X-Ray Study of Ejecta-DominatedSupernova Remnants with Suzaku
Hiroya Yamaguchi
Department of Physics, Graduate School of Science, Kyoto University
Kitashirakawa Oiwake-cho, Sakyo-ku, Kyoto, 606-8502, Japanhiroya@cr.scphys.kyoto-u.ac.jp

This thesis was submitted to the Department of Physics,
Graduate School of Science, Kyoto University
on January 2008
in partial fulfillment of the requirements
for the degree of Doctor of Philosophy in physics.

Abstract
We present the results of spectral and spatial studies of the young supernova remnants (SNRs:
SN 1006, E0509–67.5, RCW 86, Tycho, and N103B), using the data obtained by the Suzaku satellite.
Almost all of the heavy elements which exist in the present universe have been synthesized instars, and ejected to interstellar space by supernova (SN) explosions. Therefore, we are motivatedto study the nucleosynthesis and explosion mechanisms of SNe, in order to understand chemical
evolution of the universe. The young SNRs are bright and dominated by the emission from ejecta,
and hence they provide fruitful information about the elements synthesized by the SNe.
From SN 1006, we detected the K-shell emission from Fe, for the first time. The full-band
spectrum extracted from southeast of the remnant was well-fitted with a model consisting of three
optically thin thermal non-equilibrium ionization plasmas and a non-thermal component. Two of
the thermal components are due to ejecta, and have different ionization ages. The other thermal
component is associated with the swept-up interstellar medium (ISM). The Fe-K emission was
found in the ejecta with the lower ionization age.
The full-band spectrum of E0509–67.5 could be well-described by a composite model which
was very similar to the best-fit model for the SN 1006 spectrum, but an ISM component was not
required. The relative abundances of the ejecta component strongly suggest that this remnant is a
Type Ia SN origin.
In the analysis of the northwestern shell of RCW 86, we revealed for the first time the morphology of the Fe-K emission in this region. The origin of the emission was found to be Fe-rich ejecta
heated by the reverse shock very recently. The morphology and the plasma age of the ejecta suggest that the evolution of the reverse shock in RCW 86 is quite different from position to position,
because the density of the ISM around this SNR has a large spatial variation.
Analyzing the high-quality Fe-K line spectra extracted from Tycho and N103B, we gave thetight constraints to the ionization ages of Fe-K ejecta in each SNR. Additionally, we detected Cr-K
lines from both the SNRs, for the first time.
In all of the Type Ia SNRs which we observed (SN 1006, E0509–67.5, Tycho, and N103B), we
found that the ionization ages of Fe (in the ejecta) are significantly lower than those of the lighter
elements. It suggests that the Fe ejecta have been heated by reverse shock more recently than theother elements, because it concentrates toward the center of the remnant. At least in the Type Ia
SN explosion, the ejecta mixing effect inside the progenitor star is suggested to be almost negligible.

Contents
1 Introduction

1

2 Supernovae and Nucleosynthesis

3

2.1

Nucleosynthesis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

3

2.1.1

Stellar Nucleosynthesis . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

3

2.1.2

Explosive Nucleosynthesis . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

4

2.2Classification of SNe . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

5

2.3Type Ia SNe . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . .

7

2.4Type II SNe . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10

3 Supernova Remnants
3.1

13

Evolution of SNRs . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 13
3.1.1

Free Expansion Phase . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 13

3.1.2

Adiabatic Phase (Sedov Phase) . . . . . . . . . . . . . . . . . . . . . . . . . . 14

3.1.3

Radiative Cooling Phase . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15

3.1.4

Transition Phase Between the Free Expansion and Sedov Stages . . . . . . . 16

3.2

Shock Wave Heating . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 19

3.3

Thermal X-Ray Emission from SNRs . . . . . . . . . . . . . . . . . . . . . . . . . . . 20

3.4

3.3.1

Bremsstrahlung . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 21

3.3.2

Line Emission . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 22

Non-Equilibrium State . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 24
3.4.1

Thermal Non-Equilibrium . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 24i

ii

CONTENTS
3.4.2
3.5

Non-Equilibrium Ionization (NEI) . . . . . . . . . . . . . . . . . . . . . . . . 25

Cosmic-Ray and Non-Thermal Emission . . . . . . . . . . . . . . . . . . . . . . . . . 26
3.5.1

Fermi Acceleration . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 27

3.5.2

Synchrotron Emission . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 28

4 Instruments

31

4.1

Overview of Suzaku . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 31

4.2

XRT . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 33

4.3

XIS

. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 36

4.3.1

Overview . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 36

4.3.2

Characteristics and Performances . . . . . . . . . . . . . . . . . . . . . . . . . 39

4.4

HXD . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 43

5 SN 100647

5.1

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

5.2

Suzaku Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51

5.3

Overall Structure . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51

5.4

5.3.1

Thin-Thermal Plasma Map . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51

5.3.2

Spectrum of the SN 1006 SE Region . . . . . . . . . . . . . . . . . . . . . . . 53

5.3.3

Iron Line Map . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 53

5.3.4

Energy and Width of the Emission Lines . . . . . . . . . . . . . . . . . . . . 53

Spectral Structure in the Narrow Bands . . . . . . . . . . . . . . . . . . . . . . . . . 55
5.4.1

The 5–10 keV Band Spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . 55

5.4.2

The 0.4–1.1 keV Band Spectra . . . . . . . . . . . . . . . . . . . . . . . . . . 56

5.4.3

The 1.2–2.8 keV Band Spectra . . . . . . . . . . . . . . . . . . . . . . . . . . 58

5.5

Model Fit of the Full Band Spectra . . . . . . . . . . . . . . . . . . . . . . . . . . . . 61

5.6

Discussion for SN 1006 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 65
5.6.1

Origin of the Plasmas . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 65

5.6.2

Relative Abundance in the Ejecta. . . . . . . . . . . . . . . . . . . . . . . . 65

CONTENTS

iii

5.6.3

ISM Density and Ion Temperature . . . . . . . . . . . . . . . . . . . . . . . . 66

5.6.4

The Power-Law Component . . . . . . . . . . . . . . . . . . . . . . . . . . . . 66

5.6.5

Summary for SN 1006 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 68

6 E0509–67.5

69

6.1

Previous Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 69

6.2

Suzaku Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 70

6.3

Spectral Structure in the Narrow Bands . . . . . . . . . . . . . . . . . . . . . . . . . 72
6.3.1

The 4–10 keV Band Spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . 72

6.3.2

The 1.1–4.0 keV Band Spectrum . . . . . . . . . . . . . . . . . . . . . . . . . 73

6.4

Broadband Spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 76

6.5

Discussion for E0509–67.5 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 80
6.5.1

Origin of the Components . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 80

6.5.2

Relative Abundance and Mass of the Ejecta . . . . . . . . . . . . . . . . . . . 80

6.5.3

Contribution of the ISM . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 81

6.5.4

Summary for E0509–67.5 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 83

7 RCW 8685

7.1

Previous Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 85

7.2

Suzaku Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 88

7.3

Analysis and Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 89

7.4

7.3.1

XIS Image . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 89

7.3.2

Spectrum in the Hard X-Ray Band . . . . . . . . . . . . . . . . . . . . . . . . 90

7.3.3

Full Band Spectrum . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 92

Discussion for RCW 86. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95

7.4.1

Low Temperature Plasma . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95

7.4.2

Fe-K Line and High Temperature Plasma . . . . . . . . . . . . . . . . . . . . 96

7.4.3

Non-Thermal Emission

7.4.4

Unified Picture for All the Components . . . . . . . . . . . . . . . . . . . . . 97

. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 97

iv

CONTENTS
7.4.5

Summary for RCW 86 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 99

8 Fe-K Emission in Other SNRs
8.1

8.2

8.3

101

Tycho . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 101
8.1.1

Previous Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 101

8.1.2

Suzaku Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 102

8.1.3

Analysis and Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 102

N103B . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 107
8.2.1

Previous Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 107

8.2.2

Suzaku Observation . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 107

8.2.3

Analysis and Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 107

Summary of the Chapter . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 110

9 Discussion

111

9.1

Summary of the Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 111

9.2

Nucleosynthesis in Type Ia SNe . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 112

9.3

Ejecta Distribution . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 113

9.4

ISM Density and Evolutional State of the SNRs . . . . . . . . . . . . . . . . . . . . . 116

9.5

Future Works . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 119

10 Conclusion

121Chapter 1

Introduction
Figure 1.1 shows abundance ratio of elements (1 ≤ Z ≤ 30) in the solar photosphere, called“solar abundance” (Anders & Grevesse 1988), which is almost consistent to the average element
abundance ratio of the present universe. Immediately after the beginning of the universe, elements
other than hydrogen and helium (and tiny light elements such as lithium) did not exist. At present,
however, various heavy elements exist abundantly, as shown in figure 1.1. In addition to the elements
shown in this figure, the universe consists of more than 90 kinds of elements up to uranium, andabout 280 kinds of their isotopes. Almost all of these main heavy elements were synthesized instars, and were ejected to interstellar space by supernova (SN) explosions. Nucleosynthesis (nuclear

Fig. 1.1.—: Abundances of elements (relative to H) in the solar photosphere (Anders & Grevesse 1988).

1

2

CHAPTER 1. INTRODUCTION

fusion) is performed very actively just at the moment of the SN explosion as well. Therefore, we
are motivated to study the abundance ratio and total amounts of the heavy elements ejected byeach SN in order to understand chemical evolution of the universe.
Although SN explosions are extensively studied by optical observations, since the density of SN
ejecta immediately after the explosion is quite high, the light emitted from inside matters cannot
reach to the stellar surface directly. Therefore, we can only obtain the information of the synthesized
elements laid on the surface layer of SNe. However, the remnants of SNe, supernova remnants(SNRs), compose high temperature plasmas by interacting with interstellar medium (ISM), andprovide X-rays (or light in other wavelength) for a few tens of thousands of years after the explosion.
Since the plasma is optically thin, all the elements included in the SNRs can be observed by X-rays
directly at the same time. This is the most remarkable advantage of the X-ray observation on
SNRs.
Especially, young SNRs whose ages are a few hundred or about a thousand years are not only
bright but also dominated by the emission from ejecta (compared with the emission from sweptup ISM). Therefore, the observation on the young SNRs is the best method to investigate theabundances of the elements synthesized by the SNe.
Since the abundances of SNRs can be known by measuring the intensities of the emission lines
from highly ionized heavy elements, the detector sensitivity for the line spectrum is essential to theaccurate measurements. Suzaku, the newest Japanese X-ray astronomy satellites, has good energy
resolution, high efficiency, and low background level for diffuse sources. Especially, it is superior tothe other X-ray missions, in the sensitivity for the K-shell line emission from Fe, one of the main
products of SN explosions.
This thesis focuses on the study of X-ray emission from young SNRs utilizing the excellent
performances of Suzaku. In chapters 2 and 3, we describe the reviews of SNe and SNRs, respectively. The basic characteristics of the instruments on board Suzaku which we used in this work
are summarized in chapter 4. The following chapters provide the first results on the analyses ofindividual SNRs with Suzaku: SN 1006 in chapter 5, E0509–67.5 in chapter 6, RCW 86 in chapter 7, and Tycho and N103B in chapter 8. SN 1006, RCW 86, and Tycho are historical SNRs
in our Galaxy (the Milky Way). while E0509–67.5 and N103B are the young SNRs in the Large
Magellanic Cloud. Combining all the works, we develop the discussions in chapter 9, and we finallysummarize our findings in chapter 10.

Chapter 2

Supernovae and Nucleosynthesis
A Supernova (SN), an explosion of a star at the end of its lifetime, is one of the most dramatic
events in the universe. The brightness of SN is comparable to that of whole galaxy, and its explosion
energy is typically larger than 1051 ergs. It is believed that SN explosions occur per ∼30 years in
our Galaxy (e.g., van den Bergh & Tammann 1991). SN explosions synthesize various heavy
elements and release them into interstellar space. Furthermore, the shock wave generated by the
SN explosions produce density fluctuation of interstellar matter (ISM), which promote new starformation. Therefore, SNe play a very important role in the chemical evolution of the universe.
This chapter is devoted to explain the basic characteristics of SNe and their nucleosynthesis.

2.1

Nucleosynthesis

2.1.1

Stellar Nucleosynthesis

A star is maintained by balance between the thermal pressure and the gravity. Since stellar
temperature (T ) is in inverse proportion to a radius (R) of the star, if gravitational contraction
progresses and the central temperature of the star exceeds the critical value, a thermo-nuclear
fusion reaction will start.
When the central temperature is ∼ 10 7 K, hydrogen is ignited via p-p chain process:
p+p →

2

H + e+ + νe + 1.44 MeV

3He + γ + 5.49 MeV

He +3 He →4

He + 2p + 12.86 MeV ,

2
3

H+p →

while the central temperature exceeds ∼ 2 × 10 7 K, the hydrogen ignition proceeds dominantly via
CNO cycle:
p +12 C →13

N + γ + 1.95 MeV
34

CHAPTER 2. SUPERNOVAE AND NUCLEOSYNTHESIS

H
He
C+O
O+Ne+Mg
Si
Fe
Fig. 2.1.—: Onion-like structure in massive, evolved stars just prior to core collapse. (The scale is notexact.)
13

N →

13

C + e+ + νe + 1.37 MeV

14

N + γ + 7.54 MeV

14

p+ N →

15

O + γ + 7.35 MeV

15

15

N + e+ + νe + 1.86 MeV

12

C +4 He + 4.96 MeV .

p +13 C →O →p +15 N →Eventually, helium is synthesized from hydrogen by each process. The former is dominant in thesun. Stars spend most of their lifetime with the hydrogen burning, and hence they are called
main-sequence stars.
When the supply of hydrogen at the core is exhausted, the outward pressure is no longer
created. Then, the core begins to collapse, causing a rise in temperature and pressure whichbecomes high enough to ignite helium, and carbon and oxygen are synthesized by helium burning.
This precess progresses quickly in stars whose mass is larger than ∼ 2M . If the stellar massexceeds ∼ 8M , carbon burning takes place, and O+Ne+Mg core is produced. If the mass is larger
than ∼ 10M , silicon is synthesized by O burning, and Fe is synthesized by Si burning. As a result,the star has an onion-like structure as shown in figure 2.1, at the final stage of its evolution. Sucha massive star finally becomes a core-collapse SN (section 2.4).

2.1.2

Explosive Nucleosynthesis

The elements synthesized by explosive nuclear fusion depend on the temperature and density. The relation between dominant nucleosynthesis processes and temperature are summarizedin table 2.1. Since the explosive nucleosynthesis must take place quickly before the temperature
decreasing, it requires higher temperature than the stellar nucleosynthesis.
When the maximum temperature is ∼ 2 × 10 9 K, explosive carbon burning takes place:12

C +12 C →4

He +20 Ne ,

5

2.2. CLASSIFICATION OF SNETable 2.1:: Explosive nucleosynthesis.

Main products
Ne, O, Mg, Si
O, Si, S, Ar, Ca
Si, S, Ar, Ca
Fe, Co, Ni, Zn

C & Ne burning
O burning
Incomplete Si burning
Complete Si burning

Ignition temperature (K)2×109
3–4×109
4×10 95×10 9and explosive neon burning takes place, at the same temperature. Main reactions of the Ne burning
are as follows:
20
20
24

Ne + γ →4

He +16 O

Ne +4 He → γ +24 Mg

Mg +4 He → γ +28 Si .

If the maximum temperature is 3–4×10 9 K, explosive oxygen burning takes place:16
16

O +16 O →4

He +28 Si

O +16 O → γ +32 S .

Furthermore, 36 Ar and 40 Ca are synthesized by the reaction between 32 Si and 4 He. If the temperature exceeds ∼ 4 × 109 K, explosive silicon burning mainly takes place. Silicone and sulfur (and soon) absorb the nucleus of a helium and synthesize heavy elements, such as chromium, iron, nickel,
and zinc. In particular, a large amount of 56 Ni are synthesized in the complete Si burning process.
As described in the following section, the radioactive decay of 56 Ni (through 56 Co to 56 Fe) generates
the luminosity of SNe.

2.2

Classification of SNe

The fundamental classification scheme for SNe is traditionally based on their optical spectra
near maximum light, as given in figure 2.2. This classification is not reflecting the difference in theorigins of SNe as below-mentioned. If a optical spectrum lacks absorption lines of hydrogen Balmer
series, the SN is classified as Type I. A spectrum contains the hydrogen lines, on the other hand,
the SN is classified as Type II (Minkowski 1939; 1940). Type I SNe are sub-classified according to
the presence or absence of the strong Si II absorption feature at 6150 Å (e.g., Elias et al. 1985). The
SNe with strong Si line spectrum are defined as Type Ia, while the other Type I SNe have further
subclasses: Type Ib whose spectrum contains He lines, and Type Ic without He lines (Wheeler etal. 1987). On the other hand, subclasses of Type II SNe are distinguished by the shape of theirlight curves. Those with a pronounced “plateau”, producing a nearly constant luminosity ∼100days, are termed Type II-P, while those with a nearly linear decline in magnitude with time from6

CHAPTER 2. SUPERNOVAE AND NUCLEOSYNTHESIS

Maximum Light Spectrano H

HType I
strong Si

Type II

weak Si

Light curve

Ib / Ic

Type Ia
He

Type II-P

Type II-L

othersno He

Core-collapse
Type Ib

Type IcφGRB

Fig. 2.2.—: Classification sheme for supernovae based on the early spectra and light curve features.

Fig. 2.3.—: Light curves for Type Ia, Ib, II-P, and II-L supernovae (Wheeler et al. 1990).

2.3. TYPE IA SNE

7peak are termed Type II-L (e.g., Barbon et al. 1973). The typical light curves of each type of SNeare shown in figure 2.3. Although SN 1987A, occurred in the Large Magellanic Cloud (LMC) on
1987, is classified into Type II SN, the shape of its light curve is agreed with neither Type II-P nor
II-L.
The progenitors of Type Ia SNe are now believed to be near-Chandrasekhar mass (∼ 1.4M )white dwarfs, while Type II progenitors are believed to be massive stars (M >∼ 10M ). Details ofType Ia and II are described in sections 2.3 and 2.4, respectively. Although the Type Ib spectrum
lacks H absorption lines, the property of their light curves is very similar to that of Type II SNe.
Furthermore, Type Ib SNe are observed only in spiral galaxies and star forming regions. Therefore,
the progenitors of Type Ib SNe are thought to be massive stars that have lost their hydrogenenvelope. Wolf-Rayet stars are the possible progenitors, because they lose their outer hydrogenenvelope by means of very strong stellar wind. The progenitors of Type Ic are also believed to bemassive stars. There is some evidence that a few percent of the Type Ic SNe may be the progenitorsof γ-ray bursts (GRB).

2.3

Type Ia SNe

Type Ia SNe are observed in all types of galaxies including elliptical galaxies, where no massive
stars exist, and not concentrated in spiral arms (e.g., Wheeler 1982). The lack of absorption lines ofhydrogen in the early-time optical spectra of Type Ia SNe suggests that progenitors of Type Ia SNehave no outer hydrogen envelope. From these facts, thermonuclear explosions of mass-accreting
C+O white dwarfs in binary systems have been considered to be the most plausible models forType Ia SNe (e.g., Nomoto 1982).
−8−1
If the accretion from the companion star is as rapid as Ṁ >∼ 4 × 10 M yr , the white dwarf−6−1
can increase the C+O core mass. Especially at Ṁ >∼ 4 × 10 M yr , the accreting white dwarf
blows off a strong stellar wind, and increases the core mass by steady hydrogen burning with therate of ∼ 10−6 M yr−1 (Hachisu et al. 1996). When the mass of the white dwarf reaches about
the Chandrasekhar limit (∼ 1.4M ), the central density becomes near 2 × 10 9 g cm−3 , and thegravity exceeds the electron degeneracy pressure. As a result, the carbon burning is ignited at thecore. Nuclear fusion advances explosively, and the star is totally disrupted by the nuclear reactionenergy, leaving no neutron star nor black hole.

First, Arnett (1969) proposed that the flame front propagates at a supersonic speed (“detonation wave”). However, this theory failed because most of the material of the white dwarf are burned
to 56 Ni, which is contradictory to the fact that the strong absorption lines of Si are observed in the
spectra of Type Ia SNe. Then, the carbon deflagration model, whose flame speed is 10–30% of thesound speed, was proposed (W7 model: Nomoto et al. 1976; 1984). This model successfully reproduces the observational abundance ratio of the heavy elements, and hence is widely used as “thestandard model” for the Type Ia SNe. By the carbon deflagration, ∼ 0.6M of 56 Ni is synthesizedin the inner region of the star, and it powers the light curves of Type Ia SNe thorough radioactive

8

CHAPTER 2. SUPERNOVAE AND NUCLEOSYNTHESIS

decays via 56 Ni→56 Co→56 Fe. An e-folding time of 56 Co is ∼111 days, which corresponds to therate of the light curve decline of ∼0.01 mag day −1 .
However, the result of carbon deflagration depends on its flame speed, which is highly uncertain.
If the initial flame speed is much lower than that of W7 model, the hydrodynamical behavior isvery different. The deflagration might induce a detonation at low density layers. In the delayed
detonation model (Khokhlov 1991; Woosley & Weaver 1994) the deflagration wave is assumed to
be transformed into detonation at a certain layer during the first expansion phase. Iwamoto etal. (1999) constructed delayed detonation models with several parameterized transition densities.
The slow deflagration with a flame speed of 1.5% of the local sound speed is artificially transformed
into detonation when the density ahead of the flame decreases to 3.0, 2.2, and 1.7 × 10 7 g cm−3
(WS15DD3, WS15DD2, and WS15DD1, respectively). They found that the amount of synthesized
56 Ni (and the other elements) changes slightly.
In both models of the fast deflagration (W7) and the delayed detonation, the total explosiveenergy is estimated to be ∼1.3–1.5×10 51 ergs. Furthermore, the peak luminosity of the light curve
was almost consistent across Type Ia SNe, which led to the use of Type Ia SNe as a standard candle
to determine cosmological parameters.
Nucleosynthesis
The Mass of synthesized heavy elements by Type Ia SNe are given in table 2.2, which are
predicted by the deflagration model (W7: Nomoto et al. 1984) and the delayed detonation models(Iwamoto et al. 1999). Figure 2.4 shows the comparison between the abundance ratio of the W7
and that of the WS15DD2. The ratio predicted by the other delayed detonation model is similar
to the WS15DD2. We can see the abundances of the heavy elements such as Si and Fe (relative to
O-Ne-Mg) are larger in the delayed detonation model than in the W7 model.

92.3. TYPE IA SNE

Table 2.2:: Predicted mass of heavy elements nucleosynthesized by Type Ia supernova.

Species

Mass (M )W7∗
WS15DD1† WS15DD2†
C
3.2×10−2
5.4×10−3
9.0×10−3
N
2.6×10−9
2.8×10−4
2.7×10−4O
1.4×10−1
8.8×10−2
6.6×10−2−2−3
Ne
1.3×10
1.5×10
1.0×10−3
Mg
2.3×10−2
7.7×10−3
4.5×10−3Si
1.6×10−1
2.7×10−1
2.1×10−1
S
8.4×10−2
1.6×10−1
1.2×10−1−2−2
Ar
2.2×10
3.2×10
2.5×10−2
Ca
4.1×10−2
3.1×10−2
2.4×10−2
Cr
1.1×10−2
2.1×10−2
1.7×10−2−1−1
Fe
7.7×10
6.7×10
7.9×10−1
Ni
7.2×10−2
3.8×10−2
5.9×10−2∗ Nomoto et al. (1984).
† Iwamoto et al. (1999).

WS15DD3†
1.7×10−2
1.8×10−4
5.6×10−2
1.3×10−3
2.6×10−3
1.6×10−1
9.5×10−2
1.9×10−2
1.9×10−2
1.4×10−2
8.7×10−1
7.3×10−2

Fig. 2.4.—: Abundance ratio of elements (relative to solar values) in the carbon deflagration model (circle:
W7 model by Nomoto et al. 1984) and the delayed detonation model (triangle: WS15DD2 model by Iwamoto
et al. 1999). The ratio is normalized to oxygen.

10

2.4

CHAPTER 2. SUPERNOVAE AND NUCLEOSYNTHESISType II SNe

Since Type II SNe are observed in arms of spiral galaxies, their progenitors are believed to bemassive stars. A star with more than ∼10M finally has an onion-like structure (figure 2.1), andobtains an Fe core. Since Fe is the most stable element, the core can no longer generate nuclear
energy. However, neutrino continue taking away the energy from the star, and hence the density
and temperature in the Fe core continually increase due to the subsequent contraction. When theinternal temperature exceeds ∼ 3 × 10 9 K, photo-disintegration of Fe nuclei takes place:56
4

Fe → 13 4 He + 4n − 124.4 MeV

He → 2p + 2n − 28.3 MeV ,and the thermal energy is absorbed, as a result. Then, the star cannot support the gravitational
force any more and collapses gravitationally.
The energy explosively released by the infalling matter creates a shock wave. It propagates
outward and reaches the outermost part of the star. The total energy generated by the gravitational
collapse is ∼ 1053 ergs. However, most of the energy are carried off by neutrinos, and the kinetic
energy of the shock wave is typically 10 51 ergs. After the explosion, a neutron star or a stellar blackhole is left. The Crab pulsar (figure 2.5) is one of the most famous neutron star, which was left bya Type II SN explosion of A.D. 1054.

Fig. 2.5.—: Composite image of Crab pulsar and its nebula. Red and blue are respectively optical image
from Hubble telescope and X-ray image from Chandra.

11

2.4. TYPE II SNE

A decay rate of the Type II SN light curve is typically lower than that for the Type Ia SN.
Since the progenitor of Type II SN has a thick hydrogen envelope, the SN possesses the thermalenergy due to the shock heating over a long duration. Therefore, the light curve of most TypeII SNe shows plateau (II-P in figure 2.3). After the plateau phase, the light curve is powered by
56 Ni decay. The shape of Type II-L SN light curve is probably due to the expulsion of most of thehydrogen envelope of the progenitor star.
Nucleosynthesis
The shock wave propagating in the star elevates the temperature of the mantle, and leads explosive nucleosynthesis. The relation between the temperature (T ) behind the shock, the explosion
energy (E), and the stellar radius (R) is given asE∼

4π 3 4
R aT ,
3(2.1)

where a = 7.57 × 10−15 ergs cm−3 K−4 . Therefore, the maximum temperature, which decides thesynthesized elements (see table 2.1), is lower for the outer layer. Since the temperature of theoutermost layer is lower than that of ignition temperature of carbon burning, the abundances ofthe ejected material are also affected by the envelope, which depends on the mass of the star. The
mass and relative abundance of the synthesized elements for Type II SNe with 15M , 20M , and
25M are given in table 2.3 and figure 2.6, respectively (Thielemann et al. 1996).

12

CHAPTER 2. SUPERNOVAE AND NUCLEOSYNTHESISTable 2.3:: Predicted mass of heavy elements nucleosynthesized by core-collapse supernovae. (Thielemannet al. 1996)

Species
C
N
O
Ne
Mg
Si
S
Ar
Ca
Cr
Fe
Ni

SN II (15M )8.3×10−2
5.4×10−3
4.4×10−1
4.1×10−2
4.8×10−2
7.5×10−2
2.3×10−2
3.8×10−3
3.1×10−3
9.5×10−4
1.4×10−1
1.1×10−2

Mass (M )SN II (20M )1.1×10−1
2.7×10−3
1.5
2.6×10−1
1.8×10−1
1.0×10−1
2.5×10−2
4.5×10−3
3.9×10−3
1.1×10−3
8.0×10−2
7.4×10−3

SN II (25M )1.5×10−1
9.5×10−4
3.0
6.3×10−1
2.3×10−1
1.2×10−1
4.1×10−2
7.4×10−3
6.2×10−3
1.5×10−3
5.8×10−2
2.2×10−3

Fig. 2.6.—: Abundance ratio of elements (relative to solar values) synthesized by core-collapse supernovae with stellar mass of 15M (circle), 20M (square), and 25M (triangle), predicted by Thielemann etal. (1996). The ratio is normalized to oxygen.

Chapter 3

Supernova Remnants
A supernova remnant (SNR) is the structure resulting from the SN explosion. It consists ofejected material (ejecta) expanding from the explosion, and the swept-up ISM. Most SNRs havebeen discovered in radio survey. They are distinguished from H II regions by the relatively flat shapeof their radio spectra. Now, 265 SNRs are identified in our Galaxy (Green 2006), and about a halfof them are observed by X-rays. In this chapter, we explain the overview of physics concerning
SNRs: the macroscopic evolution of SNRs, the heating process performed by their shock wave, andthe mechanism of X-ray emission from the SNRs.

3.1

Evolution of SNRs

The evolution of SNRs can be divided into several phases. Their conditions are mainly determined by the ratio between the mass of ejecta (M Ej ) and that of swept-up ISM (MISM ). Here, we
assume that the density of the ambient medium is uniform. The detailed descriptions are shown in
Truelove & McKee (1999) for young stages (from free expansion phase through Sedov phase) andCioffi et al. (1988) for evolved stages (radiation phase).

3.1.1

Free Expansion Phase

Since MISM is negligible compared to MEj in the initial phase, the ejecta expand without
deceleration. The expansion velocity (∼ 10 4 km s−1 ) is much larger than the sound speed in theambient gas (∼ 10 km s−1 ), and hence shock wave (blast wave) precede the ejecta. The physics
of the shock wave is described in next section. Since most of the explosion energy (E) is released
kinetically (only 2–3% of the energy converted into thermal energy), the expansion velocity (v s )and the radius (Rs ) of the shock front can be expressed by
vs =

s

2E
E
= 8.5 × 108
51
MEj
10 erg13

1/2 

MEj
1.4M

−1/2

[cm s−1 ] ,

(3.1)

14

CHAPTER 3. SUPERNOVA REMNANTS
Rs = v s t ,(3.2)

where t is the time since the explosion. The mass of the ISM swept-up by the blast wave is described
as
4
(3.3)
MISM = πRs3 µmH n0 ,3
where µ, mH , and n0 are mean atomic weight of solar metallicity per hydrogen, mass of hydrogenatom, and hydrogen number density of the ISM, respectively. Since the free expansion phase continues until that MISM becomes comparable to MEj , its timescale can be approximately estimatedfrom equation 3.1, 3.2, and 3.3 to bet ∼ 2.5 × 102

E1051 erg

−1/2 MEj
1.4M

5/6 

µ1.4

−1/3 n0
1 cm−3

−1/3[yr] .

(3.4)

Therefore, this phase typically continues for a few hundred years. If the ambient density is low,
however, the timescale becomes longer. SN 1987A is believed to be in this phase.

3.1.2

Adiabatic Phase (Sedov Phase)

When MISM becomes much larger than MEj (typically MISM >∼ 10MEj ), the blast wave isdecelerated by the swept-up ISM. However, energy loss via radiation is still negligible comparedwith its initial energy. Therefore, the SNR expands adiabatically. In this stage, the shock evolution
can be approximated with that generated by a point explosion in a uniform medium, which is well
explained by a self-similar solution derived by Sedov (1959). Therefore, this phase is often called
as “Sedov phase”. The pattern of the gas flow can be characterized only by two parameters, E andn0 . The radius (Rs ) and velocity (vs ) of the blast wave, and the mean temperature just behind theshock front (Ts ) are written as
Rs = 4 × 1019
vs =



t
104 yrt
dRs= 5 × 107
4
dt
10 yr


Ts = 3 × 10

6t
104 yr2/5 

E
1051 erg

−3/5 

−6/5 

1/5 

E
51
10 ergE
51
10 ergn0
1cm−3

1/5 

2/5 

−1/5

n0
1cm−3

n0
1cm−3

−1/5

−2/5

[cm]

(3.5)

[cm s−1 ]

(3.6)

[K] .

(3.7)

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. Therefore, we can determine T s and n0 from an X-ray
spectrum of the SNR. If the distance to the SNR is known, R s is also determined observationally.
Then, the age of the SNR (t) and the explosion energy (E) can be derived according to equation 3.5
and 3.7.
The Sedov phase continues for a few ten thousand years. Until the end of this phase, ∼ 70%of the initial explosion energy is converted into thermal energy of swept-up ISM (Chevalier 1974).

3.1. EVOLUTION OF SNRS

15

Fig. 3.1.—: Cooling coefficient for an optically-thin thermal plasma with solar abundances (Anders &Grevesse 1989) as a function of plasma temperature. The individual contributions from major elements areshown with dashed lines, and black solid line shows the total value.

3.1.3

Radiative Cooling Phase

In the Sedov phase, the temperature behind the shock (T s ) decreases with time, according
to equation 3.7. Then, the cooling rate gradually increases, and hence the evolution is no longer
adiabatic. As shown in figure 3.1, when the temperature drops to about (2–3)×10 5 K, electrons
efficiently recombine with the heavy elements which are dominant in solar abundance ISM, such ascarbon and oxygen, and the gas radiates ultraviolet emission lines. Since the pressure in the shell
becomes lower as the temperature decreases, the materials in the shell are further compressed, andcool more rapidly. It provides a dense shell, while the interior is still hot and expands adiabatically.
Thus, the shell is driven by the pressure of the internal gas according to P V γ = const., where P
and V are the mean pressure and the volume of the internal gas respectively, and γ is the specificheat ratio (= Cp /CV ). For the non-relativistic single atomic gas, γ = 5/3. Then, the shock expands
with the time dependency ofRs ∝ t2/7 .
(3.8)
This stage is called pressure-driven snowplow phase (McKee & Ostriker 1977). Cygnus Loop is
believed to be in this phase.
As the temperature cools down further, the pressure can be ignored and the cool shell expands
at a constant radial momentum (Ms vs = const., where Ms is the mass of the shell) with the timedependency ofRs ∝ t1/4 .
(3.9)

16

CHAPTER 3. SUPERNOVA REMNANTS

This stage is called momentum-conserving snowplow phase (Cioffi et al. 1988).
When the shock speed becomes comparable to the proper motion of surrounding ISM (10–20 km s−1 ), the SNR lose its boundary between the outside and went out into the ISM. This phaseis called as “Disappearance phase”. The age of SNRs in this phase is larger than ∼ 10 6 yr.

3.1.4

Transition Phase Between the Free Expansion and Sedov Stages

In subsections 3.1.1 and 3.1.2, we considered the limited situations of M ISM  MEj and MISM MEj (although energy loss is negligible). However, the transition phase between these two stages,
when MISM is comparable to MEj , is very important, because the most of young SNRs, such as
Cas A and SN 1006, are probably in this stage. Figure 3.2 shows the realistic evolution of the blast
wave between the self-similar limits of the free expansion phase and the Sedov phase.
In this transition phase, the interaction with ISM cannot be ignored, although it is not yet
dominant. The swept-up ISM pushes back on the ejecta, which are thus decelerated. It causes
another shock wave to propagate inward through the ejecta, called as “reverse shock” (McKee
1974). Therefore, the ejecta are compressed and heated by the reverse shock. The boundary
between the ISM (shocked by the blast wave) and the ejecta (shocked by the reverse shock) is
called the “contact discontinuity”. The schematic view of SNRs in the transition phase is shownin figure 3.3. Only between the two shock waves, the material is hot enough to emit X-rays andthis forms the bright shell of young SNRs. In the central region, ejecta is not hot and is still
freely expanding. Such ejecta cannot emit X-rays, and hence its information cannot been received

Fig. 3.2.—: Position of the blast wave shock as a function of time from a simulation of uniform ejecta
expanding into a uniform ISM (Truelove & McKee 1999).

17

3.1. EVOLUTION OF SNRS

blast wave
(forward shock)shocked ejecta

shocked ISM

contact discontinuity
free expanding(unshocked) ejecta

reverse shock

ISM

Fig. 3.3.—: Schematic view of SNRs in the transition phase. Only the shocked material (between theforward and reverse shocks) is hot enough to emit X-rays.

Fig. 3.4.—: Composite image of E0102–72, an SNR in the Small Magellanic Cloud. Red and blue
respectively show the shocked ISM and the shocked ejecta. Green structures are the optical filaments.
Credit: X-ray (NASA/CXC/SAO); optical (NASA/HST); radio: (ATNF/ ATCA)

18

CHAPTER 3. SUPERNOVA REMNANTS

directly. Figure 3.4 shows a typical composite image of E0102–72, a young SNR in the Small
Magellanic Cloud (Gaetz et al. 2000). Blue shows X-rays from ejecta shocked by reverse shock,while red shows optical light from swept-up ISM.
As shown in figure 3.5, not only blast wave but also reverse shock initially propagate outward(as seen by an outside observer). After M ISM becomes greater than MEj , the reverse shock propagates back to the centre. In the case that uniform ejecta expand in homogeneous ISM, approximate
shock trajectories in the moderate transition phase (0.5 < t ∗ < 2.2, in figure 3.5) are given as
Rb∗ = (1.42t∗ − 0.254)2/5(3.10)

Rr∗ = t∗ (0.779 − 0.106t∗ − 0.533 ln t∗ )(3.11)

for the blast wave and reverse shock, respectively (Truelove & McKee 1999). R ∗ and t∗ are dimensionless values of radius and time normalized by these characteristic quantities: R ∗ ≡ R/Rch andt∗ ≡ t/tch . Rch and tch are those when MISM = MEj is achieved. The shock trajectories for the casesthat the ejecta distribution is not uniform but is power-law or exponential like are calculated by
Truelove & McKee (1999) and Dwarkadas & Chevalier (1998) in details.
Note that, however, they do not treat effects of the convective or Rayleigh-Taylor instability.
This instability appears when the blast wave is still inside the progenitor star and is believed to leadto mixing and clumping of the ejecta (e.g., Chevalier & Klein 1978; Muller & Arnett 1982), andarise again in the SNR evolution, at the contact discontinuity (e.g., Gull 1973; 1975). Therefore,
geometry of SNRs is generally complicated by irregularities in both the ejecta and ISM: shock
fronts are not exactly spherical (see figure 3.4).

Fig. 3.5.—: Shock positions (left) and velocities (right) for blast wave and reverse shock in remnants of
uniform ejecta. Note that the reverse shock velocity is measured in the frame of the unshocked ejecta aheadof it. Both figures are taken from Truelove & McKee (1999).

19

3.2. SHOCK WAVE HEATING

vsv1 = 0ISM

vs' = 0v2

SNR

Observer rest frame

vu = vsvd

ρu, puρd, pdISM

SNR(upstream)(downstream)

Shock rest frameFig. 3.6.—: A schematic view around the shock front in observer rest frame (left) and Shock rest frame(right). v, ρ, and p, are the velocity, density, and pressure, respectively.

3.2

Shock Wave Heating

When the shell of an SNR expands supersonically, a shock wave is formed, and the gas iscompressed and heated. We here take a coordinate system in which the shock is at rest, and define
upstream and downstream as figure 3.6. The tangential component of the gas velocity is assumed
to be zero on both sides. Then, conservation of mass, momentum, and energy across the shock
front can be expressed respectively asρu vu = ρ d vd
(3.12)
ρu vu2 + pu = ρd vd2 + pd

(3.13)

1 2
1
v + wu = vd2 + wd ,(3.14)
2 u
2where ρ, v, p, and w are the density, velocity, pressure, and enthalpy of both upstream (subscriptu) and downstream (subscript d), respectively. If the fluid is ideal gas, the enthalpy is given as
w = Cp T =γp,(γ − 1)ρ

(3.15)

where γ is the specific heat ratio (= C p /CV ). Therefore, from equations 3.12, 3.13, 3.14, and 3.15,
relations between the physical condition of the upstream and the downstream (Rankine-Hugoniot
relations) are obtained asρuvd(γ + 1)pu + (γ − 1)pd==
(3.16)
ρdvu (γ − 1)pu + (γ + 1)pd
pu (γ + 1)ρu − (γ − 1)ρd=pd
(γ + 1)ρd − (γ − 1)ρu

Td
pd ρu
pd (γ + 1)pu + (γ − 1)pd==.
Tu pu ρd pu (γ − 1)pu + (γ + 1)pd

(3.17)(3.18)

20

CHAPTER 3. SUPERNOVA REMNANTS

For the velocities of the gas, we obtain1[(γ − 1)pu + (γ + 1)pd ]2ρu(3.19)

1 [(γ + 1)pu + (γ − 1)pd ]2.
2ρu (γ − 1)pu + (γ + 1)pd

(3.20)

vu2 =
vd2 =

Assuming the strong shock (pd /pu  1), equations 3.16, 3.18, 3.19, and 3.20 can be respectively
rewritten as
ρu
vd
γ −1
=
=
(3.21)
ρdvu γ + 1
pd γ − 1Tu=Tdpu γ + 1(3.22)(γ + 1)pd
2ρu(3.23)(γ − 1)2 pd.
2(γ + 1)ρu

(3.24)

vu2 =
vd2 =

Here, vu equals to vs , the shock velocity at the rest frame of the upstream (see figure 3.6), Using
equations 3.21 and 3.24, therefore, we can derive the mean post-shock temperature:kTd = µmH

pd 2(γ − 1)=µmH vs2 ,ρd(γ + 1)2(3.25)

where k, µ, and mH are the Boltzmann constant, the mean atomic weight, and the hydrogen mass,
respectively. For the non-relativistic single atomic gas (γ = 5/3), the ratio (among the upstream
and downstream) of the density and velocity are derives as ρ d /ρu = 4 and vd /vu = 1/4, respectively.
Moreover, the post-shock temperature is described askTd =3µmH vs2 .
16(3.26)The shock wave heating can be understood as the energy conversion from the kinetic energy to theinternal (thermal) energy.

3.3

Thermal X-Ray Emission from SNRs

In this section, we summarize thermal X-ray emission mechanisms concerning SNRs. They
tell us a lot of informations about the SNR itself. Since the plasma in SNRs are optically-thin, themain emission processes which we observe are bremsstrahlung continuum and line emissions.

21

3.3. THERMAL X-RAY EMISSION FROM SNRS

3.3.1

Bremsstrahlung

Charged particles in the Coulomb field of other charged particles emits radiation called
bremsstrahlung or free-free emission. The most intense emission in neutral plasma is thebremsstrahlung from electrons by the Coulomb field of ions. The emission power of the emission
per unit time, volume, and frequency is given bydW
16π 2 e6= √ 3 2 ne ni Z 2 gf f (v, ν) ,dV dtdν
3 3c me ve(3.27)

where ne , ni , e, me , ve , c, Z, and gf f are electron and ion density, charge unit, electron mass
and velocity, light speed, and Gaunt factor, respectively. Gaunt factor is a certain function of theenergy of electron and of the frequency of emission, reviewed in Bressaard & van de Hulst (1962),
and Karzas and Latter (1961).
The shocked ISM and ejecta are heated up to ∼ 10 7 K and become thin-thermal plasma−3(n <∼ 1 cm ). When the electron temperature is kT e , the probability dP that a particle has avelocity in the velocity range dve is estimated under the assumption of an isotropic distribution ofvelocities to be!
me ve22
dP ∝ v exp −dv .
(3.28)
2kTe
Combining equations 3.27 and 3.28, and performing the integral, the emission power of plasma in
an isotropic thermal distribution (thermal bremsstrahlung) is derived asdW
dV dtdν2π 1/2 −1/2 2
32πe6
TeZ ne ni e−hν/kTe gf f=
3mc3 3kme= 6.8 × 10−38 Z 2 ne ni Te−1/2 e−hν/kTe gf f [ergs s−1 cm−3 Hz−1 ],(3.29)(3.30)

where gf f (Te , ν) is a velocity averaged Gaunt factor. For typical thermal plasma emitting X-rays,
gf f =−1/2

The factor Te



3 kTe
π hν

−1/2

.

(3.31)

comes from the fact that dW/dV dtdν ∝ v −1 .

Integrating equation 3.29, the total power per unit volume emitted by thermal bremsstrahlung
is obtained asdW
dtdV

32πe6 2πkTe 1/2 2
Z ne ni gB
3hmc3 3me= 1.4 × 10−27 Z 2 ne ni Te1/2 gB [ergs s−1 cm−3 ],

=





(3.32)(3.33)The density and/or temperature of the plasma are higher, the power of the bremsstrahlung emission
becomes higher. Therefore, we can investigate the densities (of electrons and ions) and electron
temperature from the intensity and shape of the thermal continuum spectrum.

22

3.3.2

CHAPTER 3. SUPERNOVA REMNANTS

Line Emission

The transitions between two discrete quantum levels often accompany photon emission called
line emission or bound-bound emission. In the pure Coulomb cases for hydrogen atom or otherhydrogen-like ions (which have only one captured electron), such as He II, O VIII, and Si XIV, a fairly
complete discussion of transition rates can be given by quantum theory. For hydrogen atom, theenergy of a photon absorbed or emitted in a transition between two discrete levels with principal
quantum numbers n and n0 is given by
Ebb = Ry



1
1
− 02
2
n
n



,

(3.34)

where Ry is Rydberg constant (= 13.6 eV). For heavy elements, E bb roughly becomes2

Ebb ∼ Z Ry



1
1
− 02
2
n
n



.

(3.35)

These H-like transitions are named as “Lyman series” after its discoverer, Theodore Lyman. The
transitions are named sequentially by Greek letters: transition of 1s–2p is Lyα (H-like Kα), 1s–3p
is Lyβ (H-like Kβ), 1s–4p is Lyγ (H-like Kγ), and so on, where s and p mean azimuthal (orbital
angular momentum) quantum numbers (l) are 0 and 1, respectively. The line energies of Lyman
series of major elements are listed in table 3.1.
If the ions have two or more electrons, the transition process becomes relatively complicated.
For example, He-like ions (which have two electrons) emit three most intense lines: resonance (1s 2
1 S – 1s2p 1 P ), forbidden (1s2 1 S – 1s2s 3 S ), and inter-combination (1s 2 1 S – 1s2p 3 S )
0
1
0
1
0
2,1

Table 3.1:: K-shell emission line energies of H-like and He-like ions.

Element

Line energy (eV)H-like ion
He-like ion∗
Lyα Lyβ Lyγ
Kα(r) Kα(f ) Kα(i) Kβ
Kγ
C
367
436
459
308
299
304
355
371
N
500
593
625
431
420
426
498
522
O
654
774
817
574
561
569
666
698
Ne
1022 1211 1277
921
905
914
1073 1127
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
Fe6966 8266 8732
6702
6641
6670 7798 8217
∗ r, f , and i indicate resonance, forbidden, and inter-combination lines.

3.3. THERMAL X-RAY EMISSION FROM SNRS

23

Fig. 3.7.—: Ion fractions of O (top-left), Si (top-right), S (bottom-left), and Fe (bottom-right) as a functionof temperature. Plasma condition is assumed to be in collisional ionization equilibrium (CIE) state. In eachpanel, red, green, and blue, represent He-like, H-like, fully ionized ions, respectively. Black represent ionswhose ionization state is lower than He-like.

24

CHAPTER 3. SUPERNOVA REMNANTS

transition lines. The relative intensities of the these lines depend on the temperature because thecollisional excitation rates have not the same dependence with temperature among the three lines.
The line energies of He-like ions are listed in table 3.1. Those of He-Kβ (1s 2 –1s3p) and He-Kγ(1s2 –1s4p) are also listed.
The ionization states depend on the electron temperature. Figure 3.7 shows the ion fractionsof O, Si, S, and Fe as a function of temperature, assuming that the plasma reaches collisional
ionization equilibrium (CIE) condition (see subsection 3.4.2). When the ionization states are lower
than the He-like, the energies of emission lines become lower than those of He-like ion.
Ionized thin-thermal plasma are collisionally excited by free electrons. Since the lifetime of theexcited states are much shorter than the average timescale of collision between the electrons andions, we can assume that the rate of photon emission is the same as that of excitation. Therefore,
the line intensity is proportional to n e ni , which determines the probability of collisions between
electrons and ions (ni is the ion density). We can hence derive relative abundances by comparing
line intensities from different elements. Moreover, if the plasma is dominated by hydrogen, theintensity of bremsstrahlung is proportional to n e nH , and hence we can derive absolute abundances(with respect to H) by comparing the intensities of the line and continuum.

3.4

Non-Equilibrium State

3.4.1

Thermal Non-Equilibrium

The hot gas in SNRs is heated by “collisionless shocks”, generally propagating in extremely−3
low density (<∼ 1 cm ), where the shock transition occurs on a length scale much shorter than
the typical particle mean free path to Coulomb collisions (Laming 2000). It is suggested that thecollisionless shock of young SNRs do not produce electron-ion temperature equilibration. According
to equation 3.26, the temperature in the downstream plasma for each particle species a is
kTa =3ma vs2 ,16(3.36)

with ma the particle mass, vs the shock velocity. Therefore, in the limit of true collisionless plasma,
the particles will have temperature proportional to their masses: protons will have a temperature
mH /me ∼ 1800 times that of the electrons. In realistic conditions, Coulomb equilibration will takeplace, with rated∆TZ 2 ∆T= −0.13,(3.37)dne t
A Te3/2
where ∆T = Ti − Te , and Z and A are respectively charge and atomic mass of the ions (e.g., Spitzer
1962; Laming 2001). Assuming Ti > Te , we can integrate equation 3.37 to get
Te =Z 2 Ti ne t
5
× 0.13
2
A Z +1!2/5.

(3.38)

25

3.4. NON-EQUILIBRIUM STATEIn the case that the ions are hydrogen atoms (Z = A = 1), we obtainkTH = 2.4 × 1011



ne t
1 cm−3 s

−1 

kTe
1 keV



[keV] .

(3.39)

Therefore, timescale (teq ) for electrons and protons to reach equilibrium (energy equipartition) isroughly estimated to be teq ∼ 1011 n−1
e s. This corresponds to a few or ten thousand years in aplasma of ne = 1 cm−3 .
Although the heating process in collisionless shocks is not well understood, a circumstantial
evidences for the non-equilibration have been obtained by the observations of young SNRs, such as
Cas A, Tycho, and SN 1006: there are no object whose electron temperature exceeds 5 keV, whereasa typical shock velocity of ∼ 2000–5000 km s −1 should give rise to a mean plasma temperature of> 10 keV (e.g., Hughes et al. 2000b). Ion temperature can be measured from the line broadening∼
due to thermal Doppler effect. However, this measurement is very difficult using present X-ray
detectors because of their poor energy resolution. In future, ion temperature of thermal plasma inyoung SNRs will be able to be measured with micro-calorimeter on board NeXT satellite, whoseenergy resolution will be ∼ 5 eV (FWHM) in whole X-ray band.

3.4.2

Non-Equilibrium Ionization (NEI)

After the shock heating, the initial ionization state of the shocked material is considered tobe nearly neutral. The ions are gradually losing electrons one-by-one through collisions with freeelectrons until their ionization states are in equilibrium with the thermal electrons (CIE condition).
The characteristic timescale for a plasma to reach CIE was derived by Masai (1984; 1994).
The collisional ionization rate equation for an element of atomic number Z is written as,dfz= Sz−1 fz−1 + (Sz + αz )fz + αz+1 fz+1d(ne t)Z
X

fz = 1 (z = 0, 1, ..., Z) ,(3.40)(3.41)

z=0

where fz is the ionization fraction of the element which is ionized (z − 1) times, and S z and αz
respectively represent the ionization and recombination rate coefficients for the z-th ionized ion.
In the CIE condition, dfz /d(ne t) = 0. Therefore, the timescale (tion ) to reach the CIE is estimatedwith the zero-th order approximation to be,
ne tion ≈Z
X(Sz + αz )−1

z=0≈ [(Sz + αz )min|Sz −αz | ]−1≈ 1012 [cm−3 s] ,(3.42)

where (Sz + αz )min|Sz −αz | is the absolute value of the diagonal element that gives the minimum
difference between Sz and αz (Masai 1993). The timescale is nearly independent of Z and T e , but

26

CHAPTER 3. SUPERNOVA REMNANTS

depends only on ne . Therefore, “ionization age” (or “ionization parameter”) is defined as τ = n e t,which is often used as a key diagnostic of the non-equilibrium ionization (NEI) state. As given byequation 3.42, τ is typically required to be ≥ 10 12 cm−3 s for full ionization equilibrium. Therefore,
plasmas in most of young SNRs are in the NEI condition.

3.5

Cosmic-Ray and Non-Thermal Emission

Cosmic-rays are particles with very high energy up to 10 20 eV. The energy density of cosmicrays (∼ 1 eV cm−3 ) is larger than those of starlight, Galactic magnetic field and cosmic micro-wave
background. Since the discovery of cosmic-rays (Hess 1912), the origin and mechanism of theacceleration have been one of the biggest long-standing problems.
As shown in figure 3.8, the spectrum of cosmic-rays shows a single power-law up to the knee
energy (∼ 1015.5 eV), where the gyro radius of electrons in the typical interstellar magnetic field ismuch smaller than the galactic radius. Therefore, cosmic-rays below the knee energy are likely to beof galactic origin. The discoveries of synchrotron X-rays from SN 1006 have indicated that electrons
are accelerated close to the knee energy (Koyama et al. 1995). Besides SN 1006, synchrotron X-ray
emissions were discovered from other SNRs: e.g., RX J1713.7.3946 (Koyama et al. 1997; Slane etal. 1997), Cas A (Vink et al. 2000a; Hughes et al. 2000a), and Tycho’s SNR (Hwang et al. 2002).

Fig. 3.8.—: Energy spectrum of cosmic protons (Sokolsky 1989). The vertical axis is written in differential
flux times E 2.5 .

27

3.5. COSMIC-RAY AND NON-THERMAL EMISSION

downstream
vd

upstreamvu

test particle

E1

E0

E0'

E0'

Fig. 3.9.—: A schematic view of first order Fermi acceleration.

Recently, a significant number of SNRs with synchrotron X-rays (called “non-thermal SNRs”) havebeen discovered in Galactic plane (e.g., Ueno 2005). These discoveries provide good evidence forthe cosmic-ray acceleration at the shocked shell of SNRs.
In this section, we explain the mechanism of shock acceleration performed at the SNR rims
and synchrotron emission from accelerated electrons.

3.5.1

Fermi Acceleration

The most plausible process of the cosmic-ray acceleration is first order Fermi acceleration (Bell1978; Blandford & Ostriker 1978), which takes place at shock front of SNRs. Similarly to section 3.2,
we define the fluid velocities of upstream and downstream as v u and vd (vu > vd ), respectively. Here,we consider a test particle in the shock fluid. (Schematic view is shown in figure 3.9.) The particle
scattered back elastically in the downstream changes its energy from E 0 into E00 as
E00 ' γd2 (1 − βd )2 E0 ,(3.43)q

where γd ≡ 1 − βd2 and βd ≡ vd /c. c is the light speed. Conversely, the particle energy changes in
the upstream by a scattering as
E1 ' γu2 (1 + βu )2 E00 ' (1 − 2βd )(1 + 2βu )E0 ,(3.44)

where we assume γu ' 1, γd ' 1, and βd  1, βu  1, which indicates that the shock speed is
non-relativistic. Since βu is larger than βd , particles obtain energy from the shock fluid during
each round-trip between the two regions. The scattering occur through the interactions between
charged particles and the magnetic field, not through Coulomb collisions. In the case of 3-dimension,
particles are scattered in various direction, therefore equation 3.44 is somewhat modified. The mean
particle energy after n round trips is given as (Bell 1978)
4
En = E0 · expn(βu − βd )
3




.

(3.45)

28

CHAPTER 3. SUPERNOVA REMNANTS

Bell (1978) also calculated the probability (≡ η s ) that the scattered particles escape from theacceleration region for each round trip asηs = 4vd,
vp(3.46)

where vp is a particle velocity. Hence the probability that a particle makes n round trips is givenasPn = (1 − ηs )n = (1 − 4βd )n ,(3.47)
where we assumed vp = c (particle is relativistic). Using the equation 3.47, we can calculate thedifferential energy spectrum of accelerated particles as
N (E) ∝

β +2βddPn− u∝ E βu −βd .
dEn(3.48)In the case of βd = 14 βu (equation 3.16 for γ = 5/3), we obtain an equation of
N (E) ∝ E −2 .

3.5.2

(3.49)

Synchrotron Emission

Charged particles emit synchrotron radiation in magnetic field. The synchrotron radiation
power per unit frequency from a single electron of energy E in a magnetic field B is given as(Rybicki & Lightman 1979)√ 3 
3 e B sinαωP (ω, α) =F,
(3.50)
2
2π me cωcwhere α and F (x) are respectively the pitch angle and the function given by Rybicki & Lightman(1979). The peak frequency ωc is
3γ 2 eB sin α
ωc =.
(3.51)
2me c
Note that protons radiate little synchrotron radiation since they are as 1800 times heavy as electrons. Integrating over ω, we found that total emission power is written as
4
Psync = σT cβ 2 γ 2 UB .
3

(3.52)

Here σT = 8πr02 /3 is the Thomson cross section, and U B (= B 2 /8π) is the magnetic field energy
density.
Assume that the number density of electrons is represented by a power-law distribution expressed in the form
N (E)dE = CE −p dE ,(3.53)orN (γ)dγ = Cγ −p dγ .

(3.54)

29

3.5. COSMIC-RAY AND NON-THERMAL EMISSION

The quantity C can vary with pitch angle and the like. Then, the total power radiated per unit
volume and unit frequency is given by the integral of N (γ)dγ times the single particle radiation
formula over all energies or γ. Thus, we have
Ptot (ω) = CZP (ω)γ

−p

dγ ∝

Z

F



ω
γ −p dγ .
ωc

(3.55)Let us change variables of integration to x ≡ ω/ω c , noting ωc ∝ γ 2 ;Ptot (ω) ∝ ω −(p−1)/2

Z

F (x)x(p−3)/2 dx .

(3.56)The integral of this equation becomes constant, then we have
Ptot (ω) ∝ ω −(p−1)/2 ,(3.57)

so that the spectral index of synchrotron radiation (s) is related to the particle distribution index
p by
p−1
.
(3.58)
s=
2
Therefore, power-law emission is an indicator of synchrotron emission, namely accelerated electrons.
In the case of first order Fermi acceleration which gives an energy/ index of electron to be p ∼ 2(equation 3.49), the spectrum index of the synchrotron emission is 0.5 (photon index = 1.5). Thisvalue is almost consistent with the canonical value of the non-thermal radio spectrum from SNR
shell.

Chapter 4

Instruments
4.1

Overview of Suzaku

Suzaku (Mitsuda et al. 2007) is the fifth Japanese X-ray astronomy satellites, developed by aJapanese–US collaboration. It was launched from Uchinoura Space Center (USC) on 10 Jury 2005,
and successfully put into a near-circular orbit at ∼570 km altitude with an inclination angle of 31 ◦ .
Figure 4.1 shows a schematic of view of Suzaku in orbit. The spacecraft length is 6.5 m along thetelescope axis after deployment of the extensible optical bench (EOB).
Suzaku carries three scientific detector systems: Hard X-ray Detector (HXD: Takahashi etal. 2007), X-ray Imaging Spectrometer (XIS: Koyama et al. 2007), and X-ray Spectrometer (XRS:

Fig. 4.1.—: Schematic view of the Suzaku satellite in orbit. On the top, the X-ray telescopes (XRT) canbe seen. This picture is taken from Mitsuda et al. (2007)

31

32

CHAPTER 4. INSTRUMENTS

Kelley et al. 2007). The XIS and XRS are placed in the focal plane of X-Ray Telescope (XRT:
Serlemitsos et al. 2007). The configuration of the systems are shown in Figure 4.2, where XRT-I
and XRT-S are telescopes for the XIS and XRS, respectively.
Although the XRS, one of the main instruments of Suzaku, became not operational because of
liquid helium loss, all the other detectors are still working properly. Therefore, Suzaku retains its
excellent X-ray sensitivity in a wide-energy range of 0.2–600 keV, with moderate energy resolution
and a low background level. We show in following sections the characteristics of each operating
instrument.

Fig. 4.2.—: Side view of Suzaku with the internal structures after the extensible optical bench deployment.
This picture is taken from Mitsuda et al. (2007).

33

4.2. XRTFig. 4.3.—: Wolter type I reflecting optics.

4.2

XRT

Five sets of the XRT were developed jointly by NASA/GSFC, Nagoya University, Tokyo
Metropolitan University, and ISAS/JAXA. One of them is equipped for the XRS (called XRTS), and the others are equipped for the XIS (XRT-I). Since the XRT-S is no more functional,
hereafter the XRT-I is called simply “XRT”. All of the four XRT are co-aligned to image the same
region of the sky.
Telescope Design
The XRT is configured to form a Wolter type I grazing-incidence reflective optics (figure 4.3).
Incident X-rays are reflected twice by parabolic surface and hyperbolic surface, and are condensed
to the focal plane, when their incident angle is less than a certain value (see table 4.1). Since theprojected area per reflector is very small, the mirrors are formed by a large number of thin-foils inorder to enlarge the effective area. The design parameters and basic performance of the XRT aresummarized in table 4.1.
Characteristics
The left panel of figure 4.4 shows the energy dependence of the total on-axis effective areaof the four XRT, including the detector efficiency, compared with those of other X-ray missions.
The XRT provide a large efficiency with a moderate imaging capability in the energy range of
0.2–12 keV. The effective area of the four XRT modules is much larger than that of Chandra inalmost whole energy band, and comparable to that of XMM-Newton above ∼5 keV. As shownin the right panel of figure 4.4, The effective area also depends on the off-axis angle, because thereflection coefficient decreases with increasing the X-ray incident angle. This dependence is called“vignetting”, calculated by the ray-tracing simulator (Ishisaki et al. 2007), and confirmed in the

34

CHAPTER 4. INSTRUMENTS

Table 4.1:: Design parameters and performance of the XRT (Serlemitsos et al. 2007).

Focal length
4.75 mSubstrate
Al (152µm)
Reflector
Au (>0.1 µm)
Number of nesting
175
Inner / outer diameter 118 mm / 399 mm
Incident angle0.◦ 18–0.◦ 60
Geometrical area
873 cm2
Field of view∗
200 (@ 1 keV) / 140 (@ 7 keV)
Effective area
450 cm2 (@ 1.5 keV) / 250 cm2 (@ 7 keV)
Half Power diameter
2.0 0∗ Diameter of the area within which the effective area is more than
50% of the on-axis value.

orbit by the observations of the Crab Nebula (Serlemitsos et al. 2007).
Since the angular resolution of the XRT is limited, the image of a point source becomes blurredon its focus. The Point-Spread Function (PSF) and Encircle-Energy Fraction (EEF) are shown inleft and right panel of figure 4.5, respectively. The half power diameter (HPD), defined as thediameter within which the X-ray flux is 50% of the total flux of a point source, is ∼2 arcmin foreach XRT. This value is significantly improved from ASCA (∼ 3. 0 6).
The response of the XRT is described by Ancillary Response Files (ARFs). In all spectral
analyses described in following chapters, we use ARFs generated utilizing xissimarfgen software(Ishisaki et al. 2007).

4.2. XRT35

Fig. 4.4.—: Left: Total effective area of the four XRT modules compared with that of XMM-Newton
and Chandra. Right: Vignetting of the XRT using the data of the Crab Nebula. The model curves were
calculated with the ray-tracing simulator. The left and right figures are taken from Mitsuda et al. (2007)
and Serlemitsos et al. (2007), respectively.

Fig. 4.5.—: Point-Spread Function (PSF) and Encircle-Energy Fraction (EEF) of the XRT (for XIS2) inthe focal plane. Both figures are taken from Serlemitsos et al. (2007).

36

4.3

CHAPTER 4. INSTRUMENTS

XISFour sensors of the XIS were developed by the collaboration of Kyoto University, Osaka
University, ISAS/JAXA, Massachusetts Institute of Technology (MIT), Rikkyo University, and
RIKEN (The Institute of Physical and Chemical Research). The XIS employs X-ray sensitive silicon charge-coupled devices (CCD), similar to that used in the ASCA SIS, Chandra ACIS, and
XMM-Newton EPIC. Coupled with the XRT, the XIS have imaging capability in the energy rangeof 0.2–12 keV. The photograph and schematic view of the XIS sensor (CCD+camera body) are
shown in figure 4.6.

4.3.1

Overview

Specifications of the CCD
The CCD chip has an imaging area of 1024×1024 pixels and two frame-store regions of
512×1024 pixels, as shown in figure 4.7. The pixel size of the imaging area is 24 µm × 24 µm,
while that of the frame-store regions is 21 µm × 13.5 µm. The CCD has four readout nodes (Node

Baﬄe

Hood

230

Electric
Valve

Door

Door Calibration
Source
CCD

Wall Calibration
Source70

Optical Blocking
Filter (OBF)
Cover Shield

Thermo-Electric
Cooler (TEC)

Heat Sink

Flexible Print

Fig. 4.6.—: The picture (left) and schematic view (right) of XIS sensor. Both pictures are taken from
Koyama et al. (2007).

37

4.3. XIS1024 pixels

1024 pixels







Charge Injection Gate








Imaging Area 
ABCD

No physical boundaries


bewteen the segments 

A, B, C, and D







Y

X

illuminated by 55Fe

Frame-store Region

256
pixelsRead-out Nodes

Fig. 4.7.—: Schematic view of CCD chip. The CCD consists of four segments (A, B, C, and D), each witha dedicated read-out node.

A–D) and each node reads the signals from 256 columns of the chip. Therefore the CCD is divided
to four segments (Segment A–D). Photoelectrons generated in the imaging area are simultaneouslytransfered to the frame-store regions, and read by the nodes in order. The time required to transfer
one pixel is ∼0.25 µs. The CCD chip is cooled to –90 ◦ C by the Thermo-Electric Cooler (TEC)
which is controlled by the Thermal Controller Electronics (TCE). For on-orbit calibration, eachsensor has 55 Fe calibration sources. Two sources are located on the side wall of the housing andilluminate two corners (top of the Segment A and D) of the imaging area. Since the CCD is sensitive to optical and UV light, the XIS have optical blocking filters (OBF), made of aluminum andpolyimide film, 20 mm above the imaging area (see the right panel of figure 4.6). The OBF has
a low transmission coefficient for optical light (≤ 5 × 10 −5 ), but is transparent to X-rays (≥80%
Table 4.2:: Design parameters and performance of the XIS (Koyama et al. 2007).

Field of view
Energy range
FormatEnergy resolution
Readout noiseTime resolution

180 × 180
0.2–12 keV
1024×1024 pixels∼130 eV (FWHM; @ 5.9 keV)
∼2.5 electrons
8 s (normal) / 7.8 ms (P-sum)

38

CHAPTER 4. INSTRUMENTS

above 0.7 keV).
The design parameters are summarized in table 4.2. Note that the XIS has two CCD clocking
modes: normal mode and parallel-sum (P-sum) mode. Normal mode has imaging capability withtime resolution of 8 sec. On the other hand, P-sum mode has good time resolution (7.8 msec), but
it has no imaging capability. During all observations described in this thesis, the XIS were operated
in the normal clocking mode.
Event Detection
When an X-ray photon is absorbed in a depletion layer of the CCD chip, it produces a charge
cloud typically smaller than the pixel size of the CCD. Thus the charge does not occupy more than
2×2 pixels. In contrast, a charged particle excites free electrons along its track, producing a larger(than 2×2 pixels) charge cloud. In order to distinguish X-ray events from charged particle events,
we use the Grade method (figure 4.8). An event is recognized when a pixel has a pulse-height (PH)between the lower and upper thresholds and the PH is a local maximum among adjacent pixels.
Then the surrounding 3×3 pixels are compared to the split threshold, and the event is classified
into the Grades given in figure 4.8. We regard the events of Grade 0, 2, 3, 4, and 6 as the X-ray
events. On the other hand, all of the events which split over the 2×2 pixels region are classified into
Grade 7 and regarded as background events, because they are usually formed by charged particle
events.

Fig. 4.8.—: Classification of the “Grade”. Grade 0, 2, 3, 4, and 6 are regarded as X-ray events. “PH”
indicates pulse-height which is proportional to the energy of incident X-ray.

4.3. XIS

39

Fig. 4.9.—: Schematic picture of the response function (the XIS response to a monochrome X-ray). It
consists of (1) main peak, (2) sub peak, (3) triangle component, (4) Si escape peak, (5) Si fluorescent peak,and (6) constant component. This figure is taken from Matsumoto et al. (2006)

4.3.2

Characteristics and Performances

Response Function
Even if incident X-rays are monochrome, the spectral shape is not a simple Gaussian but ismore complicated as is shown schematically in figure 4.9. It consists of (1) main peak, (2) sub
peak, (3) triangle component, (4) silicon escape peak, (5) silicon fluorescent peak, and (6) constantcomponent. In the energy range below the Si K-edge (1.839 keV), components (4) and (5) arenot seen. The physical interpretations these components are described in Matsumoto et al. (2006).
This spectral response to monochrome X-rays is called “response function”, which was calibrated
by our ground experiments. The response function is included in the response matrix files (RMFs),
which can be made using xisrmfgen software (Ishisaki et al. 2007).
Quantum Efficiency and Energy Resolution
The XIS consist of three Front-Illuminated (FI) CCDs and one Back-Illuminated (BI) CCD.
The left panel of figure 4.10 shows the quantum efficiency (QE) of the FI and BI, which weremodeled on the results of our ground calibration. Since the BI has no gate structure obstructing
incident X-rays on its surface, it has much higher QE below ∼4 keV than the FI. However, the QEof the BI for hard X-rays is lower than that of the FIs, because the depletion layer of the BI is

40

CHAPTER 4. INSTRUMENTS

Fig. 4.10.—: Left: Quantum efficiency of the FI CCD (solid line) and the BI CCD (dash line), which weremodeled based on the results of the ground calibration. Right: Energy resolution of the FI CCD (square)
and the BI CCD (circle), which were also the results of the ground calibration.

thinner than that of the FIs (∼ 75 µm and ∼ 45 µm for the FI and BI, respectively). The energy
resolution is comparable between the FI and BI, as shown in the right panel of figure 4.10. Bothachieve a pre-flight energy resolution of ∼130 eV (FHWM of the main peak component) at 5.9 keV.
Non X-Ray BackgroundFigure 4.11 shows non X-ray background (NXB) spectra of the XIS and EPIC on board
XMM-Newton, normalized by their solid angle of the field of view (FOV) and by their effective
area. It hence represents surface brightness of the background, and gives a measure of the S/Nratio for diffuse objects. The NXB spectra of the XIS and EPIC are respectively extracted fromthe night Earth data and “CLOSED” data (taken with the filter wheel in the closed position). The
background level of the EPIC-MOS is ∼5 times higher than that of the XIS-FI and is comparable
to that of the XIS-BI above ∼6 keV, where the NXB is dominant in the background spectra.
Although the background continuum level of the EPIC-pn is lower than that of EPIC-MOS, the
NXB spectrum of the EPIC-pn has very strong emission lines (Ni, Cu, and Zn) above 7 keV. The
intensity of Al-Kα line at 1.49 keV in the XIS spectra is also much weaker than that in the EPIC
spectra. Therefore, it is proved that the XIS-FI has the lowest background level (the highest S/Nratio) for diffuse objects. Furthermore, the EPIC NXB frequently shows “background flare”, andthe NXB count rate varies by a factor of about 100. Since the orbit altitude of Suzaku is much
lower than XMM-Newton and Chandra, Suzaku is less influenced by solar flares, and hence thebackground level of the XIS is very stable. Therefore, the XIS is very sensitive to diffuse dimsources especially in the hard X-ray band. The details of the XIS background properties can befound in Yamaguchi et al. (2006) and Tawa et al. (2008).

4.3. XIS

41

Fig. 4.11.—: Comparison of the NXB spectra normalized by the solid angle of the FOV and by the effective
area between the Suzaku XIS and the XMM-Newton EPIC. The black, red, green, and blue represent the
XIS-FI, XIS-BI, EPIC-MOS, and EPIC-pn, respectively.

Degradation of Performances
Since the launch, the energy resolution of the XIS is gradually degrading due to the radiation
damage (Koyama et al. 2007). Figure 4.12 shows the on-orbit time histories of the energy resolution
at 5.9 keV. The apparent linear trend (∼50 eV/yr) is due to charge loss caused by increasing of
charge transfer inefficiency (CTI). The details of the CTI are described in Nakajima et al. (2008).
The calibration results on the degradation of the energy resolution are reflected to the RMF (Ishisaki
et al. 2007).
The low energy QE had been degraded by the contamination (due to out-gassing from thesatellite) lied on the OBF (Koyama et al. 2007). The contamination rate is unexpectedly high, andthe rate is different from sensor to sensor. Moreover, the thickness of the contamination varies withposition on the OBF. We therefore need to include additional, time-varying low energy absorption.
This is given as a function of both the observation date after the XIS door-open, and of detector
coordinates (specifying the position on the OBF). The time history of the contamination is shownin the left panel of figure 4.13, which is measured by observing the SNR E 0102–72 and the isolated
neutron star RX J1856.5–3754. Although the chemical composition for the contamination material
is not conclusively determined, it is assumed that the contaminant contains C and O with thenumber ratio C/O ∼ 6, from the overall spectral shape in the low energy absorption for the X-ray
sources and the best guess for the out-gassing source in Suzaku. The off-axis absorption is determined utilizing the diffuse X-rays from the day Earth rim. The radial profile of the contamination isshown in the right panel of figure 4.13. These calibration results are reflected to the ARF (Ishisaki

42

CHAPTER 4. INSTRUMENTS

et al. 2007). Note that, however, the calibration accuracy is relatively poor in large off-axis angle(>∼ 8 arcmin) region. Furthermore, the calibration for the FI is less accurate than that for the BI(Koyama et al. 2007).

Fig. 4.12.—: Time history of the energy resolution (FWHM) measured with Mn-Kα line (5.9 keV) fromthe 55 Fe calibration sources. Top and bottom indicate the XIS0 and XIS1, respectively. This figure is takenfrom Koyama et al. (2007).

Fig. 4.13.—: Left: The time history of the contamination measured at the center of the OBF. Right:
The radial profile of the contamination of the XIS-BI. Circle and cross represent the data for one month andfive months after the door-open, respectively. Both figures are taken from Koyama et al. (2007).

4.4. HXD

4.4

43

HXD

The HXD was developed by the collaboration of the University of Tokyo, Aoyama Gakuin
University, Hiroshima University, ISAS/JAXA, Kanazawa University, Saitama University, Stanford
Linear Accelerator Center (SLAC), and RIKEN. It is a non-imaging, collimated hard X-ray scintillating instrument sensitive in the 10–600 keV band. Since the main goal of this thesis is the
spectroscopy in the energy band below 10 keV, we do not use the data obtained with the HXD. In
this section, therefore, we show only basic characteristics of this instrument.
The HXD sensor is a compound-eye detector instrument (figure 4.14), consisting of 16 main
detectors (arranged as a 4×4 array) and the surrounding 20 crystal scintillators for active shielding.
Each unit actually consists of two types of detectors: a GSO/BGO phoswich counter, and 2mmthick PIN silicon diodes located inside the well, but in front of the GSO. The PIN diodes aremainly sensitive below ∼60 keV, while the GSO/BGO scintillator is sensitive above ∼30 keV. The
characteristics of the HXD are summarized in table 4.3. The outer anti-coincidence scintillators
can be used as a wide-field hard X-ray detector, which is referred as the Wide-band All-sky Monitor(WAM). This can be used to detect bright X-ray transients, γ-ray bursts, and solar flares.
Figure 4.15 shows the background counting rate as a function of energy for the 10–500 keV
region. The background rate of the HXD on board Suzaku is the lowest among the existing missions
for most X-ray energies. The sensitivity in this energy region is essentially limited by the accuracyof background estimation. At present, the background model for the HXD-PIN can be reproducedwith an accuracy of 5% of the background level (Kokubun et al. 2007).

44

CHAPTER 4. INSTRUMENTS

Fig. 4.14.—: Schematic view of HXD. This picture is taken from Takahashi et al. (2007).

Table 4.3:: Performance of the HXD (Takahashi et al. 2007).

Field of view
Bandpass
Energy resolution (FWHM)
Effective area
Time resolution0
0 <4.◦ 5 × 4.◦ 5 (>∼ 100 keV) / 34 × 34 (∼ 100 keV)
10–70 keV (PIN) / 40–600 keV (GSO)−1/2∼3.0 keV (PIN) / 7.6E MeV % (GSO)∼160 cm2 (@ 20 keV) / ∼260 cm2 (@ 100 keV)
61 µs or 31 µs

4.4. HXD

45

Fig. 4.15.—: In-orbit detector background of the PIN/GSO compared with those of the RXTE and
BeppoSAX. The dotted lies indicate 1 Crab, 100 mCrab, and 10 mCrab intensities. This picture is takenfrom Kokubun et al. (2007).

Chapter 5

SN 1006
SN 1006 is famous as the prototype of non-thermal SNRs, owing to the first discovery ofsynchrotron X-rays from the shell of the remnant (Koyama et al. 1995). However, its thermal
emission from shock heated plasma also has a very important aspect. As explained below, theevolutional state of SN 1006 is probably younger than any other historical Galactic SNRs, due tothe extremely low ISM density around SN 1006. Thermal X-ray emission from a young SNR isdominated by the radiation from ejecta, as described in chapter 3. Therefore, a young SNR iswell-suited for investigation the abundances of synthesized elements and thermal evolution of theejecta in a very early phase of SNR evolution. For this reason, we choose the ‘youngest’ SNR,
SN 1006, as the main target of this thesis. We summarize, in this chapter, the results of Suzaku
observation and analysis on SN 1006 (mainly following Yamaguchi et al. 2008a). Unless otherwise
stated, figures in this chapter are taken from Yamaguchi et al. (2008a).

5.1

Previous Results

SN 1006 (G327.6+14.6) is the remnant of the historical SN explosion on May 1, A.D. 1006,
which was observed and recorded in all over the world. The SN was very bright, was comparedwith the Moon, and remained visible for perhaps several years. The records survive from Japan,
China, Korea, Iraq, Syria, Egypt, Italy, Switzerland, France, and possibly Spain. Its remnant was
first cataloged by Bolton et al. (1964) as a radio source and identified (as the remnant of SN 1006)
by Gardner & Milne (1965). The energy index in radio band is measured by Allen et al. (2001) tobe α = 0.57. The distance to SN 1006 is determined to be 2.2 kpc by optical observations (Winkler
et al. 2003); we assume this value here.
In X-ray band, SN 1006 was first observed by Einstein satellite (Becker et al. 1980). Using
ASCA, Koyama et al. (1995) discovered synchrotron X-rays from shells of the SNR, that indicatesthe existence of high energy electrons up to ∼100 TeV. The Chandra results of the small-scale
structure have been successfully explained by a diffusive shock acceleration model (DSA); the first47

48

CHAPTER 5. SN 1006

Fig. 5.1.—: XMM-Newton image of SN 1006, taken from Rothenflug et al. (2004). Red, green, and blue
are 0.5–0.8 keV, 0.8–2.0 keV, and 2.0–4.5 keV, respectively. The areas dominated by non-thermal emission
appear white. North is up and east is to the left.

order Fermi mechanism is working at the shock front of SNRs (Bamba et al. 2003; Yamazaki
et al. 2004). Rothenflug et al. (2004) observed whole region of the remnant with XMM-Newton(figure 5.1), and concluded that the anisotropic morphology of the non-thermal emission is dueto the ambient magnetic field running southwest to northeast along the Galactic plane. Although
the detection of inverse Compton TeV γ-rays with CANGAROO was reported by Tanimori etal. (1998), this result had been denied by HESS observatory (Ksenofontov et al. 2005).
Based on the historical record, SN 1006 is widely regarded to be one of the Galactic type Ia
SNRs (Schaefer 1996). As explained in chapter 2, The current SN models predict that Fe production
in Type Ia SNe is far larger than that of core-collapse SNe (e.g., Nomoto et al. 1984; Iwamoto etal. 1999). Therefore, measuring the Fe abundance is one of the essential clues to classify the SN
type of any specific remnant. The K-shell X-ray emission lines from ionized Fe offer the most direct
information for abundance determination. In fact, strong Fe-Kα lines had been observed from
Tycho’s SNR (e.g., Hwang et al. 1998; Decourchelle et al. 2001) and many of the other Type Ia
SNRs.
In SN 1006, Vink et al. (2000b) suggested the possible presence of an Fe-K emission line at
6.3±0.2 keV from the BeppoSAX MECS spectrum. However, the detection was not clear, because

5.1. PREVIOUS RESULTS

49

the energy resolution of the MECS (a gas scintillation counter) was relatively poor (FWHM∼500 eV@6 keV). Other recent X-ray missions, Chandra and XMM-Newton, have not succeeded in detecting
an Fe-K line. Cold Fe in the interior of SN 1006 is known to exist based on ultraviolet (UV)
absorption studies (e.g., Wu et al. 1993; Hamilton et al. 1997; Winkler et al. 2005). Figure 5.2
indicates the positions of background UV sources (with circle mark) on the Hα image of SN 1006.
As shown in figure 5.3, blue and red-shifted Fe II absorption lines were detected in the spectra of
the background stars, labeled ‘Q2’ and ‘SM star’ in figure 5.2, and interpreted as being due tounshocked Fe in the interior of the remnant. These results show that portions of the Fe-rich ejecta
are still expanding freely, and have not yet been overrun by the reverse shock. However the amountof Fe inferred in these studies (< 0.16 M ) is much less than the amount (0.6–0.8 M ) predicted
to be produced in the thermonuclear disintegration of a ∼1.4 M white dwarf.
In the early phases of supernova remnant evolution, even the shock-heated plasma is far fromthermal equilibrium in terms of either the ionization or particle (electron and ion) temperatures, as
explained in chapter 3. The thermal X-ray spectrum of SN 1006 suggests an ionization timescaleof τ = ne t ∼ 2 × 109 cm−3 s (Vink et al. 2000b), which is far from the full ionization equilibrium,and lower than nearly all other Galactic SNRs. Moreover, the electron temperature just behindthe northwest (NW) rim of the remnant has been estimated to be kT e ∼ 0.7 keV (Long et al. 2003)or ∼ 1.5 keV (Vink et al. 2003), which are much lower than the ion temperature expected from theshock velocity of the NW rim (∼2980 km s −1 : Ghavamian et al. 2002). Therefore, non-equilibrationof ion and electron temperatures is present in SN 1006 as well.
The extreme non-equilibrium state and high shock velocity are likely due to the low densityof the ambient gas, because SN 1006 (b = +14.6) is further from the Galactic plane than other
Galactic SNRs, such as Cas A (b = −2.1) and Tycho’s SNR (b = +1.4). Although Kepler’s SNR(b = 6.8) is nearly as far above the plane as SN1006 (when their respective distances are included),
Kepler appears to be evolving into a dense circumstellar medium. Therefore, the evolutionary stateof SN 1006 may be the lowest of these youngest Galactic SNRs, even though the real age is ∼3
times older than the others.

50

CHAPTER 5. SN 1006

Fig. 5.2.—: Background UV sources are marked on Hα image of SN 1006, which shows the faint nonradiative filaments that delineate the primary shock around most of the shell (Winkler et al. 2003). The
cross indicates the geometric center of the optical shell. This image is taken from Winkler et al. (2005).

Fig. 5.3.—: Top: Normalized (to the continuum level) UV spectra of Q2 (labeled in figure 5.2) in the
wavelength region of prominent Fe II lines. It shows clear evidence for broad absorption features. Middle:
Same to top but the spectra of SM star. Similar absorption profile is seen. Bottom: the Q1 spectrum. It
shows no evidence for broad Fe II lines; the narrow lines are probably of interstellar origin. This figure istaken from Winkler et al. (2005).

51

5.2. SUZAKU OBSERVATION
Table 5.1:: Log of Suzaku observations on SN 1006 and North Ecliptic Pole.

RegionNE
SW

Obs. ID
Date
100019020
2005/09/09
100019030
2005/09/10
100019050
2006/01/26
SE
500016010
2006/01/30
NW
500017010
2006/01/31∗
NEP
500026010
2006/02/10∗ North Ecliptic Pole, used as background data.

5.2

Exposure (ks)
43
31
31
52
53
82

Suzaku Observation

Four pointings were made of SN 1006 with Suzaku during the performance verification (PV)phase. The observation ID, date, and exposure time are listed in table 5.1. During all of theseobservations, the XIS were operated in the normal full-frame clocking mode with the editing modeof 3 × 3 or 5 × 5. We employed the cleaned revision 0.7 data, and used the HEADAS software
version 6.0.4 and XSPEC version 11.3.2 for the data reduction and analysis. The X-ray data taken
during low cut-off rigidity (≤ 6 GV) were excluded so as to minimize any possible uncertainty inthe non X-ray background (NXB). The total effective exposure times after these screening were
∼50 ks, for each of the pointings. The response matrix files (RMF) and ancillary response files(ARF) were made using xisrmfgen and xissimarfgen (Ishisaki et al. 2007) version 2006-10-17.

5.3

Overall Structure

5.3.1

Thin-Thermal Plasma Map

The field of view (FOV) of the XIS is ∼ 18 0 × 180 , and hence a four-point mosaic can completely
cover SN 1006 (about 300 diameter). Since we intended to study the thermal plasma, we searched
for the region of the remnant with the most prominent thermal emission. We hence made a mosaic
image in the O VII Kα line band (0.55–0.6 keV), because this line is a major component of thethermal plasma in SN 1006. The result is shown in figure 5.4. We can see bright emission inthis band in the northeast (NE) and southwest (SW) regions of the remnant. However, these aredominated by non-thermal emission (Koyama et al. 1995), from which it is difficult to extract
thermal spectra. Other regions are dominated by thermal emission. In particular, as shown infigure 5.4, the southeast quadrant (here SN 1006 SE) is brighter than the NW. Therefore, ourstudy concentrated on this region. The spectral results of non-thermal rims (NE and SW) are
reported by Bamba et al. (2008). The FOV of the XIS on SN 1006 SE is outlined by the black solid
square in figure 5.4.

52

CHAPTER 5. SN 1006

-41:40:00

50:00-42:00:00

10:00
04:00

15:03:00

02:00

Fig. 5.4.—: XIS-BI mosaic image of SN 1006 in the 0.55–0.6 keV band. Exposure and vignetting effects
are corrected. The coordinates (R.A. and Dec) are referred to epoch J2000.0. The black square shows the
FOV of the XIS on the SE region of SN 1006 where this chapter is concentrated. The blue squares show theFOVs on the other three quadrants (NE, NW, and SW).

Fig. 5.5.—: Background-subtracted XIS spectra extracted from the whole SE quadrant (SN 1006 SE). The
black and red points represent the FI and BI spectra, respectively.

5.3. OVERALL STRUCTURE

5.3.2

53

Spectrum of the SN 1006 SE Region

We extracted the spectrum of the entire SN 1006 SE region, excluding the two corners in the
FOV that contain calibration source emission. For the background, we used the North Ecliptic Pole
(NEP) data. Although we have been monitoring and correcting the increase of the Charge-Transfer
Inefficiency (CTI) for the XIS (section 4.3; Koyama et al. 2007), recovery of the energy resolution
cannot be made. However, we can ignore the difference of the spectral resolution between the
SN 1006 SE and NEP observations, because the latter observation was made only ten days after
the former (see table 5.1). To minimize the uncertainty due to the background subtraction, inparticular that of NXB, we applied the same data-screening criteria to both the SN 1006 SE and
NEP observations, and took background spectra from the same detector coordinates as the source
regions after excluding point-like sources detected in the NEP data. The background-subtracted
spectra are shown in figure 5.5. Since the data from the three FIs are nearly identical, we merged
those individual spectrum to improve the photon statistics.
As shown in figure 5.5, we found clear K-shell (Kα) lines from Ar, Ca, and Fe, for the firsttime. With a power-law plus Gaussian-line fit, we determined the line center energy of the Fe-Kα
to be ∼6.43 keV. This energy constrains the Fe ionization state to be approximately Ar-like (Fe IX)
and Ne-like (Fe XVII) . The lack of Fe-L emissions (which are generally seen in 0.7–1.5 keV band)also indicates that Fe is below the Fe XVII state.

5.3.3

Iron Line Map

We show in figure 5.6 an image in a relatively narrow band (6.33–6.53 keV) that contains the
Fe-Kα line. Data from the observation of the NW quadrant are included. In the other two regions,
NE and SW, the non-thermal continuum flux at 6.33–6.53 keV is very high and therefore we cannot
extract reliable information on the Fe-Kα flux. We do not show the images of these regions.
We can see that the Fe-Kα flux is enhanced at the southern part of the remnant (outlined
in red with a ellipse), except for the NE and SW quadrants where the non-thermal emission isdominant. The mean surface brightness at 6.33–6.55 keV within the elliptical region is 8.5 (± 0.5) ×
10−9 photons cm−2 s−1 arcmin−2 , while that outside it (only in the SE and NW quadrants) is
4.6 (± 0.3) × 10−9 photons cm−2 s−1 arcmin−2 . In order to study the thin-thermal spectrum withthe best S/N ratio for Fe-K line, we extracted the X-ray spectrum from within the elliptical region,
excluding the corner of the calibration sources. The background subtraction was made in the same
way as that of the full-field spectrum. The results are given in figure 5.11. Hereafter, all detailed
analyses are made using this spectrum.

5.3.4

Energy and Width of the Emission Lines

In order to study the line features, we fitted the spectra extracted from the elliptical region
with a phenomenological model; a power-law for the continuum and Gaussians for the emission

54

CHAPTER 5. SN 1006

-41:40:00

50:00-42:00:00

10:00
04:00

15:03:00

02:00

Fig. 5.6.—: XIS intensity map at the Fe-Kα line energies (6.33–6.53 keV). Exposure and vignetting effects
are corrected. The data from the three FIs are combined. Two corners of the calibration sources are removed.
The two black squares indicate FOVs of the XIS (on SN 1006 SE and NW). Since the non-thermal continuum
flux at 6.33–6.53 keV is very high in the regions of NE and SW, we cannot extract reliable information of
the Fe-Kα flux, and hence we do not show the images of these regions. The red ellipse shows the region
where we extracted the spectra for a detailed analysis.

Table 5.2:: The center energies and widths of the emission lines.

Line
Center energy∗ (eV) Width† (eV)Mg-Kα1338 (1337–1340)
< 5.4
Si-Kα
1815 (1813–1816)
40 (38–42)
S-Kα
2361 (2355–2365)
60 (54–65)
Ar-Kα3010 (2991–3023)
< 50
Ca-Kα
3692 (3668–3724)
< 57
Fe-Kα
6430 (6409–6453)< 60∗ Errors (statistical only) are given in theparentheses (see text).
† One standard deviation (1σ).

5.4. SPECTRAL STRUCTURE IN THE NARROW BANDS

55

lines. The best-fit central energies and widths for the emission lines are shown in table 5.2. Since
the absolute energy calibration error is ±0.2% above 1 keV (Koyama et al. 2007), the relevant
elements are uniquely identified from the best-fit central energy. These are also listed in table 5.2.
Except for Mg, the central energies are significantly lower than that of the respective He-like Kαlines (see table 3.1). We also note that the widths of the Gaussians identified with Si-Kα and S-Kα
are significantly broader than the instrumental energy resolution at those energies. Figure 5.9awould help to compare the widths of the data with those of narrow lines. In subsection 5.4.3, we
discuss this matter in detail.

5.4

Spectral Structure in the Narrow Bands

In order to study the plasma characteristics in relation to the various elements, we at firstdivided the spectrum into three representative energy bands: the 0.4–1.1 keV band for the O andNe (light elements) lines, the 1.2–2.8 keV band for the Mg, Si, and S (medium elements) lines,
and the 5–10 keV band for the Fe (heavy element) line. With the fitting to these individual band
spectra, we constrained the plasmas including light, medium, and heavy elements separately.

5.4.1

The 5–10 keV Band Spectra

We fitted the 5–10 keV spectrum with a solar abundance (Anders and Grevese 1989) VNEI
model. Since the NEIvers 2.0 plasma code does not include K-shell emission lines for ions below
the He-like state, we reverted to the NEIvers 1.1 code. If we fix the Fe abundance to be solar,

Fig. 5.7.—: XIS spectra in the 5–10 keV band with the VNEI model. The black and red data points
represent the FI and BI spectra, respectively. In the BI spectrum, the energy band above 8 keV is ignored.

56

CHAPTER 5. SN 1006

then the model cannot reproduce the Fe-Kα profile or flux, although the χ 2 /d.o.f. of 25/28 is
acceptable. Allowing the Fe abundance to be free results in a greatly improved best-fit χ 2 /d.o.f. of
10/27. The best-fit kTe , ne t, and Fe abundance are 3.8 (1.7–26) keV, 6.1 (2.2–10) ×10 9 cm−3 s,and 6.5 (>1.9) solar, respectively. The Fe abundance is significantly enhanced relative to solar,
even if the continuum flux in this band is assumed to be composed of pure thermal emission. The
best-fit spectrum is shown in figure 5.7.

5.4.2

The 0.4–1.1 keV Band Spectra

Figure 5.8 shows the 0.4–1.1 keV band spectra. This energy band includes Kα lines of thelight elements, with dominant emission from the oxygen K-shell line series. The spectra were fittedwith a VNEI model, allowing the abundances of C, N, O, and Ne to be free parameters. Also, we
let the abundances of Ca and Fe to be free, because the L-shell lines of these elements fall into
this energy band (0.4–1.1 keV). Interstellar absorption was fixed to a hydrogen column density of
6.8×1020 cm−2 (Dubner et al. 2002). In the initial fits, we found a significant inconsistency between
the FI and BI spectra near the energy of the O-edge (0.54 keV). This may be due to incomplete
calibration information for the contamination layer on the optical blocking filter (OBF) of the XIS.
Since the calibration of the OBF contamination for the BI is more accurate than that for the FI (seesection 4.3), we decided to retain the BI data across this energy band and ignore the 0.5–0.63 keVband in the FI spectrum. The results are shown in figure 5.8a, with the best-fit temperature andionization parameter of kTe = 0.58 (0.56–0.59) keV and ne t = 6.7 (6.6–6.8) ×109 cm−3 s.
In figure 5.8a, we find an apparent disagreement between the data and model in the ∼0.7–
0.85 keV energy band. Usually, this energy band is dominated by the iron L-shell line transitionsof 3s − 2p (∼730 eV) and 3d − 2p (∼830 eV) from Fe XVII. Although the fluxes of these twolines in NEI models are nearly equal over a wide range of plasma conditions, extrapolation to theextremely low ionization states that we see here in SN 1006 is quite uncertain. In particular, forthe young Type Ia SNR E0509–67.5, Warren and Hughes (2004) found that the Fe L-shell emission
was dominated by an emission from a line feature near 0.73 keV, which was not sufficiently strong
in the spectral models and had to be included as a separate Gaussian component. This 0.73 keVdata excess (compared from model) is confirmed in a Suzaku spectrum of E0509–67.5 (chapter 6).
Based on the strength of other Fe-L emission lines and the weakness of O emission in E0509–67.5
the additional line feature at 0.73 keV was confidently associated with Fe-L emission.
Although we cannot eliminate the possibility that the residual comes from faint Fe-L emission,
here we consider a different explanation. The best-fit spectral model of the oxygen lines are shownin figure 5.8b, in which the black and red solid lines are K-shell transitions from O VII and O VIII,
respectively. In most astrophysical plasmas, the best-fit temperature of ∼0.6 keV gives weak K-shelllines of O VII compared to those of O VIII. On the other hand, in lower temperature plasmas (e.g.,kTe ∼ 0.1 − 0.2 keV) where the O VII Kα line is dominant, the line fluxes decrease rapidly along theK-shell transition series (Kα,β,γ,δ,,ζ, etc.). Therefore, K-shell lines in the higher transitions canbe safely ignored, and hence conventional NEI codes do not include O VII K-shell transition lines

5.4. SPECTRAL STRUCTURE IN THE NARROW BANDS

57

Fig. 5.8.—: (a) XIS spectra in the 0.4–1.1 keV band fitted with a VNEI model. The black and red data
points represent the FI and BI spectra, respectively. (b) Spectrum in photon space corresponding to themodel used in (a), but only showing O lines. It is binned every 2 eV. The black and red lines represent
O VII series and O VII series, respectively. The dotted lines are additional transitions of Kδ (714 eV), K
(723 eV), and Kζ (730 eV) of O VII added as separate Gaussians (see text). (c) Same spectra as (a), but
for fits including Gaussians representing Kδ (714 eV), K (723 eV), and Kζ (730 eV) of O VII (see text). The
residuals in the energy band 0.7–0.85 keV seen in (a) are largely removed.

58

CHAPTER 5. SN 1006

higher than Kδ. This is the reason that no oxygen line in the ∼0.7–0.85 keV band is present in
the best-fit model (figure 5.8b). In SN 1006, however, the ionization timescale is low, while kT e
is moderate, so the fluxes of higher level K-shell transitions from O VII may be relatively strong
and cannot be ignored. To account for them we added Gaussians at 714 eV, 723 eV, and 730 eV
to represent the O VII Kδ, K, and Kζ lines, respectively. At a plasma temperature of 0.6 keV,
the flux ratio of O VIII Ly to Lyδ predicted by the NEIvers 1.1 code is ∼0.5 (see figure 5.8b).
In the absence of a detailed calculation we merely assumed that the flux ratio of the O VII lines
follow the same pattern, namely K/Kδ = Kζ/K = 0.5. The best-fit results with these additionallines, see figure 5.8c, are a reasonably good fit. The additional artificial lines are also shown as
the dotted lines in figure 5.8b. The temperature and ionization parameter are almost the same as
those obtained without the additional higher K-shell transitions of O VII.
The best-fit intensity ratio of O VII Kβ to O VIII Lyα is ∼1.4, which is consistent with the resultof the XMM-Newton/RGS observation of the X-ray knot localized at the NW rim of SN 1006 (Vink
et al. 2003; O VII Kβ/O VIII Lyα ∼ 1.6). This agreement may justify our approach to include theK-lines manually. Also, this suggests that the O VII K-shell line series are key spectral components
for all over the SN 1006 plasma.

5.4.3

The 1.2–2.8 keV Band Spectra

This band includes K-shell lines of intermediate-mass elements (Mg, Si, and S). As shown intable 5.2, we can see significant broadening of the Si and S Kα lines. Therefore, the 1.2–2.8 keV band
spectra cannot be reproduced with a single-component plasma model. In fact, we found significant
disagreement between the data and model around the Si-Kα and S-Kα lines (figure 5.9a) in the
one-component VNEI fits. The lack of data in the 1.83–1.85 keV band is due to current calibration
errors of the XIS; there is a small gap in energy (∼10 eV) at the Si K-edge, which is not implemented
in the current response function. The apparent line broadening of the Si Kα line, however, is muchlarger (∼40 eV) than this gap. Therefore, the disagreement is not due to a calibration error, but
is real.
If the intrinsic broadening of the Si-Kα line (∼ 40 eV; see table 5.2) is due to thermal Doppler
broadening, the ion (silicon) temperature must be ∼ 13 MeV, which requires a shock velocity of∼ 1.5 × 104 km s−1 , (derived from equation 3.36). Ghavamian et al. (2002) determined the shock
velocity (from measurements of the Hα line width) to be ∼2900 km s −1 in the NW portion of
SN 1006. This region is where the blast wave is interacting with significant amounts of interstellar
matter, compared to other parts of the rim. In fact, the average size of the remnant (∼15. 0 5 radius)
and the well-known age and distance allow us to determine a mean expansion speed for SN 1006
to be ∼104 km s−1 . The true current blast wave speed in SN 1006 is almost surely bracketed bythese two values. Furthermore, the reverse shock, which is likely to be the heating source for the Si
that we see, typically moves into the ejecta at only a fraction of the blast wave speed. Therefore,
we consider it unlikely that the line broadening is due to thermal Doppler broadening.
Another possible explanation proposes that the emission consists of several thermal plasma

59

5.4. SPECTRAL STRUCTURE IN THE NARROW BANDS

components. We fitted the spectra with 2-VNEI models, in which we allowed for different kT evalues, but the same ne t between components. However, no combination of parameters would fit
the data, even if we let the abundances of Mg, Si, and S be free parameters. Not only 2-VNEI,
but also 2-VPSHOCK models failed to fit the data when the two components were forced to havethe same ionization timescale. All of the models that we tried, but rejected, are summarized intable 5.3. The addition of a power-law component to these models also did not help to improve thefittings.
Next we tried two-component VNEI models with the same electron temperature, but different
ne t values. The best-fit reduced χ2 /d.o.f. was greatly improved to 401/346 (see table 5.3, forcomparison). Models with different electron temperature and ionization timescale (2-kT e and 2ne t) gave no significant improvement of the reduced χ 2 . We thus conclude that a 2-component
VNEI model with different ne t values is necessary to fit the medium element plasma band. The
best-fit parameters and spectra are given in table 5.4 and figure 5.9b, respectively.
Table 5.3:: Best-fit parameters and χ2 values of the spectral fitting in the 1.2–2.8 keV band with the various
models.

Model
1-VNEI
2-VNEI
1-VPSHOCK
2-VPSHOCK

kTe1
(keV)
1.5
0.18
1.6
0.36

kTe2
(keV)
–
20
–
12

ne t
(cm−3 s)4.5×109
8.4×109
8.3×109
1.0×1010

χ2 /d.o.f.
848/351706/346726/351597/346

Table 5.4:: Best-fit parameters of the spectral fitting in the 1.2–2.8 keV band with 2-component VNEI
model with different ne t values.

Parameter
NH (cm−2 )kTe (keV)
Mg
SiS
ne t (cm−3 s)nH ne V (cm−3 )χ2 /d.o.f.

Component 1
Component 2
20
6.8×10 (fixed)
1.1 (1.0–1.2)
4.2 (3.5–5.2)
6.1 (3.5–13)
19 (15–24)
15 (13–24)
24 (19–34)
23 (17–31)
1.3 (1.1–1.7) ×1010 7.9 (6.2–9.8) ×108
2.8 (2.3–3.4) ×1055 1.9 (1.6–2.3) ×1056
401/346 = 1.16

60

CHAPTER 5. SN 1006

Fig. 5.9.—: (a) XIS spectra in the 1.2–2.8 keV band fitted with the one-component VNEI model. The
black and red represent the FI and BI, respectively. (b) Same to (a), but fitted with the two-component
VNEI.

61

5.5. MODEL FIT OF THE FULL BAND SPECTRA

5.5

Model Fit of the Full Band Spectra

In section 5.4, we discuss how we separately derived spectral parameters for the three-energy
bands that represent typical plasma conditions for the main light elements (O and Ne), medium
elements (Mg, Si, and S), and heavy element (Fe). A summary is given in table 5.5. We used
these parameters as the initial values to search for plasma parameters in the full energy band of
0.3–10 keV; the resulting overall best-fit parameters must be consistent with the data over theentire XIS energy band.
Based on table 5.5, we assume that the model that describes the full spectral range must consistof, at least, three plasma components: (1) high-kT e (∼1 keV) with high-ne t (∼ 1010 cm−3 s), (2)similarly high-kTe with low-ne t (∼ 109 cm−3 s), and (3) low-kTe (∼ 0.6 keV) with medium-ne t(∼ 7 × 109 cm−3 s). First, we consider whether or not the full band spectra can be well reproducedwith only these components. We also consider a broader range of kT e and ne t values for the Fe-K
emission.
As shown in table 5.4, components (1) and (2) are significantly enhanced in Mg, Si, and S(relative to solar). In the present spectral fits, we made the assumption that these components
are composed purely of metals without any admixture of hydrogen or helium. This assumption
arose because hydrogen and helium emit only continuum emission in the X-ray band. Given thecomplexity of our spectral model, it is just not possible with the XIS data to obtain reliable
estimates for the level of continuum emission from the three plasma components separately. As afurther assumption, we fixed the abundance of component (3) to the solar value. (Anticipating theresults given below, we found that most of the continuum emission above 1 keV actually comesfrom a hard power-law like spectral component.) Operationally, we fixed the oxygen abundance inboth the pure-metal components (i.e., nos. 1 and 2) to a large value (1 × 10 4 ), and fitted for theabundances of other elements relative to oxygen. The C and Ni abundances were assumed to be
the same as Ne and Fe, respectively. The emission measure (EM ) in the plasma code is defined(in XSPEC) as nH ne V , even if the plasma is dominated by heavy elements. However, the oxygen
density can be calculated as nO = (8.51 × 10−4 nH )×104 , where the numerical value is the solar
abundance of oxygen from Anders and Grevesse (1989). Details of this method can be found in
Vink et al. (1996).

Table 5.5:: Plasma components determined from the narrow band spectra.

kTe (keV)ne t (cm−3 s)
0.58 (0.56–0.59) 6.7 (6.6–6.8)×10 9
1.1 (1.0–1.2)
1.3 (1.1–1.7)×10 10
1.1 (1.0–1.2)
7.9 (6.2–9.8)×108
5.0–10Fe
3.8 (1.7–26)
6.1 (2.2–10)×10 9∗ The plasma component identifications described in text.

Band (keV)
0.4–1.1
1.2–2.8

Major elements
O, Ne
Mg, Si, S

Component∗
(3)
(1)
(2)

62

CHAPTER 5. SN 1006

Fig. 5.10.—: XIS spectra in the 0.33–10 keV band fitted with the three-component plasma model (seetext). The black and red points represent the FI and BI data, respectively. The model fails to describe thecontinuum emission above ∼4 keV.

Since the NEIvers 1.1 spectral model does not include K-shell emission lines from Ar andlow ionization states of Ca, we added Gaussians at the energies of their Kα transitions: 3.02 keV(Ar) and 3.69 keV (Ca) (see table 5.2), For component (1) we additionally freely varied the Caabundance, which effectively allowed for the Ca L-shell lines to contribute at low energies (< 0.5
keV). For component (2), we tied the Ca abundance to that of Fe, which was allowed to vary freely.
We also added Gaussians for the upper-level oxygen transitions at 0.714, 0.723, and 0.730 keV withthe same reason as described in subsection 5.4.2.
We fitted the FI data in the energy band from 0.33 to 10 keV (0.33 to 8 keV for the BI data).
We ignored the band below the C-edge at 0.28 keV, because of the large uncertainty in calibration
due to contamination on the OBF (see section 4.3). Energies of 0.5–0.63 keV (only for FI) and
1.83–1.85 keV were also ignored for the same reasons as noted in subsections 5.4.2 and 5.4.3. Since
the absolute gain of the XIS has an uncertainty of ±5 eV (Koyama et al. 2007), we allowed for asmall offset in the photon-energy to the pulse-height gain relationship.
With this model and assumptions, the best-fit χ 2 /d.o.f. obtained was 1081/833. The fitted
spectra are shown in figure 5.10. All line features are well fitted, even Fe-K, suggesting that
including additional thermal components would not be justified. However, there is a systematic
data excess in the continuum at energies >∼ 4 keV. We include a power-law component to this model,
since we cannot reject the possibility that there is a non-thermal component here that is similar tothe bright NE and SW rims (albeit much less intense). The fit with the additional power-law wassignificantly improved with χ2 /d.o.f. = 996/831. The best-fit parameters are given in table 5.6. As
already noted, the gain may have an uncertainty of ±5 eV. The fitting process requires a gain offset
of 4.8 eV for both FI and BI, within the allowable range. The best-fit model in the full energy
band is shown in figure 5.11.

5.5. MODEL FIT OF THE FULL BAND SPECTRA

Table 5.6:: Best-fit parameters of the broadband spectral fitting.
Component
Absorption
VNEI 1 (Ejecta 1)

ParameterValue
NH (cm )6.8×1020 (fixed)kTe (keV)
1.2 (1.1–1.4)
Abundance (104 solar)C
0.19∗ (0.08–0.30)
N
0 (fixed)O
1.0 (fixed)Ne
0.19∗ (0.08–0.30)
Mg
4.1 (3.7–4.7)
Si
17 (15–20)
S
23 (20–27)
Ca
11 (9.6–13)
Fe
0.68† (0.62–0.76)
Ni0.68† (0.62–0.76)−3
ne t (cm s)
1.4 (1.2–1.6) ×1010−3nH ne V (cm )2.7 (2.4–3.0) ×1051
VNEI 2 (Ejecta 2) kTe (keV)
1.9 (1.5–2.6)
Abundance (104 solar)C0.33‡ (0.23–0.45)
N
0 (fixed)O
1.0 (fixed)
Ne
0.33‡ (0.23–0.45)
Mg
3.0 (2.2–3.9)Si
10 (8.6–12)S
12 (9.1–16)
Ca18§ (5.1–44)
Fe18§ (5.1–44)
Ni18§ (5.1–44)−3
ne t (cm s)
7.7 (6.7–9.2) ×108
−3nH ne V (cm )1.8 (1.3–2.5) ×1052
NEI (ISM)
kTe (keV)
0.51 (0.31–0.55)
ne t (cm−3 s)
5.8 (5.7–6.1) ×109
−3nH ne V (cm )1.1 (1.0–1.2) ×1056
Power-lawΓ
2.9 (2.8–3.0)
Normk (photons cm−2 s−1 )6.3 (5.5–7.2) ×10−4
Gaussian lines to complement the incomplete VNEI code (see text)
Line
Center Energy(keV)(fixed)
Norm (photons cm−2 s−1 )O VII-Kδ
0.714
6.4 (6.1–6.6) ×10−4
O VII-K
0.723
3.2 ×10−4 #O VII-Kζ
0.730
1.6 ×10−4 ∗∗
Ar-K
3.01
5.7 (4.4–7.0) ×10−6
Ca-K
3.69
2.4 (1.3–3.5) ×10−6
Gain
Offset (eV)–4.82
χ /d.o.f.
996/831 = 1.20
∗ † ‡, , and § represent linked parameters.
k
The differential flux at 1 keV.
#
Fixed to 50% of the normalization of O VII-Kδ.
∗∗
Fixed to 50% of the normalization of O VII-K.
−2

63

64

CHAPTER 5. SN 1006

Fig. 5.11.—: XIS spectra in the 0.33-10 keV band which were extracted from the red elliptical region shownin figure 5.6. The black and red data points represent the FI and BI spectra, respectively. The solid lines,
colored green, blue, light blue, orange, and magenta show the components of the ISM, Ejecta 1, Ejecta 2,
Power-law, and additional Gaussians of the best-fit model for the BI spectrum. Note that the Gaussian likestructures (magenta) at ∼1.2 keV and ∼1.9 keV are Si escape peaks (see figure 4.9 in section 4.3) of theadded Gaussians at 3.02 keV and 3.69 keV due to the XIS response.

5.6. DISCUSSION FOR SN 1006

5.6

65

Discussion for SN 1006

In section 5.5, we analyzed the full energy band spectrum extracted from the elliptical region
shown in figure 5.6, and found that it could be described adequately with a parameterized model
including three thermal plasmas in non-equilibrium ionization (VNEI 1, VNEI 2, and NEI, in
table 5.6) and one power-law component (Power-law, in table 5.6). Two of the plasmas, VNEI 1
and VNEI 2, were assumed to have non-solar abundances, and the other, NEI, was assumed tohave solar composition. In the following we discuss the implications of our spectral results.

5.6.1

Origin of the Plasmas

The NEI model, assumed to have solar abundances, produces most of the observed low-energy
X-rays, particularly the Kα lines from O VII, O VIII, and Ne IX (see figure 5.11). This component
is rather uniformly extended over the entire remnant, as can be seen in the O VII line band map(figure 5.4). We associate this component with the swept-up ISM.
The other plasmas, VNEI 1 and VNEI 2, with their non-solar elemental abundance ratios, areplausibly of ejecta origin. Since VNEI 1 has a larger ionization parameter than the other plasmas,
we suggest that this plasma was heated by reverse shock in the early stage of remnant evolution(here Ejecta 1). On the other hand, VNEI 2 has an extremely low ionization parameter, and hence
should have been heated much more recently (here Ejecta 2). We have firmly detected Fe-Kα
emission for the first time. The low ionization state of Fe, plus its overabundance in the Ejecta 2
component, is generally consistent with the Type Ia origin of SN 1006.

5.6.2

Relative Abundance in the Ejecta

We compare our best-fit relative abundances with the predicted nucleosynthesis yield of thewidely-used W7 Type Ia SN model (section 2.3; Nomoto et al. 1984). For Ejecta 1, although theabundances of C, Ne, Mg, Si, and S relative to O are broadly consistent with the W7 model, Ca
and Fe fall far below their predicted values, as shown in figure 5.12a. In the case of Ejecta 2,
on the other hand, the heavy elements, albeit with large errors, are broadly consistent with theabundance pattern from the W7 model, as shown in figure 5.12b. These results, along with thedifference in the ionization timescale between the components just mentioned, are consistent witha layed composition of the ejecta with the higher-Z elements more concentrated toward the centerof SN 1006.
For this simple interpretation, one expects the Fe line emission from the low-ionization
timescale component to be spatially located interior to the Mg and Si emission from the highionization component, i.e., near the center of SN 1006. As shown in figure 5.6, though, the Fe flux
appears to peak close to the southeastern rim. In summary, our two-component spectral model forthe ejecta is a highly simplified view of what is surely a complex multi-n e t and multi-kTe structurethat varies throughout the interior of the SNR. A longer observation, particularly of the inner

66

CHAPTER 5. SN 1006

regions of SN 1006, will help to improve the statistical accuracy of the Fe-K line flux.

5.6.3

ISM Density and Ion Temperature

The plasma parameters given in table 5.6 were derived from the southeast solid ellipse of
figure 5.6, which is π × 6.0 1 × 4.0 4 = 84 arcmin2 in area. This corresponds to 3.3 × 10 38 cm2 , fromwhich we estimated the emission volume to be V = (3.3 × 10 38 )3/2 = 6.0 × 1057 cm3 . Therefore,
the EM of the heated ISM derived from the spectral fits corresponds to an electron density of n e =0.15 f −0.5 cm−3 , where f is the filling factor. Since the age of SN 1006 is ∼1000 yr, the ionization
timescale can be roughly estimated as n e t ∼ 4.7 × 109 cm−3 s, which is almost consistent with thebest-fit value of the ionization parameter of the NEI (ISM) component. Such a low ionization
age suggests that the temperatures of the ions and electrons may also be far from equilibrium.
According to equation 3.39, under the assumption that the initial ratio between kT e and kTHjust at the shock front is the ratio of the electron and proton masses, the proton temperature isestimated to be,ne tkTH = 7.8
5.8 × 109 cm−3 s


−1 

kTe
0.51 keV

5/2[keV] .

This result is consistent with that of the Hα observation of T e /TH ≤ 0.07 (Ghavamian et al. 2002).
The extreme non-equilibrium state of the plasma is due to the low density of the ambient medium.
According to equations (2) and (5) of Ferrière (2001), the density of H I and H II at the Z-height of
SN 1006 (Z = 550 pc) is ∼0.03 cm−3 . This is consistent with our estimate of the ambient densityof nH /4 ∼ 0.03 cm−3 . Recently, Acero et al. (2007) obtained the gas density around the SE rim tobe ∼ 0.05 cm−3 using XMM-Newton data, which is consistent with our result.

5.6.4

The Power-Law Component

The low density of the ambient medium will allow the velocity of the shock front to remainhigh for a relatively long time. This may be one reason why this remnant shows such efficientparticle acceleration up to very high energies (∼ 10 14 eV), observed as power-law (synchrotron)
emission from the NE and SW rims (Koyama et al. 1995). The photon index from these rims
is Γ ∼ 2.7–2.9. Now here, away from the bright rims, we also detect a power-law component in
the hard energy band (the Power-law, in table 5.6) with a best-fit index of Γ ∼2.8–3.0, which isconsistent with the NE and SW bright rim emission. Since the power-law emissions from the rims
are most likely to be synchrotron X-rays, the power-law component newly found from the ellipse infigure 5.6 is also likely to be synchrotron X-ray emission. The surface brightness is estimated to be∼ 6 × 10−15 ergs cm−2 s−1 arcmin−2 in the 2–10 keV band. This value is ∼ 50-times lower than thepeak emission at the NE and SW rims. If similar to the NE rim, the power-law component should
show thin filamentary structures (Bamba et al. 2003). However, XMM-Newton (Rothenflug et al.
2004) and Chandra (Hughes et al., in preparation) observations found little evidence for filamentary
X-ray rims along the SE rim.

5.6. DISCUSSION FOR SN 1006

67

Fig. 5.12.—: Metal abundances as a function of atomic number derived from the spectral fitting. The
data points of (a) and (b) represent those of Ejecta 1 and Ejecta 2, respectively. The solid lines show theabundance relative to oxygen calculated in the W7 model for a Type Ia supernova by Nomoto et al. (1984).

68

CHAPTER 5. SN 1006

5.6.5

Summary for SN 1006

We have analyzed high-quality spectra obtained with Suzaku of a region in the southern partof SN 1006, selected because it is bright in Fe-Kα. The results and interpretations are summarized
as follows:
1. The spectrum can be described by a model with at least three NEI thermal plasmas and one
power-law component.
2. The fits yield different ionization parameters of n e t ∼ 8 × 108 cm−3 s (low), ∼ 6 × 109 cm−3 s(medium), ne t ∼ 1010 cm−3 s (high).
3. The low-ne t plasma are highly overabundant in heavy elements, in which we found Kα lines
from Fe, for the first time. The abundance pattern is consistent with that of Type Ia SN
ejecta.
4. The abundance of the medium-ne t plasma is assumed to be solar, and we associate thiscomponent with the shocked ISM. Although oxygen is not overabundant, Kα lines of O VII
and O VIII from this plasma appear to dominate the thermal emission from SN 1006.
5. The high-ne t plasma is overabundant in medium elements, like Mg, Si, and S, but heavy
element like Fe. Our interpretation is that this plasma also has an ejecta origin, and thecomposition is dominated by lower atomic number species.
6. Temperature equilibrium between ions and electrons has not yet been achieved. The proton
temperature of the shocked ISM is ∼15-times higher than the electron temperature.
7. The spectrum from this region contains a power-law component. The photon index, Γ ∼2.9,
is similar to that of the northeast and southwest bright rims, suggesting that this componentis also of synchrotron origin.

Chapter 6

E0509–67.5
E0509–67.5 is probably the youngest Type Ia SNR in the Large Magellanic Cloud (LMC).
(SN 1987A is younger, but is a core-collapse type.) Therefore, X-rays from E0509–67.5 are dominated by the radiation from the ejecta, similarly to SN 1006. Since the LMC is a face-on galaxy
and is located in the high Galactic latitude region, the X-rays can be observed without almost
suffering from the interstellar absorption. Furthermore, the distance to the LMC is well-known to
be 51 kpc (Feast et al. 2002), and hence the amounts of synthesized elements in the SNR ejecta
can be investigated with a high reliability. In this chapter, we summarize the results of Suzaku
observation on E0509–67.5, and discuss the nucleosynthesis of its progenitor.

6.1

Previous Results

E0509–67.5 was discovered by an X-ray survey of the LMC with the Einstein Observatory(Long et al. 1981). With an Hα observation, the remnant was found to have a very circular shelllike structure (Tuohy et al. 1982), as shown in figure 6.1. It shows no detectable emission in O III
and Si II, indicating a Balmer-dominated remnant (Smith et al. 1991; 1994). Hα studies also found−1<
it has a very fast shock velocity of >∼ 2000 km s , and a young age of ∼ 1500 yr (Smith et al. 1991).
Radio observations obtained a spectral index of α = 0.46 (Mathewson et al. 1983) and a differential
flux of 0.066 Jy at 1 GHz (Hendrick & Reynolds 2001), of which the origin is suggested to besynchrotron emission.
In X-rays, ASCA found, for the first time, its spectrum has strong K-shell emissions of Si
and S, and L-shell emissions of Fe, with much weaker O, Ne, and Mg K-shell emission (Hughes
et al. 1995). This fact strongly suggests that E0509–67.5 has a Type Ia SN origin. With a great
spatial resolution, Chandra revealed the clear shell-like structure in the X-ray image (figure 6.2),
and obtained its inner and outer radii to be 12.8 arcsec and 15.2 arcsec, respectively (Warren &Hughes 2004). Furthermore, they concluded that the age is less than 860 yr, if the constant-density
ejecta profile is assumed. The most serious issue raised by the Chandra observation is that the69

70

CHAPTER 6. E0509–67.5

Fig. 6.1.—: Hα image of E0509–67.5, taken from Tuohy et al. (1982). North is up and east is to the left.

ionization age of Fe is higher than those of other elements. This picture is not similar to the Suzaku
result on SN 1006 which we derived in chapter 5. However, the efficiency and energy resolution of
Chandra is relatively poor, especially in the energy band of Fe-K emission. Therefore, it is possiblethat the ionization age was not measured exactly. We hence observed E0509–67.5 to perform an
accurate plasma diagnosis.
−3
The density of the ambient gas is known to be very low ( <∼ 0.02 cm : Tuohy et al. 1982;< 0.05 cm−3 : Warren & Hughes 2004), similarly to SN 1006. Therefore E0509–67.5 is also suited∼
to the study of the X-ray spectral property in the very early phase of SNR evolution.

6.2

Suzaku Observation

The Suzaku observation of the SNR E0509–67.5 was made on 2006 August 18 (Observation
ID = 501041010) with the targeted position at (RA, Dec) = (05:09:33, –67:30:54). The XIS wereoperated in the normal full-frame clocking mode. We employed the cleaned revision 1.2 data (almost
consistent with the revision 0.7 data used in the analysis of SN 1006), and used the HEADAS
software version 6.0.4 and XSPEC version 11.3.2 for the data reduction and analysis. The X-ray
data taken during low cut-off rigidity (≤ 6 GV) were excluded to minimize positional dependenceof the NXB. After the screening, the effective exposure time was obtained to be ∼45 ksec. The
response matrix files (RMF) and ancillary response files (ARF) were made using xisrmfgen andxissimarfgen (Ishisaki et al. 2007) version 2006-10-17.
Figure 6.3 shows the XIS image in 0.5–5.0 keV band. Since the diameter of the SNR (∼ 30 00 )is much smaller than the angular resolution of the XRT, (∼ 2 0 : Serlemitsos et al. 2007), the SNR
can be regarded as a point-like source spread by the effect of the PSF (see section 4.2).

6.2. SUZAKU OBSERVATION

71

Fig. 6.2.—: Chandra image of E0509–67.5, taken from Warren & Hughes (2004). Red, green, and blue
contain emission from 0.2–0.69 keV, 0.69–1.5 keV, and 1.5–7 keV, respectively. North is up and east is tothe left.

72

CHAPTER 6. E0509–67.5

20:00.0

25:00.0-67:30:00.0

35:00.0

40:00.0

30.0

11:00.0

30.0 5:10:00.0

30.0

09:00.0

30.0

08:00.0 07:30.0

Fig. 6.3.—: XIS 0.5–5.0 keV image. The coordinates (RA, Dec) refer to epoch J2000.0. The field of viewof the XIS is indicated by the black square, where the data from the four XIS are combined. The green circle
and annulus are the source and background regions, respectively.

6.3

Spectral Structure in the Narrow Bands

We extracted XIS spectra from a circular region of 3 0 radius centered on E0509–67.5 shownin figure 6.3. The X-ray flux in this region is ∼95% of the total flux of the source (see the rightpanel of figure 4.5 in section 4.2). The background spectra were extracted from the annular region
around the source with inner and outer radii of 3 0 and 60 , respectively. The background region is
also shown in figure 6.3. Since the data from the three FIs are nearly identical, we merged theindividual spectrum to improve the photon statistics. The background-subtracted full-band spectra
are shown in figure 6.6.
At first, we divide the spectra into narrow energy bands and study the plasma characteristics
in relation to the various elements.

6.3.1

The 4–10 keV Band Spectrum

In order to determine the center energy of an Fe-Kα line, we first fitted the 4–10 keV spectrawith a power-law (for the continuum) and a Gaussian (for the emission line). The best-fit centroidof the Fe-K line was obtained to be 6436 (6406–6466) eV.
We then fitted the spectra with a solar abundance (Anders & Grevese 1989) VNEI model. If

6.3. SPECTRAL STRUCTURE IN THE NARROW BANDS

73

Fig. 6.4.—: XIS spectra in the 4–10 keV band fitted with the VNEI model. Black and red data points
represent the FI and BI spectra, respectively.

we fix the Fe abundance to be solar, the model cannot reproduce the Fe-Kα profile and returns an
unacceptable χ2 /d.o.f. of 52/28. Allowing the Fe abundance to be free, the fit was greatly improved
with best-fit χ2 /d.o.f. of 14/27. The best-fit kTe , ne t, and Fe abundance are 3.2 (2.5–9.0) keV,
7.1 (1.5–13) ×109 cm−3 s, and 18 (>7.1) solar, respectively. Since the Fe abundance is significantly
enhanced relative to the solar value, we can conclude the Fe-K emission originates from the ejecta
heated by a reverse shock. The best-fit spectrum is shown in figure 6.4.

6.3.2

The 1.1–4.0 keV Band Spectrum

We can see, in the 1.1–4.0 keV band spectra, strong Kα lines of Si and S as well as Kβ ofthese elements and Kα of Mg, Ar, and Ca. We ignored the data in 1.83–1.85 keV because of thesame reason as noted in subsection 5.4.3 (calibration issue). First, we fitted the 1.1–4.0 keV spectrawith one-component VNEI model. Since K-shell emission lines from Ar are not included in the
NEIvers 1.1 spectral code, we added a Gaussian at ∼ 3.1 keV. (The center energy of the line was
fixed but tuned to obtain a best-fit.) In this narrow band fitting, interstellar absorption was fixed
to a hydrogen column density of 6.5 × 10 20 cm−2 , which is close to the value obtained by Warren& Hughes (2004). Then, the best-fit χ 2 /d.o.f. obtained was an unacceptable value of 920/350.
Although we tried to fit the spectra with one-component VPSHOCK model, this model also failed
to fit with χ2 /d.o.f. value of 696/350. The results of the VNEI and VPSHOCK fits are summarizedin table 6.1, while the spectra fitted with the one-component VNEI are shown in figure 6.5a.
The most remarkable disagreement between the data and model (in figure 6.5a) is found in
the center energy of the S-Kα line (around ∼ 2.45 keV); the center energy predicted by the model
is significantly lower than that of the data. In the 1.1–4.0 keV band, the most prominent emission

74

CHAPTER 6. E0509–67.5

Fig. 6.5.—: (a) The XIS spectra in the 1.1–4.0 keV band fitted with the one-component VNEI model. The
black and red data points represent the FI and BI spectra, respectively. A Gaussian at 3.1 keV is added inorder to represent the Ar-K line which is not included in the VNEI model. (b) The same to (a) but fittedwith the two-component VNEI model. A Gaussian at 2.12 keV is also added to represent the Kβ emission
line from low-ionized Si.

6.3. SPECTRAL STRUCTURE IN THE NARROW BANDS

75

Table 6.1:: Best-fit parameters of the spectral fittings in the 1.1–4.0 keV band with one-component VNEIand VPSHOCK models.

ParameterkTe (keV)
Mg
Si
S
Ca
ne t (cm−3 s)χ2 /d.o.f.

VNEI
2.2
0.51
14
6.7
5.3
7.6 × 109920/350 = 2.63

VPSHOCK
1.6
0.40
10
8.5
8.5
1.7 × 1010
696/350 = 1.99

Table 6.2:: Best-fit parameters of the spectral fitting in the 1.1–4.0 keV band with 2-component VNEI
model with different ne t values.

Parameter
Component 1
Component 2−2
20
NH (cm )6.5×10 (fixed)kTe (keV)
1.6 (1.5–1.7)
Mg
0.37 (0.32–0.40)
0 (fixed)
Si
5.2 (4.9–5.5)
25 (24–27)
S
14 (13–15)
0 (fixed)
Ca
6.7 (4.5–9.3)
0 (fixed)
ne t (cm−3 s)
1.7 (1.6–1.8) ×1010
4.0 (3.7–4.2) ×109
nH ne V (cm−3 )1.40 (1.37–1.44) ×1058
5.96 (5.68–6.30) ×1057
Gaussian lines to compensate for the incompleteness of the VNEI code
Line
Center Energy(keV)
Norm (photons cm −2 s−1 )Si-Kβ ∗
2.12
3.4 (2.3–4.6) ×10−6
Ar-Kα
3.10
1.7 (1.0–2.4) ×10−6
χ2 /d.o.f.
328/346 = 0.95
∗ Kβ line from low-ionized Si (lower than He-like).

line is Si-Kα, and hence the best fit value of n e t is mainly determined by the center energy of thisline. Therefore, this disagreement suggests that the mean ionization age (n e t) of Si is lower thanthat of S.
In subsection 5.4.3, we have shown the middle band (1.2–2.8 keV) spectrum of SN 1006 canbe well fitted with the two-component VNEI model which has the common electron temperature
but different ne t values. Here, we applied the same method and tried to fit with a two-component
VNEI model. Then, the fitting was greatly improved with acceptable χ 2 /d.o.f. value of 354/344.
Although the good fit was obtained with this two-component model, we found a data excess still

76

CHAPTER 6. E0509–67.5

remaining around the energy of 2.12 keV. We considered this excess is due to a Kβ emission from
low-ionized Si, because the NEIvers 1.1 spectral code lacks the Kβ emission line from Si whose
ionization state is lower than He-like. Therefore, we added a Gaussian at 2.12 keV (the best-fit
center energy was fine-tuned) and fitted again. The best-fit parameters and spectra are given
in table 6.2 and figure 6.5b, respectively. Note that although we first treated the abundances of
Mg, Si, S, and Ca in both components as free parameters, all of them in the lower-n e t component(Component 2, in table 6.2) did not need significant values, except the Si abundance. Therefore, we
finally fixed the abundance of Mg, S, and Ca of the lower-n e t component to 0. The χ2 /d.o.f. value(328/346) was significantly reduced from the result without the additional Gaussian at 2.12 keV.
Moreover the intensity ratio of the 2.12 keV Gaussian to the Si-Kα in the lower-n e t component is∼ 7% (see figure 6.5b), which is a suitable value for the Kβ/Kα ratio for the Si in the low ionization
state.

6.4

Broadband Spectrum

In section 6.3, we separately derived spectral parameters for the energy bands including Kshell emissions of intermediate-mass elements (Mg, Si, S, Ar, and Ca) and heavy element (Fe).
The summary is listed in table 6.3. We did not try to fit the O-Ne band spectrum in the previous
section, because the E0509–67.5 spectrum shows very weak K-shell lines from light elements (Oand Ne) differently from the spectrum of SN 1006. We use these parameters as the initial values
to search for the plasma parameters in the entire XIS energy band of 0.4–10 keV.
Based on table 6.3, we assume the model that describes the full spectral range must consistof, at least, two plasma components with (1) high-n e t (∼ 2 × 1010 cm−3 s) and (2) low-ne t (∼4 × 109 cm−3 s). As shown in table 6.2, the component (1) is significantly enhanced in Si, S,
and Ca, relative to solar. This fact suggests that these elements originate from the ejecta of
the SNR. Therefore we make an assumption that this component is composed purely of metals
without hydrogen and helium. We fixed the oxygen abundance of the component (1) to a large
value (1 × 104 ) and determined the abundances of other elements relative to oxygen with a fitting.
The C and Ni abundances were assumed to be the same as Ne and Fe, respectively. As alreadynoted in section 5.5, the emission measure (EM ) in the plasma code is defined (in XSPEC) as

Table 6.3:: Plasma components determined from the narrow band spectra.

Major elementskTe (keV)ne t (cm−3 s)
Mg, Si, S, Ca
1.6 (1.5–1.7) 1.7 (1.6–1.8)×10 10Si
1.6 (1.5–1.7) 4.0 (3.7–4.2)×109
4–10
Fe
3.2 (2.5–9.0) 7.1 (1.5–13)×10 9∗ The plasma component identifications described in text.
Band (keV)
1.1–4.0

Component∗
(1)
(2)
(2)

6.4. BROADBAND SPECTRUM

77

nH ne V even if the plasma is dominated by heavy elements. The oxygen density can be calculated
as nO = (8.51 × 10−4 nH ) × 104 , where the numerical value is the solar abundance of oxygen from
Anders & Grevesse (1989). The component (2), on the other hand, is enhanced in Si and Fe only.
Therefore, we assume that this component is composed purely of Si, Fe, and Ni without any otherheavy elements as well as hydrogen and helium. (Although there are no apparent feature of Ni Kand L-shell lines, we make the same assumption to component (1); the Ni abundance is equal to theFe.) We thus fixed the Si abundance of the component (2) to 1 × 10 6 and fitted for the abundanceof Fe (and Ni) relative to Si. The abundances of the other elements are fixed to 0.
Warren & Hughes (2004) concluded the continuum flux is not dominated by hydrogen thermal bremsstrahlung radiation but non-thermal synchrotron emission, since an ambient density−3(nH <∼ 0.05 cm ) derived from the dynamical state of the remnant is too low to contribute to thecontinuum flux. We, here, follow their conclusion, and hence included a power-law component forthe continuum and fitted the spectra in the energy band from 0.4 to 10 keV. With the same reasons
as noted in subsection 6.3.2, we added Gaussians at 2.12 keV (for Si-Kβ) and 3.10 keV (for Ar-Kα).
With this model, the best-fit χ2 /d.o.f. was obtained to be an unacceptable value of 1131/719.
We found the most remarkable disagreement between the data and model around 0.73 keV. In
subsection 5.4.2, we found in the SN 1006 spectrum the disagreement at this energy as well, andpointed out this was plausibly due to the higher K-shell transitions of O VII (i.e., Kδ, , ζ, etc.),
which were not included the spectral code we used. In the case of the E0509–67.5 spectrum, however,the same explanation cannot be applied because of the weak intensity of O VII series. Therefore
we concluded this disagreement is purely due to Fe L-shell emission and added a single Gaussian
at 0.73 keV (the best-fit energy is fine-tuned) in order to compensate for the incompleteness of thespectral code. Warren & Hughes (2004) also found the 0.73 keV enhancement in the Chandra ACIS
spectrum of E0509–67.5 and pointed out the necessity of adding a separate Gaussian component.
Then, the fit with the additional Gaussian was significantly improved with χ 2 /d.o.f. = 789/718.
The best-fit parameters and the full band spectra fitted with the best-fit model are given in table 6.4
and figure 6.6, respectively.

78

CHAPTER 6. E0509–67.5

Table 6.4:: Best-fit parameters of the broadband spectral fitting.

Component
Absorption
VNEI 1 (Ejecta 1)

Parameter

ValueNH
5.8 (4.6–7.6) ×1020
kTe (keV)
3.7 (2.6–4.2)
4
Abundance (10 solar) C
1.4∗ (1.2–1.5)
N
0 (fixed)O
1.0 (fixed)Ne
1.4∗ (1.2–1.5)
Mg
0.92 (0.77–1.1)
Si
17 (16–19)
S
35 (33–40)
Ca
16 (13–20)
Fe
2.0† (1.8–2.1)
Ni2.0† (1.8–2.1)
ne t (cm−3 s)
1.4 (1.3–1.5) ×1010
nH ne V (cm−3 )3.65 (3.48–3.84) ×1053
VNEI 2 (Ejecta 2) kTe (keV)
3.5 (3.3–3.8)
Abundance (104 solar) Si
100 (fixed)Fe62‡ (52–85)
Ni62‡ (52–85)−3
ne t (cm s)
3.5 (3.1–3.6) ×109
nH ne V (cm−3 )1.45 (1.36–1.51) ×1053
Power-lawΓ
3.7 (3.5–3.8)§−2−1
Norm (photons cm s )3.7 (3.4–3.9) ×10−4
Gaussian lines to compensate for the incompleteness of the VNEI code
Line
Center Energy(keV)
Norm (photons cm −2 s−1 )Fe-L blends
0.733
4.1 (3.7–4.4) ×10 −4
Si-Kβ
2.12
3.0 (1.9–4.0) ×10−6
Ar-Kα
3.10
2.1 (1.5–2.8) ×10−62
χ /d.o.f.
789/718 = 1.10
∗ , † , and ‡ represent linked parameters.
§ The differential flux at 1 keV.
(cm−2 )

6.4. BROADBAND SPECTRUM

79

Fig. 6.6.—: XIS spectra of E0509–67.5 in the 0.4–10 keV band. The black and red data points represent
the FI and BI spectra, respectively. The solid lines colored blue, light blue, orange, and magenta show thecomponents of the Ejecta 1, Ejecta 2, Power-law, and additional Gaussians of the best-fit model for the FI
spectrum. The green solid line indicates an estimated flux of an ISM component assuming the density ofambient medium to be 0.05 cm−3 (see text in subsection 6.5.3). kTe , ne t, and a metal abundance of thiscomponent are assumed to be 1 keV, 1 × 1010 cm−3 s, and 0.3 solar, respectively. Note that this component
is not included in the spectral fitting.

80

6.5

CHAPTER 6. E0509–67.5

Discussion for E0509–67.5

In section 6.4, we analyzed the broadband spectrum of E0509–67.5, and found that it could be
described with a parameterized model including two thermal plasmas in non-equilibrium ionization(VNEI 1 and VNEI 2, in table 6.4) and one power-law component (Power-law in table 6.4). In
following, we discuss the implications of our spectral results.

6.5.1

Origin of the Components

Both of the thin-thermal plasmas, VNEI 1 and VNEI 2, are overabundant in heavy elements.
Therefore we interpreted they have ejecta origin. Since VNEI 1 has a larger ionization parameter
than the other, this plasma must be heated by the reverse shock in the early stage of the remnant
evolution (hereafter Ejecta 1). On the other hand, VNEI 2 has relatively low ionization age.
Therefore it should have been heated by the reverse shock very recently (hereafter Ejecta 2). We
found that the Ejecta 2 consists of Si, Fe, and Ni only.
The Fe-K emission is mainly produced by the Ejecta 2 component (see figure 6.6), whichsuggests the mean ionization age of Fe is lower than the other elements (e.g., S and Ca) whose
K-shell emission is purely produced by the Ejecta 1. This result differs from the Chandra result
derived by Warren & Hughes (2004): the plasma consisting of the Fe ejecta has a higher ionization
timescale (and a much higher electron temperature) than the other elements. They used the
VPSHOCK spectral code for their fitting and obtained the best-fit kT e and ne t for the Fe ejecta to
be 10.0 keV and 3.4 × 1010 cm−3 s, respectively. However, we found these spectral parameters with
the VPSHOCK model predicts the Fe-K energy centroid of 6489 eV, which is significantly higher
than our result (6436 ± 30 eV: subsection 6.3.1). Because of its limited effective area, the Chandra
spectrum has no sufficient photon statistics of Fe-K line (see figure 4 of Warren & Hughes 2004).
Therefore, the ionization timescale for the Fe ejecta would be determined only from the profile of
Fe-L blends, of which the atomic physics is not as well understood, and would be overestimated as
the result.
The power-law component is most likely to be a synchrotron emission. The photon index ofΓ = 3.7 (3.5 − 3.8) is almost consistent with the Chandra result of Warren & Hughes (2004). As
shown in figure 6.2, the filamentary hard emission (blue) is found at the outermost of the remnant.
It suggests the electrons are accelerated up to more than TeV energy at the blast wave of the−1
remnant with a very high shock velocity ( >∼ 3600 km s : Tuohy et al. 1982), similarly to SN 1006
(Koyama et al. 1995, Bamba et al. 2003).

6.5.2

Relative Abundance and Mass of the Ejecta

The low abundances of O and Ne compared to those of intermediate-mass elements (i.e., Si
and S) suggest E0509–67.5 is a product of a Type Ia SN. Therefore we compare our best-fit relative
abundances with the predicted nucleosynthesis yield of the Type Ia SN deflagration model (W781

6.5. DISCUSSION FOR E0509–67.5
Table 6.5:: Mass of major elements in the ejecta for several assumed filling factors.

f ∗
0.5
0.8
O (M )
0.14
0.18
Si (M )0.56
0.47
S (M )
0.19
0.24
Fe (M )
0.64
0.44
Total (M )
1.68
1.50
∗ Filling factor for the Ejecta 1.

0.9
0.19
0.41
0.26
0.33
1.38

0.95
0.20
0.37
0.27
0.26
1.27

0.99
0.20
0.31
0.27
0.16
1.12

model: Nomoto et al. 1984). The data points in figure 6.7a indicate the relative abundances inthe Ejecta 1, while the solid line shows those of W7 model. Since the model does not reproduce
our result well, we also compare the data with the delayed detonation model (WS15DD3 model)
predicted by Iwamoto et al. (1999), which is shown as the dashed line in figure 6.7a. We can see thelatter model has the better reproducibility of the abundance ratio between O, Mg, and S. However,
the abundances of Si, Ca, and Fe relative to O are much lower than the predicted values. This
result suggests these elements are concentrated on the inner layer of the remnant and hence did
not heated enough in the early stage of the SNR evolution. Indeed, Si and Fe are enhanced in
the Ejecta 2, which is interpreted as the ejecta heated more recently. Although the Ne abundance(relative to O) is much larger than the predicted value, note that it is not so reliable because the
Ne K-shell lines (at the energies of 0.9–1.1 keV) are very weak and embedded in the Fe-L blends.
As shown in figure 6.2, E0509–67.5 has a clear shell-like structure. By the image fitting, Warren& Hughes (2004) determined the inner and outer radii of the shells to be 12.8 arcsec and 15.2 arcsec,
respectively. These values correspond to r in = 9.4 × 1018 cm and rout = 1.1 × 1019 cm, from which
3 − r 3 )/3 = 2.4 × 1057 cm3 . Here, we define thewe estimate the emission volume to be V = 4π(r inoutfilling factor for the Ejecta 1 as f , and assume that for the Ejecta 2 as (1 − f ). The mass (sum ofEjecta 1 and 2) of each major element for several values of f is estimated as shown in table 6.5.
(Note that the elemental mass values do not depend on the assumed oxygen or silicon abundances
relative to hydrogen.) The relative abundances for the several f are also shown in figure 6.7b.
If we assume f <∼ 0.8, the estimated total mass of the ejecta exceeds the Chandrasekhar limitof a white dwarf. Therefore, the filling factor for the Ejecta 1 may larger than ∼ 0.8. Then the Femass must be <∼ 0.4M , which is significantly lower than the predicted value of ∼ 0.6–0.8M (e.g.,
Iwamoto et al. 1999). This result suggests a large amount of Fe ejecta still remain at the inner
region of the remnant, and hence have not been heated by the reverse shock yet.

6.5.3

Contribution of the ISM

Differently from the result of SN 1006 (section 5.5), our best-fit model for the E0509–67.5
broadband spectrum does not need the contribution of an emission from the ISM heated by the

82

CHAPTER 6. E0509–67.5

Fig. 6.7.—: (a) Metal abundances of the Ejecta 1 as a function of atomic number derived from the spectral
fitting. The solid and dashed lines show the abundance relative to oxygen calculated in the W7 (Nomoto
et al. 1984) and WS15DD3 (Iwamoto et al. 1999) models for a Type Ia supernova. (b) Abundance ratio
for the sum of the Ejecta 1 and 2 , when the filling factor for the Ejecta 1 (see text) is assumed to be 0.5(circle), 0.9 (triangle), and 0.99 (star). The dashed line is the same to (a).

6.5. DISCUSSION FOR E0509–67.5

83

blast wave. This result is consistent with Warren & Hughes (2004). Here, we estimate the flux of the
ISM emission and compare it with that of the full-band spectrum. The optical studies constrained−3
the density of the ambient medium to n H <∼ 0.02 cm (Tuohy et al. 1982). Furthermore, it was−3also confirmed by the X-ray study (nH <∼ 0.05 cm : Warren & Hughes 2004). If we assume thatthe ambient medium is homogeneous with the density of 0.05 cm −3 , the EM of the swept-up ISM
is estimated to be 5.6 × 1055 cm−3 . The spectrum with this EM is shown as the green solid line infigure 6.6, assuming kTe , ne t, and a metal abundance to be 1 keV, 1 × 10 10 cm−3 s, and 0.3 solar(the abundance of the LMC: Dickey & Lockman 1990), respectively. The flux of this component
is found to be much lower than the other components. Therefore, the contribution of the ISM
emission is negligible relative to the whole SNR spectrum, independently of the actual kT e and ne tvalues of this component.

6.5.4

Summary for E0509–67.5

We have analyzed the high-quality spectra of E0509–67.5 which is the youngest and the brightest SNR in the LMC, obtained with Suzaku. The results and interpretations are summarized asfollows:
1. The spectrum can be described by a model with at least two NEI thermal plasmas and one
power-law component.
2. The fits yield different ionization parameters of n e t ∼ 3.5 × 109 cm−3 s (low), and ne t ∼1.4 × 1010 cm−3 s (high).
3. The Fe-K emission is mainly produced by the low-n e t plasma. The center energy of the Fe-Kαline is ∼ 6.44 eV.
4. The high-ne t plasma is overabundant in Si and S. However, the abundance of Fe is much
lower than the predicted value for Type Ia SN ejecta.
5. The spectrum contains a power-law component with the photon index of Γ ∼3.7, suggestingthat this component has a synchrotron origin.
6. The contribution of the emission from the swept-up ISM is much smaller than those of theother components.

Chapter 7

RCW 86
As shown in chapters 5 and 6, the ejecta in SN 1006 and E0509–67.5, which are evolving in thinambient gas, have extremely low ionization ages. These facts suggest that progresses of heating
and ionization processes in SNR ejecta are strongly influenced by the densities of the circumstellar
matter. RCW 86, the oldest historical SNR, is known to be expanding in non-uniform ISM (spatial
variation of the ambient density is quite large: e.g., Vink et al. 1997). Therefore, it is expectedthat the plasma ages of the ejecta are different from position to position even in one SNR. In thischapter, we summarize the results of Suzaku observation and analysis on RCW 86 northeasternshell, where the largest non-uniformity (in the entire SNR) of the ISM density is suggested. Unlessotherwise stated, figures in this chapter are taken from Yamaguchi et al. (2008b), which this chapter
mainly follows.

7.1

Previous Results

RCW 86 (G315.4–2.3) is one of the historical Galactic supernova remnants (SNRs), a possible
remnant of supernova in AD 185 (Stephenson & Green 2002). It was first discovered by a radio
survey observation (Mills et al. 1961). The energy index in radio band is measured by Caswell etal. (1975) to be α = 0.6. Dickel et al. (2001) presented the imaging and polarimetry of this SNR
at a frequency of 1.34 GHz with the Australia Telescope Compact Array (ATCA), and found thatthe southwest (SW) rim is colliding with dense gas just now. Figure 7.1 shows the 1.34 GHz radio
contour (Dickel et al. 2001) of RCW 86 overlaid by Hα image (gray scale: Smith 1997). Theirmorphologies are different from each other. The distance to RCW 86 was determined to be 2.8 kpc
by optical observations (Rosado et al. 1996); we assume this value in this chapter.
In X-ray band, RCW 86 was discovered by Naranan et al. (1977) with Apollo space-craft. A
higher energy spectrum (2–20 keV) with the Ginga satellite was explained by a high temperatureplasma of kTe >∼ 4 keV (Kaastra et al. 1992). The ASCA observations discovered synchrotron
X-rays from the SW and northeast (NE) shells of RCW 86 (Bamba et al. 2000; Borkowski et85

86

CHAPTER 7. RCW 86

Fig. 7.1.—: 1.34 GHz radio contours and Hα gray scale (Smith 1997) of RCW 86. This picture is takenfrom Dickel et al. (2001).

al. 2001), which indicates that the energy of the electrons reaches up to ∼100 TeV by shock
acceleration, similar to the mechanism in SN 1006 (Koyama et al. 1995). Another remarkablediscovery with ASCA is the ∼6.4 keV line in the spectra of the SW (Vink et al. 1997) and NE rims
(Tomida 1999) of RCW 86. Since the center energy of the line is consistent with Kα from neutral
Fe, a fluorescent origin had often been proposed (Vink et al. 1997; Tomida 1999). However, Bambaet al. (2000) and Borkowski et al. (2001) showed that the X-ray spectrum of the SW shell could be
represented by a three-component model; two thin-thermal plasmas with low (∼0.3–0.8 keV) andhigh (>5 keV) temperatures, and a non-thermal component. The Fe-Kα line was explained by thehigh-temperature plasma with an extremely low ionization parameter (τ = n e t) of ∼ 109 cm−3 s.
Rho et al. (2002) observed the SW shell with Chandra, and spatially resolved the lowtemperature plasma and non-thermal emission. Since the low-temperature plasma is spatiallycorrelated with optical Hα emission, they concluded the origin to be a blast wave. On the other
hand, the non-thermal emission is localized at the inner region of the low-temperature plasma;hence, the origin was suggested to be reverse shock. They found, furthermore, that the Fe-Kα
emission correlated with the non-thermal emission, and suggested that the Fe-Kα line originatesfrom a high-temperature plasma of Fe-rich ejecta heated by the synchrotron-emitting reverse shock.
However, the morphology of the Fe-Kα emitting region could not be determined.
Recently, the SW region of RCW 86 was observed with Suzaku. Utilizing the good sensitivity and spectroscopic performances of the XIS in the energy band of Fe-K lines for diffuse
sources (section 4.3), Ueno et al. (2007) determined the center energy of the Fe-Kα line to be7.1. PREVIOUS RESULTS

876404 (6400–6407) eV, which indicates the emission is from low-ionized iron less than Si-like (Fe XIII).
Furthermore, they revealed the Fe-K line morphology for the first time and discovered its distribution is spatially different from that of the non-thermal emission. They thus confirmed that theorigin of the Fe-Kα is not fluorescence caused by supra-thermal electrons, nor non-thermal X-rays,
but concluded that a more likely origin is a high-temperature ejecta in an extremely ionization
non-equilibrium.
The NE region of RCW 86 also exhibits an Fe K-shell emission (albeit the flux is less than thatof SW) as well as the soft thermal and the non-thermal X-rays (ASCA: Tomida 1999). Therefore,
this region would be another good place to solve the puzzle concerning the origin of the Fe-Kα
related to the thermal and non-thermal emission in RCW 86. Chandra and XMM-Newton revealed
that the soft-thermal and non-thermal emission filaments join smoothly along the outer shell (seefigure 7.2; Vink et al. 2006). They, however, failed to detect Fe-Kα from this region, possibly due
to the limited sensitivity and the high background level near the Fe K-shell energy. We therefore
observed the NE shell with Suzaku to see if the Fe-Kα is really present or not. If present, we will
study its morphology and spectra to reveal the origin.

Fig. 7.2.—: Left: XMM-Newton image of RCW 86. Red, green, and blue represent 0.5–1.0 keV, 1.0–
2.0 keV, and 2.0–6.6keV, respectively. North is up and east is to the left.
Right: Chandra image ofnortheast rim (indicated with the white rectangle in the XMM-Newton image) of RCW 86. The color issame to the left. Credit: ESA/Univ. of Utrecht/Vink et al. (left), NASA/CXC/Univ. of Utrecht/Vink etal. (right).

88

CHAPTER 7. RCW 86

Fig. 7.3.—: HESS image of RCW 86, but a preliminary result. White contours correspond to 3, 4, 5, 6σ
significance. This picture is taken from Hoppe et al. (2007).

Furthermore, TeV γ-rays from RCW 86 were detected by HESS very recently (figure 7.3:
Hoppe et al. 2007). The morphology of the γ-ray emission was confirmed to be shell-like, similarly
to RX J1713.7–3946 (Aharonian et al. 2007a) and RX J0852.0–4622 (Aharonian et al. 2007b).
At present, the origin of the γ-ray is not determined to be hadronic or leptonic. The wide-band
spectral studies of the non-thermal emission will help to reveal the origin.

7.2

Suzaku Observation

The Suzaku observation of the NE region of RCW 86 was made on 2006 August 12 (Observation
ID = 501037010). The XIS were operated in the normal full-frame clocking with the 3 × 3 or 5 × 5
editing mode. We employed cleaned revision 1.2 data, and used the HEADAS software version
6.0.4 and XSPEC version 11.3.2 for the data reduction and analysis. After screening, the effective
exposure time was obtained to be ∼53 ksec. The response matrix files (RMF) and ancillary response
files (ARF) were made using xisrmfgen and xissimarfgen (Ishisaki et al. 2007) version 2006-10-17.

89

7.3. ANALYSIS AND RESULTS

7.3

Analysis and Results

7.3.1

XIS Image

10:00

10:00

We show the Fe-Kα intensity map in figure 7.4 in gray scale. This map was made from the6.3–6.5 keV (the Fe-Kα band) image by subtracting the continuum level in this band, which wasestimated with a power-law model (with photon index of 3) fitting for the 5.0–6.2 keV continuum
spectrum. For comparisons with the soft and hard X-ray fluxes, we overlaid intensity contoursof the 0.5–1.0 keV band (figure 7.4 left) and the 3.0–6.0 keV band (figure 7.4 right). No spatial
correlation between the Fe-Kα and the hard X-ray emission was found. Moreover, the soft Xray emission is systematically shifted outward with respect to the Fe-Kα emission. As we show insubsection 7.3.3, the 0.5–1.0 keV emission is dominated by an optically thin-thermal plasma with anelectron temperature of ∼0.3 keV, and the Ne IX-Kα line is the major emission of this component.
We then made projection profiles of the Ne IX-Kα line band (0.87–0.94 keV), and the Fe-Kα band
(6.3–6.5 keV), in the position given in figure 7.4 left. The result is shown in figure 7.5. We can seethat the peak of Fe-Kα is significantly shifted to the inner region from the soft emission (Ne IX-Kα).
With the fitting of a Gaussian-plus-constant model, the peak positions (right ascensions) of thesoft plasma and Fe-Kα are found to be separated by 2 arcmin.

NE-62:20:00-62:20:00

BGD

30:00

30:00

BGD

East
46:00

14:44:00

46:00

14:44:00

Fig. 7.4.—: 0.5–1.0 keV (left) and 3.0–6.0 keV (right) intensity contours overlaid on the Fe-K map (grayscale) smoothed with a Gaussian kernel of σ = 1.1 arcmin. The coordinates (RA and Dec) refer to epoch
J2000.0. The Field of view of the XIS is indicated by the black squares, where the data from the four XIS are
combined, but that of the four corners, irradiated by the 55 Fe calibration sources, are removed. The green
ellipses and the triangles are the source and background regions, respectively. The blue rectangle indicatesthe region where the projection profiles are extracted (see figure 7.5).

90

CHAPTER 7. RCW 86

Fig. 7.5.—: The projection profiles (along Right ascension) of the Ne IX-Kα band (cross; 0.87–0.94 keV)and the Fe-Kα band (circle; 6.3–6.5 keV) for the blue rectangular region given in the left panel of figure 7.4.
The bin-size of the former corresponds to ∼ 0.0 28, while that of the latter is ∼ 0.0 69.

7.3.2

Spectrum in the Hard X-Ray BandFor a quantitative study, we extracted the representative XIS spectra from two elliptical regions: the East and NE (see figure 7.4), where the soft X-ray and Fe-Kα, and the hard X-ray flux
are the brightest, respectively. We did not divide the former two diffuse emissions because the separation between them is comparable to the angular resolution of the XRT (∼ 2 0 : see section 4.2), andthey are overlapping each other. The solid angles of these ellipses are π × 4. 0 2 × 3.0 1 = 41 arcmin2
and π × 3.0 5 × 1.0 6 = 18 arcmin2 , respectively. The background data were taken from the outsideof the remnant (the solid triangles in figure 7.4). For the background spectra, we excluded thepoint-like sources at (14:46:03, –62:22:30) and (14:46:03, –62:19:01), and the CCD corners whichcontain the 55 Fe calibration source.
We first determined the Fe-Kα line parameters. The background-subtracted spectra in the
3–10 keV band in the East region are shown in figure 7.6. Since the data from the three FIs arenearly identical, we merged the individual spectrum to improve the photon statistics.
We detected a clear Fe-Kα line in the spectra. We then fitted the spectra with a power-law(for the continuum) and a Gaussian (for the emission line). The best-fit parameters are given intable 7.1. The centroid energy of the Fe-Kα line is slightly higher than that of the emission linefrom neutral iron (6400 eV). Although the absolute energy calibration error is reported to be as
accurate as ±0.2%, above 1 keV (Koyama et al. 2007), we checked the energy scale of each XIS
sensor using the 55 Fe calibration sources. We then confirmed that the peak energies of Mn-Kα lines
were consistent with the laboratory value of 5895 eV within the statistical errors for all of the XIS
sensors. We therefore conclude that the energy scale error in the Fe-K line band is less than the

91

7.3. ANALYSIS AND RESULTS

Fig. 7.6.—: XIS 3–10 keV spectra extracted from the East region. Black and red represent the FI (sumof the three sensors) and BI, respectively. The best-fit power-law and Gaussian models are shown by solidlines.

Table 7.1:: Best-fit parameters of 3–10 keV fittings.

Component
Power-law

ParameterEastΓ
3.18 (3.02–3.31)∗
Flux
1.10 (1.07–1.14) ×10−4∗
Surface brightness
2.68 (2.61–2.78) ×10−6
GaussianCenter (eV)6424 (6404–6444)†
Width (eV)
<90
‡
Flux
7.1 (5.8–8.3) ×10−6‡
Surface brightness
1.7 (1.4–2.0) ×10−7
χ2 /d.o.f.
146/192 = 0.76
∗ Photon flux (photons cm−2 s−1 ) and surface brightness(photons cm−2 s−1 arcmin−2 ) in the 3–10 keV band.
† One standard deviation (1σ).
‡ Total flux (photons cm−2 s−1 ) and surface brightness(photons cm−2 s−1 arcmin−2 ) in the line.

NE
2.96 (2.88–3.05)
2.44 (2.39–2.48) ×10−4
1.36 (1.33–1.38) ×10−56424 (fixed)
0 (fixed)<8.1 ×10−7<4.8 ×10−8
204/234 = 0.87

92

CHAPTER 7. RCW 86

statistical error. The best-fit energy of 6424 (6404–6444) eV constrains the Fe ionization state tobe approximately between Ar-like and Ne-like.
For a comparison, we extracted the spectra from the NE region in the same way as that of
the East, and fitted the spectra with a power-law plus a Gaussian. Since we could see no clear the
Fe-Kα line, we fixed the center of the Gaussian to be 6424 eV, and determined the upper limitof the line flux. The result is given in table 7.1. The surface brightness of the Fe-Kα and thepower-law component in the NE are at least ∼3-times smaller and ∼5-times higher than those in
the East region.

7.3.3

Full Band Spectrum

Figure 7.7 shows the 0.4–10 keV spectra from the regions of East and NE. The data reductionis the same as described in subsection 7.3.2. We can clearly see several emission lines below ∼2 keVin the East spectra. This suggests that the soft X-rays are dominantly from an optically thinthermal plasma. Therefore, we fitted the spectra from the East region with a thin-thermal plasma
model and a power-law component. For the thermal plasma component, we used a VPSHOCK(variable abundance plane-parallel shock) model with the NEIvers 1.1 code (since the NEIvers 2.0
code does not include K-shell emission lines from low ionized Fe, we do not use this new version).
Note that if we use a VNEI model which is used in the other chapters, the result and interpretation
do not change essentially. We treated the abundances relative to solar (Anders and Grevesse 1989)
to be free parameters, with the following elements in the parenthesis fixed to be the same: (C,N, and O), (Si and S), and (Ca, Fe, and Ni). The data in the 1.83–1.85 keV band were ignored
because of the current calibration errors of the XIS; there is a small energy gap (∼10 eV) at the
Si K-edge that is not implemented in the current response function. Since the absolute gain of the
XIS has an uncertainty (especially for low-energy X-rays), we allowed a small offset in the energy
scale, independently between the FI and BI.
This 2-component model gives nice fits to the soft X-ray band including all the emission lines
and the hard continuum emission with a best-fit kT e value of the thermal component of 0.38 keV
and a power-law index of 2.9. However it cannot reproduce the Fe-K line flux, with the unacceptableχ2 /d.o.f. of 1222/1054. As given in subsection 7.3.2, the center energy of the Fe-K line is slightly
but significantly higher than that of Fe I, which suggests that the origin of the Fe-K line is due to a
low ionized thin-thermal plasma, but is not the fluorescence of neutral Fe. We hence added anotherthin-thermal component with a higher temperature and a low ionized parameter. The abundancesin this plasma were fixed to solar, except for Fe. Then, the best-fit χ 2 /d.o.f. was significantly
reduced to be 1093/1050. The best-fit parameters and models are respectively shown in table 7.2,
and in figure 7.7 with solid lines.
We next fitted the spectra of the NE region. We can see emission lines in the soft X-ray bandin these spectra. This suggests that the NE also includes a low-temperature plasma. However,the line profiles are too weak to obtain a reasonable fit with many free parameters of the lowtemperature plasma. We therefore fixed the soft plasma parameters as those of the East region.

7.3. ANALYSIS AND RESULTS

93

Fig. 7.7.—: (a) 0.4–10 keV spectra of the East region. Black and red represent the FI and BI, respectively.
The components of the best-fit model for the FI are shown with the solid lines: green, light blue, and blue
are respectively low-temperature plasma (VPSHOCK 1 in table 7.2, high-temperature plasma (VPSHOCK 2
in table 7.2), and power-law. (b) the spectra of the NE region. The colors are same to (a). Since these
spectra have no significant Fe-K line, the high-temperature plasma component did not be added.

94

CHAPTER 7. RCW 86
Table 7.2:: Best-fit parameters of full-band model fittings.

Component
Absorption
VPSHOCK 1Parameter
NH (cm−2 )kTe (keV)
C, N, ONe
Mg
Si, S
Ca, Fe, Niτ (cm−3 s)
ne nH V (cm−3 )kTe (keV)Feτ (cm−3 s)
ne nH V (cm−3 )ΓNorm∗Offset (eV)
Offset (eV)East
4.1 (3.8–4.4) ×1021
0.33 (0.30–0.37)
0.15 (0.11–0.24)
0.33 (0.26–0.50)
0.31 (0.28–0.48)
0.31 (0.19–0.43)
0.23 (0.17–0.36)
7.7 (5.5–11) ×1010
3.7 (2.4–5.8) ×1057
VPSHOCK 2
1.8 (1.4–2.3)
32 (14–90)
2.3 (1.8–2.9) ×109
2.0 (1.1–2.6) ×1056
Power-law
2.89 (2.45–3.28)
1.1 (0.72–1.9) ×10−3
Gain (FI)–3.2
Gain (BI)–5.4
2
χ /d.o.f.
1093/1050 = 1.04
∗ The differential flux (photons cm −2 s−1 ) at 1 keV.

NE
4.4 (4.2–4.5) ×1021
1.2 (0.89–1.6)
0.15 (fixed)
0.33 (fixed)
0.31 (fixed)
0.31 (fixed)
0.23 (fixed)
6.7 (5.9–7.7) ×109
2.4 (2.0–2.9) ×1056––––
2.78 (2.74–2.82)
3.3 (3.2–3.5) ×10−3
–3.2
–5.4
1138/1084 = 1.05

Since no Fe-Kα was found (see subsection 7.3.2), we did not include the thermal component withhigher temperature. Thus, normalizations of the low-temperature plasma and power-law component
are free parameters. To obtain the best-fit, we also fine-tuned a temperature and an ionization
parameter in the low-temperature plasma and the photon index in the power-law component. The
best-fit χ2 /d.o.f. is an acceptable value of 1138/1084. The best-fit parameters and spectra are given
in table 7.2 and figure 7.7.
The thermal plasma parameters derived for the NE differ from the values reported by Vink
et al. (2006); we find a higher τ and a lower kT e . Indeed, kTe and τ are correlated in the sensethat both a higher kTe or a higher τ give a higher ionization. Note that Vink et al. (2006) useda different spectral fitting code (SPEX). Both fits agree about the ionization state of the plasma,
but differ concerning the relative contributions of kT e and τ . For the NE region, the exact value
for kTe is difficult to determine, since most of the continuum has a non-thermal origin.

7.4. DISCUSSION FOR RCW 86

7.4

95

Discussion for RCW 86

From morphology and spectral analyses, the X-rays in the northeastern quadrant of RCW 86
can be separated into three distinct components: the low temperature (VPSHOCK 1) and high temperature (VPSHOCK 2) thermal components both in ionization non-equilibrium, and the power-law
component. We separately discuss the origin of these components, and finally try to give a unified
picture about them.

7.4.1

Low Temperature Plasma

As shown in figures 7.4 and 7.5, the soft X-rays are dominated in the outer shell of theEast region. The spectrum is well represented by a thin-thermal plasma with a low temperatureof kTe ∼ 0.3 keV. Since the Hα filament was reported to be localized near at the East region(figure 7.1: Smith 1997), the low-temperature plasma (VPSHOCK 1) is likely to be due to theinterstellar medium (ISM) heated by a blast wave.
The metal abundances that we obtained are significantly lower than the solar values. A subsolar abundance is also reported in the rim region of the Cygnus Loop (e.g., Miyata et al. 2007),
where an ISM component is dominant. Therefore, shocked ISM spectra may generally show such
low abundances. In the case of RCW 86, Bocchino et al. (2000) found the thermal spectra of thesoutheast and north rim show lower metal abundances than the solar values, and claimed that
this result is due to metal depletion behind the shock. Since the abundances of our result arealmost consistent with those of Bocchino et al. (2000), the same interpretation would be accepted.
However, we should note that the complexity of our three-component model makes it difficult todetermine the absolute values of the abundances. Indeed, if we fix the Ne abundance to be the solar,
and fit the full band spectra of the East region, the best-fit can be obtained with an acceptableχ2 /d.o.f. (1103/1051), and almost the same relative abundances to those of table 7.2. An exact
separation of the low and high-temperature components and the power-law component is necessary
for a certain determination of the absolute metal abundances.
The solid angle of the East ellipse in figure 7.4 (= 41 arcmin 2 ) corresponds to 2.5×1038 cm2
at 2.8 kpc. Assuming the depth of the emitting region to be 8.5 pc, about half of the radius of
the remnant, the emission volume is estimated to be V = 6.6 × 10 57 cm3 . Therefore, the emission
measure (EM = ne nH V ) of the VPSHOCK 1 in this region corresponds to a proton density of−0.5
nH = 0.68 fE1
cm−3 , where fE1 is the filling factor for this component. Since the soft thermal
emission concentrates at the outer thin rim in the East ellipse (see figure 7.4), the density of the
low-temperature plasma is likely to be much larger than 0.68 cm −3 .
Similarly, the density of the VPSHOCK 1 in the NE region is estimated to be n H =−0.5
0.26 fN
cm−3 , where fN 1 is the filling factor for the VPSHOCK 1 in the NE. Although sev1
eral uncertainties remain (e.g., the filling factors, the temperature difference between the East andNE), the density of the swept-up ISM may be larger at the East rim than at the NE rim. The
variation of the ionization parameter is consistent with this result, namely τ (= n e t) in the East is

96

CHAPTER 7. RCW 86

larger than that of the NE.

7.4.2

Fe-K Line and High Temperature Plasma

By imaging (subsection 7.3.1) and spectral (subsection 7.3.2) analyses, we proved that the
Fe-Kα emission has no correlation with the hard X-ray continuum. Therefore, the origin of the
Fe-Kα line is not due to fluorescence caused by supra-thermal electrons, nor non-thermal X-rays.
It is located behind the blast wave (VPSHOCK 1) and is over-abundant in Fe. We hence proposethat the origin of a high-temperature plasma (VPSHOCK 2) is a Fe-rich ejecta heated by a reverse
shock.
The ionization parameter in VPSHOCK 2, determined from the center energy of Fe-Kα, isextremely low, τ ∼ 2.3 × 109 cm−3 s. The electron density of the ejecta component is estimated−0.5
to be ne = 0.19 fE2
cm−3 , where fE2 is the filling factor for the VPSHOCK 2 in the East
region. We then estimate the elapsed time since the ejecta was heated by the reverse shock:
0.5 sec < 380 yr. This is much less than the age of ∼1800 yr (from thet = τ /ne = 1.2 × 1010 fE2∼
historical record), and hence the ejecta must have been heated very recently.
0.5 M , which is reasonable as a portion of the ejecta,
The estimated mass of Fe is ∼ 0.07fE2
and some of the Fe-ejecta may still remain in the interior of the remnant, which has not yet beenheated by the reverse shock. However, we have to note the uncertainty in the mass determination.
Although we show in subsection 7.3.3 that the spectrum of the East region is well reproduced by
the three-component model, we cannot conclude this model is unique. We therefore, estimated thesystematic error on the Fe mass by fitting the 3–10 keV spectrum with (1) only one VPSHOCK
model (without a power-law), or (2) a pure Fe plasma and a power-law. In both cases, we fixed
the values of kTe and τ to 1.8 keV and 2.3 × 109 cm−3 s, the best-fit parameters given in table 7.2.

Fig. 7.8.—: 3–10 keV spectra of the East region fitted with only one VPSHOCK component (1), and apure Fe plasma VPSHOCK and a power-law (2). The colors are same to figure 7.7. The Fe mass is estimated
0.5
0.5
to be 0.06fE2
M and 0.6fE2
M for cases of (1) and (2), respectively.

7.4. DISCUSSION FOR RCW 86

97

The best-fit spectra for both cases are shown in figure 7.8. The Fe mass was, then, obtained to be
0.5 M and 0.6f 0.5 M for cases of (1) and (2), respectively. These values can be regarded as
0.06fE2
E2
the lower and upper limits of the Fe mass in the East region. Thus, there is an uncertainty of oneorder of magnitude.

7.4.3

Non-Thermal Emission

Since the non-thermal X-rays are well represented by a power-law with photon indices of∼ 2.8, synchrotron radiation is the most plausible origin. The synchrotron filament of the NE rim
is spatially connected from the soft thermal rim (the East) along the outer shell of RCW 86 (seefigure 7.2). Therefore, this filament may be produced by the blast wave.
In the case of the SW rim of RCW 86, Rho et al. (2002) and Ueno et al. (2007) suggested thatthe non-thermal filament is not a blast wave, but reverse shock, because the non-thermal filamentis inside the soft thermal emission on the SNR shell. Also the blast-wave velocity at the SW shell(∼800 km s−1 : Rosado et al. 1996; ∼600 km s−1 : Ghavamian et al. 2001) is too slow to accelerate
electrons by the standard diffusive shock acceleration theory, and to emit intense synchrotron Xrays. Since the gas density of the NE shell is ∼10-times lower than that of the SW shell (Pisarski
et al. 1984), it is possible that the velocity of the NE blast wave is much higher than that of thepresent SW shell. Vink et al. (2006) interpreted that the velocity of the NE non-thermal filamentshould be ∼2700 km s−1 . This high speed can accelerate electrons to more than TeV energy to
produce synchrotron X-rays in the blast wave shock.

7.4.4

Unified Picture for All the Components

We propose a unified picture to explain all of the results in figure 7.9. A few hundred years ago,
blast waves in the East and NE rims would be expanding with the same high velocity. However
the blast wave in the East collided with a dense medium very recently, and the forward shock
decelerated rapidly. At that same time, the reverse shock behind the East rim began to moveinward to the interior of the remnant, and it heated the Fe-rich ejecta. Since RCW 86 may be aremnant in the OB association (Westerlund 1969), a candidate for the “dense medium” is either acavity wall surrounding the remnant, as suggested by Vink et al. (1997), or a molecular cloud. We
found in the NANTEN results (Matsunaga et al. 2001) evidence of a cloud just at the position of
the East rim of RCW 86 (l=315.7, b=–2.4), as shown in figure 7.10. We strongly propose a deeper
radio observation to reveal whether the shell of the remnant is truly interacting with the cloud.
The forward shock in the NE rim, on the other hand, is still expanding in a tenuous region,and hence keeps a high shock velocity. The reverse shock behind the NE may also expand withhigh velocity, and hence has not yet reached the Fe-rich ejecta layers. This may explain the absenceof Fe-K emission behind the NE region.

98

CHAPTER 7. RCW 86

Fig. 7.9.—: Schematic view of the interpreted structure of RCW 86 NE region.

Fig. 7.10.—: NANTEN CO (J = 1–0) map of the RCW 86 northeast region (gray scale: Matsunaga et
al. 2001), and XIS 0.5–1.0 keV image (color), shown in Galactic coordinates.

7.4. DISCUSSION FOR RCW 86

7.4.5

99

Summary for RCW 86

We have analyzed Suzaku/XIS data obtained from the northeast quadrant of RCW 86. The
results and interpretations are summarized as follows:
1. The Fe-Kα emission is enhanced at the inner region from the soft thermal rim (the East
region), and has no correlation with the hard X-ray non-thermal filament (the NE region).
2. The centroid of the Fe-Kα line (6404–6444 eV) and the lack of Fe-L lines constrain the Fe
ionization state to be approximately between Ar-like (Fe IX) and Ne-like (Fe XVII).
3. The spectra extracted from the soft X-rays and Fe-Kα enhanced region (the East) are wellfitted with two thin-thermal plasmas, which have different temperature of kT e ∼ 0.3 keV and∼ 1.8 keV and one power-law with a photon index of Γ ∼ 2.9.
4. The spectra extracted from the hard X-ray filament (the NE) can be represented by a lower
temperature plasma and a power-law.
5. The lower temperature plasma has sub-solar metal abundances. The origin of this component
would be a blast-shocked ISM.
6. The higher temperature plasma includes an over-solar abundant iron, and has an extremely
low plasma age of τ = ne t ∼ 2.3 × 109 cm−3 s. Hence it produces nearly 6.4 keV Kα lines.
This component may be Fe-rich ejecta heated by reverse shock very recently.
7. The power-law component, which is the brightest in the NE region, can be regarded assynchrotron emission.
8. The shock front at the East rim would recently have collided with dense ISM, after which theblast wave decelerated rapidly, and the reverse shock began to move toward the center of theremnant.
9. The blast wave at the NE rim may still be expanding in a tenuous region. Therefore, it keeps
a high shock velocity for efficient particle acceleration.

Chapter 8

Fe-K Emission in Other SNRs
In this chapter, we analyze the Suzaku data of Tycho and N103B, which are also young but
slightly more evolved than the SNRs discussed in the previous chapters. Since both of Tycho
and N103B show very strong emission lines, their spectral properties and morphologies of theintermediate-weight elements (i.e., Si or S) have been relatively well-studied. This chapter, therefore, concentrates on the studies of their spectra in the energy band above 5 keV including Fe-K
emission. As proved in the previous chapters, Suzaku is superior to the other X-ray missions inthe sensitivity for the Fe-K band. Utilizing this merit, we perform the Fe-K spectroscopy with thehighest quality which had ever been obtained.

8.1

Tycho

8.1.1

Previous Results

The Tycho’s SNR (G120.1+1.4, 3C10) recorded by Tycho Brahe in 1572 is generally considered
to be a historical Type Ia SNR (e.g., Baade 1945). The distance to Tycho was determined tobe 2.3 kpc by Green (1984); we assume this value in this section. The average energy index
in radio band is measured to be α = 0.5 (Katz-Stone et al. 2000). From the the width of thebroad Hα emission, Ghavamian et al. (2001) derived the blast wave velocity of the east rim to be∼ 2000 km s−1 . However, the X-ray proper motion study by Hughes (1996; 2000) using ROSAT
showed the average blast wave velocity is much higher, 4600 ± 400 km s −1 . FCRAO and Nobeyama
CO data showed that a large molecular cloud surrounds the SNR along the northeastern boundary(Lee et al. 2004). Therefore, we can consider that the blast wave in this region is decelerated bythis molecular cloud.
Hwang & Gotthelf (1997) found many emission-line features of Mg, Si, S, Ar, Ca, and Fe,
using data obtained by ASCA. All narrow band images of these emission lines showed the shell-likemorphology, but they are clearly distinct from each other. Especially, the radial profile of the Fe-K

101

102

CHAPTER 8. FE-K EMISSION IN OTHER SNRS

line peaks at a smaller radius than the those of the other emission lines and continuum emission.
This fact suggests an incomplete mixing of the Si and Fe layers. The average ionization age of
the Si and S ejecta is reported to be (0.8–1.3) × 10 11 cm−3 s (Hwang & Gotthelf 1997), whereasthat of Fe ejecta is suggested to be much lower (∼ 10 9 cm−3 s: Hwang et al. 1998). The highresolution XMM-Newton images confirmed that the Fe-K emission clearly peaks at a smaller radius(Decourchelle et al. 2001). Furthermore, the XMM-Newton image showed the contact discontinuity
between the ejecta and ISM is partly distorted by the Rayleigh-Taylor instability.
The cosmic ray acceleration has also been studied for a long time. Hard X-rays were detected
at energies up to 20 keV with Ginga (Fink et al. 1994); and 30 eV with RXTE (Petre et al. 1999).
Using Chandra data, Hwang et al. (2002) showed the spectra at the SNR rim can be describedby a non-thermal model, but could not reject the possibility of the thermal origin. Bamba etal. (2005) discovered very thin filaments on the outer edge of the SNR, and concluded they havethe non-thermal origin.

8.1.2

Suzaku Observation

The Suzaku observation of the Tycho’s SNR was made on 2006 June 27. The observation
ID and exposure time were 500024010 and ∼ 102 ks, respectively. In addition, the backgroundobservation was made on 2006 June 29 (observation ID = 500025010) with the targeted position at(RA, Dec) = (00:36:54, 64:17:42), the nearby sky of Tycho. The exposure time was ∼ 51 ks. The
Galactic latitudes of the source and background regions are almost consistent with each other.
The XIS were operated in the normal full-frame clocking mode, during both of the observations.
We employed the cleaned revision 1.2 data, and used the HEADAS software version 6.0.4 andXSPEC version 11.3.2 for the data reduction and analysis. The response matrix files (RMF) andancillary response files (ARF) were made using xisrmfgen and xissimarfgen (Ishisaki et al. 2007)
version 2006-10-17.

8.1.3

Analysis and Results

Figure 8.1 shows the XIS 5.0–8.0 keV image of the Tycho’s SNR (left), and the backgroundregion (right). First, we extracted the spectrum of the entire SNR from a 4. 0 3-radius circular region
shown with a large green circle in figure 8.1 (left). For the background, we used the spectrum takenfrom the elliptical region shown in figure 8.1 (right). Since the data of the three FIs are nearly
identical, we merged the individual spectrum to improve the photon statistics. The backgroundsubtracted full-band spectra are shown in figure 8.2.
Although Kα and Kβ lines of intermediate-mass elements (Mg, Si, S, Ar, and Ca) are clearly
divided, we concentrate on the spectral analysis in the energy band above 5 keV where Fe-K lines
are included. In this energy band, Cr-Kα and Fe-Kβ as well as Fe-Kα are detected. Therefore,
we study the features of these lines by fitting the 5–12 keV spectra with a power-law (for the103

8.1. TYCHO

30:00.0

20:00.0

25:00.0

15:00.0

NW
64:20:00.0

10:00.0

15:00.0

05:00.0

Center

63:55:00.0

27:00.0

30.0

26:00.0

30.0 0:25:00.0

BGD

10:00.0

64:00:00.0

30.0

24:00.0 23:30.0

05:00.0

30.0

38:00.0

30.0 0:37:00.0

30.0

36:00.0

30.0

35:00.0

Fig. 8.1.—: Left: XIS 5.0–8.0 keV image of the Tycho’s SNR. The large circle (green), the small circle(red), and the ellipse (yellow) show the regions where we extracted the spectra. The coordinates (RA andDec) refer to epoch J2000.0. The Field of view of the XIS is indicated by the black square, where the datafrom the four XIS are combined, but that of the four corners, irradiated by the 55 Fe calibration sources, areremoved. Right: XIS 5.0–8.0 keV image of the background sky. The green ellipse show the region for the
background spectrum.

Fig. 8.2.—: Background-subtracted XIS spectra extracted from the whole region of the Tycho’s SNR.
Black and red represent the FI and BI, respectively.

104

CHAPTER 8. FE-K EMISSION IN OTHER SNRS

Table 8.1:: Center energies and widths of the emission lines in the 5–12 keV spectrum of the Tycho’s SNR.

Line
Center energy (eV) Width∗ (eV)Cr-Kα 5480 (5464–5496)
0Fe-Kα
6444 (6443–6446)
61 (59–62)
Fe-Kβ7104 (7087–7118)
0∗ One standard deviation (1σ).

continuum) and Gaussians (for the lines). The best-fit central energies and widths are given intable 8.1. Although the central energy of the Fe-Kα is marginally higher than those of the SN 1006
and E0509–67.5 spectra (chapter 5 and 6), it is still much lower than that of He-like Fe-Kα (seetable 3.1). Since an atomic physics for ionized Cr has not yet been well-studied at present, theionization state of the Cr cannot be determined from the central energy. However, the ionization
state must be extremely low, because the central energy of the Cr-Kα in the Tycho spectrum isfound to be significantly lower than that of a Cr-Kα in a spectrum of N103B (see section 8.2), of
which the ionization age is much higher than that of Tycho.
Ionization Age of the Fe Plasma
In order to estimate the ionization age of the plasma which produces Fe-K emission, we fitted
the 5–12 keV spectra using a VNEI model. As seen in the previous sections, continuum fluxes
above 5 keV of young SNRs are often explained by a non-thermal emission. Also in the case of
Tycho, there is a high possibility that the origins of the hard continuum and the line emissions
are different from each other. Indeed, bright non-thermal filaments had been discovered from theforward shock region of Tycho (Hwang et al. 2002; Bamba et al. 2005). Therefore, we tried to fit
the continuum with a power-law, and the Fe-K line with a pure Fe plasma (a VNEI model). Since
we showed in previous chapters Fe-K lines are typically reproduced by a 2–4 keV NEI plasma, we
now fixed the electron temperature of the VNEI component to 3 keV, as the assumption. Since
the NEI spectral code does not include the K-shell lines from Cr and Kβ of Fe in lower ionization
states than He-like, we added Gaussians at 5.48 keV and 7.10 keV (see table 8.1).
Then, the best-fit χ2 /d.o.f. was obtained to be 2081/1124. The fitted spectra are shown infigure 8.3a. We can see the significant disagreement between the data and model around the Fe-Kαline. Anticipating the result shown below, we found that this disagreement can be explained by thesuperposition of red- and blue-shifted lines due to proper motions of the Fe ejecta. Hence, we fitted
Fe-Kα line with 2-component VNEI models, which have the same values of red and blue shifts aswell as the same values of the other parameters (i.e., kT e , ne t, and emission measure). Then, thefit was significantly improved with χ 2 /d.o.f. of 1275/1123. The best-fit parameters are given intable 8.2. Note that, in this table, the emission measure is given as n e nFe V for the sum of thered- and blue-shifted components, where V is the emitting volume. The fitted spectra are shown105

8.1. TYCHO

Fig. 8.3.—: (a) XIS 5–12 keV spectra of the whole region of the Tycho’s SNR, fitted with the onecomponent VNEI, power-law, and Gaussians (see text). Black and red represent the FI and BI, respectively.
(b) Same to (a), but fitted using the two-component VNEI. The red- and blue-shifted VNEI components areshown by magenta and blue solid lines, respectively.

Table 8.2:: Best-fit parameters of the fitting for the 5–12 keV spectrum of the Tycho’s SNR.

Component
Power-law
VNEI (Fe)

Cr-Kα
Fe-Kβ

ParameterΓNorm∗ (photons cm−2 s−1 )kTe (keV)ne t (cm−3 s)z†
ne nFe V ‡ (cm−3 )Center Energy (keV)
Norm§ (photons cm−2 s−1 )Center Energy (keV)
Norm§ (photons cm−2 s−1 )Value
2.70 (2.67–2.74)
7.5 (7.1–7.9) × 10−2
3.0 (fixed)
8.0 (7.8–8.2) × 109
7.9 (7.8–8.3) × 10−3
1.20 (1.19–1.21) × 1055
5.48 (fixed)
1.8 (1.6–2.2) × 10−5
7.10 (fixed)
2.2 (2.0–2.6) × 10−5
1275/1123 = 1.14

χ2 /d.o.f.
∗ The differential flux at 1 keV.
† Red and blue shifts.
‡ Sum of two components of the VNEIs (see text).
§ Total flux in the Gaussians.

106

CHAPTER 8. FE-K EMISSION IN OTHER SNRS

in figure 8.3b, in which the best-fit model is shown as solid lines colored magenta (the red-shifted
VNEI), blue (the blue-shifted VNEI), and black (the other components).
The ionization parameter of the Fe-K emitting plasma is obtained to be n e t ∼ 8 × 109 cm−3 s.
In order to investigate how ne t depends on kTe , we fixed kTe to several values (other than 3 keV)and fitted the spectra. However, we found that the value of n e t does not change so much, whenkTe is higher than 2 keV.
If we fit the continuum with a bremsstrahlung model instead of the power-law, the good
fit was also obtained with the electron temperature and emission measure of the bremsstrahlungcomponent of kTe = 5.1 keV and ne nH V = 2.2 × 1058 cm−3 , respectively. The values of the otherparameters (in the VNEI or Gaussian components) did not change from those listed in table 8.2.
Evidence of Doppler Shift
We extracted spectra from two regions: the Center and Northwest (NW) represented in figure 8.1 (left). For the background, we used the spectra taken from the elliptical region shown infigure 8.1 (right) again. The background-subtracted 5–12 keV spectra of the Center and NW are
shown in figures 8.4a and 8.4b, respectively. We can see that the width of the Fe-Kα line is larger in
the Center spectrum (a) than in the NW spectrum (b). In order to confirm it quantitatively, we fit
these spectra with the same model what we used for the whole SNR spectrum. Then, the best-fit
models are obtained as shown by the solid lines in figure 8.4. The obtained red and blue shift (z)for the Center spectra was 1.1 (1.0–1.2) × 10 −2 . This value corresponds to the expansion velocityof ∼ 3300 km s−1 . On the other hand, z for the NW spectra was obtained to be 6.8 (5.3–7.5)
× 10−3 , which is significantly lower than the value for the Center. Therefore, we can conclude
that the shock-heated Fe ejecta is now expanding with the velocity of ∼ 3300 km s −1 , and the line

Fig. 8.4.—: XIS 5–12 keV spectra of the Center (a) and NW (b) of the Tycho’s SNR. Black and red
represent the FI and BI, respectively. Magenta and blue solid lines show the red- and blue-shifted VNEI
components.

8.2. N103B

107

broadening is due to the superposition of the Doppler shifts.

8.2

N103B

8.2.1

Previous Results

N103B is the fourth brightest X-ray SNR in the LMC. It was first identified as a SNR by
Mathewson et al. (1983) based on the non-thermal nature of its radio spectrum and the relative
strengths of Si II and Hα emission lines. From its small size (3 pc-radius), Hughes et al. (1995)
estimated the age to be ∼ 1500 yrs. Chandra data showed that the remnant is much brighter on
the western rim than toward the East (Lewis et al. 2003), this is possibly due to density contrasts
between the East and West.
The remnant is located on the northeastern edge of an H II region DEM 84, ∼ 40 pc from theyoung star cluster NGC 1850. Furthermore, the optical spectrum is not dominated by the strong
Balmer lines which are often associated with the remnants of Type Ia SN (Tuohy et al. 1982).
For long years, therefore, it was considered that N103B had a massive star progenitor (e.g., Chu& Kennicutt 1988). However, Hughes et al. (1995) found strong emission lines from Si, S, Ar,
Ca, and Fe, while no emission from O, Ne, or Mg in ASCA spectrum. They hence concluded
that N103B is the result of a Type Ia SN. On the other hand, the high energy resolution spectraof the Reflection Grating Camera (RGS) on board XMM-Newton showed prominent O, Ne, andMg emission (van der Heyden et al. 2002). They hence claimed that N103B was the result of acore-collapse SN. However, Lewis et al. (2003) analyzed Chandra data and considered that the O
emission was associated with the ISM rather than the ejecta. Thus, the classification of N103B isstill somewhat uncertain.

8.2.2

Suzaku Observation

The Suzaku observation of the SNR N103B was made on 2005 August 30 with the targeted
position at (RA, Dec) = (05:08:58, –68:43:10). The observation ID and exposure time are 100013010
and ∼ 33 ks, respectively. The XIS were operated in the normal full-frame clocking mode. We
employed the cleaned revision 1.2 data, and used the HEADAS software version 6.0.4 and XSPEC
version 11.3.2 for the data reduction and analysis. The response matrix files (RMF) and ancillary
response files (ARF) were made using xisrmfgen and xissimarfgen (Ishisaki et al. 2007) version
2006-10-17.

8.2.3

Analysis and Results

Since the diameter of the SNR (∼ 3000 ) is much smaller than the angular resolution of the
XRT, (∼ 20 : Serlemitsos et al. 2007), N103B can be regarded as a point-like source. Therefore, we108

CHAPTER 8. FE-K EMISSION IN OTHER SNRS

Fig. 8.5.—: (a) Background-subtracted XIS spectra of N103B. Black and red represent the FI and BI,
respectively. (b) Same to (a), but in 1.5–5.0 keV.
Table 8.3:: Center energies and widths of the emission lines in the 5–10 keV spectrum of N103B.

Line
Center energy (eV) Width∗ (eV)Cr-Kα 5600 (5520–5665)
0Fe-Kα6550 (6539–6561)
95 (81–105)∗ One standard deviation (1σ).

extracted XIS spectra from a circular region of 3 0 radius centered on the SNR. The background
spectra were extracted from the annular region around the source with inner and outer radii of 3 0
and 60 , respectively. Since the data of the three FIs are nearly identical, we merged the individual
spectrum to improve the photon statistics. The background-subtracted full-band spectra are shownin figure 8.5a. Figure 8.5b shows the same spectra but in the 1.5–5.0 keV band, where Lyα linesof H-like Si and S as well as K-shell lines of the He-like ions are clearly seen. This implies theionization age of the plasma in N103B is older than those in SN 1006, E0509–67.5, and Tycho,
because the spectra of these SNRs show no feature of the H-like lines of the intermediate-masselements (see figure 8.2, for example).
Similarly to section 8.1, we concentrate on the spectral analysis in the energy band above
5 keV. The XIS 5–10 keV spectra are shown in figure 8.6. In addition to Fe-Kα, we detected the
Cr-Kα line, for the first time. We fitted the 5–10 keV spectra with a power-law (for the continuum)
and Gaussians (for the emission lines), and then determined the center energies and the widths ofthe lines as given in table 8.3. We found the Kα of both of Cr and Fe have higher central energies
than those of the Tycho spectrum (for comparison, see table 8.1).

8.2. N103B

109

Ionization Age of the Fe Plasma
We fitted the 5–10 keV spectra using a VNEI model to determine the ionization age of the
Fe-K emitting plasma. Similarly to the previous section, we included a power-law to reproduce thecontinuum flux, and fitted the Fe-K line with a VNEI model which consists purely of Fe ions andelectrons. The electron temperature of the VNEI was assumed to be 3 keV. Since the NEI code
does not include Cr-Kα lines, we added Gaussians at 5.60 keV.
Then, the best-fit χ2 /d.o.f. was obtained to be 57/45. Although the fit was acceptable froma statistical point of view, we found that the model could not reproduce the Fe-Kα profile, as
shown in figure 8.6a. On the analogy of the result obtained in section 8.1, we considered this line
broadening is due to the superposition of Doppler shifts. We, therefore, fitted the Fe-Kα line with
2-component VNEI models, which have the same values of red and blue shifts as well as the samevalues of the other parameters. Then, the fit was significantly improved with χ 2 /d.o.f. = 22/44.
The best-fit parameters are given in table 8.4. Note that emission measure is given as n e nFe V
for the sum of the red- and blue-shifted components, where V is emitting volume. The fitted
spectra are shown in figure 8.6b, in which the best-fit model is shown as solid lines colored magenta(the red-shifted VNEI), blue (the blue-shifted VNEI), and black (the other components). The red(or blue) shift was obtained to be z ∼ 0.01, which corresponds to the proper motion velocity of∼ 3000 km s−1 . Note that, however, the spectrum is what integrated with the emission from thewhole SNR. Therefore, the actual expansion velocity (of the shock-heated Fe-rich ejecta) is possibly
higher than 3000 km s−1 .
We fitted the spectra fixing kTe to several values other than 3 keV, and investigated thedependence of ne t to kTe . However, we found that the value of n e t does not change so much, when

Fig. 8.6.—: (a) XIS 5–10 keV spectra of N103B, fitted with the one-component VNEI, power-law, andGaussian. Black and red represent the FI and BI, respectively.
(b) Same to (a), but fitted using thetwo-component VNEI. The red- and blue-shifted VNEI components are shown by magenta and blue solidlines, respectively.

110

CHAPTER 8. FE-K EMISSION IN OTHER SNRS
Table 8.4:: Best-fit parameters of the fitting for the 5–10 keV spectrum of N103B.

Component
Power-law
VNEI (Fe)Cr-KαParameterΓNorm∗ (photons cm−2 s−1 )kTe (keV)ne t (cm−3 s)z†
ne nFe V ‡ (cm−3 )Center Energy (keV)
Norm§ (photons cm−2 s−1 )Value
3.1 (2.4–4.0)
1.7 (0.55–9.1) × 10−3
3.0 (fixed)
3.5 (3.2–4.0) × 1010
1.1 (0.93–1.2) × 10−2
1.5 (1.3–1.7) × 1054
5.60 (fixed)
9.3 (2.0–16) × 10−7
22/44 = 0.50

χ2 /d.o.f.
∗ The differential flux at 1 keV.
† Red and blue shifts.
‡ Sum of two components of the VNEIs (see text).
§ Total flux in the Gaussian.

kTe is higher than 2 keV.
If we fit the continuum with a bremsstrahlung model instead of the power-law, the good
fit was also obtained with the electron temperature and emission measure of the bremsstrahlungcomponent of kTe = 3.5 keV and ne nH V = 4.3 × 1058 cm−3 , respectively. The values of the otherparameters (in the VNEI or Gaussian components) did not change from those listed in table 8.4.

8.3

Summary of the Chapter

We have analyzed Suzaku/XIS data obtained from the Tycho’s SNR and N103B. The newfacts revealed in this chapter are summarized as follows:
1. The ionization ages of the Fe-K emitting plasmas in the Tycho’s SNR and N103B wererespectively determined to be ne t = 8.0 (7.8–8.2) × 109 and ne t = 3.5 (3.2–4.0) × 1010 , which
are the tightest constraints obtained so far.
2. Cr-Kα lines were detected from both the SNRs, for the first time.
3. Evidences of Doppler shift due to the expansion of the Fe-rich ejecta were found from Fe-Kαlines of the spectra of both the SNRs.

Chapter 9

Discussion
9.1

Summary of the Results

In the previous chapters, we analyzed the X-ray data of several young SNRs (SN 1006, E0509–
67.5, RCW 86, Tycho, and N103B) obtained with the newest X-ray astronomical satellite Suzaku.
Thanks to the superior performances of Suzaku, especially the sensitivity and the spectral resolution
in the Fe-K band, we successfully demonstrated a lot of new results. From SN 1006, we firmly
detected the K-shell emission from Fe, for the first time, and found that the Fe ionization state isquite low. We accurately determined the ionization age of Fe ejecta in E0509–67.5, and showed
it was significantly lower than what was previously reported. The Fe-K line morphology in the
RCW 86 northeastern region was revealed for the first time, and the origin of the Fe-K emission
was determined to be the Fe-rich ejecta. From the Fe-K spectra of Tycho and N103B, we found theevidences of Doppler shift due to the ejecta expansion. Furthermore, we detected the Cr-K lines
from these SNRs, for the first time.
In particularly the analyses of SN 1006 and E0509–67.5, we performed the full-band fittings
by logically dividing the spectra into several components: the ejecta (heated by reverse shock), theISM (heated by the blast wave), and the non-thermal component. As the results, we found that theionization ages (ne t) of Fe-rich ejecta (in both the SNRs) are significantly different from those of theejecta abundant with the other elements (i.e., Si, S). Here, we summarize in table 9.1 the physicalproperties of all the SNRs which we observed. In this table, n e t[Si] and ne t[Fe] respectively means
the ionization ages of Fe-rich and Si-rich ejecta, and n 0 indicates the average ambient density. The
values which are not obtained in this thesis referred to other papers (see footnote in table 9.1). We
will use these values to develop further discussions in sections 9.3 and 9.4.

111

112

CHAPTER 9. DISCUSSION
Table 9.1:: Summary of the results.

ne t[Si] (cm−3 s)
ne t[Fe] (cm−3 s)n∗0 (cm−3 ) Age (yr)SN 1006
1.4 (1.2–1.6) × 1010
7.7 (6.7–9.2) × 108
0.03
1000
10
< 0.05†< 850†
E0509–67.5 1.4 (1.3–1.5) × 10
3.5 (3.1–3.6) × 109
∼
∼
Tycho
1.0 (0.8–1.3) × 1011 ‡ 8.0 (7.8–8.2) × 109
1.0§
430
12
k
10
N103B
> 10
3.5 (3.2–4.0) × 10?
∼ 1500#9∗∗
RCW 86?
2.3 (1.8–2.9) × 10
Various1820∗ Ambient density. † Warren & Hughes (2004); ‡ Hwang & Gotthelf (1997);§ Badenes et al. (2006); k Lewis et al. (2003); # Hughes et al. (1995);∗∗ Spatial variation of the ambient density is quite large.

9.2

Radius (pc)
9.5
3.7
2.8
3.7
18

Nucleosynthesis in Type Ia SNe

SN 1006, E0509–67.5, and Tycho are surely classified as the remnants of Type Ia SN (e.g.,
Schaefer 1996; Hughes et al. 1995; Baade 1945). Indeed, we showed (in sections 5.6 and 6.5) thatthe relative abundances in the ejecta components of SN 1006 and E0509–67.5 are almost consistent
with the predicted nucleosynthesis yield of the carbon deflagration model (W7: Nomoto et al. 1984),
except for the high-Z elements such as Fe (see figures 5.12a and 6.7a). The reason of the lack of the
Fe abundances will be discussed again in section 9.3. The W7 model is widely used as the standard
model for the nucleosynthesis in Type Ia SNe. Strictly speaking, however, we can see that theabundance ratios of Si/O and S/O which we obtained are slightly higher than the predicted values.
These results suggest that the amount of O ejected by Type Ia SN explosion is smaller than thepredicted value by the deflagration model.
A more efficient burning (relative to the deflagration model) of C and O at the surface of awhite dwarf is predicted by the delayed detonation model (e.g., Iwamoto et al. 1999). The delayed
detonation in Type Ia SNe was first proposed by Khokhlov et al. (1991). This theoretical model
assumes a slow initial deflagration at the center of a white dwarf and a deflagration-to-detonationtransition (DDT) at some stage of the thermonuclear explosion. Due to the supersonic detonation
wave, larger amounts of intermediate-weight elements (i.e., Si, S, and Ca) and Fe-group elements
are expected to be synthesized, and only a few light elements (i.e., C, O, and Ne) are left withoutburning (see table 2.2 and figure 2.4). In the case of Tycho, Badenes et al. (2006) found that
the X-ray spectra obtained with XMM-Newton and Chandra are well reproduced by an optimized
delayed detonation model. Therefore, the relative abundances of the elements synthesized by TypeIa SN explosions may be generally similar to what predicted by the delayed detonation model rather
than the pure deflagration model.
However, we should note that the delayed detonation models constructed by Iwamoto etal. (1999) and Badenes et al. (2006) are only one-dimensional models. They did not include any

9.3. EJECTA DISTRIBUTION

113

asymmetrical effects, Rayleigh-Taylor and Kelvin-Helmholtz instabilities, and turbulence, whichshould be essentially taken into consideration. Some simulations have explored delayed detonation
explosions in three dimensions (e.g., Gamezo et al. 2005). However, since the mechanism of the
DDT had not been understood, they artificially determined the central densities (of a white dwarf)
when the DDT occurs. Under the present condition, the delayed detonation is only one of the phenomenological model to explain the abundances of the produced elements obtained by observationsof Type Ia SNe and their remnants.

9.3

Ejecta Distribution

Figure 9.1 shows the plot of the ne t[Si] and ne t[Fe] of each SNR listed in table 9.1. RCW 86 is
not included because it shows no evidence of Si line emission which originates from the ejecta. We
can clearly see that ne t[Fe] are significantly lower than n e t[Si] in the cases of all SNRs. Although
the dispersion is large, ne t[Fe] are approximately consistent with 10% of n e t[Si]. A possible scenario
to explain this fact is: when the Si ejecta began to be heated by a reverse shock, most of the Feejecta had not arrived at the shock front yet. After the volume of the SNRs became about 10 times
larger, the Fe ejecta gradually began to be heated. As a matter of fact, since n e is not uniform
in each SNR and depends on the ionization state of the elements, it is difficult to quantitatively
determine when the Si and Fe were heated. However, it is at least suggested that the Fe-rich ejecta
layers were distributed inward of the Si-rich layers.

Fig. 9.1.—: Ionization parameters of Si ejecta (ne t[Si]) and Fe ejecta (ne t[Fe]) for each SNR. The solidand dashed lines represent ne t[Si] = ne t[Fe] and ne t[Si] = 10 ne t[Fe], respectively.

114

CHAPTER 9. DISCUSSION

SN 1006

heated Si & others

unshocked Fe

heated Fe

reverse shock

Fig. 9.2.—: Schematic view of ejecta and reverse shock structures in SN 1006.

For this interpretation, we can expect that the emission from Fe in these SNRs are spatiallylocated interior to the emission from Si and the other elements. In the case of Tycho, it wasreported that the Fe-K emission is more centrally distributed than the K-shell emission of otherelements (Hwang et al. 1998; Decourchelle et al. 2001). On the other hand, in SN 1006, we could
not find the interior distribution of Fe-K emission, but it was concentrated at the southern part of
the SNR rim (see figure 5.6). This result suggests that most of the Fe ejecta in SN 1006 has not yet
been heated by the reverse shock except for the southern rim. We show in figure 9.2 a schematic
view of the SNR structure derived from our interpretation. Similarly to SN 1006, in E0509–67.5
and Tycho, a large amount of Fe ejecta may remain distinct from other ejecta in the inner layer
and thus be too cool to provide enough X-ray emission. This is the possible reason why the Feabundances we obtained from these young Type Ia SNRs are much lower than the predicted value.
In the theory for a Type Ia SN explosion, the lower expansion velocity of 56 Ni relative tothose of the other elements is expected, as shown in figure 9.3 taken from Iwamoto et al. (1999).
(Note that this figure shows the deflagration case, but the delayed detonation models produce
qualitatively similar distribution.) Although this theoretical expectation is derived from an onedimensional simulation, nevertheless our observational results qualitatively support it. Therefore,
the ejecta mixing effect inside the progenitor star due to the Rayleigh-Taylor instability is suggested
to be almost negligible at least in the Type Ia SN explosion.
By contrast, in several core-collapse SNRs, for example, Cas A (Hughes et al. 2000a; Hwang etal. 2000; 2004), and G292.0+1.8 (Park et al. 2002; 2007), the abundance structure is quite complex;
significant mixing effects are suggested. Figure 9.4 shows an X-ray image of Cas A obtained by
Chandra million-second observation (Hwang et al. 2004). Red, Blue, and Green colors represent
Si-K band (1.78–2.0 keV), Fe-K band (6.52–6.95 keV), and continuum emission (4.2–6.4 keV),
respectively. In the southeastern side of the remnant, we can see that the Fe ejecta lies exterior

9.3. EJECTA DISTRIBUTION

115

Fig. 9.3.—: Composition structure of Type Ia SNe after the freezing of nuclear reactions for the case of
the W7 model. The horizontal axis is the distance from the stellar center indicated by the mass coordinate;
the corresponding expansion velocity is shown at the top of the panel. This figure is taken from Iwamoto etal. (1999).

Fig. 9.4.—: Chandra million-second image of Cas A. Red, Blue, and Green contain emission from 1.78–
2.0 keV (Si-K band), 6.52–6.95 keV (Fe-K band), and 4.2–6.4 keV band (continuum), respectively. North isup and east is to the left. Credit: NASA/CXC/GSFC/Hwang et al. (2004).

116

CHAPTER 9. DISCUSSION

to the lighter Si ejecta; this structure was first reported by Hughes et al. (2000a). Since 56 Ni is
synthesized (by complete Si burning) at deeper region in the progenitor star than Si is synthesized(by O burning), this ejecta distribution requires that a spatial overturning of a significant portionof the SN core. As the source of such large-scale mixing, neutrino-driven convection for initiatingcore-collapse SN explosions is proposed (Herant et al. 1994; Burrows et al. 1995).
The contrast between the ejecta distribution of Type Ia SNRs and that of core-collapse SNRs
may be useful to presume progenitors of the other SNRs. In chapter 7, we found that the Fe-K
emission observed from RCW 86 northeastern region has the ejecta origin. If the lower-Z elements(e.g., O or Si) distribute exterior to the Fe ejecta, we must have observed emission from them.
However, we did not find any evidence of such elements in the ejecta component. Therefore, we
consider that the inversion of the ejecta occurred in RCW 86, similarly to Cas A and G292.0+1.8,
and hence the lower-Z elements have not yet been heated enough. Indeed, since RCW 86 is spatiallycoincident with an OB association (Westerlund 1969), RCW 86 is likely to be the remnant of acore-collapsed SN.
N103B (section 8.2) is also located in a star forming region. Therefore, the progenitor ofthis SNR is still debatable. However, Lewis et al. (2003) found that the Fe-K emission is locatedcentrally, as seen in Tycho (Hwang & Gotthelf 1997), while the Si-K emission is more prominent
in exterior regions of the remnant. Qualitatively, such ejecta distribution is more similar to thatseen in Tycho than that seen in Cas A and G292.0+1.8. Furthermore, as we show in figure 9.1,
the relation between the ne t[Si] and ne t[Fe] is also similar to those of the other Type Ia SNRs (i.e.,
SN 1006, E0509–67.5, and Tycho). For these reason, we strongly propose that N103B is a Type Ia
remnant. However, van der Heyden et al. (2002) claimed that N103B has a Type II origin because
the XMM-Newton spectra showed prominent O-K emission. In this thesis, we analyzed the Suzaku
spectra of N103B only in the energy band above 5 keV, but Suzaku has good sensitivity and energy
resolution also in the energies below 1 keV (Koyama et al. 2007). Utilizing these performances, wemust investigate the amounts of O and other lower-Z elements to determine the origin of this SNR,in the future work.

9.4

ISM Density and Evolutional State of the SNRs

The ages of SN 1006 and Tycho are well-known from their historical records: 1000 yr and
430 yr, respectively. Although the real age of SN 1006 is more than two times older than that of
Tycho, the ionization age (both of the n e t[Si] and ne t[Fe]) is significantly younger in the former thanthe latter, as shown in figure 9.1. This may be due to the extremely low density of ISM aroundSN 1006. According to the equations (1), (2), and (5) of Ferrière (2001), The space-averaged
number densities of H2 (hydrogen in molecular form), H I (neutral atomic hydrogen), and H II
(ionized hydrogen) as the functions of Z-height (the distance from the Galactic plane) are given as
figure 9.5. The distance to SN 1006 (from the solar system) is 2.2 kpc (Winkler et al. 2003), and
the Galactic latitude of this SNR is b = +14.6. They correspond to Z = 550 pc, where the average
ISM density is 0.03 cm−3 . This is fully consistent with the value which we determined in chapter 5

117

9.4. ISM DENSITY AND EVOLUTIONAL STATE OF THE SNRS

Fig. 9.5.—: Space-averaged number densities of hydrogen in the molecular form (H2 : red), neutral atomic
hydrogen (H I: green), and warm ionized hydrogen (H II: blue) as the functions of the distance from the
Galactic plane, according to the approximate equations given by Ferrière (2001). The black curve represents
the sum of the three components. The Z-heights of SN 1006 and Tycho are also indicated.

(also listed in table 9.1). On the other hand, the distance to Tycho is 2.3 kpc (Green 1984), and
its Galactic latitude is b = +1.4, corresponding to Z = 56 pc. Since it is very near to the Galactic
plane, the ISM density of 1 cm−3 is much higher than that around SN 1006. Therefore, it seems
that the ionization of the ejecta in Tycho was quickly developed in the very early phase of theremnant evolution.
As described in section 3.1, the lower an ambient density is, the slower a dynamical evolutionof an SNR becomes. Using the equations given by Truelove & McKee (1999), we calculated the
explosion energies (E), average radii of the reverse shock (R r ), and velocities of the blast wave andreverse shock (vb and vr , for the observer rest frame) of SN 1006 and Tycho. We simply assumed
the uniform density profiles of the ejecta and ambient medium, and included the values of theage, ambient density (n0 ), and blast wave radius (Rb ), listed in table 9.1. The results are given intable 9.2. In the both cases of SN 1006 and Tycho, we found that the reverse shocks still propagate

Table 9.2:: Dynamical parameters of SN 1006 and the Tycho’s SNR.

SN 1006
Tycho

E (1051 ergs)
1.4
0.61

Rb (pc)9.5
2.8

Rr (pc)
6.6
2.0

vb (km s−1 )
5100
3600

vr (km s−1 )
2700
2000

118

CHAPTER 9. DISCUSSION

outward with high velocities, namely both the SNRs are in the early stage of the transition phasebetween the free expansion and Sedov phases. Even if we assumed the exponential ejecta density
profiles (instead of uniform distributions) and calculated the dynamical parameters following theequations given by Dwarkadas & Chevalier (1998), the results did not change essentially. Note
that, the blast wave velocity of SN 1006 northwest (NW) rim (∼ 2890 km s −1 : Ghavamian et al.
2002) determined by the Hα observation is much lower than our estimation given in table 9.2. Thisis probably due to that the ISM density around NW rim is locally high. Indeed, Hα emission isthe brightest at the NW rim (see figure 5.2), and the ambient density of this region is estimated tobe 0.3–0.5 cm−3 (Dubner et al. 2002).
We again use the equations given by Truelove & McKee (1999), and expected the future
dynamical evolutions of SN 1006 and Tycho. Figure 9.6 shows the expected average radii of theblast wave (solid lines) and the reverse shock (dashed lines) as the functions of ages of each SNR.
In the case of SN 1006, since the ambient density is extremely low, the reverse shock will not arriveat the center of the remnant until ∼ 3000 years after. Therefore, the transition phase (between thefree expansion stage and the Sedov stage) will continue for a long time. On the other hand, in thecase of Tycho, the reverse shock will begin to propagate inward ∼ 400 years after. When the age
will be ∼ 1800 yr, the reverse shock will arrive at the SNR center, and hence all of the ejecta will
be heated completely.
In chapter 7, we reported the results of the Suzaku observation on the RCW 86 northeast
quadrant, and revealed for the first time the morphology of the shock-heated Fe-rich ejecta in thisregion. Then, we found that the radius of the reverse shock behind the East rim (where the ambientdensity is large) is smaller than that behind the NE rim (which expands in a tenuous region). It isvery interesting that the evolution of the reverse shock is quite different from position to positioneven in one SNR, if the ambient density has a large spatial variation. The additional observation on

Fig. 9.6.—: (a) Expected shock evolution of the blast wave (solid line) and the reverse shock (dashed line)
in SN 1006. Uniform density distributions of the ejecta and ambient medium are assumed. The dotted line
indicates the present. (b) The same to (a), but for the Tycho’s SNR.

119

9.5. FUTURE WORKS

the other region of RCW 86 will be helpful for the deeper understanding of the dynamical evolutionof SNR shock in inhomogeneous interstellar space.

9.5

Future Works

Construction of the Plasma CodeDue to the excellent energy resolution of Suzaku, we newly detected many emission lines: Kα
from low ionized Ar and Ca (in SN 1006 and E0509–67.5), Kβ from low ionized Si (in E0509–67.5)
and Fe (in Tycho), Cr-Kα (in Tycho and N103B), and possibly O-K series in the higher transitions(in SN 1006). Since these lines have quite low emissivities and hence can be ignored in mostcases, they have not been included in the plasma codes which we used in this thesis. However,in the case of young ejecta-dominated SNRs, contributions of these lines are never negligible as
we showed. Furthermore, they are very important informations to determine more accurately thethermal condition of the plasma and the abundances of synthesized elements. It is necessary tounderstand the atomic physics in low ionization states, and to construct a more detailed plasma
code.
Measurement of Ion Temperature
We found that in the young SNRs with low ambient densities, such as SN 1006 or E0509–
67.5, both the swept-up ISM and the heated ejecta are far from the full ionization equilibrium.
Therefore, it is expected that the equilibration between the ion and the electron temperatures havenot been achieved as well. Indeed, the electron to proton temperature ratios of ISM just behindthe forward shock are measured in several SNRs by optical observation (summarized by Rakowski2005). For example, in the case of SN 1006, the ratio is constrained to be T e /TH ≤ 0.07 (Ghavamian
et al. 2002). On the other hand, the ion temperature of ejecta can be measured from widths of theemission lines broadened by the thermal Doppler effect. However, we at present have no detector
which has enough energy resolution to measure the line broadening. In young SNRs, since mostof the energy must be stored in the ions in the ejecta, the measurement of the ion temperatures isnecessary for the study of the energetics of SNRs in early evolutional state.
X-ray micro-calorimeter on board NeXT satellite, which will be launched on 2012, has a very
high energy resolution of ∼ 5 eV (FWHM) in the whole energy band of 0.3–12 keV. On the other
hand, the line width due to a thermal Doppler broadening is described as∆E = 2(ln 2)1/2 E



2kTi
mi c2

1/2

[eV] ,(9.1)

where E, Ti , and mi are energy of the emission line, temperature and mass of the ion, respectively.
c is the light speed. Therefore, if the temperatures of Si (kT Si ) or Fe (kTFe ) are respectively higher
than 35 keV or 6 keV, their line widths are obtained to be ∆E > 5 eV, and hence they can be

120

CHAPTER 9. DISCUSSION

detected ideally. For example, when the reverse shock velocity is 2000 km s −1 (for the reverse shock
rest frame), the post-shock ion temperatures are expected to be kT Si ∼ 220 keV and kTFe ∼ 440 keV,
according to equation 3.36 and assuming the complete collisionless shock. Therefore, we will beable to measure the ion temperatures of ejecta in typical young SNRs using the micro-calorimeter.
Such observations are strongly encouraged.

Chapter 10

Conclusion
We performed an X-ray study of ejecta-dominated SNRs, SN 1006, E0509–67.5, RCW 86,
Tycho, and N103B, utilizing the good energy resolution and high sensitivity of Suzaku. The important results are summarized as follows:
1. Including the low-ionized Fe K-shell emission from SN 1006, many emission lines are firmly
detected for the first time, and their central energies are determined with high accuracy. Sincesome of these lines are not included in the plasma code which we used, a detailed study of
the atomic physics and a construction of the plasma code are much required.
2. The full-band spectra extracted from SN 1006 and E0509–67.5 are well-described by models
with several plasma components. The relative abundances of the heavy elements in the ejecta
components are almost consistent with predicted nucleosynthesis yield of the Type Ia SN.
3. The ionization ages of the Fe ejecta in the Type Ia SNRs are significantly lower than thoseof the lighter element ejecta. This fact suggests that Fe have been heated by reverse shock
more recently than the other elements. This explanation is consistent with a picture where
the ejecta are stratified by composition with Fe in the interior.
4. The extreme non-equilibrium states of the plasma in SN 1006 and E0509–67.5 are due to thelow densities of the ambient medium.
5. The ambient density of RCW 86 has a large spatial variation. Therefore, the evolution of thereverse shock in this SNR is quite different from position to position; the radius of the reverse
shock behind the dense region is smaller than that behind the tenuous region.

121

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Acknowledgments
I deeply appreciate the guidance and support by Professor Katsuji Koyama. I am very grateful
to my collaborators, especially Satoru Katsuda and Hiroshi Nakajima (Osaka University), John
P. Hughes (Rutgers University), Jacco Vink (University Utrecht), Aya Bamba (ISAS/JAXA), Akiko
Kawachi (Tokai University), Koji Mori (Miyazaki University), Junko S. Hiraga (RIKEN), and Ryo
Yamazaki (Hiroshima University). Many helpful comments were given by Shuichiro Inutsuka (Kyoto
University) and Masaomi Tanaka (University of Tokyo).
I would like to thank most of the members of the Cosmic-ray laboratory in Kyoto University,
especially Yoshiaki Hyodo and Hideyuki Mori, for their continuous discussions and technical supports. Kentaro Miuchi helped me to improve the manuscript. I also thank the Suzaku satellite, the
XIS, and all members of the Suzaku Science Working Group (SWG).
This work is supported by JSPS Research Fellowship for Young Scientists, and the Grant-inAid for the 21st Century COE ”Center for Diversity and Universality in Physics” from the Ministryof Education, Culture, Sports, Science and Technology (MEXT) of Japan.

131