スライド 1

Supernova Remnants
and
their Emission
Aya Bamba
(Aoyama Gakuin U.)
Member of astro group in Aoyama Gakuin Univ.
井上剛志：星間物理学（理論）、プラズマ物理（理論）
星形成、超新星残骸、ＧＲＢ、磁気流体シミュレーション
大平豊：高エネルギー宇宙物理学（理論）、プラズマ物理（理論）
無衝突プラズマ現象、宇宙線粒子加速、SNR、CTA、CALET
坂本貴紀：宇宙物理学、雷雲ガンマ線（実験・観測）
Swift、小型ロボット望遠鏡、ＧＲＢ、雷・スプライト/TGF
澤田真理：精密プラズマ分光（実験・観測）
Suzaku/Astro-H、超新星残骸、銀河系中心
柴田徹：宇宙線物理学（実験）
宇宙線の伝播、中間子多重発生、ＣＴＡ
田中周太：高エネルギー宇宙物理学（理論）
相対論的電子・陽電子プラズマ、パルサー、パルサー星雲、
馬場彩：高エネルギー宇宙物理学（実験・観測）
Suzaku/ASTRO-H、ＣＴＡ，超新星残骸、パルサー星雲
山崎了：高エネルギー宇宙物理学（理論）
ＧＲＢ，超新星残骸、宇宙線粒子加速、ＣＴＡ、Ｆｅｒｍｉ
吉田篤正：宇宙物理学（実験）
ＭＡＸＩ、ＣＡＬＥＴ、ＣＴＡ，ＧＲＢ、中性子星、Suzaku/Astro-H
Once a star explodes …
Anyhow, stars explode. (An-pan-man knows it ..)
We can study the remnants
even we do not know how stars explode.
Parameters:
Energy input:
Duration:
Mass:
~1051 erg
delta function
~Msun
remnants of dead stars: Supernova remnants (SNRs)
1. Evolution of SNRs and emission
2. Acceleration of particles on shocks of SNRs
and emission
Goal: SNRs makes chemical and high-E
of our diverseness
An-pan-man March
(one of the most famous
super-star)
time flies quickly,
shining stars will disappear,
so you should go ahead
with smile !
Once a star explodes …
Anyhow, stars explode. (An-pan-man knows it ..)
We can study the remnants
even we do not know how stars explode.
Parameters:
Energy input:
Duration:
Mass:
~1051 erg
delta function
~Msun
remnants of dead stars: Supernova remnants (SNRs)
1. Evolution of SNRs and emission
2. Acceleration of particles on shocks of SNRs
and emission
Goal: SNRs makes chemical and high-E
of our diverseness
0. Why X-ray observation is important for SNR study ?
X-ray detector
electrons
X-rays
ne ∝ EX
X-ray detectors basically change the X-ray
to number of electrons.
We can measure
- position
- time
- energy
simultaneously !
1. Evolution of supernova remnants
and emission
1.1. When SNRs are young
The surrounding interstellar matter can be negligible
-> All explosion energy is used to
the kinetic energy of the exploded star (ejecta)
Eexp
1
2
 M *vshock
2
when Eexp=1051erg and M*=10Msun,
vshock = 3.2x108 cm/s
(10% of the light speed)
uniform expansion free from deacceleration
radius ∝ vshockt
“free expansion phase”
only kinetic E ->
no energy dissipation
no emission
Expansion is “visible”
Tycho’s SNR (SN1572)
(Katsuda+10)
comparison of X-ray images taken on different period
-> detection of movement of shock
How long can we use free expansion approximation ?
Stars explode in interstellar matter.
ISM
The shock sweep up the ambient ISM.
When mass of swept-up mass exceed
the mass of the exploded star,
we cannot ignore the swept-up ISM.
the radius of SNR:
the swept-up volume:
the swept-up mass:
vshockt
4/3p(vshockt)3
4/3p(vshockt)3rISM
the expansion starts to stop when;
4
3
M star  p (vshock t  r ISM
3
3
3
1
3
t
vshock M * r ISM
4p
Eexp 
10 

 6.3  10  51
 10 erg 
 Eexp 

 2000 51
 10 erg 
1 / 2
(assumption: uniform density)
1 / 2
 M*

 10M sun
 M*

 10M sun



5/ 6



5/ 6
 nISM 

3 
 1cm 
 nISM 

3 
 1cm 
1 / 3
s
1 / 3
years
Observed expansion rates
(Moffett+93)
R∝tm
Shocks of famous SNRs already starts decelerate
1.2. When the shock starts decelerate
Self-similar solution by Sedov
1/ 5
 Eexp   nISM 
 
R  5.0 51
3 
 10 erg   1cm 
1 / 5


t
 3

 10 year 
1/ 5
vshock
Eexp   nISM 
dR
8
 

 2.110  51
3 
dt
 10 erg   1cm 
1
E
 nISM  3
exp
10 

T  1.5 10  51
R
3 
 10 erg  1cm 
Eexp 
8

 1.2 10  51
 10 erg 
2/5
 nISM 

3 
 1cm 
2 / 5
2/5
1 / 5
pc
R∝t0.4


t
 3

 10 year 
3 / 5
[cm / s]
v∝t-0.6
K


t
 3

 10 year 
6 / 5
K
T∝t-1.2
loss of kinetic E -> thermal E of downstream plasma
E loss by emission is still negligible
“Sedov phase” or “adiabatic phase”
X-ray gallery of young supernova remnants
Kepler’s nova
Tycho’s nova
Cassiopeia A
SN1006
beautiful fireworks in the universe
~1 event / 30 yrs
Thermal emission from the heated plasma (1)
bremsstrahlung
Since the downstream is so hot (~106-8 K or 0.1-10 keV),
the gas is almost ionized.
+
+
bending the direction
by coulomb interaction
= acceleration
+
+
+
radiation !
bremsstrahlung
Thermal emission from the heated plasma (2)
line emission
Electrons in atoms orbit around nuclei.
When they change their orbit,
they emit/absorb photon
with transition energy
-> emission lines
In the plasma, atoms are highly ionized
into around He-like or H-like ions.
He-like
H-like
X-ray emission lines
-> we can know
how much heavy elements are
distributed into interstellar medium
O
Ne Mg
Si
S
Ar
Ca Fe
(cal src)
Identification of
major heavy elements !
Tycho’s spectra by Suzaku
Searching for minor elements
There are many kinds of elements !
All should be made in stars and distributed by supernovae
Important to measure the amount of elements near iron
(produced in imcomplete Si burning)
in order to understand how heavy metals are produced
chromium
manganese
Suzaku detection of Cr and Mn emission lines
from Tycho
Suzaku 100ks observation -> detection of Cr and Mn lines !
MMn/MCr = 0.5 (0.2-0.7)
(Tamagawa+08)
First detection of emission lines from minor elements
Near-future observatories w. excellent E resolution will
detect minor elements from many SNRs.
How elements dissipate into the interstellar medium?
mixing ?
onion-like?
It should have information of
its explosion
The line distribution of Cas A by Chandra (Hwang+04)
Si
Fe
heavier elements are
located inside of lighter
elements ??
We need 3D information
Lines have doppler broadening by expansion
red-shift
blue-shift
Velocity in km/s
(Hayoto+11)
Si
S
Fe
Tycho spectrum (Suzaku)
line shift -> expansion velocity
Ar
Radial peak in arcmin
(ASCA: Hwang & Gotthelf 97)
Line broadening due to expansion
Heavier elements stay inside of the remnant.
The plasma age
Ionization is mainly by collision of ions and electrons
in SNR plasma.
Plasma in SNRs are so tenuous
and ionization takes long time.
In order to reach the equilibrium
between temperature and ionization,
nt ~ 1012 s cm-3
if n ~ 1 cm-3, t ~ 3x104 yr
Plasma before equilibrium: non equilibrium state
or ionizing
(check Yamaguchi-san’s talk)
How emission changes with different nt
1.3. When the plasma cooled down below 2MK …
Radiative cooling coefficient
(Gehrels+93)
Plasma emit more and more
-> cool down easily
-> more efficient emission
-> …
Cooling of plasma is
accelerated
with strong emission
radiative cooling phase
Emission of plasma -> taking E out from the shock
-> shock speed slows down more
hot plasma
cool
and
dense
shell
ISM
cool down -> pressure can be ignored
shell collects ISM further like snowplow
“snowplow phase”
R∝t2/7 -> R∝t1/4
mixed morphology SNRs
shells are already cold to emit X-rays
radio
X-ray
ejecta is still hot
IC443 (Keohane+)
1.4. Disappearance of SNR
The shock speed slows down more and more
-> comparable to the proper motion of surrounding ISM
(10-20 km/s)
SNRs lose the boundary between the outside
-> disappear of SNR
time scale ~ 106 yrs
2. Acceleration of particles
on shocks of SNRs
2.1. cosmic rays
very high E particles
in the universe
uCR ~ 1eV/cc
c.f. CMB
star light
magnetic field
turbulence
thermal E
knee=1015.5eV
0.3 eV/cc
< 0.3 eV/cc
0.3 eV/cc
0.3 eV/cc
0.01 eV/cc
1 CR per your fingertip
per 1 second.
one of the main components
of our Galaxy
ankle=1018.5eV
(Cronin 1999)
2.2. Shocks of SNRs are cosmic ray accelerator !
chemical abundance of cosmic rays
made through separation of
heavier elements
Be: basically made through separation
including radio isotope 14Be
typical age of CR: ~6x106 years (Garcia-Munoz+77)
~escape timescale from Galaxy
Galaxy volume:
5x1066 cm3
CR energy density: 1.6x10-12 erg cm-3
-> we need energy input to CR of 1x1040 erg s-1
E input by SNRs: 1051 erg per 30 years = 1042 erg s-1
If 1% of SN energy is injected CRs,
we can explain all of the E of CRs by SNRs.
2.3. 3 min. recipe of particle accleration (1) terminology
lab. system:
outer region
(u1 = 0)
shock front
(us)
inner region
(u2)
SN
shock system:
shock front
(us = 0)
upstream region
(uu = us)
downstream region
(ud = g-1 us)
g+1
ideal gas: g = 5/3 → ud = 1/4us
2.3. 3 min. recipe of particle accleration (2) acceleration
paticles change their direction with scattering
by magnetic field turbulence
E conversation in upstream/downstream in shock system
lab. system
uy
shock front
x
ud
uu
particles get energy always crossing shock
spectrum of particle: power-law
<- same to cosmic rays !
ux
2.4. Maximum E of particles
Particles gyrate with magnetic field (gyro motion)
The radius should be smaller than the size of the system
E
1  L  B   vshock 

Z 
 

15
10 eV 2  1 pc  1G   c 
larger system
stronger magnetic field
can accelerate particles
to higher energy
2.5. Nonthermal emission from accelerated particles (1)
synchrotron emission
electrons gyrate
by magnetic field
= acceleration !
-
-> emission
In the case of e is relativistic:
synchrotron emission
typical emission frequency:
in 1uG magnetic field,
~GeV e -> 1012 Hz (radio band)
~TeV e -> 1018 Hz (X-ray band)
(Yamaguchi+08)
Sync. X-rays
(Bamba+08)
thermal X-rays
SNRs are really hot bubbles.
Shock fronts accelerate (at least) electrons.
Spectrum of synchrotron emission
∝ B2 (ve ~ c)
When electron spectrum is power-law,
The spectrum of the synchrotron emission is,
We can know the index of electrons
from the spectrum of synchrotron emission.
CR has power-law spectrum,
we can expect the index of synchrotron emission.
2.6. Nonthermal emission from accelerated particles (2)
inverse Compton emission
hn
hn’
-
(SNR case: photons are CMB)
collision between particles and photon
-> photon get E from particles
“inverse Compton emission”
Uph: E density of scattered photon
We can measure magnetic field
2.7. Nonthermal emission from accelerated particles (3)
emission via pi-on decay
When protons with E>1GeV collide with other protons,
sometimes produce pi-0 meson.
pi-0 meson decays into two photons.
We need a lot of protons !
molecular cloud etc.
+
Only footprint
by CR protons !!
+
Summary of emission
lines
bremss.
sync.
pi-0
IC
If we can detect each component,
we can know thermal and nonthermal matter in SNRs.
SNRs are also detected in TeV gamma-ray band
(Acero+10)
shells of SN1006
is also TeV emitter
lines
bremss.
sync.
pi-0
IC
IC (e origin) ?
pi-0 (p origin) ?
No confirmation yet…
GeV emission !
(Abdo+10)
lines
bremss.
sync.
pi-0
IC
W44
pi-0 emission …?
association with molecular cloud
Only several SNRs are detected in GeV gamma-rays.
It is still unknown what makes difference.
2.8. Topics: thin filaments on shocks
Sync. X-rays forms thin filaments
on shock fronts
-> gyro-radius of electrons is so small
SN1006 NE shell
~0.01 – 0.1 pc !
0.3 – 2.0 keV
2.0 – 10.0 keV
(Bamba +03)
Why diffusion is so small ?
-frequent scattering of electrons
magnetic field turbulence scatter electrons
turbulent B -> frequent scatter -> small diffusion
-small gyro radius (= large magnetic field)
a lot of accelerated electrons
-> ~ large current ~ induced magnetic field
-> more efficient acceleration -> …
Very efficient acceleration !
(Bamba+05)
e-
shock
Energy budget on shock
From Rankine-Hugoniot relation,
kTd =
2(g-1)
2
v
s
(g+1)2
ideal gas
with E loss
~ 0.19vs2
< 0.19vs2
thermal EE of
of
thermal
downstream
downstream plasma
plasma
kinetic E
E of accelerated
particle
kinetic E
If the acceleration is very efficient,
the downstream plasma becomes
coolerE than
kinetic
the case without acceleration
3. Summary
Stars and their explosions make
the chemical and high energy diverseness
of our universe.
X-ray observations are one of the best tools
to investigate such kind of diverseness.
SNRs are Hanasaka Jiisan in our universe
(flower-blossom-old man)
Japanese old story
“Hanasaka”
When he sprays ash
onto a deadwood,
the wood becomes a
cherry tree in full bloom
Energy budget on shock
From Rankine-Hugoniot relation,
kTd =
2(g-1)
2
v
s
(g+1)2
ideal gas
~ 0.19vs2
thermal E of
downstream plasma
kinetic E
kinetic E
Power of bremsstrahlung
∝ne2
gff: gaunt factor (Bressaard+62)
When the plasma is thermallized
probability of
ve = ve ~ ve+dve
spectrum has the cut-off at hn=kT
-> we can know the temperature of plasma
temperature dependence of ionization fraction of iron
9 electrons are stripped
higher temperature
-> higher ionization
we can measure the ion kT
by measuring ionization state
Hillas diagram (Hillas84)
SNRs can accelerate particles up to ~knee energy

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