Another look at the BL Lacertae flux and spectral variability
C. M. Raiteri, M. Villata, L. Bruschini, A. Capetti, O. M. Kurtanidze, V. M. Larionov, P. Romano, S. Vercellone, I. Agudo, H. D. Aller, M. F. Aller, A. A. Arkharov, U. Bach, A. Berdyugin, D. A. Blinov, M. Böttcher, C. S. Buemi, P. Calcidese, D. Carosati, R. Casas, W.-P. Chen, J. Coloma, C. Diltz, A. Di Paola, M. Dolci, N. V. Efimova, E. Forné, J. L. Gómez, M. A. Gurwell, A. Hakola, T. Hovatta, H. Y. Hsiao, B. Jordan, S. G. Jorstad, E. Koptelova, S. O. Kurtanidze, A. Lähteenmäki, E. G. Larionova, P. Leto, E. Lindfors, R. Ligustri, A. P. Marscher, D. A. Morozova, M. G. Nikolashvili, K. Nilsson, J. A. Ros, P. Roustazadeh, A. C. Sadun, A. Sillanpää, J. Sainio, L. O. Takalo, M. Tornikoski, C. Trigilio, I. S. Troitsky, G. Umana
aa r X i v : . [ a s t r o - ph . C O ] S e p Astronomy&Astrophysicsmanuscript no. bl08˙le c (cid:13)
ESO 2018October 22, 2018
Another look at the BL Lacertae flux and spectral variability
Observations by GASP-WEBT, XMM-Newton, and Swift in 2008–2009 ⋆ C. M. Raiteri , M. Villata , L. Bruschini , A. Capetti , O. M. Kurtanidze , V. M. Larionov , , , P. Romano ,S. Vercellone , I. Agudo , , H. D. Aller , M. F. Aller , A. A. Arkharov , U. Bach , A. Berduygin , D. A. Blinov ,M. B ¨ottcher , C. S. Buemi , P. Calcidese , D. Carosati , R. Casas , , W.-P. Chen , J. Coloma , C. Diltz , A. DiPaola , M. Dolci , N. V. Efimova , , E. Forn´e , J. L. G ´omez , M. A. Gurwell , A. Hakola , T. Hovatta , ,H. Y. Hsiao , , B. Jordan , S. G. Jorstad , E. Koptelova , S. O. Kurtanidze , A. L¨ahteenm¨aki , E. G. Larionova ,P. Leto , E. Lindfors , R. Ligustri , A. P. Marscher , D. A. Morozova , M. G. Nikolashvili , K. Nilsson ,J. A. Ros , P. Roustazadeh , A. C. Sadun , A. Sillanp¨a¨a , J. Sainio , L. O. Takalo , M. Tornikoski , C. Trigilio ,I. S. Troitsky , and G. Umana (A ffi liations can be found after the references) ABSTRACT
Aims.
In a previous study we suggested that the broad-band emission and variability properties of BL Lacertae can be accounted for by a doublesynchrotron emission component with related inverse-Compton emission from the jet, plus thermal radiation from the accretion disc. Here weinvestigate the matter with further data extending over a wider energy range.
Methods.
The GLAST-AGILE Support Program (GASP) of the Whole Earth Blazar Telescope (WEBT) monitored BL Lacertae in 2008–2009at radio, near-IR, and optical frequencies to follow its flux behaviour. During this period, high-energy observations were performed by XMM-Newton, Swift, and Fermi. We analyse these data with particular attention to the calibration of Swift UV data, and apply a helical jet model tointerpret the source broad-band variability.
Results.
The GASP-WEBT observations show an optical flare in 2008 February–March, and oscillations of several tenths of mag on a few-daytime scale afterwards. The radio flux is only mildly variable. The UV data from both XMM-Newton and Swift seem to confirm a UV excessthat is likely caused by thermal emission from the accretion disc. The X-ray data from XMM-Newton indicate a strongly concave spectrum,as well as moderate ( ∼ ff erence between the source SEDs in 2008 and 1997can be explained in terms of pure geometrical variations. The outburst state occurred when the jet-emitting regions were better aligned with theline of sight, producing an increase of the Doppler beaming factor. Conclusions.
Our analysis demonstrates that the jet geometry can play an extremely important role in the BL Lacertae flux and spectral variability.Indeed, the emitting jet is probably a bent and dynamic structure, and hence changes in the emitting regions viewing angles are likely to happen,with strong consequences on the source multiwavelength behaviour.
Key words. galaxies: active – galaxies: BL Lacertae objects: general – galaxies: BL Lacertae objects: individual: BL Lacertae – galaxies: jets
1. Introduction
Blazars are active galactic nuclei whose extreme properties arethought to be owing to their relativistic jets pointing toward us.BL Lacertae, the prototype of the “BL Lac objects” blazar class,has been the target of many campaigns by the Whole EarthBlazar Telescope (WEBT) collaboration since 1999. The tensof thousands of optical-to-radio data collected by the WEBTallowed us to study its multiwavelength flux variability, colourbehaviour, and the correlations among flux variations in di ff erentbands, and revealed a possible periodicity of the radio out-bursts. The results have been published by Villata et al. (2002); Send o ff print requests to : C. M. Raiteri ⋆ The radio-to-optical data presented in this paper are storedin the GASP-WEBT archive; for questions regarding their avail-ability, please contact the WEBT President Massimo Villata( [email protected] ). Ravasio et al. (2002); B¨ottcher et al. (2003); Villata et al.(2004b,a); Bach et al. (2006); Papadakis et al. (2007);Villata et al. (2009b); Larionov et al. (2010).In a recent paper, Raiteri et al. (2009) analysed the mul-tiwavelength data from the 2007–2008 WEBT campaign, in-cluding three pointings by XMM-Newton. The XMM-Newtondata revealed a UV excess, which was interpreted to be dueto thermal emission from the accretion disc, as well as a spec-tral curvature in the X-ray band. The authors constructed spec-tral energy distributions (SEDs) of BL Lacertae correspondingto various epochs where the source was in di ff erent brightnessstates, using both their own data and data from the literature.They applied the inhomogeneous, rotating helical jet model byVillata & Raiteri (1999, see also Raiteri et al. 1999, Raiteri et al.2003, Ostorero et al. 2004) to fit the SEDs, and suggested thatthe broad-band spectral properties of BL Lacertae may resultfrom the combination of two synchrotron emission componentswith their self inverse-Compton emission, plus a thermal com-
1. M. Raiteri et al.: Another look at the BL Lacertae flux and spectral variability
Fig. 1. R -band light curve bythe GASP collaboration fromFebruary 2008 to February2009. Data are not correctedfor the host galaxy contamina-tion or Galactic extinction. Theparticipating observatories aremarked with di ff erent symbolsand colours. The total numberof data points is indicated in theupper right. ponent from the disc. Subsequently, Capetti et al. (2010) anal-ysed optical spectra acquired in the same period with the 3.56m Telescopio Nazionale Galileo (TNG). They found a broad H α emission line, with luminosity of ∼ × erg s − and FWHMof ∼ − , even brighter than that found in 1995–1997by Vermeulen et al. (1995) and Corbett et al. (1996, 2000). Thisfavours the hypothesis that the UV excess is caused by ther-mal emission from the accretion disc, the most likely sourceof ionising photons for the broad line region. The multiwave-length data available for the Raiteri et al. (2009) analysis lackedsimultaneous information in the γ -ray band, so that the inverse-Compton spectral region was poorly constrained. But in 2008the Fermi satellite was able to detect BL Lacertae (Abdo et al.,2010a), even if in a low state compared to the past detections bythe Compton Gamma Ray Observatory (CGRO, Hartman et al.1999, Bloom et al. 1997). In the same period, observations in theUV and X-ray bands were performed by Swift, while in the op-tical, near-IR, mm and cm radio bands the source was monitoredby the GLAST-AGILE Support Program (GASP) of the WEBT.This o ff ered the unique opportunity to study the source emissionover a very extended spectral range. The results of this furtherinvestigation e ff ort on BL Lacertae are presented in this paper.
2. GASP observations
The GASP was born in 2007 as a WEBT project, with the aimof monitoring a list of 28 γ -ray loud blazars in the optical, near-IR, mm, and cm radio bands during the γ -ray observations ofthe AGILE and Fermi (formerly GLAST) satellites (see e.g.Villata et al., 2008, 2009a). Data are collected periodically bythe WEBT President, who checks the consistency of the vari-ous datasets. The GASP light curves are then available for mul-tiwavelength studies, mostly in the framework of the GASPcollaboration with the AGILE and Fermi research teams. TheGASP data presented in this paper were taken at the observato-ries listed in Table 1.The optical data were calibrated with respect to a com-mon choice of reference stars in the same field of the source(Bertaud et al. 1969 in U and B bands; Fiorucci & Tosti 1996 in http://heasarc.gsfc.nasa.gov/docs/cgro/ http://agile.iasf-roma.inaf.it/ http://fermi.gsfc.nasa.gov/ Table 1.
List of optical, near-IR, and mm–cm radio observatoriescontributing data to this work.
Optical and near-infrared
Observatory Tel. size Bands[cm]Abastumani, Georgia 70 R Armenzano, Italy 35
BRI
Armenzano, Italy 40
BVRI
Calar Alto, Spain a R Campo Imperatore, Italy 110
JHK
Crimean, Ukraine 70
BVRI
El Vendrell, Spain 20 R Kitt Peak (MDM), USA 130
UBVRI
L’Ampolla, Spain 36 R Lulin, Taiwan 40 R New Mexico Skies, USA 30
VRI
Roque (KVA), Spain 35 R Sabadell, Spain 50 R St. Petersburg, Russia 40
BVRI
Talmassons, Italy 35
BVR
Teide (BRT), Spain 35
BVR
Tuorla, Finland 103 R Valle d’Aosta, Italy 81
BVRIRadio
Observatory Tel. size Frequencies[m] [GHz]Mauna Kea (SMA), USA 8 × b a Calar Alto data were acquired as part of the MAPCAT (MonitoringAGN with Polarimetry at the Calar Alto Telescopes) project. b Radio interferometer including 8 dishes of 6 m size. V , R , and I ). The source photometry was evaluated from a circu-lar region with an 8 arcsec aperture radius, while the backgroundwas taken in a surrounding annulus with 10 and 16 arcsec radii.In this way the measure is essentially seeing-independent and alldatasets are a ff ected by the same contamination from the lightof the host galaxy. Raiteri et al. (2009) estimated that with theabove prescriptions the contamination amounts to about 60%of the host total flux density, which is 0.36, 1.30, 2.89, 4.23,5.90, 11.83, 13.97, and 10.62 mJy in the U , B , V , R , I , J , H ,
2. M. Raiteri et al.: Another look at the BL Lacertae flux and spectral variability
Fig. 2.
June–November 2008 light curves of BLLacertae from UV to near-IR after correctionfor the host galaxy contribution, but not for theGalactic extinction. Data taken by the GASP-WEBT collaboration are plotted as blue circles,UVOT data as red triangles. The UVOT u , b ,and v light curves have been shifted to matchthe ground-based U , B , and V ones (see text fordetails). and K bands, respectively. When converting magnitudes intoflux densities, we corrected for the Galactic extinction accord-ing to the Cardelli et al. (1989) laws, using R V = .
1, the stan-dard value for the di ff use interstellar medium, and A B = . R -band light curvefrom February 2008 to February 2009 built with GASP data,which are not corrected for the host galaxy contribution here.A noticeable flare was observed at the beginning of the period,in 2008 February–March; afterwards both the average bright-ness level and the variability amplitude decreased. However,the source remained active, its brightness oscillating by severaltenths of magnitude on a few-day time scale. This is not an un-usual behaviour for BL Lacertae (see e.g. Raiteri et al., 2009,who reported on a 0.9 mag brightening in 24 hours).Optical data at other wavelengths as well as near-IR data areshown in Fig. 2 for the period June–November 2008 (the UV data displayed in the figure are presented in Sect. 3). The hostgalaxy contribution has been subtracted to distinguish the be-haviour of the active nucleus. This reveals that the brightnessevolution follows the same trend in the various bands, but mag-nitude variations are more pronounced at higher frequencies,which is a common feature of BL Lac objects. For example, thebrightness increase following the almost symmetric dip aroundJD ∼ ∼ B , V , R , and I bands, respectively. Wenotice that the source redshift is z = . α emission line enters the tailsof the R and I passbands. However, referring to Capetti et al.(2010), one can estimate that its flux contribution is only a fewthousandths of that of the continuum. Hence, the presence of theline cannot a ff ect the variability in these bands, at least when thesource brightness is at these levels. Another interesting exam-ple of fast variability is the rise of ∼ .
3. M. Raiteri et al.: Another look at the BL Lacertae flux and spectral variability JD = . R band, which unfortu-nately was not observed in other bands.Radio flux densities at cm–mm wavelengths are displayedin Fig. 3 together with de-reddened and host-galaxy subtracted R -band optical flux densities. The former are complemented bydata from the VLA / VLBA Polarization Calibration Database (PCD). The flux variation amplitude appears to decrease fromthe highest to the lowest frequencies, as usual. One interest-ing feature is the fast radio flare that is visible in the 37 GHzlight curve at JD ∼ ff ected by large uncertainties dueto unfavourable weather conditions; but an increase of the radioflux is visible also at 43, 22, and 14.5 GHz, giving strength tothe possibility that these events are correlated. We also noticethat there is neither a contemporaneous nor a delayed clear radiocounterpart to the optical flare observed at the beginning of theperiod (JD ∼
3. Swift observations
The Swift satellite observed BL Lacertae in 2008 August,September, and October, for a total of 24 pointings. In particular,from August 20 to September 9 a daily sampling was obtained.
The Ultraviolet / Optical Telescope (UVOT; Roming et al. 2005)onboard the Swift spacecraft acquires data in the optical v , b ,and u bands, as well as in the UV filters uvw uvm
2, and uvw uvotmaghist andthen binned the results. These values were compared with thoseobtained by first summing the frames acquired in the same bandwith uvotimsum , and then performing the aperture photometrywith the task uvotsource . We verified that the two methods areequivalent.The final UVOT light curves are shown in Fig. 2. We sub-tracted the host galaxy contribution, taking into account thatwith the 5 arcsec aperture radius we used for the photometry,about 50% of the host flux was included. We adopted the hostgalaxy flux densities given by Raiteri et al. (2009, see also Sect.2); these authors also discussed that the host contribution can beconsidered negligible in the UV. Fig. 3.
Optical flux densities ( R band, top panel), after correc-tion for Galactic extinction and host galaxy contamination, com-pared to radio flux densities at di ff erent frequencies. Blue cir-cles represent GASP data; red diamonds indicate data from theVLA / VLBA PCD. The vertical green line corresponds to theXMM-Newton pointing of May 16–17; the yellow strip high-lights the period of Swift observations, from 2008 August 20 toOctober 2.The comparison between the UVOT u , b , and v data and the U , B , and V data taken by the GASP observers reveals that ano ff set is present between the space light curves and the ground-based ones. We estimated mean o ff sets U − u = . B − b = . V − v = − .
05. The UVOT light curves shown in Fig. 2have been shifted accordingly. Taking into account that the av-erage UVOT colour indices of BL Lacertae are u − b ∼ − . b − v ∼ .
8, the above o ff sets disagree with those derivedby Poole et al. (2008) for the objects on which they based theirphotometric calibration of UVOT, i.e. Pickles stars and GRBmodels. Indeed, these objects have a di ff erent spectral shape,so that the Poole et al. (2008) calibrations may not hold for BLLacertae.The UVOT data confirm the variability trend traced by theground-based ones, extending it to UV frequencies. This in-dicates that the variability mechanism a ff ecting the near-IR–optical emission, which is dominated by beamed synchrotronradiation, can also produce flux changes in the UV, where a con-tribution from the synchrotron emission is thus expected, besidesa possible contribution from thermal disc radiation.
4. M. Raiteri et al.: Another look at the BL Lacertae flux and spectral variability
Fig. 4.
Top: Observed spectrum of BL Lacertae in the optical–UV band. Red diamonds refer to the mean UVOT spectrum re-sulting from 16 epochs analysed in this paper. Blue circles rep-resent the average spectrum obtained from the three observa-tions of the OM instrument onboard XMM-Newton in 2007–2008 (Raiteri et al., 2009), normalised to the mean UVOT spec-trum in the v band. The solid line is the log-parabolic fit usedin the calibration procedure. Bottom: Optical–UV SEDs of BLLacertae. Blue circles and red diamonds are derived from theOM and UVOT average spectra shown in the top panel by usingstandard prescriptions to obtain dereddened flux densities. Blacksquares represent the mean UVOT SED after recalibration as ex-plained in the text. In both panels the filter labels are centred onthe standard UVOT λ e ff (Poole et al., 2008).The derivation of the source intrinsic flux densities for fur-ther analysis (see Sect. 5) requires some attention. In their pa-per on the photometric calibration of UVOT, Poole et al. (2008)give e ff ective wavelengths of 5402, 4329, 3501, 2634, 2231, and2030 Å for the v , b , u , uvw uvm
2, and uvw λ e ff of the UV filters will be longer forvery red spectra. Moreover, they provide count-rate-to-flux con-version factors for both Pickles stars and GRB models, but inthe UV bands their validity range is limited to b − v = . b − v = .
03, respectively, while BL Lacertae has b − v ∼ . ff ective wavelengths λ e ff and count-rate-to-flux conversion factors CF for the UVOT filters by fold-ing the BL Lacertae spectrum with their e ff ective areas (seePoole et al., 2008). We first built a composite observed spec-trum of BL Lacertae by combining a mean OM spectrum (ob-tained from the three XMM-Newton pointings of 2007–2008,Raiteri et al. 2009) with an average UVOT spectrum (resulting from 16 UVOT observing epochs analysed in this paper). Tocompensate for the di ff erent brightness state, we increased theOM flux densities by ∼
6% so that the two spectra match inthe V band. The composite spectrum is shown in Fig. 4 (toppanel), together with its log-parabolic fit that we used in thefolding procedure. The resulting e ff ective wavelengths (see Eq.8 in Poole et al. 2008) are: 5439, 4381, 3500, 2776, 2295, and2225 Å for the v , b , u , uvw uvm
2, and uvw . × − erg cm − s − Å − from the v to the uvw < ∼ uvw
1) and CF( uvw .The new λ e ff would produce a decrease of extinction in the uvw uvm uvw λ ∼ ff ective area and BL Lacertae spectrum, similarly to whatwas done above for the λ e ff and CF: A Λ = . R d λ E Λ ( λ ) F λ ( λ ) 10 A ( λ ) / . R d λ E Λ ( λ ) F λ ( λ ) , (1)where A Λ is the extinction in the Λ band, E Λ is the e ff ective areaof that band, and F λ is the source flux density. The result is aGalactic extinction of 1.10, 1.44, 1.74, 2.40, 3.04, and 2.92 magfrom the v to the uvw λ e ff . Thebottom panel of Fig. 4 shows the mean SED obtained after recal-ibration of the UVOT data according to our procedure. For com-parison, we also show the OM and UVOT SEDs derived from theaverage spectra shown in the top panel, for which the amount ofGalactic extinction was calculated from the Cardelli et al. (1989)law at the standard λ e ff . Notice that the recalibration process hasshifted λ e ff ( uvw
2) redward so much that it overlaps with the stan-dard λ e ff ( uvm uvw uvw ff erent redshifts (see e.g. Villata et al., 2008; Raiteri et al.,2008; D’Ammando et al., 2009). Moreover, it seems to confirmthe UV excess claimed by Raiteri et al. (2009) that was ascribedto thermal emission from the accretion disc, even if this excessmay be less pronounced than indicated by the OM data. Ouranalysis highlights the importance of calculating the amount ofextinction in the critical UV bands, close to the 2175 Å bump,by folding the Galactic mean extinction law through the e ff ec-tive area curves and source spectrum. In any case, as pointed outby Fitzpatrick & Massa (2007), one has to keep in mind that theuse of an average dereddening curve implies a significant errorowing to the scatter of Galactic extinction curves. Notice that our CF( uvw
1) corresponds to that given by Poole et al.(2008) for the Pickles stars. 5. M. Raiteri et al.: Another look at the BL Lacertae flux and spectral variability c oun t s / s / k e V Swift−XRT10.5 2 5−2−1012 χ Energy (keV)
Fig. 5.
Swift-XRT spectrum of BL Lacertae on 2008 August 29.The bottom panel shows the deviations of the observed data fromthe folded model (a single power law with fixed absorption) inunits of standard deviations.
The X-ray Telescope (XRT; Burrows et al. 2005) data were pro-cessed with version 0.12.3 of the xrtpipeline task containedin the FTOOLS package, applying standard screening criteria.Inspection of the light curves revealed that the count rate waslow, from 0.15 to 0.24 counts s − , so that observations were per-formed in photon counting mode, and no correction for pile-up was necessary. Source and background spectra were ex-tracted with xselect from a circular region of 20 pixel (47arcsec) radius centred on the source and from a surroundingannulus of 30 and 50 pixel radii, respectively. We used ver-sion 011 of the response matrix available in the HEASARCcalibration database (CALDB), and calculated the ancillary re-sponse file with xrtmkarf , using the exposure map created by xrtpipeline . The source spectra were binned with grppha tohave a minimum of 20 counts in each bin, and they were finallyanalysed with version 12.5.1 of the Xspec task, using the energychannels greater than 0.3 keV.Spectral analysis was performed for each observation fol-lowing Raiteri et al. (2009): we first fitted a single power lawwith free absorption , and then fixed the Galactic absorption to N H = . × cm − , which takes into account both atomic andmolecular column density. Statistics is not good enough to eval-uate if a double power law model can improve the fit. The resultsof spectral fitting on XRT data are reported in Table 2 for all ob-servations with an exposure longer than 3 ksec; Col. 1 gives thedate and start time of the observation; Col. 2 its duration; Col. 3the hydrogen column; Col. 4 the power law photon index; Col. 5the 1 keV flux density; Col. 6 the χ /ν (and degrees of freedom).One spectrum (August 29) is shown in Fig. 5.Fits with free absorption resulted in a very variable N H ,which is unlikely to correspond to a real change of absorption.The average and median values are 3.46 and 3 . × cm − ,respectively, confirming that the value assumed for the Galacticabsorption is quite reasonable. Hence, we favoured the secondmodel, whose χ /ν is usually smaller than in the N H -free case, We adopted the Tuebingen-Boulder ISM absorption model(Wilms et al., 2000).
Fig. 6.
Details of the multiwavelength behaviour of BL Lacertaein the period around the Swift observations. The X-ray flux den-sity (at 1 keV, top) is compared to that in the UV ( uvw R band), near-IR ( K band), and radio (37 GHz, bottom)frequency range.and that produces results with smaller errors (because of onedegree of freedom more). In only two cases (August 23 andSeptember 2) a double power law model with absorption fixedto the Galactic value clearly improved the fit.The photon index Γ ranges from 1.92 to 2.25, indicatinga spectrum that oscillates from moderately hard to moderatelysoft. The average value is 2.07, with standard deviation of 0.08.To understand whether these spectral changes correspond to realvariations or are owing to noise, we recall the definition of themean fractional variation F var = √ σ − δ /< f > (Peterson,2001), which is commonly used to characterise variability. Here < f > is the mean value of the variable we are analysing, σ its variance, and δ the mean square uncertainty. In our case, σ = .
006 is smaller than δ = . . µ Jy,with a mean value of 1.75 and standard deviation of 0.24. In thiscase F var = .
11, and the variations can be considered reliable.Multiwavelength light curves of BL Lacertae in the periodaround the Swift observations are shown in Fig. 6. The sourcebehaviour at 1 keV di ff ers from the common trend characterisingthe UV, optical, and near-IR bands. In particular, the X-ray fluxpeaks when the near-IR–UV fluxes reach a minimum. However,there are also similarities, like the flux increase at the beginningof the common observing period, and the final decrease. Thismay indicate that the 1 keV flux behaviour sometimes is relatedto the brightness changes that occur at lower wavelengths, whilein other cases another variability mechanism prevails. Indeed,according to Raiteri et al. (2009) this frequency domain receivesthe variable contribution of two di ff erent emission components(see also Sect. 5).
6. M. Raiteri et al.: Another look at the BL Lacertae flux and spectral variability
Table 2.
Spectral fitting to the XRT data from Swift observations in 2008. Only exposures longer than 3 ksec are considered. Foreach epoch the first line reports the result of the single power law model with free absorption, while the second line shows thatobtained when fixing the Galactic absorption to N H = . × cm − . Start Exp N H Γ F χ /ν ( ν )[s] [10 cm − ] [ µ Jy]2008-08-20 @ 15:19:01 5072 3 . + . − . . + . − . . + . − . . ± .
12 1 . ± .
12 0.84 (30)2008-08-21 @ 11:53:00 5200 3 . + . − . . + . − . . + . − . . ± .
10 1 . ± .
13 1.01 (40)2008-08-22 @ 00:45:00 5374 3 . + . − . . + . − . . + . − . . ± .
10 1 . ± .
14 0.88 (37)2008-08-23 @ 00:51:01 4859 1 . + . − . . + . − . . + . − . . ± .
13 1 . ± .
14 1.34 (34)2008-08-24 @ 13:50:00 5118 3 . + . − . . + . − . . + . − . . ± .
10 1 . ± .
14 0.93 (42)2008-08-25 @ 09:07:01 5013 3 . + . − . . + . − . . + . − . . ± .
09 1 . ± .
12 0.97 (46)2008-08-26 @ 09:16:00 4756 3 . + . − . . + . − . . + . − . . ± .
09 1 . ± .
14 0.64 (46)2008-08-27 @ 00:07:00 4197 2 . + . − . . + . − . . + . − . . ± .
11 1 . ± .
15 0.63 (38)2008-08-27 @ 23:59:00 3002 3 . + . − . . + . − . . + . − . . ± .
14 1 . ± .
18 0.68 (21)2008-08-29 @ 12:48:00 5828 3 . + . − . . + . − . . + . − . . ± .
08 1 . ± .
13 1.02 (59)2008-08-30 @ 08:05:00 5044 4 . + . − . . + . − . . + . − . . ± .
08 2 . ± .
14 1.13 (51)2008-08-31 @ 08:11:01 5100 3 . + . − . . + . − . . + . − . . ± .
09 2 . ± .
14 0.93 (52)2008-09-01 @ 08:17:01 5491 3 . + . − . . ± .
14 1 . + . . . ± .
08 1 . ± .
12 0.99 (55)2008-09-02 @ 05:11:01 4077 3 . + . − . . + . − . . + . − . . ± .
11 1 . ± .
14 0.98 (35)2008-09-03 @ 14:55:00 4791 3 . + . − . . + . − . . + . − . . ± .
10 1 . ± .
13 1.00 (45)2008-09-04 @ 07:06:00 5032 4 . + . − . . + . − . . + . − . . ± .
10 1 . ± .
14 1.01 (38)2008-09-05 @ 09:02:00 4491 4 . + . − . . + . − . . + . − . . ± .
10 2 . ± .
15 1.12 (40)2008-09-06 @ 05:58:31 4676 3 . + . − . . + . − . . + . − . . + . − . . ± .
14 1.13 (33)2008-09-08 @ 13:46:01 4860 2 . + . − . . + . − . . + . − . . ± .
13 1 . ± .
13 0.96 (30)2008-09-09 @ 02:37:00 4501 4 . + . − . . + . − . . + . − . . + . − . . ± .
13 1.17 (28) 7. M. Raiteri et al.: Another look at the BL Lacertae flux and spectral variability
Fig. 7.
UV and X-ray light curves obtained by the OM and EPICinstruments onboard XMM-Newton on 2008 May 16–17. Black,red, and blue filled circles represent UV W
1, UV M
2, and UV W
4. XMM-Newton observations
The X-ray Multi-Mirror Mission (XMM) - Newton satellite ob-served the source during revolution 1545, on 2008 May 16–17,with a total exposure of ∼
134 ks. Data were processed with theScience Analysis System (SAS) package version 9.0.
The Optical Monitor (OM; Mason et al., 2001) onboard XMM-Newton is a 30-cm telescope carrying six optical / UV filters, andtwo grisms. BL Lacertae observations in May 2008 consisted of10 subsequent exposures in UV W
1, followed by 9 in UV M W
2. All exposures were ∼ omichain to reduce the data and the tasks omsource and omphotom to derive the source magnitude. Theerror on the aperture photometry is 0.03, 0.04, and 0.09 magfor the UV W
1, UV M
2, and UV W W = .
45, UV M = .
25, and UV W = . ff ectivewavelengths of the OM filters (2910, 2310, and 2120 Å for theUV W
1, UV M
2, and UV W −3 c oun t s / s / k e V MOS1+MOS2+pn XMM−Newton: May 16−17, 20081 100.5 2 5−202 χ Energy (keV)
Fig. 8.
EPIC spectrum of BL Lacertae on 2008 May 16–17;black squares, red triangles, and green diamonds representMOS1, MOS2, and pn data, respectively. The bottom panelshows the deviations of the observed data from the folded model(a double power law with fixed absorption) in units of standarddeviations.de-reddened magnitudes into flux densities was done with re-spect to Vega.
The European Photon Imaging Camera (EPIC) onboard XMM-Newton includes three detectors: MOS1, MOS2 (Turner et al.,2001), and pn (Str¨uder et al., 2001). All instruments were usedwith a thin filter. The two MOS cameras observed in small-window imaging mode, while pn was used in timing mode.We followed the standard prescription to reduce the data, in-cluding filtering of high background periods with a threshold of0.35 counts s − for MOS, but with a stricter threshold of 0.1counts s − for pn.For both MOS1 and MOS2, we created a filtered sky image,and extracted the source counts from a 50 arcsec radius circularregion, while background was evaluated in a circle on an ex-ternal CCD. As for pn, we extracted the source counts from astrip between RAWX =
35 and 39, and the background from twostrips at columns 24–28 and 48–52. To get the most reliable andbest calibrated events, we used the FLAG == < = ff ects were not a ff ecting the MOS datawith the epatplot task.Through the grppha task of the FTOOL package we binnedthe source spectra to have a minimum of 25 counts in each binand then analysed them together by means of the Xspec taskof the XANADU package. Only spectral bins corresponding toenergies between 0.3 and 10 keV for MOS1 and MOS2 and inthe range 0.5–10 keV for pn were considered, because they haveboth a better calibration and a higher signal-to-noise ratio.The three EPIC spectra were analysed together by first fit-ting a single power law with free absorption, and then fixing theGalactic absorption to N H = . × cm − . We also tried adouble power law with the same Galactic absorption. The re-sults are shown in Table 3 (see also Fig. 8). In agreement withRaiteri et al. (2009), the χ /ν suggests that a single power law
8. M. Raiteri et al.: Another look at the BL Lacertae flux and spectral variability
Table 3.
Spectral fitting to the EPIC (pn + MOS1 + MOS2) datafrom the XMM-Newton observations of 2008 May 16–17. Thefirst line reports the result of the single power law model withfree absorption; the second shows that obtained when fixing theGalactic absorption to N H = . × cm − , and the third thatfrom a double power law model with the same Galactic absorp-tion. N H Γ F χ /ν ( ν )[10 cm − ] [ µ Jy]2.60 ± . ± .
01 1 . ± .
01 1.09 (2030)3.40 1 . ± .
007 1 . ± .
008 1.48 (2031)3.40 2 . + . − . , 1 . + . − . . + . − . , 0 . + . − . with Galactic absorption of N H = . × cm − does not rep-resent a good fit to the data. As for the other two fits, the doublepower law seems to better fit the data, which is also confirmedby a very low F-test probability of ∼ . × − . This implies astrong spectral curvature.To check for possible flux variations, we extracted X-raylight curves from the same source and background regions de-fined for the spectra, with the same selection expressions. Weconsidered only the events in the time intervals free from highbackground and belonging to the 0.3–10 keV energy range forMOS1 and MOS2, and 0.5–10 keV for pn. The source countswere corrected for the background and then binned in one hourintervals. The results are shown in Fig. 7, which also displaysthe behaviour of the background to check the reliability of theflux variations. The background increased significantly only inthe last 6 hours. Just before, at JD ∼ .
05, there is asmall flare clearly visible in all three light curves. Mean sourcerates for the whole period are 1.63, 0.61, and 0.64 counts s − forpn, MOS1, and MOS2, respectively, with standard deviations of0.15, 0.04, and 0.03 counts s − . This means fractional variations F var of 7%, 6%, and 4%, which do not change significantly if weexclude the last 6 hours. Hence, we can conclude that the X-rayflux of BL Lacertae is mildly variable on an hour time scale.
5. Modelling the SED
Figure 9 shows the broad-band SED of BL Lacertae in di ff er-ent brightness states. The SED corresponding to 2008 May 16–17 includes the XMM-Newton UV and X-ray data analysed inSect. 4. In order to avoid o ff sets caused by source variability, theOM spectrum was constructed with the last UV W M M W γ -ray satellitedetected BL Lacertae; the Fermi data we plotted in Fig. 9 werederived from Abdo et al. (2010a). Table 4.
Main parameters of the helical jet model for the fit toboth the 2008 faint-state SED and the 1997 outburst-state SED.For each epoch, “low” and “high” refer to the low- and high-energy synchrotron plus self inverse-Compton components, re-spectively.
SED 2008 SED 1997Parameter Low High Low High ζ ◦ ◦ ◦ ◦ a ◦ ◦ ◦ ◦ ψ ◦ ◦ . ◦ ◦ φ ◦ ◦ ◦ ◦ log ν ′ s (0) 14 17.8 14 17.8 c min , max l min − . − . − . − . l max − . − . − . − . γ max (0) 3.7 3.8 3.7 3.8 c γ l γ − . − . − . − . α Γ c s , c l s , c − − − − The August 2008 SEDs indicate a faint, synchrotron-dominated state of the source that we fitted with the rotating he-lical jet model by Villata & Raiteri (1999, see also Raiteri et al.1999, Raiteri et al. 2003, Ostorero et al. 2004). This model hasbeen used by Raiteri et al. (2009) to fit the broad-band SED ofBL Lacertae in December 2007 – January 2008. Their main find-ing was that the BL Lacertae broad-band SED cannot be ex-plained by a single synchrotron component plus its self inverse-Compton emission. Indeed, the very strong historical X-ray vari-ability requires an additional synchrotron (plus self inverse-Compton) component. Moreover, the UV excess suggests ther-mal contribution from the accretion disc (see Sect. 3.1). TheSED analysed by Raiteri et al. (2009) lacked simultaneous γ -raydata, which made it impossible to constrain the inverse-Comptonemission of the high-energy component. The 2008 August SEDin Fig. 9 now o ff ers us the possibility to perform a more detailedanalysis.In addition, we also display in Fig. 9 the broad-band SEDcorresponding to the big outburst of July 1997, which showeda considerable inverse-Compton dominance. The X-ray spectraplotted in the figure are the result of the combined analysis ofthe ASCA and RXTE data by Tanihata et al. (2000). Becauseof the very strong variability of the source in that period, theauthors distinguished between a low state, which was well fit-ted by a single power law model, and a flare state, for whichthe best fit was obtained with a double power law model. Thislast fit resulted in a very strong spectral curvature . In July 1997observations in the γ -ray band were performed by CGRO. Thedata from the EGRET instrument onboard CGRO in Fig. 9 weretaken from Bloom et al. (1997), while those from the OSSE de-tector were derived from the High Energy Astrophysics ScienceArchive Research Center (HEASARC). The low-frequency in-formation is from the WEBT archive; the range of optical fluxvariation in the period is indicated. This outburst state of thesource was fitted with the same rotating helical model that weused to fit the faint state of 2008. We notice that Tanihata et al. (2000) adopted a Galactic total ab-sorption of N H = . × cm − . http://heasarc.gsfc.nasa.gov/
9. M. Raiteri et al.: Another look at the BL Lacertae flux and spectral variability
Fig. 9.
Broad-band SEDs of BL Lacertae inAugust 2008 (blue) and July 1997 (red). The2008 SED is built with UV and X-ray datafrom two epochs of Swift observations (Sect.3), 2008 August 25 and 30, together withsimultaneous optical-to-radio data from theGASP-WEBT collaboration, and contempora-neous γ -ray data from Fermi (from Abdo et al.,2010a). The 1997 SED includes EGRET datafrom Bloom et al. (1997), OSSE data fromthe HEASARC archive, ASCA + RXTE spec-tra from Tanihata et al. (2000), while low-frequency data are from the WEBT archive.Solid lines represent model fits; we distin-guish the low-energy emission component (dot-ted lines) from the high-energy one (dashedlines); the contribution by an accretion disc of ∼ × erg s − is marked witha dotted-dashed line. We also show in black theSED corresponding to the XMM-Newton ob-servations of May 2008 (Sect. 4); both the sin-gle power law with free absorption and doublepower law with Galactic absorption fits to theEPIC spectra are displayed. In performing the model fits our aim was to see whether itwas possible to reproduce the high and low states by changingthe geometrical configuration only. Moreover, we took into ac-count the results by Larionov et al. (2010), who explained thelong-term BL Lacertae optical and near-IR variability in termsof variations of the Doppler boosting factor due to changes ofthe viewing angle of the emitting region.The resulting model parameters are reported in Table 4,while the corresponding fits are shown in Fig. 9. We also in-cluded blackbody radiation from an accretion disc with a lu-minosity of 5 × erg s − and a temperature of ∼ ζ = ◦ , covering an angle a = ◦ . The maximum Lorentzfactor of the relativistic electrons is log γ max (0) = .
7, while thebulk Lorentz factor of the plasma in the jet is
Γ =
7. The high-energy (UV–X-ray and related inverse-Compton) emission com-ponent comes from a helix portion with a pitch angle ζ = ◦ ,covering an angle a = ◦ . The maximum Lorentz factor of therelativistic electrons is log γ max (0) = .
8, while the bulk Lorentzfactor is
Γ =
13. These parameters, as well as those definingthe power laws according to which the maximum and minimumemitted frequencies and the flux densities decrease with distancefrom the jet apex, are maintained fixed. The di ff erence betweenthe fits to the 1997 and 2008 SEDs is only due to a variation ofthe orientation of the jet emitting regions, through only two geo-metric parameters: the angle between the helix axis and the lineof sight ψ and the rotation angle φ . The outburst state requiresa better alignment of the emitting regions with the line of sight,which implies ψ approaching the helix pitch angle, and smallerrotation angles φ .
6. Discussion
The satisfactory fits that we obtained in the previous section forboth the outburst and faint states of BL Lacertae give strengthto our interpretation of the source SED in terms of two syn-chrotron plus self inverse-Compton emission components. In theSED, the synchrotron peak of the low-energy component falls inthe near-IR band, and its inverse-Compton reaches a maximum
Outburst state Faint state
Line of sight
Fig. 10.
Sketch of our helical jet model during both the 2008faint state and the 1997 outburst state. The angle between the jetaxis and the line of sight, ψ , has been multiplied by a factor 10with respect to the values given in Table 4 for clarity. We distin-guish the inner region, emitting the high-energy synchrotron plusself inverse-Compton component (purple-blue) from the outerzone, where the low-energy emission component is produced(green-yellow-red).in the 1–50 MeV energy range. The synchrotron and inverse-Compton peaks of the high-energy component occur in the far-UV–soft-X-ray band and in the energy range 0.2–5 GeV, respec-tively. Whether these two components come from two distincthelices or from di ff erent regions inside the same helical jet is notclear. We consider it more likely that there is a unique jet, wherethe high-energy component comes from a region closer to theemitting jet apex than the low-energy one. We notice that a dou-ble synchrotron component is not an unusual interpretation forthe blazar emission properties, as it has been proposed also forMkn 421 (Donnarumma et al., 2009) and 3C 454.3 (Ogle et al.,2010).
10. M. Raiteri et al.: Another look at the BL Lacertae flux and spectral variability
Moreover, the fits show that the whole range of BL Lacertaemultiwavelength variability can be interpreted in terms of orien-tation e ff ects. Although the rotating helical jet model we haveadopted in the previous section is not a physically completemodel, but more a phenomenological approach, it has the ad-vantage of taking into account variations of the orientation ofthe emitting regions with respect to the line of sight, with conse-quent changes of the Doppler beaming factor. This is an aspectthat is usually neglected by theoretical models of blazar emis-sion, which explain flux and spectral changes uniquely in termsof energetic processes inside the jet.Our interpretation is in line with previous results.Marscher et al. (2008) analysed the evolution of the BL Lacertaeoptical polarization during 2005, and suggested that the plasmaflows along helical streamlines. According to Villata et al.(2009b), the long-term optical and radio behaviour of BLLacertae suggests a scenario where the emitting plasma flowsalong a rotating helical path in a curved jet. This rotating he-lical structure could be caused by orbital motion in a binaryblack hole system, coupled with the interaction of the plasmajet with the surrounding medium. Indeed, the binary black holescenario could explain the periodicity observed in the radio lightcurves of BL Lacertae (Villata et al., 2004a, 2009b), the discov-ery of a precessing jet nozzle with the VLBA (Stirling et al.,2003), and possibly the parsec-to-kiloparsec jet misalignment(see e.g. Kharb et al., 2010). Moreover, the analysis of the BLLacertae spectral evolution in 2000–2008 by Larionov et al.(2010) favoured a picture where the optical and near-IR flux andcolour variability can be explained by a variable viewing angleof the emitting region. These authors also suggested that a fractalhelical structure may be at the origin of the di ff erent time scalesof variability.The values of the angles ζ = ◦ and ψ = . ◦ , as well asthe Lorentz factor Γ = ζ = ◦ found for the high-energycomponent indicates that this inner jet helical region would bemore twisted than the outer, lower-energy one. In practice, ourmodel results indicate a helical jet whose axis is bent betweenthe X-ray and optical regions by about 2 ◦ (see the values of ψ inTable 4) and that is more wrapped near the apex and then tendsto relax with a decreasing pitch angle. The di ff erent orientationsassumed by such a jet in 2008 and 1997 are sketched in Fig. 10,where all the angles are strongly increased for clarity.The thermal emission component that we added to the he-lical jet model is justified by the UV excess found in the OMdata from XMM-Newton, and (though with less evidence) in theUVOT data from Swift. As discussed in the present paper and inRaiteri et al. (2009), the amount of this excess strongly dependson both the Galactic extinction and instrument calibration, but itis not easy to cancel it out completely. In any case, Capetti et al.(2010) showed that after twelve years from the first detectionof the H α broad emission line by Vermeulen et al. (1995) andCorbett et al. (1996, see also Corbett et al. 2000), the H α line isstill there, even more luminous than before. This suggests that adisc is also there to photoionise the broad line region.Photons coming from the disc or broad line region could thenenter the jet, and be inverse-Compton scattered, giving rise toother high-energy emission components that are sometimes in-voked to account for the SED properties of blazars. In particu-lar, the 1997 outburst state has previously been interpreted byMadejski et al. (1999) in terms of three emission components:synchrotron, synchrotron self-Compton, and Comptonisation of the broad emission line flux. Similar results were obtained byB¨ottcher & Bloom (2000) and by Ravasio et al. (2002). Our“geometrical” interpretation does not require these external-Compton emission components, which are not expected to con-tribute if the jet emission regions are parsecs away from thecentral black hole (see e.g. Sikora et al., 2008; Marscher et al.,2010; Abdo et al., 2010b). Acknowledgements.
We acknowledge Ann E. Wehrle for useful comments.This research has made use of NASA’s Astrophysics Data System. TheTorino and Palermo teams acknowledge financial support by the Italian SpaceAgency through contract ASI-INAF I / / / / ST08 / References
Abdo, A. A., Ackermann, M., Agudo, I., et al. 2010a, ApJ, 716, 30Abdo, A. A., Ackermann, M., Ajello, M., et al. 2010b, Nature, 463, 919Bach, U., Villata, M., Raiteri, C. M., et al. 2006, A&A, 456, 105Bertaud, C., Dumortier, B., Veron, P., et al. 1969, A&A, 3, 436Bessell, M. S., Castelli, F., & Plez, B. 1998, A&A, 333, 231Bloom, S. D., Bertsch, D. L., Hartman, R. C., et al. 1997, ApJ, 490, L145B¨ottcher, M. & Bloom, S. D. 2000, AJ, 119, 469B¨ottcher, M., Marscher, A. P., Ravasio, M., et al. 2003, ApJ, 596, 847Burrows, D. N., Hill, J. E., Nousek, J. A., et al. 2005, Space Science Reviews,120, 165Capetti, A., Raiteri, C. M., & Buttiglione, S. 2010, ArXiv e-printsCardelli, J. A., Clayton, G. C., & Mathis, J. S. 1989, ApJ, 345, 245Clements, S. D., Smith, A. G., Aller, H. D., & Aller, M. F. 1995, AJ, 110, 529Corbett, E. A., Robinson, A., Axon, D. J., & Hough, J. H. 2000, MNRAS, 311,485Corbett, E. A., Robinson, A., Axon, D. J., et al. 1996, MNRAS, 281, 737D’Ammando, F., Pucella, G., Raiteri, C. M., et al. 2009, A&A, 508, 181Donnarumma, I., Vittorini, V., Vercellone, S., et al. 2009, ApJ, 691, L13Fiorucci, M. & Tosti, G. 1996, A&AS, 116, 403Fitzpatrick, E. L. & Massa, D. 2007, ApJ, 663, 320Hartman, R. C., Bertsch, D. L., Bloom, S. D., et al. 1999, ApJS, 123, 79Hufnagel, B. R. & Bregman, J. N. 1992, ApJ, 386, 473Kharb, P., Lister, M. L., & Cooper, N. J. 2010, ApJ, 710, 764Larionov, V. M., Villata, M., & Raiteri, C. M. 2010, A&A, 510, A93Madejski, G. M., Sikora, M., Ja ff e, T., et al. 1999, ApJ, 521, 145Marscher, A. P., Jorstad, S. G., D’Arcangelo, F. D., et al. 2008, Nature, 452, 966Marscher, A. P., Jorstad, S. G., Larionov, V. M., et al. 2010, ApJ, 710, L126Mason, K. O., Breeveld, A., Much, R., et al. 2001, A&A, 365, L36Ogle, P. M., Wehrle, A. E., Balonek, T., & Gurwell, M. A. 2010, ArXiv e-printsOstorero, L., Villata, M., & Raiteri, C. M. 2004, A&A, 419, 913Papadakis, I. E., Villata, M., & Raiteri, C. M. 2007, A&A, 470, 857Peterson, B. M. 2001, in Advanced Lectures on the Starburst-AGN Connection,ed. I. Aretxaga, D. Kunth, & R. M´ujica (Singapore: World Scientific), 3Poole, T. S., Breeveld, A. A., Page, M. J., et al. 2008, MNRAS, 383, 627Raiteri, C. M., Villata, M., Capetti, A., et al. 2009, A&A, 507, 769Raiteri, C. M., Villata, M., Chen, W. P., et al. 2008, A&A, 485, L17Raiteri, C. M., Villata, M., Tosti, G., et al. 1999, A&A, 352, 19Raiteri, C. M., Villata, M., Tosti, G., et al. 2003, A&A, 402, 151Ravasio, M., Tagliaferri, G., Ghisellini, G., et al. 2002, A&A, 383, 763Roming, P. W. A., Kennedy, T. E., Mason, K. O., et al. 2005, Space ScienceReviews, 120, 95Schlegel, D. J., Finkbeiner, D. P., & Davis, M. 1998, ApJ, 500, 525
11. M. Raiteri et al.: Another look at the BL Lacertae flux and spectral variability
Sikora, M., Moderski, R., & Madejski, G. M. 2008, ApJ, 675, 71Stirling, A. M., Cawthorne, T. V., Stevens, J. A., et al. 2003, MNRAS, 341, 405Str¨uder, L., Briel, U., Dennerl, K., et al. 2001, A&A, 365, L18Tanihata, C., Takahashi, T., Kataoka, J., et al. 2000, ApJ, 543, 124Tornikoski, M., Valtaoja, E., Ter¨asranta, H., & Okyudo, M. 1994a, A&A, 286,80Tornikoski, M., Valtaoja, E., Ter¨asranta, H., et al. 1994b, A&A, 289, 673Turner, M. J. L., Abbey, A., Arnaud, M., et al. 2001, A&A, 365, L27Vermeulen, R. C., Ogle, P. M., Tran, H. D., et al. 1995, ApJ, 452, L5Villata, M. & Raiteri, C. M. 1999, A&A, 347, 30Villata, M., Raiteri, C. M., Aller, H. D., et al. 2004a, A&A, 424, 497Villata, M., Raiteri, C. M., Gurwell, M. A., et al. 2009a, A&A, 504, L9Villata, M., Raiteri, C. M., Kurtanidze, O. M., et al. 2004b, A&A, 421, 103Villata, M., Raiteri, C. M., Kurtanidze, O. M., et al. 2002, A&A, 390, 407Villata, M., Raiteri, C. M., Larionov, V. M., et al. 2008, A&A, 481, L79Villata, M., Raiteri, C. M., Larionov, V. M., et al. 2009b, A&A, 501, 455Wilms, J., Allen, A., & McCray, R. 2000, ApJ, 542, 914 INAF, Osservatorio Astronomico di Torino, Italy Dipartimento di Fisica Generale, Universit`a di Torino, Italy Abastumani Observatory, Mt. Kanobili, Georgia Astron. Inst., St.-Petersburg State Univ., Russia Pulkovo Observatory, St. Petersburg, Russia Isaac Newton Institute of Chile, St.-Petersburg Branch INAF-IASF Palermo, Italy Institute for Astrophysical Research, Boston University, MA, USA Instituto de Astrof´ısica de Andaluc´ıa, CSIC, Granada, Spain Department of Astronomy, University of Michigan, MI, USA Max-Planck-Institut f¨ur Radioastronomie, Bonn, Germany Tuorla Observatory, Dept. of Physics and Astronomy, Univ. ofTurku, Finland Department of Physics and Astronomy, Ohio Univ., OH, USA INAF, Osservatorio Astrofisico di Catania, Italy Osservatorio Astronomico della Regione Autonoma Valle d’Aosta,Italy Armenzano Astronomical Observatory, Italy Inst. de Ci`encies de l’Espai (CSIC-IEEC), Spain Agrupaci´o Astron`omica de Sabadell, Spain Graduate Institute of Astronomy, National Central University,Taiwan Observatori El Vendrell, Spain INAF, Osservatorio Astronomico di Roma, Italy INAF, Osservatorio Astronomico di Collurania Teramo, Italy Harvard-Smithsonian Center for Astrophysics, Cambridge, MA,USA Aalto University Mets¨ahovi Radio Observatory, Finland Department of Physics, Purdue University, USA Lulin Observatory, National Central University, Taiwan School of Cosmic Physics, Dublin Institute For Advanced Studies,Ireland Circolo Astrofili Talmassons, Italy Finnish Centre for Astronomy with ESO (FINCA), University ofTurku, Finland30