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Fund. Physics & Astrophysics of Supernova Remnants. Lecture #1 What SNRs are and how are they observed Hydrodynamic evolution on shell-type SNRs Microphysics in SNRs – electron-ion equ Lecture #2 Microphysics in SNRs - shock acceleration Statistical issues about SNRs Lecture #3

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Fund physics astrophysics of supernova remnants l.jpg
Fund. Physics & Astrophysicsof Supernova Remnants

  • Lecture #1

    • What SNRs are and how are they observed

    • Hydrodynamic evolution on shell-type SNRs

    • Microphysics in SNRs – electron-ion equ

  • Lecture #2

    • Microphysics in SNRs - shock acceleration

    • Statistical issues about SNRs

  • Lecture #3

    • Pulsar wind nebulae


Order of magn estimates l.jpg
Order-of magn. estimates

  • SN explosion

    • Mechanical energy:

    • Ejected mass:

      • VELOCITY:

  • Ambient medium

    • Density: Mej~Mswept when:

      • SIZE:

      • AGE:


Classical radio snrs l.jpg

Tycho – SN 1572

“Classical” Radio SNRs

  • Spectacular shell-like morphologies

    • comparedto optical

    • polarization

    • spectral index(~ – 0.5)

      BUT

  • Poor diagnostics on the physics

    • featureless spectra (synchrotron emission)

    • acceleration efficiencies ?


Slide4 l.jpg

90cm Survey4.5 < l < 22.0 deg(35 new SNRs found;Brogan et al. 2006)

Blue: VLA 90cm Green: Bonn 11cmRed: MSX 8 mm

  • Radio traces both thermal and non-thermal emission

  • Mid-infrared traces primarily warm thermal dust emission

A view of Galactic Plane


Snrs in the x ray window l.jpg
SNRs in the X-ray window

  • Probably the “best” spectral range to observe

    • Thermal:

      • measurement of ambient density

    • Non-Thermal:

      • synchrotron-emitting electrons are near the maximum energy (synchrotron cutoff)


X ray spectral analysis l.jpg
X-ray spectral analysis

  • Low-res data

    • Overall fit with thermal models

  • High-res data

    • Abundances of elements

    • Single-line spectroscopy!


Shell type snr evolution a classical and wrong scenario l.jpg
Shell-type SNR evolutiona “classical” (and wrong) scenario

Isotropic explosion and further evolution

Homogeneous ambient medium

Three phases:

  • Linear expansion

  • Adiabatic expansion

  • Radiative expansion

Isotropic

Homogeneous

Linear

Adiabatic

Radiative


Basic concepts of shocks l.jpg

r

V

shock

Strong shock

If

Basic concepts of shocks

  • Hydrodynamic (MHD) discontinuities

  • Quantities conserved across the shock

    • Mass

    • Momentum

    • Energy

    • Entropy

  • Jump conditions(Rankine-Hugoniot)

  • Independent of the detailed physics


Forward and reverse shocks l.jpg

Forward

shock

Density

Reverse

shock

Radius

Forward and reverse shocks

  • Forward Shock: into the CSM/ISM(fast)

  • Reverse Shock: into the Ejecta (slow)


Dimensional analysis and self similar models l.jpg
Dimensional analysisand Self-similar models

  • Dimensionality of a quantity:

  • Dimensional constants of a problem

    • If only two, such that M can be eliminated, THEN evolution law follows immediately!

  • Reduced, dimensionless diff. equations

    • Partial differential equations (in r and t) then transform into total differential equations (in a self-similar coordinate).


Early evolution l.jpg
Early evolution

  • Linear expansion only if ejecta behave as a “piston”

  • Ejecta with and

  • Ambient medium with and

  • Dimensional parameters and

  • Expansion law:


A self similar model l.jpg
A self-similar model

(Chevalier 1982)

  • Deviations from “linear” expansion

  • Radial profiles

    • Ambient medium

    • Forward shock

    • Contact discontinuity

    • Reverse shock

    • Expanding ejecta


Evidence from sne l.jpg
Evidence from SNe

  • VLBI mapping (SN 1993J)

  • Decelerated shock

  • For an r-2 ambient profileejecta profile is derived


The sedov taylor solution l.jpg
The Sedov-Taylor solution

  • After the reverse shock has reached the center

  • Middle-age SNRs

    • swept-up mass >> mass of ejecta

    • radiative losses are negligible

  • Dimensional parameters of the problem

  • Evolution:

  • Self-similar, analytic solution (Sedov,1959)


The sedov profiles l.jpg

Density

Pressure

Temperature

The Sedov profiles

  • Most of the mass is confined in a “thin” shell

  • Kinetic energy is also confined in that shell

  • Most of the internal energy in the “cavity”


Thin layer approximation l.jpg
Thin-layer approximation

  • Layer thickness

  • Total energy

  • Dynamics

Correct value:1.15 !!!


What can be measured x rays l.jpg

from spectral fits

What can be measured (X-rays)

… if in the Sedov phase


Testing sedov expansion l.jpg

Deceleration parameter

Tycho SNR (SN 1572) Dec.Par. = 0.47

SN 1006 Dec.Par. = 0.34

Testing Sedov expansion

Required:

  • RSNR/D(angular size)

  • t(reliable only for historical SNRs)

  • Vexp/D(expansion rate, measurable only in young SNRs)


Other ways to measure the shock speed l.jpg
Other ways to “measure”the shock speed

  • Radial velocities from high-res spectra(in optical, but now feasible also in X-rays)

  • Electron temperature from modelling the (thermal) X-ray spectrum

  • Modelling the Balmer line profile in non-radiative shocks (see below)


End of the sedov phase l.jpg
End of the Sedov phase

  • Sedov in numbers:

  • When forward shock becomes radiative: with

  • Numerically:


Beyond the sedov phase l.jpg

Internal energy

Kinetic energy

Beyond the Sedov phase

  • When t>ttr, energy no longer conserved.What is left?

  • “Momentum-conservingsnowplow” (Oort 1951)

  • WRONG !! Rarefied gas in the inner regions

  • “Pressure-driven snowplow” (McKee & Ostriker 1977)


Numerical results l.jpg

2/5

2/7=0.29

1/4=0.25

Numerical results

(Blondin et al 1998)

0.33

ttr

Blondin et al 1998


An analytic model l.jpg
An analytic model

Bandiera & Petruk 2004

  • Thin shell approximation

  • Analytic solution

H either positive (fast branch)

limit case: Oort or negative (slow branch)

limit case: McKee & Ostriker

H,K from initial conditions


Inhomogenous ambient medium l.jpg
Inhomogenous ambient medium

  • Circumstellar bubble (ρ~ r -2)

    • evacuated region around the star

    • SNR may look older than it really is

  • Large-scale inhomogeneities

    • ISM density gradients

  • Small-scale inhomogeneities

    • Quasi-stationary clumps (in optical) in young SNRs (engulfed by secondary shocks)

    • Thermal filled-center SNRs as possibly due to the presence of a clumpy medium


Collisionless shocks l.jpg
Collisionless shocks

  • Coulomb mean free path

    • Collisional scale length (order of parsecs)

    • Larmor radius is much smaller (order of km)

  • High Mach numbers

    • Mach number of order of 100

  • MHD Shocks

    • B in the range 10-100 μG

  • Complex related microphysics

    • Electron-ion temperature equilibration

    • Diffusive particle acceleration

    • Magnetic field turbulent amplification


Electron ion equilibration l.jpg

(Cargill and Papadopoulos 1988)

(Spitzer 1978)

Electron & Ion equilibration

  • Naif prediction, for collisionless shocks

  • But plasma turbulence may lead electrons and ion to near-equilibrium conditions

  • Coulomb equilibration on much longer scales


Optical emission in sn1006 l.jpg
Optical emission in SN1006

  • “Pure Balmer” emissionin SN 1006

  • Here metal lines are missing (while they dominate in recombination spectra)

    • Extremely metal deficient ?


Non radiative emission l.jpg
“Non-radiative” emission

  • Emission from a radiative shock:

    • Plasma is heated and strongly ionized

    • Then it efficiently cools and recombines

    • Lines from ions at various ionization levels

  • In a “non-radiative” shock:

    • Cooling times much longer than SNR age

    • Once a species is ionized, recombination is a very slow process

  • WHY BALMER LINES ARE PRESENT ?


The role of neutral h l.jpg
The role of neutral H

(Chevalier & Raymond 1978, Chevalier, Kirshner and Raymond 1980)

  • Scenario: shock in a partially neutral gas

  • Neutrals, not affected by the magnetic field, freely enter the downstream region

  • Neutrals are subject to:

    • Ionization (rad + coll)[LOST]

    • Excitation (rad + coll)Balmer narrow

    • Charge exchange (in excited lev.)Balmer broad

  • Charge-exchange cross section is larger at lower vrel

  • Fast neutral component more prominent in slower shocks


H alpha profiles l.jpg

(Kirshner, Winkler and Chevalier 1987)

(Hester, Raymond and Blair 1994)

Cygnus Loop

H-alpha profiles

  • MEASURABLE QUANTITIES

  • Intensity ratio

  • Displacement (not if edge-on)

  • FWHM of broad component (Ti !!)

  • FWHM of narrow component

  • (T 40,000 K – why not fully ionized?)


Snr 1e 0102 2 7219 l.jpg

Optical

X-rays

Radio

SNR 1E 0102.2-7219

(Hughes et al 2000, Gaetz et al 2000)

  • Very young and bright SNR in the SMC

  • Expansion velocity (6000 km s-1, if linear expansion)measured in optical (OIII spectra) and inX-rays (proper motions)

  • Electron temperature~ 0.4-1.0 keV, whileexpected ion T ~ 45 keV

  • Very smallTe/Ti, orTimuch less than expected?Missing energy in CRs?


Lectures 2 3 l.jpg
Lectures #2 & #3

  • Shock acceleration

    • The prototype: SN 1006

    • Physics of shock acceleration

    • Efficient acceleration and modified shocks

  • Pulsar Wind Nebulae

    • The prototype: the Crab Nebula

    • Models of Pulsar Wind Nebulae

    • Morphology of PWN in theory and in practice

    • A tribute to ALMA


The strange case of sn1006 l.jpg

Tycho with ASCA

Hwang et al 1998

The “strange case” of SN1006

“Standard”X-ray spectrum


Thermal non thermal l.jpg
Thermal & non-thermal

  • Power-law spectrum at the rims

  • Thermal spectrum in the interior


Diffusive shock acceleration l.jpg

shock

flow speed

X

Diffusive shock acceleration

  • Fermi acceleration

    • Converging flows

    • Particle diffusion(How possible, in acollisionless plasma?)

  • Particle momentum distributionwhere r is the compression ratio (s=2, if r = 4)

  • Synchrotron spectrum

  • For r = 4, power-law index of -0.5

  • Irrespectively of diffusion coefficient

(in the shock reference frame)


The diffusion coefficient l.jpg
The diffusion coefficient

  • Diffusion mean free path(magnetic turbulence)(η > 1)

  • Diffusion coefficient


And its effects l.jpg
…and its effects

  • Acceleration time

  • Maximum energy

  • Cut-off frequency

    • Naturally located near the X-ray range

    • Independent of B


Basics of synchrotron emission l.jpg
Basics of synchrotron emission

  • Emitted power

  • Characteristic frequency

  • Power-law particle distribution

  • If then

  • Synchrotron life time


Sn 1006 spectrum l.jpg
SN 1006 spectrum

  • Rather standard( -0.6)power-law spectrum in radio(-0.5 for a classical strong shock)

  • Synchrotron X-rays below radio extrapolationCommon effect in SNRs(Reynolds and Keohane 1999)

  • Electron energy distribution:

  • Fit power-law + cutoff to spectrum:

    “Rolloff frequency”


Measures of rolloff frequency l.jpg
Measures of rolloff frequency

  • SN 1006 (Rothenflug et al 2004)

  • Azimuthal depencence of the breakChanges in tacc? or in tsyn? ηof order of unity?


Dependence on b orientation l.jpg
Dependence on B orientation?

  • Highly regular structure of SN 1006.Barrel-like shape suggested (Reynolds 1998)

  • Brighter where B is perpendicular to the shock velocity?

Direction of B ?


Radio x ray comparison l.jpg
Radio – X-ray comparison

(Rothenflug et al 2004)

  • Similar pattern (both synchrotron)

  • Much sharper limb in X-rays (synchrotron losses)


Slide44 l.jpg

  • Evidence for synchrotron losses of X-ray emitting electrons

  • X-ray radial profile INCONSISTENT with barrel-shaped geometry (too faint at the center)


3 d geometry polar caps l.jpg

Ordered magnetic field

(from radio polarization)

3-D Geometry. Polar Caps?

Polar cap geometry:

electrons accelerated

in regions with quasi-parallel field

(as expected from the theory)


Statistical analysis l.jpg
Statistical analysis

(Fulbright & Reynolds 1990)

Expected morphologies in radio

Polar cap SNR

(under variousorientations)

Barrel-like SNR

(under variousorientations)


The strength of b l.jpg
The strength of B ?

  • Difficult to directly evaluate the value of the B in the acceleration zone.νrolloffis independent of it !

  • “Measurements” of B must rely on some model or assumption


Very sharp limbs in sn 1006 l.jpg

Chandra

ASCA

Very sharp limbs in SN 1006


B from limb sharpness l.jpg
B from limb sharpness

(Bamba et al 2004)

Profiles of resolved non-thermal X-ray

filaments in the NE shell of SN 1006

Length scales 1” (0.01 pc) upstream20” (0.19 pc) downstream

Consistent withB ~ 30 μG


A diagnostic diagram l.jpg

rolloff

tsync> tacc

 > Bohm

A diagnostic diagram

  • Acceleration timetacc = 270 yr

  • Derivation of the diffusioncoefficients:u=8.9 1024 cm2s-1d=4.2 1025 cm2s-1(Us=2900 km s-1)to compare withBohm=(Emaxc/eB)/3


Non linear shock acceleration l.jpg
Non-linear shock acceleration

  • Such high values of B are not expected in the case of pure field compression(3-6 μG in the ISM, 10-20 μG in the shock – or even no compression in parallel shocks)

  • Turbulent amplification of the field?

  • Possible in the case of efficient shock acceleration scenario: particles, streaming upstream, excite turbulence

(e.g. Berezhko; Ellison; Blasi)


Shock modification l.jpg

Dynamical effects of theaccelerated particles ontothe shock structure(Drury and Voelk 1981)

Shock modification

  • Intrinsically non linear

  • Shock precursor

  • Discontinuity (subshock)

  • Larger overall compression factor

  • Accelerated particle distribution is no longer a power-law


Deviations from power law l.jpg

Blasi Solution

Thermal

Deviations from Power-Law

  • In modified shocks,acc. particles withdifferent energiessee different shockcompression factors.Higher energy Longer mean free path Larger compress.factor Harder spectrum

  • Concavity in particledistribution.(also for electrons)

Standard PL


Gamma ray emission l.jpg

Synchrotron

νFν

IC

γ-ray

Radio

X-ray

Gamma-ray emission

  • Measurement of gamma-ray emission, produced by the same electrons that emit X-ray synchrotron, would allow one to determine the value of B.


Slide55 l.jpg

(Ellison et al 2000)

  • On the other hand, there is another mechanism giving Gamma-ray emission

    • accelerated ions

    • p-p collisions

    • pion production

    • pion decay (gamma)

  • Lower limit for B

  • Need for “targets”(molecular cloud?)

  • Efficiency in in accelerating ions?(The origin of Cosmic rays)


Tev telescopes generation l.jpg
TeV telescopes generation

  • H.E.S.S. Cherenkov telescopes

  • Observations :

  • RX J0852.0-4622(Aharonian et al 2005)

  • Upper limits on SN 1006(Aharonian et al 2005)

  • RX J1713.7-3946(Aharonian et al 2006)


Observ of rx j0852 0 4622 l.jpg
Observ. of RX J0852.0-4622

  • Good matching between X-rays and gamma-rays

  • CO observation show the existence of a molecular cloud

  • Pion-decay scenario slightly favoured. Nothing proved as yet


Indirect tests on the crs l.jpg

(Blondin and Ellison 2001)

(Decourchelle et al 2000)

Indirect tests on the CRs

  • Some “model-dependent” side effects of efficient particle acceleration

  • Forward and reverse shock are closer, as effect of the energy sink

  • HD instabilities behavior depends on the value of eff


Shock acceleration efficiency l.jpg
Shock acceleration efficiency

  • Theory predicts (~ high) values of the efficiency of shock acceleration of ions.

  • Little is known for electrons

  • Main uncertainty is about the injection process for electrons

    • Shock thickness determined by the mfp of ions (scattering on magnetic turbulence)

    • Electrons, if with lower T, have shorter mfps

    • Therefore for them more difficult to be injected into the acceleration process


The d relationship l.jpg
The Σ–D Relationship

  • Empirical relation

    • SNR surface brightness, in radio

    • SNR diameter

    • Any physicalreason forthis relation ?

(Case & Bhattacharya 1998)


A basic question l.jpg
A basic question

  • Is the correlation representative of the evolution of a “typical object”?

  • Or is, instead, the convolution of the evolution of many different objects?

  • Theorists attempts to reproduce it.

Berezhko & Voelk 2004


Dependence on ambient density l.jpg
Dependence on ambient density

(Berkhuijsen 1986)

  • Primary correlations are D-n, and Σ-n

  • Diff. ISM conditions


The prototype l.jpg

The Crab Nebula

Optical

Thermal filaments

Amorphous compon.

Radio

Filled-center nebula

No signs of shell

X-rays

More compact neb.

Jet-torus structure

X-rays

Crab Nebula - radio

Crab Nebula – Ha + cont

The “Prototype”


The crab nebula spectrum l.jpg

-0.8

-1.1

-1.5

-0.3

The Crab Nebula spectrum

(apart from optical filaments and IR bump)

Synchrotron emission

  • Radio

  • Optical

  • Soft X-rays

  • Hard X-rays


Some basic points l.jpg
Some basic points

  • Synchrotron efficiency

    • 10-20% of pulsar spin-down power

  • Powered by the pulsar

  • High polarizations (ordered field)

  • No signs of any associated shell.


Basics of synchrotron emission67 l.jpg
Basics of synchrotron emission

  • Emitted power

  • Characteristic frequency

  • Power-law particle distribution

  • If then

  • Synchrotron life time


Simple modelling l.jpg
Simple modelling

(Pacini & Salvati 1973)

  • Homogeneous models (no info on structure)

  • Magnetic field evolution

    • Early phases (constant pulsar input)

    • Later phases (most energy released)


Slide69 l.jpg

Power-law injection

With upper energy cutoff

Continuum injection

link to the pulsar spin down

Particle evolution (adiabatic vs synchrotron losses)

Evolutionary break

Adiabatic regime

(-0.3 in radio)

Synchrotron-dominated regime

(-0.8 in optical)


Kennel coroniti model 1984 l.jpg

Pulsar magnetosphere

ISM

Pulsar wind

Stellarejecta

Termination shock

Pulsar Wind Nebula

Kennel & Coroniti model (1984)

Basics of “Pulsar Wind Nebula” scenario

  • Pulsar magnetosphere

  • Pulsar wind

  • Termination shock

  • Pulsar Wind Nebula

  • Interface with theejecta (CD, FS)

  • Stellar ejecta

  • Interface with theambient medium(RS, CD, FS)

  • Ambient medium (either ISM or CSM)


The ingredients l.jpg
The ingredients

  • Pulsar wind

    • super-relativistic

    • magnetized(toroidal field)

    • isotropic

  • Termination shock

    • mass conservation

    • magnetic flux cons.

    • momentum cons.

    • energy cons.

      where (specific enthalpy)


Large and small limits l.jpg
Large and small σ limits

  • Large σ

    • weak shock

    • flow stays super-relativistic

    • neither field, nor density jump

    • inefficient in converting kinetic into thermal energy

  • Small σ

    • strong shock

    • flow braked to mildly relativistic speed

    • both field and density increase

    • kinetic energy efficienly converted


Mhd evolution in the nebula l.jpg
MHD evolution in the nebula

  • Steady solution (flow timescale << SNR age)

    • number flux cons. - magnetic flux cons.

    • momentum cons. - energy cons.

  • Asymptotic velocity !!!

    • no solution for V∞=0

    • outer expansion Vext~1500 km s-1 (for the Crab Nebula)

    • then σ~3 10-3

    • size of termination shock, from balance of wind ram pressure and nebular pressure

      Rn~10 arcsec

      (wisps region)


Radial profiles l.jpg
Radial profiles

  • Inner part with:

  • Outer part with:

  • Equipartition in the outer part:


Do we expect what observed l.jpg
Do we expect what observed?

  • Injected particles

    • power-law, between a min and a max energy

    • only 1 free parameter (n2 and p2 from the jump conditions at the termination shock)

    • plus wind parameters (L, σ and γ1 )

  • Energy evolution during radial advection


Best fit solution l.jpg
Best-fit solution

  • Parameters:

  • Fit to:


Problems ia l.jpg
Problems -Ia

  • The sigma paradox

    • A value is required, in order to get an effective slowing-down of the flow, and a high (10-20 %) synchrotron conversion efficiency

    • BUT the (magnetically driven) pulsar wind cannot have been produced with a low σ .

    • With a normal MHD evolution, the value of σ must keep constant from the acceleration region till the termination shock.


Problems ib l.jpg
Problems - Ib

  • A POSSIBLE WAY OUT

    • A tilted pulsar generates a striped wind.


Problems ic l.jpg
Problems -Ic

  • Magnetic reconnection in the wind zone (if possible) would dissipate the field.(Coroniti 1990)

  • Reconnection in the wind zone does not efficiently destroy the field. Reconnection at the termination shock is more effective.

    (Lyubarski & Kirk 1991)


Problems iia l.jpg
Problems - IIa

  • The unexpected radio emission

    • Predicted radio flux is far lower (a factor ~100) than observed.

    • No easy way to cure it. Little freedom on the particle number. Total power is fixed: more particles mean a lower γ1.

    • Radio emitting electrons as a relict. Was the Crab much more powerful in the past? Ad hoc. All PWNe are radio emitters.


Problems iib l.jpg
Problems IIb

  • Can it be “Diffusive synchrotron radiation”?(Fleishman & Bietenholz 2007)Turbulence spectral index ν.

  • Theory only for a fully turbulent field

    • Total spectrumis reproduced

    • But observedpolarization isnot explained


Non spherical structure l.jpg
Non-spherical structure

(Begelman & Li 1992)

  • Particle, moving passively along field lines (flow motion assumed to be irrotational)

  • Axisymmetric nebular field structure

  • Steady state solutions


Slide84 l.jpg

pulsar axis

van der Swaluw 2003

3C 58

MHD simulations


Elongated structures of pwne l.jpg

pulsar spin

Elongated structures of PWNe

3C 58

G5.4-0.1

G11.2-0.3

Crab Nebula


Details of the structure l.jpg

counter-jet

torus

knot

inner ring

jet

Details of the structure

Crab Nebula

Vela


Jet sizes l.jpg

4’ = 6 pc

40” =

0.4 pc

Crab Nebula (Weisskopf et al 2000)

PSR B1509-58 (Gaensler et al 2002)

13” = 0.2 pc

80” = 0.8 pc

3C 58 (Slane et al. 2004)

Vela Pulsar (Pavlov et al. 2003)

Jet sizes


Simulating pwne l.jpg
Simulating PWNe

(Komissarov, 2006; Del Zanna et al 2004, 2006)

  • Relativistic MHD codes

  • Modelling a PWN like the CrabVelocity Magnetization Max Energy


Surface brightness maps l.jpg
Surface brightness maps

Jet-Torus structure


Ingredients l.jpg
Ingredients

  • Wind parameters

    • magnetization (still small, but not too much)σ~0.02 – 0.1aaa

    • wind anisotropy ( γeq~10 γpol )

    • “filling” the jets (since B = 0)


Pwn ejecta interaction l.jpg

Reverse Shock

PWN Shock

Forward Shock

Pulsar

Termination

Shock

Pulsar Wind

Unshocked Ejecta

Shocked Ejecta

Shocked ISM

PWN

ISM

PWN-ejecta interaction

  • PWNe are confined by the associated shell-like SNR

  • Not only the SNR is detectable (like in the Crab)

  • In the Crab NebulaUV emissionassociated with aslow shock (againstthe SN ejecta)


A tribute to alma l.jpg
A TRIBUTE TO ALMA

  • SNRs and PWNs are mostly non-thermal in that spectral range.

    • no use of spectral capabilities

    • use of high spatial resolution, + wide field, + photometric stability (extended sources)

  • Is mm-submm a “new band” for SNRs, or just an extension of the radio range?

  • A study of the Crab Nebula(extension of a former work, Bandiera et al 2002)


What has been done already l.jpg

-0.28

-0.20

Spectral map

230 GHz map

What has been done already

  • Comparison of 1.3 mm (230 GHz) images (with IRAM 30-m telescope, 10” res) and radio (20 cm, VLA) maps


A further emission component l.jpg
A further emission component

  • Radio spectral index: -0.27

  • Concave spectral index from radio to mmReal effector artifact?(absolutephotometry)

  • Evidence foran additionalemission component


Component b l.jpg
Component B

  • Image obtained optimizing the subtraction of amorphous part, and filaments, of radio image (PSF matched), with best-fit weights.


The subtracted components l.jpg
The subtracted components

  • Amorphous component: consistent with an extension of the spectrum to mm, with the radio spectral index (-0.27).

  • Filaments: consistent with spectral bending (νb~80 GHz).

  • Morphologically, component B resembles more the Crab in the optical than in the radio (ALTHOUGH, in the mm range, electrons of Component B do not lose energy significantly by synchrotron).


The integrated spectrum l.jpg
The integrated spectrum

  • Radio comp (A)

  • Component B,with low freqcutoff.

  • Evidence higherthan from theerror bar.

  • Components Aand B coexistin the optical.


Physical scenario l.jpg
Physical scenario

  • Number of particles in Component B:Ntot~ 2 1048.

  • Consistent with Kennel & Coroniti)

  • Filament magnetic fields ~6times higher than the rest AND particle do not diffuse in/out of filaments (κ<100 κB).


With alma l.jpg
With ALMA

  • The same analysis, with a resolution 100 times higher.

  • Detailed mapping of Component B.

  • Separation of comp A and B also through differences in the polarization patterns.

  • Analysis of the spectral bending in individual filaments, and possibly even across the filament (B estimates).

  • Mapping B in filaments (aligned? ordered?)


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