XLth Rencontres de Moriond, March 2005 Possible Dark Matter Signals from Antiprotons, Positrons, X-rays and Gamma-rays Ullrich Schwanke (Humboldt University, Berlin) Remember to change spectrum plots and add Dermer et al. XLth Rencontres de Moriond, March 2005

Overview Introduction: Signatures of Dark Matter (DM) Search for positron and antiproton signals The HEAT balloon experiment Gamma-ray Astronomy 511 keV annihilation line (Integral) Diffuse gamma-ray emission (EGRET) Gamma-rays from the Galactic centre (H.E.S.S.) Summary and Outlook

Precision Cosmology WMAP Excess of total matter density over baryonic matter density is strongest argument for DM. Experimental evidence: cosmic microwave background (e.g. WMAP) Distance-luminosity relation for supernovae Primordial nucleosynthesis Galaxy distribution WMAP

Dark Matter Searches What is the exact nature of dark matter ? (mass, quantum numbers, couplings, spatial distribution) Direct searches look for interactions of DM particles with matter. Collider experiments spin-(in)dependent scattering with target nuclei, record transferred energy, direction of nucleus Controlled experimental environment. Covered by later talks. Indirect searches look for secondaries: annihilation products of DM particles Reasonable candidates: Antiprotons Positrons Gammas Neutrinos This talk

Antiprotons, Positrons and Gammas Extraterrestrial sources. Detection in orbit/atmosphere. Potentially large amount of DM (~entire Milky Way). Competition from less exotic production mechanisms Modelling of Milky Way required. Antiprotons Propagation effects Expect energy spectrum with cut-off at mass of DM particle Positrons Similar to antiprotons, lower range Gammas Directional information can be correlated with (dark) matter density in the Milky Way Gamma-line(s) would be unique signature. GLAST Simulation 

Search for Antiprotons and Positrons Historic claims for a sizable fraction of positrons/antiprotons in the cosmic radiation Experimental challenge: small fraction of e+/p-, wealth of background with opposite charge Good particle ID required 1987 Im unteren Energiebereich kann man Gammastrahlen am besten oberhalb der Erdatmosphäre direkt mit kleinen Teilchendetektoren auf Satelliten messen. Als Beispiel ist hier das Compton Gamma Ray Observatory dargestellt, das uns eine Fülle neuer Erkenntnisse beschert hat. Die Detektortechnologie ist restlos von den Elementarteilchenphysikern übernommen, die ebenfalls Photonen und andere elementaren Teilchen an Beschleunigerexperimenten vermessen müssen. Die Satellitentechnologie stößt allerdings bei Photonenergien von ca. 10 GeV an eine unüberwindliche Schranke. Die Detektoren müssten oberhalb dieser Grenze prohibitiv groß werden, da der Fluss von so hochenergetischen Gammastrahlen aus dem Weltall mit wachsender Energie sehr schnell kleiner wird. BESS, CAPRICE, High-Energy Antimatter Telescope, ... HEAT BESS

HEAT-e and HEAT-pbar Two flights: 1994 and 1995 One flight: 2000 Im unteren Energiebereich kann man Gammastrahlen am besten oberhalb der Erdatmosphäre direkt mit kleinen Teilchendetektoren auf Satelliten messen. Als Beispiel ist hier das Compton Gamma Ray Observatory dargestellt, das uns eine Fülle neuer Erkenntnisse beschert hat. Die Detektortechnologie ist restlos von den Elementarteilchenphysikern übernommen, die ebenfalls Photonen und andere elementaren Teilchen an Beschleunigerexperimenten vermessen müssen. Die Satellitentechnologie stößt allerdings bei Photonenergien von ca. 10 GeV an eine unüberwindliche Schranke. Die Detektoren müssten oberhalb dieser Grenze prohibitiv groß werden, da der Fluss von so hochenergetischen Gammastrahlen aus dem Weltall mit wachsender Energie sehr schnell kleiner wird. Two flights: 1994 and 1995 One flight: 2000

Positron Fraction Confirmed by two different instruments (HEAT-e and HEAT- pbar) Near solar maximum (1995 and 1995) and solar minimum (2000) Different vertical geomagnetic cutoffs: ~1 GeV (1995) and ~4 GeV (1994, 2000) 1987 Im unteren Energiebereich kann man Gammastrahlen am besten oberhalb der Erdatmosphäre direkt mit kleinen Teilchendetektoren auf Satelliten messen. Als Beispiel ist hier das Compton Gamma Ray Observatory dargestellt, das uns eine Fülle neuer Erkenntnisse beschert hat. Die Detektortechnologie ist restlos von den Elementarteilchenphysikern übernommen, die ebenfalls Photonen und andere elementaren Teilchen an Beschleunigerexperimenten vermessen müssen. Die Satellitentechnologie stößt allerdings bei Photonenergien von ca. 10 GeV an eine unüberwindliche Schranke. Die Detektoren müssten oberhalb dieser Grenze prohibitiv groß werden, da der Fluss von so hochenergetischen Gammastrahlen aus dem Weltall mit wachsender Energie sehr schnell kleiner wird.

Interpretation of the Positron Fraction Neutralino DM inefficient generation of positrons increase annihilation rate by clumping Kaluza-Klein Dark Matter viable positron source for mass range 300..400 GeV e+ diffusion parameters D. Hooper, hep-ph/0409272 (Annihilation rate normalized to data)

Antiproton Fraction and Flux 1987 Im unteren Energiebereich kann man Gammastrahlen am besten oberhalb der Erdatmosphäre direkt mit kleinen Teilchendetektoren auf Satelliten messen. Als Beispiel ist hier das Compton Gamma Ray Observatory dargestellt, das uns eine Fülle neuer Erkenntnisse beschert hat. Die Detektortechnologie ist restlos von den Elementarteilchenphysikern übernommen, die ebenfalls Photonen und andere elementaren Teilchen an Beschleunigerexperimenten vermessen müssen. Die Satellitentechnologie stößt allerdings bei Photonenergien von ca. 10 GeV an eine unüberwindliche Schranke. Die Detektoren müssten oberhalb dieser Grenze prohibitiv groß werden, da der Fluss von so hochenergetischen Gammastrahlen aus dem Weltall mit wachsender Energie sehr schnell kleiner wird. Some claimed excesses in the past Measurements seem to be consistent with purely secondary production of antiprotons Primary antiproton flux from annihilation of a 964 GeV MSSM neutralino (P. Ullio, astro-ph/9904086 (1999))

Outlook PAMELA (launch ~2005) Space-bore experiments (AMS 02, PAMELA) will allow for much more stringent searches Much better duty cycle than balloon experiments Impact of solar environment can be studied in greater detail

X-Rays and Gamma-Rays Integral Soft g-rays: < 1 MeV Very high energy -rays: > 100 GeV Air-Cherenkov Telescopes H.E.S.S. Whipple/Veritas MAGIC CANGAROO High energy g-rays: 10 MeV – 100 GeV EGRET, GLAST The atmosphere is opaque to gamma-rays so satellites needed – except at very high energies were the particle cascades initiated by gamma- rays are sufficiently energetic to generate Cherenkov light (~5 GeV)

Galactic 511 keV Annihilation Line Accurate tracer of galactic positrons. Thermalization of positrons required. Various detections since initial discovery in 1973. Agreement on absolut flux, no time dependence Morphology less clear (halo + galactic disk component, galactic positron fountain?) e+e- Instrument Year Flux (10-3 cm-2 s-1) Centroid (keV) Width (keV) HEAO-3 79-80 1.130.13 510.920.23 1.6+0.9-1.6 GRIS 88 and 92 0.880.07 2.50.4 HEXAGONE 89 1.000.24 511.330.41 2.90+1.10-1.01 TGRS 95-97 1.070.05 510.980.10 1.810.54

New Data: Integral and SPI launched in Oct 02 SPectromètre Integral 16° FoV (FWHM) 20 keV – 10 MeV 2 keV energy resolution (at 1 MeV) 2° angular resolution

Observations of the Galactic Centre 12  Data not released yet Flux Energy (keV) Gaussian Model (10° FWHM) Measurement relies on accurate subtraction of instrumental annihilation line Flux and intrinsic line width compatible with earlier mesurements Azimuthally symmetric galactic bulge component with FWHM=9° centred at GC Rate Galactic longitude (°)

Interpretation and Outlook Dark Matter Interpretation Light DM particles (1-100 MeV) Agrees with DM relic density Rather flat halo Other Interpretations Supernovae Wolf-Rayet Stars Neutron stars, pulsars Cosmic rays ...and (of course) Black holes Will more data (better morphology) really help? Flux()/Flux(0) C. Boehm et al., astro-ph/0309686

X-Rays and Gamma-Rays Integral Soft g-rays: < 1 MeV Very high energy -rays: > 100 GeV Air-Cherenkov Telescopes H.E.S.S. Whipple/Veritas MAGIC CANGAROO High energy g-rays: 10 MeV – 100 GeV EGRET, GLAST The atmosphere is opaque to gamma-rays so satellites needed – except at very high energies were the particle cascades initiated by gamma- rays are sufficiently energetic to generate Cherenkov light (~5 GeV)

Diffuse Gamma-Ray Emission CGRO (1991-2000) EGRET 20 MeV – 30 GeV energy resolution 20% angular resolution: 1.3° at 1 GeV 0.4° at 10 GeV

EGRET Gamma-Ray Data Subtraction of 271 EGRET point sources  Diffuse gamma-ray emission remains Right now, EGRET data (and more) can be described by scenarios with and without DM S. D. Hunter et al. Astrophys. J. 481, 205 (1997) Solution without DM: Strong, Moskalenko & Reimer, Astrophys. J. 613, 962 (2004) Solution with DM: W. de Boer, hep-ph/0408166 (2004); W. de Boer, Herold, Sander & Zhukov, hep-ph/0408166 (2004)  See W. de Boer‘s Talk tomorrow

1) Solution without Dark Matter (30.5°<l<179.5°, 180.5°<l<330.5°) 0 decay 1.0-2.0 GeV Inverse Compton Bremsstrahlung Extragalactic Gamma-Ray Background GALPROP: Numeric evaluation of Diffusion-Loss-Equations. Input: B/C (to fix proton diffusion), local cosmic ray spectra, measured distributions of atomic, molecular and ionized H. Describes (anti)proton and electron/positron data, too.

2) Solution with Dark Matter Explains EGRET data with a photon component from neutralino annihilation Sets limit on WIMP mass in 50-100 GeV range Determines halo structure (isothermal halo i.e. not cuspy) DM signal compatible with supersymmetry for boost factors of ~20 (-30°<l<+30°) E>0.5 GeV Neutralino annihilation See W. de Boer‘s Talk Backgrounds

X-Rays and Gamma-Rays Integral Soft g-rays: < 1 MeV Very high energy -rays: > 100 GeV Air-Cherenkov Telescopes H.E.S.S. Whipple/Veritas MAGIC CANGAROO High energy g-rays: 10 MeV – 100 GeV EGRET, GLAST The atmosphere is opaque to gamma-rays so satellites needed – except at very high energies were the particle cascades initiated by gamma- rays are sufficiently energetic to generate Cherenkov light (~5 GeV)

Ground-based g-ray Observatories VERITAS (10/2006) MAGIC (08/2004) Čerenkov-Teleskope haben sich im letzten Jahrzehnt zur technischen Reife entwickelt. Zur Zeit werden weltweit vier neue Observatorien auf- bzw. ausgebaut, die die Beobachtung des Südhimmels (mit dem galaktischen Zentrum) und des Nordhimmels erlauben. Da die Teleskope nur in mondlosen Nächten beobachten können (sonst werden die Schauersignale von gewöhnlichem Licht überstrahlt), hat man weiterhin darauf geachtet, die Teleskope in Ost-West-Richtung so zu arrangieren, dass fast 24 Stunden am Tag beobachtet werden kann. Dies ist sehr wichtig, denn die Quellen der kosmischen Strahlen, wie z.B. Röntgenbinäre oder Blasare, haben sich als sehr variabel herausgestellt, was uns tiefe Einblicke in die innere Dynamik dieser Objekte gibt. H.E.S.S. (12/2003) CANGAROO III (03/2004)

The Imaging Cherenkov Technique Focal Plane ~ 10 km Particle Shower 5 nsec Intensity  Shower Energy At 100 GeV ~ 10 Photons/m2 (300 – 600 nm) ~ 120 m Image Orientation  Shower Direction Image Shape  Primary Particle

Stereoscopic Imaging Intersection of image axes gives precise shower direction

Performance The Crab Nebula Duty cycle: 1000h per year Trigger threshold: 40 – 100 GeV Angular resolution is a few arcminutes (~0.1°, stereo) Collection area: 50000 m2 Relative energy resolution ~20% Factor 102 improved sensitivity EGRET H.E.S.S. 30 sec 1 night 1 year Cas A 2002 Crab 1989 H.E.S.S. 2004

Observations of the Galactic Centre H.E.S.S. Field of View (5°)

The Dynamical Centre: Sgr A* 3  106 solar mass black hole Very low luminosity Highly variable non-thermal emission in IR and X-ray Extremely compact source < 0.1 milliarcseconds in mm. Surrounded by supernova- remnant Sgr A East and H II region Sgr A West Sgr A East Sgr A* 3‘ MPE / R. Genzel et al.

H.E.S.S. Result (2003) 17 hours of data Taken with 2 telescopes during construction of the array 160 GeV threshold 11 signal from close to Sgr A* Point-like source See A&A 425, L13-16 (2004)

Position Contours from Hooper et al. 2004 Chandra GC survey NASA/UMass/D.Wang et al. Chandra GC survey NASA/UMass/D.Wang et al. CANGAROO (80%) CANGAROO (80%) H.E.S.S. (95%) H.E.S.S. Whipple (95%) Whipple (95%) Contours from Hooper et al. 2004

Position: Compatible with Sgr-A* H.E.S.S. “HESS J1745-290” 95% 68% Chandra F. Banagoff et al.

Energy Spectrum HESS: dN/dE  E-2.2 Flux > 160 GeV: 5 % of Crab flux CANGAROO: dN/dE  E-4.6 ~ 1 Crab

H.E.S.S 2004 Data 50 h of data with full 4 telescope array Significance of HESS J1745-290 is 35 s Position, flux and spectrum compatible New source detected in the same field of view

Interpretations of the TeV Signal from the Galatic Centre Particle Acceleration near the Black Hole Sgr A*: F. Aharonian & A. Neronov, astro-ph/0408303 (2004); Atoyan & Dermer, astro-ph/0401243 (2004). Particle Acceleration in the supernova remnant Sgr A East: Crocker et al. astro-ph/0408183 (2004) Dark Matter Annihilation: D. Horns, astro-ph/0408192; Bergström et al., astro-ph/0410359

1) Particle Acceleration close to Sgr A* Low luminosity of Sgr A*  ~10 TeV photons can escape It has been suggested that Sgr A* is spinning at a good fraction of the maximum possible speed. Rotation in a magnetic field produces a huge electro-magnetic field Acceleration of protons to 1018 eV (?) VHE gamma-rays via curvature radiation or hadronic interactions Acceleration of electrons (?) TeV Gamma-rays via Inverse Compton Scattering More efficient than proton acceleration Or acceleration at shocks in the accretion disk TeV radiation via: p + p  p+/-, p0  gg

VHE g-rays from Sgr A* ? Log E (eV) Aharonian et al. 2004 Log E (eV) Data can be explained as radiation of accelerated protons… or electrons close (<10 Rg) to Sgr A* Need simultaneous X-ray data to test

2) Particle Acceleration in Sgr A East Spectral index measured by H.E.S.S. close to expectation from Fermi acceleration Sgr A East is a powerful SNR 10,000 years old Compact (~3 arcmins) Energy: 4 x 1052 erg Crocker et al. explain overabundance of cosmic rays from the GC around 1018 eV Flux normalization from H.E.S.S. (or a nearby EGRET source) under the assumption of pp induced p0 decay Explains particle acceleration up to the ankle (3 1018 eV)

Association with CR Anisotropy? EGRET p+p  0+X  n+X Log (dF/dE / cm-2 s-2 eV-1) Fit H.E.S.S. AGASA (1018 eV) Log (E/eV) Crocker et al 2004, astro-ph/0408183

3) DM Interpretation: Spectrum CANGAROO Spectrum consistent with a 1.1 TeV neutralino-type WIMP HESS Spectrum requires a mass > 12 TeV Most models favour a < 2 TeV WIMP Requires high DM density and/or cross section Kaluza-Klein DM requires large boost factors (>103) DM interpretation cannot be ruled out Wimp annihilation spectra have a cutoff at ~(0.2…0.3) M

DM Interpretation: Morphology Morpholgy not constrained (yet) by current H.E.S.S. Data Data favour a steep cuspy dark matter profile (well, for 100% DM) =1.1 =1.0 With better statistics, DM contribution might be separable from (then recognised) ordinary sources

Summary and Outlook For antiprotons and positrons, future space- borne experiments will do a lot better than balloon experiments. 511 keV line: Interpretation? GLAST (5/2007) will provide improved sensitivity for E<100 GeV Search for gamma-lines and continuum. Very high-energy gamma-rays Better cross-calibration of experiments. Multi-wavelength campaigns. Extend spectrum to higher energies, improve source localization and understanding of Galactic Centre region. Observation of other DM candidates (e.g. dwarf galaxies orbiting the Milky Way) GLAST