Transcript Consequences of Neutrino Emission from a Phase

K S Cheng
Department of Physics
University of Hong Kong
Collaborators:
W.M. Suen (Wash. U)
Lap-Ming Lin (CUHK)
T.Harko & R. Tian (HKU)
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It is proposed that strange matter is the most stable form of matter in high
density (e.g. Witten 1984)
Strange matter can be formed in various astrophysical situations, e.g. early
Universe (Witten 84), in the core of proto-neutron stars (e.g. Takahara et al.
85), accreting binaries (Cheng & Dai 96) etc.
However,the exact phase transition process is still an open question. It can
begin with a single quark seed at the center of the star and grow to the
entire star via either slow combustion or fast detonation.
Using the thermodynamics equilibrium, conservation of Baryon and
conservation of charge, Glendenning (1992) shows that hybrid stars, which
contain a mixture of quark droplets and normal matter, is more possible.
1. Introduction
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In our study we did not consider the detail formation process from
normal matter to quark matter. We simply assume that a neutron
star suddenly undergoes a phase-transition. We use a 3D Newtonian
hydrodynamic code to study the consequences of phase-transitioninduced collapse.
This code solves a set of non-viscous Newtonian fluid flow
equations, which describe the motion of fluid inside the star by
using a standard high-resolution shock capturing scheme with
Riemann solver. The code is developed by the Washington group and
has been applied to study various neutron star dynamical problems,
e.g. (Gressman et al. 1999,2002; Lin & Suen 2004, Lin et al. 2006).
2. Numerical Simulation
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ρ
υi
P
mass density of the fluid
Cartesian components of the velocity
fluid pressure
2.1 Basic equations
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τ
total energy density
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ε
Φ
internal energy per unit mass of the fluid
Newtonian potential
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The system is completed by specifying an equation of
state P = P(ρ, ε)
2.1 Basic equations
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At t=0-, the initial configuration of star is determined by hydrostatic
equilibrium with an assumed neutron star EOS
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At t=0+, we switch the EOS to the following form:
ρtr is chosen when Pq =0
and
2. Numerical Simulation
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ρ =ρtr
Dotted line – Glendenning 1992
Solid line – our approximated model
Pressure profiles before and
after phase transition
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Freq. is roughly scaled
as ρ1/2
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figure has no phase transition, it indicates that the
oscillation is unlikely due to numerical fluctuation.
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Definition of
neutrinosphere
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τeff effective neutrino optical depth
κeff effective opacity
Janka (2001)
3.1 Neutrinosphere 10
Balantekin & Yuksel (2005)
3.2 Neutrino luminosity 11
Temperature at the neutrinosphere
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Density at the neutrinosphere
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Time evolution of temperature and
density profiles
3.1 Neutrinosphere 14
Goodman et al. (1987)
4.3 e± pair production rate 15
Annihilated e± pair energy luminosity has the same pulsation-like time
evolution. The efficiency of neutrino converting into e± is very low most of the
time except it increases to almost 100% at the peak of some pulses (very high
neutrinosphere temperature (Tian 2008).
4.3 e± pair production rate 16
Requirement
4.4 Mass ejection 17
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Whether the absorbed energy is
enough to eject mass is
determined discretely at each time
slice with an interval of 0.0075 ms
Suppose at T1 a layer of mass is
ejected; until T2 no mass is ejected
The outer edge of the ejected mass
distribution can reach the speed of
light
e± pairs is created outside the star;
the ejected mass will absorb e±
pairs created in certain area and be
accelerated
4.4 Mass ejection 18
(ms)
(ms)
4.4 Mass ejection 19
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Profiles
◦ Complicated and
irregular
◦ Multi-peaked or
single-peaked
Durations (T)
◦ ~ 5 ms to ~ 5 103 s,
Typically ~ a few
seconds
Variabilities (T)
◦ ~ 1 ms, even ~ 0.1
ms, Typically ~ 10-2
T
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Extremely intensity neutrino pulses can result from the
phase-transition-induced collapse of neutron stars due
to density and temperature oscillation; the
temperatures of these pulse neutrinos can be as high
as 10-20MeV, which is significantly higher than the
non-oscillating case (~5MeV).
These high energy neutrinos can enhance the efficiency
of electron/positron pair creation rate, which may blow
off part of surface material from the stars and
accelerate them to extremely relativistic speed, which
result in gamma-ray bursts when they collide with
each other or with ISM.
Summary and Discussion 22