Cosmological and astrophysical probes of dark matter annihilation

Transcription

1 Cosmological and astrophysical probes of dark matter annihilation Institute for Cosmic Ray Research, University of Tokyo Kazunori Nakayama J.Hisano, M.Kawasaki, K.Kohri and KN, Phys.Rev.D79,063514(2009)[ ] J.Hisano, M.Kawasaki, K.Kohri and KN, Phys.Rev.D79,043516(2009)[ ] M.Kawasaki, K.Kohri and KN, to appear in Phys.Rev.D [ ] J.Hisano, KN, and M.J.S.Yang, Phys.Lett.B678,101(2009)[ ] ACP IPMU (2009/07/02)

2 Contents First part PAMELA/Fermi results & constraints on DM annihilation scenario Second part Neutrino signals from GC Diffuse gamma-ray background

3 First Part : Quick Summary

4 Energy content of the Universe after WMAP What is the dark matter? Can it be detected? Collider Direct detection Indirect detection DM-nucleon Scattering DM annihilation Cosmic Ray Signals

5 Detecting dark matter signal DM + DM e ±, γ, p, ν,... e ± γ ν We are here

6 PAMELA observation )) - )+!(e + ) / (!(e !(e Positron fraction Muller & Tang 1987 MASS 1989 TS93 HEAT94+95 CAPRICE94 AMS98 HEAT00 Clem & Evenson 2007 excess in cosmic-ray positron flux PAMELA Energy (GeV) Adriani et al.,arxiv:

7 ATIC/PPB-BETS observations 1,000 excess in electron+positron flux E e 3.0 dn/dee (m 2 s 1 sr 1 GeV 2 ) ATIC BETS, PPB-BETS HEAT AMS ,000 Energy (GeV) J.Chang et al. Nature (2008)

8 Fermi LAT collaboration, Fermi observation Inconsistent with ATIC results. Still there may be excess.

9 HESS observation sr -1 ) 2 m -2 s -1 dn/de (GeV 3 E 10 2! E ± 15% ATIC PPB-BETS Kobayashi H.E.S.S. H.E.S.S. - low-energy analysis Systematic error Systematic error - low-energy analysis Broken power-law fit Consistent with Fermi results Energy (GeV) HESS collaboration,

10 Dark Matter : Decay or Annihilate Decaying DM DM need not be completely stable. DM lifetime with can explain PAMELA. Annihilating DM τ sec Flux n DM τ cm 3 s 1 DM may have weak scale annihilation cross section. Cross section with σv cm 3 s 1 can explain PAMELA. Flux n 2 DM σv cm 3 s 1

11 Positron fraction Total flux [GeV 2 m 2 s 1 sr 1 ] χχ µ + µ : m χ = 1.2TeV, σv = cm 3 s 1

12 Positron fraction Total flux [GeV 2 m 2 s 1 sr 1 ] χχ τ + τ : m χ = 1.5TeV, σv = cm 3 s 1

13 It is important to investigate Relation to other signals Gamma-rays from Galactic center Diffuse extragalactic gamma-rays Anti-protons Synchrotron radiation Neutrinos from Galactic center γ ν e ±

14 Σv in cm 3 sec Σv in cm 3 sec Μ Μ GC VLT GC radio GC radio PAMELA DM mass in GeV Μ Μ Sgr dsph Γ GC Γ GR Γ Σv in cm 3 sec ρ DM (r) 1/r PAMELA DM mass in GeV GR Γ Sgr dsph Γ Σv in cm 3 sec ρ DM (r) 1/(1 + r 2 ) GC VLT GC radio Τ Τ PAMELA DM mass in GeV GC radio Sgr dsph Γ DM mass in GeV GC Γ GR Γ DM DM Τ Τ, isothermal profile PAMELA GC Γ GR Γ Sgr dsph Γ Constraints on DM ann models. Sensitively depends on DM halo profile. G.Bertone et al.,

15 What we have done : Neutrino-induced muon flux from Galactic center Diffuse extragalactic gamma-rays from dark matter annihilation Both give useful constraints on DM models, rather insensitive to DM halo profile J.Hisano, M.Kawasaki, K.Kohri and KN, arxiv: M.Kawasaki, K.Kohri and KN, arxiv: J.Hisano, KN, and M.J.S.Yang, arxiv:

16 Μ Μ Τ Τ GC VLT GC VLT Σv in cm 3 sec PAMELA Sgr dsph Γ GC Γ GR Γ Σv in cm 3 sec PAMELA Sgr dsph Γ GC Γ GR Γ Σv in cm 3 sec GC radio DM mass in GeV GC radio Μ Μ GR Γ PAMELA Sgr dsph Γ Σv in cm 3 sec GC radio DM mass in GeV GC radio DM DM Τ Τ, isothermal profile PAMELA GC Γ GR Γ Sgr dsph Γ DM mass in GeV DM mass in GeV G.Bertone et al.,

17 Μ Μ Τ Τ Σv in cm 3 sec GC VLT PAMELA Sgr dsph Γ GC Γ GR Γ ν Σv in cm 3 sec GC VLT PAMELA Sgr dsph Γ GC Γ GR Γ Σv in cm 3 sec GC radio DM mass in GeV GC radio Μ Μ GR Γ PAMELA Sgr dsph Γ Σv in cm 3 sec GC radio DM mass in GeV GC radio DM DM Τ Τ, isothermal profile PAMELA GC Γ GR Γ Sgr dsph Γ DM mass in GeV DM mass in GeV G.Bertone et al.,

18 Σv in cm 3 sec GC VLT Μ Μ PAMELA Sgr dsph Γ GC Γ GR Γ ν Σv in cm 3 sec GC VLT Τ Τ PAMELA Sgr dsph Γ GC Γ GR Γ diffuse γ Σv in cm 3 sec GC radio DM mass in GeV GC radio Μ Μ GR Γ PAMELA Sgr dsph Γ Σv in cm 3 sec GC radio DM mass in GeV GC radio DM DM Τ Τ, isothermal profile PAMELA GC Γ GR Γ Sgr dsph Γ DM mass in GeV DM mass in GeV G.Bertone et al.,

19 Second part : Neutrino Flux J.Hisano, M.Kawasaki, K.Kohri and KN, Phys.Rev.D79,043516(2009)[ ] J.Hisano, KN, and M.J.S.Yang, Phys.Lett.B678,101(2009)[ ]

20 Neutrino Signal from DM Annihilation χχ W + W, b b, l + l,... e ±, γ, p, ν,... Ritz, Seckel (88),Kamionkowski (90),... Bertone,Nezri,Orloff,Silk (04), Yuksel,Horiuchi,Beacom,Ando(07) Interaction inside the Earth ν detector µ from GC Earth Search for up-going muons Limits from Super-K

21 SK limit on upward muon flux from GC direction x Flux Limit (cm -2 sec -1 ) Baksan Limit IMB Limit Kamiokande Limit MACRO Limits Super-K Limits Cone Half Angle From Galactic Center (Degrees) S.Desai et al., Phys.Rev.D70, (2004)

22 Muon flux from DM N µ = de νµ df νµ de νµ f(e νµ ) (a) Neutrino flux from DM: µ df νµ de νµ SK detector (b) Probability of ν µ µ f(e νµ ) : ν Rock

23 (a) Neutrino flux from GC df νµ de νµ = R ρ 2 8πm 2 ( F dn (ν µ) F σv F de νµ ) J Ω Neutrino spectra : dn (ν µ) F de νµ = i ( dn (ν i) F P νi ν µ de νi ) E νi =E ν µ DM halo profile dependent part : J Ω = dω Ω l.o.s. Neutrino oscillation ( ) 2 dl(ψ) ρ(l) R ρ Typical value of J Ω J 2 Ω Ω NFW isothermal

24 (b) Probability of ν µ µ f(e νµ ) de µ dσ νµ p µx de µ n (rock) p R(E µ ) Cross section : Number density of proton in the rock : n (rock) p G2 F s π E ν µ = 1.3N A cm 3 ν µ µ W N N Muon range : R(E µ )

25 Energy loss of muon in matter Dutta, Reno, Sarcevic, Seckel, Phys.Rev.D63, (2001) de µ dx = α(e µ) β(e µ )E µ Ionization loss : α(e µ ) 2 MeVcm 2 g 1 Radiative loss : (Brems, pair creation,...) β(e µ ) 10 6 cm 2 g 1 Typical propagation distance : R µ E µ α(e µ )ρ rock 1 km(e µ /1TeV)

26 Muon range in the rock Propagation distance [km] 10 1 all loss ionization loss E Muon energy [GeV] E µ 1 TeV E µ 1 TeV R µ 1 km(e µ /1TeV) deviation from linearity

27 Probability of ν µ µ f(e νµ ) f(e νµ ) E 2 ν µ de µ dσ νµ p µx de µ n (rock) p R(E µ ) G2 F s π E ν µ E µ Higher energy neutrinos are more likely converted into muon Monochromatic neutrino : χχ ν ν is constrained more severely than secondary neutrino : χχ µ + µ, µ ν µ ν e e

28 Limits from SK : Annihilation into left-handed leptons is not favored. Annihilate into left handed leptons Annihilate into right handed leptons (ν ν + l L l+ R ) (l R l+ L ) J.Hisano, M.Kawasaki, K.Kohri, KN (2008)

29 ν µ ν µ Contour : Muon flux ( cm 2 s 1 ) σv J Ω [cm 3 s 1 ] Gray : current SK bound (for θ = 5 ) µ + µ Total muon flux N µ = de νµ df νµ de νµ f(e νµ ) σv J Ω [cm 3 s 1 ] m 2 χ N µ const.

30 ν µ ν µ Contour : Muon flux ( cm 2 s 1 ) σv J Ω [cm 3 s 1 ] PAMELA & Fermi Gray : current SK bound (for θ = 5 ) σv J Ω [cm 3 s 1 ] µ + µ PAMELA & Fermi Total muon flux N µ = de νµ df νµ de νµ f(e νµ ) m 2 χ N µ const.

31 Lesson from neutrino Construct a DM model which fits PAMELA/Fermi data (either ann or decay) Check if your model produce monochromatic neutrinos with similar rate or not If yes, your model may conflict with SK bound irrespective of DM density profile Check carefully the SK bound!

32 Τ Τ Σv in cm 3 sec PAMELA GR Γ Sgr dsph Γ DM mass in GeV G.Bertone et al.,

33 Τ Τ Σv in cm 3 sec PAMELA GR Γ Sgr dsph Γ DM mass in GeV G.Bertone et al., P.Maede et al.,

34 Τ Τ Σv in cm 3 sec PAMELA GR Γ Sgr dsph Γ DM mass in GeV G.Bertone et al., P.Maede et al., Neutrino constraint becomes standard.

35 Τ Τ Σv in cm 3 sec PAMELA GR Γ Sgr dsph Γ Caution : This figure considers only secondary neutrino from tau decay. If DM annihilates into line neutrino, constraint becomes DM mass more in GeVstringent G.Bertone et al., P.Maede et al., Neutrino constraint becomes standard.

36 Possible improvement at SK High-energy neutrino-induced muons are detected through Cherenkov light Energy of each muon is not measured However, SK can distinguish muon events by event shape : shower and non-shower Higher energy muons more likely observed as showering muon DM-originated neutrinos more likely produce shower events than atmospheric neutrinos

37 Simulation 3 kind of muon events : Showering Through-going shower mu dn/d(loge) 5000 Non-showering Stopping Through-going nonshower mu Stopping mu Probability for shower E! (GeV) S.Desai et al., Astropart.Phys.29,42 (2008)

38 Flux [GeV -1 cm -2 s -1 sr -1 ] atmos : Honda et al.,2005 DM atmos BG E ν [GeV]

39 atmos : Honda et al.,2005 Flux [GeV -1 cm -2 s -1 sr -1 ] Non-shower muon DM atmos BG Shower muon E ν [GeV]

40 atmos : Honda et al.,2005 Flux [GeV -1 cm -2 s -1 sr -1 ] Non-shower muon DM atmos BG Shower muon E ν [GeV] Shower muon events contain relatively large contribution from DM-produced neutrino

41 Ratio between Shower muon and Total muon atmospheric : N shower µ /N total µ 0.12

42 ν µ ν µ Contour : Muon flux ( cm 2 s 1 ) σv J Ω [cm 3 s 1 ] Gray : current SK bound (for θ = 5 ) µ + µ A factor improvement is expected on the annihilation cross section σv J Ω [cm 3 s 1 ] May soon reach to PAMELA/Fermi region? Hisano, KN, Yang

43 ν µ ν µ Contour : Muon flux ( cm 2 s 1 ) σv J Ω [cm 3 s 1 ] PAMELA & Fermi Gray : current SK bound (for θ = 5 ) σv J Ω [cm 3 s 1 ] µ + µ PAMELA & Fermi A factor improvement is expected on the annihilation cross section May soon reach to PAMELA/Fermi region? Hisano, KN, Yang

44 Comments on IceCube Huge detector High statistics Located at South Pole cannot see Galactic center through upward muons Use downward muons? Atmospheric muon BG is 10^6 larger than DM signal

45 A planned extension : DeepCore Primary purpose : better sensitivity on low-energy neutrino Inner detector with denser instrumentation Use original detector as muon veto Remove atmospheric muon BG S.Seo, Talk at Dark2009

46 Expected sensitivity of DeepCore (5yr) 5000 Μ Μ Boost Σ Limit PAMELA M Χ GeV Spolyar, Buckley, Freese, Hooper, Murayama,

47 Summary DM interpretation of PAMELA/Fermi Confirm/constrain by other signals Neutrino-induced muon Flux Useful constraints on annihilating/decaying DM. Shower/non-shower separation may be a useful way to extract DM information. SK III, DeepCore, KM3NeT

48 Second Part II : Diffuse gamma-rays M.Kawasaki, K.Kohri and KN, to appear in Phys.Rev.D [ ]

49 Gamma-Ray Flux

50 Gamma-Ray Flux Galactic center

51 Gamma-Ray Flux Galactic center Extra Galactic diffuse

52 Continuum Gamma-Rays from DM ann. Internal Brems. Final state charged particle always emit photon. χχ l + l χχ l + l γ χ χ e + e γ Cascade decay χχ τ + τ, W + W hadrons(π ±, π 0, ρ,... ) 2γ

53 Extra-Galactic component Ullio, Bergstrom, Edsjo, Lacey (2002) Dominant contribution is summation over the DM ann. in external clustering objects γ We are here

54 . [ dφγ de ] ext = σv 8π ρ 2 m m 2 χ dz(1 + z) 3 H(z) dn γ de 2 (z) 2 (z) : Enhancement factor ( 2 (z) = 1 : homogeneous DM) 2 (z) dmm dn(z) dm drρ 2 M (r) Number of clustering objects : Press-Schechter theory Press, Schechter (1974) Sheth, Mo, Tormen (2001) Universal DM halo profile (Moore, NFW,...)

55 Enhancement factor 2 (z) Moore! 2 (z) Moore NFW Burkert ρ(r) 1 r 1.5 (1 + r 1.5 ) NFW ρ(r) 1 r(1 + r) z ρ(r) Burkert 1 (1 + r)(1 + r 2 ) About 10^5-10^6 enhancement for DM annihilation rate

56 Gamma-rays from 10 < b < 90 Extragalactic component is comparable to Galactic component

57 e ± µ ± M.Kawasaki, K.Kohri, KN, τ ± W ±

58 Constraints on annihilation cross section Cuspy profile : extragalactic is weaker than GC bound Cored profile : extragalactic is stronger than GC bound Fermi will soon make the bound stronger

59 Effects of Inverse Compton scattering CMB photon e ± + γ CMB e ± + γ Profumo, Jeltema, Belikov, Hooper, τ + τ Second peak around E (IC) γ γ 2 ee CMB 0.1 GeV ( mdm 1 TeV ) 2 More stringent constraint

60 Summary DM interpretation of PAMELA/Fermi Confirm/constrain by other signals, as Neutrino-induced muon Flux Useful constraints on annihilating/decaying DM. SK III, DeepCore, KM3NeT Gamma-ray Flux Both Galactic and extra-galactic gamma-rays may be significant in DM ann scenario. Fermi

61 Back-up Slides

62 Astrophysical source Pulsar (single or sum) Gamma-ray burst Sum of all pulsars Hooper, Blasi, Serpico, arxiv: Profumo, arxiv: K.Ioka, arxiv: Geminga pulsar Hooper, Blasi, Serpico, arxiv:

63 DM or Astrophysics? Anisotropy in CR flux I min I max Point source δ = I max I min I max + I min K(E) Rc 0.01% ( 1 kpc R ) ( E ) GeV

64 Anisotropy from nearby Pulsar Fermi 5yr Hooper, Blasi, Serpico, arxiv:

65 !=2, MED diffusion setup, ST model B Vela Monogem (B ) Loop I Positron Fraction 0.1 Geminga Cygnus L Energy [GeV] S.Profumo, arxiv:

66 Press-Schechter theory dn dm = 2 π ρ m δ c M [ 1 dσ(m) σ(m) 2 dm ] exp [ ] δ2 c 2σ 2. (M) fraction of mass in halos heavier than M fraction of mass in each mass interval M : halo mass σ(m) : dispersion of density field δ c : critical overdensity σ 2 (M) = 1 2π 2 W 2 (k M )P (k)k 2 dk 10-2 z = 0 z = 2 z = Ullio et al. (2002) log 10 (M / M!! ) Predict number of collapsed objects with mass M