counterparts of gw sourcesf… · • extended duration x-ray cooling: from gas falling back to r...
TRANSCRIPT
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Péter Mészáros Pennsylvania State University
X- and gamma-raycounterparts of GW sources
Gravitational Wave Physics & Astronomy WorkshopGWPAW2015, Osaka, June 2015
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STELLAR MASS SYSTEMS
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Short GRB- DNS inspiral GW
3
If SGRB are indeed DNS or BH-NS mergers, A-LIGO/A-VIRGO should find few/year
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Mészáros
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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
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Mészáros
Kobayashi & Mészáros, 02, ApJ 589:861
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Mészáros
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Mészáros
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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)
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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
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LIGO-Virgo SR6-SR5 GRB search
11
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
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Collapsar GRB GW
12
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
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BH-torus in GRB collapsar :
Papaloizu-Pringle instability:
big quadrupole
13
But:
Kiuchi, Shibata et al, 2011, PRL 106:251102
Detectable at 100 Mpc...?
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A-LIGO & A-VIRGO
GW/GRB coincid. expected
• Grey: GW-all GRB
• Blue: GW- GBM
• Red: GW-BAT
Clark et al. 1409.8149
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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
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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
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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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Mészáros
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
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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
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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
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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
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LARGE MASS RATIO SYSTEMS
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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)
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Tidal Disrup,on Event
A star ventures inside the ,dal radius of a black hole and is torn apart(slide credit M. Begelman, 2015)
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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)
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Tidal forces …
... unbind ~half the debris
… throw the other half into highly eccentric orbits
Semi-‐major axis:
(Slide : M. Begelman, 2015)
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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
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Simula,ons by Guillochon & Ramirez-‐Ruiz 2013
Very super-Eddington!
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Messy: intersecting streams & shocks
t = 2t0 t = 6t0
Based on numerical simualtions of Shiokawa et al 2015 ApJ, 804:85
30
But, with increasing numerical detail:
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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
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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
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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)
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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
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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)
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↓ 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)
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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
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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
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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 :
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DOUBLE MASSIVE SYSTEMS
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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:
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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 !
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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)
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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
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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
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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…
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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
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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,
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Mészáros
counterpart luminosity uncertainty
• Stellar merger: ~ fairly robust
• mBH tidal disr.: ~ fairly uncertain
• MBHB merger: ~quite uncertan
However:
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• 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:
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AMON
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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
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IceCube Collaboration Meeting/Munich Fall 2013AMON
AMON Members
53
Triggering Observatories Follow-up Observatories
http://amon.gravity.psu.edu/participants.shtml
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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 :
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Thanks!
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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
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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
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Einstein Telescope f ~ 1-104 Hzsensit~10x ALIGO
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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
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DECIGO (Slides: Masaki Ando)
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Sensitivities
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GW science goals