kTe ∼1.8 (1.6–2.0) keV 9 . On theother hand, the bremsstrahlung shape indicates that kTe is ∼4.5 (4.4-4.5) keV. We regard that
this inconsistency is possibly due to the non-thermal contribution. In fact, Tamagawa et al. (2009)
reported the existence of non-thermal emission of Tycho using the wideband spectra of the Suzaku
XIS and HXD. They detected X-ray emission up to 30 keV with a power-law index of 2.69, andconcluded that the emission comes from synchrotron radiation. We should emphasize that oursemi-phenomenological model can in principle separate the non-thermal components, by ﬁxing
the electron temperature of the bremsstrahlung to that from the Fe Kβ/Kα line intensity ratio.
However, further details of the non-thermal component in Tycho is beyond the scope of this thesis.

8

We used ﬁgure 4.10.
We derive this value using ﬁgure 4.10. We also conﬁrm this from the same semi-phenomenological model ﬁttingin the 1.7–5 keV band. The electron temperatures derived from the Kβ/Kα ﬂux ratios of Si and S indicate ∼2.4 keVand ∼1.9 keV, respectively. All the Kβ/Kα ratios indicate the electron temperature is around 2 keV.
9

8.1. SPECTRAL FEATURES

91

W49B

Tycho
(a)

(a)

(b)

(b)

Fig. 8.3.—: Semi-phenomenological model ﬁttings of W49B (left) and Tycho (right) spectra. (a) XIS
FI spectrum ﬁtted with a bremsstrahlung continuum and lines. We can see a large bump above 9 keV in
W49B. A rising slope above 7.5 keV in the lower panel (residual) of Tycho is likely due to the non-thermal
contribution (see text). (b) Same spectrum as (a), but with the radiative recombination continuum of Fe.

CHAPTER 8. DISCUSSION ON RECOMBINATION-DOMINANT SNRS

92

Table 8.1—: Best-ﬁt spectral parameters in ﬁgure 8.3.

W49B
bremsstrahlung
RRC

Fe He-RRC

lines
Cr Kα
Mn Kα
Fe HeαFe LyαNi KαFe HeβFe Heγ-∞Edge energy [keV]
8.830 (ﬁxed)
Center energy [keV]
5.655 (ﬁxedc)
6.162 (ﬁxedc)
6.667 (6.666-6.667)
6.961 (ﬁxedc)
7.722 (7.712-7.733)
7.845 (ﬁxedd)
8.347 (8.323-8.374)

kTe [keV]
1.67 (1.66-1.68)
kTe [keV]
1.32 (1.19-1.60)
Width [keV]
0 (ﬁxed)0 (ﬁxed)0.36 (0.35-0.38)
0 (ﬁxed)0 (ﬁxed)0 (ﬁxed)0.22 (0.19-0.25)

Normalizationa
2.98 (2.93-3.01)
Normalizationb
11.6 (10.8-12.6)
Normalizationb
2.89 (2.61-3.15)
0.99 (0.77-1.20)
91.2 (90.4-92.0)
1.37 (1.17-1.56)
2.29 (2.05-2.52)
3.92 (3.60-4.18)
4.14 (3.75-4.51)

Tycho
bremsstrahlung
RRC

Fe He-RRC

lines
Cr Kα
Fe Kα
Fe LyαFe KβNi KαEdge energy [keV]
8.830 (ﬁxed)
Center energy [keV]
5.471 (5.462-5.483)
6.432 (6.431-6.432)
6.966 (ﬁxedf )7.106 (7.096-7.123)
7.470 (ﬁxedg)

kTe [keV]
Normalizationa
4.48 (4.45-4.52)e1.00 (1.00-1.01)
kTe [keV]
Normalizationb
4.48 (4.45-4.52)e1.20 (0.99-1.40)
Width [keV]
Normalizationh
0.060 (0.058-0.061)h 2.79 (2.56-3.11)
0.060 (0.058-0.061)h 66.8 (66.3-67.2)
0.060 (0.058-0.061)h 0.09 (0.00-0.36)
0.060 (0.058-0.061)h 2.93 (2.66-3.20)
0.060 (0.058-0.061)h 0.82 (0.55-1.09)



Note. — The uncertainties in the parentheses are the 90% conﬁdence range. a The unit is 10−13 ne nH dV /(4πD2 )(cm−5 ), where ne , nH , V , and D are the electron and hydrogen densities (cm−3 ), emitting volume (cm3 ), and distanceto the source (cm), respectively. b The unit is 10−5 cm−2 s−1 . c Fixed to the value in table 5.2. d Fixed to the ideal
value at kTz = 2.7 keV. e Linked each other. f Fixed to the ideal value. f Fixed to the ideal value of neutral Ni.
hLinked each other.

8.1. SPECTRAL FEATURES

93

The semi-phenomenological model ﬁtting have clearly revealed the spectral features of the
newfound RD SNR as below.
1. Strong recombination structures (RRC and cascade lines).
2. Strong Lyα line.
3. High line center energy for Heα. 10 .
1) Remarkable fact is that the electron temperatures of W49B and Tycho are nearly the same(∼1.6-1.8 keV) 11 . Nevertheless, there are signiﬁcant spectral diﬀerences; W49B has a very strong
recombination structure above 8 keV, while no sign in Tycho. Thus, the two types of SNRs, the RD
and ID SNRs, are clearly distinguished from the presence of the strong recombination structure.
2) We can see intense Lyα emission in the W49B spectrum from Si, S, Ar, Ca, and even Fe(ﬁgure 8.2). On the contrary, we can hardly see them in the Tycho spectrum. The Fe Lyα/Heα
ﬂux ratio is 0.015 (0.013-0.017) for W49B, while only an upper limit (0.0054) is derived for Tycho.
The high Lyα/Heα ﬂux ratio indicates high ionization temperature. In fact, kTz derived from thisratio is ∼2.6 (2.4-2.7) keV for W49B, while ≤0.85 keV for Tycho.

Table 8.2—: Observed Kα line center energies (eV).

Cr
Mn
Fe
Ni

W49B
5655 (5646-5663)
6162 (6142-6183)
6667 (6666-6667)
7722 (7712-7733)

IC 4435575 (5438-5697)
not detected
6676 (6656–6696)
not detectedTycho
5471 (5462-5483)
5950 (5900-6000) *6432 (6431-6432)
7470 (ﬁxed to the neutral value)* From Tamagawa et al. (2009).

10

Note that we here use the word ”Heα” as the mixture of the lines from He-, Li-, Be-like and lower ionized ions.
The Fe Kβ/Kα ﬂux ratio is 0.043 (0.040-0.046) for W49B, while 0.044 (0.040-0.048) for Tycho. This consistent
result strongly suggests the electron temperature for both SNRs are nearly the same.
11

94

CHAPTER 8. DISCUSSION ON RECOMBINATION-DOMINANT SNRS

3) In table 8.2, we summarize the observed Heα 12 center energies of Cr, Mn, Fe, and Ni for
W49B, IC 443, and Tycho 13 . All the centroids of Tycho is signiﬁcantly lower than those of IC 443
and W49B. The high line centroids for the RD SNRs indicate the highly-ionized state of the plasma.
Note that we ﬁnd a marginal Kα sign of Cr and robust one of Fe from IC 443 for the ﬁrst time(ﬁgure 8.4). We ﬁt the spectrum of IC 443 in the energy band of 5.0-7.5 keV with a power-law and
two Gaussian functions. The best-ﬁt ﬂuxes are 1.7 (0.0–4.1) ×10−6 photons cm−2 s−1 for Cr and1.46 (1.17–1.77) ×10−5 photons cm−2 s−1 for Fe. Although the Cr emission is marginal, χ2 /dof is
improved from 36.4/74 to 34.8/72 by introducing the Cr Kα line.

IC 443

Cr
Kα

+ FI
+ BI

Fe Kα

Fig. 8.4.—: NXB subtracted XIS FI and BI spectra of IC 443 extracted from the circular region inﬁgure 6.3. We ﬁt the spectra with a power-law and two Gaussians which represent Cr and Fe Kα lines.

12

As we see later, the line centroids of the Tycho spectrum clearly indicate the underionized state than those from
He-like ions. However, here we describe the Kα lines as Heα for convenience.
13
As for G359.1-0.5, these lines are not detected because the plasma temperature is relatively low and the source
emission is highly contaminated by the Galactic center X-ray emission.

8.1. SPECTRAL FEATURES

95

:%

+HOLNH/LOLNH%HOLNH

7\FKR%OLNH,&
1HXWUDO

Fig. 8.5.—: Weighted line center energies of Fe and Cr Kα lines for the observed recombinationdominant SNRs (W49B and IC 443) and ionization-dominant SNR (Tycho). We also show the theoretical center energies of in diﬀerent ionization states with triangles. (Mewe et al. 1985, Yang et al. 2009,http://www.camdb.ac.cn/e/spectra/spectra search.asp)

In ﬁgure 8.5, we plot the Kα centroids of Fe and Cr for these SNRs. We also show thetheoretical center energies of these lines in diﬀerent ionization states. In the case of W49B andIC 443, the line centroids indicate that the Fe ions are ionized into nearly He-like. In the caseof Tycho, on the contrary, centroids are signiﬁcantly lower than those from B-like. Thus, highcentroids of these Heα lines, which indicate the presence of highly-ionized atoms, are characteristicof the RD plasma.

CHAPTER 8. DISCUSSION ON RECOMBINATION-DOMINANT SNRS

96

6.9
25

IXOO\LRQL]HG+OLNH+HOLNH

6.8

charge state

Charge state of Fe

20

IC 443

W49B
6.7

15
6.6
10

Tycho

K-line energy

6.5

5
6.4
0
0.05

0.1

0.5

1

Intensity-weighted center energy (keV)30

5Temperature (keV)
Fig. 8.6.—: Average ionization state and K-line energy of Fe as a function of CIE plasma temperature. The
K-line enrgy is intensity-weighted for the lines including satellites. We also show the observed Fe Kα center
energy of W49B, IC 443 and Tycho. Here we derive kTz from the observed intensity-weighted centroid.
Using this temperature, we estimate that major ionization states of Fe are Si- to F-like for Tycho, while Lito He-like for IC 443 and W49B. We calculated the Kα centroid for W49B by taking the weighted center
between Heα and Lyα. This ﬁgure is taken from Masai et al. (2002).

Although the line center is a function of both kTe and kTz , it mainly reﬂects the ion fraction(i.e. kTz ). Hence we can roughly estimate kTz from the intensity-weighted center energy assuming
the CIE state. The result is shown in ﬁgure 8.6. We assume the systematic uncertainty of energy
scale is 0.1 % following Ozawa et al. (2009b). From the observed center energy, we can derivekTz of Fe ions as 0.16–0.40 keV for Tycho, 1.5–3.8 keV for IC 443, and 1.7–2.9 keV for W49B. Thisvalue is consistent with that derived from RRC/Heα or Lyα/Heα ﬂux ratios for W49B.

8.2. CHARACTERISTICS

8.2

97

Characteristics

We summarize the electron temperature, age, distance, radius, and X-ray morphology of thethree RD SNRs (W49B, IC 443 and G359.1-0.5) and ID SNR (Tycho) in table 8.3. It should be
noted that all the RD SNRs show the awkward and center-ﬁlled morphology in X-rays (mixedmorphology SNRs). We should also note the high-energy gamma-rays are detected from W49B
(GeV gamma-rays by Fermi) and IC 443 (GeV by EGRET, AGILE, & Fermi and TeV by MAGIC
and VERITAS: Esposite et al. 1996; Tavani et al. 2009; Albert et al. 2007; Acciali et al. 2009).
Also, TeV gamma-rays are detected by HESS near G359.1-0.5 (Aharonian et al. 2006). We alsomention the OH maser is detected from G359.1-0.5 (Yusef-Zadeh et al. 1996) and IC 443 (Hewitt
et al. 2006). Yusef-Zadeh et al. (2003) showed there is a strong association between mixedmorphology SNRs and maser-emitting SNRs. Hewitt et al. (2009) show the correlation between
maser-emitting SNRs and a class of GeV- to TeV-energy gamma-ray sources. There might berelations between the mixed-morphology, maser-emitting, high-energy accelerator, and RD SNRs.
Table 8.3—: Characteristics of the observed SNRs.

SNR
W49B
IC 443
G359.1-0.5
Tycho

kTe a (keV)
1.5
0.7
0.4
1.8

age (year) D b (kpc)
not known
11.4
1.5
3000d, 4000e
not known
8.5
∼440
2.8

R c (pc)
8.3
11
20
3.3

X-ray morphology
elongated, center-ﬁlled
elongated, center-ﬁlled
elongated, center-ﬁlled
almost spherical, shell-likea

Derived from our analyses using the Suzaku data. b Distance to the source. References are Gwinn et al. (1992)
for W49B, Welsh & Sallmen et al. (2003) for IC 443, Uchida, Morris, & Yusef-Zadeh (1992) for G359.1-0.5,
and Ruiz-Lapuente (2004) for Tycho. The distance are under debate for some of these SNRs. For example, theuncertainty is 8–12 kpc for W49B and 2-4 kpc for Tycho. Note that the physical parameters in table 8.4 arefunction of the distance, hence there is signiﬁcant uncertainty on the resultant values. c Radius of the source. We
assumed the distance in this table, and major axes of 2.5, 25, 7.9, and 4.0 arcmin for W49B, IC 443, G359.1-0.5,
and Tycho, respectively. d Petre et al. (1988). e Troja et al. (2008).

Table 8.4—: Physical parameters of the observed SNRs.

SNR
W49B
IC 443
G359.1-0.5
Tycho** *dominant process
recombination
recombination
recombination
ionization

density (cm−3 )
1.9
1.0
0.23
2.3

thermal energy* (1050 ergs)
9.6
0.15**
1.6
0.52

mass* (M )
110
3.7**
68
5.0

√
√
*Note that if we consider the ﬁlling factor f , it is multiplied by 1/ f for density, and f for energy and mass. We
assume a sphere with radii of 2.5, 7.5, 4.0 arcmin for W49B, IC 443, and Tycho, respectively, and an ellipsoidal
shape with radii of 7.9, 4.9, and 4.9 arcmin for G359.1-0.5. **Note that our estimation is based on the result
from a part of the SNR. The energy and mass are expected to be dozens of times larger for the whole remnant.
***We use the result of NEI (thin thermal underionized plasma code) ﬁtting for Tycho.

98

CHAPTER 8. DISCUSSION ON RECOMBINATION-DOMINANT SNRS

We then derive the physical parameters of the plasma in table 8.4, assuming the distancein table 8.3. Because the observed ﬂux of the thermal component (VAPEC) is proportional to(1/4πD2 ) nH ne dV , we can derive the density of the plasma (n) assuming the hydrogen density(nH ) and electron density (ne ) are equal (nH ne n) in a hydrogen-dominant and uniform density
plasma. The total energy of the plasma is obtained using the equation; E = 3kT /2×(nH +ne )×V
3kT nV . The total mass of the plasma is estimated by the equation M = nV mH , where mH is thehydrogen mass. Note that we cannot cover the whole remnant for IC 443, hence the total energy
and mass would be dozens of times larger. We should note that the estimated mass is sub-solar for
Tycho, while several tens of solar for the RD SNRs. Because of the high metallicity for W49B (∼3
solar), IC 443 (∼2 solar), and G359.1-0.5 (<∼10 solar), we can safely assume most of the plasma
originate from the ejecta. Hence, the plausible progenitors of the RD SNRs are massive stars.
In ﬁgure 8.7, we show the electron and ionization temperatures for the observed SNRs. We also
show the ASCA results from Kawasaki et al. (2005). We should note that the electron temperaturesof W49B and IC 443 are signiﬁcantly lower than the ASCA results. This is because we take theRRC contribution into account.
(Fe) (Ca)W49B

Overionized(recombining)

Suzaku(Our results)

ASCA
IC 443 (S)G359.1-0.5(Si)(Kawasaki et al. 2005)

Underionized(ionizing)
Tycho (Fe)Fig. 8.7.—: Electron and ionization temperatures of the SNRs observed with Suzaku (black circle). Gray
points indicate ASCA results. We derived kTe from the bremsstrahlung shape for W49B, IC 443, andG359.1-0.5, and from the Fe Kβ/Kα ﬂux ratio for Tycho. We obtained kTz from the RRC/Heα ratio for
W49B and IC 443, Lyα/Heα ratio for G359.1-0.5, and Kα line center for Tycho. Parentheses indicate theelement from which we derive kTz . We should note that kTe of W49B and IC 443 become lower by taking
RRC into account (arrows).

8.3. FORMATION MECHANISMS

8.3

99

Formation Mechanisms

We have clearly found the overionized or recombining plasmas in SNRs but the formation
mechanism is an open question. We here mention the evolution of the ionization degree in SNRs,
and discuss the possible formation mechanisms of the anomalous plasma state.

Time Evolution of Ionization Degree in Normal SNRs
We ﬁrst mention normal evolution of ionization degree in SNRs. Soon after the SN explosion,
elements in the progenitor are heated and almost fully ionized due to a passage of the blast wave
going through the progenitor. Expansion after the shock passage cools down the plasma. The
density is high enough that highly ionized ions soon recombine with electrons. Assuming the typicaldensity in this phase (∼ 108 cm−3 ) and typical recombination coeﬃcient (∼ 10−12 –10−11 cm3 s−1 ),recombination timescale is estimated to be ∼103 − 104 seconds. In this way, highly ionized ionsrecombine within one day.
After a few days of the SN explosion, the plasma is again heated by the radiative decay ofIn this phase, the temperature of the center of the
4progenitor is ∼several × 10 K and atoms are in the several-ionized state. On the other hand, thetemperature outside the photosphere is ∼several × 103 K, and atoms are in the ﬁrst- or secondionized state.
56 Ni which is synthesized in the explosion 14 .

After ∼100 days, the whole SN region becomes optically thin. Almost all the ejecta atoms
are in the ﬁrst- or second-ionized state in this phase. This ionization state of the ejecta lasts untila reverse shock passes the plasma. As the ejecta expands, the reverse shock is formed by theinteraction with the ambient matter and it propagates toward the ejecta core from the contact
interface. Then, thermal energy is transferred from ions to electrons through Coulomb collisions
in the post-shock region. Heated electrons gradually ionize atoms, as we mentioned in chapter 1.
For a ﬁducial density (1 cm−3 ) of SNR plasmas, electrons are heated up to ∼ keV in ∼102 years(Masai 1994), while it takes ∼104 years (or nt ∼ 1012 cm−3 s) for collisional ionization to reachits equilibrium state 15 . Therefore, X-ray bright young SNRs are thought to exhibit under-ionized(kTe > kTz ) or ionizing states.

Ni decays to 56 Co in ∼6 days (half life), and 56 Co to 56 Fe in ∼77 days (half life).
Equipartition timescale is of the same order as ionization timescale.

14 56
15

100

CHAPTER 8. DISCUSSION ON RECOMBINATION-DOMINANT SNRSFormation Processes of Recombining Plasma
As we see above, the overionized or recombining plasma state cannot be expected in theordinary evolution of SNRs. Hence we need a particular process to explain the coexistence of thehighly-ionized atoms and low-temperature electrons as below.
1. Increase the ionization degree of atoms.
2. Decrease the thermal energy of electrons.
1) One possible scenario to increase the ionization degree is photoionization. Since no strong
radiation source is found near the observed RD SNRs, very recent photoionization is unlikely.
Instead, strong photoionization at the moment of the SN explosion may be possible.
X-ray afterglow is an example, which occurs subsequently to a gamma-ray burst (GRB). It is
known that some of GRBs are associated with the SN explosion (Type Ic) of massive (∼40 M )stars, called hypernovae; examples are GRB980425, GRB030329, and GRB031203. Perna & Loeb
(1998) suggest that the X-ray and UV components of the afterglow of GRB create an ionized bubbleof radius ∼100 pc in the surroundings of the density of 1 cm−3 . It is possible that X-ray afterglowof these hypernovae irradiates and ionizes surrounding matter in the very initial phase of the SNR
evolution.
Similar example is an X-ray ﬂash, which is less energetic than a GRB. It emits mainly X-rays
instead of gamma-rays. Some of X-ray ﬂashes are thought to be associated with the SN explosionof ∼20 M stars; an example is XRF060218.
Another example of photoionization is break-out X-rays from normal SNe. Colgate (1968)predicts a soft X-ray ﬂash to occur when the shock wave breaks out of the progenitor. Soderberg
et al. (2008) discovered the luminous X-ray outburst from SN 2008D, and interpret it as the shock
break-out. Candidates to emit break-out X-rays are type Ib and Ic SNe 16 .
These mechanisms would make atoms highly-ionized in the initial phase of the SNR evolution.
If the highly-ionized atoms still survive, we observe the overionized or recombining plasmas. We
discuss this scenario later from the view points of the timescale and metallicity.

16

Type II SNe are expected to emit ultraviolet photons rather than X-rays because of large radii of their progenitors.
Note that typical temperature (T ) is expected to be T ∝ R−3/4 E 1/4 , where R is a radius of the progenitor and E isthe explosion energy.

8.3. FORMATION MECHANISMS

1012) One mechanism to decrease the electron temperature is thermal conduction. This possibility
is discussed by Kawasaki et al. (2002 Ph.D.thesis) for IC 443. They showed that IC 443 is composedof a central hot (kTe ∼ 1.0 keV) region and a surrounding cool (kTe ∼ 0.2 keV) shell, and interpreted
that the hot interior have cooled via thermal conduction to the cool exterior. Under the several
reasonable boundary conditions 17 , they estimated the conduction timescale (tcond ) in which thetemperatures of the hot and cold plasma components to be comparable to each other. They
approximately estimate the age of the plasma to be (0.3–1.0) ×104 years, which is consistent with
the SNR age (∼4000 years : Troja et al. 2008). Hence they claim that the timescale of thermal
conduction is shorter than that of recombination, and thus overionized plasma is formed in IC 443.
Based on our results, however, the conduction timescale is estimated to be


tcond

1.5 × 104ne
1.0 cm−3kTe
0.7 keV

−5/2

yr ,

(8.1)

following equation 5 of Kawasaki et al. (2002). Assuming both kTz and kTe are e-folding with theircharacteristic timescales 18 , the age of the plasma (tage ) is estimated fromtexp(− tage
)
kTz
rec
∼
t
age
kTe
exp(− t)

1.1 keV.
0.7 keV

(8.2)

cond

Here, trec is the characteristic timescale to reach ionization equilibrium. In the case of S at atemperature around ∼1 keV, trec ∼ 3 × 104 (n/1 cm−3 ) years (equation 3 of Kawasaki et al. 2002;
Masai 1994). Then we obtain tage ∼ 1.4 × 104 years. Because it is far longer than the SNR age,
this scenario cannot work on our new results. Thus, it is diﬃcult to explain the overionized orrecombining plasma with the thermal conduction process.

They estimate the scale length of the temperature gradient is ∼(1–2)×1019 cm, which corresponds to the distance
between the hot-cool plasma boundary and the edge of the SNR.
18
This is the same assumption with Kawasaki et al. (2002).
17

CHAPTER 8. DISCUSSION ON RECOMBINATION-DOMINANT SNRS

102ᶆᶃᵿᶒᶃᶂᴾᵿᶌᶂ
ᶆᶇᶅᶆᶊᶗᴾᶇᶍᶌᶇᶘᶃᶂᴾᵡᵱᵫ
ᶓᶌᶆᶃᵿᶒᶃᶂᴾᵡᵱᵫ
ᶃᶈᶃᶁᶒᵿ
ᵱᵬຜẟ଩ԗཋឋ
ᶂᶃᶌᶑᶃᴾᵡᵱᵫᘔએඬᶐᵿᶐᶃᶄᶇᶃᶂᴾᵧᵱᵫ

ᶀᶊᵿᶑᶒᴾᶕᵿᶔᶃᶐᵿᶎᶇᶂᴾᵿᶂᶇᵿᶀᵿᶒᶇᶁᴾᶃᶖᶎᵿᶌᶑᶇᶍᶌ

Fig. 8.8.—: Schematic illustration of the rarefaction scenario of Itoh & Masai (1989). When supernova
explosion occurs in the dense circumstellar matter, the matter is heated and highly ionized by the blast
wave. Subsequent rapid adiabatic expansion cools the shocked gas quickly.

1) & 2) The other possibility is the combination of 1) and 2). As a candidate of the formation
mechanisms of recombining plasma, we may propose the rarefaction process which is discussed by
Itoh & Masai (1989). Figure 8.8 shows a schematic illustration of this process. In this scenario, a
SN19 explodes in the dense circumstellar matter (CSM) which is formed by the progenitor’s stellar
wind in its supergiant phase. Because of the presence of CSM, the gas is shock-heated to hightemperature and rapidly ionized in the early phase of the SNR evolution. Subsequent outbreak of
the shock wave through the dense CSM into a low-density interstellar medium (ISM) caused drastic
adiabatic expansion of the shocked gas. If the CSM is much denser than the ambient ISM, thetimescale of adiabatic cooling is much shorter than that of recombination, and hence the remnant
becomes overionized in this stage. The recombination process dominates radiation, and the X-ray
spectrum is characterized by strong recombination lines and continua.
Itoh & Masai (1989) investigate the dynamical evolution and non-equilibrium X-ray emissionof the SNR using a spherically-symmetric hydrodynamic code. The time evolution of physical quantities is shown in ﬁgure 8.9 with parameters taken in the calculations. In this example (density of
the CSM is ∼ 600 cm−3 ), recombination timescale is ∼ 101 –102 years for typical recombination coeﬃcients (∼ 10−12 –10−11 cm3 s−1 ). Thus, the highly ionized state could remain for several decades
until the blast wave breaks out of the dense CSM. We can see the X-ray luminosity and electron
temperature decrease drastically by rarefaction due to rapid adiabatic expansion, as seen in ﬁgure 8.9. High Fe Kα center energy of W49B and IC 443 can be explained with this scenario. We
also show the time evolution of the shocked-matter radial proﬁle and X-ray spectrum in ﬁgure 8.10.
We can see a strong RRC of Fe at 200 years after the explosion. Thus, rarefaction is a possible
process to produce the recombining plasma.
In the following, we discuss the origin of the observed recombining plasma from the view pointsof timescale, metallicity, spatial distribution, and surrounding environment.
19Type II SNe are assumed, because their progenitors are expected to blow slow winds with high mass-loss rates
in their red supergiant phase.

8.3. FORMATION MECHANISMS

103

D

0RGHO)EF

EODVWVKRFNIURQWVWDUWRIUDUHIDFWLRQ

0RGHO(EODVWVKRFNIURQW

RXWHUVXUIDFHRI&60

VXUIDFHRIHMHFWD

H

G,&
:%

7\FKRFig. 8.9.—: Time evolution of physical quantities with (Model E) and without (Model F) CSM (Itoh &Masai 1989). (a) Model parameters. Msn , Esn , rsn , Mcs , rcs , and N0 indicate the initial supernova ejecta
mass, kinetic energy, ejecta radius, CSM mass, CSM radius, and ISM density. They assume the uniform
ejecta density up to the core radius rc , and the density ∝ r−n at radius rc < r < rsn with a constant n. (b)Time evolution of the radii, (c) X-ray luminosity, (d) electron temperatures, and (e) emissivity-weighted line
energy of Fe Kα. We can see that the luminosity and electron temperature rapidly decrease by rarefaction(c and d). Wiggles of the curves in (d) are due to ﬁnite zoning of the remnant in the calculations, and are
spurious. We show the observed line center energy in (e) for reference.

104

CHAPTER 8. DISCUSSION ON RECOMBINATION-DOMINANT SNRS0RGHO(LQLWLDOUHYHUVHVKRFND \HDU&60

HMHFWDE \HDU,60

F \HDU

)H+HȘ
)H55&

Fig. 8.10.—: Shocked-matter radial proﬁles (left) and spectra (right) at (a) 41 years, (b) 201 years, and (c)990 years after the explosion (Itoh & Masai 1989), where v,N,T , and Te are the velocity, density, electron-ion
averaged temperature, and electron temperature, respectively. The spectra is shown for total emission, line
emission, bremsstrahlung (B), recombination radiation (R), and two-photon emission (T). The total andline spectra are convolved with a Gaussian to a resolution of 0.1E 0.5 keV (FWHM), where E is the photon
energy in units of keV. The distance and the interstellar hydrogen column density to the SNR are taken to
be 3 kpc and 1022 cm−2 , respectively. We can see strong RRC and cascade structures of Fe in the spectraof (b) and (c), which is similar to the Suzaku W49B spectrum.

8.3. FORMATION MECHANISMS

105

Recombination Timescale
One question is that highly-ionized ions which were formed in the initial or early phase of
the SNR evolution can survive until today or not. The lifetimes of the fully-stripped ions (τ ) areroughly estimated from
1
,
(8.3)
τ
αne
where α is the total K-shell recombination rate coeﬃcient.
We obtain
4.8 ×
3
−3
−1
3
−3
−1
10 (ne /1.0 cm ) years for Si and 3.4 × 10 (ne /1.0 cm ) years for S, respectively. Notethat these values are somewhat shorter than the actual timescale for the plasma to reach the CIE
state. Nevertheless, the estimated lifetimes are comparable to the age of IC 443 (∼3×103 –4×103
years; Petre et al. 1998; Troja et al. 2008). The overionized plasma can, therefore, still survive atpresent. Thus, both photoionization and rarefaction are plausible processes to be responsible forthe formation of the overionized plasma.
We also discuss the timescale for the other SNRs. The lifetimes of highly-ionized ions are
3.5 × 104 (ne /0.23 cm−3 )−1 years for H-like Si, 6.5 × 104 (ne /0.23 cm−3 )−1 years for H-like S,
and 5.3 × 103 (ne /1.9 cm−3 )−1 years for H-like Fe. Because no information of the age is available
for W49B and G359.1-0.5, We cannot directly compare these lifetimes with the age of these SNRs.
However, from the size and temperature, we can reasonably assume W49B and G359.1-0.5 are
middle-aged SNRs (∼103 –104 years). Thus the highly-ionized ions could survive until today.

Metallicity of the Recombining Plasma
The abundances of the RD SNRs are ∼3 solar for W49B (Fe & Ni), ∼2 solar for IC 443 (Si &S), and <10 solar for G359.1-0.5 (Si & S). All these are signiﬁcantly higher than the solar values,
implying at least part of the recombining plasma is the ejecta origin. We here discuss this point.
In the case of photoionization at the explosion, temporary strong X-rays (afterglow of the
GRB, X-ray ﬂash, or breat-out X-rays) irradiate and ionize the matter outside the progenitor; itis diﬃcult to ionize the ejecta elements. Because typical abundance of the CSM is ∼1 solar, we
can say it is unlikely to form a high metallicity recombining plasma by the photoionization process
only.
On the other hand, it is possible to form overionized states of ejecta elements in the case of
the rarefaction scenario. Collision between the ejecta and CSM produce a reverse shock in theinitial phase of the SNR evolution (initial reverse shock in ﬁgure 8.10). This initial reverse shockheats the ejecta and then atoms are highly ionized. The subsequent blast wave break cools theshocked matter rapidly leaving the ionization state high. Thus, the recombining plasma of theejecta elements could be formed.

106

CHAPTER 8. DISCUSSION ON RECOMBINATION-DOMINANT SNRSᵧᵱᵫᵰᶃᶁᶍᶋᶀᶇᶌᶇᶌᶅ
ᵮᶊᵿᶑᶋᵿ
ᵧᶍᶌᶇᶘᶇᶌᶅ
ᵮᶊᵿᶑᶋᵿ

ᵡᵱᵫᶃᶈᶃᶁᶒᵿᵤᶍᶐᶕᵿᶐᶂᴾᶑᶆᶍᶁᶉ
Fig. 8.11.—: Schematic view of the plasma state in the rarefaction model. In the early phase of the SNR
evolution, shock-heated electrons ionize the ejecta and dense CSM and cools rapidly when the blast wave
breaks out of the CSM into the low density ISM. If there is no anisotropy in the ejecta or the stellar wind(CSM), the strongly recombining plasma are expected to form in a ring shape of CSM and part of ejecta.

Spatial Distribution
In the photoionization scenario of GRB afterglow or X-ray ﬂash, it is likely that the RD region
is anisotropically formed, such as in a bipolar shape. W49B might be a possible candidate 20 . In
the case of break-out X-rays, the overionized region is expected to be a spherically-symmetric shell,
which is not the case for W49B.
In the rarefaction model, we can expect the plasma structure as ﬁgure 8.11. The strongly
recombining plasma is produced mainly in the heated CSM, which was once highly-ionized in theearly phase of the SNR evolution. The heated ejecta could be recombining as we mentioned in thediscussion of the metallicity. In this scenario, the strongly recombining plasma would form in a
spherically-symmetric shell or sphere shape.
Our results of IC 443 are consistent with the rarefaction scenario. As we see in ﬁgure 6.11, thestrongly recombining plasma seems to form around the ejecta ring area which is revealed by XMMNewton (ﬁgure 6.2). We nevertheless cannot strongly conclude that it is spherically-symmetricshell like, because current two ﬁelds of view of Suzaku XIS do not cover the whole reverse shockarea. Thus, observations of the whole remnant are strongly required.
As we see in ﬁgure 5.12, the highly overionized region of W49B is not spherically-symmetric.
Instead, it shows asymmetrical structure which extends east and west. The strongly recombiningor highly overionized region is located in the west. It seems the rarefaction scenario cannot work
for W49B. However, if we interpret the explosion had occurred asymmetrically to the west, thescenario is still plausible. This is probable because there is a dense molecular cloud at the east endof W49B (see ﬁgure 5.2).
20Note that the photoionization senario is unlikely in perspective of metallicity.

8.3. FORMATION MECHANISMS

107

We expect that the highly overionized region corresponds to the steep gradient area in theelectron temperature and X-ray intensity from the rarefaction model. This picture fairly wellcoincide with our results of W49B and IC 443 (see ﬁgure 6.10 and 5.11, for example).

Surrounding Environment
We propose that the most plausible scenario for the formation of the overionized plasma is thedrastic cooling in the rarefaction process. In this case, some dense CSM should preexist aroundthe supernova progenitor at the moment of explosion. Because such CSM can be formed by astellar wind, the RD SNRs would originate from gravitationally collapsed supernovae 21 of massive
stars 22 . In this case, there might be a connection between RD SNRs and star-forming regions ormolecular clouds. We discuss this point.
IC 443 is located near the Gem OB1 association, which is one of the most massive molecular
cloud complexes in the Galaxy. Also, the progenitor of IC 443 is suggested to be a massive starwith strong stellar wind activity (Braun & Strom 1986; Meaburn et al. 1990). These facts strongly
support the existence of the dense CSM around IC 443 at the time of its explosion.
It is likely that W49B is located near the star forming region, W49A. The composite image(ﬁgure 5.2) is suggestive of bipolar structures surrounding massive stars. We can see the molecular
cloud in the east side of the remnant (ﬁgure 5.2). Thus, it is natural that we consider W49B as aremnant of gravitationally collapsed massive star.
G359.1-0.5 is surrounded by a molecular cloud (Uchida, Morris, and Yusef-Zadeh 1992).
Because it is located in the Galactic center region, there are many massive star-forming regions.
Accordingly, the progenitor of G359.1-0.5 was possibly a massive star.

21

Especially Type II SNe (the core collapse of red supergiants) because progenitors of Type II SNe blow slow stellar
winds with high mass-loss rates.
22
Note that this picture is also supported by the estimated mass of the RD SNRs (table 8.4).

108

8.4

CHAPTER 8. DISCUSSION ON RECOMBINATION-DOMINANT SNRS

Future Works

Research for Other RD SNRs and RD Regions
We have discovered the RD SNRs, but the sample is still limited to three cases. In order toinvestigate their characteristics and origin, we should further explore other SNRs with high statics.
The RD SNRs might have center-ﬁlled morphology in X-rays, exhibit strong Lyα lines in theirspectra, and associate with star-forming regions or dense molecular clouds. With these criteria,
the candidates would be W44, W28, 3C391, and 3C397. Another key to verify the formation
mechanisms is the spatial distribution of the RRC. Thus, the observation of the whole IC 443
region is strongly required.

High resolution X-ray spectroscopy
Although we ﬁrmly detected the recombination continuum and lines, we cannot separately
detected the ﬁne structures with the energy resolution of the X-ray CCD. For example, we cannot
clearly derive the recombination continua and lines. We also cannot distinguish the resonance,inter-combination, and forbidden lines.
X-ray micro-calorimeter on board Astro-H, which is the next Japanese X-ray satellite expected
to be launched in 2014, has a very high energy resolution of 7 eV between 0.3–12 keV. If we use thisdevice, we could completely divide the continuum and line emissions (ﬁgure 8.12). We, furthermore,
could separately detect the emission lines of resonance, inter-combination, and forbidden. As we
mention in chapter 4, the line intensity ratio between the resonance and forbidden is diﬀerent
among the collisional-dominant and recombination-dominant cases. Thus, X-ray micro-calorimeter
would reveal a new aspect for the study of the recombination-dominant plasma.

)H+HȘ
/\Ș
&DVFDGHOLQHV+H55&

Fig. 8.12.—: Simulated Astro-H micro-calorimeter spectrum of W49B in the 5–12 keV band.

Chapter 9

Conclusion
We have performed a recombination X-ray study from three SNRs (W49B, IC 443, and G359.10.5) utilizing the good energy resolution and high sensitivity of Suzaku. The important results are
summarized as follows:
1. Strong radiative recombination continua are ﬁrmly detected for the ﬁrst time in supernova
remnants. The detected continuum structures originate from highly-ionized iron atoms for
W49B, wheras silicon and sulfur atoms for IC 443 and G359.1-0.5. Recombination lines ofiron due to the cascade decays of electrons are also detected from W49B for the ﬁrst time.
2. For all these remnants, electron temperatures independently determined from thebremsstrahlung and recombination continuum shapes are consistent with each other for allthese SNRs. Volume emission measures for both components are the same order. Hence, all
the X-ray spectra of the observed SNRs can be explained by a single recombining plasmacomponent in the highly overionized state.
3. Spectra of these recombination-dominant SNRs all exhibit the strong recombination continuum structures, strong H-like Kα line intensity, and high center energies for He-like Kα lines.
All these SNRs have awkward and center-ﬁlled morphology in X-ray.
4. Spatial distribution of W49B and IC 443 both indicates signiﬁcant anisotropy of therecombination-dominant plasma. The region with strong recombination features corresponds
to a high gradient area of electron temperature and surface brightness. Combined with thespatial distribution, timescales of highly-ionized atoms, metallicity of recombining plasma,and surrounding environment of the remnants, the plausible candidate for strongly recombining plasma of a supernova explosion in a dense circumstellar matter is proposed.

109

Appendix A

Suzaku Observation of Tycho
Tycho was observed twice 1 with Suzaku in the large proposal project. Because the pointing
positions are almost the same, we merged these two data sets. The observation log is given intable A.1. We used the Lockman Hole observation 2 as background data. Three of the XISs
(two for FIs and one for a BI) were operated in the normal full-frame clocking mode with a SCI
technique during all the observations. Data were cleaned using processing version 2.2.11.24. We
used HEASOFT version 6.5.1 for data reduction and XSPEC version 11.3.2 for spectral analysis.
After screening with the standard criteria, the net integration time is ∼400 ks for the source regionand ∼70 ks for the background region.
Figure A.1a shows the XIS image of Tycho in the 0.5–7 keV band. We extract the source
spectrum from the circular region with a radius of 4 arcmin. Background spectrum was also
extracted from the Lockman Hole data with the same region in the detector coordinates. The
source spectrum is made by subtracting the background spectrum. In ﬁgure A.1b, we show the
XIS 0 spectra. We can see several prominent lines of K-shell emission from Ne, Mg, Si, S, Ar, Ca,
and Fe. The spectra of the two FI sensors are merged to improve the statistics. During the ﬁtting,
we ﬁx the interstellar extinction on Tycho to a hydrogen column density of 7 × 1021 cm−2 with thesolar elemental abundances, following Cassam-Chenai et al. (2007).

1

Suzaku also observed Tycho in June 2006 for 50 ksec. This is the data taken without SCI, and the energy gain
and response are signiﬁcantly diﬀerent from those of the SCI data. We therefore use only the SCI data.
2
This is the ﬁeld with no bright X-ray sources.

111

APPENDIX A. SUZAKU OBSERVATION OF TYCHO

112

Table A.1—: Observation log of Tycho and Lockman Hole.

∗(a)TargetStart Date

Obs.ID

Tycho ∗
Tycho ∗
Lockman Hole∗∗

2008-08-04
2008-08-11
2008-05-18

503085010
503085020
103009010

R. A.
Dec.
(J2000.0)00h 25m 20s 64◦ 08 1800h 25m 20s 64◦ 08 18
10h 51m 45s 57◦ 16 00Used as source data. ∗∗ Used as background data.

(b)

20:00.0

Fe-L

S

15:00.0'HFOLQDWLRQ

Exposure
(ks)
313
103
73

Ca

10:00.0

Ne Mg

Fe

Si

05:00.0

Ar
64:00:00.0

27:00.0

30.0

26:00.0

30.0

0:25:00.0

30.0

24:00.0

23:30.0

5LJKWDVFHQVLRQFig. A.1.—: (left) Suzaku XIS image in the 0.7–5.0 keV band of Tycho shown on a linear intensity scale.
Data from the three active XISs are combined. The white circle indicates the source region. The whitesquares are XIS ﬁeld of views for two observations. (right) XIS 0 spectrum extracted from a white circle in
the left ﬁgure. The energies of the prominent emission lines from speciﬁc elements are labeled. We show
the raw source spectrum in black, background spectrum in light gray, and background-subtracted source
spectrum in dark gray. Note that the background contribution is negligible.

Appendix B

Abbreviation• APEC: Astrophysical Plasma Emission Code(thin thermal plasma code in the ionization equilibrium state)• CCD: Charge Coupled Device• CFCI: Checker-ﬂag Charge Injection(one of the charge injection techniques of the Suzaku XIS)• CIE: Collisional Ionization Equilibrium• CSM: Circumstellar Matter• CTI: Charge Transfer Ineﬃciency• CXB: Cosmic X-ray Background• dof: Degree of Freedom• EOB: Extendable Optical Bench (equipped on Suzaku)• FOV: Field of View
• FWHM: Full Width at Half Maximum• He-RRC: Radiative Recombination Continuum accompanied by electron recombination with
H-like ions into the ground state of He-like ions
• Heα: He-like Kα(In the same way, Heβ and Heγ mean He-like Kβ and He-like Kγ, respectively.)• H-RRC: Radiative Recombination Continuum accompanied by electron recombination withfully-ionized ions into the ground state of H-like ions
• HXD: Hard X-ray Detector (onboard the Suzaku satellite)• ID: Ionization-dominant• ISM: Interstellar Matter
113

114

APPENDIX B. ABBREVIATION• Lyα: H-like Kα(In the same way, Lyβ and Lyγ mean H-like Kβ and H-like Kγ, respectively.)• NXB: Non X-ray Background• RD: Recombination-dominant• RRC: Radiative Recombination Continuum• SCI: Spaced-row Charge Injection(one of the charge injection techniques of the Suzaku XIS)• SN: Supernova• SNR: Supernova Remnant• XIS: X-ray Imaging Spectrometer (onboard the Suzaku satellite)• XRT: X-ray Telescope (onboard the Suzaku satellite)

References
Acciari, V. A., et al. 2009, ApJ, 700, 1034
Aharonian, F., et al. 2006, Physical Review Letters, 97, 221102
Albert, J., et al. 2007, ApJ, 664, L87
Anders, E., & Grevesse, N. 1989, Geochim. Cosmochim. Acta, 53, 197
Badnell, N. R. 2006, ApJS, 167, 334
Bamba, A., et al. 2009, ApJ, 691, 1854
Bamba, A., Yokogawa, J., Sakano, M., & Koyama, K. 2000, PASJ, 52, 259
Bautz, M. W., et al. 2004, Proc. SPIE, 5501, 111
Bocchino, F., & Bykov, A. M. 2003, A&A, 400, 203
Braun, R., & Strom, R. G. 1986, A&A, 164, 193
Brogan, C. L., & Troland, T. H. 2001, ApJ, 550, 799
Cassam-Chenaı̈, G., Hughes, J. P., Ballet, J., & Decourchelle, A. 2007, ApJ, 665, 315
Chevalier, R. A. 1974, ApJ, 188, 501
Cioﬃ, D. F., McKee, C. F., & Bertschinger, E. 1988, ApJ, 334, 252
Claussen, M. J., Frail, D. A., Goss, W. M., & Gaume, R. A. 1997, ApJ, 489, 143
Colgate, S. A. 1968, Canadian J. Phys., 46, 476.
Cornett, R. H., Chin, G., & Knapp, G. R. 1977, A&A, 54, 889
Cox, D. P., et al. 1999, ApJ, 524, 179
Denoyer, L. K. 1978, MNRAS, 183, 187
Dickel, J. R., Williamson, C. E., Mufson, S. L., & Wood, C. A. 1989, AJ, 98, 1363
Downes, D., Goss, W. M., Schwarz, U. J., & Wouterloot, J. G. A. 1979, A&AS, 35, 1
Duin, R. M., & van der Laan, H. 1975, A&A, 40, 111
115

116

REFERENCES

Esposito, J. A., Hunter, S. D., Kanbach, G., & Sreekumar, P. 1996, ApJ, 461, 820
Furuzawa, A., et al. 2009, ApJ, 693, L61
Gaensler, B. M., et al. 2006, ApJ, 648, 1037
Giovanelli, R., & Haynes, M. P. 1979, ApJ, 230, 404
Gray, A. D. 1994, MNRAS, 270, 835
Green D. A. 2009, http://www.mrao.cam.ac.uk/surveys/snrs/
Gwinn, C. R., Moran, J. M., & Reid, M. J. 1992, ApJ, 393, 149
Hewitt, J. W., Yusef-Zadeh, F., Wardle, M., Roberts, D. A., & Kassim, N. E. 2006, ApJ, 652, 1288
Hewitt, J. W., Yusef-Zadeh, F., & Wardle, M. 2009, ApJ, 706, L270
Hwang, U., Petre, R., & Hughes, J. P. 2000, ApJ, 532, 970
Ishisaki, Y., et al. 2007, PASJ, 59, 113
Itoh, H., & Masai, K. 1989, MNRAS, 236, 885
Iwamoto, K., et al. 1999, ApJS, 125, 439
Kawasaki, M. Ph.D. thesis; The University of Tokyo
Kawasaki, M. T., Ozaki, M., Nagase, F., Masai, K., Ishida, M., & Petre, R. 2002, ApJ, 572, 897
Kawasaki, M., Ozaki, M., Nagase, F., Inoue, H., & Petre, R. 2005, ApJ, 631, 935
Keohane, J. W., Reach, W. T., Rho, J., & Jarrett, T. H. 2007, ApJ, 654, 938
Kelley, R. L., et al. 2007, PASJ, 59, 77
Kokubun, M., et al. 2007, PASJ, 59, 53
Koyama, K., et al. 2007a, PASJ, 59, 23
Koyama, K., et al. 2007b, PASJ, 59, 245
Krause, O., Tanaka, M., Usuda, T., Hattori, T., Goto, M., Birkmann, S.,& Nomoto, K. 2008,
Nature, 456, 617
Kushino, A., et al. 2002, PASJ, 54, 327
LaMarr, B., et al. 2004, Proc. SPIE, 5501, 385
Leahy, D. A. 2004, AJ, 128, 1478
Masai, K. 1994, ApJ, 437, 770
Masai, K., Dogiel, V. A., Inoue, H., Schönfelder, V., & Strong, A. W. 2002, ApJ, 581, 1071
Matsumoto, H., et al. 2006, Proc. SPIE, 6266

REFERENCES

117

Mazzotta, P., Mazzitelli, G., Colafrancesco, S., & Vittorio, N. 1998, A&AS, 133, 403
Meaburn, J., Whitehead, M. J., Raymond, J. C., Clayton, C. A., & Marston, A. P. 1990, A&A,
227, 191
Mewe, R., Gronenschild, E. H. B. M., & van den Oord, G. H. J. 1985, A&AS, 62, 197
Miceli, M., Decourchelle, A., Ballet, J., Bocchino, F., Hughes, J. P., Hwang, U., & Petre, R. 2006,
A&A, 453, 567
Mitsuda, K., et al. 2007, PASJ, 59, 1
Moﬀett, D. A., & Reynolds, S. P. 1994, ApJ, 437, 705
Nakajima, H., et al. 2008, PASJ, 60 , 1
Nakayama, M., & Masai, K. 2001, ApJ, 375, 328
Nomoto, K., Thielemann, F.-K., & Yokoi, K. 1984, ApJ, 286, 644
Olbert, C. M., Clearﬁeld, C. R., Williams, N. E., Keohane, J. W., & Frail, D. A. 2001, ApJ, 554,
L205
Ozawa, M., et al. 2008, Proc. SPIE, 7011,
Ozawa, M., Koyama, K., Yamaguchi, H., Masai, K., & Tamagawa, T. 2009a, ApJ, 706, L71
Ozawa, M., et al. 2009b, PASJ, 61, 1
Perna, R., & Loeb, A. 1998, ApJ, 503, L135
Petre, R., Szymkowiak, A. E., Seward, F. D., & Willingale, R. 1988, ApJ, 335, 215
Porquet, D., Mewe, R., Dubau, J., Raassen, A. J. J., & Kaastra, J. S. 2001, A&A, 376, 1113
Prigozhin, G. Y., et al. 2004, Proc. SPIE, 5501, 357
Prigozhin, G. Y., et al. 2008, IEEE Trans. Elec. Device., 55 , 2111
Pye, J. P., Becker, R. H., Seward, F. D., & Thomas, N. 1984, MNRAS, 207, 649
Radhakrishnan, V., Goss, W. M., Murray, J. D., & Brooks, J. W. 1972, ApJS, 24, 49
Rasmussen, A. P., Behar, E., Kahn, S. M., den Herder, J. W., & van der Heyden, K. 2001, A&A,
365, L231
Rho, J., & Petre, R. 1998, ApJ, 503, L167
Rho, J., Jarrett, T. H., Cutri, R. M., & Reach, W. T. 2001, ApJ, 547, 885
Ruiz-Lapuente, P. 2004, ApJ, 612, 357
Sedov, L. I. 1959, Similarity and Dimensional Methods in Mechanics, New York: Academic Press,
1959

118

REFERENCES

Serlemitsos, P. J., et al. 2007, PASJ, 59, 9
Shelton, R. L. 1999a, ApJ, 521, 217
Shelton, R. L., Cox, D. P., Maciejewski, W., Smith, R. K., Plewa, T., Pawl, A., & Różyczka, M.
1999b, ApJ, 524, 192
Soderberg, A. M., et al. 2008, Nature, 453, 469
Smith, R. K., Brickhouse, N. S., Liedahl, D. A., & Raymond, J. C. 2001, ApJ, 556, L91
Takahashi, T., et al. 2007, PASJ, 59, 35
Tamagawa, T., et al. 2009, PASJ, 61, 167
Tavani, M., et al. 2010, ApJ, 710, L151
Tawa, N., et al. 2008, PASJ, 60, 11
Thielemann, F.-K., Nomoto, K., & Hashimoto, M.-A. 1996, ApJ, 460, 408
Troja, E., Bocchino, F., & Reale, F. 2006, ApJ, 649, 258
Troja, E., Bocchino, F., Miceli, M., & Reale, F. 2008, A&A, 485, 777
Truelove, J. K., & McKee, C. F. 1999, ApJS, 120, 299
Tsunemi, H., Yamashita, K., Masai, K., Hayakawa, S., & Koyama, K. 1986, ApJ, 306, 248
Tucker, W. H., & Gould, R. J. 1966, ApJ, 144, 244
Uchida, K., Morris, M., & Yusef-Zadeh, F. 1992, AJ, 104, 1533
Uchiyama, H., et al. 2009, PASJ, 61, 9Wargelin, B. J., Beiersdorfer, P., Neill, P. A., Olson, R. E., & Scoﬁeld, J. H. 2005, ApJ, 634, 687
Warren, J. S., et al. 2005, ApJ, 634, 376
Welsh, B. Y., & Sallmen, S. 2003, A&A, 408, 545
Yamaguchi, H., Ozawa, M., Koyama, K., Masai, K., Hiraga, J. S., Ozaki, M., Yonetoku, D. 2009,
ApJ, 705, 6
Yang, X. J., Tsunemi, H., Lu, F. J., & Chen, L. 2009, ApJ, 692, 894
Yang, X. J., Tsunemi, H., Lu, F. J., & Chen, L. 2009, ApJ, 692, 894
Yusef-Zadeh, F., et al. 1996, The Galactic Center, 102, 151
Yusef-Zadeh, F., Wardle, M., Rho, J., & Sakano, M. 2003, ApJ, 585, 319

Acknowledgments
Without the support from my powerful collaborators, this thesis could not have been completed. My deepest appreciation goes to Katsuji Koyama who has been patiently encouraging me
throughout my graduate course. I have learned a lot from him, not only the basis of X-ray astronomy, but also prospective attitude for the research and science. I express my gratitude to Hiroya
Yamaguchi in RIKEN for his very helpful support. I thank the ”encouraging person” Kuniaki
Masai for giving me important suggestion and willingly oﬀering me a lot of plasma radiation codes.
His prompt response always surprised me.
My gratitude is for Masaomi Tanaka in University of Tokyo and Nozomu Tominaga in Kohnan
University. I have learned a lot on supernovae and nucleosynthesis from them. Referring to the
W49B proposal, I thank Toru Tamagawa. His easygoing attitude and nice smile helped me a lot.
As for the calibration work of the Suzaku XIS, I was deeply helped by Hiroshi Nakajima, Hironori
Matsumoto, Hideki Uchiyama, and Katsuji Koyoma. I thank the contribution of all the XIS teammembers.
I would like to thank all the members in the Cosmic-ray laboratory group in Kyoto University.
I am grateful to Takeshi Go Tsuru and Hironori Matsumoto whose meticulous comments wereof enormous help. I am grateful to Hideki Uchiyama for so much kindness, Makoto Sawada andShinya Nakashima for checking the thesis and providing the ﬁgure, and Takao Ohnishi for helping
with G359.1-0.5.
Finally, I express supreme thanks to all of my friends and family. I especially thank Sou
Matsuura, Satoshi Okuzumi and Keisuke Onuma. It was my blissful moment to discuss the physics,
science, and views of life. I was very shocked by Teruaki Enoto, Yudai Suwa, Masaomi Tanaka,
Hajime Takami, and Shunsaku Horiuchi. It was very inspiring and interesting moment to talk tothem. It was a great honor to be a member of ”Four Devas” with Asami Hayato, Ryoko Nakamura,and Kaori Onda. I wish to express my sincere gratitude to my family, mother, father, grandmother,and sister, for their continuous and generous help.
I acknowledge ﬁnancial support from the Japan Society for the Promotion of Science. This
work is supported by the Grant-in-Aid for the Global COE Program ”The Next Generation of
Physics, Spun from Universality and Emergence”, Challenging Exploratory Research, from the
Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan.

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