Supernova Remnantsand
their Emission
Aya Bamba(Aoyama Gakuin U.)
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
Member of astro group in Aoyama Gakuin Univ.
Once a star explodes …
We can study the remnants even we do not know how stars explode.
Parameters:Energy input: ~1051 ergDuration: delta functionMass: ~Msun
1. Evolution of SNRs and emission2. Acceleration of particles on shocks of SNRs
and emission
remnants of dead stars: Supernova remnants (SNRs)
Anyhow, stars explode. (An-pan-man knows it ..)
Goal: SNRs makes chemical and high-Eof our diverseness
time flies quickly,shining stars will disappear,so you should go ahead
with smile !
An-pan-man March(one of the most famous
super-star)
Once a star explodes …
We can study the remnants even we do not know how stars explode.
Parameters:Energy input: ~1051 ergDuration: delta functionMass: ~Msun
1. Evolution of SNRs and emission2. Acceleration of particles on shocks of SNRs
and emission
remnants of dead stars: Supernova remnants (SNRs)
Anyhow, stars explode. (An-pan-man knows it ..)
Goal: SNRs makes chemical and high-Eof our diverseness
0. Why X-ray observation is important for SNR study ?
X-ray detector
X-rays
electrons
ne ∝ EX
X-ray detectors basically change the X-rayto number of electrons.
We can measure- position- time- energy
simultaneously !
1. Evolution of supernova remnantsand 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)
2*exp 2
1shockvME
when Eexp=1051erg and M*=10Msun,vshock = 3.2x108 cm/s
(10% of the light speed)
uniform expansion free from deaccelerationradius ∝ vshockt “free expansion phase”
only kinetic E -> no energy dissipationno 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: vshocktthe swept-up volume: 4/3 (p vshockt)3
the swept-up mass: 4/3 (p vshockt)3rISM
the expansion starts to stop when;
yearscm
n
M
M
erg
E
scm
n
M
M
erg
E
Mvt
tvM
ISM
sun
ISM
sun
ISMshock
ISMshockstar
3/1
3
6/5
*
2/1
51
exp
3/1
3
6/5
*
2/1
51
exp10
31
*3
3
110102000
11010103.6
4
3
3
4
(assumption: uniform density)
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
Kyear
t
cm
n
erg
E
KRcm
n
erg
ET
scmyear
t
cm
n
erg
E
dt
dRv
pcyear
t
cm
n
erg
ER
ISM
ISM
ISMshock
ISM
5/6
3
5/2
3
5/2
51
exp8
31
351
exp10
5/3
3
5/1
3
5/1
51
exp8
5/2
3
5/1
3
5/1
51
exp
10110102.1
110105.1
]/[10110
101.2
101100.5
R∝t0.4
v∝t-0.6
T∝t-1.2
loss of kinetic E -> thermal E of downstream plasmaE loss by emission is still negligible“Sedov phase” or “adiabatic phase”
Kepler’s nova
Cassiopeia A
X-ray gallery of young supernova remnants
beautiful fireworks in the universe~1 event / 30 yrs
Tycho’s nova
SN1006
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 directionby 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 ionizedinto around He-like or H-like ions.
He-like H-like
X-ray emission lines-> we can knowhow much heavy elements are distributed into interstellar medium
Tycho’s spectra by Suzaku
Si
S
ArCa
Fe
(cal src)
MgNeO
Identification of major heavy elements !
Searching for minor elements
There are many kinds of elements !All should be made in stars and distributed by supernovaeImportant 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 linesfrom Tycho
(Tamagawa+08)
Suzaku 100ks observation -> detection of Cr and Mn lines !MMn/MCr = 0.5 (0.2-0.7)
First detection of emission lines from minor elementsNear-future observatories w. excellent E resolution willdetect minor elements from many SNRs.
How elements dissipate into the interstellar medium?
onion-like? mixing ?
The line distribution of Cas A by Chandra (Hwang+04)
Si Fe
It should have information of its explosion
heavier elements are located inside of
lighter elements ??We need 3D information
Fe
Si
Radial peak in arcmin(ASCA: Hwang & Gotthelf 97)
Velocity in km/s (Hayoto+11)
S
Ar
Line broadening due to expansionHeavier elements stay inside of the remnant.
Lines have doppler broadening by expansion
blue-shiftred-shift
Tycho spectrum (Suzaku)line shift -> expansion velocity
The plasma age
Ionization is mainly by collision of ions and electronsin SNR plasma.
Plasma in SNRs are so tenuousand ionization takes long time.
In order to reach the equilibriumbetween temperature and ionization,
nt ~ 1012 s cm-3
if n ~ 1 cm-3, t ~ 3x104 yr
Plasma before equilibrium: non equilibrium stateor 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 isacceleratedwith strong emission
radiative cooling phase
ISM
Emission of plasma -> taking E out from the shock-> shock speed slows down more
hot plasma
coolanddens
e shell
cool down -> pressure can be ignoredshell collects ISM further like snowplow“snowplow phase”
R∝t2/7 -> R∝t1/4
mixed morphology SNRs
shells are already cold to emit X-rays
IC443 (Keohane+)
ejecta is still hot
radioX-ray
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 particleson shocks of SNRs
2.1. cosmic rays
(Cronin 1999)
knee=1015.5eV
ankle=1018.5eV
very high E particlesin the universe
uCR ~ 1eV/cc
1 CR per your fingertipper 1 second.
one of the main componentsof our Galaxy
c.f. CMB 0.3 eV/cc star light < 0.3 eV/cc magnetic field 0.3 eV/cc turbulence 0.3 eV/cc thermal E0.01 eV/cc
2.2. Shocks of SNRs are cosmic ray accelerator !
chemical abundance of cosmic rays
made through separation ofheavier elements
Be: basically made through separationincluding 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:
shock system:shock front
(us = 0)
upstream region(uu = us)
downstream region
(ud = us)g-1g+1
ideal gas: g = 5/3 → ud = 1/4us
shock front(us)
outer region(u1 = 0)
inner region(u2)
SN
paticles change their direction with scatteringby magnetic field turbulence
E conversation in upstream/downstream in shock system
particles get energy always crossing shockspectrum of particle: power-law
ux
uy
ud uu
lab. systemshock front
x
2.3. 3 min. recipe of particle accleration (2) acceleration
<- same to cosmic rays !
2.4. Maximum E of particles
c
vZ
G
B
pc
L
eV
E shock
112
1
1015
Particles gyrate with magnetic field (gyro motion)The radius should be smaller than the size of the system
larger systemstronger magnetic field
can accelerate particles to higher energy
2.5. Nonthermal emission from accelerated particles (1)synchrotron emission
-
electrons gyrateby 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)
SNRs are really hot bubbles.Shock fronts accelerate (at least) electrons.
thermal X-rays
(Yamaguchi+08)
Sync. X-rays
(Bamba+08)
When electron spectrum is power-law,
The spectrum of the synchrotron emission is,
We can know the index of electronsfrom the spectrum of synchrotron emission.
CR has power-law spectrum,we can expect the index of synchrotron emission.
∝ B2 (ve ~ c)
Spectrum of synchrotron emission
2.6. Nonthermal emission from accelerated particles (2)inverse Compton emission
-
hn
-
hn’
collision between particles and photon-> photon get E from particles“inverse Compton emission”
Uph: E density of scattered photon
(SNR case: photons are CMB)
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
sync. IC
pi-0
lines
bremss.
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
sync. IC
pi-0
lines
bremss.
IC (e origin) ?pi-0 (p origin) ?
No confirmation yet…
W44
GeV emission !(Abdo+10)
sync. IC
pi-0
lines
bremss.
pi-0 emission …?association with molecular cloudOnly 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 filamentson shock fronts
-> gyro-radius of electrons is so small
0.3 – 2.0 keV2.0 – 10.0 keV
SN1006 NE shell
(Bamba +03)
~0.01 – 0.1 pc !
Why diffusion is so small ?
-frequent scattering of electronsmagnetic field turbulence scatter electronsturbulent B -> frequent scatter -> small diffusion
-small gyro radius (= large magnetic field)a lot of accelerated electrons-> ~ large current ~ induced magnetic field-> more efficient acceleration -> …
e-
shock
Very efficient acceleration !(Bamba+05)
From Rankine-Hugoniot relation,
kinetic E
thermal E of downstream plasma
kinetic E
kTd = mvs22(g-1)
(g+1)2ideal gas ~ 0.19vs
2
Energy budget on shock
thermal E ofdownstream plasma
E of acceleratedparticle
kinetic E
with E loss < 0.19vs2
If the acceleration is very efficient,the downstream plasma becomes cooler
thanthe case without acceleration
3. Summary
Stars and their explosions makethe chemical and high energy diverseness
of our universe.
X-ray observations are one of the best toolsto investigate such kind of diverseness.
SNRs are Hanasaka Jiisan in our universe(flower-blossom-old man)
Japanese old story“Hanasaka”
When he sprays ashonto a deadwood,the wood becomes acherry tree in full bloom
From Rankine-Hugoniot relation,
kinetic E
thermal E of downstream plasma
kinetic E
kTd = mvs22(g-1)
(g+1)2ideal gas ~ 0.19vs
2
Energy budget on shock
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
higher temperature-> higher ionization
we can measure the ion kTby measuring ionization state
9 electrons are stripped
Hillas diagram (Hillas84)
SNRs can accelerate particles up to ~knee energy