List of Figures

View Image Figure 1:
Cartoon showing standard formation channels for close NS-NS binaries through binary stellar evolution. Image reproduced from [178].
View Image Figure 2:
Cartoon picture of a compact binary coalescence, drawn for a BH-BH merger but applicable to NS-NS mergers as well (although NSs are generally assumed to be non-spinning). Image adapted from Kip Thorne.
View Image Figure 3:
Isodensity contours and velocity profile in the equatorial plane for a merger of two equal-mass NSs with MNS = 1.4M ⊙ assumed to follow the APR model [3] for the NS EOS. The hypermassive merger remnant survives until the end of the numerical simulation. Image reproduced by permission from Figure 4 of [144], copyright by APS.
View Image Figure 4:
Isodensity contours and velocity profile in the equatorial plane for a merger of two equal-mass NSs with MNS = 1.5M ⊙ assumed to follow the APR model [3] for the NS EOS. With a higher mass than the remnant shown in Figure 3, the remnant depicted here collapses promptly to form a BH, its horizon shown by the dashed blue circle, absorbing all but 0.004% of the total rest mass from the original system. Image reproduced by permission from Figure 5 of [144], copyright by APS.
View Image Figure 5:
Isodensity contours and velocity profile in the equatorial plane for a merger of two unequal-mass NSs with M1 = 1.3M ⊙ and M2 = 1.6M ⊙, with both assumed to follow the APR model [3] for the NS EOS. In unequal-mass mergers, the lower mass NS is tidally disrupted during the merger, forming a disk-like structure around the heavier NS. In this case, the total mass of the remnant is sufficiently high for prompt collapse to a BH, but 0.85% of the total mass remains outside the BH horizon at the end of the simulation, which is substantially larger than for equal-mass mergers with prompt collapse (see Figure 4). Image reproduced by permission from Figure 6 of [144], copyright by APS.
View Image Figure 6:
Dimensionless binding energy E ∕M b 0 vs. dimensionless orbital frequency M Ω 0, where M0 is the total ADM (Arnowitt–Deser–Misner) mass of the two components at infinite separation, for two QE NS-NS sequences that assume a piecewise polytropic NS EOS. The equal-mass case assumes MNS = 1.35M ⊙ for both NSs, while the unequal-mass case assumes M1 = 1.15M ⊙ and M2 = 1.55M ⊙. The thick curves are the numerical results, while the thin curves show the results from the 3PN approximation. The lack of any minimum suggests that instability for these configurations occurs at the onset of mass shedding, and not through a secular orbital instability. Image reproduced by permission from Figure 16 of [305], copyright by AAS.
View Image Figure 7:
Mass-shedding indicator ( ∂(lnh)) (∂(lnh)) χ ≡ ∂r eq∕ ∂r pole vs. orbital frequency M0 Ω, where h is the fluid enthalpy and the derivative is measured at the NS surface in the equatorial plane toward the companion and toward the pole in the direction of the angular momentum vector, for a series of QE NS-NS sequences assuming equal-mass components. Here, χ = 1 corresponds to a spherical NS, while χ = 0 indicates the onset of mass shedding. More massive NSs are more compact, and thus able to reach smaller separations and higher angular frequencies before mass shedding gets underway. Image reproduced by permission from Figure 19 of [305], copyright by AAS.
View Image Figure 8:
Isodensity contours for QE models of NS-NS binaries. In each case, the two NSs have masses M1 = 1.15M ⊙ (left) and M2 = 1.55M ⊙ (right), and the center-of-mass separation is as small as the QE numerical methods allow while able to find a convergent result. The models assume different EOS, resulting in different central concentrations and tidal deformations. Image reproduced by permission from Figures 9 – 12 of [305], copyright by AAS.
View Image Figure 9:
Approximate energy spectrum dEGW ∕df derived from QE sequences of equal-mass NS-NS binaries with isolated ADM masses MNS = 1.35M ⊙ and a Γ = 2 EOS, but varying compactnesses (denoted M ∕R here), originally described in [303]. The diagonal lines show the energy spectrum corresponding to a point-mass binary, as well as values with 90%, 75%, and 50% of the power at a given frequency. Asterisks indicate the onset of mass-shedding, beyond which QE results are no longer valid. Image reproduced by permission from Figure 2 of [98], copyright by APS.
View Image Figure 10:
Type of final remnant corresponding to different EOS models. The vertical axis shows the total mass of two NSs. The horizontal axis shows the EOSs together with the corresponding NS radii for MNS = 1.4M ⊙. Image reproduced by permission from Figure 3 of [134], copyright by APS.
View Image Figure 11:
Evolution of the total rest mass Mtot of the remnant disk (outside the BH horizon) normalized to the initial value for NS-NS mergers using a Γ = 2 polytropic EOS with differing mass ratios and total masses. The order of magnitude of the mass fraction in the disk can be read off the logarithmic mass scale on the vertical axis. The curves referring to different models have been shifted in time to coincide at tcoll. Image reproduced by permisison from Figure 5 of [240], copyright by IOP.
View Image Figure 12:
Fluid density isocontours and magnetic field distribution (in a plane slightly above the equator) immediately after first contact for a magnetized merger simulation. The cavities at both trailing edges are attributed to magnetic pressure inducing buoyancy. Image reproduced by permission from Figure 1 of [6], copyright by APS.
View Image Figure 13:
Evolution of the density in a NS-NS merger, with magnetic field lines superposed. The first panel shows the binary shortly after contact, while the second shows the short-lived HMNS remnant shortly before it collapses. In the latter two panels, a BH has already formed, and the disk around it winds up the magnetic field to a poloidal geometry of extremely large strength, ∼ 1015 G, with an half-opening angle of 30°, consistent with theoretical SGRB models. Image reproduced by permission from Figure 1 of [241], copyright by AAS.
View Image Figure 14:
Dimensionless GW strain Dh ∕m 0, where D is the distance to the source and m 0 the total mass of the binary, versus time for four different NS-NS merger calculations. The different merger types become apparent in the post-merger GW signal, clearly indicating how BH formation rapidly drives the GW signal down to negligible amplitudes. Image reproduced by permission from Figures 5 and 6 of [134], copyright by APS.
View Image Figure 15:
Effective strain at a distance of 100 Mpc shown as a function of the GW frequency (solid red curve) for the same four merger calculations depicted in Figure 14. Post-merger quasi-periodic oscillations are seen as broad peaks in the GW spectrum at frequencies f GW = 2 – 4 kHz. The blue curve shows the Taylor T4 result, which represents a particular method of deducing the signal from a 3PN evolution. The thick green dashed curve and orange dot-dashed curves depict the sensitivities of the second-generation Advanced LIGO and LCGT (Large Scale Cryogenic Gravitational Wave Telescope) detectors, respectively, while the maroon dashed curve shows the sensitivity of a hypothetical third-generation Einstein Telescope. Image reproduced by permission from Figures 5 and 6 of [134], copyright by APS.
View Image Figure 16:
Comparison between numerical waveforms, shown as a solid black line, and semi-analytic NNLO EOB waveforms, shown as a red dashed line (top panel). The top panels show the real parts of the EOB and numerical relativity waveforms, and the middle panels display the corresponding phase differences between waveforms generated with the two methods. There is excellent agreement between with the numerical waveform almost up to the time of the merger as shown by the match of the orbital frequencies (bottom panel). Image reproduced by permission from Figure 14 of [15], copyright by APS.
View Image Figure 17:
Evolution of a binary strange star merger performed using a CF SPH evolution. The “spiral arms” representing mass loss through the outer Lagrange points of the system are substantially narrower than those typically seen in CF calculations of NS-NS mergers with typical nuclear EOS models. Image reproduced by permission from Figure 4 of [35], copyright by APS.
View Image Figure 18:
Summary of potential outcomes from NS-NS mergers. Here, Mthr is the threshold mass (given the EOS) for collapse of a HMNS to a BH, and Q M is the binary mass ratio. ‘Small’, ‘massive’, and ‘heavy’ disks imply total disk masses Mdisk ≪ 0.01M ⊙, 0.01M ⊙ ≲ Mdisk ≲ 0.03M ⊙, and Mdisk ≳ 0.05M ⊙, respectively. ‘B-field’ and ‘J-transport’ indicate potential mechanisms for the HMNS to eventually lose its differential rotation support and collapse: magnetic damping and angular momentum transport outward into the disk. Spheroids are likely formed only for the APR and other stiff EOS models that can support remnants with relatively low rotational kinetic energies against collapse. Image reproduced by permission from [282], copyright by APS.