1. Millisecond Magnetars as GRB Central Engines Todd Thompson The Ohio State University B. Metzger, N. Bucciantini, E. Quataert, P. Chang, J. Arons Thompson et al. 2004 Metzger et al. 2007, 2008a,b Bucciantini et al. 2006, 2007, 2008, 2009 Usov 1992 Thompson 1994 Wheeler et al. 2000

2. “A nagging question in all these models is what produces the ‘observed’ ultra-relativistic flow? How are 10 -5 M sun of baryons accelerated to an ultra-relativistic velocity with  ~ 100 or larger? Why is the baryonic load so low? Why isn’t it lower? There is no simple model for that. An ingenious theoretical idea is clearly needed here.” Piran’s questions circa 1999.

3. The Basic Character of GRBs What is required? –Lorentz Factor: 100 <  < 1000 (if internal shocks) –Timescale: ~ 10 - 100s (+ late-time activity +variability) –Total Energy: ~ 10 51 ergs, highly asymmetric, supernovae

4. The Basic Character of GRBs What is required? –Lorentz Factor: 100 <  < 1000 (if internal shocks) –Timescale: ~ 10 - 100s (+ late-time activity +variability) –Total Energy: ~ 10 51 ergs, highly asymmetric, supernovae  Kinetic luminosity ~ 10 51 ergs / 10s ~ 10 50 ergs s -1  Mass loss rate ~ 10 50 ergs s -1 /  / c 2 ~ 6  10 -7 M sun s -1

5. The Basic Character of GRBs What is required? –Lorentz Factor: 100 <  < 1000 (if internal shocks) –Timescale: ~ 10 - 100s (+ late-time activity +variability) –Total Energy: ~ 10 51 ergs, highly asymmetric, supernovae  Kinetic luminosity ~ 10 51 ergs / 10s ~ 10 50 ergs s -1  Mass loss rate ~ 10 50 ergs s -1 /  / c 2 ~ 6  10 -7 M sun s -1 Piran:“Why is the baryonic load so low? Why isn’t it lower? There is no simple model for that.”

6. The Answer The characteristic mass-loss rate in the wind driven from the surface of a newly-born neutron star is (Duncan et al 1986; Qian & Woosley 1996) for a timescale of 10 -100 seconds. If a millisecond magnetar, the wind has kinetic luminosity of 10 50 ergs/s, with  ~ 10 2 - 10 3, (or magnetization of  > 10 2 - 10 3 ).

7. The Progenitor Instability & Collapse Collapse Bounce, shock formation Prompt Shock Stalls (~ ms) R ~ 100 km Shock Revival & Explosion ~ 0.5 - 1 s KH cooling epoch, wind. Bounce Cooling Epoch

8. 300 km Burrows et al. (1995) Collapse  Bounce  Stall  Explosion  Wind  Cooling

9. Wind Emergence: Common to All SN Explosions Supernova Shockwave Wind Scheck et al. (2006) Qian & Woosley (1996)

10. What about rotation and magnetic fields? Espanek, 1999

11. Initially, not magnetically- dominated, only ~10s later. Magnetar-strength fields dominate dynamics in the cooling/wind epoch! Rotation? C. Thompson (1994); T. Thompson (2003), T. Thompson et al. (2004) During Supernova t ~ 0 - 0.5 s Magnetic Field: Sub-Dominant  Dominant t ~ 1 s t ~ 10 s

12. Stage I: Non-relativistic - Magneto-Centrifugal Winds RARA Schatzman (1962); Weber & Davis (1967) - Radial B-field forces wind material to corotate to R A. At large distances, B  dominates. - Increases both V ∞ and J & E & M loss rates. - Efficient spindown mechanism. - R A increases as L (t) decreases in time.

13. RLRL RARA Stage II: Relativistic - Magneto-Centrifugal Winds - R L = c/  bounds R A. - As R A approaches R L, V ∞ approaches c.

14. RLRL RARA Stage II: Relativistic - Magneto-Centrifugal Winds - R L = c/  bounds R A. - As R A approaches R L, V ∞ approaches c. -  = B 2 /(4  c 2 ) (~  ∞ ) if efficient conversion field dissipation. (Drenkhahn & Spruit 02; Arons 07)

15. RLRL RARA Stage II: Relativistic - Magneto-Centrifugal Winds - R L = c/  bounds R A. - As R A approaches R L, V ∞ approaches c. -  = B 2 /(4  c 2 ) (~  ∞ ) if efficient conversion field dissipation. (Drenkhahn & Spruit 02; Arons 07) - L (t) dictates evolution!

16. Neutron Star Birth in Real Time Theoretical calculations of proto-neutron star cooling provide expectations for the time evolution of luminosity. Also, SN 1987A! Pons et al. (1999)

17. Evolutionary Calculations Supernova Explosion Phase

18. Evolutionary Calculations Supernova Explosion Phase

19. Evolutionary Calculations Supernova Explosion Phase

20. Evolutionary Calculations Supernova Explosion Phase

21. Evolutionary Calculations Non-Relativistic Relativistic Ultra-Relativistic Force-Free Supernova Explosion Phase

22. What we’ve done We solve the free 1-D equatorial non-relativistic neutrino-heated magneto-centrifugal wind problem for a ms-proto-magnetar. –Derive wind structure (T, , V, X ), the spindown rate, the mass loss rate -- given a monopole field geometry. (Thompson et al. 2001, 2004; Metzger et al. 2007, 2008a,b) Solve the free 1-D and 2-D axisymmetric relativistic MHD wind problem in the regimes  1 (relativistic, magnetically dominated). –Assess 1-D assumptions! –Calculate spindown explicitly when  >> 1, but with significant mass-loss (spindown is not “force-free,” not “vacuum” dipole,  is not infinity, not Usov 1992!). –Calculate interaction of wind with the overlying progenitor. (Bucciantini et al. 2006, 2007, 2008, 2009)

23. Evolutionary Calculations (Metzger et al. 2007)

25. Conclusion: Millisecond proto-magnetars drive winds with Exactly in the range needed for GRBs. What about asymmetry?

26. Jet Production Wind hits spherical SN shockwave in ~ few seconds. Creates reverse shock, contact discontinuity. Residual B phi squeezes (pinch force) flow. Net polar acceleration. Analogous to models of pulsar wind nebulae by Li & Begelman ‘92; Chevalier & Luo ‘93. See also Konigl & Granot ‘02; Komissarov & Lyubarsky 03,04; Uzdensky & MacFadyen ‘07; Nagataki ‘09).

27. Jet Production (Bucciantini et al. 07, 08, 09) Wind shocks on the outward-going SN shockwave at ~10 4 km, toroidal B exerts a pinch force that drives jet. Models are Poynting flux-dominated at breakout from progenitor.  jet  20  wind ~ 20 jet supernova

28. Jet Production (Bucciantini et al. 09) Magnetic dissipation in outflow is uncertain: consider both high & low magnetization limits. –Low-  :  wind (~R L ) ~  jet (~ R star ) –High-  : If  (~R L ) ~ 60, then  jet (~ R star ) ~ 10 -15 (but, accelerating), magnetic field ~5 times equipartition. –Jet production robust in either case. All of the spindown energy goes to the jet, virtually none goes to the spherical “supernova.” For weaker winds, the reverse shock propagates back into the outflow, coupling the wind & bubble dynamics: not a “free” wind (see also Uzdensky & MacFadyen; Komissarov & Barkov). Worries: 3D stability, entrainment, etc. (e.g., McKinney & Blandford)

29. Nucleosynthesis (Bucciantini et al. 09) For most cases, not hot enough for significant excess 56 Ni production. However, exact prediction is sensitive to timing of wind-bubble interaction. However, we expect 10 -2 M sun of Ne & O & C with v ~ 0.1 - 0.2c along the poles. The power of the wind (B, P) is directly connected with abundance of nucleosynthetic products, their velocity distribution. Compare w/ Nagataki et al; Maeda et al; Mazzali et al

30. For magnetar B & ms P, everything works for GRBs. How? –Weak physics sets baryon loading of the wind. –B & P set Lorentz factor, magnetization, & kinetic luminosity. –Neutrino diffusion time sets characteristic timescale. –Interaction of wind with progenitor leads to collimation. Robust to most uncertainties (e.g, magnetic dissipation) Nucleosynthesis connected to timing, wind power, etc. –Association with supernovae is free. –Non-Association with supernovae? Maybe AIC of WD. This model predicts well-defined (testable) correlations between these quantities in an average sense. They are not independent. The model appeals to the idea that neutron stars are probably born with a range of B & P, implying a connection with normal supernovae (e.g., Soderberg). Diversity: B(P). If you believe in the “Collapsar” mechanism, you (should) believe the ms-magnetar mechanism: angular momentum requirements. Summary

31. Wind & Spindown: –Couple physics more realistically to the supernova explosion, feedback effects on wind and magnetar torque in weak-wind limit –Does time dependence of outflow at magnetar surface imprint itself on the outflow dynamics of the jet at the progenitor surface? Jets: Stability in 3D? Entrainment? Lightcurves: –Use Lorentz factor, magnetization, opening angle, kinetic luminosity to predict prompt lightcurves (model-dependent). Try to make contact with observations. See Metzger et al. 2008: SGRBEEs!) Late-time activity: power-source for plateau? Flares? (Dall’Osso et al. today; talk by Margutti) LIMITATIONS: Energetics (e.g., Cenko et al. today): Maximum energy ~ few x 10 52 ergs (e.g., 090926A) The Present & Future

33. The observation of a long-duration GRB does not imply black hole birth (necessarily).

34.  Energy  Lorentz Factor  Timescale  Collimation Lightcurves?

35. Picture: WD/WD merger or AIC of a WD to a rapidly rotating magnetar. SGRB: ~ 0.1M sun disk accretes on ~ 0.1-1s. Polar GRB as in NS-NS merger calculations. (e.g., Dessart et al. 07) EE: Remaining magnetar spins down, drives wind. A Model for Short-Duration GRBs with “Extended Emission” Gehrels et al. 2006 Gal-Yam et al. 2006 GRB 060614

36. Proto-magnetar winds can account for EE:s Short GRBs with Extended Emission Metzger et al. (2008a)

37. Characteristics of the model: –As time increases, degree of relativity increases. –As time increases, energy loss rate decreases. –A magnetar is left behind. Late-time activity? Flares? –Maximum total GRB energy: ~ 5  10 52 ergs. Gravity waves. (e.g., Cenko et al. 2009) Some questions: –Magnetic dissipation? –What is the nature of SN-less long-duration GRBs? AIC? –What is the nature of the short-duration GRBs with extended emission? AIC? –Can we predict prompt lightcurves? (model-dependence) –With energy & Lorentz factor & jet opening angle, can we calculate character of afterglow? Yes. More