Abstract
Neutron stars provide a window into the properties of dense nuclear matter. Several recent observational and theoretical developments provide powerful constraints on their structure and internal composition. Among these are the first observed binary neutron star merger, GW170817, whose gravitational radiation was accompanied by electromagnetic radiation from a short γ-ray burst and an optical afterglow believed to be due to the radioactive decay of newly minted heavy r-process nuclei. These observations give important constraints on the radii of typical neutron stars and on the upper limit to the neutron star maximum mass and complement recent pulsar observations that established a lower limit. Pulse-profile observations by the Neutron Star Interior Composition Explorer (NICER) X-ray telescope provide an independent, consistent measure of the neutron star radius. Theoretical many-body studies of neutron matter reinforce these estimates of neutron star radii. Studies using parameterized dense matter equations of state (EOSs) reveal several EOS-independent relations connecting global neutron star properties.
Highlights
Composition controls the heat flow into the atmosphere [5]
Based on lattice quantum chromodynamics (QCD) calculations, it is expected that matter at very high densities consists of asymptotically free quark matter, but since these calculations [9] are valid only at densities larger than about 40ρs, and the maximum density in a neutron star is less than about 10ρs [10], it is not clear whether such a phase exists in neutron star interiors
Laboratory nuclei provide important details of the equation of state (EOS) up to ρs, and heavy-ion collisions add some constraints for densities up to about 5ρs, but only for hot, nearly symmetric nuclear matter (SNM)
Summary
The PNM EOS is significant for two reasons It establishes strong constraints on the nuclear symmetry energy, and second, it closely approximates the neutron star EOS from densities of about ns/2, where the heterogeneous crust containing nuclei makes the transition to homogeneous nucleon matter, up to 2–3ns or even higher, where hyperons, a kaon or pion condensate, or quark matter might appear. There are two ways to define the nuclear symmetry energy: It can be viewed as the difference S(nB) between the energies of PNM and SNM at a given density, or it can be related to the lowest-order coefficient S2(nB) in a Taylor expansion of the energy in the neutron excess 1 − 2x: S(nB ) = EB(nB, 0) − EB(nB, 1/2), S2(nB ). Where EPNM and PPNM are the PNM energy and pressure, respectively, and B ࣕ −EB(ns, 1/2) 16 MeV is the bulk binding energy at saturation
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