Simulations of neutron star mergers: Status and prospects

Transcription

1 Simulations of neutron star mergers: Status and prospects David Radice 1,2 1 Research Associate, Princeton University 2 Taplin Member, Institute for Advanced Study First multi-messenger observations of a neutron star merger and its implications for nuclear physics INT 18-72R March 13, 2018

2 NS merger: roadmap

3 Binary NS inspiral

4 Tidal effects in NS mergers 10 Part of the orbital energy goes into tidal deformation Q ij = 2 E ij Accelerated inspiral Imprinted on the gravitational waves Constrains dimensionless tidal parameter 2 = 2 M 5 R5 M 5

5 ~ Inspiral modeling 13 f min = 10Hz, = 50 M c /(1/4) 3/5 = 2.5 M sun, ~ = 0 M c /(1/4) 3/5 = 2.7 M sun, ~ = 0 M c /(1/4) 3/5 = 2.9 M sun, ~ = 0 M c /(1/4) 3/5 = 2.5 M sun, ~ = 1200 M c /(1/4) 3/5 = 2.7 M sun, ~ = 1200 M c /(1/4) 3/5 = 2.9 M sun, ~ = 1200 PRL 114, (2015) P H Y S I C A f max [Hz] From Kawaguchi, Kiuchi+ (2017) f min = 10Hz, f max = 1000Hz, = 50 From Bernuzzi+ (2015) FIG. 3 (color online). Phasing and amplitude compariso three representative models. Waves are aligned on a timewi the bottom panels indicate the crossing of the TEOB Resum Precision modeling over many orbits, see also Hinderer+ (2016), Dietrich+ (2017), Kiuchi+ (2017) Open issues: spin, last GW cycles before merger

6 Constraints from GW PRL 119, (2017) P H Y S I C A L R E V I E W L E T T E R S week ending 20 OCTOBER 2017 " FIG. 5. Probability density for the tidal deformability parameters of the high and low mass components inferred from the detected signals using the R post-newtonian 5 model. Contours enclosing 90% and 50% of the probability density are overlaid (dashed lines). The diagonal dashed line indicates the Λ 1 ¼ Λ 2 boundary. The Λ 1 and Λ 2 parameters characterize the size of the tidally induced mass deformations of M = 16 (M A + 12M B )M 4 (A) A 2 each 5 star and are proportional 13 to k 2 ðr=mþ (M 5 A. Constraints + M B are ) 5 +(A $ B) apple 800 shown for the high-spin scenario jχj 0.89 (left panel) and for the low-spin jχj 0.05 (right panel). As a comparison, we plot predictions for tidal deformability given by a set of representative equations of state [ ] (shaded filled regions), with labels following [161], all of which support stars of 2.01M. Under the assumption that both components From are LIGO/Virgo neutron stars, we collaboration, apply the function PRL ΛðmÞ 119, prescribed by that (2017) equation of state to the 90% most #

7 Prompt-BH formation

8 Simulation results ( ) M B DR, Perego, Zappa, ApJL 852:L29 (2018)

9 Simulation results ( ) M B DR, Perego, Zappa, ApJL 852:L29 (2018)

10 EOS constraints (I) excluded M [M ] excluded R [km] From Bauswein, Just+ (2017)

11 EOS constraints (II) M disk + Mej [M ] AT2017gfo tbh [ms] L BHBLf DD2 LS220 SFHo See also Bauswein ApJL 850:L34 DR, Perego, Zappa, ApJL 852:L29 (2018)

12 Hypermassive NSs z

13 GW-driven phase max [10 15 g cm 3 ] 4 2 LS LS LS LS DD DD DD DD LGW [10 55 erg s 1 ] SFHo SFHo SFHo SFHo u [ms] Bernuzzi, DR+, PRD 94: (2016)

14 Postmerger peak frequency f 2 0 f spiral f peak M sun 2.7 M sun h eff,x (20 Mpc) Type I adligo f peak [khz] M sun Type II Type III From Bauswein+ 2015ET f [khz] 2 From Bauswein R 1.6 [km] Post-merger signal has a characteristic peak frequency fpeak correlates with the NS radius Small statistical uncertainty, systematics not understood yet See also Takami+ 2014; Rezzolla & Takami 2016; Dietrich+ 2016; Bose+ 2017

15 Extreme-density physics 2.5 BHB DD2 Neutron stars in binaries have masses clustered around ~1.35 M M [M ] M =1.6 M Phase transition at high-density not constrained by the inspiral Can we probe the equation of state of nuclear matter at the highest densities? n max /n nuc 2 Yes, with the postmerger signal See also Bauswein+ 2011, 2013, 2015, Read+ 2013, Hotokezaka+ 2013, Takami+ 2014, Bernuzzi+ 2015, Clark+ 2014, 2016, Bose+ 2017, Chatziioannou 2017, DR, Bernuzzi, Del Pozzo+, ApJL 842:L10 (2017)

16 Gravitational waveform s] DD h+ (D = 100 Mpc) f [khz] f [khz] DD t t mrg [ms] BHB BHB DD2 DD2 DD t t mrg [ms] t t mrg [ms] log h e (f)(d = 100 Mpc) DR, Bernuzzi, Del Pozzo+, ApJL 842:L10 (2017)

17 End of GW-driven phase EGW/(M ) J/(M 2 ) Zappa, Bernuzzi, DR+, PRL in press (2018)

18 Viscous evolution to collapse

19 Angular momentum transport t visc = 1 t visc 15 ms See also: Shibata & Kiuchi 2017; Kiuchi, Kyotoku DR ApJL:838 L2 (2017)

20 Angular momentum transport t visc = 1 t visc 15 ms See also: Shibata & Kiuchi 2017; Kiuchi, Kyotoku DR ApJL:838 L2 (2017)

21 Angular momentum transport t visc = 1 t visc 15 ms Delayed collapse! See also: Shibata & Kiuchi 2017; Kiuchi, Kyotoku DR ApJL:838 L2 (2017)

22 Gravitational waves E GW [M c 2 ] `mix = 0 `mix = 5m `mix = 25 m `mix = 50 m t t mrg [ms] See also: Shibata & Kiuchi 2017; Kiuchi, Kyotoku DR ApJL:838 L2 (2017)

23 Gravitational waves E GW [M c 2 ] `mix = 0 `mix = 5m `mix = 25 m How large is the turbulent viscosity? `mix = 50 m How do hypermassive NS evolve over many viscous timescales? Can we distinguish long- and short-lived hypermassive NSs? t t mrg [ms] See also: Shibata & Kiuchi 2017; Kiuchi, Kyotoku DR ApJL:838 L2 (2017)

24 Viscous evolution to equilibrium

25 Long-lived remnants DD Mb [M ] M [M ] ( )M M J [Gc 1 M 2 ] DR, Perego, Bernuzzi, Zhang, in prep.

26 Viscous evolution Mb [M ] Disk ejecta Remnant ejecta RNS DD2 ( ) M M J [Gc 1 M 2 ] See also Fujibayashi, Kiuchi+ (2017) DR, Perego, Bernuzzi, Zhang, in prep.

27 Excess gravitational mass BHB DD2 LS220 SFHo M [M ] DD2 ( )M M M b /M RNS DR, Perego, Bernuzzi, Zhang, in prep.

28 Stable or unstable? From Kaplan, Ott, O Connor+ (2014)

29 Stable or unstable? ' 0.1 M!!! From Kaplan, Ott, O Connor+ (2014)

30 The remnant of GW Fig. 1. The strength of the red and blue KN signatures of a BNS merger depends on the compact remnant which forms immediately after the merger; the latter in turn depends on the total mass of the original binary or its remnant, Mtot, relative to the maximum NS mass, Mmax. A massive binary (Mtot & Mmax ) results in a prompt collapse to a BH; in such cases, the polar shock-heated ejecta is negligible and the accretion disk outflows are weakly irradiated by neutrinos, resulting in a primarily red KN powered by the tidal ejecta (left panel). By contrast, a very low mass binary Mtot. 1.2Mmax creates a long-lived SMNS, which imparts its large rotational energy & 1052 erg to the surrounding ejecta, imparting relativistic expansion speeds to the KN ejecta or producing an abnormally powerful GRB jet (right panel). In the intermediate case, 1.2Mmax. Mtot Mmax a HMNS or short-lived SMNS forms, which produces both blue and red KN ejecta expanding at mildly relativistic velocities, consistent with observations of GW From Margalit & Metzger 2017 Long-lived unlikely the maximum ral (Hinderer et SMNS al. 2010; Damour & Nagar 2010; > Damour limit 2011),on where the proportionality factor NS k 1.3mass 1.6 et al. 2012; Favata 2014; Read et al. 2013; Del Pozzo et al. 2013; Agathos et al. 2015; Lackey & Wade 2015; Chatziioannou et al. 2015) and for quasi-periodic oscillations of the post-merger remnant (e.g. Bauswein & Janka 2012; Bauswein et al. 2012; Clark et al. 2014; Bauswein See etalso Rezzolla+, & Stergioulas 2015; Bauswein al. 2016). Searches on timescales of tens of ms to. 500 s post-merger revealed is greater for smaller values of the NS compactness, Cmax = (GMmax /c2 R1.6 ), where R1.6 is the radius of a 1.6M NS (e.g. Bauswein et al. 2013). For slightly less massive binaries with Mtot. Mth, the merger instead produces a hyper-massive neutron star (HMNS), which Shibata+, (2017) is supportedruiz+ from collapse by di erential rotation (and, potentially, by thermal support). For lower values of

31 The origin of the elements R-Process Are neutron star mergers the site of the r-process?

32 Ejection mechanisms [g/cm 3 ] x [km] See also Metzger+2008; Wanajo+2014; Fernandez+2014; Metzger+2014; Perego+2014; Martin+2015; Sekiguchi+2015,2016; Foucart+2016; Siegel Ye DR, Galeazzi+ MRAS 460:3255 (2016)

33 Dynamic ejecta: role of neutrinos Mass fraction SFHo: ( ) M ; cooling only sin 2 ( ) 10 2 Electron fraction Polar angle Mass fraction Perego, DR, Bernuzzi, ApJL:850 L37

34 Dynamic ejecta: role of neutrinos Mass fraction SFHo: ( ) M ; cooling and heating sin 2 ( ) 10 2 Electron fraction Polar angle Mass fraction Perego, DR, Bernuzzi, ApJL:850 L37

35 Conclusions Numerical relativity is essential in the age of multimessenger astronomy Do we really understand the outcome of NS mergers? Neutrinos play a crucial role for nucleosynthesis and EM counterparts

36 Thank you!