Download Circumstellar interaction of supernovae and gamma-ray bursts

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) ~ 1051 ergs.
(To realise hugeness of the energy, the
energy emitted in the atmospheric nuclear
explosion is ~ 1 MT ≈ 4x1022 ergs.)
SUPERNOVAE
Thermonuclear Supernovae
Core Collapse
Supernovae
Two kinds of supernova explosions
Thermonuclear
Supernovae
Core collapse
Supernovae
•Type Ia
•Type II, Ib, Ic
•No remnant
remaining
•Neutron star or Black
hole remains
•Less massive
progenitor (4-8 MSolar)
•More massive
progenitor (> 8 MSolar)
•Found in elliptical and
Spiral galaxies
•Found only in Spiral
arms of the galaxy
(Young population of
stars)
Energy scales in various explosions
Chemical explosives
~10-6 MeV/atom
Nuclear explosives
~ 1MeV/nucleon
Novae explosions
few MeV/nucleon
Thermonuclear explosions
few MeV/nucleon
Core collapse supernovae
100 MeV/nucleon
Based on
optical spectra
Classification
H (Type II)
(Various typesIIn, IIP, IIL, IIb
etc.)
No H (Type I)
Si (Type Ia)
No Si (6150Ao)
He (Type Ib)
No He (Type Ic)
Crab
Tycho
Kepler
Cas A
Circumstellar
matter
Density
Not to scale
Radius
SN explosion centre
Photosphere
Outgoing ejecta
Reverse shock shell
Contact discontinuity
Forward shock shell
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)
SPACE TELESCOPES
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
VLA
GMRT
Observations of SN 1993J at
meter and shorter wavelengths
Date of
observation
Frequency
GHz
Flux density
mJy
Rms
mJy
Dec 31, 01
0.239
57.8 ± 7.6
2.5
Dec 30, 01
0.619
47.8 ± 5.5
1.9
Oct 15, 01
1.396
33.9 ± 3.5
0.3
Jan 13, 02
1.465
31.4 ± 4.28
2.9
Jan 13, 02
4.885
15.0 ± 0.77
0.19
Jan 13, 02
8.44
7.88 ± 0.46
0.24
Jan 13, 02
14.97
4.49 ± 0.48
0.34
Jan 13, 02
22.49
2.50 ± 0.28
0.13
Flux density (mJy)
Composite radio spectrum on day 3200
 = 0.6
GMRT
VLA
Frequency (GHz)
Synchrotron Aging
Due to the efficient synchrotron
radiation, the electrons, in a
magnetic field, with high
energies are depleted.
.
4
dE
2
e


2
2
2

B sin E


4 7
3m c
 dt  Sync
b
Q(E)E-g
N(E)=kE-g
N(E)
steepening of spectral index from =(g-1)/2 to g/2 i.e. by 0.5
Ecut off
3e
2
 
B
sin

E
3 5
4

m
c
.
E
1
 2
bB t
Composite radio spectrum on day 3200
Flux density (mJy)
break =4 GHz
c2= 7.3
per 5 d.o.f.
17
R= 1.8x10 cm
B= 38±17 mG
 = 0.6
GMRT
VLA
Frequency (GHz)
c2= 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
V
E
 dE 
 E


R
t
 dt  Adia
R
V
Ejecta velocity
Size of the SN
Energy gain due to diffusive Fermi acceleration
E EV
E(R / t)
 dE 





tc
20 
20 
 dt  Fermi
2
4( v1  v 2 )

3v
4 
tc 
v
 1
1 



 v1 v 2 
2
v1 Upstream velocity
v 2Downstream velocity
  Spatial diffusion
coefficient of the test
particles across
ambient magnetic
field
vParticle velocity
E
E

 2 2
2 2
1
dE / dt Total ( R t / 20  ) E  bB E  t E
For
 t
and
B  B0 / t
(Fransson & Bjornsson,
1998, ApJ, 509, 861)
Break frequency
.
 R
1 / 2
1/ 2 
 break  B 
t
 2t 
 20 

3
0
.
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
B
U
. mag
( 2g 13)

4
6
 8.5 10  5.0 10
(Chevalier, 1998, ApJ, 499, 810)
4
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 times
a
Million trillion
supernova
as
bright as source
sun
Brightest
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
Swift was launched in 2004
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
  dE 
g

 N   qE

t E  dt total 
 E g
N ( E , t )   (g 1)
/(g  1)
E
E  Ebreak
E  Ebreak
Form of synchrotron spectral distribution
  (g 1) / 2
I   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
v
t
  4 3 
1
r 

M  3


v
GM
2 P
 4r

t
M
r2
To find similarity solution, we substitute velocity, density and pressure
into the spherically symmetric adiabatic gas dynamics equations
r
A
v  U ( )  1 m U ( )
t
t
b 2m
  n t G ( )
A
b 2m r 2
b  2 m  2 (1 m ) 2
P nt
( )  n  2 t
 ( )
2
A
t
A
where
r
 m
At
This reduces the partial differential equations to
G '
U  m   U '3U  2m  0
G
G '
  U ' (U  m)   '2  U (U  1)  0
G
G '
 '
(g  1)(U  m) 
(U  m)  2(U  1)  (g  1)2m  0
G

where
( )
 ( ) 
G ( )
and
n3
m
n2
g 1
A
G0 02
g 1
b
a 
1
n2
Hugoniot conditions
g 1
Gi 
g 1
g  1  2m
Ui 
g 1
(g  1)2(1  m) 2
i 
(g  1) 2
2m
U0 
g 1
(g  1)
0 
(g  1) 2