Circumstellar interaction of supernovae and gamma-ray bursts Poonam Chandra

National Radio Astronomy Observatory & University of Virginia

Supernovae

Calcium in our bones Oxygen we breathe Iron in our cars

SUPERNOVA

Death of a massive star Violent explosions in the universe Energy emitted (EM+KE) ~ 10 51 ergs.

(To realise hugeness of the energy, the energy emitted in the atmospheric nuclear explosion is ~ 1 MT ≈ 4x10 22 ergs.)

SUPERNOVAE

Thermonuclear Supernovae

Supernovae Core Collapse

Two kinds of supernova explosions

Thermonuclear Supernovae Core collapse Supernovae •Type Ia •No remnant remaining •Less massive progenitor (4-8 M Solar ) •Found in elliptical and Spiral galaxies •Type II, Ib, Ic •Neutron star or Black hole remains •More massive progenitor (> 8 M Solar ) •Found only in Spiral arms of the galaxy (Young population of stars)

Energy scales in various explosions Chemical explosives Nuclear explosives Novae explosions Thermonuclear explosions Core collapse supernovae ~10 -6 MeV/atom ~ 1MeV/nucleon few MeV/nucleon few MeV/nucleon 100 MeV/nucleon

H (Type II) (Various types IIn, IIP, IIL, IIb etc.) Based on optical spectra Classification Si (Type Ia) No H (Type I) He (Type Ib) No Si (6150A o ) No He (Type Ic)

Crab Kepler Tycho Cas A

Not to scale Circumstellar matter SN explosion centre Photosphere Outgoing ejecta Reverse shock shell Contact discontinuity Forward shock shell

Radius

Shock Formation in SNe

: Blast wave shock : Ejecta expansion speed is much higher than sound speed.

Shocked CSM : Interaction of blast wave with CSM . CSM is accelerated, compressed, heated and shocked.

Reverse Shock Formation : Due to deceleration of shocked ejecta around contact discontinuity as shocked CSM pushes back on the ejecta.

Chevalier & Fransson, astro-ph/0110060 (2001)

Circumstellar Interaction Shock velocity of typical SNe are ~1000 times the velocity of the (red supergiant) wind. Hence, SNe observed few years after explosion can probe the history of the progenitor star thousands of years back.

Interaction of SN ejecta with CSM gives rise to radio and X-ray emission

•

Radio emission from Supernovae

: Synchrotron non-thermal emission of relativistic electrons in the presence of high magnetic field.

•

X-ray emission from Supernovae

: Both thermal and non-thermal emission from the region lying between optical and radio photospheres .

X-ray emission from supernovae Thermal X-rays versus Non-thermal X-rays

X-rays from the shocked shell

Inverse Compton scattering (non-thermal)

X-rays from the clumps in the CSM (thermal)

Swift XMM

RADIO TELESCOPES

Radio Emission in a Supernova Radio emission in a supernova arises due to synchrotron emission, which arises by the

ACCELERATION OF ELECTRONS

in presence of an

ENHANCED MAGNETIC FIELD

.

?

?

SN 1993J

Date of Explosion

:

28 March 1993 Type

:

IIb Parent Galaxy

:

M81 Distance

:

3.63 Mpc

Giant Meterwave Radio Telescope

235 MHz map of FOV of SN1993J M82 M81 1993J

Observations of SN 1993J at meter and shorter wavelengths Date of observation Dec 31, 01 Dec 30, 01 Oct 15, 01 Jan 13, 02 Jan 13, 02 Jan 13, 02 Jan 13, 02 Jan 13, 02 Frequency GHz 0.239

0.619

1.396

1.465

4.885

8.44

14.97

22.49

Flux density mJy 57.8 ± 7.6

47.8 ± 5.5

33.9 ± 3.5

31.4 ± 4.28

15.0 ± 0.77

7.88 ± 0.46

4.49 ± 0.48

2.50 ± 0.28

Rms mJy 2.5

1.9

0.3

2.9

0.19

0.24

0.34

0.13

Composite radio spectrum on day 3200  = 0.6

Frequency (GHz)

GMRT VLA

Synchrotron Aging

Due to the efficient synchrotron radiation, the electrons, in a magnetic field, with high energies are depleted.

.

dE dt Sync

  2

e

4 3

m

4

c

7

B

2 sin 2 

E

2

b

Q(E)  E g N(E)=kE g steepening of spectral index from  =( g -1)/2 to g /2 i.e. by 0.5

E cut



off

 1

bB

2

t

.

  3

e

4 

m

3

c

5

B

sin 

E

2

E

Composite radio spectrum on day 3200 R= 1.8x10

17 cm B= 38 ±17 mG  break =4 GHz c 2 = 7.3 per 5 d.o.f.

 = 0.6

Frequency (GHz)

GMRT VLA c 2 = 0.1 per 3 d.o.f.

Synchrotron Aging in SN 1993J Synchrotron losses Adiabatic expansion Diffusive Fermi acceleration

Energy losses due to adiabatic expansion

dE dt Adia

 

V R E

 

E t R

Ejecta velocity

V

Size of the SN

Energy gain due to diffusive Fermi acceleration

dE dt Fermi

 

E t c

 2

EV

20   

E

(

R

/

t

) 2 20     4 ( v 1  v 2 ) 3 v

t c

 4   v   1 v 1  1 v 2   v v 1   Upstream velocity 2 Downstream velocity Spatial diffusion coefficient of the test particles across ambient magnetic field v Particle velocity

  

dE

/

E dt



Total

 (

R

2

t

 2 / 20   )

E E



bB

2

E

2 

t

 1

E

For  

t

and

B



B

0 /

t

(Fransson & Bjornsson, 1998, ApJ, 509, 861)

Break frequency

.

.



break



B

0  3  

R

2 20  

t

 1 / 2  2

t

1 / 2   2

Magnetic field independent of equipartition assumption & taking into account adiabatic energy losses and diffusive Fermi acceleration energy gain

B=330 mG

U rel

.

U mag



B

 ( 2 g 4  13 )  8 .

5  10  6  5 .

0  10  4 (Chevalier, 1998, ApJ, 499, 810)

ISM magnetic field is few microGauss. Shock wave will compress magnetic field at most by a factor of 4, still few 10s of microGauss. Hence

magnetic field inside the forward shock is highly enhanced

, most probably due to instabilities Equipartition magnetic field is 10 times smaller than actual B, hence magnetic energy density is 4 order of magnitude higher than relativistic energy density

Gamma-ray burst

They were discovered serendipitously in the late 1960s by U.S. military

satellites

which were on the look out for Soviet nuclear testing in violation of the

atmospheric

nuclear test ban treaty. These satellites carried gamma ray detectors since a nuclear explosion produces gamma rays.

Gamma-Ray Burst

How explosive???

Even 100 times brighter than a

Brightest source of Cosmic Gamma Ray Photons

Gamma-ray bursts

Long-duration bursts:

 Last more than 2 seconds.  Range anywhere from 2 seconds to a few hundreds of seconds (several minutes) with an average duration time of about 30 seconds .

Short-duration bursts:

 Last less than 2 seconds.  Range from a few milliseconds to 2 seconds with an average duration time of about 0.3 seconds (300 milliseconds).

In universe, roughly 1 GRB is detected everyday.

GRB Missions BATSE BeppoSAX

GRB interaction with the surrounding medium Often followed by "afterglow" emission at longer wavelengths ( X-ray , UV , optical , IR , and radio ).

GRB properties

Afterglows made study possible and know about GRB GRB are extragalactic explosions.

Associated with supernovae They are collimated.

They involve formation of black hole at the center.

If collimated, occur much more frequently.

GRB 070125

 Brightest Radio GRB in Swift era.

 Detected by IPN network.

 Followed by all the telescopes in all wavebands in the world.

 Detection in Gamma, X-ray, UV, Optical, Infra-red and radio.

 Jet break around day 4.

 Still continuing radio observations.

GRB 070125

THANKS!!!!

First order Fermi acceleration V1 Vs V2

Boltzmann Equation in the presence of continuous injection



N



t

  

E

 

dE dt total N

  

qE

 g

N

(

E

,

t

) 

E

 g  ( g

E

 1 ) /( g  1 ) Form of synchrotron spectral distribution

E E

 

E E break break I

     ( g  1 ) / 2  g / 2      

break break

Kardashev, 1962, Sov. Astr. 6, 317

Self-similar solutions Equations of conservations in Lagrangian co-ordinates for the spherically symmetric adiabatic gas dynamics are 

r



t

 

M

 

v



t



v

4  3 4 

r

2

r

3 

P



M

  1  

GM r

2

To find similarity solution, we substitute velocity, density and pressure into the spherically symmetric adiabatic gas dynamics equations

v



t r U

(  ) 

A t

1 

m



U

(  )  

P



b A n t

 2

m G

(  )

b A n t

2

m r

2  (  ) 

t

2

b A n

 2

t

 2

m

 2 ( 1 

m

)  2  (  ) where  

r At m

This reduces the partial differential equations to 

G

' 

U



G G



G G

' '  ( g

G

  

m U

 ' (

U

 1 )(

U

  

U

 '  3

U m

)

m

)      ' 2

m

 '  2  (

U

  0 

U m

)  (

U

 1 ) 2 (

U

  1 ) 0  ( g  1 ) 2

m

 0  where (  )   ( 

G

(  ) )

m

 and

n

 3

n

 2

A

   g g  1  1

G

0  0 2

b a

  1

n

 2

Hugoniot conditions

G i U i

  g g g  1  1  1  g  1 2

m



i

 ( g  1 ) 2 ( 1  ( g  1 ) 2

m

) 2

U

0  0   g 2

m

 1 ( ( g g  1 )  1 ) 2