Cosmic Ray Acceleration in Supernova Remanants Vladimir Ptuskin IZMIRAN, Russia

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Cosmic Ray Acceleration in Supernova Remanants Vladimir Ptuskin IZMIRAN, Russia. N cr ~ 10 -10 cm -3 - total number density w cr ~ 1.5 eV/cm 3 - energy density E max ~ 3x10 20 eV - max. observed energy δ cr ~ 10 -3 at 10 12 - 10 14 eV - anisotropy - PowerPoint PPT Presentation

Transcript of Cosmic Ray Acceleration in Supernova Remanants Vladimir Ptuskin IZMIRAN, Russia

  • Cosmic Ray Acceleration in Supernova Remanants

    Vladimir Ptuskin

    IZMIRAN, Russia

  • ulsarNcr ~ 10-10 cm-3 - total number densitywcr ~ 1.5 eV/cm3 - energy densityEmax ~ 3x1020 eV - max. observed energycr ~ 10-3 at 1012 - 1014 eV - anisotropy rg ~ 1E/(Z31015 eV) pc - Larmor radius

    cosmic ray halo

    interacting

    galaxies

    M87

    GRB

    stellar

    wind

    SNR

    pulsar

    Galactic

    disk

    GC

    close

    binary

    bubble

    Sun

  • energy balance Ginzburg & Syrovatskii 1964required source power 31038 erg/(s kpc2)SN kinetic energy 21039 erg/(s kpc2)(Wsn=1051 erg, Gal = 0.03 yr-1local SN rate 50 Myr-1kpc-2)~ 15% - efficiency of CR acceleration in SNRsother Galactic accelerators: pulsars [21050 (10 ms/)2 erg], stellar winds [21038 erg/s kpc2], Galactic GRBs [1051 erg/105 yr], micro quasars, Galactic Center acceleration by external shock: a) normal composition after correction on atomic properties (FIP, volatility) b) delay between nuclear synthesis and acceleration (Soutoul test: 59Ni 59Co, high obs. 59Co/56Fe gives t > 105yr Leske 1993)

  • diffusive shock accelerationSNRFermi 1949, Krymsky 1977, Bell 1978ushD(p)shockaverage gain of momentumdistributionfunction(test particles)time ofaccelerationCR intensityresonantdiffusion kres~1/rgLarmor radius

  • maximum energycondition of acceleration,critical Pecklet number(parameter of modulation)SNRWsn=1051erg ismn0=1cm-3maximum value-typical in interstellar mediumdiffusion should be anomalously slow near the shock(upstream and downstream)cosmic ray streaming instability in shock precursorBell 1978, Lagage & Cesarsky 1983, McKenzie & Vlk 1982, Achterberg 1983,Vlk et al. 1988, Fedorenko 1990, Bell & Lucek 2000, 2001

  • Nagano & Watson 2000Bohm limitgalacticextra-galactic?kneestandard assumption B ~ Bism Bohm diffusion

  • Berezhko &Elliison 1999nonlinear shock modification by cosmic ray pressure for high Mach shocks Axford 1977, 1981Eichler 1984Berezhko et al. 1996Malkov et al. 2000not power law spectrum at the shock

  • 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 emissionMHz = 4.6 BGEe,GeV2E = 50 MeV 30 GeV (100 GeV for IR) = 1.9 2.5We = 1048 1049 ergGinzburg &Syrovatskii 1964Shklovsky 1976nonthermal X-rayskeV = 1 BG(Ee/120 TeV)2 max ~ 100 TeV

    SN1006 Koyama et al. 1995Cas A Allen et al. 1997RX J1713-39 Koyama et al. 1997RX J0852-46 (Vela jr) Slane et 2001-rays (0) = 30-3000 MeV Cygni, IC443Esposito et al. 1996Sturner & Dermer 1996TeV rayselectrons/protonsmax ~ 100 TeV

    SN1006 Tanimori et al 1998RX J1713 Muraishi et al. 2000 Aharonian et al. 2004Cas A Aharonian et al. 2001RX J0852-46 (Vela jr)G338.3-0.0; G23.3-0.3; G8.7-0.1 Aharonian et al. 2005 esynchrotroneinverse Compton = 0(Ee/mec2)2p0SNRnot confirmedby HESS (2004) !

  • confrontation with observationsproblems:Galactic sources should work up to (1-3)1018 eV (Fe ?) no VHE gamma-rays from not very young SNRs tsnr 3103 yr average cosmic ray source spectrums = 2.1 - 2.4 (depending on propagation model)

  • Ptuskin & Zirakashvili 2003Wsn = 1051 erg, Bism = 5 G, n0 = 0.4 cm-3 cr = 0.5, = 0.04, a = 0.3under extreme conditions:Emax 1017Z(ush/3104km/s)2 (cr/0.5)Mej1/3n1/6 eV Bmax 103(ush/3104km/s)n1/2 Gstrong cosmic-ray streaming instability (B B0), Bell & Lucek 2000, 2001 - non-linear wave interactions of Kolmogorov type in shock precursorPtuskin & Zirakashvili 2003, 2005B > B0B < B0maximum momentum of accelerated protonsabandonment of Bohm limit hypotheses

  • average source spectrumspectrum atthe shock instantaneousSNR luminosityin run-awaycosmic raysaveragecosmic-raysourcespectrumadiabatic stage Q ~ crsnWsnp-4 (Sedov) - universal spectrum !ejecta-dominated stageSNII in RSG wind: Q ~ p-6.5 at star~ r -10SNI in uniform medium: Q ~ p-7.0 (Chevalier Nadyozhin)SN ratestep functiondelta function

  • hot bubble0.013 cm-3, 3Gism R=60pc n=1cm-3 denseRSG windWeaver et al. 1977Chevalier & Liang 1989KASCADESNIIRoth et al. 2003 Eknee 61015 Z eV, ~ crWsnM1/2(Mejuw)-1

    Emax 41016 Z eV at tmin = 7 days star~ r-10

    M=10-5uw=10km/sRw=2pcPtuskin & Zirakashvili 2004expected break of all particle spectrum = 0.5

  • Nagano & Watson 2000galacticextra-galactic?kneedispersion of SNs? reacceleration?early transition to extragalactic CRs?2nd knee

  • Reacceleration by multiple shocks Reacceleration in plerionsSNRSNRSNRpulsar windSNR = 41015Z eV 1019Z eV

    Bell 1991, 2000, Berezhko 1993u= Bur/cOB association: u=3103 km/s, B=10-5 G, R=30 pcf ~ 1/p3ta ~ R/(Fshu) at Di < uR ~ D/(Fshu2) at Di > uRR uEmax ~ 1017Z eV

    Axford & Ip 1991, Bykov & Toptygin 1990, 2001Klepach et al. 2000terminationshockCrab pulsarfew msec pulsar

  • SummaryMaximum energy of accelerated particles strongly depends on SNR age in the presence of cosmic-ray streaming instability accompanied by non-linear wave dissipation. Emax can reach 1017Z eV in very young SNRs (with corresponding increase of random magnetic field to up to 10-3 G) and may fall down to less than 1011Z eV at the end of Sedov stage. Standard estimate of Emax based on the Bohm limit calculated for interstellar magnetic filed strength is not justified.This gives a clue to understanding why SNRs are not bright in very high energy -rays at t > 3103 yr.

    Average source spectrum ~ p-4 up to ~ 61015Z eV is formed during adiabatic (Sedov) stage of SNR evolution provided constant fraction of incoming gas momentum flux goes to cosmic ray pressure at the shock. Steep power-law spectrum above this energy is produced at the preceding ejecta-dominated stage. The knee observed at 41015 eV may mark the transition from ejecta-dominated to adiabatic evolution of SNR shocks which accelerate cosmic rays.

  • strong streaming instability and non-linear wave interactions in shock precursor ( ):abandonment of Bohm limit hypothesesPtuskin & Zirakashvili 2003eq. for cosmic rays(1D, u=const)eq. for mhd waves(wk is spectral energy density)supersonicconvectiongrowth rateDfin agreement withBell & Lucek 2001lineardampingnonlinear waveinteractions ofKolmogorov type~ kB(>k)/(4)1/2Verma et al. 1996eq. for maximummomentum

  • diffusion coefficient:growth rate:

  • streaming instability in shock precursor(no damping)Alfven velocitycosmic-ray pressurewave energy densityweak random field:strong random field:characteristic velocity of waves~ 0.5 for very strong shock