Hirschegg, 17.01.2017

Compact binary mergers:
what else apart from gravitational waves?

Stephan Rosswog
collaborators for this project:
A.Goobar, U. Feindt, O. Korobkin, G. Martinez-Pinedo, J. Sollerman, M.-R. Wu⇒ “Detectability of compact binary merger macronovae”, arXiv:1611.09822

Outline• I. Introduction• Glueing fields together• GWs: EM needed• Types of expected EM• II. Overview mass loss and nucleosynthesis• Mechanisms• Dynamic ejection + nucleo• ν-winds + nucleo• Other types of “winds”• Rate constraints from nucleo• a) mass in Galaxy• b) constraints nucleo + sGRB + GW

• III. Macronovae• Our model• Explored parameter space• Impact of mass formula• Detection prospects• “Tanvir+ event”• IV. Summary

I.1 Neutron star mergers: “glueing together separate fields”

I.1 Neutron star mergers: “glueing together separate fields”
nucleosynthesis
Stellar (binary) evolution

Chemical enrichmentof the Cosmos

Gravitational wave detection

Radioactively poweredtransients (“macronovae”)(short) Gamma-Ray Bursts

I.2 Gravitational wave astronomy
The first gravitational wave detection

GW150914⇒ “physics from gravitational wave signal”⇒ “astronomy/astrophysics from the electromagnetic (EM) transient”

I.2 Gravitational wave astronomy
The first gravitational wave detection

GW150914⇒ “physics from gravitational wave signal”⇒ “astronomy/astrophysics from the electromagnetic (EM) transient”

I.2 Gravitational wave astronomy
The first gravitational wave detection

GW150914⇒ “physics from gravitational wave signal”⇒ “astronomy/astrophysics from the electromagnetic (EM) transient”
location in the sky essentially unknown,~ 600 deg2 error region==> electromagnetic (EM)-transientneeded for sky location!

I.3 Which EM-transients can we expect toaccompany a gravitational wave chirp?
for binary black holes:• IF the merger occurs in gaseous environment:
post-merger black hole has mass reduced by△m, has received “kick”
S.R. 2017“BH kicked through disk”

parameters of this illustrative example:• m1=29 M⊙ & m2= 36 M⊙• △m= 3 M⊙
• mdisk= 0.01 M⊙• v_kick= 500 km/s

I.3 Which EM-transients can we expect toaccompany a gravitational wave chirp?
for binary black holes:• IF the merger occurs in gaseous environment:
post-merger black hole has mass reduced by△m, has received “kick”
S.R. 2017“BH kicked through disk”

parameters of this illustrative example:• m1=29 M⊙ & m2= 36 M⊙• △m= 3 M⊙
• mdisk= 0.01 M⊙• v_kick= 500 km/s

for compact binaries involving neutron stars:(1) Short Gamma-Ray Bursts- collimated into ~8o ⇒ detect ~ 1 out of 70 bursts
time scales ~ second(2) Radioactively powered transients (“macronovae”)- de-compressed neutron star matter forms heavy nuclei- “cloud of radioactive, expanding matter”- isotropic emission
time scales ~ days

(3) Radio flares- from dissipation of kinetic energy in ambient medium
time scales ~ months

II. Mass loss and nucleosynthesis

II.1 Mass loss channels for NSNS and NSBH:• dynamic ejecta: by hydrodynamics and gravitational torques
nsns: few 10-3 to few 10-2 M⊙ (mass ratio!)
nsbh: up to ~0.2 M⊙ (mass ratio, bh spin χ)II. Mass loss and nucleosynthesis

II.1 Mass loss channels for NSNS and NSBH:“dynamic”• dynamic ejecta: by hydrodynamics
and gravitational torques
nsns: few 10-3 to few 10-2 M⊙ (mass ratio!)
nsbh: up to ~0.2 M⊙ (mass ratio, bh spin χ)II. Mass loss and nucleosynthesis

II.1 Mass loss channels for NSNS and NSBH:“dynamic”• dynamic ejecta: by hydrodynamics
and gravitational torques
nsns: few 10-3 to few 10-2 M⊙ (mass ratio!)
nsbh: up to ~0.2 M⊙ (mass ratio, bh spin χ)• ν-driven winds: massive neutron star present: ~0.01 M⊙bh + torus: ≪ ~0.01 M⊙

L⌫ ⇠ 1053 erg/shE⌫ i ⇠ 15 MeV(from Perego et al. 2014)

II. Mass loss and nucleosynthesis

II.1 Mass loss channels for NSNS and NSBH:“dynamic”• dynamic ejecta: by hydrodynamics
and gravitational torques
nsns: few 10-3 to few 10-2 M⊙ (mass ratio!)
nsbh: up to ~0.2 M⊙ (mass ratio, bh spin χ)• ν-driven winds: massive neutron star present: ~0.01 M⊙
bh + torus: ≪ ~0.01 M⊙

L⌫ ⇠ 1053 erg/shE⌫ i ⇠ 15 MeV(from Perego et al. 2014)• magnetically-driven winds:
either from disks and/or massive “neutron star”(Siegel et al. 2014)

II. Mass loss and nucleosynthesis

II.1 Mass loss channels for NSNS and NSBH:“dynamic”• dynamic ejecta: by hydrodynamics
and gravitational torques
nsns: few 10-3 to few 10-2 M⊙ (mass ratio!)
nsbh: up to ~0.2 M⊙ (mass ratio, bh spin χ)L⌫ ⇠ 1053 erg/s• ν-driven winds: massive neutron star present: ~0.01 M⊙
bh + torus: ≪ ~0.01 M⊙

hE⌫ i ⇠ 15 MeV(from Perego et al. 2014)• magnetically-driven winds:
either from disks and/or massive “neutron star”• torus unbinding (viscous diss., nuclear recomb.):~ 20% of initial torus mass

(Siegel et al. 2014)(Fernandez & Metzger 13, Just+ 15)• initial torus mass can be large!
⇒ could plausibly unbind ~0.02 … 0.1 M⊙(from Giacomazzo et al. 2013)

II. Mass loss and nucleosynthesis

II.1 Mass loss channels for NSNS and NSBH:“dynamic”• dynamic ejecta: by hydrodynamics
and gravitational torques
nsns: few 10-3 to few 10-2 M⊙ (mass ratio!)
nsbh: up to ~0.2 M⊙ (mass ratio, bh spin χ)L⌫ ⇠ 1053 erg/s• ν-driven winds: massive neutron star present: ~0.01 M⊙
bh + torus: ≪ ~0.01 M⊙“winds”: - stay longer closer the central neutrino source
winds:fraction driven towards equilibrium• magnetically-driven⇒ electron& Woosley
either(Qianfromdisks96)and/or massive “neutron star”• torus unbinding (viscous diss., nuclear recomb.):~ 20% of initial torus mass
- different(Fernandez &velocities
Metzger 13, Just+ 15)

hE⌫ i ⇠ 15 MeV(from Perego et al. 2014)~ 0.3(Siegel et al. 2014)• initial torus mass⇒canbe large!
different
nucleosynthesis⇒ could plausibly unbind ~0.02 … 0.1 M⊙⇒ different EM-transients(from Giacomazzo et al. 2013)

II.2 Nucleosynthesis for dynamic ejecta
Example:Modelled Physics• 3D Lagrangianhydrodynamics;
SPH with 6 million
particles (e.g. S.R. 2016)• Newtonian self-gravity• GW-backreaction force• Nuclear equation ofstate (Shen et al. 1998)• opacity-depent neutrino
cooling and weakinteractions/change ofcomposition(S.R. & Liebendoerfer 2003)neutron star binary (1.4 and 1.5 Msol , no stellar spins)

II.2 Nucleosynthesis for dynamic ejecta
Example:Modelled Physics• 3D Lagrangianhydrodynamics;
SPH with 6 million
particles (e.g. S.R. 2016)• Newtonian self-gravity• GW-backreaction force• Nuclear equation ofstate (Shen et al. 1998)• opacity-depent neutrino
cooling and weakinteractions/change ofcomposition(S.R. & Liebendoerfer 2003)

neutron star binary (1.4 and 1.5 Msol , no stellar spins)

II.2 Nucleosynthesis for dynamic ejecta
Example:neutron star binary (1.4 and 1.5 Msol , no stellar spins)

Modelled Physics• 3D Lagrangianhydrodynamics;
SPH with 6 million
particles (e.g. S.R. 2016)• Newtonian self-gravity• GW-backreaction force• Nuclear equation ofstate (Shen et al. 1998)• opacity-depent neutrino
cooling and weakinteractions/change ofcomposition(S.R. & Liebendoerfer 2003)⇒ dynamic ejecta are extremely neutron-rich
Ye ~ 0.04, i.e. only 4 protons per 100 nucleons;
ejected mass: ~ 0.01 Msol
average ejecta velocity <v∞>= 0.12 c

r-process calculations for dynamic ejecta(Korobkin, S.R, Arcones, Winteler, MNRAS 426, 1940 (2012) )T9⇢[g/cm3 ]Winnet network(Winteler 2012)

5 831 isotopes

r-process calculations for dynamic ejecta(Korobkin, S.R, Arcones, Winteler, MNRAS 426, 1940 (2012) )T9⇢[g/cm3 ]Winnet network(Winteler 2012)

5 831 isotopes

Resulting abundances•“Punchlines”
1. all binaries reach the “platinum peak” (A=195) without any tuning,
not much below A ~ 130
2. the abundance pattern is extremely robust with respect to astrophysics:
A. all mass ratios (even NSBH-systems) deliver the same pattern
B. Ye ≲ 0.25 needed for 3rd r-process peak (“platinum peak”)
3. sensitivity to nuclear physics ⇒ e.g. talks Eichler, Giuliani, Surman, Wu
4. tidal dynamic ejecta component delivers very low-Ye material (Ye~ 0.04),
likely complemented by higher Ye components (ν-driven wind, disk
evaporation etc.)

II.3 Nucleosynthesis neutrino-driven winds(Perego, et al. 2014, Martin, et al. 2015)

Questions:• neutrino properties?
• impact ν-oscillations ⇒ talk Gail McLaughlin• neutrino-driven winds:
mass loss?
geometry?
nucleosynthesis?
resulting radioactivelypowered transients?

L⌫ ⇠ 1053 erg/shE⌫ i ⇠ 15 MeV(from Perego et al. 2014)⇒ 0.2 ≲ Ye ≲ 0.4
ρv

Ye entropy

II.3 Nucleosynthesis neutrino-driven winds(Perego, et al. 2014, Martin, et al. 2015)

Questions:• neutrino properties?
• impact ν-oscillations ⇒ talk Gail McLaughlin• neutrino-driven winds:
mass loss?
geometry?
nucleosynthesis?
resulting radioactively
powered transients?

L⌫ ⇠ 1053 erg/shE⌫ i ⇠ 15 MeV(from Perego et al. 2014)⇒ 0.2 ≲ Ye ≲ 0.4
ρv

Ye entropy

nucleosynthesis in neutrino-driven winds(Martin, Perego, Arcones, Thielemann, Korobkin, S.R. (2015))

Winnet network(Winteler 2012, 5831 isotopes)

complementary r-process in the range 80 ≲ A ≲ 130

nucleosynthesis in neutrino-driven winds(Martin, Perego, Arcones, Thielemann, Korobkin, S.R. (2015))

Winnet network(Winteler 2012, 5831 isotopes)

complementary r-process in the range 80 ≲ A ≲ 130

II.4 Rate constraints
In the Milky Way:
all r-process: 19 000M

r-process A > 80: 14 000 M

r-process A > 130: 2530 MII.4 Rate constraints
In the Milky Way:
all r-process: 19 000M

r-process A > 80: 14 000 M

r-process A > 130: 2530 M

What does this imply for the production rate?
Mr = mej ⇥ R ⇥ ⌧Gal

total r-mass

typical ejecta massMr
• degeneracy mej ⇥ R =⌧Gal

average rate

age of Galaxy

all r-process:r-process A > 80:
r-process A > 130:

1.9 ⇥ 106M

yr
6M
1.4 ⇥ 10yr
7M
2.5 ⇥ 10yr

⇒ “rare high-mass (e.g. NSNS merger) or
frequent low-mass events (like supernovae) ?”Rate constraints
10log(required ejecta mass [Msol])

2100

100010 000-1-3[yr Gpc ]aLIGO nsbh rate compendium
LIGO nsbh O2
LIGO nsbh O3

LIGO nsbh O1

ALL r-processr-process A > 80
r-process A> 130

0nsbh Foucart+ 14
nsns Bauswein+ 13
nsbh Kyutoku+ 13

nsns Rosswog 13-2
nsns Hotokezaka+ 13

LIGO nsns O3-4

LIGO nsns O2
LIGO nsns O1
aLIGO nsns rate compendium-6-7-6-4-5-1-3-2-1log10 ( R [yr gal. ] )from SR++ (2016).

Rate constraints
10log(required ejecta mass [Msol])

2

100

1000

10 000-1-3[yr Gpc ]aLIGO nsbh rate compendium
LIGO nsbh O2
LIGO nsbh O3

LIGO nsbh O1

ALL r-processr-process A > 80
r-process A> 130

0pop. synth. nsns
Chruslinska+16
nsbh Foucart+ 14
nsns Bauswein+ 13
nsbh Kyutoku+ 13

nsns Rosswog 13-2nsns Hotokezaka+ 13

LIGO nsns O3-4

LIGO nsns O2
LIGO nsns O1
aLIGO nsns rate compendium-6-7-6-4-5-1-3-2-1

log10 ( R [yr gal. ] )from SR++ (2016).

Rate constraints
10log(required ejecta mass [Msol])

2

100

1000

10 000-1-3[yr Gpc ]aLIGO nsbh rate compendium
LIGO nsbh O2
LIGO nsbh O3

LIGO nsbh O1

ALL r-processr-process A > 80
r-process A> 1300
pop. synth. nsns
Chruslinska+16
nsbh Foucart+ 14
nsns Bauswein+ 13
nsbh Kyutoku+ 13

nsns Rosswog 13-2
sGRB rate
Petrillo + 16
nsns Hotokezaka+ 13

LIGO nsns O3-4

LIGO nsns O2
LIGO nsns O1
aLIGO nsns rate compendium-6-7-6-4-5-1-3-2-1

log10 ( R [yr gal. ] )from SR++ (2016).

III. Macronovae (MN)
use here two simple models, i.e. no full radiative transfer (⇒ talk Ryan Wollager):macronova model I:(Grossman++ 14)• spher. symmetric, homologously expanding ρ-profile

macronova model ΙI:
like model I, BUT:

• “radiating volume”: above “diffusion surface” (τdyn = τdiff)• nuclear heating rate direct from nuclear network• “Finite Range Droplet Model” (FRDM) for nuclear
mass (Moeller++ 95)• either FRDM mass model or
Duflo-Zuker (DZ31, Duflo-Zuker 95)• constant thermalization used f= 0.5• time-dependent thermalization
efficiencies for each species
according to the model developedby Barnes et al. 16• BB-spectrum with effective temperature of thephotosphere (τph=2/3)• opacities:• dynamic ejecta: κ= 10 cm2/g (Kasen+13, Barnes+13)• “wind” ejecta: κ= 1 cm2/g (Grossman+14)

III.2 Explored parameter space
Dynamic ejecta nsns- & nsbh mergers

lanthanide + actinides > 20% of mass

mej up to ~0.04 M⊙
q=0.78compare withJ0453+1559: q= 0.75

mej up to ~0.16 M⊙

ns13ns13

from SR++ (2016)

(parametrized) winds“ν-wind inspired”“disk evaporation inspired”
explore “unknown”
parameter spaceparameters of wind model:
0.01 … 0.2 M⊙• mass:0.01 … 0.50 c• velocity:• electron fraction: 0.25 … 0.35
15 kB/bar.
• entropy:
Log(kinetic energy)
from SR++ (2016)

(broad-brush) nucleosynthesis:“winds”“dynamic ejecta”“weak r-process, A ≲ 130”

“strong r-process, A ≳ 130”

III.3 How big is the impact of nuclear physics
on potential observability?
from SR++ (2016)

comparison nuclear mass formulae:
Finite Range Droplet Model (FRDM)
vs.
Duflo Zuker 31 param. (DZ31)(at example ns12 + ns14 at t= 100 days)⇒ substantially larger amounts of trans-lead elements for Duflo Zuker 31(as first discussed in Barnes et al. (2016) ⇒ talk Jennifer Barnes)

comparison networks:comparison of mass formulae:

heating rate for same mass formulaheating rate for FRDM vs DZ31
(for same network (Mendoza-Temis++ 2015))

from SR++ (2016)⇒ at relevant times (days): ~ order of magnitude larger heating
rate due to α-decays of trans-lead elements for Duflo Zuker 31thermalization efficiency FRDM

thermalization efficiency DZ31

from SR++ (2016)⇒ at ~ 10 days ~factor of two larger for Duflo Zuker 31

broadband filters

impact on optical/near IR light-curves:
our brightest nsns-example: 1.4 M⊙ + 1.8 M⊙

grizyMacronova model 1:• original Grossman++(2014)• FRDM mass model• efficiency = 0.5 = const.

LSST~380 - 580 nm~520 - 700 nm~650 - 850 nm~780 - 950 nm~950 - 1050 nm

2MASS
J ~1.1 - 1.4 micron
H ~1.5 - 1.8 micron
K ~2.0 - 2.4 micron• “faint and fast in blue”
Macronova model 2:
like model 1, but• DZ31• time-dep. efficiencies• “bright in NIR forweeks”• “up to a magnitude
brighter for DZ 31”

III.4 What are the detection prospects?

~ 10 detections peryear VISTA/LSST

from SR++ (2016)squares: Zwicky Transient Facility (ZTF; for depth 21 mag)circles:
60 second exposures VISTA (J & K band) and LSST (grizy)
diamonds: 180 second exposures VISTA (J & K band) and LSST (grizy)

III.5 Comparison with “macronova candidate” GRB130603B• possibly first detection of “macronova” in the aftermath ofshort GRB 130603B
9 days30 days after burst

sGRB130603B (Tanvir et al. 2013):
optical- T90≈ 0.18 s, z= 0.356- nIR-transient, present at≈9 days, but faded away
after 30 daysnIR• If interpretation correct:- short Gamma-ray Bursts caused by compact binary mergers- compact binary mergers are a major source of rapid neutron capture elements- isotropic macronovae promising accompanying signature for “chirp” GW signals

black: HST-band F160W (1.5 µm)
nsns

N1: dimmest nsns (1.2 +1.2 M⊙)
N5: brightest nsns (1.4+1.8 M⊙)

red: HST-band F606W (peak: 5.9 µm)

nsbh

B1: nsbh (1.4 +7.0 M⊙, χ= 0.7)
B2: nsbh (1.4 +7.0 M⊙, χ= 0.9)
B3: nsbh (1.2 +7.0 M⊙, χ= 0.9)

⇒ several models are “near” the observed properties(e.g. nsns 1.4 + 1.8 M⊙,
nsbh 1.2 M⊙ + 7.0 M⊙, χ= 0.9,
wind mej= 0.10 M⊙ , vej= 0.10 c, Ye= 0.25)

many models much dimmer

wind

wind7: mej= 0.05 M⊙, vej= 0.05 c,
Ye= 0.25
wind8: mej= 0.05 M⊙, vej= 0.01 c,
Ye= 0.30
wind9: mej= 0.10 M⊙, vej= 0.10 c,
Ye= 0.25

V. Summary• EM-transients crucial for GW-source localization• EM-transients ⇔ matter ejection ⇔ r-process nucleo. ⇔ sGRBs• various ejection channels:• dynamic (very low Ye component ⇒ “strong r-process” ⇒ large opacities ⇒ “late + red”)• various types of “winds” (higher Ye ⇒ “weak” r-process ⇒ lower opacities ⇒ “early + blue”)• nuclear physics has substantial impact on observability, e.g. α-decaying trans-leadnuclei (DZ31) impacts on• heating rates• thermalization efficiencies⇒ substantially brighter lightcurves• detection prospects: LSST could see ~10 nsns-events yr-1
• several of our dyn. ejecta and wind models are good candidates for “Tanvirevent” (prov. DZ31, trans-lead α-decay, standard opacities), but many will be very
difficult to observe• caveats of current MN model:I) opacities
II) wait for full radiative transfer