Magnetic Fields inSupernova Remnants

Kashi & Urumqi, 2005 Sept. 7th-14th

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.)

Evolution of SNRs(based on Woltjer 1972)

log

Rad

ius

log Time

R t R t2/5 R t2/7 R t1/4

FreeExpansion

Adiabatic Radiation RadiationSedov internal pressure momentum

merg

ing into

the inte

rste

llar

mediu

m

Magnetic Field and Evolution of SNRs

Magnetic pressure number

RH = magnetic pressure = B0

2/8 476

B02(mGs)

.

dynamic

pressure 1/20vs

2 n0(cm-

3)vs2(100km/s)

100

10

1

0.1

0.01

RH

10Gs 100Gs 1mGs 10mGs

B0

10-8 dyne cm-2

10-7 dyne cm-2

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)

Observation

of Magnetic

Fields

Faraday rotation angle:

rot(rad) = RM(rad/m2) (m)2

Rotation measure: RM(rad/m2) = 8.1105 N(cm-3) B‖(G) dz(pc)

(rad) = 0(rad) + RM(rad/m2)(m)2 +n

G127.1+0.5=11cm E-Vectors = 6cm

Ambiguity of Rotation Measure HB9 100-m-RT

+

(rad) = 0.2+114 (m)2

6cm

11cm 21cm

Ambiguity of Rotation Measure HB9 100-m-RT

+

(rad) = 0.2+114 (m)2

6cm

11cm 21cm

S1476cm

Urumqi25m-RT

TP + B-Field + Pulsar ( )

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

Young Shells

Tycho 10.55 GHzTP +B-Field 100-m-RT

Fine structure at 15 arcsec scale (0.2 pc) VLA 5 GHz (Wood et al., 1992)

Tycho’s SNR

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

Magnetic Field Direction in SNRs

(Whiteoak & Gardner 1968)

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

Evolved Shells

CTB1 10.55 GHzTP+B-Field 100-m-RT

The Orientation of bilateral SNRs and the Galactic Magnetic Field

G127.1+0.5 HC30 G93.3+6.9

Magnetic Field Direction in SNRs

(Whiteoak & Gardner 1968)

Magnetic Field Direction in G179.0+2.5

= 6cmTP + E-Vectors

Old SNR with radial B-Field!!

Filled-center SNRs (Tau A)

VLA 21cm/6cm, (Bietenholz & Kronberg 1990)

100-m-RT 32GHz,(Reich 2002)

G21.5-0.9

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

Depolarization

Polarization degree:

P(%) = 3+3 sin B02 / (B0

2 + Br2), (Burn

1966)

3+3

=2r 2.83r

R=1

r

8.1 105 n B║ r(rad)n(cm-

3)B(Gs)r(pc)

Variation of total power rVariation of pol. Int. Sedov equations + strong

shock n0, B0, E0, tage, Vshock r

I

Magnetic Field Strength 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:

= relative radiating volumeR = radius (arcmin)d = distance (kpc)S1GHz = flux density (Jy)B = magnetic induction (µGs)

Tycho ~ 0.2 mGG127.1+0.5 ~ 12G

Bmin = 199 -2/7 R-6/7 d-2/7 S1GHz

2/7

(Pacholczyk 1970)

RHTycho 0.1

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)

W44 (Claussen et al. 1997)

0.28±0.09mG

W51C (Brogan et al. 2000)

1.5±0.05mG 1.9±0.10mG

OH 1720 Zeeman Data• 10 sources observed• Magnetic fields between 0.1 and a few

mG• W44:

• W51C Magnetic pressure 10-7 dyne cm-2

Dynamic pressure: 1/20Vs2 2 10-7

dyne cm-2

Magnetic pressure: B2/8 3 10-9 dyne cm-2

Thermal pressure: nkT 6-8 10-9 dyne cm-2

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.

Thank You

On Xinwei day the new star fadedaway (China, 14th century B.C.)