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 youngThe surrounding interstellar matter can be negligible

-> All explosion energy is used tothe 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: 4/3p(vshockt)3

the swept-up mass: 4/3p(vshockt)3rISM

the expansion starts to stop when;

(

yearscmn

MM

ergE

scmn

MM

ergE

Mvt

tvM

ISM

sun

ISM

sun

ISMshock

ISMshockstar

3/1

3

6/5

*

2/1

51exp

3/1

3

6/5

*

2/1

51exp10

3 1*

3

3

110102000

11010103.6

4334

rp

rp (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

Kyeart

cmn

ergE

KRcmn

ergE

T

scmyeart

cmn

ergE

dtdRv

pcyeart

cmn

ergE

R

ISM

ISM

ISMshock

ISM

5/6

3

5/2

3

5/2

51exp8

31

351exp10

5/3

3

5/1

3

5/1

51exp8

5/2

3

5/1

3

5/1

51exp

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

SiS

ArCa

Fe

(cal src)

MgNeO

Identification of major heavy elements !

Searching for minor elementsThere 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)

SAr

Line broadening due to expansionHeavier elements stay inside of the remnant.

Lines have doppler broadening by expansionblue-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-1

g+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

cvZ

GB

pcL

eVE shock

1121

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 shocksSync. X-rays forms thin filaments

on 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 electrons

magnetic 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 = vs22(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 = vs22(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