Volume 612, April 2018
H.E.S.S. phase-I observations of the plane of the Milky Way
Article Number A14
Number of page(s) 8
Section Stellar structure and evolution
Published online 09 April 2018

© ESO 2018

1 The SS 433/W50 system

SS 433 (RA 19h11m49.57s, Dec 4°5857.9) is the first binary system containing a stellar-mass compact object in which relativistic jets were discovered (Abell & Margon 1979; Fabian & Rees 1979). Located at a distance of 5.5 ± 0.2 kpc (Blundell & Bowler 2004; Lockman et al. 2007), SS 433 is an eclipsing X-ray binary system containing a black hole that is most likely ~10–20 M (Margon 1984) orbiting a ~30 M A3–7 supergiant star in a circular orbit with radius 79–86 R (Fabrika 2004). SS 433 is extremely bright with a bolometric luminosity of Lbol ~ 1040 erg s−1 (Cherepashchuk 2002) peaking at ultraviolet wavelengths. The source displays the most powerful jets known in our Galaxy, with Ljet ≳ 1039 erg s−1 (Dubner et al. 1998; Margon 1984; Marshall et al. 2002), ejected at a relativistic velocity of 0.26 c (Margon & Anderson 1989). The jets show a precessional period of ~162.4 days with a half opening angle of θpre ≈ 21° with respect to the normal to the orbital plane, with precessional phase Ψpre = 0 defined as the phase with the maximum exposure of the accretion disk to the observer (Fabrika 1993). The inclination of the jets with respect to the line of sight subtends an angle of i ≈ 78° (Eikenberry et al. 2001). This value is, however, time dependent owing to precession.

Both the high luminosity of SS 433 and its enormous jet power are thought to be a consequence of the persistent regime of supercritical accretion onto the compact object via Roche lobe overflow at a rate of ~10−4 M yr−1. In addition, SS 433 is one of the only two X-ray binary systems in which the presence of baryons in their jets has been found (Kotani et al. 1994); the other system is 4U 1630-47 (Díaz Trigo et al. 2013). Clouds of plasma with baryonic content propagate along ballistic trajectories to large distances without appreciable deceleration. At growing distances from the source, the collimated jets (with opening angle of ≈1.2°) can be distinguished in the X-ray, optical, and radio bands. The X-ray jets give rise to lines of highly ionized heavy elements (Kotani et al. 1994; Marshall et al. 2002). The emission is produced by hot gas (T ~108 K), which cools due to expansion and radiative losses while propagating outwards. Plasma at T ~ 107 K is, however, still observed at large distances along the jet, indicating that a continuous source of heating is required to maintain the observed emission (Migliari et al. 2002). At radio wavelengths, the observed synchrotron flux density is about 1 Jy at 1 GHz with luminosities reaching ~4 × 1032 erg s−1.

SS 433 is surrounded by the radio shell of W50, which is a large 2° × 1° nebula catalogued as SNR G39.7−2.0 (Green 2006). Its present morphology is thought to be the result of the interaction between the jets of SS 433 and the surrounding medium (Goodall et al. 2011). This scenario is supported by the position of SS 433 at the centre of W50, the elongation of the nebula in the east-west direction along the axis of precession of the jets (forming the so-called “ears” of W50; see Safi-Harb & Ögelman 1997), the presence of radio, IR, optical and X-ray emitting regions also aligned with the jet precession axis, and the structure of the magnetic field through the observation of linearly polarized radio emission in the SS 433/W50 region (Farnes et al. 2017). At a distance of ~1020 cm, or about ~30 pc, the outflowing jets are decelerated and, together with an enhanced intensity of the emission at radio wavelengths, large-scale X-ray lobes are observed. The extended X-ray emission is mostly of non-thermal origin and generally much softer than the emission from the central source (hard X-ray emission is found however in the eastern interaction regions; see Safi-Harb & Ögelman 1997). At the position of the maximum of the X-ray extended emission, optical filaments perpendicular to the jet precession axis are found. Spectral analysis shows that these filaments are formed by the sweeping up of the interstellar gas and display a proper motion of 50–90 km s−1 (Zealey et al. 1980; Fabrika 2004). At larger distances, beyond the W50 length scales, the presence of molecular clouds aligned in the direction of SS 433 jets has been reported (Yamamoto et al. 2008). These clouds, which extend for ~250 pc at a distance of 5 kpc, may have formed through the interaction of SS 433 jets with the interstellar H I gas, which would imply that SS 433 jets are more extended by a factor of ~3 than the observed X-ray jets (Yamamoto et al. 2008).

2 Gamma-ray emission and absorption processes in SS 433/W50

SS 433 is an exceptional laboratory to test theoretical predictions of high (100 MeV < E < 100 GeV) and very high-energy (VHE; E > 100 GeV) emission produced in microquasar jets (see e.g. Levinson & Blandford 1996; Atoyan & Aharonian 1999; Kaufman Bernadó et al. 2002; Bosch-Ramon et al. 2006; Orellana et al. 2007; Reynoso et al. 2008b; Bosch-Ramon & Khangulyan 2009). In a leptonic framework, gamma rays could be produced through inverse Compton (IC) scattering of ambient photon fields, which are dominated in this case by the supergiant companion star and the UV and mid-IR emission from the extended accretion disk (Gies et al. 2002; Fuchs et al. 2006). In addition, synchrotron-self Compton emission and the interaction of accelerated electrons with jet ions through relativistic Bremsstrahlung processes could also generate VHE fluxes (Aharonian & Atoyan 1998). In a hadronic scenario, interactions of relativistic protons in the jet produce gamma rays through π0 decay (see e.g. Reynoso et al. 2008b, for a detailed study applied to SS 433). The target ions could be provided both by the companion and disk winds or by the pool of thermal protons outflowing within the jets.

Gamma rays, when produced in the inner regions of SS 433, can be strongly attenuated (see e.g. Reynoso et al. 2008a). Both the donor star and compact object are thought to be embedded in a thick extended envelope (Zwitter et al. 1991), which forms as a result of the supercritical accretion rate onto the compact object, and provides a dense low-energy UV and mid-IR radiation field in which VHE photons are absorbed. In addition, about 10−4 M yr−1 are expelled within the ~30° half opening angle subtended by the envelope (Fabrika & Sholukhova 2008). Therefore, absorption of VHE gamma rays can also occur due to gamma-nucleon interactions through photo-pion and photo-pair production processes, whereas the photon field of the companion can also effectively reduce the gamma-ray flux through pair creation (Reynoso et al. 2008a). Such strong absorption is expected to occur along ~80% of the SS 433 precession cycle with a maximum at precession phase Ψpre ≈ 0.5, and during the regular eclipses of the compact object by the donor star at orbital phases ϕorb ≈ 0.

The interaction regions between the SS 433 jets and the surrounding W50 nebula could also produce VHE gamma-ray emission, for example through IC scattering of cosmic microwave background (CMB) photons by electrons accelerated at the eastern (e1, e2, e3) and western (w1, w2) termination regions (Safi-Harb & Ögelman 1997; Aharonian & Atoyan 1998; Bordas et al. 2009), or through pp interactions if protons are efficiently accelerated at the interaction shocks (Heinz & Sunyaev 2002; Bosch-Ramon et al. 2005). Non-thermal emission from these regions has indeed been observed from radio to X-rayenergies (Brinkmann et al. 1996, 2007; Dubner et al. 1998; Safi-Harb & Ögelman 1997; Safi-Harb & Petre 1999; Fuchs 2002).

Recently, high-energy gamma-ray emission from a source associated with SS 433/W50 has been reported from the analysis of ~5 yr of Fermi-LAT archival data (Bordas et al. 2015). The relatively large point spread function (PSF) of the Fermi-LAT at the sub-GeV energies in which the source is detected, larger than ~ 1.5°, prevents from an accurate localization of the emitter. At VHEs, SS 433/W50 remains so far undetected. Upper limits have been reported by the CANGAROO-II and HEGRA Collaborations, at the level of 2–3% of the Crab nebula flux above 800 GeV (Hayashi et al. 2009; Aharonian et al. 2005), following extended observation campaigns including both the central system and the jet/medium interaction regions. The selected dates of these observations, however, did not account for the gamma-ray absorption affecting the inner system. Such constraints were instead accounted for in observations of SS 433 performed by VERITAS in 2007 and MAGIC in 2008, for which upper limits are reported in Saito et al. (2009) and Guenette & for the VERITAS Collaboration (2009), respectively.

In this work, a search for VHE emission from the microquasar SS 433 with the Major Atmospheric Gamma Imaging Cherenkov telescopes (MAGIC) and High Energy Spectroscopic System (H.E.S.S.) Imaging Atmospheric Cherenkov Telescopes (IACTs) is reported, following dedicated observations of the source spanning several years and taken at orbital/precession phases where gamma-ray absorption should be minimal. The SS 433/W50 interaction regions are also investigated, for which a wider data set is used that is not restricted to the low-absorption phases criterion applied to the study of the inner system. The observations and analysis results are described in Sect. 3 and are later discussed in Sect. 4.

3 VHE observations, analysis, and results

3.1 H.E.S.S. and MAGIC observations

Observations of SS 433/W50 were conducted with the H.E.S.S. and MAGIC Cherenkov telescope arrays. The H.E.S.S. is an IACT array located in the Khomas highland of Namibia (23°S, 16°E, 1800 m above the sea level). In its first phase, the system consisted of four identical 13 m diameter imaging Cherenkov telescopes, covering a field of view (FoV) of about 5° diameter (Bernlöhr et al. 2003). A major upgrade took place in 2012 with the addition of CT-5, which is a 28 m diameter telescope at the centre of the array. The data presented in this paper makes use of observations taken during the H.E.S.S.-I phase only (CT1-4 configuration). In this configuration, H.E.S.S. is able to detect at a 5σ statistical significance level a source with ~ 0.7% of the Crab nebula flux in 50 h of observations (Aharonian et al. 2006a). The MAGIC is a stereoscopic system of two 17 m diameter Cherenkov telescopes located at the observatory El Roque de Los Muchachos (28°N, 18°W, 2200 m above the sea level) on La Palma, Canary Islands, Spain. Each telescope is composed of a pixelized camera with a FoV of 3.5°. The sensitivity of MAGIC at the time of the observations was ~0.76% ± 0.03% of the Crab nebula flux in 50 h above 290 GeV (Aleksić et al. 2012a).

The H.E.S.S. and MAGIC data set includes observations of SS 433/W50 taken in 2006, 2007, 2009, 2010, and 2011. Whereas the 2006 and 2007 H.E.S.S. observations were part of the H.E.S.S. Galactic Plane Survey (HGPS; H.E.S.S. Collaboration 2018b), both the MAGIC and H.E.S.S. campaigns in 2009, 2010, and 2011 were dedicated to SS 433. Making use of the ephemeris provided by Goranskij (2011), the latter were scheduled at times in which the source was found at orbital and precession phases Ψ where absorption of its putative VHE emission is expected to be at its minimum, Ψpre = 0.9–0.1 (Reynoso et al. 2008a). The total H.E.S.S. effective exposure time on SS 433/W50 amounts to 45 h of data after standard quality selection cuts. A total effective exposure time of 8.7 h is available for the central system after selecting low-absorption precession/orbital phases. In May and June 2010, MAGIC performed observations of SS 433 in stereo mode for 10 h. After quality cuts, 7.8 h of good data remained. A summary of the H.E.S.S. and MAGIC observations of SS 433 is collected in Table 1.

The H.E.S.S. observations of the central system were performed at zenith angles ranging between 28° to 54°, with an average of 34°, while MAGIC observed at zenith angles between 24° and 30°. The observations in both H.E.S.S and MAGIC were performed in wobble mode (Fomin et al. 1994), with an offset of 0.4° for MAGIC and 0.7° for H.E.S.S., respectively, away from the source position to simultaneously evaluate the background. This observation mode allowed imaging of not only the central binary system, but also the eastern (e1, e2, e3) and western (w1, w2) interaction regions with the W50 nebula. In the case of H.E.S.S., the interaction regions were also observed as part of the HGPS programme (Aharonian et al. 2006b; H.E.S.S. Collaboration 2018b), at zenith angles between 25° to 57°, with an average of 35°–38° depending on the region. The total exposure time varies from region to region (see Table 2).

Table 1

Observations of SS 433 performed by H.E.S.S. and MAGIC telescopes.

Table 2

Integral H.E.S.S. and MAGIC flux upper limits derived for SS 433 during low-absorption orbital/precessional phases and for the eastern/western interaction regions indicated in Fig. 1 using all available data.

3.2 Analysis

Data analysis was performed following the standard analysis procedure for each of the two instruments (see Aharonian et al. 2006a, for H.E.S.S.; and Aleksić et al. 2016, for MAGIC analysis details). The imaging technique is based on the parameterization of the images formed in the camera plane in order to extract the information contained in the shower with the Hillas parameters (Hillas 1985). The signal extraction was performed by the reconstruction and calibration of the size and arrival time of the Cherenkov pulses. The event reconstruction was obtained by image cleaning and shower parameterization, whereas the signal and background discrimination and energy estimation were obtained by comparison of the Hillas parameters with look-up tables for a given shower intensity and impact distance (see Aharonian et al. 2006a; Aleksić et al. 2012a), or by training an algorithm to perform gamma/hadron separation via the random forest (RF) method (Albert et al. 2008). The event direction was derived in stereoscopic observations from the intersection of the major axes of the shower images in multiple cameras. Finally, the signal was extracted geometrically from the angular distance θ2 ; i.e. the angular distance between the source position and the estimated source position for a given event. The signal is then determined by an upper cut in these angles, since gamma rays are reconstructed with small angles and the background follows a featureless, almost-flat distribution. For the H.E.S.S. analysis, an independent cross-check with the model analysis technique (de Naurois & Rolland 2009) was performed, making use of an independent calibration procedure of the raw data, with both the analysis chains providing compatible results. Standard cuts were used, where a cut of 60 photoelectrons on the intensity of the extensive air showers is applied, providing a mean energy threshold of ~287 GeV for the analysis reported here. The energy threshold reached by MAGIC is 150 GeV. A point-like source was assumed for the analysis of SS 433. The interaction regions display extended emission at lower energies. To account for such extension, the MAGIC and H.E.S.S. analyses were optimized assuming a source radius (θ-cut) of 0.05°, 0.17°, 0.25°, 0.07°, and 0.07° for e1, e2, e3, w1, and w2, respectively, derived from the extension of these regions observed at X-ray energies (see e.g. Safi-Harb & Ögelman 1997; Safi-Harb & Petre 1999; Aharonian et al. 2005, and references therein).

3.3 Results

The H.E.S.S. and MAGIC observations reported here do not show any significant VHE emission either for the central source SS 433 or for any of the interaction regions with the W50 nebula e1, e2, e3, w1, and w2 (see Fig. 1). Integral upper limits Rolke et al. (2005) have been calculated at E ≥ 300 GeV and at E ≥ 800 GeV; the latter allows for a comparison with previous results on the source reported by the HEGRA (Aharonian et al. 2005) Collaboration. The results are summarised in Table 2. A day-by-day analysis of the H.E.S.S. and MAGIC data sets was also performed without any signature of significant emission, which could suggest a flaring episode during the dates of observation.

The H.E.S.S. and MAGIC observations were used to compute the differential flux upper limits for the VHE emission from the central binary system at orbital/precession phases where absorption should be at its minimum. These limits were computed through a maximum-likelihood ratio test applied to the combined data sets obtained by both observatories. Events in the signal region (nON) and in the background control regions (nOFF) from each instrument are collected in addition to the ratio of the areas in the signal and the background regions (α), effective area (Aeff) and effective observing time teff corresponding to the observations of each instrument. A likelihood profile is then computed in each studied energy bin (ΔEi) for both the signal and background distributions. Systematic uncertainties are accounted for through the inclusion of additional likelihood profiles for the distributions of α, Aeff , and energy resolution, assuming systematics at the level of δα = 10%, δAeff = 15%, and δEi = 15% for the measurements of these quantities by each instrument (Aharonian et al. 2006a; Aleksić et al. 2016). The inclusion of these systematics results in an enhancement by ~15% to 30% on the final combined differential flux upper limit values, depending on the studied energy bin. To obtain the final combined differential flux upper limits, a likelihood ratio test is employed assuming a given range of values for the normalization factor of the gamma-ray differential spectrum, N0 . From the maximum of the likelihood profile, a 95% confidence interval for the differential upper limit in each energy bin ΔEi is derived through d N∕dE = N0 × EΓ, where a fixed spectral index Γ = 2.7 was assumed. The final differential upper limits are shown in Fig. 2, both for each instrument and the combined values, together with the Crab nebula flux, for reference, and the theoretical predictions on the gamma-ray flux from SS433 expected at low-absorption precession phases Ψ ∈ [0.9, 0.1] by Reynoso et al. (2008b).

thumbnail Fig. 1

Significance map, derived from the H.E.S.S. data, for the FoV centred at the position of SS 433/W50 at E ≥ 287 GeV. GB6 4.85 GHz radio contours (white, from Gregory et al. 1996) are superimposed. Cyan circles indicate the positions of the interaction regions e1, e2, e3 (eastern “ear”) and w1, w2 (western “ear”). The bright source located north-west of SS 433 is MGRO J1908+06 (Abdo et al. 2007).

thumbnail Fig. 2

Differential flux upper limits (95% C.L.) from SS 433 obtained with MAGIC (blue), H.E.S.S. (green) and a combination of both telescopes (red) assuming a power-law with a spectral index Γ = 2.7 for the differential gamma-ray flux. The predicted differential gamma-ray flux from Reynoso et al. (2008b) for precessional phases Ψpre ∈ [0.9, 0.1] in which absorption of VHE emission should be at its lower level is also displayed (dashed orange), together with the Crab Nebula flux, for reference (from Aleksić et al. 2012b).

4 Discussion

The H.E.S.S. and MAGIC observations reported here do not show any significant signal of VHE emission from SS 433/W50. The variable absorption of a putative VHE gamma-ray flux emitted from the inner regions of the binary system, which could be responsible for this non-detection, is accounted for in this study by selecting observationscorresponding to precession/orbital phases where this absorption should be at its minimum. The combinationof the MAGIC and H.E.S.S. observations in addition provides a relatively wide coverage of the relevant precession phases from 2006 to 2011. If a long-term super-orbital variability exists in SS 433 with timescales of ~few years, for example related to a varying jet injection power or the changing conditions of the absorber in the surroundings of the central compact object, such variability does not result in an enhancement of the TeV flux up to the detection level of current IACTs.

While SS 433 remains undetected at VHE, the system displays non-thermal emission at lower energies along the jets and/or at the SS 433/W50 interaction regions, which ensures the presence of an emitting population of relativistic particles in the system. In particular for the eastern nebula interaction sites, the observed synchrotron X-ray emission implies the presence of for example up to multi-TeV electrons in these regions (Safi-Harb & Petre 1999).

By considering in detail the photon and matter fields both from the companion star and accretion/circumstellar disks, gamma-ray fluxes from SS 433/W50 have been predicted at a level of ~ 10−12–10−13 ph cm−2 s−1 (see e.g. Band & Grindlay 1986; Aharonian & Atoyan 1998; Reynoso et al. 2008b). Reynoso et al. (2008b) consider in particular pp interactions between relativistic and cold protons in SS 433 jets during low-absorption precession/orbital phases, producing gamma-ray fluxes at Eγ ≥ 800 GeV during these precession phases at a level of ΦVHE ≈ 2.1 × 10−12 ph cm−2 s−1. The general framework used to derive the relativistic proton distribution in Reynoso et al. (2008b) has been revised by Torres & Reimer (2011), who report significant deviations of these proton fluxes for jets displaying large Lorentz factors and/or small viewing angles, for example blazar jets and gamma-ray bursts. In SS 433, with a moderate jet Lorentz factor of 1.036 (v = 0.26c; Abell & Margon 1979) and a relatively large jet viewing angle, ~78° (Eikenberry et al. 2001), the correction factor on the fluxes predicted by Reynoso et al. (2008b) could be affected at the level of ~20%. The gamma-ray flux predicted by Reynoso et al. (2008) depends on the efficiency in transferring jet kinetic energy to the relativistic proton population, qp, which is treated in their model as a free parameter. Using the HEGRA upper limits to the VHE gamma-ray flux from SS 433, qp is constrained to be ≤2.9 × 10−4. This upper limit is obtained however under the assumption that the HEGRA observations took place during a complete precessional cycle. With the H.E.S.S. and MAGIC upper limits reported here, a more stringent constraint on the fraction of power carried by relativistic protons in the SS 433 jets is obtained, qp ≲ 2.5 × 10−5.

In a hadronic scenario, gamma rays from π0-decay should also be accompanied by neutrinos from the decay of charged pions. The IceCube Collaboration analyzed the data around the position of SS 433/W50 with no significant detection of the source. An upper limit at 90% confidence level (C.L.) on the muon neutrino flux at 1 TeV is set at TeV cm−2 s−1 (Aartsen et al. 2014), using four years of data and assuming an E−2 flux distribution. The model of Reynoso et al. (2008b) also predicts the neutrino flux emitted by this system with a neutrino differential flux at 1 TeV of ϕν = 2 × 10−12 cm−2 s−1 averaged over all precessional phases. The IceCube upper limits can be used to put a limit on qp of ~ 3.3 × 10−5, which is marginally less restrictive than the value obtained with the gamma-ray observations. However, the gamma-ray and neutrino flux estimates of Reynoso et al. (2008b) are based on a proton acceleration efficiency of η = tacctgyr ~ 0.07; tacc is the acceleration time and tgyr = ceBEp, where B is the magnetic field in the accelerator region and Ep is the proton energy. Accounting for adiabatic and radiation losses, these authors derive a maximum energy for relativistic protons of Ep ≤ 103 TeV. Different values for the magnetic field, target proton densities, and/or adiabatic expansion velocities in the acceleration region would also imply variations in predicted gamma-ray and neutrino fluxes.

At the interaction regions of the jets of SS 433 with the surrounding W50 nebula, the X-ray spectra from the extended lobes are well represented by a power-law model (Moldowan & Safi-Harb 2005); a synchrotron origin for this emission has been suggested (Safi-Harb & Ögelman 1997; see also Safi-Harb & Petre 1999) that would imply the presence of electrons with energies up to ~50 TeV in those regions. The VHE gamma-ray emission from the SS 433/W50 interaction regions was first considered by Aharonian & Atoyan (1998), who estimated gamma-ray fluxes at a level of ~ 10−12 ph cm−2 s−1 for the eastern e3 region produced by electrons scattering off CMB photons. Bordas et al. (2009) also considered the non-thermal emission produced in microquasar jets/ISM interaction regions, providing gamma-ray flux estimates as a function of the kinetic power of the jets, age of the system, and particle density of the environment. The application of this model to SS 433/W50 yielded fluxes at the level of ~10−13 erg cm−2 s−1 for E > 250 GeV for an assumed distance to the system of 5.5 kpc (Bordas et al. 2010), which are roughly at the level of the upper limitsreported here. However, as noted in Aharonian & Atoyan (1998; see also discussion in Safi-Harb & Petre 1999; and Aharonian et al. 2005), electrons accelerated at the interaction region shock interface could lose most of their energy mainly through synchrotron emission for ambient magnetic fields at or above ~ 10 μG, preventing an effective channelling through IC scattering that is relevant for the production of gamma rays at high and very high energies. The integral flux upper limits for the interaction regions shown in Table 2, together with the assumption that the same high-energy electron population is responsible for the observed (synchrotron) X-ray emission and the putative gamma-ray fluxes, can be used to constrain the magnetic fields present in the shocked SS 433 jets/ISM interaction regions. Rowell & HEGRA Collaboration (2001; see also Aharonian et al. 2005) make use of HEGRA upper limits obtained for the e3 interaction region (ΦHEGRA ≤ 2.1 × 10−12 ph cm−2 s−1) and the predictions by Aharonian et al. (1997) to derive a lower limit on the post-shock magnetic field in this region of ~ 13 μG. Using the upper limits reported here, a more constraining lower limit on the magnetic field of 20–25 μG is obtained.

The huge kinetic luminosity and baryonic matter transported by the SS 433 jets and the presence of the surrounding target material provided by the disk wind and/or the W50 nebula render pp interactions at those larger scales a good TeV emission mechanism as well, as shown for even more modest energy budgets (see e.g. Bosch-Ramon et al. 2005). Bordas et al. (2015) found a gamma-ray signal from the direction of SS 433/W50. The steadiness of the flux and the derived spectral properties, with the gamma-ray emission extending only from 200 MeV to 800 MeV, prompted Bordas et al. (2015) to suggest a jet-medium interaction scenario for the observed emission. A cut-off power law was needed to fit the Fermi-LAT spectrum with cut-off energies of a few GeV. The upper limits reported here are therefore fully consistent with the Fermi-LAT extrapolation of the fitted spectra. If the MeV/GeV emission is produced by relativistic particles in SS 433 jets, as suggested in Bordas et al. (2015), the acceleration mechanism may be only relatively efficient, thereby preventing a significant detection of the system in the VHE regime.

The upper limits reported here for SS 433 and those obtained on the steady emission from other well-established Galactic microquasars (e.g. Cyg X-1, Cyg X-3, GRS 1915+105, MWC 656, and Cir X-1; see e.g. Nicholas & Rowell 2008; H.E.S.S. Collaboration 2009; H.E.S.S. Collaboration 2018a; Saito et al. 2009; Aleksić et al. 2010, 2015) imply that if their jets are also baryon loaded as in SS 433 (see also the case of 4U1630-47, Díaz Trigo et al. 2013), their contribution to the Galactic cosmic-ray flux must be limited to relatively low energies in the GeV domain, as more efficient proton acceleration is constrained by the lack of VHE gamma-ray emission.


The authors would like to thank M. M. Reynoso for sharing data and information from his paper Reynoso et al. (2008b). 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 B.M.B.F. and M.P.G., the Italian INFN and INAF, the Swiss National Fund SNF, the ERDF under the Spanish MINECO (FPA2015-69818-P, FPA2012-36668, FPA2015-68378-P, FPA2015-69210-C6-2-R, FPA2015-69210-C6-4-R, FPA2015-69210-C6-6-R, AYA2015-71042-P, AYA2016-76012-C3-1-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 the Croatian Science Foundation (HrZZ) Project 09/176 and the University of Rijeka Project, by the DFG Collaborative Research Centers SFB823/C4 and SFB876/C3, and by the Polish MNiSzW grant 745/N-HESS-MAGIC/2010/0. The support of the Namibian authorities and the University of Namibia in facilitating the construction and operation of H.E.S.S. is gratefully acknowledged, as is the support by the German Ministry for Education and Research (BMBF), the Max Planck Society, the German Research Foundation (DFG), the French Ministry for Research, the CNRS-IN2P3 and the Astroparticle Interdisciplinary Programme of the CNRS, the UK Science and Technology Facilities Council (STFC), the IPNP of the Charles University, the Czech Science Foundation, 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.


  1. Aartsen, M. G., Ackermann, M., Adams, J., et al. 2014, ApJ, 796, 109 [NASA ADS] [CrossRef] [Google Scholar]
  2. Abdo, A. A., Allen, B., Berley, D., et al. 2007, ApJ, 664, L91 [NASA ADS] [CrossRef] [Google Scholar]
  3. Abell, G. O., & Margon, B. 1979, Nature, 279, 701 [NASA ADS] [CrossRef] [Google Scholar]
  4. Aharonian, F. A., & Atoyan, A. M. 1998, New Astron. Rev., 42, 579 [NASA ADS] [CrossRef] [Google Scholar]
  5. Aharonian, F. A., Atoyan, A. M., & Kifune, T. 1997, MNRAS, 291, 162 [NASA ADS] [CrossRef] [Google Scholar]
  6. Aharonian, F., Akhperjanian, A., Beilicke, M., et al. 2005, A&A, 439, 635 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  7. Aharonian, F., Akhperjanian, A. G., Bazer-Bachi, A. R., et al. 2006a, A&A, 457, 899 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  8. Aharonian, F., Akhperjanian, A. G., Bazer-Bachi, A. R., et al. 2006b, ApJ, 636, 777 [NASA ADS] [CrossRef] [Google Scholar]
  9. Albert, J., Aliu, E., Anderhub, H., & et al. 2008, Nucl. Instr. Meth. Phys. Res. A, 588, 424 [Google Scholar]
  10. Aleksić, J., Antonelli, L. A., Antoranz, P., et al. 2010, ApJ, 721, 843 [Google Scholar]
  11. Aleksić, J., Alvarez, E. A., Antonelli, L. A., et al. 2012a, Astropart. Phys., 35, 435 [NASA ADS] [CrossRef] [Google Scholar]
  12. Aleksić, J., Alvarez, E. A., Antonelli, L. A., et al. 2012b, A&A, 540, A69 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  13. Aleksić, J., Ansoldi, S., Antonelli, L. A., et al. 2015, A&A, 576, A36 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  14. Aleksić, J., Ansoldi, S., Antonelli, L. A., et al. 2016, Astropart. Phys., 72, 76 [NASA ADS] [CrossRef] [Google Scholar]
  15. Atoyan, A. M., & Aharonian, F. A. 1999, MNRAS, 302, 253 [NASA ADS] [CrossRef] [Google Scholar]
  16. Band, D. L., & Grindlay, J. E. 1986, ApJ, 311, 595 [NASA ADS] [CrossRef] [Google Scholar]
  17. Bernlöhr, K., Carrol, O., Cornils, R., et al. 2003, Astropart. Phys., 20, 111 [NASA ADS] [CrossRef] [Google Scholar]
  18. Blundell, K. M., & Bowler, M. G. 2004, ApJ, 616, L159 [NASA ADS] [CrossRef] [Google Scholar]
  19. Bordas, P., Bosch-Ramon, V., Paredes, J. M., & Perucho, M. 2009, A&A, 497, 325 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  20. Bordas, P., Bosch-Ramon, V., & Paredes, J. M. 2010, Int. J. Mod. Phys. D, 19, 749 [NASA ADS] [CrossRef] [Google Scholar]
  21. Bordas, P., Yang, R., Kafexhiu, E., & Aharonian, F. 2015, ApJ, 807, L8 [NASA ADS] [CrossRef] [Google Scholar]
  22. Bosch-Ramon, V., & Khangulyan, D. 2009, Int. J. Mod. Phys. D, 18, 347 [NASA ADS] [CrossRef] [Google Scholar]
  23. Bosch-Ramon, V., Aharonian, F. A., & Paredes, J. M. 2005, in Astrophysical Sources of High Energy Particles and Radiation, eds. T. Bulik, B. Rudak, & G. Madejski, AIP Conf. Ser., 801, 196 [NASA ADS] [CrossRef] [Google Scholar]
  24. Bosch-Ramon, V., Romero, G. E., & Paredes, J. M. 2006, A&A, 447, 263 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  25. Brinkmann, W., Pratt, G. W., Rohr, S., Kawai, N., & Burwitz, V. 2007, A&A, 463, 611 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  26. Cherepashchuk, A. 2002, Space Sci. Rev., 102, 23 [NASA ADS] [CrossRef] [Google Scholar]
  27. de Naurois, M., & Rolland, L. 2009, Astropart. Phys., 32, 231 [NASA ADS] [CrossRef] [Google Scholar]
  28. Díaz Trigo, M., Miller-Jones, J. C. A., Migliari, S., Broderick, J. W., & Tzioumis, T. 2013, Nature, 504, 260 [NASA ADS] [CrossRef] [Google Scholar]
  29. Dubner, G. M., Holdaway, M., Goss, W. M., & Mirabel, I. F. 1998, AJ, 116, 1842 [NASA ADS] [CrossRef] [Google Scholar]
  30. Eikenberry, S. S., Cameron, P. B., Fierce, B. W., et al. 2001, ApJ, 561, 1027 [NASA ADS] [CrossRef] [Google Scholar]
  31. Fabian, A. C., & Rees, M. J. 1979, MNRAS, 187, 13P [NASA ADS] [CrossRef] [Google Scholar]
  32. Fabrika, S. 2004, Astrophys. Space Phys. Rev., 12, 1 [NASA ADS] [Google Scholar]
  33. Fabrika, S. N. 1993, MNRAS, 261, 241 [NASA ADS] [Google Scholar]
  34. Fabrika, S. N., & Sholukhova, O. 2008, in Microquasars and Beyond, 52 [Google Scholar]
  35. Farnes, J. S., Gaensler, B. M., Purcell, C., et al. 2017, MNRAS, 467, 4777 [NASA ADS] [CrossRef] [Google Scholar]
  36. Fomin, V. P., Stepanian, A. A., Lamb, R. C., et al. 1994, Astropart. Phys., 2, 137 [NASA ADS] [CrossRef] [Google Scholar]
  37. Fuchs, Y. 2002, ArXiv e-prints [arxiv:astro-ph/0207429] [Google Scholar]
  38. Fuchs, Y., Koch Miramond, L., & Ábrahám, P. 2006, A&A, 445, 1041 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  39. Gies, D. R., McSwain, M. V., Riddle, R. L., et al. 2002, ApJ, 566, 1069 [NASA ADS] [CrossRef] [Google Scholar]
  40. Goodall, P. T., Alouani-Bibi, F., & Blundell, K. M. 2011, MNRAS, 414, 2838 [NASA ADS] [CrossRef] [Google Scholar]
  41. Goranskij, V. 2011, Peremennye Zvezdy, 31, 5 [NASA ADS] [Google Scholar]
  42. Green, D. W. E. 2006, Cavendish Laboratory (Cambridge, UK), 8661, 2 [Google Scholar]
  43. Gregory, P. C., Scott, W. K., Douglas, K., & Condon, J. J. 1996, ApJS, 103, 427 [NASA ADS] [CrossRef] [Google Scholar]
  44. Guenette, R., & for the VERITAS Collaboration 2009, ArXiv e-prints [arxiv:0908.0714] [Google Scholar]
  45. Hayashi, S., Kajino, F., Naito, T., et al. 2009, Astropart. Phys., 32, 112 [NASA ADS] [CrossRef] [Google Scholar]
  46. Heinz, S., & Sunyaev, R. 2002, A&A, 390, 751 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  47. H.E.S.S. Collaboration, (Acero, F., Aharonian, F., et al.) 2009, A&A, 508, 1135 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  48. H.E.S.S. Collaboration (Abdalla, H., Abramowski, A., et al.) 2018a, A&A, 612, A10 (H.E.S.S. SI) [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  49. H.E.S.S. Collaboration (Abdalla, H., Abramowski, A., et al.) 2018b, A&A, 612, A2 (H.E.S.S. SI) [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  50. Hillas, A. M. 1985, International Cosmic Ray Conference, 3, 445 [NASA ADS] [Google Scholar]
  51. Kaufman Bernadó, M. M., Romero, G. E., & Mirabel, I. F. 2002, A&A, 385, L10 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  52. Kotani, T., Kawai, N., Aoki, T., et al. 1994, PASJ, 46, L147 [NASA ADS] [Google Scholar]
  53. Levinson, A., & Blandford, R. 1996, ApJ, 456, L29 [NASA ADS] [CrossRef] [Google Scholar]
  54. Lockman, F. J., Blundell, K. M., & Goss, W. M. 2007, MNRAS, 381, 881 [NASA ADS] [CrossRef] [Google Scholar]
  55. Margon, B. 1984, ARA&A, 22, 507 [NASA ADS] [CrossRef] [Google Scholar]
  56. Margon, B., & Anderson, S. F. 1989, ApJ, 347, 448 [NASA ADS] [CrossRef] [Google Scholar]
  57. Marshall, H. L., Canizares, C. R., & Schulz, N. S. 2002, ApJ, 564, 941 [NASA ADS] [CrossRef] [Google Scholar]
  58. Migliari, S., Fender, R., & Méndez, M. 2002, Science, 297, 1673 [NASA ADS] [CrossRef] [PubMed] [Google Scholar]
  59. Moldowan, A., & Safi-Harb, S. 2005, JRASC, 99, 141 [Google Scholar]
  60. Nicholas, B., & Rowell, G. 2008, in eds. F. A. Aharonian, W. Hofmann, & F. Rieger, AIP Conf. Ser., 1085, 245 [NASA ADS] [CrossRef] [Google Scholar]
  61. Orellana, M., Bordas, P., Bosch-Ramon, V., Romero, G. E., & Paredes, J. M. 2007, A&A, 476, 9 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  62. Reynoso, M. M., Christiansen, H. R., & Romero, G. E. 2008a, Astropart. Phys., 28, 565 [NASA ADS] [CrossRef] [Google Scholar]
  63. Reynoso, M. M., Romero, G. E., & Christiansen, H. R. 2008b, MNRAS, 387, 1745 [NASA ADS] [CrossRef] [Google Scholar]
  64. Rolke, W. A., López, A. M., & Conrad, J. 2005, Nucl. Instr. Meth. Phys. Res. A, 551, 493 [Google Scholar]
  65. Rowell, G. P., & HEGRA Collaboration 2001, ArXiv eprints [arxiv:astro-ph/0104288] [Google Scholar]
  66. Safi-Harb, S., & Ögelman, H. 1997, ApJ, 483, 868 [NASA ADS] [CrossRef] [Google Scholar]
  67. Safi-Harb, S., & Petre, R. 1999, ApJ, 512, 784 [NASA ADS] [CrossRef] [Google Scholar]
  68. Saito, T. Y., Zanin, R., Bordas, P., et al. 2009, ArXiv eprints [arxiv:0907.1017] [Google Scholar]
  69. Torres, D. F., & Reimer, A. 2011, A&A, 528, L2 [NASA ADS] [CrossRef] [EDP Sciences] [Google Scholar]
  70. Yamamoto, H., Ito, S., Ishigami, S., et al. 2008, PASJ, 60, 715 [NASA ADS] [Google Scholar]
  71. Zealey, W. J., Dopita, M. A., & Malin, D. F. 1980, MNRAS, 192, 731 [NASA ADS] [CrossRef] [Google Scholar]
  72. Zwitter, T., Calvani, M., & D’Odorico, S. 1991, A&A, 251, 92 [NASA ADS] [Google Scholar]

All Tables

Table 1

Observations of SS 433 performed by H.E.S.S. and MAGIC telescopes.

Table 2

Integral H.E.S.S. and MAGIC flux upper limits derived for SS 433 during low-absorption orbital/precessional phases and for the eastern/western interaction regions indicated in Fig. 1 using all available data.

All Figures

thumbnail Fig. 1

Significance map, derived from the H.E.S.S. data, for the FoV centred at the position of SS 433/W50 at E ≥ 287 GeV. GB6 4.85 GHz radio contours (white, from Gregory et al. 1996) are superimposed. Cyan circles indicate the positions of the interaction regions e1, e2, e3 (eastern “ear”) and w1, w2 (western “ear”). The bright source located north-west of SS 433 is MGRO J1908+06 (Abdo et al. 2007).

In the text
thumbnail Fig. 2

Differential flux upper limits (95% C.L.) from SS 433 obtained with MAGIC (blue), H.E.S.S. (green) and a combination of both telescopes (red) assuming a power-law with a spectral index Γ = 2.7 for the differential gamma-ray flux. The predicted differential gamma-ray flux from Reynoso et al. (2008b) for precessional phases Ψpre ∈ [0.9, 0.1] in which absorption of VHE emission should be at its lower level is also displayed (dashed orange), together with the Crab Nebula flux, for reference (from Aleksić et al. 2012b).

In the text

Current usage metrics show cumulative count of Article Views (full-text article views including HTML views, PDF and ePub downloads, according to the available data) and Abstracts Views on Vision4Press platform.

Data correspond to usage on the plateform after 2015. The current usage metrics is available 48-96 hours after online publication and is updated daily on week days.

Initial download of the metrics may take a while.