Péter Mészáros Pennsylvania State University

X- and gamma-raycounterparts of GW sources

Gravitational Wave Physics & Astronomy WorkshopGWPAW2015, Osaka, June 2015

STELLAR MASS SYSTEMS

Short GRB- DNS inspiral GW

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If SGRB are indeed DNS or BH-NS mergers, A-LIGO/A-VIRGO should find few/year

Mészáros

Mészáros Hei08

3 Usual Phases of Rotating Collapse

• In-spiral (binaries, or core blobs)

• Merger - central condensation + disk, subject to instabilities (again blobs?)

• Ring-down

Mészáros

Kobayashi & Mészáros, 02, ApJ 589:861

Mészáros

Mészáros

X-Gamma counterpart: short GRB 050724

GRB 050724

SWIFT EM observations:

- Prompt, hard MeV emission t < 2 s - Sometimes (30%) followed by a

~100 s hard X-ray tail- Also standard (Swift) XR afterglow

with (50%) steep decay, followed by shallow plateau+ flares, followed by (100%) traditional power law decay

very likely a DNS merger

(more in Edo Berger;s talk)

Mészáros

GRB 090510

Two notable features of traditional “short hard bursts”:1) A spectrum extending to GeV, in a second, hard component2) Delay between GeV emission onset and MeV trigger- GeV ag

Fermi LAT/GBM observations

LIGO-Virgo SR6-SR5 GRB search

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Abadie & LIGO-Virgo collaborat. 2012, ApJ 760:12

←Best strain noise spectra: DDNS<16 Mpc, DBHNS <28 Mpc

Search for GW from 154 GRB in 2009-2010 from LIGO SR-6 & Virgo SR-5

(before adv. upgrades)

See also: Aasi+14 (IPN), PRL113, 011102

Collapsar GRB GW

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3

FIG. 3: Snapshots of the meridional density distribution withsuperposed velocity vectors in model u75rot1 taken at varioustimes. The top left panel (note its special spatial range) showsa snapshot from 10ms after bounce. The top right and bot-tom left panels show the point of PNS instability and the timeat which the AH first appears, respectively. The bottom rightpanel, generated with a separate color range, shows the hy-peraccreting BH at ⇠ 15ms after its formation. All colormapshave density isocontours superposed at densities (from outerto inner) of ⇢ = (0.1, 0.25, 0.5, 0.75, 1.0, 2.5, 5.0)⇥1010 g cm�3.

horizon (AH) appears within ⇠1 ms and quickly engulfsthe entire PNS. With the PNS and pressure support re-moved, postshock material and the shock itself immedi-ately subside into the nascent BH. The bottom panel ofFig. 2 shows the evolution of BH mass and dimensionlessspin a? in all models. The former jumps up as the AHswallows the PNS and postshock region, then increasesat the rate of accretion set by progenitor structure andis largely una↵ected by rotation at early times. The di-mensionless spin reaches a local maximum when the BHhas swallowed the PNS core, then rapidly decreases assurrounding lower-j material plunges into the BH. Thisis a consequence of the drop of j at a mass coordinateclose to the initial BH mass (cf. Fig. 1). Table I summa-rizes for all models the values of a? at its peak and at thetime we stop the LR run.

In Fig. 3, we plot colormaps of the density in the merid-ional plane of the spinning model u75rot1 taken at var-ious postbounce times. The rotational flattening of thePNS is significant and so is the centrifugal double-lobedstructure of the post-BH-formation hyperaccretion flow.The latter is unshocked and far sub-Keplerian with in-flow speeds of up to 0.5c near the horizon. The flow willbe shocked again only when material with su�ciently

FIG. 4: Top: GW signals h+,e emitted by the rotating modelsas seen by an equatorial observer and rescaled by observerdistance D. Bottom: Spectrogram of the GW signal emittedby the most rapidly spinning model u75rot2.

high specific angular momentum to be partly or fully cen-trifugally supported reaches small radii (cf. [16]). Basedon progenitor structure, our choice of rotation law, andthe assumption of near free fall, we estimate that thiswill occur after ⇠1.4 s, ⇠2.4 s, ⇠3.9 s in model u75rot2,u75rot1.5, u75rot1, respectively. At these times, theBHs, in the same order, will have a mass (a?) of ⇠8 M�(0.75), ⇠14 M� (0.73), and ⇠23 M� (0.62).GW Signature.—The top panel of Fig. 4 depicts the

GW signals emitted by our rotating models. Due to theassumed octant symmetry, GW emission occurs in thel = 2, m = 0 mode. The nonrotating model leads toa very weak GW signal and is excluded. At bounce, astrong burst of GWs is emitted with the typical signalmorphology of rotating core collapse (e.g., [27]) and thepeak amplitude is roughly proportional to model spin.Once the bounce burst has ebbed, the signal is domi-nated by emission from turbulence behind the shock. Itis driven first by the negative entropy gradient left by thestalling shock and then by neutrino cooling, whose e↵ectmay be overestimated by our simple treatment. Interest-ingly, the signal strength increases with spin. This is notexpected in a rapidly spinning ordinary 2D CCSN, sincea positive j gradient in the extended postshock regionstabilizes convection. In our models, the postshock re-gion is considerably smaller and shrinks with postbouncetime. The driving entropy gradients are steeper and thechange of j in the postshock region is smaller. Also, incontrast to 2D, our 3D models allow high-mode nonax-isymmetric circulation. We surmise that the combinationof these features with increasing spin (feeding greater cir-

C. Ott et al, 2011, PRL106:161103

← Model u75rot2Use 75 M⊙ rot. prog.model Woosley-Heger 02, 10-4 Zsun, 3+1 GR calculationEGW=3.4 10-7 M⊙ , fc=807 Hz

Undetectable unless in Milky Way

Chaotic infall: very small quadrupole

BH-torus in GRB collapsar :

Papaloizu-Pringle instability:

big quadrupole

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But:

Kiuchi, Shibata et al, 2011, PRL 106:251102

Detectable at 100 Mpc...?

A-LIGO & A-VIRGO

GW/GRB coincid. expected

• Grey: GW-all GRB

• Blue: GW- GBM

• Red: GW-BAT

Clark et al. 1409.8149

Mészáros

Swift Era: frequent X-ray afterglow plateau in many GRBs (long, short):

Magnetar phase = GW source? • It is one of the

explanations for Swift X-ray plateaus (→energy injection)

• If so, magnetar must be fast rotating (collapsar paradigm)

• Fast rotation → bar instability?

• If so → GW emiss.

Corsi & Mészáros ’09ApJ 702:1171

Mészáros

GW + EM dipole losses

• Upper:

• Red: EM dipole energy losses ;

• Dot-dash: GW losses without EM loss term

• Solid black: GW losses with EM loss term

• Lower:

• Surface fluid effective angular velocity Ωeff/π, where Ωeff= Ω -Λ (pattern minus peculiar) along a Riemann seq.(e.g. Lai-Shapiro)

GW: with pattern Ω - EM: from frozen-in surface field

Corsi &

Bar instability → rotating ellipsoid

Mészáros

GW & EM loss effects

• Black-solid: GW+EM

• Black-dash-dot: GW only

• Blue-dot: Virgo nom.

• Purple dash: adv. LIGO/Virgo

Upper: GW amplitude hc @ d=100 Mpc, for:

Lower: GW signal frequency. for:

- Black-solid: GW + EM losses-Black-dash: GW losses (only)

Corsi & Meszaros 09

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Thus: GW counterpart in GRB may be X-ray plateau phase (caused by temp. magnetar)

• It is one of the explanations for Swift X-ray plateaus (→energy injection)

• If so, magnetar must be fast rotating (collapsar paradigm)

• Fast rotation → bar instability?

Corsi & Mészáros ’09ApJ 702:1171

Magnetar likelier in DNS

• Also, DNS merger likely to lead to wind-like mass ejection, mostly isotropic

• SGRB only seen along jet (kilonova, macronova, etc. - - many authors)

• BUT: jet radiation may be scattered by wind leading to isotropic X-ray plateau

And in fact….

Kisaka, Ioka & Nakamura, 1506.02030

DNS merger isotropic XR

• The scattered jet radiation is in XR, large angle (~isotr.)

• Shocked ejecta produces isotr. opt/IR flux

Kisaka et al, 1506.02030

DNS Magnetar plateau XR light curve & detectability

• Dashed: black plataeau XR emission, jet angle 0.1, rad., @ dist. 100 Mpc

• Solid red: isotrop. scattered XR emission, scatt. efficiency 10-3

Kisaka et al, 1506.02030

LARGE MASS RATIO SYSTEMS

Tidal disruption of star by MBH

• When star’s self-gravity ~ BH differential attraction Gm*/R*

2 ~ GMh [1/r2 -1/(r+R*)2] ~ (GMh/r2)(2R*/r) → Gravitational tidal disruption, within the tidal radius rt , r < rt ~(2m*/Mh)1/3 R* ~ 5.1012 M6

1/3(R*/R⊙)(M*/M⊙)-1/3 , • Compared to Schwarzschild radius,

Rg=2GMh/c2=3.1011M6 cm, rt/Rg ~24 (R*/R⊙)(M*/M⊙)-1/3 M6

-2/3

(Rees, 1988, Nat. 333, 523)

Tidal  Disrup,on  Event

A  star  ventures  inside  the  ,dal  radius  of  a  black  hole  and  is  torn  apart(slide credit M. Begelman, 2015)

Disruption vs. swallowing• Different stellar types have

tidal radii rt comparable to the Schwarzschild radii Rg of different MBH masses

• (e.g. solar type stars are swallowed whole by a 108 M⊙ MBH, since rt <Rg)

• For a 106 M⊙ solar type stars are disrupted, but light He stars and WDs are swallowed whole

• Further swallowing helped by periastron distance rp << rt → penetration factor βp=rt/rp >>1

• He and WDs are preferentially disrupted by BHs of smaller mass <105M⊙, where rt > Rg

(Kobayashi et al, 2004, ApJ, 615:855)

   Tidal  forces  …

...  unbind  ~half  the  debris  

…  throw  the  other  half  into     highly  eccentric  orbits  

Semi-­‐major  axis:

(Slide : M. Begelman, 2015)

Disruption simulation

Solar-type stellar orbits, ↑ , for βp = rt/rp = 1, 5, 10 penetration parameter in Schwarzschild black hole of 106 M⊙

SPH numerical calculation, βp =10 → snapshots at 8 different instants t=-335,-236,-138,-40,50,157,255,353

(Kobayashi et al, 2004, ApJ, 615:855)

SPH - GR

Simula,ons  by  Guillochon  &  Ramirez-­‐Ruiz  2013

Very super-Eddington!

Messy: intersecting streams & shocks

t = 2t0 t = 6t0

Based on numerical simualtions of Shiokawa et al 2015 ApJ, 804:85

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But, with increasing numerical detail:

Resulting accretion onto BH

• Black line approx. the dM/dt~ t-5/4

analytical behavior expected for the fall-back

• But- actual details dep, on combinat. of penetration depth and viscous (alpha) parameter of accr. flow

Shiokawa et al 2015 ApJ, 804:85

Radiation (X-ray)

• At the simplest level, if periastron rp < rt (βp >>1)

→ compression, shock, heating, with Eth ~Gm*/r* ~1048(M*/M⊙)2(R*/R⊙)-1 erg

• kT~ Gm*mp/r* ~ 1 (M*/M⊙)(R*/R⊙)-1 keV,

→ prompt X-ray flare/outburst expected, • Orbital time Δt ~ r*/vp ~10 (M*/M⊙)-1/6 (R*/R⊙)3/2 M6

-1/3 s

• Opacity τT~ M*σT/4πmpR*2 ~1010 (M*/M⊙)(R*/R⊙)-2

• During Δt , radiation from the heated star can diffuse out from a depth D~(c Δt R*/τT )1/2

Two phases: disruption flare and extended fall-back

X-ray flare & extended emission

• Prompt flare: Lx,pr~(Eth/Δt)(cΔt/R*τT)1/2 ~< 1042 (M*/M⊙)17/12(R*/R⊙)1/4 M6

-1/6 erg/s

Δtx,pr ~Δt ~10 (M*/M⊙)-1/6 (R*/R⊙)3/2 M6-1/3 s → Lx ~ t-1/2

• Extended duration X-ray cooling: from gas falling back to rp , a fraction f0.1=f/0.1 self-intersects:

Lx,ext~ 1045 f0.1 M61/6 (M*/M⊙)7/3 (R*/R⊙)-5/2 erg/s

Δtext ~ 10 M6 (M*/M⊙)-1 (R*/R⊙)3/2 day → Lx,ext ~ t-5/3

(Kobayashi, Laguna, Phinney, Mészáros, ApJ, 615:855)

Tidal disruption GW characteristics

• Dimensionless strain h at distance D : h ~ (G/ D c4) dQ2/dt2 ~(Rg/D)(v/c)2

~ 2.10-22βp(D/10 Mpc)-1M62/3(M*/M⊙)4/3 (R*/R⊙)-1

• Wave frequency: for penetration βp =rt/rp , f ~ (GMh/rp

3)1/2~6.10-4βp3/2(M*/M⊙)1/2(R*/R⊙)-3/2 Hz

→ eLISA- frequency is OK for solar stars, but it is not sensitive enough

however : DECIGO — sensitive for He stars, WDs

Disruption of Solar type stars, Mh=106 M , a=0

Average temperature near periastron Gravitational strain at D=20 Mpc, SPH: solid line, point particle: dashed

(Kobayashi, Laguna, Phinney, Mészáros, ApJ, 615:855)

↓ Helium star and ↓ Kerr MBH

↑ He star, Schwarzschild hole (a=0), βp=1, temperature and strain, polarizations: h+: thick, hx:t hin

Solar star, Kerr hole, a=1, βp=5. Prograde: solid, retrograde: dashed SPH: thick; point particle: thin

(Kobayashi, Laguna, Phinney, Mészáros, ApJ, 615:855)

He and WD stars @ 10 Mpc

Type M/M⊙ R/R⊙ f h Fx,fl <

He 0.5 8.0-2 0.02 9.9-22 7.7-10

WD 0.4 1.67-2 0.17 3.5-21 3.8-9

WD 0.8 1.0-2 0.54 1.5-20 2.2-8

WD 1.4 3.3-3 3.68 9.4-20 2.1-7

f x βp3/2 , h x βp(D/10 Mpc)-1M6

2/3 , Fx,f x (M*/M⊙)17/12(R*/R⊙)1/4 (D/10 Mpc)-2 M6

-1/6 erg/cm2/s

eLISA&

DECIGO

● For Mbh ≥ 105 M⊙

● DECIGO detect GW ofHe★, WD★ to 100 Mpc,

● with X-ray flare fluxes Fx,fl ≥ 10-11 -10-12 erg/cm2/s

eLISA

TDE in a wind

• Only small fraction of returning debris can accrete

• The rest makes a tenuous unbound slow lateral outflow (wind)→reprocess

• Can trap inner XR/UV rad’n at late stages after flare, →low Teff optical TDEs

• For MBH>107 M wind is ionized, → can result in XR bright flares

Metzger & Stone, 1506.03453

Another factor affects emerging radiation :

DOUBLE MASSIVE SYSTEMS

Massive BH binary

(MBHB)

fmerg ~ (G Mtot /a3merg)1/2 ~ 4x10-3 M6-1 Hz

hmerg ~(4π GMtot/c4 D) a2merg f2merg ~1.26x10-16 M6 (D/10Mpc)-1

→eLISA target — and also EM X-ray counterparts!

What the space-time does:

Massive BH binary (MBHB)

but:Accretion disk around two BHs →M★ primary, ms secondary, rs separationrH Hill radius

Kocsis et al 2012, MN 427:2689 Mind the gap !

What does the gas do?Tries to accrete (can’t help it);

Not simple accretion disc !

Disk shrinking initially keeps up with binary shrinking

• At radius where the visc. stress ~ grav. torque:

• →Disk is truncated at an inner edge redge ~ 2a , where a=semi-major axis (gas pile up, “flood gate”)

• As binary separation shrinks, inner edge moves in, keeping redge ~ 2a

BUT:▪ tgw =(3/64)(c5a4/G3M3)(1+q)2/q ~ a4

▪ tvisc=(2/3) r2 /ν(r) ~ a2 (with r~a, ν(r)~const)

- As “a” shrinks, eventually tgw < tvisc

- disk decouples at rdec, - MBHB merges at ~tgw/4 after tdec

- Disk shrinks in tshr~tvisc ~tgw(dec), arrives at merged BH ~(3/4)tshr after merger → rapid accretion →

→ strong X-ray emission (AGN disk)(Milosavljevic & Phinney 05 ApJ 622:L93)

EM spectrum & XR detect.

(Milosavljevic & Phinney 05 ApJ 622:L93)

- Eddington limited disk spectra for mergers of a*=0.99 and Mbh=104, 105, 106 M⊙ (R to L)

- Before decoupling: disk inner edge large → spectrum soft (UV/O, thin lines)

- After BHs merge and disk reaches center → hard spectrum (X-rays, thick lines)

- Chandra detect to z~1; future XR mission z~10

Actually, merger & disk shrinking dynamics can be more complicated

• Before decoupling, inflow rate in gap (dM/dt)gap may be up to ~0.1 (dM/dt)disk

• Different shrink/merge regimes dep. on mass ratio, gas infall rate, etc.

• Have different gap types and migration regimes

• EM spectr. luminosity in units of LEdd=1043 erg/s for an Mbh=105 M⊙ ↗ , hotter than single disk due to binary heating Kocsis et al 2012, MN 427:2689

Outer disk

Inner disk

MBHB inspiral & accr.: even

more complex?• 2D num.calculations show:

pre-merger cavity lopsided, with narrow accretion streams and minidiscs

• Pre-merger accretion rate is larger due to the narrow accr. streams and mini-discs

• Post-merger accretion rate lower, because of a smooth single disc

Farris et al 2015 MN 447:L80

Surf. brightness

Surf. density

Before At merger After

And in fact…

2D MBHBinsp. & accr.

• Top: spectrum before (right curves, hard) and after merger (left curves, soft).

• Harder XR spectrum before merger due to shock-heated minidiscs

Farris et al 2015 MN 447:L80

Disc spectrum @ various epochs

Light curve @ various energies

Mészáros

• Stellar DNS-BHNS merger: SGRB, magnetar(?), plateau(?)- strong gamma/X counterparts, appreciable optical/IR

• Medium (105-107 M⊙) massive BH stellar tidal disruption (large mass ratio): substantial XR, stronger O/IR counterpart?

• Massive BH binary: strong XR changing to strong optical (or vice-versa)?

GW XR-Gamma Counterparts :

Thus,

Mészáros

counterpart luminosity uncertainty

• Stellar merger: ~ fairly robust

• mBH tidal disr.: ~ fairly uncertain

• MBHB merger: ~quite uncertan

However:

• A network of various complementary messenger facilities

• http://amon.gravity.psu.edu/amon_system.shtml

• Trigger: sub-threshold signal in 2 or more ≠ messengers

Synergy between all 4 messengers together:

An agglomerator of counterparts:

AMON

IceCube Collaboration Meeting/Munich Fall 2013AMON

What AMON Will Do•Accept (sub-threshold) data from triggering observatories, especially data that cannot be used standalone, e.g.:

•single muon neutrinos from IceCube,•two-interferometer coincidences from LIGO-Virgo,•sub-threshold Swift BAT signals, etc.

•Streamline administrative overhead:

•Run temporal and spatial coincidence algorithms provided by users

•Issue “alerts” when coincidences are found to initiate closer investigation or OFU

•Help usher in the era of multimessenger astronomy!52

→

status quofuture

IceCube Collaboration Meeting/Munich Fall 2013AMON

AMON Members

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Triggering Observatories Follow-up Observatories

http://amon.gravity.psu.edu/participants.shtml

X,G counterparts

• GRBs (short): strong X- and gamma counterparts of GW

• GRB (long): strong X- and gamma, but GW: weak? medium?

• Tidal disruption of star by 105-106 MBH: adv. space detector

• MBHB inspiral, merger: strong X-ray emission likely, characteristics however still being debated.

Summary :

Thanks!

GW facilities• LIGO , VIRGO : ~10-103 Hz

• Geo600, TAMA300 : ~10-103 Hz

• Einstein Tel. (undergr.): 1-104 Hz

• KAGRA (Kamioka) : 1-104 Hz

• DeciGo/BBO (space): 10-2 - 102 Hz

• eLISA/NGO (space): 10-4 - 1 Hz

SGRB-LGRB both GW detectable?

57SGRB detectable — but LGRB: more rarely?

Kobayashi & Mészáros 02, ApJ 589, 861

DNS binary merger simult with a SGR: - GW from isnpiral, merger, ringdown

LGRG collapsar: rotation → assume instability →break-up into 2+ inspiraling blobs → GW

D=220 Mpc D=27 Mpc

Einstein Telescope f ~ 1-104 Hzsensit~10x ALIGO

eLISA/NGO

• Space interferometer, three satellites: one mother and two daughters, separated by 106 km in an equilateral triangle, at constant distance from Earth, in the ecliptic plane, 1 AU from Sun, 20o behind Earth, f =10-4 Hz-1 Hz

DECIGO (Slides: Masaki Ando)

Sensitivities

GW science goals