1. Magnetic reconnection in stars: fast and slow D. J. Mullan University of Delaware, Newark DE USA

2. Flares in stars Flare: a transient increase in brightness In the Sun, flares occur in magnetic regions. Flare stars are known to have strong surface fields. Flares derive their energy from magnetic fields, Magnetic energy accumulates slowly, is released rapidly How to release magnetic energy? Magnetic reconnection (M.R.) Stars: is there a progression in the properties of M.R. along the main sequence?

3. Berger et al (2010): L Hα /L bol vs. Spec. type Notation fl. qus. up.lt Fl.; Qu.;U.L.

4. Berger et al (2010): L X /L bol vs. Spec. type

5. Berger et al (2010): L(radio) vs L(X-ray)

6. Why is L X proportional to L rad in F-(early)M stars? Guedel and Benz (1993): “common origin scenario” X-ray emission and radio emission rely on the same (or closely related) populations of electrons Electrons are relativistic (at least mildly so) First they interact with B  synch radio Then thermalize in ambient gas  X-rays

7. Towards later spectral types: ≥M7-8 Hα emission diminishes rel. to L bol : reduced deposition of mechanical energy in the chromosphere X-ray emission diminishes rel. to L bol : reduced deposition of mechanical energy in the corona But L(rad) increases rel. to L bol : an electron population survives which emits radio but not X-rays

8. Magnetic reconnection Berger et al: a decrease occurs in the efficiency of chromospheric and coronal heating at M6-M8 An increase occurs in the efficiency of radio emission also at M6-M8 What could cause main sequence stars to undergo systematic changes in chr/cor/radio as the spectral type increases?

9. Magnetic Reconnection (1): resistivity is dominant Transverse thickness δ of plasma sheet where reconnection occurs δ = δ SP = √(ηc 2 Δ /4πV A ) η=electrical resistivity; outflow V A =Alfven sp. δ SP = fn(T, N e, L, B) ( SP = Sweet and Parker,1958)

10. Ohm’s law E + v x B / c = ηJ + J x B /(nec) Convection resistance Hall effect Two regimes: If Convection=resistance  Sweet-Parker reconnection (resistivity dominates) If Convection = Hall effect  Hall reconnection dominates: two fluids are involved, with particles M 1 /M 2 >> 1

11. Reconnection: two regimes (1) Sweet-Parker reconnection: dominant wave mode = Alfven in the ions: speed of the wave V A is the same at all length scales Measured V A in solar active regions: CoMStOC (1988-1994)  no larger than a few tens of Mm/sec Elec. energy at v=30 Mm/sec is ≈ 3 keV

12. Reconnection: regime (2) (2) Hall reconnection: dominant wave mode = Whistlers: Wave speed increases at shorter length scales This difference makes Hall reconnection faster than S-P: models yield factor of 10 6 enhancement in reconnection rate Electrons escape at V ae  E e > 300 keV

13. Reconnection: 2 length-scales S-P reconnection occurs on diffusive length scale δ SP = √(ηc 2 Δ /4πV A ) Hall reconnection occurs on ion inertial length scale : d i = c/ω pi (ω pi = plasma frequency in the ions)  Electrons are magnetized, ions are not

14. The Hall hypothesis for flares Transition from slow to fast reconnection is predicted to occur when a reconnection site evolves to a condition where δ SP = d i Onset of a flare occurs when this conditions is first satisfied in an A.R. Theoretical basis of the Hall hypothesis: computer modeling Observational basis?

15. Conditions in stellar flares Data base: EUVE Observe stars at energies from 25 eV to 200 eV Good for observing flares: non-flaring stars (kT= few eV) emit little Flares: kT = 1 keV

16. Flare light curve analysis Observe: (i) τ d (decay time-scale) (ii) EM (Emiss. meas. at peak of flare) Assume: radiative cooling time is comparable to conductive cooling time Derive N, T, L (loop length) Calculate B from B 2 = 16πN kT (Mullan et al. 2006) 140 flares: N, L, B, T: wide ranges (10 3,10 3, 60, 15)

17. Stellar flare data: two length scales 140 stellar flares Single instrument Knowing T, N e, B, L Evaluate δ SP, d i Plot! Flare conditions are consistent with δ SP = d i i.e. when the Hall effect sets in

18. Flares in stars Hall effect onset brings significant ordering to the properties of stellar flares Flare build-up: Sweet-Parker slow reconnection Flare onset: when the SP diffusion region becomes as thin as the ion inertial length, Hall reconnection sets in Reconnection becomes rapid: FLARE

19. The Hall effect triggers a flare Reconnection occurs in two phases: (1) Sweet-Parker (slow): δ SP > d i (2) Hall effect (10 6 times faster): δ SP < d i Some active regions never flare: why not? Conditions never lead to δ SP as small as d i But slow reconnection leads to some enhanced coronal heating. T(A.R.) > T(diff. cor. =1.7-1.8 MK )

20. Further testing stars for fast (Hall) reconnection Two length scales: δ SP and d i Both depend on local parameters: N, T, L, B: Evaluate δ SP and d i in parameter space Limit T: (i) “hot” (corona) (ii) “cool atmos.” Resistivity: (i) Spitzer (ii) Kopecky (1958)

21. Hall reconnection onset: in parameter space: hot corona

22. Stars with hot coronae 150 representative pts in parameter space 90% of pts in “phase space” lie below the line δ SP = d i Reconnection in 90% of “coronal stars” is fast E e > 300 keV (“common origin hypothesis”) Flares in stars with spectral types G, K, and early M have L(rad) and L X

23. Hall reconnection onset in parameter space: cool atmos.

24. Stars with cool atmospheres 90% of points in “parameter space” lie above the line δ SP = d i Reconnection in 90% of “stars” is slow Flares are rare in stars later than M6-M7 No (nearly-)relativistic electrons to emit synch radio or heat ambient gas to X-ray emitting temps. But….

25. Radio emission: electron cyclotron maser (ECM) Melrose and Dulk (1982): maser is driven by a loss-cone distribution Condition: ω pe << Ω e Electron beam with speed v/c=0.1, density 10 7 cm -3 has ECM growth rate = 10 8 sec -1 : saturates in 100r e (< 1 km) v/c = 0.1  v = 30 Mm/sec (E = 3 keV: non-relativistic!)

26. ECM growth rates Treumann (2006): shell distribution: even faster growth rates for ECM emission Tang & Wu (2009): for E= 10’s of keV, ECM growth rates = 10 9 sec -1 TW: for E<3 keV, ECM growth rate decreases rapidly

27. ECM Radio emission in stellar atmospheres Electrons moving at speeds V of ≥30 Mm/sec are capable of significant ECM emission At a slow reconnection site (S-P): charged particles emerge at speed V A Where, in the n,B,L plane, does V A have values ≥30 Mm/sec ?

28. Hall reconnection onset: in the n,B,L plane: cool atmos.

29. Slow reconnection: an effective source of ECM Even if reconnection is slow in >M7 stars (i.e. no bona fide flares), in 70% of stars electrons can be ejected at speeds of ≥30 Mm/sec Effective source of ECM radio emission Coolest dwarfs are only rarely (10%) sites of flares (i.e. fast reconnection) but can be effective (70%) sites of coherent radio emission (ECM)