Multiwavelength observations of a VHE gamma-ray flare from PKS 1510–089 in 2015

Multiwavelength observations of a VHE gamma-ray flare from PKS 1510–089 in 2015

M. L. Ahnen ETH Zurich, CH-8093 Zurich, Switzerland    S. Ansoldi Università di Udine, and INFN Trieste, I-33100 Udine, Italy also at the Department of Physics of Kyoto University, Japan    L. A. Antonelli INAF National Institute for Astrophysics, I-00136 Rome, Italy    C. Arcaro Università di Padova and INFN, I-35131 Padova, Italy    A. Babić Croatian MAGIC Consortium, Rudjer Boskovic Institute, University of Rijeka, University of Split - FESB, University of Zagreb - FER, University of Osijek,Croatia    B. Banerjee Saha Institute of Nuclear Physics, 1/AF Bidhannagar, Salt Lake, Sector-1, Kolkata 700064, India    P. Bangale Max-Planck-Institut für Physik, D-80805 München, Germany    U. Barres de Almeida Max-Planck-Institut für Physik, D-80805 München, Germany now at Centro Brasileiro de Pesquisas Físicas (CBPF/MCTI), R. Dr. Xavier Sigaud, 150 - Urca, Rio de Janeiro - RJ, 22290-180, Brazil    J. A. Barrio Universidad Complutense, E-28040 Madrid, Spain    W. Bednarek University of Łódź, PL-90236 Lodz, Poland    E. Bernardini Deutsches Elektronen-Synchrotron (DESY), D-15738 Zeuthen, Germany
Humboldt University of Berlin, Institut für Physik Newtonstr. 15, 12489 Berlin Germany
   A. Berti Università di Udine, and INFN Trieste, I-33100 Udine, Italy also at University of Trieste    B. Biasuzzi Università di Udine, and INFN Trieste, I-33100 Udine, Italy    A. Biland ETH Zurich, CH-8093 Zurich, Switzerland    O. Blanch Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    S. Bonnefoy Universidad Complutense, E-28040 Madrid, Spain    G. Bonnoli Università di Siena, and INFN Pisa, I-53100 Siena, Italy    F. Borracci Max-Planck-Institut für Physik, D-80805 München, Germany    T. Bretz Universität Würzburg, D-97074 Würzburg, Germany now at Ecole polytechnique fédérale de Lausanne (EPFL), Lausanne, Switzerland    R. Carosi Università di Siena, and INFN Pisa, I-53100 Siena, Italy    A. Carosi INAF National Institute for Astrophysics, I-00136 Rome, Italy    A. Chatterjee Saha Institute of Nuclear Physics, 1/AF Bidhannagar, Salt Lake, Sector-1, Kolkata 700064, India    P. Colin Max-Planck-Institut für Physik, D-80805 München, Germany    E. Colombo Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    J. L. Contreras Universidad Complutense, E-28040 Madrid, Spain    J. Cortina Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    S. Covino INAF National Institute for Astrophysics, I-00136 Rome, Italy    P. Cumani Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    P. Da Vela Università di Siena, and INFN Pisa, I-53100 Siena, Italy    F. Dazzi Max-Planck-Institut für Physik, D-80805 München, Germany    A. De Angelis Università di Padova and INFN, I-35131 Padova, Italy    B. De Lotto Università di Udine, and INFN Trieste, I-33100 Udine, Italy    E. de Oña Wilhelmi Institute for Space Sciences (CSIC/IEEC), E-08193 Barcelona, Spain    F. Di Pierro INAF National Institute for Astrophysics, I-00136 Rome, Italy    M. Doert Technische Universität Dortmund, D-44221 Dortmund, Germany    A. Domínguez Universidad Complutense, E-28040 Madrid, Spain    D. Dominis Prester Croatian MAGIC Consortium, Rudjer Boskovic Institute, University of Rijeka, University of Split - FESB, University of Zagreb - FER, University of Osijek,Croatia    D. Dorner Universität Würzburg, D-97074 Würzburg, Germany    M. Doro Università di Padova and INFN, I-35131 Padova, Italy    S. Einecke Technische Universität Dortmund, D-44221 Dortmund, Germany    D. Eisenacher Glawion Universität Würzburg, D-97074 Würzburg, Germany    D. Elsaesser Technische Universität Dortmund, D-44221 Dortmund, Germany    M. Engelkemeier Technische Universität Dortmund, D-44221 Dortmund, Germany    V. Fallah Ramazani Finnish MAGIC Consortium, Tuorla Observatory, University of Turku and Astronomy Division, University of Oulu, Finland    A. Fernández-Barral Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    D. Fidalgo Universidad Complutense, E-28040 Madrid, Spain    M. V. Fonseca Universidad Complutense, E-28040 Madrid, Spain    L. Font Unitat de Física de les Radiacions, Departament de Física, and CERES-IEEC, Universitat Autònoma de Barcelona, E-08193 Bellaterra, Spain    C. Fruck Max-Planck-Institut für Physik, D-80805 München, Germany    D. Galindo Universitat de Barcelona, ICC, IEEC-UB, E-08028 Barcelona, Spain    R. J. García López Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    M. Garczarczyk Deutsches Elektronen-Synchrotron (DESY), D-15738 Zeuthen, Germany
   M. Gaug Unitat de Física de les Radiacions, Departament de Física, and CERES-IEEC, Universitat Autònoma de Barcelona, E-08193 Bellaterra, Spain    P. Giammaria INAF National Institute for Astrophysics, I-00136 Rome, Italy    N. Godinović Croatian MAGIC Consortium, Rudjer Boskovic Institute, University of Rijeka, University of Split - FESB, University of Zagreb - FER, University of Osijek,Croatia    D. Gora Deutsches Elektronen-Synchrotron (DESY), D-15738 Zeuthen, Germany
   D. Guberman Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    D. Hadasch Japanese MAGIC Consortium, ICRR, The University of Tokyo, Department of Physics and Hakubi Center, Kyoto University, Tokai University, The University of Tokushima, Japan    A. Hahn Max-Planck-Institut für Physik, D-80805 München, Germany    T. Hassan Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    M. Hayashida Japanese MAGIC Consortium, ICRR, The University of Tokyo, Department of Physics and Hakubi Center, Kyoto University, Tokai University, The University of Tokushima, Japan    J. Herrera Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    J. Hose Max-Planck-Institut für Physik, D-80805 München, Germany    D. Hrupec Croatian MAGIC Consortium, Rudjer Boskovic Institute, University of Rijeka, University of Split - FESB, University of Zagreb - FER, University of Osijek,Croatia    G. Hughes ETH Zurich, CH-8093 Zurich, Switzerland    K. Ishio Max-Planck-Institut für Physik, D-80805 München, Germany    Y. Konno Japanese MAGIC Consortium, ICRR, The University of Tokyo, Department of Physics and Hakubi Center, Kyoto University, Tokai University, The University of Tokushima, Japan    H. Kubo Japanese MAGIC Consortium, ICRR, The University of Tokyo, Department of Physics and Hakubi Center, Kyoto University, Tokai University, The University of Tokushima, Japan    J. Kushida Japanese MAGIC Consortium, ICRR, The University of Tokyo, Department of Physics and Hakubi Center, Kyoto University, Tokai University, The University of Tokushima, Japan    D. Kuveždić Croatian MAGIC Consortium, Rudjer Boskovic Institute, University of Rijeka, University of Split - FESB, University of Zagreb - FER, University of Osijek,Croatia    D. Lelas Croatian MAGIC Consortium, Rudjer Boskovic Institute, University of Rijeka, University of Split - FESB, University of Zagreb - FER, University of Osijek,Croatia    E. Lindfors Finnish MAGIC Consortium, Tuorla Observatory, University of Turku and Astronomy Division, University of Oulu, Finland    S. Lombardi INAF National Institute for Astrophysics, I-00136 Rome, Italy    F. Longo Università di Udine, and INFN Trieste, I-33100 Udine, Italy also at University of Trieste    M. López Universidad Complutense, E-28040 Madrid, Spain    P. Majumdar Saha Institute of Nuclear Physics, 1/AF Bidhannagar, Salt Lake, Sector-1, Kolkata 700064, India    M. Makariev Inst. for Nucl. Research and Nucl. Energy, BG-1784 Sofia, Bulgaria    G. Maneva Inst. for Nucl. Research and Nucl. Energy, BG-1784 Sofia, Bulgaria    M. Manganaro Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    K. Mannheim Universität Würzburg, D-97074 Würzburg, Germany    L. Maraschi INAF National Institute for Astrophysics, I-00136 Rome, Italy    M. Mariotti Università di Padova and INFN, I-35131 Padova, Italy    M. Martínez Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    D. Mazin Max-Planck-Institut für Physik, D-80805 München, Germany also at Japanese MAGIC Consortium    U. Menzel Max-Planck-Institut für Physik, D-80805 München, Germany    R. Mirzoyan Max-Planck-Institut für Physik, D-80805 München, Germany    A. Moralejo Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    E. Moretti Max-Planck-Institut für Physik, D-80805 München, Germany    D. Nakajima Japanese MAGIC Consortium, ICRR, The University of Tokyo, Department of Physics and Hakubi Center, Kyoto University, Tokai University, The University of Tokushima, Japan    V. Neustroev Finnish MAGIC Consortium, Tuorla Observatory, University of Turku and Astronomy Division, University of Oulu, Finland    A. Niedzwiecki University of Łódź, PL-90236 Lodz, Poland    M. Nievas Rosillo Universidad Complutense, E-28040 Madrid, Spain    K. Nilsson Finnish MAGIC Consortium, Tuorla Observatory, University of Turku and Astronomy Division, University of Oulu, Finland now at Finnish Centre for Astronomy with ESO (FINCA), Turku, Finland    K. Nishijima Japanese MAGIC Consortium, ICRR, The University of Tokyo, Department of Physics and Hakubi Center, Kyoto University, Tokai University, The University of Tokushima, Japan    K. Noda Max-Planck-Institut für Physik, D-80805 München, Germany    L. Nogués Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    S. Paiano Università di Padova and INFN, I-35131 Padova, Italy    J. Palacio Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    M. Palatiello Università di Udine, and INFN Trieste, I-33100 Udine, Italy    D. Paneque Max-Planck-Institut für Physik, D-80805 München, Germany    R. Paoletti Università di Siena, and INFN Pisa, I-53100 Siena, Italy    J. M. Paredes Universitat de Barcelona, ICC, IEEC-UB, E-08028 Barcelona, Spain    X. Paredes-Fortuny Universitat de Barcelona, ICC, IEEC-UB, E-08028 Barcelona, Spain    G. Pedaletti Deutsches Elektronen-Synchrotron (DESY), D-15738 Zeuthen, Germany
   M. Peresano Università di Udine, and INFN Trieste, I-33100 Udine, Italy    L. Perri INAF National Institute for Astrophysics, I-00136 Rome, Italy    M. Persic Università di Udine, and INFN Trieste, I-33100 Udine, Italy also at INAF-Trieste and Dept. of Physics & Astronomy, University of Bologna    J. Poutanen Finnish MAGIC Consortium, Tuorla Observatory, University of Turku and Astronomy Division, University of Oulu, Finland    P. G. Prada Moroni Università di Pisa, and INFN Pisa, I-56126 Pisa, Italy    E. Prandini Università di Padova and INFN, I-35131 Padova, Italy    I. Puljak Croatian MAGIC Consortium, Rudjer Boskovic Institute, University of Rijeka, University of Split - FESB, University of Zagreb - FER, University of Osijek,Croatia    J. R. Garcia Max-Planck-Institut für Physik, D-80805 München, Germany    I. Reichardt Università di Padova and INFN, I-35131 Padova, Italy    W. Rhode Technische Universität Dortmund, D-44221 Dortmund, Germany    M. Ribó Universitat de Barcelona, ICC, IEEC-UB, E-08028 Barcelona, Spain    J. Rico Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    T. Saito Japanese MAGIC Consortium, ICRR, The University of Tokyo, Department of Physics and Hakubi Center, Kyoto University, Tokai University, The University of Tokushima, Japan    K. Satalecka Deutsches Elektronen-Synchrotron (DESY), D-15738 Zeuthen, Germany
   S. Schroeder Technische Universität Dortmund, D-44221 Dortmund, Germany    T. Schweizer Max-Planck-Institut für Physik, D-80805 München, Germany    S. N. Shore Università di Pisa, and INFN Pisa, I-56126 Pisa, Italy    A. Sillanpää Finnish MAGIC Consortium, Tuorla Observatory, University of Turku and Astronomy Division, University of Oulu, Finland    J. Sitarek University of Łódź, PL-90236 Lodz, Poland Corresponding authors: J. Sitarek (jsitarek@uni.lodz.pl), J. Becerra González, E. Lindfors, F. Tavecchio, M. Vazquez Acosta    I. Å nidarić Croatian MAGIC Consortium, Rudjer Boskovic Institute, University of Rijeka, University of Split - FESB, University of Zagreb - FER, University of Osijek,Croatia    D. Sobczynska University of Łódź, PL-90236 Lodz, Poland    A. Stamerra INAF National Institute for Astrophysics, I-00136 Rome, Italy    M. Strzys Max-Planck-Institut für Physik, D-80805 München, Germany    T. Surić Croatian MAGIC Consortium, Rudjer Boskovic Institute, University of Rijeka, University of Split - FESB, University of Zagreb - FER, University of Osijek,Croatia    L. Takalo Finnish MAGIC Consortium, Tuorla Observatory, University of Turku and Astronomy Division, University of Oulu, Finland    F. Tavecchio INAF National Institute for Astrophysics, I-00136 Rome, Italy    P. Temnikov Inst. for Nucl. Research and Nucl. Energy, BG-1784 Sofia, Bulgaria    T. Terzić Croatian MAGIC Consortium, Rudjer Boskovic Institute, University of Rijeka, University of Split - FESB, University of Zagreb - FER, University of Osijek,Croatia    D. Tescaro Università di Padova and INFN, I-35131 Padova, Italy    M. Teshima Max-Planck-Institut für Physik, D-80805 München, Germany also at Japanese MAGIC Consortium    D. F. Torres ICREA and Institute for Space Sciences (CSIC/IEEC), E-08193 Barcelona, Spain    N. Torres-Albà Universitat de Barcelona, ICC, IEEC-UB, E-08028 Barcelona, Spain    T. Toyama Max-Planck-Institut für Physik, D-80805 München, Germany    A. Treves Università di Udine, and INFN Trieste, I-33100 Udine, Italy    G. Vanzo Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    M. Vazquez Acosta Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    I. Vovk Max-Planck-Institut für Physik, D-80805 München, Germany    J. E. Ward Institut de Fisica d’Altes Energies (IFAE), The Barcelona Institute of Science and Technology, Campus UAB, 08193 Bellaterra (Barcelona), Spain    M. Will Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    M. H. Wu Institute for Space Sciences (CSIC/IEEC), E-08193 Barcelona, Spain    D. Zarić Croatian MAGIC Consortium, Rudjer Boskovic Institute, University of Rijeka, University of Split - FESB, University of Zagreb - FER, University of Osijek,Croatia    R. Desiante (MAGIC Collaboration) Università di Udine, and INFN Trieste, I-33100 Udine, Italy    J. Becerra González Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain now at NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA and Department of Physics and Department of Astronomy, University of Maryland, College Park, MD 20742, USA    F. D’Ammando Dip. di Fisica e Astronomia, Università di Bologna, Viale Berti Pichat 6/2, I-40127 Bologna, Italy INAF – Istituto di Radioastronomia, Via Gobetti 101, I-40129 Bologna, Italy    S. Larsson (Fermi-LAT Collaboration) KTH Royal Institute of Technology, Department of Physics, AlbaNova, SE-10691 Stockholm, Sweden Oskar Klein Centre for Cosmoparticle Physics, AlbaNova, SE-10691 Stockholm, Sweden    C. M. Raiteri INAF, Osservatorio Astrofisico di Torino, via Osservatorio 20, I-10025 Pino Torinese, Italy    R. Reinthal Tuorla Observatory, Department of Physics and Astronomy, University of Turku, Finland    A. Lähteenmäki Aalto University Metsähovi Radio Observatory, Metsähovintie 114, 02540 Kylmälä, Finland    E. Järvelä Aalto University Metsähovi Radio Observatory, Metsähovintie 114, 02540 Kylmälä, Finland Aalto University Department of Radio Science and Engineering, P.O. BOX 13000, FI-00076 AALTO, Finland.    M. Tornikoski Aalto University Metsähovi Radio Observatory, Metsähovintie 114, 02540 Kylmälä, Finland    V. Ramakrishnan Aalto University Metsähovi Radio Observatory, Metsähovintie 114, 02540 Kylmälä, Finland    S. G. Jorstad IAR, Boston University, 725 Commonwealth Ave, Boston, 02215, USA St.Petersburg State University, Universitetsky prospekt, 28, St. Petersburg, 198504, Russia    A. P. Marscher IAR, Boston University, 725 Commonwealth Ave, Boston, 02215, USA    V. Bala IAR, Boston University, 725 Commonwealth Ave, Boston, 02215, USA    N. R. MacDonald IAR, Boston University, 725 Commonwealth Ave, Boston, 02215, USA    N. Kaur Physical Research Laboratory, Ahmedabad 380009, Gujrat, India Indian Institute of Technology, Gandhinagar 382355, Gujrat, India    Sameer Physical Research Laboratory, Ahmedabad 380009, Gujrat, India Department of Astronomy and Astrophysics, The Pennsylvania State University, 532-D, Davey Laboratory, University Park, PA 16802, USA    K. Baliyan Physical Research Laboratory, Ahmedabad 380009, Gujrat, India    J. A. Acosta-Pulido Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    C. Lazaro Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    C. Martínez-Lombilla Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    A. B. Grinon-Marin Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    A. Pastor Yabar Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    C. Protasio Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain    M. I. Carnerero Inst. de Astrofísica de Canarias, E-38200 La Laguna, Tenerife, Spain Universidad de La Laguna, Dpto. Astrofísica, E-38206 La Laguna, Tenerife, Spain INAF, Osservatorio Astrofisico di Torino, via Osservatorio 20, I-10025 Pino Torinese, Italy    H. Jermak Department of Physics, Lancaster University, Lancaster, LA1 4YW, UK. Astrophysics Research Institute, Liverpool John Moores University, Brownlow Hill, Liverpool, L3 5RF, UK.    I. A. Steele Astrophysics Research Institute, Liverpool John Moores University, Brownlow Hill, Liverpool, L3 5RF, UK.    V. M. Larionov St.Petersburg State University, Universitetsky prospekt, 28, St. Petersburg, 198504, Russia Pulkovo Observatory, St. Petersburg, Russia    G. A. Borman Crimean Astrophysical Observatory, Crimea    T. S. Grishina St.Petersburg State University, Universitetsky prospekt, 28, St. Petersburg, 198504, Russia
Received ; accepted
Key Words.:
galaxies: active – galaxies: jets – gamma rays: galaxies – quasars: individual: PKS 1510-089
Abstract

Context:PKS 1510–089 is one of only a few flat spectrum radio quasars detected in the VHE (very-high-energy,  GeV) gamma-ray band.

Aims: We study the broadband spectral and temporal properties of the PKS 1510–089 emission during a high gamma-ray state.

Methods:We performed VHE gamma-ray observations of PKS 1510–089 with the MAGIC telescopes during a long high gamma-ray state in May 2015. In order to perform broadband modelling of the source, we have also gathered contemporaneous multiwavelength data in radio, IR, optical photometry and polarization, UV, X-ray and GeV gamma-ray ranges. We construct a broadband spectral energy distribution (SED) in two periods, selected according to VHE gamma-ray state.

Results:PKS 1510–089 has been detected by MAGIC during a few day-long observations performed in the middle of a long, high optical and gamma-ray state, showing for the first time a significant VHE gamma-ray variability. Similarly to the optical and gamma-ray high state of the source detected in 2012, it was accompanied by a rotation of the optical polarization angle and the emission of a new jet component observed in radio. However, due to large uncertainty on the knot separation time, the association with the VHE gamma-ray emission cannot be firmly established. The spectral shape in the VHE band during the flare is similar to the ones obtained during previous measurements of the source. The observed flux variability sets for the first time constraints on the size of the region from which VHE gamma rays are emitted. We model the broadband SED in the framework of the external Compton scenario and discuss the possible emission site in view of multiwavelength data and alternative emission models.

Conclusions:

1 Introduction

PKS 1510–089 is a bright flat spectrum radio quasar (FSRQ) located at the redshift of (Tanner et al., 1996). The source is one of only six objects firmly classified as a FSRQ from which gamma-ray emission has been detected in the very-high-energy (VHE,  GeV) range (Abramowski et al., 2013). Moreover, one of the highest recorded apparent speed of superluminal motion, up to , has been seen in the ultrarelativistic jet of PKS 1510–089 (Jorstad et al., 2005). Like in many other FSRQs, the GeV gamma-ray emission of PKS 1510–089 is strongly variable (Abdo et al., 2010; Saito et al., 2013; Aleksić et al., 2014). The doubling-time scales of the PKS 1510–089 flares observed in the GeV range go down to 1 h (Saito et al., 2013).

Most of the FSRQs have been detected in the VHE gamma-ray range during (usually short) flares (see, e.g., Albert et al., 2008; Aleksić et al., 2011; Ahnen et al., 2015). Since 2013 MAGIC performs regular monitoring of PKS 1510–089. Interestingly, until 2015, no variability was seen in PKS 1510–089 in VHE gamma rays; neither in H.E.S.S (Abramowski et al., 2013) nor in MAGIC (Aleksić et al., 2014) observations. One should note, however, that both VHE gamma-ray detections happened during long periods of enhanced optical and GeV gamma-ray activity. Hence no low-state VHE gamma-ray emission has been established so far from PKS 1510–089.

In May 2015, a strong flare of PKS 1510–089 was observed in GeV gamma-rays by the Large Area Telescope (LAT) on board the Fermi satellite. The source showed also at this time high activity in optical (Jankowsky et al., 2015; Mirzoyan, 2015) and IR bands (Sameer et al., 2015; Carrasco et al., 2015). The high state triggered further MAGIC observations, which led to the detection of an enhanced VHE gamma-ray activity from the source (Mirzoyan, 2015). The VHE gamma-ray emission has been also observed by the H.E.S.S. telescope (Zacharias et al., 2016). In May 2016 another flare happened (de Naurois, 2016; Mirzoyan, 2016), with an even stronger VHE gamma-ray flux than in May 2015. The May 2016 flare will be discussed in a separate paper.

In this paper we report on the observations of PKS 1510–089 during the May 2015 flare. In Section 2 we shortly introduce the instruments which provided multiwavelength data and describe the data reduction procedures. In Section 3 we present the multiwavelength behaviour of the source. Section 4 is devoted to the interpretation of the data in the framework of an external Compton model. The most important results are summarized in Section 5.

2 Instruments, observations and data analysis

During the May 2015 outburst PKS 1510–089 was observed by various instruments in a broad range of frequencies (from radio up to VHE gamma rays). In this section we introduce the different instruments and data sets and explain the data analysis procedure.

2.1 Magic

MAGIC is a system of two Imaging Atmospheric Cherenkov Telescopes with a mirror dish diameter of 17 m each. They are located in Canary Island of La Palma ( N,  W), at the height of 2200 m a.s.l. (Aleksić et al., 2016a). As PKS 1510–089 is a southern source, only observable at zenith angle , the corresponding trigger threshold is  GeV (Aleksić et al., 2016b), about 1.7 times larger than for the low zenith observations.

The MAGIC telescopes observed PKS 1510–089 for 5.4 hours between 18 and 24 of May, 2015 (MJD 57160–57166). The data have been analyzed using MARS, the standard analysis package of MAGIC (Zanin et al., 2013; Aleksić et al., 2016b). As part of the data set was affected by Calima111Calima is a dust wind originating in Saharian Air Layer. we have applied a correction for the atmosphere transmission based on LIDAR information (Fruck & Gaug, 2015).

2.2 Fermi-Lat

Fermi-LAT monitors the gamma-ray sky every 3 h in the energy range from to beyond (Atwood et al., 2009). An analysis of the (publicly available) Pass 8 SOURCE class events was performed for a Region of Interest (ROI) of radius centered at the position of PKS 1510–089. In order to reduce contamination from the Earth Limb, a zenith angle cut of was applied. The analysis was performed with the ScienceTools software package version v10r0p5 using the P8R2_SOURCE_V6222 http://fermi.gsfc.nasa.gov/ssc/data/analysis/documentation/Cicerone/Cicerone_LAT_IRFs/IRF_overview.html instrument response function and the gll_iem_v06 and iso_P8R2_SOURCE_V6_v06 models333 http://fermi.gsfc.nasa.gov/ssc/data/access/lat/BackgroundModels.html for the Galactic and isotropic diffuse emission (Acero et al., 2016), respectively.

An unbinned likelihood analysis was applied using gtlike, including in the model all 3FGL sources (Acero et al., 2015) within from PKS 1510–089. The spectral indices and fluxes were left free for sources within , while sources from to have their parameters fixed to their catalog value. A first unbinned likelihood fit was performed for the events collected within almost 4 months of data from March 22, 2015 to July 19, 2015 (MJD 57103–57223) in the energy range between 100 MeV and 800 GeV. The sources with test statistic (TS; Mattox et al., 1996) below 5 were removed from the model. Next, the optimized output model was used to produce the light curves and spectra of PKS 1510–089 in different time bins (from 1 day to 3 hours) and energy ranges (E100 MeV, E1 GeV). For the calculation of the light curves, all sources were fixed in the model except PKS 1510–089 for which both the flux normalization and the spectral index were left free and modeled as a power-law. For the calculation of the spectral points, in addition also the spectral index of PKS 1510–089 was fixed to its best fit value during the considered time period for which the spectral points are estimated. The normalization of the Galactic and isotropic diffuse emission models was left to vary freely during the calculation of both the light curves and the spectra.

2.3 Swift-XRT and UVOT

The multi epochs (16 individual pointings) event-list obtained by the X-ray Telescope (XRT) (Burrows et al., 2004) on-board of Swift satellite during the period of May 11 to May 25, 2015 (MJD 57153–57167) with the total exposure time of 26.6 ks were downloaded from the publicly available database table SWIFTXRLOG (Swift-XRT Instrument Log). The individual exposures ranged from 0.6 to 4 ks. They were processed using the HEASOFT package version 6.18. All the observations from this period had been performed in photon counting (PC) mode. The source region was defined as a circle of 20 pixel (47”) radius at the center of the source, while the background region was defined by an annulus centered at the source with inner and outer radii of 40 (94”) and 80 pixels (188”), respectively. The source and background spectra were extracted using XSELECT task (v2.4c). The source spectrum count rate does not exceed 0.5 counts/s in any of the observation epochs. Therefore, no pile-up correction was needed. For the light curve analysis we have combined 3 pairs of epochs (MJD 57157, 57161, 57163) separated by 2 hr.

The xrtexpomap task (v0.2.7) was used to correct the flux loss due to the fact that some of the CCD pixels were not used during the data collection. The xrtmkarf task (v0.6.3) took into account vignetting and bad pixels. The grppha task was used to group source spectra in a way that each bin contains 20 counts. XSPEC task (v12.9.0i) was used to calculate the flux and power law model spectral parameters using fixed equivalent Galactic hydrogen column density of [cm] (Kalberla et al., 2005).

The second instrument on board the Swift satellite, the Ultraviolet/Optical Telescope (UVOT, Poole et al., 2008) was used to monitor the flux of the source in the 180–600 nm wavelength range. Following Raiteri et al. (2010) we used an iterative procedure for the data calibration, where the effective wavelength, counts-to-flux conversion factor and Galactic extinction for each filter were calculated by taking into account the filter effective area and source spectral shape. Out of the 16 pointings in the investigated time period, 8 were taken with a full set of filters (v, b, u, w1, m2, w2). For the SED modelling we used pointings on MJD 57160 (with u and w2 filters) contemporaneous with the MAGIC flaring state, and on MJD 57165 (all filters available) during the post-flare MAGIC observations.

2.4 Optical photometry and polarization

PKS 1510–089 is regularly monitored as part of the Tuorla blazar monitoring program444http://users.utu.fi/kani/1m in R-band using a 35 cm Celestron telescope attached to the KVA (Kunglinga Vetenskapsakademi) telescope located at La Palma. The data analysis was performed with the semi-automatic pipeline using the standard analysis procedures (Nilsson et al. in prep). The differential photometry was performed using the comparison star magnitudes from Villata et al. (1997). The magnitudes were corrected for the Galactic extinction using values from Schlafy & Finkbeiner (2011).

The optical polarization observations were performed with a number of instruments: Nordic Optical Telescope (NOT), Steward Observatory, Perkins Telescopes, RINGO3, AZT-8 and LX-200. The NOT polarimetric observations were done with ALFOSC in the R-band using the standard setup for linear polarization observations (lambda/2 retarder followed by a calcite). The data were analyzed using the standard procedures with semi-automatic software as in Hovatta et al. (2016). The Steward polarimetric observations were obtained as part of an ongoing monitoring program of gamma-ray-bright blazars in support of the Fermi mission555http://james.as.arizona.edu/~psmith/Fermi. The observations were performed in the 5000-7000 Å  band. The data analysis pipeline is described in Smith et al. (2009). Polarimetric R-band observations were also provided by the 1.8 m Perkins telescope at Lowell Observatory equipped with PRISM (Perkins Reimaging System). The data analysis was done following the standard procedures as in Chatterjee et al. (2008). Polarization observations were also taken with the RINGO3 polarimeter (Arnold et al., 2012) on the fully robotic and autonomous Liverpool Telescope on La Palma, Canary Islands (Steele et al., 2004) as part of the Liverpool blazar monitoring campaign (see Jermak et al., 2016), in collaboration with the Monitoring AGN Polarimetry at the LIverpool Telescope (MAPLIT) program. Simultaneous observations in the ’blue’, 350– 640 nm; ’green’, 650–760 nm; and ’red’, 770– 1000 nm passbands were taken using a rapidly rotating (once per 4 seconds) polaroid which modulates the incoming beam of light in 8 rotor positions. For this work we use only the 650–760 nm measurements, which are the closest to the R-band used in the mentioned above polarization instruments. The beam is simultaneously split by 2 dichroic mirrors into three electron multiplying CCD cameras (EMCCDs). The combination of the flux from the 8 rotor positions using equations from Clarke & Neumayer (2002) can be used to find the linear Stokes parameters, which were used to calculate the degree and angle of polarization. Finaly, additional optical polarimetric data are reported here from the 70cm AZT-8 telescope (Crimea) and the 40cm LX-200 telescope (St.Petersburg), both equipped with nearly identical imaging photometers-polarimeters (Larionov et al., 2008). Polarimetric observations were performed using two Savart plates rotated by relative to each another. By swapping the plates, the observer can obtain the relative Stokes and parameters from the two split images of each source in the field. Instrumental polarization was found via stars located near the object under the assumption that their radiation is unpolarized. The electric vector position angle (EVPA) was corrected for the ambiguity by minimizing the difference to the closest data point unless there is a gap of over 14 days.

2.5 Infrared

PKS 1510–089 is monitored by a number of IR instruments. We have used the publicly available data in B, V, R, J and K bands from the Small and Moderate Aperture Research Telescope System (SMARTS) instrument located at Cerro Tololo Interamerican Observatory (CTIO) in Chile. The data reduction and calibration is described in Bonning et al. (2012). We converted the magnitudes into flux units using Bessell et al. (1998) and corrected for the interstellar dust absorption following Schlafy & Finkbeiner (2011).

The observations at Teide Observatory (Canary Islands) were obtained with the 1.52 m Carlos Sanchez Telescope (TCS), using the near-infrared camera CAIN during the nights of MJD 57162–57174. This camera is equipped with a 256 256 pixels NICMOS-3 detector providing a scale of 1”/pixel. Data were acquired in the three filters J, H and Ks. Observations were performed using a 5-point dither pattern (repeated twice) in order to facilitate a proper sky background subtraction. At each point, the exposure time was about 1 min, split in individual exposures of 10 s in the J filter and 6 s in the H and Ks filters to avoid saturation by sky brightness. Image reduction was performed with the caindr package under the IRAF environment666Image Reduction and Analysis Facility, http://iraf.noao.edu/. Data reduction includes flat-fielding, sky subtraction, and the shift and combination of all frames taken in the same dither cycle. Photometric calibration was made based on field stars from the 2MASS catalogue (Cutri et al., 2003). The photometric zero point was determined for each frame by averaging the offset between the instrumental and the 2MASS magnitudes of the catalogue. Deviant stars were excluded and typical errors remained below 5%.

We have also used IR photometry data obtained with a 1.2m telescope of Mt Abu InfraRed Observatory (MIRO), India, mounted with the Near Infrared Camera and Spectrograph (NICS) equipped with 1024x1024 HgCdTe Hawaii array detector. The field of view is 8’8’ with a pixel scale of 0.5”/pixel. The observations on PKS 1510–089 were performed with a 4-position dither with offsets of 30 arcsec, keeping the comparison stars777https://www.lsw.uni-heidelberg.de/projects/extragalactic/charts/1510-089.html 1, 2, 3, 4 and 6 in the field of the source. The sky and dark contributions were removed using these dithered images and aperture photometry was performed using standard procedures under IRAF (see Banerjee & Ashok (2012) for details of data reduction and analysis). The source magnitudes in J, H and Ks bands were calibrated using correction factors obtained using weighted average of the standard values of comparison stars mentioned above.

2.6 Radio

Radio monitoring observations were performed with Metsähovi Radio Telescope operating at 37 GHz frequency. The instrument and data reduction procedures are described in Teräesranta et al. (1998) and Aleksić et al. (2014).

The quasar PKS 1510–089 has been observed within a sample of gamma-ray blazars that the Boston University (BU) blazar group monitors with the Very Long Baseline Array (VLBA) approximately monthly at 43 GHz (the VLBA-BU-BLAZAR project). The observations of PKS 1510–089 are usually performed during 9 short ( min) scans within a span of 7-8 hours. The data were calibrated at the VLBA DiFX correlator and reduced using the Astronomical Image Process System (AIPS) and Difmap software packages, as described in Jorstad et al. (2005). The calibrated data are available online (www.bu.edu/blazars/VLBAproject.htm). We analyzed the data obtained from February 2015 to April 2016 (12 epochs). We modeled the total intensity images by components with circular gaussian brightness distributions. For each component we determined: flux density, distance and position angle (PA) with respect to the VLBI core888PA is measured starting from the positive direction of Declination axis (PA) increasing in the positive direction of Right Ascension axis (PA), and size. 43 GHz core is expected to be located at the distance of 6.5 pc from the central engine of PKS 1510–089 (see Pushkarev et al., 2012; Aleksić et al., 2014). A map of the parsec scale jet of the quasar formed from 20 stacked images over 6 years of VLBA observations at 43 GHz is plotted in Fig. 1 (individual images can be found at the BU blazar group website http://www.bu.edu/blazars/VLBA_GLAST/1510.html)

Figure 1: A stacked map of 43 GHz total (black contours) and polarized (color scale) intensity images of the inner, pc-scale jet of PKS 1510–089 with the direction of electric field vector polarization denoted by black line segments (contour levels are indicated on the top, polarized flux levels are shown by the color bar to the right). All images have been convolved with the same Gaussian beam, shown in the lower right corner.

The image shows the VLBI core, which is the brightest compact feature located at the southeast end of the jet. The core is used as a reference point in the stacking procedure, since it is assumed to be stationary. The stacked image reveals the full opening angle of the jet, as well as the location of the jet axis. As can be inferred from Fig. 1, the jet axis is along PA, while the projected opening angle is . The core at 43 GHz is only partially optically thick, which is supported by synchronous optical and radio core polarization variability in a number of blazars (D’Arcangelo et al., 2007; D’arcangelo et al., 2009). Fig. 1 reveals a shift between the total and polarized intensity peaks and the complex structure of the polarized emission in the core. This favors the hypothesis that the core of PKS 1510–089 is a recollimation shock, which has been inferred previously from the polarization structure in other blazars (e.g. Cawthorne et al., 2013).

3 Results

In Fig. 2 we show the multiwavelength light curve of PKS 1510–089 during the May 2015 outburst in the investigated period of MJD 57151–57174.

Figure 2: Multiwavelength light curve of PKS 1510–089 during the May 2015 flare. From top to bottom: Nightly gamma-ray flux above 150 GeV from MAGIC (the dashed line shows the average emission in Feb-Apr 2012, Aleksić et al., 2014); Fermi-LAT flux above 0.1 GeV in 6 h binning, and the corresponding spectral index (the dashed line shows the average emission from the 3FGL catalog, Acero et al., 2015); X-ray spectral flux (filled circles) and spectral index (empty circles) measured by Swift-XRT; polarization percentage and polarization angle measured by NOT, Steward, Perkins, RINGO3, AZT-8 and LX-200 (see legend); optical emission in R band (KVA, SMARTS) and UV emission in w2-band (Swift-UVOT); IR emission in J band (SMARTS, MIRO-NICS, TCS); radio observations by Metsähovi at 37 GHz. Data from IR up to UV are corrected for the Galactic absorption. The red and blue shaded regions show the Period A and Period B, respectively, for which the spectral modelling is performed.

Following the observations of MAGIC, we define two observation periods: Period A (MJD 57160-57161) and Period B (MJD57164-57166). The multiwavelength SED of both periods is investigated.

3.1 Magic

The MAGIC light curve (see top panel of Fig. 2) shows clear variability, with the highest flux observed during the two nights of Period A. The hypothesis of constant flux during all 5 observation nights of MAGIC can be clearly rejected, with a chance probability of . Even allowing for a 20% variable systematic uncertainty on individual night fluxes (motivated by variable systematic uncertainty estimate given in Aleksić et al., 2016b, rescaled to a softer source) we still obtain a small value of chance probability of for the flux to be constant. The flux during Period A is times larger than the one observed during the previous detection by MAGIC in 2012 (Aleksić et al., 2014). In the following observations during Period B, the VHE gamma-ray flux decreased to a level consistent with the detection in 2012.

In order to search for a possible short time variability we binned the light curve during the Period A into 20 min bins (see Fig. 3).

Figure 3: Light curve above 150 GeV obtained with the MAGIC telescopes during the flare in Period A. The fluxes are computed in 20 min bins. The black line shows the constant flux fit, . The gray band shows the 95% C.L. interval allowing for 20% variable systematic uncertainty (see text for details).

No variability is detected at such a time scale. Fitting the light curve with a constant flux hypothesis we obtain . We estimate the maximum variability which can be hidden by the uncertainties of the measurement by computing for each 20 min light curve bin a 95% C.L. interval on the flux using Rolke et al. (2005) prescription. We include a 20% variable systematic uncertainty in those calculations. By comparing the least constraining upper edge of the 95% C.L. interval with the average flux from those two nights we obtained that the flux did not increase by more than a factor 3.5 on time scales of 20 min.

For spectral analysis we have combined the two nights of Period A. The obtained VHE gamma-ray spectrum of PKS 1510–089 is shown in Fig. 4.

Figure 4: Spectral energy distribution of PKS 1510–089 constructed from the MAGIC data gathered on MJD 57160 and 57161 (Period A). The observed spectrum is shown in red empty squares with inclined hashing, and the EBL-deabsorbed spectrum according to Domínguez et al. (2011) model with magenta filled circles with vertical hashing. The hashing shows the uncertainty of the forward folding with a power law. The SED constructed from the data taken on MJD between 57164 and 57166 (Period B) is shown as empty (observed) and filled (EBL-deabsorbed) blue circles. For comparison, MAGIC measurements performed in Feb-Mar 2012 and H.E.S.S. measurements from Mar 2009 (Abramowski et al., 2013) are shown as gray diamonds and stars respectively.

It can be described by a power law, , with and . The spectral parameters are obtained using a forward folding method (Albert et al., 2007). The statistical uncertainty on the spectral index is much larger than the typical systematic uncertainty of of the observations performed with the MAGIC telescopes. The systematic uncertainty on the flux normalization does not include the uncertainty of the energy scale of MAGIC which for this data set we estimate as , slightly larger than given in Aleksić et al. (2016b), due to the need of LIDAR correction of Calima-affected data. Correcting for the absorption of TeV gamma-rays due to the interaction with the extragalactic background light according to Domínguez et al. (2011) model, an intrinsic spectrum with normalization of and index are obtained. The spectral shape is marginally consistent (but with large uncertainties) with the previous measurements by H.E.S.S. (observed slope , Abramowski et al., 2013) and MAGIC (intrinsic slope , Aleksić et al., 2014).

For comparison, we have also reconstructed the average flux from the MAGIC measurements performed during Period B. The combination of weaker emission and observations performed during higher atmospheric transmission results in the energy range of the reconstructed spectrum shifted to lower energies. The observed flux in Period B is at a similar level as the one detected by MAGIC in 2012. We obtained the observed and intrinsic VHE spectral slopes of PKS 1510–089 during Period B of and respectively.

3.2 Fermi-Lat

The GeV gamma-ray flux of PKS 1510–089 is highly variable in the investigated period. A few individual flares are visible, with time scales of a few days. Due to the short-term gamma-ray variability, in order to get a spectrum comparable to the MAGIC observations, for SED analysis we selected the events observed by the LAT within 6 h centered in each of the MAGIC observations. We have calculated two GeV spectra, which correspond to the two different states of the source contemporaneously to the MAGIC observations in Period A and Period B, respectively. The best description of the GeV spectrum measured by the LAT in Period A is a power-law with spectral index 2.200.07 and a flux above 100 MeV of (6.80.5) (). During Period B, the Fermi-LAT measured spectrum is best described by a power-law with a similar spectral index of 2.170.08, but significantly lower flux, (3.70.3) (). For comparison, the 3FGL flux above 100 MeV is (0.940.01) (Acero et al., 2015). Therefore PKS 1510–089 has reached a factor of 7 to 4 times its average flux over the first 4 years of Fermi-LAT observations during the two epochs for which the spectral analyses has been performed. The GeV flux was, however, still a factor of about 2 smaller than the daily peak flux observed in 2011 (Saito et al., 2015). The Fermi-LAT spectrum during both Period A and B was slightly harder than in the neighbouring days (see also Section 3.5). In the case of Period B there is a weak hint of a spectral curvature. The likelihood ratio test gives preference ( for one additional degree of freedom) of the log-parabola shape over the power-law spectral model. The corresponding photon index can be described as .

To characterize the variability in the Fermi-LAT light curve we calculated the Power Density Spectrum (PDS) both for the 2015 flare epoch (MJD 57143.9375–57182.9375) with a 3 h binning as well as for a mission lifetime at the time of the analysis (MJD 54682.655–57484.655) with 1 day binning. In both cases the flux calculation was performed in the 0.1-800 GeV energy range.

The estimated white noise level, based on the data error values was subtracted from each PDS. The PDS is rebinned in logarithmic frequency intervals, and normalized to variance per frequency unit divided by the square of the mean flux.

Figure 5: Power Density Spectrum normalized to variance per frequency unit for the mission long Fermi-LAT light curve (black) and for the 2015 flare epoch (magenta).

The PDS level, computed with such normalization, is similar for the 2015 flare and for the mission lifetime light curve (see Fig. 5). This suggests that the fractional variability is the same and presumably driven by the same variability process. The shape of the overall PDS can be described by a power law for frequencies above 0.01 day. A power-law fitted to the 0.007-0.5 day frequency range PDS gives an index of 1.14  0.07 for full data set and 0.97  0.30 for the 2015 flare. The uncertainty value of the fit is based on the scatter of the measurements with respect to the fitted line only and therefore may be underestimated. The power law index is similar to that of other FSRQs (Ackermann et al., 2011). The stochastic nature of the variability together with the limited observation length lead to a large uncertainty in the PDS at the lowest frequencies.

3.3 Swift-Xrt

The X-ray flux, measured by Swift-XRT, shows a gradual decrease of the observed flux during the studied period. Except for the first point at MJD 57153.6, which happened before the two large Fermi-LAT flares the X-ray flux can be much better described by an exponential decline () than a constant value (). Similarly, the corresponding X-ray spectral indicies are better described by a linear softening of the spectrum () than a constant ().

The flux is marginally higher during MJD 57156-57162, when a broad Fermi-LAT flare happened. XRT observations on MJD 57166-57168, during the next broad Fermi flare did not show a clear increase of the X-ray flux. For spectral analysis we combined the 4 pointings taken during Period A, each of them being 6-8 hr distant from the MAGIC observations. The X-ray spectrum in Period A can be marginally well described () with a power law with an index of . Swift-XRT observations performed during Period B resulted in a spectrum which can be well fitted with a power law with a significantly softer index of with , or by a curved spectrum with an index of with .

3.4 IR, Optical and UVOT

The optical-UV SED of PKS 1510–089 consists of an unpolarized, quasi-stable accretion disk component and non-thermal jet emission. The variability and polarization of the jet component might be diluted by the accretion disk component. The effect is strongest in the UV bands and weak in the R-band, where the polarization is measured. It does not affect the timing of the observed polarization variability.

The optical emission of PKS 1510–089 during the investigated period is clearly variable (with a factor 2 difference between the lowest and highest flux). The optical variability does not strictly follow the gamma-ray one. In the optical R-band, and to a lesser extent also in UV w2-band, the flux was slowly raising throughout May 2015. Similarly as in the optical range, the IR flux doubled in a days time before Period A. At the end of the observation period it returned to the pre-flare state.

Throughout the investigated period, a smooth rotation of optical EVPA by is happening. The rotations of optical polarization angle accompanied also the 2009 and 2012 gamma-ray flaring states (Marscher et al., 2010; Aleksić et al., 2014). However, the rotations of the EVPA seem very common in PKS 1510–089, e.g., recent work by Jermak et al. (2016) identified a rotation also in the 2011 data, therefore further data are needed to firmly associate them with the emission of VHE gamma rays. The low percentage of polarization seems to be typical for this source (Jermak et al., 2016). Indeed, during Period A the polarization percentage is low (). It triples between the period A and B. Also in the few days before the Period A (i.e. during the raising part of the Fermi-LAT flare that culminated with detection of VHE gamma-ray emission by MAGIC) a higher polarization was observed. The polarization behaviour during the 2015 flaring period agrees with what one would expect from a knot following a spiral path through a mainly toroidal magnetic field (Marscher et al., 2010). An alternative explanation might be the light travel time effects within an axisymmetric emission region pervaded by predominately helical magnetic field (Zhang et al., 2015). The evolution of the polarization is probably further complicated by superposition of individual flares seen in Fermi-LAT.

3.5 Radio

PKS 1510–089 shows moderate variability in Metsähovi observations performed at 37 GHz in May 2015. The sampling is, however, rather sparse, and especially the local peaks of the GeV flux are mostly not covered by the observations. The observations on MJD 57161, during Period A are burdened with a larger uncertainty due to adverse atmospheric conditions.

Fig. 6 shows the light curve of the core from February 2015 to April 2016.

Figure 6: VLBA light curves of the core (circles) and knot (stars) at 43 GHz. Vertical dashed line and gray shaded region show the zero separation epoch of and its uncertainty.

The light curve reveals a significant gradual increase of flux in the second half of 2015 that is most likely connected with a disturbance (knot ) detected in the VLBA images starting in December 2015 (Fig. 7).

Figure 7: Total intensity images of PKS1510-089 core region at 43 GHz, with a global peak intensity of  Jy/beam and 0.15 mas FWHM circular Gaussian restoring beam (bottom right circle). The solid and dashed lines designate positions of the VLBI core and respectively across the epochs.

Knot is bright and relatively slow, with an apparent speed =5.31.4 c. According to currently available data, the knot was ejected on MJD 5723052 (see Fig. 8).

Figure 8: Separation of knot from the core.

A similar behavior has also been observed during a high gamma-ray state in Feb-Apr 2012, when the emerging of a new radio knot () from the core was associated with a VHE outburst (Aleksić et al., 2014). The large uncertainty in the ejection time of does not allow us to associate it firmly with a particular GeV flare, as the source showed activity in Fermi-LAT close to the time of separation (note e.g. the hard flare on MJD 57245 in Fig. 9).

Figure 9: Fermi-LAT light curve above 0.1 GeV (top panel) binned in 3 day intervals and the corresponding evolution of spectral index (middle panel) and the optical R-band flux observed by KVA (bottom panel). Red and blue hashed bands show the Period A and B respectively, while the gray band shows the uncertainty on the knot separation time.

Interestingly, the position angle of is +50, while the typical projected direction of the pc-scale jet is 30 (see Fig. 1), along which so far 7 knots have been observed (see e.g. Jorstad et al., 2005; Aleksić et al., 2014). On the other hand, knot , ejected in October 2011 (before the 2012 VHE outburst), had a similar PA to the one observed now for close to the core, but turned toward the usual pc-scale jet direction a few months later. We note that is a factor of a few times slower than , and so may eventually follow a similar trajectory. The slower apparent speed and brighter flux of suggest that the velocity vector of this disturbance is closer to the line of sight than , causing it to have a higher Doppler factor.

4 Modelling of the spectral energy distribution

The gamma-ray emission of FSRQs is usually explained as the effect of the inverse Compton scattering of electrons on a radiation field external to the jet (see, e.g. Sikora et al., 1994; Ghisellini et al., 2010), the so-called external Compton (EC) scenario. The radiation field can originate from the accretion disk, broad line region (BLR) or the dust torus (DT). This scenario has been applied to explain the emission of PKS 1510–089 in its previous flaring episodes (Abdo et al., 2010). The origin of the radiation field is closely connected with the location of the emission region. Moreover, the observed VHE gamma-rays escaping from the emission region suggest that the emission region is located outside the BLR in order to escape the absorption by pair production process on BLR photons (Abramowski et al., 2013; Aleksić et al., 2014). Dotson et al. (2015) investigated the 2009 GeV flares of PKS 1510–089 using the energy dependence of the flare decay time as the diagnostic for the emission zone location. They claimed that two of the flares happened around the distance of the DT, while the other two came from the vicinity of VLBI core. On the other hand, the modelling of GeV flares seen from PKS 1510–089 in 2011 placed the emission region at the distance of  pc from the black hole, with the EC process happening on a mixture of BLR and DT radiation fields (Saito et al., 2015). A similar location ( pc distance from the black hole) with EC mainly on DT was invoked to explain the high optical and gamma-ray state observed in the beginning of 2012 (Aleksić et al., 2014). The broadband emission could be also explained by a much more distant ( pc) region for a 2-zone model in which the jet consists of the inner spine and outer sheath layer (Aleksić et al., 2014; MacDonald et al., 2015).

Figure 10: Multiwavelength spectral energy distribution of PKS 1510–089 in Period A (red symbols) and B (blue symbols). The red and the blue curves report the result of the emission model for the two periods. The black dashed and long-dashed lines show the adopted emission for the accretion disk and the dusty torus, respectively. For comparison, the dashed gray line shows the model derived for the SED in 2012 (from Aleksić et al. 2014). Historical measurements (ASDC, see http://www.asdc.asi.it/) are shown as gray points.

In Fig. 10 we present the two spectral energy distributions of PKS 1510–089 constructed from the data taken in Period A and Period B, corresponding to high and low gamma-ray flux, respectively. As can be seen, most of the flux variation ( a factor of 2–3) occurs in the Fermi-LAT and MAGIC bands. The low-energy flux (optical, X-rays) is almost constant between the two periods. Remarkably, the high energy peak during the period B has a very similar level to the 2012 high state, in spite of the IR–UV emission being a factor of higher.

We model these SEDs of PKS 1510–089 in the framework of the same one-zone model as discussed in Aleksić et al. (2014). To explain the sub-TeV emission observed by MAGIC we assume that the emission region is located beyond the BLR radius, where the external photon field is dominated by the thermal IR radiation of the DT.

For the setup of the model we assume the scaling laws and the prescriptions given in Ghisellini & Tavecchio (2009). Specifically, the radius of the BLR is given by cm and that of the torus by cm, where is the accretion disk luminosity in units of erg s. We calculate the IR radiation energy density assuming that a fraction of the disk radiation is intercepted and reprocessed by the torus heated to 1000 K. Similarly, the BLR intercepts of the disk radiation. Note that, with these prescriptions, the energy densities of the BLR and torus radiation fields do not depend on the disk luminosity, since they depends on the constant ratio . Assuming the same disk luminosity erg s as in Aleksić et al. (2014) the scaling law of Ghisellini & Tavecchio (2009) (based on reverberation mapping measurements of BLR size, see e.g. Bentz et al., 2009) allows us to infer a BLR radius of cm.

We fix the distance of the emission region from the base of the jet to cm. The emission region is filling the whole cross section of the jet, which for an assumed jet semi-aperture angle  rad, results in the emission region radius cm. Such size of the emission region is in line, even for moderate values of the Doppler factor, with the constraints set by the few day time scale variability observed by MAGIC.

The apparent superluminal motion of radio component K15 puts limits on . The large uncertainty in the separation time from the radio core does not allow us to firmly associate such limits with the speed of the emission region responsible for the emission in investigated periods A and B. Moreover, the apparent speed is measured over a much longer time scale, during which the emission region might have decelerated. On the other hand the beaming of the emission is constrained by the observed luminosity of the dominating high energy peak. It can be estimated as , where is the energy density of the external radiation field measured in the frame of reference of the blob, is the number density of the electrons, is the volume of the emission region measured in its own frame of reference and is the Doppler factor of the blob. The average squared Lorentz factor of the electrons can be approximated as if the distribution starts from and follows an index of up to the break of . Assuming the total kinetic power of a jet composed of cold protons with number density of is . Combining the two formulas, for a jet with we obtain . The observed luminosity of of the EC peak requires the jet Lorentz factor to be at least 10 even for the case of . For the modeling we apply the same values as used in Aleksić et al. (2014) for the jet bulk Lorentz factor and . For such a bulk Lorentz factor the assumed jet semi-aperture angle of  rad is broader than suggested by radio observations (see e.g. Jorstad et al., 2005). The radio observations are sensitive to the jet opening angle at a distance equal to or greater than the location of the radio core, i.e. a few pc from the base of the jet. On the other hand, the emission region assumed in the modelling is located closer to the base of the jet, where the jet opening angle can be larger, as observed e.g. in the case of M87 radio galaxy Asada & Nakamura (2012).

[1] [2] [3] [4] [5] [6] [7] [8] [9]
Period A 1
Period B 1

Table 1: Input model parameters for the models of epochs A (MJD 57160-57161) and B (57164-57166) of PKS 1510–089 in Fig. 10. [1], [2], [3] and [4]: minimum, cooling break, acceleration break and maximum electron Lorentz factor respectively. [5], [6] and [7]: slope of the electron energy distribution below , between and , and above . [8]: magnetic field [G]. [9]: normalization of the electron distribution in units of cm.

Having fixed these values, the free parameters of the model are only the intensity of the magnetic field and those describing the electron energy distribution. Hence, the observed variability according to this scenario is caused by changes in the conditions of the plasma flowing through the shock region. Since we assume that the emission occurs outside the intense radiation field of the BLR, the IC emission and the radiative losses of the emitting electrons are dominated by the scattering off the IR radiation field of the torus. As the energy density of this radiation field is relatively low, the cooling of the electrons is not very effective. A simple calculation shows that the Lorentz factor at which the cooling time equals the dynamical time is of the order of . If we assume that the electrons are injected starting from , with a power law distribution with slope , in equilibrium we would expect a break at , above which the slope of the distribution would steepen to . However, such a break could not properly describe the SED, since the required break (estimated using the X-ray and the MAGIC slopes) is larger (even taking into account Klein-Nishina effects on the spectrum in the MAGIC energy range). To reproduce the SED we therefore assume a scenario in which the electron energy distribution is a double broken power law; with a cooling break at , and a second break connected e.g. with the acceleration process at . The particles are injected into the emission region with a broken power-law energy distribution with slope and and break Lorentz factor . In equilibrium conditions, the electron energy distribution displays three power-laws with slopes from to the cooling electron Lorentz factor , from to and above . (see e.g. Ghisellini et al., 2002). Note that similar hard spectra has been also postulated in modelling of blue blazars (Ghisellini et al., 2012).

The values of these parameters required to reproduce the SEDs are reported in Table 1. In fact, the difference in the emission between Period A and B can be explained with a relatively small change in the fit parameters, mainly a slightly stronger magnetic field and lower maximal and break energies of the electrons

The model discussed here has some caveats. As typically happens in blazar modelling, the radio points overshoot the model line, which has a strong low energy cutoff due to synchrotron self-absorption. This emission is normally explained to occur from much larger regions farther along the jet. Moreover the variations of the optical emission occur on longer time scales than the flares observed in GeV. Hence, additional optical emission might be produced by the high energy electrons swept with the flow farther along the jet (up to a few pc from the base of the jet). In those regions the external radiation field density would be too low to efficiently produce high energy photons via inverse Compton process, turning synchrotron emission into the dominant radiation process. In spite of adiabatic energy losses that electrons can suffer, the total observed synchrotron radiation might still slowly increase due to aggregation of electrons from multiple individual flares. In fact, if the electrons reach the radio core (and beyond) located at the distance of they could be responsible for the emission of a new radio knot. The separation time of such a knot from the core could be estimated to occur  days after the gamma-ray flare. The extent of the core might shorten the time delay before the knot-core interaction starts. We encourage further trials of modelling of the observed high state with scenarios employing emission from larger length of the jet than a single active region.

5 Discussion and Conclusions

Using the MAGIC telescopes data we have detected enhanced VHE gamma-ray emission from the direction of PKS 1510–089 during the high optical and GeV state of the source in May 2015. It is the first time that VHE gamma-ray variability was detected for this source. The spectral shape is, however, consistent within the statistical uncertainties with the previous measurements of the source. During May 2015 the IR through UV data showed a gradual increase of the flux, while the flux in the X-ray range was slowly decreasing.

The May 2015 data revealed, similarly to the 2012 data, that the enhanced VHE gamma-ray emission occurred during the rotation of the optical polarization angle. Also, similarly to other gamma-ray flares, an ejection of a new radio component was observed, however, with an unusual position angle. This suggests that the association of VHE gamma-ray emission with rotation of EVPA and ejection of a new radio component might be a common feature of PKS 1510–089.

The source was modelled with the external Compton scenario. The evolution of the state of the source from the VHE gamma-ray flare to a lower emission (at the level of 2012 high state reported in Aleksić et al., 2014) can be explained by relatively small changes in the conditions of the plasma flowing through the emission region. The presented scenario is, however, only one possible solution. As discussed in Aleksić et al., 2014, if we assume that the VHE flaring is indeed connected to the ejection of the new component (in this case ) from the VLBA core and the rotation of the optical polarization angle, it would be natural to assume a single emission region located far outside the dusty torus. In this case the seed photons could be provided by the sheath of the jet and this scenario has been shown to provide feasible description of the previous flaring epochs of PKS 1510–089 (Aleksić et al., 2014; MacDonald et al., 2015). The VHE gamma-ray variability with time scale seen during the 2015 outburst puts constraints on the size, and therefore also on the location of the emission region. Assuming that the spine of the jet fills a significant fraction of the jet (as in Aleksić et al., 2014), the location of the emission region cannot be farther than  pc. Therefore, for placing the emission region at the radio core a high Doppler factor and a narrow jet are needed. In fact such low values of the jet extension, (Jorstad et al., 2005) and (Pushkarev et al., 2009) at the radio core are reported by the radio observations. Alternatively, the inner spine can be much narrower than the whole jet, as suggested by MacDonald et al. (2015).

To further study the connection of VHE emission with events in lower frequencies, long-term monitoring data are needed and this question will be addressed in a future publication. With the detection of this flare from PKS 1510–089, VHE gamma-ray variability (on times scales varying from tens of minutes to days) has been observed in all FSRQs known in VHE gamma rays. Fast-varying VHE gamma-ray emission is common among the brightest gamma-ray FSRQs. As it seems that most of the gamma-ray FSRQs can only be detected during these flares, it is indeed not surprising that only a handful have been detected in VHE gamma-rays.

Acknowledgements.
The MAGIC collaboration would like to thank the Instituto de Astrofísica de Canarias for the excellent working conditions at the Observatorio del Roque de los Muchachos in La Palma. The financial support of the German BMBF and MPG, the Italian INFN and INAF, the Swiss National Fund SNF, the he ERDF under the Spanish MINECO (FPA2015-69818-P, FPA2012-36668, FPA2015-68278-P, FPA2015-69210-C6-2-R, FPA2015-69210-C6-4-R, FPA2015-69210-C6-6-R, AYA2013-47447-C3-1-P, AYA2015-71042-P, ESP2015-71662-C2-2-P, CSD2009-00064), and the Japanese JSPS and MEXT is gratefully acknowledged. This work was also supported by the Spanish Centro de Excelencia “Severo Ochoa” SEV-2012-0234 and SEV-2015-0548, and Unidad de Excelencia “María de Maeztu” MDM-2014-0369, by grant 268740 of the Academy of Finland, by the Croatian Science Foundation (HrZZ) Project 09/176 and the University of Rijeka Project 13.12.1.3.02, by the DFG Collaborative Research Centers SFB823/C4 and SFB876/C3, and by the Polish MNiSzW grant 745/N-HESS-MAGIC/2010/0. The Fermi LAT Collaboration acknowledges generous ongoing support from a number of agencies and institutes that have supported both the development and the operation of the LAT as well as scientific data analysis. These include the National Aeronautics and Space Administration and the Department of Energy in the United States, the Commissariat à l’Energie Atomique and the Centre National de la Recherche Scientifique / Institut National de Physique Nucléaire et de Physique des Particules in France, the Agenzia Spaziale Italiana and the Istituto Nazionale di Fisica Nucleare in Italy, the Ministry of Education, Culture, Sports, Science and Technology (MEXT), High Energy Accelerator Research Organization (KEK) and Japan Aerospace Exploration Agency (JAXA) in Japan, and the K. A. Wallenberg Foundation, the Swedish Research Council and the Swedish National Space Board in Sweden. The Metsähovi team acknowledges the support from the Academy of Finland to our observing projects (numbers 212656, 210338, 121148, and others). This paper has made use of up-to-date SMARTS optical/near-infrared light curves that are available at www.astro.yale.edu/smarts/glast/home.php. MIRO is operated by Physical Research Laboratory, Ahmedabad with support from Dept of Space, government of India. This article is based on observations made with the 1.5 m telescope Carlos Sánchez operated by the Instituto de Astrofisica de Canarias in the Teide Observatory. The data presented here were obtained in part with ALFOSC, which is provided by the Instituto de Astrofisica de Andalucia (IAA) under a joint agreement with the University of Copenhagen and NOTSA. The Liverpool Telescope is operated on the island of La Palma by Liverpool John Moores University in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias with financial support from the UK Science and Technology Facilities Council. St.Petersburg University team acknowledges support from Russian RFBR grant 15-02-00949 and St.Petersburg University research grant 6.38.335.2015. The BU group acknowledges support by NASA under Fermi Guest Investigator grants NNX11AQ03G and NNX14AQ58G. The VLBA is an instrument of the National Radio Astronomy Observatory. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. We would like to thank the anonymous journal referee for the comments on the manuscript.

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