Hyperons & Neutron Stars

Transcription

1 Hyperons & Neutron Stars Isaac Vidaña CFC, University of Coimbra HYP2012 The 11 th International Conference on Hypernuclear & Strange Particle Physics Barcelona, October 1 st - 5 th 2012

2 In this talk I will review the role of hyperons on : v EoS & M max of Neutron Stars v Properties of Proto-Neutron Stars v Neutron Star Cooling v R-mode Instability of Neutron Stars

3 Some known facts about Neutron Stars Formed in: type II, Ib or Ic SN Mass: M ~ 1-2 M Radius: R ~ km Density: ρ ~ g/cm3 ρuniverse ~ g/cm3 ρsun ~ 1.4 g/cm3 ρearth ~ 5.5 g/cm3 Baryonic number: Nb ~ 1057 ( giant nuclei ) Magnetic field: B ~ G ( T) G Earth G Magnet 10 5 G 4.5x10 5 G Sunspots You are here!! 2.8x10 7 G Largest continuous Largest magnetic pulse field in lab. (FSU, USA) in lab. (Russia)

4 Electric field: E ~ V/cm Temperature: T ~ K Rotational period distribution è two types of pulsars: pulsars with P ~ s pulsars with P ~ ms Shortest rotational period: P B = 1.58 ms until the last discovery: PSR in Terzan 5: P J ad = 1.39 ms Accretion rates: to 10-8 M /year

5 Anatomy of a Neutron Star Equilibrium composition determined by ü Charge neutrality i q i ρ i = 0 ü Equilibrium with respect to weak interacting processes b 1 b 2 + l + ν l b 2 + l b 1 + ν l µ i = b i µ n q i µ e µ ν e ( ), µ i = ε ρ i

6 Hyperons in NS considered by many authors since the pioneering work of Ambartsumyan & Saakyan (1960) Phenomenological approaches ² Relativistic Mean Field Models: Glendenning 1985; Knorren et al. 1995; Shaffner-Bielich & Mishustin 1996, Bonano & Sedrakian 2012, ² Non-realtivistic potential model: Balberg & Gal 1997 ² Quark-meson coupling model: Pal et al. 1999, ² Chiral Effective Lagrangians: Hanauske et al., 2000 ² Density dependent hadron field models: Hofmann, Keil & Lenske 2001 Microscopic approaches ² Brueckner-Hartree-Fock theory: Baldo et al. 2000; I. V. et al. 2000, Schulze et al. 2006, I.V. et al. 2011, Burgio et al. 2011, Schulze & Rijken 2011 ² V low k : Djapo, Schaefer & Wambach, 2010 Sorry if I missed somebody

7 Hyperons are expected to appear in the core of neutron stars at ρ (2-3)ρ 0 when µ N is large enough to make the conversion of N into Y energetically favorable. µ Σ = µ n + µ e µ ν e µ Λ = µ n n + n n + Λ p + e Λ + ν e n + n p + Σ n + e Σ + ν e

8 Effect of Hyperons in the EoS and Mass of Neutron Stars stiff EoS stiff EoS soft EoS soft EoS Relieve of Fermi pressure due to the appearance of hyperons è EoS softer è reduction of the mass

9 Measured NS Masses (up to 2006) Phenomenological: M max compatible with M (Glendenning 1991) Microscopic : M max < M (Lattimer & Prakash 2007) (Schulze et al 2006)

10 Recent measurements of high masses life of hyperons more difficult PSR J (Freire et al. 2009) post-kelperian parameters: ü binary sytem (P=95.17 d) ü high eccentricity (ε=0.437) ü companion mass: ~ 1M ü pulsar mass: M =1.67± 0.11M PSR J (Demorest et al. 2010) Shapiro delay: ü binary sytem (P=8.68 d) ü eccentricity (ε=1.3 x 10-6 ) ü companion mass: ~ 0.5M ü pulsar mass: M =1.97± 0.04M

11 The hyperon puzzle Hyperons è soft (or too soft) EoS not compatible (mainly in microscopic approaches) with measured (high) masses. However, the presence of hyperons in the NS interior seems to be unavoidable. ü can YN & YY interactions still solve it? ü or perhaps hyperonic three-body forces? ü what about quark matter?

12 BHF: Schulze & Rijken 2011 V low k : Dapo, Schafer & Wambach 2010 Microscopic approaches still unable to predict larger masses M max <1.4M BHF: I. V. et al Even hyperonic 3BF cannot solve the problem 1.27 < M max <1.6M See talks of D. Logoteta & K. Tsubakihara (parallel session VIII on Thursday)

13 Situation not much clear with phenomenological approaches (Massot et al. 2012) (Weissenborn et al. 2012) (Miyatsu et al. 2012) ü χ-lm & QMC ü Hartree-Fock M max = M ü RMF ü SU(6) SU(3) ü Vary z=g 8 /g 1, α v ü φ mesons M max compatible with 1.97M ü RHF & QMC ü π & f vb M max compatible with 1.97M ü RMF ü σ 4 terms ü σ *, φ mesons M max > 2M (Bednarek et al. 2012) See talk of J. Schaffner-Bielich (parallel session VIII on Thursday)

14 Question is so open that Hyperons-NS-2012 A task force meeting-november 21-24, 2012 Copernicus Astronomical Center Warsaw, Poland Organizers: M. Bejger, P. Haensel, J. Schaffner-Bielich & L. Zdunik

15 Hyperon Stars at birth

16 Proto-Neutron Stars New effects on PNS matter: Thermal effects T MeV S / A 1 2 Neutrino trapping µ ν 0 Y e = ρ e + ρ ν e ρ B 0.4 (Janka, Langanke, Marek, Martinez-Pinedo & Muller 2006) Y µ = ρ µ + ρ ν µ ρ B 0

17 Proto-Neutron Stars: Composition Neutrino free µ ν = 0 Neutrino trapped µ ν 0 (Burgio & Schulze 2011) (Burgio & Schulze 2011) Neutrino trapped ê ê ü ü ü ü Large proton fraction Small number of muons Onset of Σ - (Λ) shifted to higher (lower) density Hyperon fraction lower in ν-trapped matter

18 Proto-Neutron Stars: EoS (Burgio & Schulze 2011) Nucleonic matter ² ν-trapping + temperature softer EoS Hyperonic matter ² ν-trapping + temperature stiffer EoS ² More hyperon softening in ν-untrapped matter (larger hyperon fraction)

19 Proto-Neutron Stars: Structure (Burgio & Schulze 2011) Nucleonic matter ν-trapping + T: reduction of M max (I. V. et al. 2003) N,Y,l,ν Hyperonic matter ν-trapping + T: increase of M max delayed BH formation N,Y,l go to BH

20 Hyperons & Neutron Star Cooling

21 Neutron Star Cooling in a Nutshell Two cooling regimes Core cools by neutrino emission Crust cools by conduction Slow Low NS mass Fast High NS mass Surface photon emission dominates at t > 10 6 yrs slow cooling fast cooling de th dt = C v dt dt = L γ L ν + H ü C v : specific heat ü L γ : photon luminosity ü L ν : neutrino luminosity ü H: heating

22 Neutrino Emission in a Nutshell Name Process Emissivity Direct URCA Modified URCA Bremsstrahlung n p + l +ν l p + l n +ν l ~ T 6 Fast N + n N + p + l +ν l N + p + l N + n +ν l ~ T 8 Slow N + N N + N +ν +ν ~ T 8 Slow Cooper pair formation n + n nn p + p pp [ ] +ν +ν [ ] +ν +ν ~ T 7 Medium

23 Hyperonic DURCA processes possible as soon as hyperons appear (nucleonic DURCA requires x p > %) ê ê Additional Fast Cooling Processes Process R (Schaab, Shaffner-Bielich & Balberg 1998) Λ p + l +ν l Σ n + l +ν l Σ Λ + l +ν l Σ Σ 0 + l +ν l Ξ Λ + l +ν l Ξ Σ 0 + l +ν l Ξ 0 Σ + + l +ν l Ξ Ξ 0 + l +ν l partner reactions generating neutrinos, Hyperonic MURCA, N+Y only N R: relative emissitivy w.r.t. nucleonic DURCA

24 Pairing Gap suppression of C v & ε by ~ e ( Δ/k B T ) 1 S 0, 3 SD 1 ΣN & 1 S 0 ΛN gap 1 S 0 ΛΛ gap (Balberg & Barnea 1998) (Wang & Shen 2010) 1 S 0 ΣΣ gap NSC97e (Zhou, Schulze, Pan & Draayer 2005) (I. V. & Tolós 2004)

25 Hyperons & the R-mode instability of Neutron Stars

26 The R-mode Instability in a Napkin Ω Kepler : Absolute Upper Limit of Rot. Freq. Instabilities prevent NS to reach Ω Kepler R-mode Instability : toroidal mode of oscillation ü restoring force: Coriolis Ω c /Ω Kepler r-mode damped by shear viscosity r-mode unstable due to GW emission r-mode damped by bulk viscosity ü emission of GW in hot & rapidly rotating NS (CFS mechanism) GW makes the mode unstable Viscosity stabilizes the mode iω (Ω) t/τ (Ω,T ) A A 0 e 1 τ (Ω,T ) = 1 τ GW (Ω) + 1 τ Viscosity (Ω,T )

27 Hyperon Bulk Viscosity ξ Y (Lindblom et al. 2002, Haensel et al 2002, van Dalen et al. 2002, Chatterjee et al. 2008, Gusakov et al. 2008, Shina et al. 2009, Jha et al. 2010, ) Sources of ξ Y : (Haensel, Levenfish & Yakovlev 2002) non-leptonic weak reactions Direct & Modified URCA strong reactions N + N N +Y N +Y Y +Y Y B + l +ν l B'+Y B'+ B + l +ν l N +Y N +Y N +Ξ Y +Y Y +Y Y +Y Reaction Rates & ξ Y reduced by Hyperon Superfluidity

28 Critical Angular Velocity of Neutron Stars r-mode amplitude: 1 τ Ω,T ( ) = 1 ( Ω) + 1 ( Ω,T) + 1 ( T) τ GW iω (Ω)t t/τ (Ω) A A o e τ ξ τ η (I.V. & Albertus 2012) è 1 τ Ω c,t ( ) = 0 r-mode instability region Ω < Ω c Ω > Ω c stable unstable As expected: smaller r-mode instability region due to hyperons BHF: NN (Av18)+NY (NSC89) (M=1.27M )

29 Take away message Hyperons in Neutron Stars ü Strong softening of EoS & reduction of NS Mass è Hyperons & Massive NS still an open question (Hyperons-NS-2012 Meeting) ü Modification of PNS properties (composition, EoS, Mass) ü Additional Fast Cooling Processes ü Reduction of r-mode instability region

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