First detection of VHE \gamma–rays from SN 1006 by H.E.S.S.

First detection of VHE –rays from SN 1006 by H.E.S.S.

HESS Collaboration    F. Acero Laboratoire de Physique Théorique et Astroparticules, Université Montpellier 2, CNRS/IN2P3, CC 70, Place Eugène Bataillon, F-34095 Montpellier Cedex 5, France    F. Aharonian Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany Dublin Institute for Advanced Studies, 5 Merrion Square, Dublin 2, Ireland    A.G. Akhperjanian Yerevan Physics Institute, 2 Alikhanian Brothers St., 375036 Yerevan, Armenia    G. Anton Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    U. Barres de Almeida University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K. supported by CAPES Foundation, Ministry of Education of Brazil    A.R. Bazer-Bachi Centre d’Etude Spatiale des Rayonnements, CNRS/UPS, 9 av. du Colonel Roche, BP 4346, F-31029 Toulouse Cedex 4, France    Y. Becherini Astroparticule et Cosmologie (APC), CNRS, Universite Paris 7 Denis Diderot, 10, rue Alice Domon et Leonie Duquet, F-75205 Paris Cedex 13, France thanks: UMR 7164 (CNRS, Université Paris VII, CEA, Observatoire de Paris)    B. Behera Landessternwarte, Universität Heidelberg, Königstuhl, D 69117 Heidelberg, Germany    M. Beilicke Universität Hamburg, Institut für Experimentalphysik, Luruper Chaussee 149, D 22761 Hamburg, Germany now at Washington University, St. Louis, USA    K. Bernlöhr Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany Institut für Physik, Humboldt-Universität zu Berlin, Newtonstr. 15, D 12489 Berlin, Germany    A. Bochow Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    C. Boisson LUTH, Observatoire de Paris, CNRS, Université Paris Diderot, 5 Place Jules Janssen, 92190 Meudon, France    J. Bolmont LPNHE, Université Pierre et Marie Curie Paris 6, Université Denis Diderot Paris 7, CNRS/IN2P3, 4 Place Jussieu, F-75252, Paris Cedex 5, France    V. Borrel Centre d’Etude Spatiale des Rayonnements, CNRS/UPS, 9 av. du Colonel Roche, BP 4346, F-31029 Toulouse Cedex 4, France    J. Brucker Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    F. Brun LPNHE, Université Pierre et Marie Curie Paris 6, Université Denis Diderot Paris 7, CNRS/IN2P3, 4 Place Jussieu, F-75252, Paris Cedex 5, France    P. Brun IRFU/DSM/CEA, CE Saclay, F-91191 Gif-sur-Yvette, Cedex, France    R. Bühler Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    T. Bulik Astronomical Observatory, The University of Warsaw, Al. Ujazdowskie 4, 00-478 Warsaw, Poland    I. Büsching Unit for Space Physics, North-West University, Potchefstroom 2520, South Africa    T. Boutelier Laboratoire d’Astrophysique de Grenoble, INSU/CNRS, Université Joseph Fourier, BP 53, F-38041 Grenoble Cedex 9, France    P.M. Chadwick University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K.    A. Charbonnier LPNHE, Université Pierre et Marie Curie Paris 6, Université Denis Diderot Paris 7, CNRS/IN2P3, 4 Place Jussieu, F-75252, Paris Cedex 5, France    R.C.G. Chaves Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    A. Cheesebrough University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K.    J. Conrad Oskar Klein Centre, Department of Physics, Stockholm University, Albanova University Center, SE-10691 Stockholm, Sweden    L.-M. Chounet Laboratoire Leprince-Ringuet, Ecole Polytechnique, CNRS/IN2P3, F-91128 Palaiseau, France    A.C. Clapson Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    G. Coignet Laboratoire d’Annecy-le-Vieux de Physique des Particules, Université de Savoie, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France    M. Dalton Institut für Physik, Humboldt-Universität zu Berlin, Newtonstr. 15, D 12489 Berlin, Germany    M.K. Daniel University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K.    I.D. Davids University of Namibia, Private Bag 13301, Windhoek, Namibia Unit for Space Physics, North-West University, Potchefstroom 2520, South Africa    B. Degrange Laboratoire Leprince-Ringuet, Ecole Polytechnique, CNRS/IN2P3, F-91128 Palaiseau, France    C. Deil Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    H.J. Dickinson University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K.    A. Djannati-Ataï Astroparticule et Cosmologie (APC), CNRS, Universite Paris 7 Denis Diderot, 10, rue Alice Domon et Leonie Duquet, F-75205 Paris Cedex 13, France thanks: UMR 7164 (CNRS, Université Paris VII, CEA, Observatoire de Paris)    W. Domainko Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    L.O’C. Drury Dublin Institute for Advanced Studies, 5 Merrion Square, Dublin 2, Ireland    F. Dubois Laboratoire d’Annecy-le-Vieux de Physique des Particules, Université de Savoie, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France    G. Dubus Laboratoire d’Astrophysique de Grenoble, INSU/CNRS, Université Joseph Fourier, BP 53, F-38041 Grenoble Cedex 9, France    J. Dyks Nicolaus Copernicus Astronomical Center, ul. Bartycka 18, 00-716 Warsaw, Poland    M. Dyrda Instytut Fizyki Ja̧drowej PAN, ul. Radzikowskiego 152, 31-342 Kraków, Poland    K. Egberts Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany Institut für Astro- und Teilchenphysik, Leopold-Franzens-Universität Innsbruck, A-6020 Innsbruck, Austria    P. Eger Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    P. Espigat Astroparticule et Cosmologie (APC), CNRS, Universite Paris 7 Denis Diderot, 10, rue Alice Domon et Leonie Duquet, F-75205 Paris Cedex 13, France thanks: UMR 7164 (CNRS, Université Paris VII, CEA, Observatoire de Paris)    L. Fallon Dublin Institute for Advanced Studies, 5 Merrion Square, Dublin 2, Ireland    C. Farnier Laboratoire de Physique Théorique et Astroparticules, Université Montpellier 2, CNRS/IN2P3, CC 70, Place Eugène Bataillon, F-34095 Montpellier Cedex 5, France    S. Fegan Laboratoire Leprince-Ringuet, Ecole Polytechnique, CNRS/IN2P3, F-91128 Palaiseau, France    F. Feinstein Laboratoire de Physique Théorique et Astroparticules, Université Montpellier 2, CNRS/IN2P3, CC 70, Place Eugène Bataillon, F-34095 Montpellier Cedex 5, France    A. Fiasson Laboratoire d’Annecy-le-Vieux de Physique des Particules, Université de Savoie, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France    A. Förster Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    G. Fontaine Laboratoire Leprince-Ringuet, Ecole Polytechnique, CNRS/IN2P3, F-91128 Palaiseau, France    M. Füßling Institut für Physik, Humboldt-Universität zu Berlin, Newtonstr. 15, D 12489 Berlin, Germany    S. Gabici Dublin Institute for Advanced Studies, 5 Merrion Square, Dublin 2, Ireland    Y.A. Gallant Laboratoire de Physique Théorique et Astroparticules, Université Montpellier 2, CNRS/IN2P3, CC 70, Place Eugène Bataillon, F-34095 Montpellier Cedex 5, France    L. Gérard Astroparticule et Cosmologie (APC), CNRS, Universite Paris 7 Denis Diderot, 10, rue Alice Domon et Leonie Duquet, F-75205 Paris Cedex 13, France thanks: UMR 7164 (CNRS, Université Paris VII, CEA, Observatoire de Paris)    D. Gerbig Institut für Theoretische Physik, Lehrstuhl IV: Weltraum und Astrophysik, Ruhr-Universität Bochum, D 44780 Bochum, Germany    B. Giebels Laboratoire Leprince-Ringuet, Ecole Polytechnique, CNRS/IN2P3, F-91128 Palaiseau, France    J.F. Glicenstein IRFU/DSM/CEA, CE Saclay, F-91191 Gif-sur-Yvette, Cedex, France    B. Glück Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    P. Goret IRFU/DSM/CEA, CE Saclay, F-91191 Gif-sur-Yvette, Cedex, France    D. Göring Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    D. Hauser Landessternwarte, Universität Heidelberg, Königstuhl, D 69117 Heidelberg, Germany    M. Hauser Landessternwarte, Universität Heidelberg, Königstuhl, D 69117 Heidelberg, Germany    S. Heinz Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    G. Heinzelmann Universität Hamburg, Institut für Experimentalphysik, Luruper Chaussee 149, D 22761 Hamburg, Germany    G. Henri Laboratoire d’Astrophysique de Grenoble, INSU/CNRS, Université Joseph Fourier, BP 53, F-38041 Grenoble Cedex 9, France    G. Hermann Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    J.A. Hinton School of Physics & Astronomy, University of Leeds, Leeds LS2 9JT, UK    A. Hoffmann Institut für Astronomie und Astrophysik, Universität Tübingen, Sand 1, D 72076 Tübingen, Germany    W. Hofmann Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    P. Hofverberg Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    M. Holleran Unit for Space Physics, North-West University, Potchefstroom 2520, South Africa    S. Hoppe Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    D. Horns Universität Hamburg, Institut für Experimentalphysik, Luruper Chaussee 149, D 22761 Hamburg, Germany    A. Jacholkowska LPNHE, Université Pierre et Marie Curie Paris 6, Université Denis Diderot Paris 7, CNRS/IN2P3, 4 Place Jussieu, F-75252, Paris Cedex 5, France    O.C. de Jager Unit for Space Physics, North-West University, Potchefstroom 2520, South Africa    C. Jahn Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    I. Jung Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    K. Katarzyński Toruń Centre for Astronomy, Nicolaus Copernicus University, ul. Gagarina 11, 87-100 Toruń, Poland    U. Katz Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    S. Kaufmann Landessternwarte, Universität Heidelberg, Königstuhl, D 69117 Heidelberg, Germany    M. Kerschhaggl Institut für Physik, Humboldt-Universität zu Berlin, Newtonstr. 15, D 12489 Berlin, Germany    D. Khangulyan Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    B. Khélifi Laboratoire Leprince-Ringuet, Ecole Polytechnique, CNRS/IN2P3, F-91128 Palaiseau, France    D. Keogh University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K.    D. Klochkov Institut für Astronomie und Astrophysik, Universität Tübingen, Sand 1, D 72076 Tübingen, Germany    W. Kluźniak Nicolaus Copernicus Astronomical Center, ul. Bartycka 18, 00-716 Warsaw, Poland    T. Kneiske Universität Hamburg, Institut für Experimentalphysik, Luruper Chaussee 149, D 22761 Hamburg, Germany    Nu. Komin IRFU/DSM/CEA, CE Saclay, F-91191 Gif-sur-Yvette, Cedex, France    K. Kosack IRFU/DSM/CEA, CE Saclay, F-91191 Gif-sur-Yvette, Cedex, France    R. Kossakowski Laboratoire d’Annecy-le-Vieux de Physique des Particules, Université de Savoie, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France    G. Lamanna Laboratoire d’Annecy-le-Vieux de Physique des Particules, Université de Savoie, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France    M. Lemoine-Goumard Université Bordeaux 1; CNRS/IN2P3; Centre d’Etudes Nucléaires de Bordeaux Gradignan, UMR 5797, Chemin du Solarium, BP120, 33175 Gradignan, France    J.-P. Lenain LUTH, Observatoire de Paris, CNRS, Université Paris Diderot, 5 Place Jules Janssen, 92190 Meudon, France    T. Lohse Institut für Physik, Humboldt-Universität zu Berlin, Newtonstr. 15, D 12489 Berlin, Germany    V. Marandon Astroparticule et Cosmologie (APC), CNRS, Universite Paris 7 Denis Diderot, 10, rue Alice Domon et Leonie Duquet, F-75205 Paris Cedex 13, France thanks: UMR 7164 (CNRS, Université Paris VII, CEA, Observatoire de Paris)    A. Marcowith Laboratoire de Physique Théorique et Astroparticules, Université Montpellier 2, CNRS/IN2P3, CC 70, Place Eugène Bataillon, F-34095 Montpellier Cedex 5, France    J. Masbou Laboratoire d’Annecy-le-Vieux de Physique des Particules, Université de Savoie, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France    D. Maurin LPNHE, Université Pierre et Marie Curie Paris 6, Université Denis Diderot Paris 7, CNRS/IN2P3, 4 Place Jussieu, F-75252, Paris Cedex 5, France    T.J.L. McComb University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K.    M.C. Medina LUTH, Observatoire de Paris, CNRS, Université Paris Diderot, 5 Place Jules Janssen, 92190 Meudon, France    J. Méhault Laboratoire de Physique Théorique et Astroparticules, Université Montpellier 2, CNRS/IN2P3, CC 70, Place Eugène Bataillon, F-34095 Montpellier Cedex 5, France    R. Moderski Nicolaus Copernicus Astronomical Center, ul. Bartycka 18, 00-716 Warsaw, Poland    E. Moulin IRFU/DSM/CEA, CE Saclay, F-91191 Gif-sur-Yvette, Cedex, France    M. Naumann-Godo Laboratoire Leprince-Ringuet, Ecole Polytechnique, CNRS/IN2P3, F-91128 Palaiseau, France    M. de Naurois LPNHE, Université Pierre et Marie Curie Paris 6, Université Denis Diderot Paris 7, CNRS/IN2P3, 4 Place Jussieu, F-75252, Paris Cedex 5, France    D. Nedbal Charles University, Faculty of Mathematics and Physics, Institute of Particle and Nuclear Physics, V Holešovičkách 2, 18000 Prague 8, Czech Republic    D. Nekrassov Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    B. Nicholas School of Chemistry & Physics, University of Adelaide, Adelaide 5005, Australia    J. Niemiec Instytut Fizyki Ja̧drowej PAN, ul. Radzikowskiego 152, 31-342 Kraków, Poland    S.J. Nolan University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K.    S. Ohm Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    J-F. Olive Centre d’Etude Spatiale des Rayonnements, CNRS/UPS, 9 av. du Colonel Roche, BP 4346, F-31029 Toulouse Cedex 4, France    E. de Oña Wilhelmi Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    K.J. Orford University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K.    M. Ostrowski Obserwatorium Astronomiczne, Uniwersytet Jagielloński, ul. Orla 171, 30-244 Kraków, Poland    M. Panter Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    M. Paz Arribas Institut für Physik, Humboldt-Universität zu Berlin, Newtonstr. 15, D 12489 Berlin, Germany    G. Pedaletti Landessternwarte, Universität Heidelberg, Königstuhl, D 69117 Heidelberg, Germany    G. Pelletier Laboratoire d’Astrophysique de Grenoble, INSU/CNRS, Université Joseph Fourier, BP 53, F-38041 Grenoble Cedex 9, France    P.-O. Petrucci Laboratoire d’Astrophysique de Grenoble, INSU/CNRS, Université Joseph Fourier, BP 53, F-38041 Grenoble Cedex 9, France    S. Pita Astroparticule et Cosmologie (APC), CNRS, Universite Paris 7 Denis Diderot, 10, rue Alice Domon et Leonie Duquet, F-75205 Paris Cedex 13, France thanks: UMR 7164 (CNRS, Université Paris VII, CEA, Observatoire de Paris)    G. Pühlhofer Institut für Astronomie und Astrophysik, Universität Tübingen, Sand 1, D 72076 Tübingen, Germany    M. Punch Astroparticule et Cosmologie (APC), CNRS, Universite Paris 7 Denis Diderot, 10, rue Alice Domon et Leonie Duquet, F-75205 Paris Cedex 13, France thanks: UMR 7164 (CNRS, Université Paris VII, CEA, Observatoire de Paris)    A. Quirrenbach Landessternwarte, Universität Heidelberg, Königstuhl, D 69117 Heidelberg, Germany    B.C. Raubenheimer Unit for Space Physics, North-West University, Potchefstroom 2520, South Africa    M. Raue Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany European Associated Laboratory for Gamma-Ray Astronomy, jointly supported by CNRS and MPG    S.M. Rayner University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K.    O. Reimer Institut für Astro- und Teilchenphysik, Leopold-Franzens-Universität Innsbruck, A-6020 Innsbruck, Austria    M. Renaud Astroparticule et Cosmologie (APC), CNRS, Universite Paris 7 Denis Diderot, 10, rue Alice Domon et Leonie Duquet, F-75205 Paris Cedex 13, France thanks: UMR 7164 (CNRS, Université Paris VII, CEA, Observatoire de Paris)    R. de los Reyes Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    F. Rieger Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany European Associated Laboratory for Gamma-Ray Astronomy, jointly supported by CNRS and MPG    J. Ripken Oskar Klein Centre, Department of Physics, Stockholm University, Albanova University Center, SE-10691 Stockholm, Sweden    L. Rob Charles University, Faculty of Mathematics and Physics, Institute of Particle and Nuclear Physics, V Holešovičkách 2, 18000 Prague 8, Czech Republic    S. Rosier-Lees Laboratoire d’Annecy-le-Vieux de Physique des Particules, Université de Savoie, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France    G. Rowell School of Chemistry & Physics, University of Adelaide, Adelaide 5005, Australia    B. Rudak Nicolaus Copernicus Astronomical Center, ul. Bartycka 18, 00-716 Warsaw, Poland    C.B. Rulten University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K.    J. Ruppel Institut für Theoretische Physik, Lehrstuhl IV: Weltraum und Astrophysik, Ruhr-Universität Bochum, D 44780 Bochum, Germany    F. Ryde Oskar Klein Centre, Department of Physics, Royal Institute of Technology (KTH), Albanova, SE-10691 Stockholm, Sweden    V. Sahakian Yerevan Physics Institute, 2 Alikhanian Brothers St., 375036 Yerevan, Armenia    A. Santangelo Institut für Astronomie und Astrophysik, Universität Tübingen, Sand 1, D 72076 Tübingen, Germany    R. Schlickeiser Institut für Theoretische Physik, Lehrstuhl IV: Weltraum und Astrophysik, Ruhr-Universität Bochum, D 44780 Bochum, Germany    F.M. Schöck Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    A. Schönwald Institut für Physik, Humboldt-Universität zu Berlin, Newtonstr. 15, D 12489 Berlin, Germany    U. Schwanke Institut für Physik, Humboldt-Universität zu Berlin, Newtonstr. 15, D 12489 Berlin, Germany    S. Schwarzburg Institut für Astronomie und Astrophysik, Universität Tübingen, Sand 1, D 72076 Tübingen, Germany    S. Schwemmer Landessternwarte, Universität Heidelberg, Königstuhl, D 69117 Heidelberg, Germany    A. Shalchi Institut für Theoretische Physik, Lehrstuhl IV: Weltraum und Astrophysik, Ruhr-Universität Bochum, D 44780 Bochum, Germany    I. Sushch Institut für Physik, Humboldt-Universität zu Berlin, Newtonstr. 15, D 12489 Berlin, Germany    M. Sikora Nicolaus Copernicus Astronomical Center, ul. Bartycka 18, 00-716 Warsaw, Poland    J.L. Skilton School of Physics & Astronomy, University of Leeds, Leeds LS2 9JT, UK    H. Sol LUTH, Observatoire de Paris, CNRS, Université Paris Diderot, 5 Place Jules Janssen, 92190 Meudon, France    Ł. Stawarz Obserwatorium Astronomiczne, Uniwersytet Jagielloński, ul. Orla 171, 30-244 Kraków, Poland    R. Steenkamp University of Namibia, Private Bag 13301, Windhoek, Namibia    C. Stegmann Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    F. Stinzing Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany    G. Superina Laboratoire Leprince-Ringuet, Ecole Polytechnique, CNRS/IN2P3, F-91128 Palaiseau, France    A. Szostek Obserwatorium Astronomiczne, Uniwersytet Jagielloński, ul. Orla 171, 30-244 Kraków, Poland Laboratoire d’Astrophysique de Grenoble, INSU/CNRS, Université Joseph Fourier, BP 53, F-38041 Grenoble Cedex 9, France    P.H. Tam Landessternwarte, Universität Heidelberg, Königstuhl, D 69117 Heidelberg, Germany    J.-P. Tavernet LPNHE, Université Pierre et Marie Curie Paris 6, Université Denis Diderot Paris 7, CNRS/IN2P3, 4 Place Jussieu, F-75252, Paris Cedex 5, France    R. Terrier Astroparticule et Cosmologie (APC), CNRS, Universite Paris 7 Denis Diderot, 10, rue Alice Domon et Leonie Duquet, F-75205 Paris Cedex 13, France thanks: UMR 7164 (CNRS, Université Paris VII, CEA, Observatoire de Paris)    O. Tibolla Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    M. Tluczykont Universität Hamburg, Institut für Experimentalphysik, Luruper Chaussee 149, D 22761 Hamburg, Germany    C. van Eldik Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    G. Vasileiadis Laboratoire de Physique Théorique et Astroparticules, Université Montpellier 2, CNRS/IN2P3, CC 70, Place Eugène Bataillon, F-34095 Montpellier Cedex 5, France    C. Venter Unit for Space Physics, North-West University, Potchefstroom 2520, South Africa    L. Venter LUTH, Observatoire de Paris, CNRS, Université Paris Diderot, 5 Place Jules Janssen, 92190 Meudon, France    J.P. Vialle Laboratoire d’Annecy-le-Vieux de Physique des Particules, Université de Savoie, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France    P. Vincent LPNHE, Université Pierre et Marie Curie Paris 6, Université Denis Diderot Paris 7, CNRS/IN2P3, 4 Place Jussieu, F-75252, Paris Cedex 5, France    J. Vink Astronomical Institute Utrecht, University of Utrecht, P.O. Box 80000, 3508TA Utrecht, The Netherlands    M. Vivier IRFU/DSM/CEA, CE Saclay, F-91191 Gif-sur-Yvette, Cedex, France    H.J. Völk Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    F. Volpe Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany    S. Vorobiov Laboratoire de Physique Théorique et Astroparticules, Université Montpellier 2, CNRS/IN2P3, CC 70, Place Eugène Bataillon, F-34095 Montpellier Cedex 5, France    S.J. Wagner Landessternwarte, Universität Heidelberg, Königstuhl, D 69117 Heidelberg, Germany    M. Ward University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K.    A.A. Zdziarski Nicolaus Copernicus Astronomical Center, ul. Bartycka 18, 00-716 Warsaw, Poland    A. Zech LUTH, Observatoire de Paris, CNRS, Université Paris Diderot, 5 Place Jules Janssen, 92190 Meudon, France
Released 2009 Xxxxx XX
Key Words.:
-rays: observations – SNR: individual (SN 1006, G327.6+14.6) – supernova remnants
offprints: naumann-godo@llr.in2p3.fr, denauroi@in2p3.fr
Abstract

Context:

Aims:Recent theoretical predictions of the lowest very high energy (VHE) luminosity of SN 1006 are only a factor 5 below the previously published H.E.S.S. upper limit, thus motivating further in-depth observations of this source.

Methods:Deep observations at VHE energies (above 100 GeV) were carried out with the High Energy Stereoscopic System (H.E.S.S.) of Cherenkov Telescopes from 2003 to 2008. More than 100 hours of data have been collected and subjected to an improved analysis procedure.

Results:Observations resulted in the detection of VHE -rays from SN 1006. The measured -ray spectrum is compatible with a power-law, the flux is of the order of 1 of that detected from the Crab Nebula, and is thus consistent with the previously established H.E.S.S. upper limit. The source exhibits a bipolar morphology, which is strongly correlated with non-thermal X-rays.

Conclusions:Because the thickness of the VHE-shell is compatible with emission from a thin rim, particle acceleration in shock waves is likely to be the origin of the -ray signal. The measured flux level can be accounted for by inverse Compton emission, but a mixed scenario that includes leptonic and hadronic components and takes into account the ambient matter density inferred from observations also leads to a satisfactory description of the multi-wavelength spectrum.

1 Introduction

The source SN 1006 is the remnant of one of the few historical supernovae. It appeared in the southern sky on 1006 May 1 and was recorded by Chinese and Arab astronomers [Stephenson & Green 2002]. The remnant of this explosion was first identified at radio wavelengths on the basis of historical evidence [Gardner & Milne 1965]. The evolution of its luminosity indicates that it is the result of a Type Ia supernova [Schaefer 1996], probably the brightest supernova in recorded history. A distance of 2.2 kpc was derived by Winkler et al. (2003) based on comparing the optical proper motion with an estimate of the shock velocity derived from optical thermal line broadening assuming a high Mach number single-fluid shock.

Contemporary interest in the very high energy (VHE) emission from supernova remnants (SNRs) has arisen due to their association as prime candidates for Galactic cosmic-ray acceleration. Firstly, Galactic SNRs have sufficient kinetic energy to explain the estimated Galactic luminosity in cosmic rays of  erg/s. Secondly, and more importantly, it has been shown that diffusive shock acceleration provides a viable mechanism which can efficiently accelerate charged particles in the blast waves of SNRs (e.g. Drury 1983; Blandford & Eichler 1987; Jones & Ellison 1991; Berezhko et al. 1996). Indeed, most shell-type SNRs are non-thermal radio emitters, which confirms that electrons are accelerated up to at least GeV energies. Moreover, the limb-brightened non-thermal radio emission traces the site of effective particle acceleration.

The source SN 1006 was also the first SNR in which a non-thermal component of hard X-rays was detected in the rims of the remnant by ASCA [Koyama et al. 1995] and ROSAT [Willingale et al. 1996], whereas the interior of the remnant exhibits a thermal spectrum with line emission. The hard featureless power-law spectrum strongly implies a synchrotron origin of the radiation, which in turn suggests that electrons can be accelerated up to energies of  TeV. Subsequent arcsecond resolution images by Chandra revealed a small-scale structure in the nonthermal X-ray filaments of the NE rim of SN 1006 [Bamba et al. 2003, Long et al. 2003], supporting the idea of high B-fields in the bright limbs of the remnant [Berezhko et al. 2002]. An analysis of the X-ray observations from XMM-Newton by Rothenflug et. al (2004) leads to the conclusion that the magnetic field in the remnant is oriented in the NE-SW direction. The synchrotron emission would then be concentrated in regions where the shock is quasi-parallel [Völk et al. 2003].

Also, -ray observations of SN 1006 were carried out by ground-based -ray telescopes. A TeV -ray signal at the level of the Crab flux was claimed by the CANGAROO-I [Tanimori et al. 1998] and CANGAROO-II [Tanimori et al. 2001] telescopes, but subsequent stereoscopic observations of the source with the H.E.S.S. telescopes in 2003 and 2004 found no evidence of VHE -ray emission and derived an upper limit of  TeV ph cm s at 99.9% confidence level [Aharonian et al. 2005]. The CANGAROO-III telescope array found only an upper limit which is consistent with the H.E.S.S. result [Tanimori et al. 2005].

The initial non-detection of SN 1006 in VHE -rays does not invalidate the hypothesis of nuclear particle acceleration in the shock. Indeed, the hadronic -ray flux is very sensitive to the ambient gas density and hence the H.E.S.S. upper limit implies a constraint on  cm [Ksenofontov et al. 2005]. Indeed, being 500 pc above the Galactic plane, the remnant is relatively isolated, and the gas density around SN 1006 was recently estimated to be around 0.085 cm [Katsuda et al. 2009]. Ksenofontov et al. (2005) furthermore showed that the lower limit for the VHE -ray flux, which is given by the inverse Compton (IC) component derived from the integrated synchrotron flux and field amplification alone, was only a factor 5 below the H.E.S.S. upper limit. These predictions promoted deep observations with the H.E.S.S. telescopes.

2 H.E.S.S. observations and analysis methods

H.E.S.S. is an array of four 13 m diameter imaging atmospheric Cherenkov telescopes situated in the Khomas Highland in Namibia at an altitude of 1800 m above sea level [Bernloehr et al. 2003, Funk et al. 2004]. The source SN 1006 was observed in 2003 with the two telescopes that were operational at that time and with the complete H.E.S.S. array in the years since. After run selection the data set comprises 130 hours (live time) of observations, of which 18 hours were taken with two telescopes only. The latter yielded a smaller effective area than the data set recorded with the full array. For that reason they are used only in morphological studies and excluded in the spectral analysis.

The data were analysed with the Model Analysis [de Naurois & Rolland 2009], in which shower images of all triggered telescopes are compared to a pre-calculated model by means of a log-likelihood minimisation. The Model Analysis does not rely on any image-cleaning procedure and uses all pixels in the camera. The noise distributions in the pixels due to the night sky background are taken into account in the model fit and result in a superior treatment of shower tails. Therefore the Model Analysis results in a more precise reconstruction and better background suppression than more conventional techniques, thus leading to improved sensitivity.

Two different sets of cuts were used: The standard cuts, including a minimum image charge of 60 photoelectrons ( GeV), cover the full energy range and are used for the spectral analysis only. The hard cuts, with a larger charge cut of 200 photoelectrons, result in an improved signal-to-background ratio at the expense of lower statistics and a higher threshold of 500 GeV. These are used for the studies of source morphology.

The results presented below have been cross-checked using the 3D Model Analysis [Lemoine-Goumard et al. 2006, Naumann-Godo et al. 2009]. Both analyses yielded consistent results.

Significant -ray emission is detected from the direction of SN 1006, concentrated in two extended regions as shown in Fig. 1. This map shows the significance over a field-of-view of with a pixel size of obtained with hard cuts using the ring background subtraction technique [Berge et al. 2007] and a small integration radius of , close to the H.E.S.S. PSF of . As the pixel size is a factor 10 smaller than the integration radius, the bins are highly correlated. In two regions of the map the significance of the H.E.S.S. observation clearly exceeds 5 .

Figure 1: H.E.S.S. -ray significance map of SN 1006 using an integration radius of . The linear colour scale is in units of standard deviations. The white solid contours correspond to the regions which contain 80% of the non-thermal X-ray emission from the XMM-Newton flux map in the 2 - 4.5 keV energy range after smearing with the H.E.S.S. PSF, shown in the inset. The white dashed circles correspond to the regions that are excluded from background determination.
Figure 2: H.E.S.S. -ray significance distribution over the full field-of-view of SN 1006 (black histogram) and excluding the circular regions around the NE and SW emission regions (red histogram). A normal distribution (red dashed line) shows that the significance distribution over the rest of the field-of-view is compatible with expectation from statistical noise fluctuations.
Figure 3: H.E.S.S. sky area with -ray significance above some threshold as a function of its value over the full field-of-view of SN 1006 (black histogram) and excluding the circular regions around the NE and SW emission regions (red histogram).

The significance distribution over the field-of-view of is shown in Fig. 2, while Fig. 3 illustrates the area corresponding to the significance above a given level. The black histogram in both figures corresponds to the full field-of-view and exhibits strong deviation from a normal distribution at large significance values. The red histogram, restricted to the part of the field-of-view outside of the white dashed circles (Fig. 1) is compatible with a normal distribution, as denoted by the red dashed line (Fig. 2). This demonstrates that the distribution of events over the field-of-view (outside the two exclusion regions) is compatible with expectation from statistical fluctuations and that systematic effects concerning background estimation are under control.

3 Morphology

Two different integration regions were defined a priori from the XMM-Newton data set [Rothenflug et al. 2004]: a map of the flux in the 2 - 4.5 keV energy range (to exclude thermal contamination) was smoothed with the H.E.S.S. PSF, and regions which contained of the flux were calculated. The two resulting regions, denoted as NE Region and SW Region, are displayed as white contours in Fig. 1 and coincide well with the regions of largest H.E.S.S. significance.

Excess event counts and significances for both regions are given in Table 1 for the two sets of cuts. The ON photons are from the regions enclosed by the solid lines in Fig. 1, while the OFF events are taken from regions of identical shape rotated in the field-of-view of the instrument around the observation position and not intersecting the exclusion regions (enclosed by dashed lines in Fig. 1). Due to varying observation positions, the number of OFF regions varies from observation to observation. Individual observation values are combined into an average normalisation factor () quoted in Table 1. Similar excess event counts and significances are observed in both regions, thus attesting to the bipolar morphology of the remnant in the TeV energy range. This is a highly constraining result, because due to the relatively uniform target density around the remnant the H.E.S.S. morphology directly reflects the distribution of high-energy particles responsible for the -ray emission.

Region ON OFF # Significance
NE, Std Cuts 4306 25421 6.67 495 7.3
NE, Hard Cuts 619 2575 6.44 219 9.3
SW, Std Cuts 3798 26523 7.615 315 4.9
SW, Hard Cuts 548 2591 7.25 191 8.7
Table 1: H.E.S.S. excess events and significances for the two regions defined from X-ray observations. is the normalisation factor between OFF and ON exposures.

Figure 4 shows the -ray H.E.S.S. excess map, produced with hard cuts and the same integration radius of , overlaid with the smoothed XMM-Newton flux contours. A striking similarity between the -ray and X-ray emission regions is found. For a quantitative analysis uncorrelated radial and azimuthal profiles of the H.E.S.S. excess events were derived and compared to the XMM-Newton profiles (Figs. 5 and 6). Again the XMM-Newton data were smoothed to match the H.E.S.S. point spread function, and the relative normalisation was adjusted to the maximum value. Within error bars, the H.E.S.S. and XMM-Newton emission profiles are almost identical, thus possibly indicating a common origin.

Figure 4: H.E.S.S. -ray image of SN 1006. The linear colour scale is in units of excess counts per . Points within are correlated. The white cross indicates the geometrical centre of the SNR obtained from XMM data as explained in the text and the dashed circles correspond to as derived from the fit. The white star shows the centre of the circle encompassing the whole X-ray emission as derived by Rothenflug et al. (2004) and the white triangle the centre derived by Cassam-Chenaï et al. (2008) from H data. The white contours correspond to a constant X-ray intensity as derived from the XMM-Newton flux map and smoothed to the H.E.S.S. point spread function, enclosing respectively 80% , 60% , 40% and 20% of the X-ray emission. The inset shows the H.E.S.S. PSF using an integration radius of .

The geometrical X-ray centre of the SNR was derived from the unsmoothed XMM-Newton data by fitting them with a Gaussian radial profile convolved with an azimuthal profile with two Gaussian components, yielding 15h2m51.1s, -41d55’32.2” as the centre of the SNR with a radius of and a thickness of . Figure 5 shows the radial profiles of H.E.S.S. and smoothed XMM-Newton excess events from the centre of the SNR. When a Gaussian is fit to the H.E.S.S. profile (Fig. 5) the shell radius is found to be and the width of the radial distribution is , which is consistent with the H.E.S.S. point spread function, thereby showing that the emission region is compatible with a thin rim.

Figure 5: Radial profile around the centre of the SNR obtained from H.E.S.S. data and XMM-Newton data in the 2 - 4.5 keV energy band smoothed to H.E.S.S. PSF.
Figure 6: Azimuthal profile obtained from H.E.S.S. data and XMM-Newton data in the 2 - 4.5 keV energy band and smoothed to H.E.S.S. PSF, restricted to radii from the centre of the SNR. Azimuth corresponds to East, corresponds to North, to West and to South.

The azimuthal profile, restricted to radii from the centre of the SNR, is shown in Fig. 6 for H.E.S.S. data and smoothed XMM-Newton data in the 2 - 4.5 keV energy band. The azimuth is defined clockwise with zero toward the East. The H.E.S.S. profile is compatible with a superposition of two Gaussian emission regions almost at from each other, respectively centred on (SW region) and (NE region) and with similar widths of and .

4 Spectral analysis

Differential energy spectra of the VHE -ray emission were derived for both regions above the energy threshold of 260 GeV. These regions correspond to 80% of the X-ray emission (after smearing with the H.E.S.S. PSF) and therefore slightly underestimate the TeV emission of the full remnant.

The spectra for the NE and SW regions are compatible with power law distributions, , with comparable photon indices and fluxes. Confidence bands for power-law fits are shown in Fig. 7 and Table 2. The integral fluxes above 1 TeV correspond to less than of the Crab flux, making SN 1006 one of the faintest known VHE sources (Table 2). The derived fluxes are well below the previously published H.E.S.S. upper limits [Aharonian et al. 2005]. The observed photon index is somewhat steeper than generally expected from diffusive shock acceleration theory and may indicate that the upper cut-off of the high-energy particle distribution is being observed; however, the uncertainties on the spectrum preclude definitive conclusions on this point.

Figure 7: Differential energy spectra of SN 1006 extracted from the two regions NE and SW as defined in Sect. 2. The shaded bands correspond to the range of the power-law fit, taking into account statistical errors.
Region photon index
NE
SW

Table 2: Fit results for power-law fits to the energy spectra.

5 Discussion

The source SN 1006 is an ideal example of a shell-type supernova remnant because it represents a type Ia supernova exploding into an approximately uniform medium and magnetic field, thereby essentially maintaining the spherical geometry of a point explosion. This can be attributed to the fact that SN 1006 is about 500 pc above the Galactic plane in a relatively clean environment, where the external gas density is rather low, 0.085 cm as indicated by Katsuda et al. (2009). Moreover, SN 1006 is one of the best-observed SNRs with a rich data-set of astronomical multi-wavelength information in radio, optical and X-rays, and all the important parameters like the ejected mass, its distance and age are fairly well-known [Cassam-Chenaï et al. 2008]. For this reason, the semi-analytical models of Truelove McKee (1999) can be approximately applied and the velocity of the shock calculated. The value of the shock velocity calculated by this means agrees well with the recent measurement in X-rays by Katsuda et al. (2009), yielding  arcsec yr in the synchrotron emitting regions (NE and SW), which corresponds to  km/s for a distance of 2.2 kpc. This does not contradict the value of  arcsec yr measured by Winkler et al. (2003) in the optical filaments, which are situated in the NW region of the remnant. All those calculations neglect the dynamic role of accelerated particles however, which is potentially quite important.

The basic model of VHE -ray production requires particles accelerated to multi-TeV energies and a target comprising photons and/or matter of sufficient density. The close correlation between X-ray and VHE-emission points toward particle acceleration in the strong shocks revealed by the Chandra observation of the X-ray filaments. Moreover, the bipolar morphology of the VHE-emission in the NE and SW regions of the remnant supports a major result of diffusive shock acceleration theory, according to which efficient injection of suprathermal downstream charged nuclear ions is only possible for sufficiently small angles between the ambient magnetic field and shock normal, and therefore a higher density of accelerated nuclei at the poles is predicted [Ellison et al. 1995, Malkov & Völk 1995, Völk et al. 2003].

Radio [Reynolds 1996] and X-ray [Bamba et al. 2008] data integrated over the full remnant were combined with VHE -ray measurements to model the spectral energy distribution of the source in a simple one-zone stationary model. For the sake of consistency, the VHE -ray energy distribution was determined from the sum of the two previously defined regions. In this phenomenological model the current distribution of particles (electrons and/or protons) is prescribed with a given spectral shape corresponding to a power law with an exponential cutoff, from which emission due to synchrotron radiation, bremsstrahlung and IC scattering on the Cosmic Microwave Background (CMB) photons is computed. The production through interactions of protons with the ambient matter are obtained following Kelner et al. (2006).

It is clear that this model oversimplifies the acceleration process in an expanding remnant, as discussed by e.g. Drury et al. (1989) and Berezhko et al. (1996). In addition one must include the uncertainties introduced by the dynamics of the ejecta, the nonuniform structure of the ambient medium and the complexities of the reaction of the accelerated particles on both the magnetic field and the remnant dynamics. This is of importance when comparing the data to the model results below.

Figure 8: Broadband SED models of SN 1006 for a leptonic scenario (top), a hadronic one (centre) and a mixed leptonic/hadronic scenario (bottom). Top: Modelling was done by using an electron spectrum in the form of a power law with an index of 2.1, an exponential cutoff at 10 TeV and a total energy of erg. The magnetic field amounts to 30 µG. Centre: Modelling using a proton spectrum in the form of a power law with an index of 2.0, an exponential cutoff at 80 TeV and a total proton energy of  erg (using a lower energy cut off of 1 GeV). The electron/proton ratio above 1 GeV was with an electron spectral index of 2.1 and cutoff energy at 5 TeV. The magnetic field amounts to 120 µG and the average medium density is 0.085 cm. Bottom: Modelling using a mixture of the above two cases. The total proton energy was  erg, with , with exponential cutoffs at 8 TeV and 100 TeV for electrons and protons respectively. The magnetic field amounts to 45 µG. The radio data [Reynolds 1996], X-ray data [Bamba et al. 2008] and H.E.S.S. data (sum of the two regions) are indicated. The following processes have been taken into account: synchrotron radiation from primary electrons (dashed black lines), IC scattering (dotted red lines), bremsstrahlung (dot-dashed green lines) and proton-proton interactions (dotted blue lines). The Fermi/LAT sensitivity for one year is shown (pink) for Galactic (upper) and extragalactic (lower) background. The latter is more representative given that SN 1006 is north of the Galactic plane.

Assuming first a purely leptonic form (Fig. 8, top), the radio and X-ray data constrain the synchrotron part of the SED in a way that the slope of the electron spectrum, which is particularly sensitive to the slope of the radio data, is bounded between 2.0 and 2.2, while the cutoff energy of electrons is limited to about 10 TeV by the X-ray data assuming a magnetic field of 30 µG. With the particle spectrum constrained by radio and X-ray data, the resulting magnetic field needs to be higher than 30 µG so that the IC emission does not exceed the measured VHE-flux. A magnetic field of 30 µG implies that assuming Bohm diffusion, electrons of 1 TeV are confined in a shell of the width of 10 arcseconds, which is much smaller than the PSF of the H.E.S.S. instrument and is therefore compatible with the radial profile shown in Fig. 5. However, while this simple leptonic scenario can account for the measured VHE -ray flux, it fails to reproduce the slope of the VHE spectrum, which is much harder than the expectations from the IC process (see Fig. 8 top). But it should be noted that non-linear Fermi shock acceleration as reviewed by Malkov & Drury (2001) usually predicts curved cosmic ray spectra with different spectral shapes for protons and electrons. There is a hint of spectral curvature observed in the case of Tycho’s and Kepler’s supernova remnants in the radio regime [Reynolds & Ellison 1992]. For SN 1006 there is also an indication of the curvature of the electron spectrum in the GeV to TeV energy range [Allen et al. 2008]. These non-linear effects, which also might well introduce a spectral curvature in the VHE regime, are not addressed by this simple model.

In a second dominantly hadronic model (Fig. 8, middle) TeV emission results from proton-proton interactions with -production and subsequent decay, whereas the X-ray emission is still produced by leptonic interactions. A rough representation of the effect of spectral curvature is included by allowing for a slightly harder spectral index for protons than for radio-emitting electrons. A lower electron fraction allows us to account for the X-ray and radio emission with a higher field value of 120 µG, which is consistent with magnetic field amplification at the shock, as indicated by the above-mentioned measurements of thin X-ray filaments. Assuming an average medium density of 0.085 cm and a proton spectral index of 2.0 with a cutoff energy of 80 TeV (inferred from the maximum energy of TeV photons), this model requires a high overall fraction of about 20%, of the supernova energy to be converted into high-energy protons. Here  erg was assumed, near the upper end of the typical range of type Ia SN explosion energies (e.g. Woosley et al. 2007), as the assumed density, observed radius and known age of SN 1006 appear to require a higher than average explosion energy. Given that the VHE emission is concentrated in polar regions of the shell, the local shock acceleration efficiency would then be several times higher than this fraction.

In a third example (mixed model), hadronic and leptonic processes contribute almost equally to the very high-energy emission. The electron spectrum is similar to the aforementioned leptonic case and the total proton energy is set to 14% of the mechanical supernova energy with the electron/proton ratio , thus leaving the magnetic field and the cutoff energy of protons the only free parameters. In the example shown in Fig. 8 (bottom panel) the magnetic field amounts to 45 µG and the cutoff energy of protons is 100 TeV. This example illustrates that in this simple one-zone case it is possible to reproduce all the multi-wavelength data on SN 1006 to a reasonable degree of accuracy including the slope of the VHE-data. While these considerations cannot exclude any of the astrophysical scenarios, they serve as a quantitative illustration of the various alternatives.

Values of total electron and proton energy, cutoff energy and magnetic field obtained in the three aforementioned cases are summarised in Table 3. These parameters yield very similar values when the NE and SW regions are adjusted independently.

Model
Leptonic - -
Hadronic
Mixed
Table 3: Parameters used in the spectral energy modelling shown in Fig. 8. Spectral indices have been fixed to and respectively for electrons and protons.

More elaborate models using e.g. a nonlinear kinetic acceleration theory [Berezhko et al. 2009] go beyond the simple approach developed here and lead to precise predictions that could be quantitatively tested against the data. Several effects which were not included in the simple model above would alter the total energy in accelerated particles required for the hadronic component. Beyond the spectral curvature mentioned previously, these include the higher compression of the target matter induced by the dynamical reaction of the accelerated particles, and consideration of the heavier nuclei composition of the accelerated hadrons instead of the pure protons assumed here. Measurements in the GeV-energy range would be pivotal to distinguish between the different scenarios. Unfortunately, the sensitivity of the Fermi Large Area Telescope for one year as given in Atwood et al. (2009) is of a factor of the order of 10 too low (depending on the model and the exact diffuse background flux) to measure the predicted flux at 1 GeV as shown in Fig.8, which makes the detection of SN 1006 by Fermi LAT rather unlikely.

6 Conclusions

Very high energy -rays from SN 1006 have been detected by H.E.S.S. The measured flux above 1 TeV is of the order of 1 of that detected from the Crab Nebula and therefore compatible with the previously published upper limit [Aharonian et al. 2005]. The bipolar morphology apparent in -rays is consistent with the non-thermal emission regions also visible in X-rays. As the VHE-shell is compatible with a scenario of thin rim emission, particle acceleration in the very narrow X-ray filaments, which are signatures of shocks, is also likely to be at the origin of the -ray signal. The measured flux level can be accounted for by inverse Compton emission assuming a magnetic field of about 30 G. A mixed scenario including leptonic and hadronic processes and taking into account the ambient matter density estimated from observation also leads to a satisfactory description of the multi-wavelength spectrum, assuming a high proton-acceleration efficiency. None of the models can be excluded at the level of modelling presented here.

Acknowledgements.
The support of the Namibian authorities and of the University of Namibia in facilitating the construction and operation of H.E.S.S. is gratefully acknowledged, as is the support by the German Ministry for Education and Research (BMBF), the Max Planck Society, the French Ministry for Research, the CNRS-IN2P3 and the Astroparticle Interdisciplinary Programme of the CNRS, the U.K. Science and Technology Facilities Council (STFC), the IPNP of the Charles University, the Polish Ministry of Science and Higher Education, the South African Department of Science and Technology and National Research Foundation, and by the University of Namibia. We appreciate the excellent work of the technical support staff in Berlin, Durham, Hamburg, Heidelberg, Palaiseau, Paris, Saclay, and in Namibia in the construction and operation of the equipment.

References

  • [Aharonian et al. 2005] F. Aharonian et al. (H.E.S.S. Collaboration), AA 437, 135 (2005).
  • [Allen et al. 2008] G. E. Allen, J. C. Houck and S. J. Sturner, ApJ 683, 773 (2008).
  • [Atwood et al. 2009] W. B. Atwood et al., ApJ 697, 1071 (2009).
  • [Bamba et al. 2003] A. Bamba, R. Yamazaki, M. Keno and K. Koyama, ApJ 589, 827 (2003).
  • [Bamba et al. 2008] A. Bamba, Y. Fukazawa, J. S. Hiraga et al., PASJ 60, S153 (2008).
  • [Berezhko et al. 1996] E. G. Berezhko, V. K. Elshin and L. T. Ksenofontov, J. Exp. Theor. Phys. 82, 1 (1996).
  • [Berezhko et al. 2002] E. G. Berezhko, L. T. Ksenofontov and H. Völk, AA 395, 943 (2002).
  • [Berezhko et al. 2009] E. G. Berezhko, L. T. Ksenofontov and H. Völk, AA 505, 169 (2009).
  • [Berge et al. 2007] D. Berge, S. Funk and J. Hinton, AA 466, 1219 (2007).
  • [Bernloehr et al. 2003] K. Bernlöhr, O. Carrol, R. Cornils et al., Astropart. Phys. 20, 111 (2003).
  • [Blandford & Eichler 1987] R. D. Blandford and D. Eichler, Phys. Rep. 154, 1 (1987).
  • [Cassam-Chenaï et al. 2008] G. Cassam-Chenaï, J. P. Hughes, E. M. Reynoso, C. Badenes and D. Moffett, ApJ 680, 1180 (2008).
  • [Drury 1983] L. O’C. Drury, Rep. Prog. Phys. 46, 973 (1983).
  • [Drury et al. 1989] L. O’C. Drury, W. J. Markiewicz and H. J. Völk, AA 225, 179 (1989).
  • [Ellison et al. 1995] D. C. Ellison, M. G. Baring and F. C. Jones , ApJ 453, 873 (1995).
  • [Funk et al. 2004] S. Funk, G. Hermann, J. Hinton et al., Astropart. Phys. 22, 285 (2004).
  • [Gardner & Milne 1965] F. F. Gardner, D. K. Milne, Astronom. J. 70, 754 (1965).
  • [Jones & Ellison 1991] F. C. Jones and D. C. Ellison, Space Sci. Rev. 58, 259 (1991).
  • [Katsuda et al. 2009] S. Katsuda, R. Petre, K. S. Long et al., ApJ 692, L105 (2009).
  • [Kelner et al. 2006] S. R. Kelner, F. A. Aharonian, V. V. Bugayov, Phys. Rev. D 74, 034018 (2006).
  • [Koyama et al. 1995] K. Koyama, R. Petre, E. V. Gotthelf et al., Nature 378, 378 (1995).
  • [Ksenofontov et al. 2005] L. T. Ksenofontov, E. G. Berezhko and H. J. Völk, AA 443, 973 (2005).
  • [Lemoine-Goumard et al. 2006] M. Lemoine-Goumard, B. Degrange and M. Tluczykont, Astropart. Phys. 25, 195 (2006).
  • [Long et al. 2003] K. S. Long, S. P.  Reynolds, J. C. Raymond et al., ApJ 586, L1162 (2003).
  • [Malkov & Völk 1995] M. A. Malkov and H.-J. Völk, AA 300, 605 (1995).
  • [Malkov & Drury 2001] M. A. Malkov and L. O’C. Drury, Rep. Prog. Phys. 64, 429 (2001).
  • [Naumann-Godo et al. 2009] M. Naumann-Godo, M. Lemoine-Goumard and B. Degrange, Astropart. Phys. 31, 421 (2009).
  • [de Naurois & Rolland 2009] M. de Naurois and L. Rolland, Astropart. Phys. 32, 231 (2009).
  • [Reynolds 1996] S. P. Reynolds, ApJ 459, L13 (1996).
  • [Reynolds & Ellison 1992] S. P. Reynolds and D. C. Ellison, ApJ 399, L75 (1992).
  • [Rothenflug et al. 2004] R. Rothenflug, J. Ballet, G. Dubner et al., AA 425, 121 (2004).
  • [Schaefer 1996] B. E. Schaefer, ApJ 459, 438 (1996).
  • [Stephenson & Green 2002] F. R. Stephenson and D. A. Green, ”Historical supernovae and their remnants”, Oxford: Clarendon Press (2002).
  • [Tanimori et al. 1998] T. Tanimori et al., ApJ L25, 135–139 (1998).
  • [Tanimori et al. 2001] T. Tanimori et al., CRR Rep 478, 33 (2001).
  • [Tanimori et al. 2005] T. Tanimori et al., 29th International Cosmic Ray Conference, ICRC 2005 Conference Proceedings, Pune, 2005, vol 4, 215.
  • [Truelove & McKee 1999] J. K. Truelove and C. F. McKee, ApJS 120, 299 (1999).
  • [Völk et al. 2003] H. Völk, E. G. Berezhko and L. T. Ksenofontov, AA 409, 563 (2003).
  • [Willingale et al. 1996] R. Willingale, R. J. West, J. P. Pye, G. C. Steward, MNRAS 278, 749 (1996).
  • [Winkler et al. 2003] P. F. Winkler, G. Gupta and K. S. Long, ApJ 585, 324 (2003).
  • [Woosley et al. 2007] S. E. Woosley, D. Kasen, S. Blinnikov and E. Sorokina, ApJ 662, 487 (2007).
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