2. The role of neutrinos in the evolution and dynamics of 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.

3. Transparent, semi-transparent 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  Mev      g/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)

4. Core collapse SN  T  10 10 K,  10 9 g/cm 3,  Y e  s  k    R  1000 km  Photodesintegration  +(A,Z)  (A  4,Z  2)+   n + 2 p  Electron captures  e  + (A,Z)  (A,Z  )+ Mcore > 1  2 Msolar

5. Infall and bounce  Infall (< 1 s) : homologous free fall neutrinos escape freely trapping (   g/cm 3 )  Bounce (   g/cm 3 )  Shock wave formation and propagation nuclei dissociation neutrino losses  Neutrino reactivation: Binding energy is 10 53 erg, SN explosion kinetic energy is 10 51 erg.  Convective overturn diffusion/emission drives SN dynamics and NS formation

6. 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

7. Metastability: -delayed collapse to BH

8. PNS’s Convective instability (Ledoux) Neutron fingers Convection Stable Shear Instability + convection may lead to rigid rotation in a few dynamical periods.

9. PNS vs. PNS from collapse from mergers  Hot (  10-50 MeV)  lepton rich Y L  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 (  10 MeV )  Deleptonized Y e <0.1  PNS + disk  ???  Probably always supramassive (short lived)  Larger T/W possible ?  Collapses to BH

10. Quasi-normal modes of PNSs

11. Long term cooling: cooling epoch After 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. C v dT/dt = -L  – L + H  Fast cooling: direct URCA, quarks, kaon or pion condensate, hyperons …       erg/cm 3 /s ; N=24-27  Standard (slow) cooling: modified URCA, bremstrahlung        erg/cm 3 /s ; N=20-21  Superfluidity slows down fast processes.

12. Neutrinos and bulk viscosity Bulk viscosity is the dominant mechanism to dissipate energy in pulsating, young NS (T=10 9 -10 10 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