The role of neutrinos in The role of neutrinos in the evolution and dynamicsthe evolution and dynamics

of neutron starsof neutron stars

José A. Pons University of Alicante (SPAIN)

Transparent and opaque regimes. NS formation and 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.

Transparent, semi-transparent Transparent, semi-transparent and opaque regimesand 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.BUTNS, SS or matter surrounding BH reach supranuclear densities and high temperaturesMevg/cm

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, 1 m)

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

Transparent: All the rest (T<1 MeV)

Core collapse SNCore collapse SN

T1010 K, 109 g/cm3 ,

Yesk

R1000 km Photodesintegration +(A,Z) (A4,Z2)+ n + 2 p Electron capturese + (A,Z) (A,Z)+

Mcore >12 Msolar

Infall and bounce Infall and bounce Infall (< 1 s) : homologous free fall neutrinos escape freely trapping (g/cm3 )Bounce (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

diffusion/emission drives SN dynamics and NS formation

Evolution:Evolution: The first minute of lifeThe 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: Metastability: -delayed -delayed collapse to BHcollapse to BH

PNS’sPNS’sConvective Convective instability instability (Ledoux)(Ledoux)

Neutron fingers Convection Stable

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

PNS vs. PNSPNS vs. PNSfrom collapse from mergersfrom collapse from mergers

Hot (10-50 MeV) lepton rich YL0.4 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 (10 MeV) Deleptonized Ye<0.1 PNS + disk ??? Probably always supramassive (short lived) Larger T/W possible ? Collapses to BH

Quasi-normal modes of Quasi-normal modes of PNSsPNSs

Long term cooling: Long term cooling: cooling epoch cooling epochAfter T drops below 1 MeV matter is transparent to neutrinos, but this does not mean that ’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 = -L – L + H

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

hyperons …

erg/cm3/s; N=24-27

Standard (slow) cooling: modified URCA, bremstrahlung

erg/cm3/s; N=20-21

Superfluidity slows down fast processes.

Neutrinos and bulk viscosity Neutrinos and bulk viscosity 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 verymuch 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 viscosityScattering ---- Compressibility ---- Shear viscosity