By Johannes O Ruoff
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Additional info for The numerical evolution of neutron star oscillations
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Concerned. However, it is in particular the oscillations of neutron stars that are very sensitive to local changes in the equation of state, which are due to the different behavior of the neutron star matter under varying pressure. It is therefore much more interesting to use realistic equations of state that take into account the underlying microphysics which determines the state of the matter as a function of pressure and temperature. For comprehensive overviews on realistic equations of state, see [4, 5, 6].
5, where we have a resolution of about 10 Hz, we have to evolve up to t = 100 ms. With a time step size that is somewhat smaller than 1·10−4ms we then need more than one million integration steps! Model 3 is unstable, which is nicely confirmed by the evolution. With the eigenvalue code Chapter 3. 2: Evolution of u for a polytropic stellar model with central density ǫ0 = 3·1015 g/cm3 . 3: The power spectrum of the above wave signal shows that at least 15 eigenmodes are present. 32). 65 ms. In the logarithmic plot of Fig.
Their respective perturbations will be denoted by α, β i , hij , and kij . 18) A = −g00 = eν . In addition to the metric perturbations, we have to describe the perturbations of the energymomentum tensor Tµν . For a perfect fluid, the only quantities that can be perturbed are energy density ǫ, pressure p, and 4-velocity uµ , whose perturbations will be denoted by δǫ, δp, and δuµ , respectively. 19) where dp =: Cs2 is the square of the sound speed Cs inside the fluid. 20) ∂t kij = −∂i ∂j α + Γkij ∂k α + δΓkij ∂k eν + α [Rij + 4π(p − ǫ)γij ] + eν δRij + 4π (p − ǫ)hij + δǫ(Cs2 − 1)γij .
The numerical evolution of neutron star oscillations by Johannes O Ruoff
by Ronald
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