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The role of neutrinos in the evolution and dynamics of neutron stars José A. PonsPowerPoint Presentation

The role of neutrinos in the evolution and dynamics of neutron stars José A. Pons

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the evolution and dynamics

of neutron stars

José A. Pons

University of Alicante (SPAIN)

Transparent and opaque regimes.

NS formation and n role in Supernovae.

Neutron stars and proto-NS.

Energetic considerations.

g-modes and convective instabilities.

Long term cooling.

All the previous issues in strange stars.

and opaque regimes

Neutrinos are weakly interacting particles, and in most astrophysical scenarios where they are produced their cross section is so low that neutrinos freely stream through matter.

BUT

NS, SS or matter surrounding BH reach supranuclear densities and high temperatures

(T> 1 Mev »1010 K, r»3´1014g/cm3)

In some cases, the mean free path becomes of the order (semitransparent) or even much shorter (opaque) than the scale of the object.

Opaque: proto-NS, proto-SS (T> 5 MeV, l» 1 m)

Semitransparent: SN envelope, NS (T=1-5 MeV).

Transparent: All the rest (T<1 MeV)

Mcore >1»2Msolar

T»1010 K, r»5´109 g/cm3 ,

Ye» 0.42, s »1-2 (k)

R»1000km

Photodesintegration

g +(A,Z) --> (A-4,Z-2)+ a

g +a ---> 2 n + 2 p

Electron captures

e- + (A,Z) --> (A,Z-1)+ n

- Infall (< 1 s) :
homologous

free fall

neutrinos escape freely

trapping (r>1012 g/cm3 )

- Bounce
(r>3 1014 g/cm3)

- Shock wave formation
and propagation

nuclei dissociation

neutrino losses

- Neutrino reactivation: Binding energy is 1053 erg, SN explosion kinetic energy is 1051 erg.
- Convective overturn

n diffusion/emission drives SN dynamics and NS formation

Evolution:The first minute of life

Mantle collapse: 0.1-1 s, heating, compression

Deleptonization: with Joule heating, maximum central T

Cooling: basically thermal neutrinos, from 50 MeV down to 1 MeV

Hot (»10-50 MeV), lepton rich

Large chemical and thermal gradients

Less compact (100 km)

No crust, no superfluid

Cold (T<1 MeV), Ye<0.1

Basically isothermal

More compact (R=10-15 km)

Solid crust, superfluid interior

Metastability: n-delayed collapse to BH

Convective instability (Ledoux)

Neutron fingers

Convection

Stable

Shear Instability + convection may lead to rigid rotation in a few dynamical periods.

from collapse from mergers

Hot (»10-50 MeV)

lepton rich YL»0.4

Non isolated !

Moderate diff. rotation

Supramassive only after accretion

T/W = 0.10-0.12

Rotation induced instabilities may appear after diffusion timescale

Less hot (»1-10MeV)

DeleptonizedYe<0.1

PNS + disk

???

Probably always supramassive (short lived)

Larger T/W possible ?

Collapses to BH

Long term cooling: n cooling epoch

After T drops below 1 MeV matter is transparent to neutrinos, but this does not mean that n’s become irrelevant. They just escape from the star as they are created. Actually, how a NS cools down during the first million years depends on neutrino emission processes in the core.

Cv dT/dt = -Lg – Ln + H

Fast cooling: direct URCA, quarks, kaon or pion condensate, hyperons …

en= 10N T96 erg/cm3/s; N=24-27

Standard (slow) cooling: modifiedURCA, bremstrahlung

en= 10N T98erg/cm3/s; N=20-21

Superfluidity slows down fast processes.

Bulk viscosity is the dominant mechanism to dissipate energy in pulsating, young NS (T=109-1010 K). Thus, the onset of dynamical instabilities, angular momentum loses, etc. during the first hours of life depend very

much on weak interaction processes.

The same processes that gives the neutrino emissivity will control viscous damping at early times.

EXAMPLE: direct URCA vs. modified URCA

BE CONSISTENT ! If you change your EOS (nuclear interaction, superfluidity, quark deconfinement) change accordingly your interaction processes and thermodynamics.

Absorption-emission ---- Specific heat ---- Bulk viscosity

Scattering ---- Compressibility ---- Shear viscosity

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