EXPECTED GAMMA-RAY EMISION FROM X-RAY BINARIES

Transcription

1 EXPECTED GAMMA-RAY EMISION FROM X-RAY BINARIES W lodek Bednarek Department of Astrophysics, Lódź, Poland

2 I have to admit I don t know? It seems to me very complicated! I will try to convince you why I have such opinion

3 Historical notes Observations General scenarios Conditions within the binary system IC e ± pair cascade processes (linear, isotropized, magnetic field driven) Effects of e ± pair energy losses in the magnetic field Dependence on the shock localization (variable stellar wind) Effects of anisotropic stellar/pulsar winds Effects of clumpy stellar wind Effects of relativistic boosting of radiation Double shock structure - acceleration of two populations of electrons Acceleration of electrons and/or hadrons?

4 Historical notes: Pulsars within binaries high energy emission (Bignami et al. 1977, Vestrand & Eichler Cyg X-3; Maraschi & Treves LS I ) TeV-PeV γ-ray emission (???): Her X-1, Vela X-1, Cyg X-3,... (lack of confirmation - Weekes 1992) Absorption of γ-rays in stellar radiation (Protheroe & Stanev 1987, Moskalenko et al. 1993) Some binaries GeV emitters (?) - EGRET error boxes LS 5039 (Paredes et al. 2000), Cyg X-3 (Mori et al. 1997), LS I (Thompson et al. 1995), Cen X-3 (Vestrand et al. 1997) IC e ± pair anisotropic cascades in the stellar radiation (Bednarek 1997, 2000) Discovery of γ-ray binaries at TeV energies (LS2883/PSR Aharonian et al. 2005; LS Aharonian et al. 2005; LS I Albert et al. 2006) Recent modelling of γ-ray emission since 2005:

5 γ-ray observations - main features Modulation of γ-ray signal from LS 5039: GeV light curve TeV light curve Figure 1: LS 5039: GeV emission from Abdo et al. (2011); TeV emission from Aharonian et al

6 Long term γ-ray variability Figure 2: LS I TeV γ-ray light curves ( ) from Acciari et al. (2011).

7 Spectral features LS I Eta Carinae Figure 3: LS I from Abdo et al. (2011); Eta Carinae from Farnier et al. (2011).

8 Three types of gamma-ray binaries Massive star + energetic pulsar: LS 2883/PSR , LS 5039, LSI Massive star + accreting black hole: Cyg X-3, Cyg X-1 (?) Two massive stars: Eta Carinae (1) The geometry of acceleration may or may not differ significantly: (2) Physical conditions rather differ significantly (V p, ξ, B): rad 1 jet rad 2 rad shock star disk star pul

9 Conditions within the binary: Massive star Magnetic field structure Wind structure tor rad dip star polar wind star equatorial wind B(R) R 3 (dip); R 2 (rad); R 1 (tor). Polar wind: v 10 3 km/s; Equatorial wind: v km/s;

10 Propagation of γ-rays within binary system Figure 4: Star: surface temperature T = K, radius R = cm, distance of the injection place D = 1.4R, E γ = 1 TeV (from Bednarek 2000). Simple scaling for stars with other parameters: τ( Eo γ S T, T, R, D, α) = S 3 T S Rτ(E o γ, T o, R o, D, α), where S T = T /T o, S R = R /R o and D in R or R o.

11 Magnetically driven: Re-directed γ-rays around B direction Types of the IC e ± pair cascade scenarios Aharonian et al. (2006), Cerutti et al. (2009); Bednarek (1997,2000,2006); Sierpowska & Bednarek (2005) linear γ isotropized magnetically driven γ e γ e γ e γ γ γ B star star star γ source γ source γ source Note: E e = 1 TeV, B = 1 G R L cm << R. Linear: γ-rays strongest to stellar limb Isotropized: Focusing of γ-rays by stellar radiation

12 Main features of the γ-ray cascades Figure 5: LS 5039 time averaged cascade spectrum: from Aharonian et al. (2006). Spectra: injected (dashed); cascade (solid); simple abs (dotted) GeV bump; TeV emission

13 Magnetically driven cascades: distribution of cascade γ-rays e γ star e γ star γ e star Figure 6: Distribution of γ-rays on the sky for injection angles: 90 o, 120 o, and 150 o (from Sierpowska & Bednarek 2005).

14 Synchrotron energy losses of e ± pairs P syn < P T IC B s < B T = 40T 2 4 G (stellar surface) Figure 7: From Bednarek (1997): T s = K. U B R 4, U rad R 2 Periastron - TeV electrons synchrotron losses important? Reason for some TeV γ-ray modulation (peri/apo)?

15 Synchrotron spectra from cascade e ± pairs Synchrotron emission: primary electrons: Bednarek & Giovannelli (2007) Synchrotron emission: secondary cascade e ± pairs (constant B): Figure 8: From Khangulyan et al. (2008); Bosch-Ramon et al. (2008)

16 Variable stellar wind: TeV emission at periastron? Change in stellar wind change in shock localization 1 2 α α 1 2 NS star shock1 shock 2 Angles α 1 and α 2 differ significantly

17 Cascade spectra for different obs. angles Figure 9: IC e ± pair cascade spectra for different obs. angles: 30 o 120 o (from Bednarek 2000)

18 Anisotropic stellar/pulsar winds Figure 10: Shock structure very complicated: from Sierpowska-Bartosik & Bednarek (2008). Complicated geometrical situations can be expected: At some phases shock structures may change drastically The shock structures may change with binary periods Shock might appear very close to the pulsar or massive star

19 Both winds aspherical Pulsar wind aspherical: e.g. Bogovalov (1999) Be stellar wind aspherical: e.g. Waters et al. (1988) polar stellar wind NS pulsar wind pulsar wind shock I pulsar wind shock II NS Be star equatorial stellar wind pulsar wind

20 Shock structures: PSR /SS2883 Figure 11: Location of the shock in PSR /SS2883: from Sierpowska-Bartosik & Bednarek (2008). Post-Shock flow can accelerate to γ 100: see Bogovalov et al. (2008) See also the case of LS I : Sierpowska-Bartosik & Torres (2009)

21 Effects of relativistic boosting of radiation See previous talk Dr G. Dubus Relativistic jet: Dubus et al. (2010a) Relativistic flow along pulsar cometary tail : Dubus et al. (2010b) Figure 12: From Dubus et al. (2010).

22 Double shock structure - two populations of electrons? Tavani & Arons (1997): PSR /SS 2883 e,p radiation radiation e,p WR Eta Carinae wind wind Figure 13: Shock structure in massive binary system Eta Carinae: from Bednarek & Pabich (2011). different conditions at the shocks (B, ξ) acceleration of electrons (hadrons?) to different maximum energies

23 Gamma-ray emission from electrons accelerated at the shocks s -1 ) dn/de / erg cm 2 log(e log(e / GeV) Figure 14: Shock structure in massive binary systems: from Bednarek & Pabich (2011). Electrons from the shock in Eta Carinae wind (solid) and WR star (dashed/

24 Two populations of electrons in pulsar/massive star binaries? E max sh,pul 63(ξ/B) 1/2 10P 100 (ξ 1 D 12 /B 12 ) 1/2 TeV. (1) E max sh,w 1.3(ξB sh ) 1/2 (D sh /T 2 4 ) 130ξ 1/2 4 B 100/T 2 4 GeV. (2) Figure 15: Energies of accelerated electrons at the pulsar, stellar shock: from Bednarek (2011, in preparation).

25 Effects of clumpy stellar wind? Figure 16: From Zdziarski et al. (2010). Clumps R cm; Pulsar wind mix (confined) with the matter and mag. field of clumps see also the model for jet-clump interaction, e.g. Owocki et al. (2009), Araudo et al. (2009)

26 Gamma-ray emission from electrons in clumpy wind Figure 17: Models: dominated by IC losses (upper), synchrotron losses (bottom): From Zdziarski et al. (2010). Pulsar: electrons with γ e 10 8 ; Stellar wind: magnetic field B 2 G. Transition between models: TeV γ-ray variability?

27 Acceleration of electrons and/or hadrons? Too strong synchrotron losses Hadronic γ-rays? (Aharonian et al. 2005) Hadronic models: e.g. Romero et al. (2003,2005); Kawachi et al. (2004); Chernyakova et al. (2006); Torres & Halzen 2007; Araudo et al. (2009); Owocki et al. (2009); Bednarek & Pabich (2011) Figure 18: Neutrino spectra from Eta Carinae (Bednarek & Pabich 2011).

28 CONCLUSION: Many effects can play essential role in formation of emission features of γ-ray binaries Very important role of geometry (processes occur aspherically) Non-steady medium (aspherical, variable, inhomogenous winds) Different radiation processes Different populations of particles Binary systems are one of the best defined but quite complicated astrophysical objects Reliable predictions of γ-ray emission features difficult