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Magnetic Fields in Supernova Remnants: Observations and Evolution

This article explores the observations and evolution of magnetic fields in supernova remnants, including historical comments, synchrotron emission, and the role of magnetic pressure and heat conduction. It also discusses the orientation of magnetic fields in different types of supernova remnants and provides examples of observed magnetic field strengths.

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Magnetic Fields in Supernova Remnants: Observations and Evolution

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  1. Magnetic Fields inSupernova RemnantsKashi & Urumqi, 2005 Sept. 7th-14th

  2. SNRs,some historical Comments • Synchrotron emission predicted by Alvén , Herlofson, Kiepenheuer • First detected as optical emission from the Crab nebula 1953 • Optical linear polarization discovered (Dombrovsky 1954) • Radio polarization from the Crab detected, (Mayer et al. 1957) On Jisi day, the 7th day of the month, a big new star appeared near the Ho star (China, 14th century B.C.)

  3. Evolution of SNRs(based on Woltjer 1972) merging into the interstellar medium R t R  t2/5 R  t2/7 R  t1/4 Free Expansion log Radius Adiabatic Radiation Radiation Sedov internal pressure momentum log Time

  4. Magnetic Field and Evolution of SNRs Magnetic pressure number RH = magnetic pressure = B02/8  476 B02(mGs) . dynamicpressure 1/20vs2 n0(cm-3)vs2(100km/s) 100 10 1 0.1 0.01 10-8 dyne cm-2 RH 10-7 dyne cm-2 B0 10Gs 100Gs 1mGs 10mGs

  5. Magnetic Field and Heat Conduction The evaporation of clouds depends on heat conduction dQ/dt = K gradT. For a typical cloud QK> 10⁸,  the low magnetic heat conduction reduces the evaporation significantly. The cloud may survive, a star may be born . QK = Kthermal 105 T(K)3 B(G)2 Kgyro n(cm-3)

  6. Observationof MagneticFields Faraday rotation angle: rot(rad) = RM(rad/m2) (m)2 Rotation measure:RM(rad/m2) = 8.1105 N(cm-3) B‖(G) dz(pc)

  7. G127.1+0.5 =11cm E-Vectors  = 6cm (rad) = 0(rad) + RM(rad/m2)(m)2 +n

  8. Ambiguity of Rotation Measure HB9 100-m-RT (rad) = 0.2+114 (m)2 + 21cm 11cm 6cm

  9. Ambiguity of Rotation Measure HB9 100-m-RT (rad) = 0.2+114 (m)2 + 21cm 11cm 6cm

  10. TP + B-Field + Pulsar ( ) S1476cmUrumqi25m-RT

  11. Types of SNRs • Young shells, historical SNRs: Tycho, SN1006, Kepler • Old shells, evolved SNRs: G127.1+0.5, G116.9+0.2, many others • Filled centered SNRs, Pulsar powered: Crab nebular, 3C58, …. • Combined SNRs

  12. Young Shells Tycho 10.55 GHz TP +B-Field 100-m-RT

  13. Tycho’s SNR Fine structure at 15 arcsec scale (0.2 pc) VLA 5 GHz (Wood et al., 1992)

  14. Young Shells • Predominantly radial field • Small scale variations (sub-pc scales) • Polarized fraction (PI/TP) 4 to 15% with local enhancements. A large fraction of random magnetic field exists (Reynolds & Gilmore 1993) • Radial field caused by external field directed towards observer (Whiteoak & Gardner, 1968) • Rayleigh-Taylor instabilities between shock and ejecta, streching of magnetic field

  15. Magnetic Field Direction in SNRs (Whiteoak & Gardner 1968)

  16. Young Shells • Predominantly radial field • Small scale variations (sub-pc scales) • Polarized fraction (PI/TP) 4 to 15% with local enhancements. A large fraction of random magnetic field exists (Reynolds & Gilmore 1993) • Radial field caused by external field directed towards observer (Whiteoak & Gardner, 1968) • Rayleigh-Taylor instabilities between shock and ejecta, streching of magnetic field

  17. Evolved Shells CTB1 10.55 GHz TP+B-Field 100-m-RT

  18. The Orientation of bilateral SNRs and the Galactic Magnetic Field G127.1+0.5 HC30 G93.3+6.9

  19. Magnetic Field Direction in SNRs (Whiteoak & Gardner 1968)

  20. Magnetic Field Direction in G179.0+2.5 • = 6cm TP + E-Vectors Old SNR with radial B-Field!!

  21. Filled-center SNRs (Tau A) 100-m-RT 32GHz, (Reich 2002) VLA 21cm/6cm, (Bietenholz & Kronberg 1990)

  22. G21.5-0.9 Nobeyama Array 22.3 GHz 100-m-RT 32 GHz, (Reich et al. 1998)

  23. Depolarization Polarization degree: P(%) = 3+3 sin  B02 / (B02 + Br2), (Burn 1966) 3+3  r (rad) n(cm-3) B(Gs) r(pc)   8.1 105 n B║  r R=1 Variation of total power  r Variation of pol. Int.   =2r 2.83r Sedov equations + strong shock I  n0, B0, E0, tage, Vshock r

  24. Assumption: Minimum total energy of electrons, protons and magnetic field. For =-2 (flux density spectral index = -0.5), and heavy particle energy 100 times electron energy, lower frequency cut 107Hz, upper cut 1011Hz: Magnetic Field Strength Bmin= 199  -2/7  R-6/7  d-2/7  S1GHz2/7 (Pacholczyk 1970)  = relative radiating volume R = radius (arcmin) d = distance (kpc) S1GHz= flux density (Jy) B = magnetic induction (µGs) Tycho ~ 0.2 mG G127.1+0.5 ~ 12G RHTycho  0.1

  25. Magnetic Field Strength:the OH Line at 1720 MHz • OH first detected (Weinreb et al. 1963) • Maser theory (Litvak et al. 1966) • Collision pumping (Elizur 1976) • OH about 100 AU behind shock front (Hollenbach & McKee 1989), (Neufeld & Dalgarno 1989) • Zeeman splitting 1.31 kHz/mG (Heiles et al. 1993), (Frail et al. 1994, W28)

  26. W44(Claussen et al. 1997) 0.28±0.09mG

  27. W51C(Brogan et al. 2000) 1.5±0.05mG 1.9±0.10mG

  28. OH 1720 Zeeman Data • 10 sources observed • Magnetic fields between 0.1 and a few mG • W44: • W51C Dynamic pressure: 1/20Vs2  2 10-7 dyne cm-2 Magnetic pressure: B2/8 3 10-9 dyne cm-2 Thermal pressure: nkT  6-810-9 dyne cm-2 Magnetic pressure  10-7 dyne cm-2

  29. Conclusions What can we learn from magnetic field observation? • Interaction of SNRs with the Galactic magnetic field • SNR parameters • In general, the dynamics of SNRs is not affected by the magnetic field • In SNRs postshock regions with strong cooling the magnetic field may have increased influence on the dynamics.

  30. Thank You On Xinwei day the new star faded away (China, 14th century B.C.)

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