Neutron-star mergers:

Transcription

1 Neutron-star mergers: Predictions by numerical relativity Masaru Shibata Center for Gravitational Physics, Yukawa Institute for Theoretical Physics, Kyoto University

2 Many people are exploring NS binaries in numerical relativity Shibata & Uryu (1999), Taniguchi Sekiguchi, Kiuchi, Kyutoku, Hotokezaka, Kawaguchi Rezzolla, Baiotti, Giacomazzo, Kastaun, Alic, Ciolfi,.. Shapiro, Liu, Etienne, Pachalidis,.. Bernuzzi, Dietrich, Nagar, Bruegmann, Gold,.. Lehner, Palenzuele, Liebling, Nielsen, Anderson,.. Bauswein, Stergioulas, Janka,.. Foucart, Duez, O Connor, Ott, Haas, Scheel, Kidder, Pfeiffer,.. Loeffler and his colleagues & many others Solid progress on understanding NS-NS/ NS-BH binary by numerical relativity

3 Outline 1. Brief introduction 2. Standard scenarios of NS-mergers 3. Gravitational waves and equations of state 4. Mass ejection 5. R-process nucleosynthesis 6. Summary

4 Many aspects of NS-NS/BH-NS One of the most promising sources of gravitational waves for LIGO/VIRGO/KAGRA Laboratory for high-density nuclear matter Promising progenitors of short-hard GRBs Possible site of r-process nucleosynthesis Laboratory for testing GR Gravitational-wave obs. + EM signals obs. + Numerical relativity will contribute to exploring all these issues?

5 2 Standard scenarios of NS- NS/BH- NS merger 2A Binary neutron stars

6 Boundary conditions from latest observations ² Binary pulsar observations suggest Ø Mass of NS in compact NS-NS is likely to be in a narrow range, m 1.35±0.15 M sun Ø Spin of NS is likely to be not very high, P rot > ~10 ms or a < ~0.05 Ø NS radius (EOS) is still uncertain, but maximum mass of NS > 2 M sun (Demorest 2010; Antoniadis 2013) à EOS of NS has to be stiff

7 Current understanding of NS-NS Merger of M sun NS with four EOSs APR4: R=11.1km ALF2: R=12.4km All EOSs satisfy M max > 2M sun H4: R=13.6km MS1: R=14.5km

8 Merger of M sun NS with four EOSs By hotokezaka Massive APR4: R=11.1km neutron stars are ALF2: remnants R=12.4km Log(ρ g/cc) Log(ρ g/cc) irrespective of EOS for canonical mass H4: R=13.6km MS1: R=14.5km

9 Possible outcomes of NS-NS mergers BH NS M > ~2.8M sun Likely for M < ~2.8M sun I.e., irrespective of EOS, threshold mass >~2.8M sun

10 2B Black hole-neutron star binaries

11 Two possibilities: Tidal disruption or not For tidal disruption, ( Self gravity of NS) < ( BH tidal force) < C M R $ BH NS (C >1) 1 C M BH 2 R NS r 3 & % r ISCO M NS ' ) ( 3 $ & % M NS M BH ' ) ( 2 $ & % R NS M NS ' ) ( 3 For tidal disruption v Large NS Radius or v Small BH mass or v High corotation spin is necessary ISCO B H B H NS NS

12 BH-NS with aligned BH spin M BH =6.75M sun a=0.75 M NS =1.35M sun R=11.1 km M BH =4.05M sun a=0 M NS =1.35M sun R=11.0 km

13 BH-NS with aligned BH spin M BH =6.75M sun a=0.75 M NS =1.35M sun R=11.1 km M BH =4.05M sun a=0 M NS =1.35M sun R=11.0 km Kyutoku et al. 2011, 2015

14 For tidal disruption of plausible BH-NS with M NS =1.35M sun, R NS ~ 12 km, & M BH > 6 M sun High BH spin is necessary > ~ 0.75 Foucart et al. (2013, 2014); Kyutoku et al. (2015) If high-mass BH, solar mass, is standard, tidal disruption is not very likely: Only quite high-spin BH can tidally disrupt NS.

15 3 Gravita:onal waves & equa:ons of state

16 Early Inspiral ( r >> R ) orb NS 3A NS-NS case Late inspiral ( ) r orb 5R NS Merger => Massive NS M~ M sun Black hole/mns + torus à GRB? Point mass + adiabatic phase Post-Newtonian Tidally deformed phase Late inspiral Dynamical & GR phase Post merger Massive NS/BH f ~ 6.5 khz h f <~1000 Hz f ~ 2 4 khz t/m

17 Gravitational waveform from NS-NS: hybrid waveform ( solar mass) 0.2 EOB (Damour, Nagar..) Numerical-relativity h c 2 D / G m Align the phase here Black: BH (by SXS) Red: SFHo (11.9 km) Green: DD2 (13.2 km) Blue: TM1 (14.5 km) Late inspiral Post merger t ret (s) Hotokezaka et al (see also efforts by Bernuzzi, 2011 )

18 2f 1/2 h(f) (D eff =100 Mpc) Spectrum The difference is determined primarily by tidal deformability Λ APR4 SFHo DD2 TMA TM1 BBH aligo Late inspiral Advanced LIGO: Zero detuned high-p Hotokezaka et al. PRD f (khz)

19 h 1 -h 2 eff =200Mpc (SNR ~ 17) APR4-SFHo Hz δλ = 400 δr =1 km TMA-TM1 DD2-TMA SFho-DD2 APR4-DD2 DD2-TM1 SFHo-TM See also Damour et al Hotokezaka et al 2016 δλ APR4-TMA SFHo-TM1 # % h 1 h 2 = 4 % $ 0 f max APR4-TM1 h 1 ( f ) h 2 f S n 2σ level 1σ level ( ) 2 ( f ) Fitting formula df & ( ( ' 1/2 Lindblom, Owen, Brown 2008

20 2f 1/2 h(f) (D eff =100 Mpc) Spectrum SNR =0.5 Information of ~10 15 g/cm 3 is reflected APR4 SFHo DD2 TMA TM1 BBH aligo Post merger Advanced LIGO: Zero-detuned high-p Hotokezaka et al. PRD f (khz)

21 Clear correlation between peak and radius Peak frequency NS radius Bauswein & Janka could be constrained with ~ 1km error Ours At one lucky event f GM R 3 Radius of 1.6 solar-mass NS

22 Caveat Merger waveforms have been computed in quite simple setting (essentially, pure hydro) Post-merger phase would be in reality determined by complicated physics Turbulence will be excited by MHD instability (e.g., Kelivin-Helmholtz instability, MRI; Kiuchi s talk) à Magnetic fields would be amplified to ~ G à Turbulent viscosity could change velocity profile, modifying waveform?? (but no detailed simulation)

23 Ψ 4 22 D m Pure hydro vs viscous hydro (α vis ~0.02) Very preliminary Need better-resolution study t ret (ms)

24 3B BH-NS: Signal of tidal disruption Tidal disruption=stiff EOS sudden shutdown Weak tidal disruption=soft EOS BH ringdown Green=Tayloy T4

25 D eff»h é»êm BH-NS Fourier spectrum For all, a=0.5, M sun p.3g3.0 p.4g3.0 p.5g3.0 p.7g3.0 p.9g3.0 c BH =0.5, Q=4, M NS =1.35 f ~ 1 2π Larger NS radius = Stiff EOS PhenomC EOB Higher mass GM NS R NS 3 ~ 1 2 khz Lackey et al. 14; Pannarale et al. 14 Mf

26 h opt (f) f 1/2 (100Mpc) Cutoff frequency is high, and SNR is small 1e-22 1e-23 1e-24 Note: D=100Mpc High Freq (3 tunings) Zero-detuned low-p Zero-detuned high-p R12i0 R13i0 R14i0 a= M sun 12km 13km 14km h 1 h 2 f ~ 1 2π f GW Foucart et al GM NS R NS 3 ~ 1 2 khz <1 even for D eff =100 Mpc 2000

27 Pressure at ρ=5*10 14 g/cc Data analysis using numerical-relativity data: ET 35.0 level (or nearby event) is necessary loghp 1 2 D D=100 Mpc L 1ê5 =5.0 L 1ê5 =4.0 L 1ê5 =3.0 c BH =0.75, Q=5, M NS =1.35M ü Adiabatic G index Lackey et al. (2014) Λ=1024 Λ=525

28 4 Mass ejec:on Metzger & Berger 2012 Jet ISM Shock (Afterglow) Optical (hours days) Radio (weeks years) Ejecta ISM Shock Radio (years) GRB (t ~ s) θ obs Kilonova Optical (t ~ 1 day) θ j BH Merger Ejecta Tidal Tail & Disk Wind v ~ c In binary merger, neutron-rich matter is ejected and it could shine (Piran s talk) à NR should clarify the ejecta properties

29 For radioactive (macronova) scenario (Li-Paczynski 98) Ejecta mass Velocity Opacity(composition) " L max ~ M % ergs/s $ # 0.01M ' & " M % at t ~ 5 days $ # 0.01M ' & 1/2 " $ # v 0.2c 1/2 % ' & " $ # 1/2 v 0.2c % ' & 1/2 " κ % $ ' # 10 cm 2 / g & " κ % $ ' # 10 cm 2 / g & 1/2 1/2 " f r-proc $ # ergs/s M = 15.0 mag m= Mpc % ' & Bright in near-infrared for 1 week after the merger

30 Observed magnitude Expected light 200Mpc Optical/Near infrared (i band) M=0.01 sola mass i band 200 Mpc 1m 1 min exposure 4m 10 min exposure by Subaru HSC (SNR=5) 8m Subaru HSC (hyper-supreme -camera) 1.75 deg 2 field of view Days after the merger Tanaka & Hotoke 2013; see also Barnes & Kasen 2013, 2016

31 Neutrino-radiation hydrodynamics simulation SFHo (R~11.9 km): M sun Rest-mass density Orbital plane Tidal torque Neutrino luminosity Neutrino wind: weak ν e ν e ν others x-z plane Shock hea:ng Sekiguchi et al

32 Ejecta mass depends on EOS : NS-NS case Soft EOS à strong gravity à SHOCK à high-mass ejection SLy APR4 Steiner ALF2 Mass ratio Tidal effect is major H4 MS1 Total mass = 2.7 solar mass Error bar for 1 < Q < 1.25 Radius of 1.35 solar mass NS Hotokezaka+ PRD 13 (See also Bauswein+ 13; Bernuzzi + 15)

33 BH-NS: disruption à ejection Tidal torque by BH induces mass ejection y y z x ejecta z x Anisotropic ejection along the orbital plane: Note: Disk wind is not taken into account. x Kyutoku et al x

34 M *ejecta (solar mass) BH-NS with NS mass 1.35M sun M NS =1.35 solar mass Soft EOS < ~0.01 M sun Data: Kyutoku et al Stiff EOS results in high mass > 0.01 M sun M BH =9.45, a=0.75 M BH =9.45, a=0.50 M BH =6.75, a= Radius of R1.35 (km) solar mass NS High BH spin is important for mass ejection

35 Mass: Dynamical ejecta properties in numerical relativity Ø NS-NS: ~ M sun depending on each mass & EOS: Soft EOS is favorable (Hotoke+ 13, Sekiguchi+ 15,16, Radice+ 16, Lehner+ 15,16) Ø BH-NS: M sun depending on each mass, BH spin, & EOS: Stiff EOS is favorable; high BH spin is also the key (Foucart+13,14, Kyutoku+15) Ø Typical velocity: c; max ~ 0.8 c Detectable for macronova from NS-NS by 8-m class telescopes

36 Other effects? Viscous hydrodynamics (α v =0.02): preliminary How large the effective viscosity? à High-resolution MHD simulation is needed See Fernandez & Metzger, Perego+, Just+ for related works

37 5 R-process nucleosynthesis Pagel (1997) N=50 1 st peak N=82 2 nd peak N=126 3 rd peak Solar abundance of elements Key = electron fraction: Y e =[p]/([n]+[p]) 37 Courtesy Sekiguchi

38 Importance of Y e =[p]/([p]+[n]) in r-process Y e ~ is critical threshold Y e < 0.22 : strong r-process nuclei with A > 130 Y e > 0.22 : weak r-process nuclei with A < st peak 2 nd peak 3 rd peak Korobkin et al Strong Ye dependence Courtesy Sekiguchi

39 High Neutrino-radiation temperature hydrodynamics γγ e + e +, n + simulation e + p +ν e SFHo (R~11.9 km): M sun Neutrino irradiation n +ν p + e Y e Electron fraction (x-y) Neutrino luminosity ν e ν e ν others Electron fraction (x-z) Sekiguchi et al. (2015) 39

40 Dynamical evolution of neutron richness Tidal torque Shock heating Mass ejec:on Neutron-rich ejecta Y e <~ 0.1 Shock heating Eject Y e >~ 0.1 Neutrino irradiation Y e ñ Y e

41 Fraction of mass Electron fraction profile: Broad solar case t - t M-6 [ms] Sekiguchi et al PRD SFHo DD2 TM Ye Ø Average depends on EOS but typically Ø Broad distribution irrespective of EOS Ø Similar results by Radice+16, Lehner+15,16

42 Consistent with solar abundance pattern Wanajo Y e <~ 0.15 Y e ~ 0.4 Y e ~ 0.2 Sekiguchi et al. 2015

43 Results for a variety of mass ratios Mass fraction Mass fraction asymmetric asymmetric SFHo DD Sekiguchi Electron fraction (Ye)

44 BH-NS merger (DD2 EOS: Y e ) M BH =5.5M sun, M NS =1.35M sun, a BH =0.75 Kyutoku, Sekiguchi

45 Still many issues remain MHD/Viscous wind from torus? How large is the effective viscosity? à High-resolution & better physics is necessary Issues in numerical relativity in the next ~ a few years Neutrino wind alone is likely to be minor effect (Sekiguchi+ 2015)

46 5 Summary Many NR simulations are ongoing for NS-NS/BH à Standard scenarios have been established. Detecting late-inspiral gravitational waves from NS-NS will constrain EOS for D eff <~ 200Mpc. Ejecta mass: ~ M sun for NS-NS, and up to 0.1 M sun for BH-NS: depends on NS- EOS: à L ~ erg/s at 1 10 days for radioactivepower nova (for κ ~ 10 cm 2 /g). à EM counterparts; luminosity depends on NS EOS & BH spin: Observation will reveal them. NS-NS merger could be the site for r-process nucleosynthesis.

47 Origin of elements Neutron star merger

48 Galactic compact NS-NS observed PSR P(day) e M(M sun ) M 1 M 2 T GW 1. B B B C J J J In globular cluster E.g., Orbital period Eccentricity Mass à Galactic merger rate ~1/10 5±1 yrs (e.g. Kalogera et al., 2007, Abadie et al. 2010) à Merger rate ~1-100/yr/(300Mpc) 3 *10 8 yrs Merger time

49 Demorest Nature 2010 The most crucial uncertainty is EOS Suggestion by X-ray obs. 2.5 Mass (solar) J Many simula+ons with many EOSs 1.0 are needed for systema+c study Strong constraint: But not strong enough Radius (km) 10 km 12 km 14 km