supernova remnants as cosmic rays accelerators
DESCRIPTION
Supernova Remnants as Cosmic Rays Accelerators. Vladimir S. Ptuskin Institute for Terrestrial Magnetism, Ionosphere and Radio Wave Propagation of the Russian Academy of Sciences (IZMIRAN), Troitsk, Moscow region 142190, Russia. N cr ~ 10 -10 cm -3 - total number density - PowerPoint PPT PresentationTRANSCRIPT
Supernova Remnants as Cosmic Rays Accelerators
Vladimir S. Ptuskin
Institute for Terrestrial Magnetism, Ionosphere and Radio Wave Propagation of the Russian Academy of Sciences
(IZMIRAN), Troitsk, Moscow region 142190, Russia
cosmic ray halo
Sun
bubble
GC close binary
Galactic disk
pulsar SNR
stellar wind
M87
GRB
interacting galaxies
ulsar
Ncr ~ 10-10 cm-3 - total number density
wcr ~ 1.5 eV/cm3 - energy density
Emax ~ 3×1020 eV - max. observed energy
Lcr ~ 5×1040 erg/s - Galactic luminosity in
CR
δcr ~ 10-3 at 1012 - 1014 eV - anisotropy
rg ~ 1E/(Z×3×1015 eV) pc - Larmor radius
source spectrum
Ncr
TQcr
E-2.7
cosmic ray densityescape time E-(0.3 … 0.6)
source spectrum E-(2.0 … 2.4)
two power laws: source spectrum + propagation
secondary species: Qcr,2 =
nvσ21N1 d, 3He, Li, Be, B … p, e+
escape length: X = ρvT
~ 10 g/cm2 at 1 GeV/nucleon
SNRSun
cosmic-ray halo
galactic disk
r
2H
flat-halo diffusion model
D
HX
2
v
Ginzburg & Ptuskin 1976 Berezinskii et al 1990 Strong & Moskalenko 1998
surface gas density 2.4 mg/cm2
GV5 s,/cm102 254.05
28 RRHD
km/s40s,/cm109.5 23.05
28 aVRHD
rigidity magnetic,v
kpc,5 5 Ze
cpR
cHH
pure diffusion
diffusion + distributed reaccele- ration in ISMJones et al 2001
Alfven velocity
Energy balance
local galactic CRenergy density 1.5 eV/cm3
needed source power 3×1038 erg/s kpc2
SN kinetic energy 2×1039 erg/s kpc2
(Wsn=1051 erg, 50 Myr-1 kpc-2)
~ 15% efficiency of CR acceleration
+ pulsars 2×1050 (10 ms/τ)2 erg+ stellar winds 2×1038 erg/s kpc2
+ Galactic GRBs 1051 erg/105 yr+ Galactic Center
SNR blast waves• SN II, SN Ib/c –core collapse of massive stars• SN Ia – thermonuclear explosion of white dwarf in binary
system Mechanical energy Wsn ~ 1051 erg (1053 for hypernova)
- Free expansion (ejecta-dominated stage): t < 300 yr, ush = 5×108 – 3×109 cm/s, R < 2 pc- Adiabatic deceleration (Sedov stage): t = 103 - 3×104 yr, ush ~ (Wsn/nism)1/5t-3/5
- Radiation cooling: t > 105 yr, R > 20 pc
Acceleration by external shock: a) “normal” composition after correction on atomic properties (FIP, volatility)
b) delay between nuclear synthesis and acceleration
yr 101.1
Co,Ni
5
59EC59
high obs. 59Co/56Fe – δt > 105 yrSoutoul et al. 1978, Leske 1993
Diffusive shock acceleration
pdt
dp sh
3
u
SNR
Fermi 1949, Krymsky 1977, Bell 1978
ush
D(p)
shock
-average gain of momentum
2
20
2
2
41
3
3
/)(
)()(
)(
res
g
sha
r
r
B
BvrD
upDt
pfpEI
pppf
distributionfunction(test particles)
time ofacceleration
CR intensity
resonantdiffusion kres~1/rg
Larmor radius
Maximum energy
10)(
pD
Ru shshcondition of acceleration,critical Pecklet number(parameter of modulation)
SNRWsn=1051erg
ismn0=1cm-3
scmPD
scmnWRu
GVism
shsh
/106
/10
23.028
25/2
05128
-maximum value
-typical in interstellar medium
diffusion should be anomalously slow near the shock
(upstream and downstream)
cosmic ray streaming instability in shock precursorBell 1978, Lagage & Cesarsky 1983, McKenzie & Vőlk 1982, Achterberg 1983,Vőlk et al. 1988, Fedorenko 1990, Bell & Lucek 2001, VSP & Zirakashvili 2003
Nagano & Watson 2000
Bohm limit
galacticextra-galactic?
knee
standard assumption δB ~ Bism
Bohm diffusion
5/1max
14max
221
10
/106
3
tE
eVZE
scmP
vrD
GV
gB
might be better for SNexplosion in progenitor windVőlk & Biermann 1988
x
u(x)
ush
ush/r
D(p)/u
precursorsubshock
upstream downstream
-∇Pcr
Nonlinear shock modification by CR pressure
nonmodified shock
xsh
cosmic ray density
Berezhko &Elliison 1999
not power law spectrum for high Mach number shocks
Axford 1977, 1981Eichler 1984Berezhko et al. 1996Malkov et al. 2000
mcp
apppf
uP
a
crshcrcr
,5.00,)(
5.0,
4
2
overall CR spectrum
Berezhko &Völk 2000
Cassiopeia A is bright at all energies of the electromagnetic spectrum. This composite image shows Cassiopeia A at many different wavelengths: radio polarization in red (VLA), X-rays in green (CHANDRA) and optical in blue (HST). Notice the outer shock, visible only in X-rays, as the thin green rim most visible at the top of the image. Also notice the bright ring which is visible at all three wavelengths, and the many different filamentary structures seen at each wavelength. The compact remains of the exploded star are visible only in X-rays, as the bright green spot slightly below and to the left of the geometric center of the bright ring.
observationsradio emissionνMHz = 4.6 BμGEe,GeV
2
E = 50 MeV – 30
GeV
(100 GeV for IR)
γ = 1.9 – 2.5
We = 1048 – 1049 erg
Ginzburg &
Syrovatskii 1964
Shklovsky 1976
nonthermal X-raysεkeV = 1 BμG(Ee/120 TeV)2
εmax ~ 100 TeV
SN1006 Koyama et al. 1995Cas A Allen et al. 1997RX J1713-39 Koyama et al. 1997RX J0852-46 (“Vela jr”) Slane et al 2001
γ-rays (π0)Ε = 30-3000 MeVγ Cygni, IC443Esposito et al. 1996Sturner & Dermer 1996
TeV γ – rayselectrons/protonsεmax ~ 100 TeV
SN1006 Tanimori et al 1998RX J1713 Muraishi et al. 2000Cas A Aharonian et al. 2001
Only upper limits on TeV γ-rays from many SNRs with
ages > 3×103 yr Buckley et al. 1998, Aharonian et al. 2002
e
γsynchrotron
e
γ inverse Comptonεγ = ε0(Ee/mec2)2p
π0
γ
SNR
SN1006
Tanimori et al. 2001
Problems:
- Galactic sources should work up to (1-3)×1018 eV (Fe ?) (reacceleration may help: Axford 1994, Bell 1992,
Bykov & Toptygin 2001, Vőlk & Zirakashvili 2004;dispersion of SN parameters: Sveshnikova 2003)
- no VHE gamma-rays from not very young SNRs tsnr ≥ 3×103 yr
(Buckley et al. 1998, Aharonian et al. 2002)
- cosmic ray source spectrum γs = 2.0 - 2.4 (depends on propagation model)
VSP & Zirakashvili 2003
Wsn = 1051 erg, Bism = 5 μG, n0 = 0.4 cm-3 ξcr = 0.5, κ = 0.04, a = 0.3
strong streaming instability and non-linear wave interactions in shock precursor:
under extreme conditions:
Emax ≈ 1017Z(ush/3×104km/s)2
×(κ/0.1)(ξcr/0.5)Mej1/3n1/6 eV
δBmax≈ 10-3 (ush/3×104km/s)n1/2 G
maximum momentum of accelerated particles:abandonment of Bohm limit hypotheses
Random field produced by cosmic-ray streaming instability
in shock precursor
22
,4
, shcrcrcr
ash uPB
wx
PV
x
wu
0,0
00,0 ,
4,
a
shcr
a
V
u
B
B
BVBB
0,0
,0 ,4
,
a
shcr
efa
V
u
B
B
BVBB
Alfven velocity
cosmic-ray pressure
wave energy density
weak random field: strong random field:
characteristic velocity of waves
Bell & Lucek 2001VSP & Zirakashvili 2003
Average source spectrum
spectrum atthe shock instantaneousSNR luminosityin run-awaycosmic rays
averagecosmic-raysourcespectrum
adiabatic stage Q ~ ξcrνsnWsnp-4 (Sedov) - universal spectrum !
ejecta-dominated stageSNII in RSG wind: Q ~ p-6.5 at ρstar~ r -10
SNI in uniform medium: Q ~ p-7
(Chevalier – Nadyozhin)
))((~
))((~
maxmax4
max32
max4
max2
ptpdt
dppRuq
ptpHppuf
shshcr
aashcr
SN rate
hot bubble0.013 cm-3, 3μG
ism R=60pc
n=1cm-3
RSG wind
Weaver et al. 1977Chevalier & Liang 1989
KASCADE
SNII
Roth et al. 2003
·Eknee ≈ 7×1015 Z eV, ~ ξcrWsnM1/2(Mejuw)-1
Emax ≈ 4×1016 Z eV at tmin = 7 days
ρstar~ r-10
∙M=10-5
uw=10km/sRw=2pc
Wsn= 1051 erg, ξcr= 0.5
VSP & Zirakashvili 2004
Other proposals on acceleration beyond the knee:
• Reacceleration by multiple shocks
• Reacceleration in plerions
SNRSNR
SNR
pulsar wind
SNR
ΩδΦ
δΦ = 4×1015Z eV – 1019Z eV
Bell 1991, 2000, Berezhko 1993
uE θ= Bφur/c
OBassociationu=3×103 km/sB=10-5 GR=30 pc
f ~ 1/p3
ta ~ R/(Fshu) at Di < uR ~ D/(Fshu2) at Di > uR
R u
Emax ~ 1017Z eV
Axford & Ip 1991, Bykov & Toptygin 1990, 2001Klepach et al. 2000
terminationshock
Crab pulsar few msec pulsar
•Galactic wind
u
R acceleration at terminationshock Jokipii & Morfill 1985, 1991
R = 300 kpc, u = 400 km/s
Emax = 3×1018Z eV
galactic disk
SNRacceleration by travelingshocks and interactionregions Völk & Zirakashvili 2004
Nagano & Watson 2000 galacticextra-galactic?
knee