DISCOVERY OF THE HARD SPECTRUM VHE γ-RAY SOURCE HESS J1641–463

arXiv:1408.5280v1 [astro-ph.HE] 22 Aug 2014 Discovery of the hard spectrum VHE γ-ray source HESS J1641−463 H.E.S.S. Collaboration, A. Abramowski1 , F. Aharonian2,3,4 , F. Ait Benkhali2 , A.G. Akhperjanian5,4 , E.O. Angüner6 , M. Backes7 , S. Balenderan8 , A. Balzer9 , A. Barnacka10,11 , Y. Becherini12 , J. Becker Tjus13 , D. Berge14 , S. Bernhard15 , K. Bernlöhr2,6 , E. Birsin6 , J. Biteau16,17 , M. Böttcher18 , C. Boisson19 , J. Bolmont20 , P. Bordas21 , J. Bregeon22 , F. Brun23 , P. Brun23 , M. Bryan9 , T. Bulik24 , S. Carrigan2 , S. Casanova25,2 , P.M. Chadwick8 , N. Chakraborty2 , R. Chalme-Calvet20 , R.C.G. Chaves22 , M. Chrétien20 , S. Colafrancesco26 , G. Cologna27 , J. Conrad28,29 , C. Couturier20 , Y. Cui21 , I.D. Davids18,7 , B. Degrange16 , C. Deil2 , P. deWilt30 , A. Djannati-Ataı̈31, W. Domainko2 , A. Donath2 , L.O’C. Drury3 , G. Dubus32 , K. Dutson33 , J. Dyks34 , M. Dyrda25 , T. Edwards2 , K. Egberts35 , P. Eger2 , P. Espigat31 , C. Farnier28 , S. Fegan16 , F. Feinstein22 , M.V. Fernandes1 , D. Fernandez22 , A. Fiasson36 , G. Fontaine16 , A. Förster2 , M. Füßling35, S. Gabici31 , M. Gajdus6 , Y.A. Gallant22 , T. Garrigoux20 , G. Giavitto37 , B. Giebels16 , J.F. Glicenstein23 , D. Gottschall21 , M.-H. Grondin38 , M. Grudzińska24 , D. Hadasch15 , S. Häffner39 , J. Hahn2 , J. Harris8 , G. Heinzelmann1 , G. Henri32 , G. Hermann2 , O. Hervet19 , A. Hillert2 , J.A. Hinton33 , W. Hofmann2 , P. Hofverberg2 , M. Holler35 , D. Horns1 , A. Ivascenko18 , A. Jacholkowska20 , C. Jahn39 , M. Jamrozy10 , M. Janiak34 , F. Jankowsky27 , I. Jung-Richardt39 , M.A. Kastendieck1 , K. Katarzyński40 , U. Katz39 , S. Kaufmann27 , B. Khélifi31 , M. Kieffer20 , S. Klepser37 , D. Klochkov21 , W. Kluźniak34 , D. Kolitzus15 , Nu. Komin26 , K. Kosack23 , S. Krakau13 , F. Krayzel36 , P.P. Krüger18, H. Laffon38 , G. Lamanna36 , J. Lau30 , J. Lefaucheur31 , V. Lefranc23 , A. Lemière31 , M. Lemoine-Goumard38 , J.-P. Lenain20 , T. Lohse6 , A. Lopatin39 , C.-C. Lu2 , V. Marandon2 , A. Marcowith22 , R. Marx2 , G. Maurin36 , N. Maxted22 , M. Mayer35 , T.J.L. McComb8 , J. Méhault38,41 , P.J. Meintjes42 , U. Menzler13 , M. Meyer28 , A.M.W. Mitchell2 , R. Moderski34 , M. Mohamed27 , K. Morå28 , E. Moulin23 , T. Murach6 , M. de Naurois16 , –2– J. Niemiec25 , S.J. Nolan8 , L. Oakes6 , H. Odaka2 , S. Ohm37 , B. Opitz1 , M. Ostrowski10 , I. Oya37 , M. Panter2 , R.D. Parsons2 , M. Paz Arribas6 , N.W. Pekeur18 , G. Pelletier32 , P.-O. Petrucci32 , B. Peyaud23 , S. Pita31 , H. Poon2 , G. Pühlhofer21 , M. Punch31 , A. Quirrenbach27 , S. Raab39 , I. Reichardt31 , A. Reimer15 , O. Reimer15 , M. Renaud22 , R. de los Reyes2 , F. Rieger2 , C. Romoli3 , S. Rosier-Lees36 , G. Rowell30 , B. Rudak34 , C.B. Rulten19 , V. Sahakian5,4 , D. Salek43 , D.A. Sanchez36 , A. Santangelo21 , R. Schlickeiser13 , F. Schüssler23 , A. Schulz37 , U. Schwanke6 , S. Schwarzburg21 , S. Schwemmer27 , H. Sol19 , F. Spanier18 , G. Spengler28 , F. Spies1 , L. Stawarz10 , R. Steenkamp7 , C. Stegmann35,37 , F. Stinzing39 , K. Stycz37 , I. Sushch6,18 , J.-P. Tavernet20 , T. Tavernier31 , A.M. Taylor3 , R. Terrier31 , M. Tluczykont1 , C. Trichard36 , K. Valerius39 , C. van Eldik39 , B. van Soelen42 , G. Vasileiadis22 , J. Veh39 , C. Venter18 , A. Viana2 , P. Vincent20 , J. Vink9 , H.J. Völk2 , F. Volpe2 , M. Vorster18 , T. Vuillaume32 , S.J. Wagner27 , P. Wagner6 , R.M. Wagner28 , M. Ward8 , M. Weidinger13 , Q. Weitzel2 , R. White33 , A. Wierzcholska25 , P. Willmann39 , A. Wörnlein39 , D. Wouters23 , R. Yang2 , V. Zabalza2,33 , D. Zaborov16 , M. Zacharias27 , A.A. Zdziarski34 , A. Zech19 , H.-S. Zechlin1 . and Y. Fukui44 , H. Sano44 , T. Fukuda44 and S. Yoshiike44 . –3– 1 Universität Hamburg, Institut für Experimentalphysik, Luruper Chaussee 149, D 22761 Hamburg, Germany 2 Max-Planck-Institut für Kernphysik, P.O. Box 103980, D 69029 Heidelberg, Germany 3 Dublin Institute for Advanced Studies, 31 Fitzwilliam Place, Dublin 2, Ireland 4 National Academy of Sciences of the Republic of Armenia, Marshall Baghramian Avenue, 24, 0019 Yerevan, Republic of Armenia 5 Yerevan Physics Institute, 2 Alikhanian Brothers St., 375036 Yerevan, Armenia 6 Institut für Physik, Humboldt-Universität zu Berlin, Newtonstr. 15, D 12489 Berlin, Germany 7 University of Namibia, Department of Physics, Private Bag 13301, Windhoek, Namibia 8 University of Durham, Department of Physics, South Road, Durham DH1 3LE, U.K. 9 GRAPPA, Anton Pannekoek Institute for Astronomy, University of Amsterdam, Science Park 904, 1098 XH Amsterdam, The Netherlands 10 Obserwatorium Astronomiczne, Uniwersytet Jagielloński, ul. Orla 171, 30-244 Kraków, Poland 11 now at Harvard-Smithsonian Center for Astrophysics, 60 Garden St, MS-20, Cambridge, MA 02138, USA 12 Department of Physics and Electrical Engineering, Linnaeus University, 351 95 Växjö, Sweden 13 Institut für Theoretische Physik, Lehrstuhl IV: Weltraum und Astrophysik, Ruhr- Universität Bochum, D 44780 Bochum, Germany 14 GRAPPA, Anton Pannekoek Institute for Astronomy and Institute of High-Energy Physics, University of Amsterdam, Science Park 904, 1098 XH Amsterdam, The Netherlands 15 Institut für Astro- und Teilchenphysik, Leopold-Franzens-Universität Innsbruck, A-6020 –4– Innsbruck, Austria 16 Laboratoire Leprince-Ringuet, Ecole Polytechnique, CNRS/IN2P3, F-91128 Palaiseau, France 17 now at Santa Cruz Institute for Particle Physics, Department of Physics, University of California at Santa Cruz, Santa Cruz, CA 95064, USA 18 Centre for Space Research, North-West University, Potchefstroom 2520, South Africa 19 LUTH, Observatoire de Paris, CNRS, Université Paris Diderot, 5 Place Jules Janssen, 92190 Meudon, France 20 LPNHE, Université Pierre et Marie Curie Paris 6, Université Denis Diderot Paris 7, CNRS/IN2P3, 4 Place Jussieu, F-75252, Paris Cedex 5, France 21 Institut für Astronomie und Astrophysik, Universität Tübingen, Sand 1, D 72076 Tübingen, Germany 22 Laboratoire Univers et Particules de Montpellier, Université Montpellier 2, CNRS/IN2P3, CC 72, Place Eugène Bataillon, F-34095 Montpellier Cedex 5, France 23 DSM/Irfu, CEA Saclay, F-91191 Gif-Sur-Yvette Cedex, France 24 Astronomical Observatory, The University of Warsaw, Al. Ujazdowskie 4, 00-478 War- saw, Poland 25 Instytut Fizyki Ja̧drowej PAN, ul. Radzikowskiego 152, 31-342 Kraków, Poland 26 School of Physics, University of the Witwatersrand, 1 Jan Smuts Avenue, Braamfontein, Johannesburg, 2050 South Africa 27 Landessternwarte, Universität Heidelberg, Königstuhl, D 69117 Heidelberg, Germany 28 Oskar Klein Centre, Department of Physics, Stockholm University, Albanova University Center, SE-10691 Stockholm, Sweden 29 Wallenberg Academy Fellow, 30 School of Chemistry & Physics, University of Adelaide, Adelaide 5005, Australia 31 APC, AstroParticule et Cosmologie, Université Paris Diderot, CNRS/IN2P3, CEA/Irfu, –5– Igor Oya−[email protected], Sabrina Observatoire de Paris, Sorbonne Paris Cité, 10, rue Alice Domon et Léonie Duquet, 75205 Paris Cedex 13, France 32 Univ. Grenoble Alpes, IPAG, F-38000 Grenoble, France CNRS, IPAG, F-38000 Grenoble, France 33 Department of Physics and Astronomy, The University of Leicester, University Road, Leicester, LE1 7RH, United Kingdom 34 Nicolaus Copernicus Astronomical Center, ul. Bartycka 18, 00-716 Warsaw, Poland 35 Institut für Physik und Astronomie, Universität Potsdam, Karl-Liebknecht-Strasse 24/25, D 14476 Potsdam, Germany 36 Laboratoire d’Annecy-le-Vieux de Physique des Particules, Université de Savoie, CNRS/IN2P3, F-74941 Annecy-le-Vieux, France 37 DESY, D-15738 Zeuthen, Germany 38 Université Bordeaux 1, CNRS/IN2P3, Centre d’Études Nucléaires de Bordeaux Gradig- nan, 33175 Gradignan, France 39 Universität Erlangen-Nürnberg, Physikalisches Institut, Erwin-Rommel-Str. 1, D 91058 Erlangen, Germany 40 Centre for Astronomy, Faculty of Physics, Astronomy and Informatics, Nicolaus Coper- nicus University, Grudziadzka 5, 87-100 Torun, Poland 41 Funded by contract ERC-StG-259391 from the European Community, 42 Department of Physics, University of the Free State, PO Box 339, Bloemfontein 9300, South Africa 43 GRAPPA, Institute of High-Energy Physics, University of Amsterdam, Science Park 904, 1098 XH Amsterdam, The Netherlands 44 Department of Physics, Nagoya University, Furo-cho, Chiku sa-ku, Nagoya, 464-8601, Japan –6– Casanova−[email protected], [email protected] Received ; accepted Draft version –7– ABSTRACT This letter reports the discovery of a remarkably hard spectrum source, HESS J1641−463, by the High Energy Stereoscopic System (H.E.S.S.) in the very-high energy (VHE) domain. HESS J1641−463 remained unnoticed by the usual analysis techniques due to confusion with the bright nearby source HESS J1640−465. It emerged at a significance level of 8.5 standard deviations after restricting the analysis to events with energies above 4 TeV. It shows a moderate flux level of φ (E > 1 TeV) = (3.64 ± 0.44stat ± 0.73sys ) × 10−13 cm−2 s−1 , corresponding to 1.8% of the Crab Nebula flux above the same energy, and a hard spectrum with a photon index of Γ = 2.07 ± 0.11stat ± 0.20sys . It is a point-like source, although an extension up to Gaussian width of σ = 0.05◦ cannot be discounted due to uncertainties in the H.E.S.S. PSF. The VHE γ-ray flux of HESS J1641−463 is found to be constant over the observed period when checking time binnings from year-by-year to the 28 min exposures timescales. HESS J1641−463 is positionally coincident with the radio supernova remnant SNR G338.5+0.1. No X-ray candidate stands out as a clear association, however Chandra and XMM-Newton data reveal some potential weak counterparts. Various VHE γ-ray production scenarios are discussed. If the emission from HESS J1641−463 is produced by cosmic ray protons colliding with the ambient gas, then their spectrum must extend up to at least a few hundred TeV. The energy released in accelerating these particles could account for the entire energy budget of the galactic cosmic ray population above a few TeV. Subject headings: gamma rays: general — cosmic rays — ISM: individual objects (SNR G338.5+0.1, SNR G338.3−0.0) –8– 1. Introduction The large field of view (FoV) of the High Energy Stereoscopic System (H.E.S.S.), together with its stereoscopic observation strategy, allowed the discovery of tens of very high energy (VHE, ≥ 0.1 TeV) γ-ray sources1 by scanning a large fraction of the Galactic plane (Aharonian et al. 2005a; Carrigan et al. 2013). With deeper exposures, more VHE γ-ray sources are detected but source confusion begins to be problematic. Complementing the spatial search for new sources, an investigation in energy bands can provide an additional powerful tool for new discoveries. In this work, it will be shown how this method allowed for the detection of a new object, HESS J1641−463 (hereafter, J1641−463), that was previously hidden in the tails of the much brighter object HESS J1640−465. Interestingly, the newly discovered source exhibits one of the hardest spectra observed in VHE γ-rays, allowing its detection at higher energies, where the two sources are clearly separated. Hereafter, the observations and the analysis technique that led to the discovery of J1641−463 are described. Finally, a discussion of plausible counterparts of this source at other wavelengths is presented. 2. H.E.S.S. observations and results H.E.S.S. is an array of five imaging atmospheric Cherenkov telescopes located in the Khomas Highland of Namibia, 1800 m above sea level. In the initial phase of the H.E.S.S. project, during which the data described here were taken, the array was composed of four 13 m diameter telescopes. Extensive air showers are measured with an average energy resolution of 15% and an event angular resolution better than 0.1◦ (Aharonian et al. 2006) for a typical energy of 1 TeV. The trigger energy threshold is about 100 GeV and increases 1 See http://tevcat.uchicago.edu/ for an updated list of VHE γ-ray sources. –9– with higher zenith angle (Funk et al. 2004). J1641−463 remained unnoticed by the standard source detection techniques due to its low brightness and its proximity to the bright source HESS J1640−465 (Abramowski et al. 2014). During a study of a possible energy-dependent morphology of HESS J1640−465, a collection of images for events with energies above a set of energy thresholds (E > 1, 2, 3, 4 and 5 TeV) was created. J1641−463 was not visible in the original images of the HESS J1640−465 FoV as those images included no energy cut in the events, and thus were dominated by the much more numerous low energy events coming from the brighter HESS J1640−465. Thanks to the improved H.E.S.S. point spread function (PSF) at higher energies, and to its hard spectrum, J1641−463 was clearly visible in the highest energy sky maps, where the contamination from HESS J1640−465 was low. This discovery triggered further H.E.S.S. observation allowing the firm establishment of a new VHE γ-ray source. The VHE γ-ray excess image obtained for E > 4 TeV is shown in Fig. 1, where the background level is estimated following the ring background model (Berge et al. 2007). The observations of the FoV around J1641−463 were carried out from 2004 to 2011, corresponding to an acceptance-corrected livetime of 72 hours, after quality selection criteria were applied as in Aharonian et al. (2006). The data were analyzed with the methods described in Aharonian et al. (2006)2 . The events were reconstructed using the Hillas parameter technique (Hillas 1995). The results were cross-checked using two independent analysis methods (Ohm et al. 2009; de Naurois & Rolland 2009). The position of J1641−463 (together with the nearby HESS J1640−465) was determined by fitting a two dimensional double-Gaussian model convolved with the H.E.S.S. PSF to the 2 The H.E.S.S. hap-12-03 analysis software package with version 32 instrument response tables was used. – 10 – two-dimensional ON-source excess event distribution for E > 4 TeV, energies at which source confusion with HESS J1640−465 is mitigated. The centroid of the Gaussian corresponding to the location of J1641−463 was found to be αJ2000 = 16h 41m 2.1s ± 3.0sstat ± 1.9ssys , δJ2000 = −46◦ 18′ 13′′ ± 35′′stat ± 20′′sys . The source is found to be point-like, but a slightly extended morphology up to a width of σ = 0.05◦ cannot be ruled out due to uncertainties in the H.E.S.S. PSF. Figure 2 shows the projection of the excess events in the rectangular region shown in Fig. 1 for different energy bands. An F-test (Martin 1971) was performed comparing the single Gaussian model fits with the double-Gaussian fits. For all the energy bands, the null hypothesis can be rejected at significance levels of 3.6 − 4.3σ, thus clearly favoring the double Gaussian model. In order to minimize the contamination from HESS J1640−465, hard cuts were used, which imply a cut on θ2 (the square of the angular difference between the reconstructed shower direction and the source position) of 0.01 deg2 , and on individual image charge in photo-electrons of 200. The source is detected with a statistical significance of 8.5σ above 4 TeV, determined by using equation (17) in Li & Ma (1983) after background suppression with the reflected background model (Berge et al. 2007). The differential VHE γ-ray spectrum of J1641−463, derived using the forwardfolding technique (Piron et al. 2001), is compatible with a power-law function dN/dE = φ0 × (E/1 TeV)−Γ with φ0 = (3.91 ± 0.69stat ± 0.78sys ) × 10−13 cm−2 s−1 TeV−1 and Γ = 2.07 ± 0.11stat ± 0.20sys for the energy range from 0.64 to 100 TeV. The flux level is φ (E > 1 TeV) = (3.64 ± 0.44stat ± 0.73sys ) × 10−13 cm−2 s−1 , corresponding to 1.8% of the Crab Nebula flux above the same energy. At those energies, the estimated total contamination from HESS J1640−465 is 15 ± 6 %, reduced at higher energies (4 ± 3 % at E > 4 TeV). A fit by a power-law with exponential cutoff is not statistically justified given – 11 – the low flux level of J1641−463. A fit to a constant value of the period-by-period3 light curve for energies above 0.64 TeV yields a χ2 /d.o.f. = 11.7/14, with a p-value of 67%. No variability can be seen in other time binnings (from year-by-year to 28 min exposures). 3. Search for counterparts at other wavelengths 3.1. Radio observations J1641−463 is found within the bounds of SNR G338.5+0.1 (Green 2009). This SNR is located at αJ2000 = 16h 40m 59s , δJ2000 = −46◦ 17.8′ , has a roughly circular morphology, and shows a flux density at 1 GHz of ≈ 12 Jy (Green 2009). A diameter between 5′ (most obvious non-thermal emission region reported in Whiteoak & Green 1996) to 9′ (Green 2009) for G338.5+0.1 is assumed in this work, the latter one displayed in Fig. 1. Kothes & Dougherty (2007) conclude that the source is located at a distance of 11 kpc4 , which implies a physical size between ≈ 16 to ≈ 30 pc. Assuming that G338.5+0.1 is in the Sedov-Taylor, the Sedov solution (see e.g. van der Swaluw 2001) is used to estimate its age: with an explosion energy of 1051 erg and the density of the external medium between 0.1 to 1 cm−3 , the age of the SNR would correspond to 1.1−3.5 kyr and 5−17 kyr for 16 pc and 30 pc diameter, respectively. The distribution of molecular gas around J1641−463 is shown in the top left inset of Fig. 2. This distribution is obtained by integrating the CO 1→0 rotational line emission, measured with NANTEN, over a range in velocity between −40 km/s to −30 km/s 3 A H.E.S.S. observing period is the period between two full moons. 4 Although Shaver & Goss (1970) report a closer distance of 5.3 kpc, in this work it is assumed a distance of 11 kpc as reported by Kothes & Dougherty (2007), which is also quoted by Green (2009). – 12 – NANTEN 21 ×10 14 H.E.S.S. 10 -46°00’ 140 6 Declination 160 3 46 41- 2 SNR G338.5+0.1 -46° 30’ SNR G338.3-0.0 J16 S S HE SNR G338.5+0.1 5 -46 0 4 J16 S S HE SNR G338.3-0.0 120 100 80 60 40 20 0 16h44m 16h42m 16h40m 16h38m Right Ascension Fig. 1.— Map of excess events with energies E > 4 TeV for the region around J1641−463 smoothed with a Gaussian of width 0.085◦ , corresponding to the 68% containment radius of instrument PSF. The white contours indicate the significance of the emission at the 5, 6, 7 and 8σ level. The black cross indicates the value and uncertainty of the best fit position of the source, the green dashed circles show the positions and approximate extensions of the two nearby SNRs, the black diamond the position of PSR J1640−4631, the dash-dotted black ellipse the 95% confidence error position of 1FHL J1640.54634, and the red box indicates the area for the extraction of the profiles shown in Fig. 2. The color scale is in units of counts per smoothing Gaussian width. The H.E.S.S. PSF is shown inside the white box. The upper left inset shows a map of the distribution of the column density of molecular hydrogen in units of cm−2 , estimated from the NANTEN CO(1−0) data, together with the H.E.S.S. significance contours. E > 1 TeV -0.3 E > 2 TeV E > 3 TeV # Excess # Excess -0.4 # Excess -0.5 80 60 40 20 0 SNR G338.3-0.0 # Excess All events # Excess 350 300 250 200 150 100 50 0 250 200 150 100 50 0 140 120 100 80 60 40 20 0 100 80 60 40 20 0 SNR G338.5+0.1 – 13 – -0.2 -0.1 0 0.1 0.2 0.3 0.4 0.5 E > 4 TeV # Excess 50 40 E > 5 TeV 30 20 10 0 -10 -0.5 -0.4 -0.3 -0.2 -0.1 0 0.1 0.2 0.3 0.4 0.5 Projected distance [°] Fig. 2.— Distribution of VHE γ-ray excess profiles and Gaussian fits (convolved with the instrument PSF) for the red rectangular slice shown in Fig. 1. Vertical lines show the position of the SNR 338.3−0.0 and G338.5+0.1. Fits using a single and a double Gaussian function are shown in dashed and solid lines respectively. Note that the energy dependence of the PSF is taken into account in the fits. – 14 – (Matsunaga et al. 2001; Mizuno & Fukui 2004). The choice of this range is motivated by the presence of dense molecular cloud clumps in the region, mapped with various NH3 emission lines with the MOPRA survey at those velocities (de Wilt et al. 2012). Using the model for the Galactic rotation curves by Kothes & Dougherty (2007) the gas is located at a distance of about 11 kpc. Assuming a ratio XCO−>NH2 = 1.5 × 1020 between the CO velocity integrated intensity and the column density of molecular gas, NH2 the total column density from the extraction region of J1641−463 is 1.7 × 1022 cm−2 . At 11 kpc the density and the total mass are about 100 cm−3 and 2.4 × 105 solar masses, respectively. 3.2. X-ray observations No candidate for an X-ray counterpart of J1641−463 was found in existing catalogs, even when extending the search radius to 0.1◦ away from the source. Two data sets from Chandra and one from XMM-Newton were thus inspected in order to search for an X-ray counterpart of J1641−463. The Chandra ObsID 11008 partially covers J1641−463, with 40 ks of exposure, while ObsID 12508 fully encloses it with 19 ks. The data-sets were processed with the CIAO package. The tool wavdetect was used to identify sources, providing 32 faint point-like or marginally extended candidates at distances smaller than 0.1◦ to the J1641−463 position. This sample was filtered by two criteria, reducing the sample to 12 candidates (see Fig. 3): first, the sources with S/N ratios below 3 were rejected. Second, a cut on the hardness ratio as defined in Elvis et al. (2009) was applied, HR = (H−S)/(H+S), where H are the counts with 2−10 keV and S the counts with 0.3−2 keV. The sources with HR ≤ 0 were excluded. Spectral fits using an absorbed power law model were performed assuming a value of NH – 15 – of 2.0 × 1022 cm−2 , corresponding to the values reported by Kalberla et al. (2005) and Dickey & Lockman (1990), in good agreement with those derived with the NANTEN data. The estimated flux densities in the 0.3−10 keV energy band result in values from 7 × 10−15 erg cm−2 s−1 (src. B) to 1.5 × 10−13 erg cm−2 s−1 (src. L). No evidence of variability was found for any of the sources after performing a one-sample Kolmogorov-Smirnov test: the probability PKS for the hypothesis of a uniform flux was PKS > 0.1. None these sources is an obvious counterpart of J1641−463 due to their low fluxes and the lack of any morphological feature that could point to such an association. The XMM-Newton ObsID 0302560201, covering the region of HESS J1640−465 (Funk et al. 2007) constitutes a partial 23.7 ksec exposure of the source area. The data set was analysed using the XMM SAS analysis task edetect chain simultaneously in all three cameras and the 5 standard energy bands. In this manner, 27 sources were found, with only one consistent with the position and upper limit to the extension of J1641−463 (See Fig. 3). This source was detected only in the pn camera and only in the energy band 0.5−1 keV with a significance of ≈ 4.6σ, and it is not detected in the Chandra data. The vignetting for this source is 0.35 in the pn camera, so the observation is very insensitive to the region of interest. Due to low statistics, calculating an HR or spectrum for this source was not possible, and it is unclear whether this may represent a counterpart. 3.3. HE Observations The only High Energy (HE, 0.1−100 GeV) source found within 0.5◦ of J1641−463 is 2FGL J1640.5−4633 (Nolan et al. 2012), also present in the 10 > GeV Fermi/LAT Catalog as 1FHL J1640.54634 (Ackermann et al. 2013), likely to be associated with HESS J1640−465 (Slane et al. 2010; Gotthelf et al. 2014) (see Fig. 1). If the spectrum of J1641−463 is extended to lower energies as a featureless power law, its HE counterpart could be confused – 16 – L K 25 G 6 J F I D A Mercer 81 C H E 3 B Fig. 3.— Chandra [1.0−10 keV] mosaic image of the field surrounding J1641−463 from the ObsIDs 11008 and 12508. The image was exposure corrected, background subtracted and smoothed with a Gaussian of width 10”. The best fit position of J1641−463 with 1σ error bars is indicated with the black cross, while the upper limit to the source extension fit is indicated by the surrounding black circle. The detected hard X-ray sources are shown as red circles. The blue dashed circles indicate the positions of the sources detected by using XMMNewton data. The dashed contours indicate the significance of the VHE γ-ray emission as shown in Fig. 1. The thick green circle shows the position of the stellar cluster Mercer 81, target of the ObsID 11008. – 17 – with 1FHL J1640.54634. However, the extrapolation of the VHE emission of J1641−463 to the Fermi/LAT energy ranges predicts a flux of (5.0 ± 2.8) × 10−11 cm−2 s−1 in the 10−500 GeV band, a factor 10 lower than the flux of 1FHL J1640.5−4634 at those energies and thus the emission of the former would be dominated by the HE counterpart of HESS J1640−465. A study to resolve such a faint, confused source is challenging and outside the scope of this work, and would not affect the conclusions of this paper. 4. Discussion Possible scenarios to explain the emission from J1641−463 include the emission from accelerated particles within a SNR, a molecular cloud illuminated by cosmic rays (CRs), a pulsar wind nebula (PWN) and a γ-ray binary. These scenarios are discussed below. If G338.5+0.1 is a young SNR it can accelerate particles up to hundreds of TeV. The left panel of Fig. 4 shows the comparison between the H.E.S.S. spectrum and the spectrum produced by accelerated protons from G338.5+0.1, interacting with the ambient gas. The predicted spectra are calculated using the parametrization of Kelner et al. (2006), assuming a proton spectrum with a power-law slope of −2.1 and multiple cutoff energies. The profile of the log-likelihood ratio test statistic (Rolke et al. 2005) was used to estimate a confidence interval of the cutoff energies, while considering the spectral index and normalization as nuisance parameters and ignoring systematic errors. The 99% confidence level (CL) lower limit on the cutoff energy corresponds to 100 TeV. This proton spectrum is one of the hardest ever inferred to explain the emission from a γ-ray source and agrees well with the prediction by diffusive shock acceleration in young SNRs (Malkov & O’C Drury 2001). Remarkably, the γ-ray spectrum of J1641−463 is harder than that observed from the young SNR RXJ1713−4936 at energies above few TeV, where a cutoff is seen (Aharonian et al. 2007). If the TeV luminosity measured by H.E.S.S. is produced by collisions of protons – 18 – with the ambient gas, then the fraction of the energy of the supernova explosion converted into hadron acceleration is estimated to be Wp = Lγ tpp ≈ 1050 n−1 , where Lγ = 4 × 1034 erg/s is the total luminosity measured by H.E.S.S. above 0.64 TeV (at 11 kpc) and tpp = (ξppσpp cn)−1 is the energy loss time of protons, with ξpp = 0.45 (Aharonian 2004). For a gas density of n = 100 cm−3 , the proton energetics is Wp ≈ 1048 erg, implying that this SNR alone could explain the Galactic CR luminosity above few TeV. If G338.5+0.1 is older (5−17 kyr, see Sec. 3.1) then VHE protons accelerated by the young SNR G338.3−0.0, positionally coincident with HESS J1640−465 (Abramowski et al. 2014) could have already reached the dense MC coincident with J1641−463. This would explain the relatively high brightness of J1641−463 in comparison with HESS J1640−465 at high energies as shown in Fig. 2 (Aharonian & Atoyan 1996; Gabici et al. 2009). In such a scenario, because of CR escape, HESS J1640−465 would no longer look like a pevatron, as the highest energy CRs would have already left (Aharonian & Atoyan 1996). The much younger adjacent G338.3−0.0 would be in this scenario a major source of CRs. Electrons of hundreds of TeV IC (inverse Compton) scattering off the cosmic microwave background photons (CMB) could explain the emission from J1641−463. These eelctrons would be accelerated either in G338.5+0.1 or in the PWN associated to the young energetic pulsar, PSR J1640−4631, discovered within the observational boundaries of HESS J1640−465 (Gotthelf et al. 2014). Even assuming a pure power law for the primary electron spectrum, the cross section for IC scattering decreases at high energies resulting in a break in the γ-ray spectrum at multi TeV energies. Such a break is not observed in the spectrum of J1641−463. The predicted IC radiation, shown in the right panel of Fig. 4, was obtained by assuming that the electron cooled spectrum is a power law of spectral index −3.14 with different cutoff energies. The 99% CL lower limit on the cutoff energy, derived as in the case of the proton model using the exact Klein-Nishina expression for the – 19 – 10-11 10 10 -12 -2 -1 -1 10 -15 10 -16 10 -17 10 -18 10 -19 TeV cm s -14 IC off CMB photons 10-12 10-13 10 -11 p-p collisions 10-13 10-14 10-15 10-16 HESS J1641-463 Cutoff 100 TeV Cutoff 200 TeV Cutoff 1000 TeV NoCutoff HESS 1713.7-3946 10-17 10-18 HESS J1641-463 Cutoff=670 TeV Cutoff=1000 TeV NoCutoff HESS 1713.7-3946 10-19 1 10 TeV 100 1 10 TeV 100 Fig. 4.— Differential γ-ray spectrum of J1641−463 together with the expected emission from p-p collisions (left) and IC off CMB photons (right). The pink area represents the 1σ confidence region for the fit to a power law model, the black data points the H.E.S.S. measured photon flux (1σ uncertainties), the arrows the 95% CL upper limits on the flux level, and the black curves the expected emission from the models, assuming different particle energy cutoff values. For comparison, the gray data points and curve represent the archival spectrum and the corresponding best fit model, respectively, of SNR RX J1713.7−3946 (Aharonian et al. 2007). – 20 – IC emission, corresponds to 670 TeV. It is extremely difficult to accelerate electrons to such energies as hundred TeV electrons suffer severe synchrotron losses in the amplified magnetic fields of acceleration sites. Both the absence of a break in the γ-ray spectrum of J1641−463 and the derived lower limit on the cutoff energy of the electron spectrum strongly disfavor the leptonic scenario. A γ-ray binary scenario could also be considered, given the point-like morphology of J1641−463 and that a similarly hard spectral index of −2.23 has been found in one of these systems (LS 5039; Aharonian et al. 2005b). An X-ray flux as low as ∼ 10−14 erg cm−2 s−1 is expected from a faint X-ray binary system similar to HESS J0632+057 (Hinton et al. 2009) assuming a distance of 11 kpc, where the lack of an obvious optical counterpart could be due to high optical extinction caused by the large distance and the position close to the Galactic plane. 5. Conclusions Deeper exposures with H.E.S.S. together with a study of the emission in various energy bands made it possible to discover a new unique VHE source, showing one of the hardest γ-ray spectra ever found at these energies, extending up to at least 20 TeV without a break. In order to explain the observed VHE γ-ray spectrum, scenarios where protons are accelerated up to hundreds of TeV at either G338.5+0.1 or G338.3−0.0, and then interact with local gas or nearby massive MCs are the most compelling ones. Other possible scenarios, such as a PWN or a γ-ray binary, are disfavored but cannot be discarded. Deeper X-ray and VHE γ-ray observations, together with a better PSF for the latter, would allow for a better identification of the source. The support of the Namibian authorities and of the University of Namibia in facilitating – 21 – the construction and operation of HESS 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. 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