Quasi-simultaneous two band optical variability of the blazars 1ES 1959+650 and 1ES 2344+514
Haritma Gaur, Alok C. Gupta, A. Strigachev, R. Bachev, E. Semkov, Paul J. Wiita, S. Peneva, S. Boeva, N. Kacharov, B. Mihov, E. Ovcharov
aa r X i v : . [ a s t r o - ph . C O ] D ec Mon. Not. R. Astron. Soc. , 1– ?? (2010) Printed 7 October 2018 (MN L A TEX style file v2.2)
Quasi − simultaneous two band optical variability of theblazars 1ES 1959+650 and 1ES 2344+514 Haritma Gaur , ⋆ , Alok C. Gupta , , A. Strigachev , R. Bachev , E. Semkov ,Paul J. Wiita , S. Peneva , S. Boeva , N. Kacharov , , B. Mihov , E. Ovcharov Aryabhatta Research Institute of Observational Sciences (ARIES), Manora Peak, Nainital – 263129, India Department of Physics, DDU Gorakhpur University, Gorakhpur - 273 009, India Institute of Astronomy and National Astronomical Observatory, Bulgarian Academy of Sciences, 72 Tsarigradsko Shosse Blvd.,1784 Sofia, Bulgaria Department of Physics, The College of New Jersey, P.O. Box 7718, Ewing, NJ 08628, USA Department of Astronomy, University of Sofia, 5 James Bourchier, 1164 Sofia, Bulgaria
Accepted ....... Received ......; in original form ......
ABSTRACT
We report the results of quasi-simultaneous two filter optical monitoring of two high-energy peaked blazars, 1ES 1959+650 and 1ES 2344+514, to search for microvariabil-ity and short-term variability (STV). We carried out optical photometric monitoringof these sources in an alternating sequence of B and R pass-bands, and have 24 and19 nights of new data for these two sources, respectively. No genuine microvariability(intra-night variability) was detected in either of these sources. This non-detection ofintra-night variations is in agreement with the conclusions of previous studies thathigh-energy peaked BL Lacs are intrinsically less variable than low-energy peaked BLLacs in the optical bands. We also report the results of STV studies for these twosources between July 2009 and August 2010. Genuine STV is found for the source1ES 1959+650 but not for 1ES 2344+514. We briefly discuss possible reasons for thedifference between the intra-night variability behaviour of high- and low-energy peakedblazars.
Key words: galaxies: active – BL Lacertae objects: general – BL Lacertae objects:individual: 1ES 1959+650, 1ES 2344+514 –optical photometry
Blazars constitute an enigmatic subclass of radio-loud activegalactic nuclei (AGNs). Blazars include both BL Lacertae(BL Lac) objects and flat spectrum radio quasars (FSRQs),as both are characterized by strong and rapid flux variabilityacross the entire electromagnetic spectrum. In addition, BLLacs show largely featureless optical continua. Blazars ex-hibit strong polarization from radio to optical wavelengthsand usually have core-dominated radio structures. Accord-ing to the orientation based unified model of radio-loudAGNs, blazar jets usually make an angle ◦ to our line-of-sight (e.g., Urry & Padovani 1995). These extreme AGNsprovide a natural laboratory to study the mechanisms of en-ergy extraction from the central super-massive black holesand the physical properties of astrophysical jets.The electromagnetic (EM) radiation from blazars ispredominantly non-thermal. At lower frequencies (throughthe UV or X-ray bands) the mechanism is almost certainly ⋆ E-mail: [email protected] synchrotron emission, while at higher frequencies the emis-sion mechanism is probably due to Inverse-Compton (IC)emission (e.g., Sikora & Madejski 2001; Krawczynski 2004).The spectral energy distributions (SEDs) of blazars havea double-peaked structure (e.g., Giommi, Ansari & Micol1995; Ghisellini et al. 1997; Fossati et al. 1998). Basedon the location of the first peak of their SEDs, ν peak ,and blazars are often sub-classified into low energy peakedblazars (LBLs) and high energy peaked blazars (HBLs)(Padovani & Giommi 1995). The first component peaks inthe near-infrared/optical for LBLs and in the UV or X-raysfor HBLs, while the second component usually peaks at GeVenergies for LBLs and at TeV energies for HBLs.Observations of blazars reveal that they are variable atall accessible timescales, from a few tens of minutes to years,and even decades, at many frequencies (e.g., Carini & Miller1992; Wagner & Witzel 1995; Gupta et al. 2004; Villataet al. 2007). These different timescales allow the variabil-ity of blazars to be broadly divided into three classes, e.g.,intra-day variability (IDV), short-term variability (STV),and long-term variability (LTV). Variations in the flux of c (cid:13) Gaur et al. a source of a couple of hundredths of a magnitude up to afew tenths of magnitude over a time scale of a day or less isknown as IDV (Wagner & Witzel 1995) or microvariabilityor intra-night optical variability. Flux changes observed overdays to a few months are often considered to be STV, whilethose taking from several months to many years are usuallycalled LTV (e.g., Gupta et al. 2004).For a better understanding of how the properties ofHBLs relate to the previously known LBLs, we need to dis-criminate among the theoretical models that try to explainthem. Optical observations offer a wealth of information onthe variability of blazars. Over the past many years, a largeamount of optical variability data for LBLs have been re-ported and rapid optical variability has been shown to becommon (Rani et al. 2010, 2011 and references therein).However, because most of the HBLs were more recentlyfound by X-ray and gamma-ray sky surveys (Perlman et al.2006, Abdo et al. 2010 and refernces therein) they have beenless extensively monitored, and the properties of the opticalvariability of HBLs are not yet clear. Preliminary observa-tions of HBLs suggested that they show statistically lesseramounts of optical variability (Heidt & Wagner 1996, 1998)and polarized light than that of LBLs (Stoke et al. 1985,1989, Jannuzi et al. 1993, 1994, Villata et al. 2000). Heidt& Wagner (1998) noticed these objects likely display differ-ent duty cycles and variability amplitudes from those of theLBLs. Romero, Cellone & Combi (1999) found such differ-ences in the case of microvariability, as just one out of threeHBLs in their sample showed micro-variations, while eightLBLs out of 12 did so. These differences were attributed tothe possible presence of stronger magnetic fields in the high-energy peaked BL Lacs which could prevent the formationof small-scale jet inhomogeneities in these objects. However,due to the relative paucity of optical monitoring for HBLs,we can not yet definitively say whether and how HBLs aredifferent from LBLs in optical variability and therefore anyadditional high quality observations of HBLs are particu-larly useful.Here we report quasi-simultaneous observations of twoHBLs, 1ES 1959+650 and 1ES 2344+514 in two opticalbands (B and R) for the first time. In addition to looking forany microvariability in these two optical bands, our observa-tions also give additional information on any microvariationsin (B − R) colour. We also have searched for short-term vari-ability timescales for these HBLs in the optical bands. Weused five telescopes located in India, Bulgaria and Greeceto monitor the optical variability of 1ES 1959+650 and 1ES2344+514 between July 2009 and November 2010.The paper is structured as follows. In Section 2, we dis-cuss the key features of these two HBLs. Section 3 describesthe observations and data reduction procedure. Section 4reports our results and we present our discussion and con-clusions in Section 5. + One of the best studied member of the HBL subclass is 1ES1959+650, with a redshift, z = 0 .
046 (Veron-Cetty & Veron2006). It was first detected at X-rays during the Slew Sur-vey made by the Einstein satellite’s Imaging Proportional Counter (Elvis et al. 1992). Based on the X-ray/radio ver-sus X-ray/optical colour-color diagram, the source was clas-sified as a BL Lac object by Schachter et al. (1993). Themass of the central black hole (BH) in 1ES 1959+650 hasbeen estimated to be ∼ × M ⊙ (Falomo et al. 2002).Heidt et al. (1999) performed a detailed study of theoptical band and found a complex structure composed ofa luminous elliptical galaxy (M R = −
23) plus a disk andan absorption dust lane. Three photometric optical valueswere found in the literature, showing a great variation ofthe source brightness in the optical band. Its brightness wasreported as: V = 16 by Marcha (1994); V = 12.8, as de-rived from digitalized plates, using B − V = 0.5 (Schachteret al. 1993); and an approximate magnitude of 13.67 ob-tained from the Cambridge Automatic Plate Measurementinstrument (Perlman et al. 1996). Kurtanidze et al. (1999)reported data taken during 1997 to 1999, finding bright-nesses in the range R = 14.50 to 14.92. Krawczynski et al.(2004) carried out multi-wavelength observations of strongflares from this source. Villata et al. (2000) observed thesource from February 29, 1996 to May 30, 1997, and recordeda rapid decrease of 0.28 mag in 4 days. More recently, Maet al. (2010) observed the source between 4 May 2003 to 19September 2003 but saw no significant optical variations.The blazar 1ES 1959+650 was further detected at X-rays with ROSAT & BeppoSAX (Beckmann et al. 2002). In2002 May, the X-ray flux of the source had increased sig-nificantly. Both the Whipple (Holder et al. 2003) and
HighEnergy Gamma Ray Astronomy (HEGRA) (Aharonian et al.2003) collaborations subsequently confirmed a higher VeryHigh Energy (VHE) γ -ray flux as well. Observing the sourcein 2000, 2001 and early 2002, the HEGRA collaborationyielded a marginal signal (Horns et al. 2002).An interesting aspect of the source activity in 2002, wasthe observations of a so-called “orphan flare” (i.e., a flare ofVHE γ -rays not accompanied by correlated increased activ-ity at other wavelengths), recorded on June 4 by the Whip-ple Collaboration (Krawczynski et al. 2004; Daniel et al.2005). The HEGRA collaboration had observed another, al-beit less significant, orphan VHE signal during moonlight 2days earlier (Tonello & Kranich 2003). Both of these flaresin VHE γ -rays, observed in the absence of high activity inX-rays, are very difficult to reconcile with the standard syn-chrotron self-Compton model (Ghisellini et al. 1998), whichis routinely very successfully employed to explain VHE γ -ray production. So, a hadronic synchrotron mirror modelwas proposed by B¨ottcher et al. (2005) to explain this or-phan TeV flare. 1ES 1959+650 was observed several timesat TeV energies (Aharonian et al. 2003, Albert et al. 2006,Tagliaferri et al. 2008, and references therein), and unsur-prisingly, this source is included in the first Fermi catalog ofAGNs (Abdo et al. 2010). + This object, at z = 0 . µ Jy was found while it had an optical brightness of V=15.5magnitude (with no galaxy subtraction). It was discovered inTeV γ -rays ( >
350 GeV) with the Whipple Observatory tele-scope and thus identified as a TeV BL Lac object (Cataneseet al. 1998). Giommi et al. (2000) reported rapid variability c (cid:13) , 1– ?? ptical IDV of HBLs in the X-ray band on a timescale of 5000 sec. The very highenergy gamma-ray emission from 1ES 2344+514 has beenobserved by the MAGIC collaboration (Albert et al. 2007)and Grube (2008) reported that this source showed a nightlyvariability in the integrated flux above 300 GeV.Miller et al. (1999) observed 1ES 2344+514 and founda positive detection of optical microvariability, with a max-imum range of about 0.08 magnitude over one night in theV-band. Measurements in R band from September 1998 toOctober 2000 by Kurtanidze and Nikolashvili (2002) haveshown LTV at the level of 0.1 mag; any intraday variabil-ity was confined within 0.05 magnitudes and no evidence ofmicrovariability at scales of hours or smaller was found. Daiet al. (2001) observed this source and reported a possiblemicro-variation of 0.14 magnitude in the V band within 26minutes, but the data quality was poor. Xie et al. (2002)monitored the source in 2000 and found no significant in-traday variation (maximum flickering amplitudes are ∆V ∼ ∼ ∼ ± Observations of these two HBLs were performed using fiveoptical telescopes, one in India, one in Greece and three inBulgaria. All of these telescopes are equipped with CCDdetectors and Johnson UBV and Cousins RI filters. Detailsof the telescopes, detectors and other parameters related tothe observations are given in Table 1. A complete log ofobservations of these two HBLs from those five telescopesare given in Tables 2 and 3.We carried out optical photometric observations duringthe period July 2009 to November 2010. The raw photomet-ric data were processed by standard methods that we nowbriefly describe. For image processing or pre-processing, wegenerated a master bias frame for each observing night bytaking the median of all bias frames. The master bias framefor the night was subtracted from all flat and source im-age frames taken on that night. Then the master flat ineach pass-band was generated by median combine of allflat frames in each pass-band. Next, the normalized masterflat for each pass-band was generated. Each source imageframe was flat-fielded by dividing by the normalized masterflat in the respective band to remove pixel-to-pixel inho-mogeneities. Finally, cosmic ray removal was done from allsource image frames. Data pre-processing used the standardroutines in Image Reduction and Analysis Facility (IRAF)and ESO-MIDAS softwares. IRAF is distributed by the National Optical Astronomy Obser-vatories, which are operated by the Association of Universities forResearch in Astronomy, Inc., under cooperative agreement withthe National Science Foundation. ESO-MIDAS is the acronym for the European Southern Obser-
We processed the data using the Dominion Astronom-ical Observatory Photometry (DAOPHOT II) software toperform the circular concentric aperture photometric tech-nique (Stetson 1987, 1992). For each night we carried outaperture photometry with four different aperture radii, i.e.,1 × FWHM, 2 × FWHM, 3 × FWHM and 4 × FWHM. On com-paring the photometric results we found that aperture radiiof 2 × FWHM almost always provided the best S/N ratio, sowe adopted that aperture for our final results.For these two HBLs, we observed three or more localstandard stars on the same fields. The magnitudes of thesestandard stars are given in Table 4. We employed two stan-dard stars from each blazar field with magnitudes similarto those of the target so as to not induce any significant er-rors due to differences in photon statistics. We used standardstars 4 and 6 for the blazars 1ES 1959+650 and stars C2 andC3 for the HBL 1ES 2344+514 and show their differentialinstrumental magnitudes for the IDV light curves. As thefluxes of the blazars and the standard stars were obtainedsimultaneously and so at the same air mass and with iden-tical instrumental and weather conditions, the flux ratiosare considered to be very reliable. Finally, to calibrate thephotometry of the blazars, we used the one standard star ofeach pair that had a colour closer to that of the blazar; star6 and star C2 were the calibrators for 1ES 1959+650 and1ES 2344+514, respectively. We adopted the magnitudes forthese stars listed in Table 4 but did not incorporate the er-rors in their standardized magnitudes when quoting blazarmagnitudes. Typical photometric errors for 1ES 1959+650are ∼ .
01 mag in each of the BVRI bands, while for 1ES2344+514 the errors in the B and V bands are ∼ .
02 magwhile in the R and I bands it is ∼ .
01 mag.
Romero et al. (2002) pointed out how an inappropriatechoice of the comparison stars used for differential photom-etry could result in spurious fluctuations in the differentiallight curve, and hence claims of spurious variability. To pre-vent this, we have selected non-variable or standard starsthat closely match the target’s magnitude and to be con-servative we have removed isolated apparently discrepantpoints from our analysis. We have quantified our resultsby employing up to four different statistics (e.g., de Diego2010). In our analysis, we have used standard stars 4 and6 as starA and starB for 1ES 1959+650 and standard starsC2 and C3 as starA and starB for 1ES 2344+514.
The variability detection parameter, C was introduced byRomero et al. (1999), and is defined as the average of C C C σ ( BL − starA ) σ ( starA − starB ) & C σ ( BL − starB ) σ ( starA − starB ) . (1) vatory Munich Image Data Analysis System which is developedand maintained by European Southern Observatoryc (cid:13) , 1– ?? Gaur et al.
Here ( BL − starA ) and ( BL − starB ) are the differentialinstrumental magnitudes of the blazar and standard starA and the blazar and standard star B, respectively, while σ ( BL − starA ), σ ( BL − starB ) and σ ( starA − starB ) are theobservational scatters of the differential instrumental mag-nitudes of the blazar and star A, the blazar and star B andstarA and star B, respectively. If C > . > C -test isnot a true statistic as it is not appropriately distributed andthis criterion is usually too conservative (de Diego 2010). We test our variability results using the standard F -test,which is a properly distributed statistic (de Diego 2010).Given two sample variances such as s Q for the blazar in-strumental light curve measurements and s ∗ for that of thestandard star, then F = s Q s ∗ . (2)The number of degrees of freedom for each sample, ν Q and ν ∗ will be the same and equal to the number of mea-surements N minus 1 ( ν = N − F value is thencompared with the F ( α ) ν Q ,ν ∗ critical value, where α is the sig-nificance level set for the test. The smaller the α value, themore improbable the result is produced by chance. If F islarger than the critical value, the null hypothesis (no vari-ability) is discarded. We have performed F -test at two sig-nificance levels (0.1% and 1%) which correspond to 3 σ and2.6 σ detections, respectively. χ -test We also performed a χ -test on the data. Given a numberof observations of a source over a given period of time, the χ statistic is expressed by χ = N X i =1 ( V i − V ) σ i , (3)where, V is the mean magnitude, and the i th observationyields a magnitude V i with a corresponding standard error σ i . Here σ i is the expected error, the error from consideringphoton noise from the source and sky, along with the CCDread-out and all possible non-systematic sources of error.Since such errors are often calculated from the usually un-derestimated values yielded by the IRAF reduction package,an error rescaling is necessary. Usually, theoretical errors aresmaller than the real errors by a factor of typically 1.3-1.75(e.g., Gopal Krishna et al. 2003, Bachev et al. 2005, Guptaet al. 2008). Our analysis of the current data indicates thisvalue is, on average, ∼ χ α,ν obtained from the χ probability function, where α is again the significance leveland ν = N − χ > χ α,ν the test indicates a larger than expected scattering of the datapoints, and hence evidence of variability. ANOVA tests are used to compare the means of a num-ber of samples. de Diego et al. (1998) used the one-wayANOVA test to investigate the variability in the light curvesof quasars. This method consists of measuring k groupsof n j =5; in short cadence observations such as ours, thesegroups should be ideally separated by 20 −
30 minutes. Vari-ance of the means of each group of five observations and themean for the dispersion within the groups are computed.Then, the ratio between these variances is calculated andmultiplied by the number of observations in each group. Thenumber obtained behaves as the F -statistic and we comparethe F -value with the critical values F ( α ) ν ,ν . In this test, if N is the number of observations and k is the number of groups( k = N/n ), the number of degrees of freedom are ν = k − ν = N − k for the errors; therefore, ν + ν = N −
1, corresponds to the degrees of freedom ofthe original data set. For a certain significance level α , if F exceeds the critical value, the null hypothesis will be re-jected. We have used the inbuilt ANOVA code available inR. Heidt & Wagner (1996) introduced the variability ampli-tude, defined as A = 100 < A > × p ( A max − A min ) − σ (%) , (4)where A max and A min are the maximum and minimumfluxes in the calibrated LCs of the blaza, < A > is theirmean, and the average measurement error of the blazar LCis σ . We use this approach to quantify any variability. We observed the source 1ES 1959+650 on 21 nights (simul-taneously in optical bands, B and R) and 3 nights in theR band alone from July 2009 to August 2010 for microvari-ability. We have performed the four tests discussed aboveon these data. The C , F , χ and ANOVA parameters neverexceed the 99% confidence level (Table 5) for B, R or (ofcourse) (B-R). So this source was stable with respect to IDVduring our observations. The differential light curves of thesource are given in figure 1 and 2 for these 24 nights. We observed the source 1ES 2344+514 on 14 nights (simul-taneously in the B and R bands) and 5 nights in the Rband from August 2009 to July 2010 for microvariability.We have tested for IDV using all four tests but no positiveresults were ever obtained in the B, R or (B-R) bands asthe C , F , χ and ANOVA results never showed significancelevels above 99% considering both stars (Table 6); althoughseveral F -test values exceed 0.99 significance with respect to c (cid:13) , 1– ?? ptical IDV of HBLs Star 1 they never do so with respect to Star 2. Hence, thissource was also very stable during each individual night ofthe observation period. The differential light curves of 1ES2344+514 are given in figure 3 and 4 for these 14 nights.
We quantify our results of short-term variability using twostatistics: C– and F-tests. Since, in ANOVA test, we usedto compare the means of a number of samples or groups (of5 observations each), which should be ideally seperated by20-30 minutes which is not possible for our short-term ob-servations. Also, χ test compares the variation of an objectwith its photometric errors. And the errors from DAOPHOTare usually underestimated. Because of which, we need to in-troduce a factor by which these errors are multiplied. Thisfactor comes from a comparison between the variation andthe errors of two intrinsically nonvariable stars. As men-tioned below, for our source 1ES 2344+514, two or perhapsmore than two out of three standard stars present in thefield are slowly varying. So, in this case, the factor is under-estimated and its not just of typically 1.3-1.75. So, χ testis not applicable here. Hence, we are not using ANOVA and χ test in short-term variability analysis. We have examined the source 1ES 1959+650 for short-termvariability during 44 nights (in B, V, R and I bands) betweenJuly 2009 and November 2010 (Figure 5).
B pass-band:
The short-term light curve of 1ES 1959+650 in the Bband is displayed in the upper panel of Figure 5. The max-imum variation noticed in the light curve of the source is1.23 magnitudes (between its brightest level at 14.99 magon JD 2455036.40207 and the faintest level at 16.22 mag onJD 2455330.40554). The values of C − and F -tests supportthe existence of short-term variation of the source in theB-band observations. We calculated the short-term variabil-ity amplitude using Equation (3) and found that source hasvaried ≈
102 per cent.
V pass-band:
The V-band short-term light curve of 1ES 1959+650is shown in the second panel of Figure 5. The maximumvariation seen is 1.06 mag (between its brightest level at14.429 on JD 2455036.40811 and the faintest level at 15.490on JD 2455330.40301). The C − and F -tests again supportthe existence of short-term variation of the source in theV-band observations. We calculated that this source has avariability amplitude of ≈
90 per cent.
R pass band:
The third panel of Figure 5 gives the R-band short-termlight curve of 1ES 1959+650. The maximum variation no-ticed in the source is 0.94 mag (between its brightest level at14.015 on JD 2455041.32782 and the faintest level at 14.959on JD 2455326.41514). The values of the C − and F -test cal-culations also support the existence of STV in our source.The amplitude of variability in the source is ≈
82 per cent.
I pass-band:
The short-term light curve of 1ES 1959+650 in the Iband is also displayed, in the fourth panel of Figure 5. The maximum variation noticed in the source is 0.85 mag (be-tween its brightest level at 13.49 mag on JD 2455036.40052and the faintest level at 14.34 mag on JD 2455330.39305).Again, the C − and F -tests both support the existence ofshort-term variations of the source in these I-band observa-tions. The amplitude of variability in the source is ≈
75 percent. (V-R) Colour:
The (V-R) colour index of 1ES 1959+650 in the (V-R)over this period comprises the bottom panel of Figure 5. Themaximum variation we saw in the source is 0.39 (betweenits colour index of 0.29 on JD 2455395.38234 and 0.68 at JD2455161.26157). Both the C − and F − tests show significant(V-R) colour variations in our observations. The amplitudeof variability in the source’s colour is ≈
36 per cent.
We have examined the source 1ES 2344+514 on 39 nights (inB,V,R and I bands) ranging from August 2009 to November2010 for short-term variability (shown in figure 6). Becauseof the variation in the standard stars present in the field ofthis blazar, we noticed variations in 1ES 2344+514 in eachband. Perhaps two out of three standard stars have slowvariations which are propagating in the source. Therefore,despite the nominally large variation in the blazar calibratedmagnitude, neither the C − nor the F − test support the exis-tence of short-term variation of the source in 1ES 2344+514.Also, the similar behaviour of standard stars were noticedin Tuorla observatory observations The maximum variation noticed in the short-term dif-ferential B-band light curve of 1ES 2344+514 is 0.80 mag(between its brightest level at 2.02 mag on JD 2455065.48144and the faintest level at 1.22 mag on JD 2455449.57756) andthe variation in the differential instrumental magnitudes ofstandard stars is 0.53. Because the stellar differential magni-tudes varied so much, none of them could be truly consideredto be of fixed brightness.The maximum variation we found in the short-termlight curve of this source in V-band is 0.73 mag (betweenits brightest level at 15.47 mag on JD 2455449.57230 andthe faintest level at 16.20 mag on JD 2455065.47806). Thevariation in the differential instrumental magnitudes of stan-dard stars is 0.45.The R-band maximum variation we found was 0.65mag (between its brightest level at 14.72 mag on JD2455449.56190 and the faintest level at 15.37 mag on JD2455070.41122). The variation in the differential instrumen-tal magnitudes of standard stars is 0.44.The maximum variation noticed in the light curve ofthe source in I-band is 0.62 mag (between its brightest levelat 13.92 mag on JD 2455449.55730 and the faintest level at14.52 mag on JD 2455063.50316). The variation in the dif-ferential instrumental magnitudes of standard stars is 0.55.Finally, the maximum difference in the (V-R) Colourfor 1ES 2344+514 is 0.22 mag (between its colour range0.62 mag on JD 2455070.41047 and 0.84 mag at JD2455065.47656). The variation in the differential instrumen-tal magnitudes of standard stars is 0.16. http://users.utu.fi/kani/1m/c (cid:13) , 1– ?? Gaur et al.
We have presented our results of quasi-simultaneous andrelatively dense temporal observations of two HBLs, 1ES1959+650 and 1ES 2344+514, in two optical bands (B andR) during 2009–2010 . These are the first extensive obser-vations of this type of these two blazars and thus adds tothe rather small amount of data on HBL microvariabilityand STV. We have observed the sources 1ES 1959+650 and1ES 2344+514 on 24 and 19 nights respectively, but no sig-nificant intra-day variability was observed during any nightfor those two targets. In addition, we can report the resultsof short-term variability studies of these sources. The blazar1ES 1959+650 has shown genuine short-term variability andit can be characterized by a preliminary fading trend of ∼ ∼ c given by B c = [4 πnm e c ( γ − / γ − , (5)where n is the local electron density, m e is the electron restmass, and γ is the central flows bulk Lorentz factor. So, if B > B c this class of instabilities will be suppressed. How-ever, if B < B c the Kelvin-Helmholtz instabilities can pro-duce significant changes in the jet morphology, and thesefeatures could be responsible for rapid variability when therelativistic shock waves interact with them. So if HBLs typi-cally possess stronger magnetic fields for given electron den-sities or lower electron densities for given field strengths,than we suggest that such a reduction in Kelvin-Helmholtzinstabilities would reduce the incidence of microvariabilityin the optical light curves.Another possible explanation for the difference betweenthe optical variability behaviour of the LBL and HBL classesof blazars can dispense with the hypothesis of suppressedKelvin-Helmholtz instabilities. We recall that the opticalband is near or above the peak of the first hump in theSED for LBLs but below that peak for HBLs. Thereforechanges in the efficiency of acceleration of, and/or in therates at which energy is radiated by, the highest energyelectrons available for synchrotron emission would have alarge and rapid effect on the emitted flux in those bands forLBLs. However, those variations would have a more mod-est and more retarded effect on optical variability in HBLs,since their optical emission is below the peak of the syn-chrotron emission. Such variations in acceleration efficiencycould arise from changes in the local number density of themost energetic electrons or the strengths of the maximumlocalized magnetic fields. While magneto-hydrodynamic in-stabilities might be the origin of such variations, so mightthe presence of turbulence in the vicinity of the shock (e.g.,Marscher, Gear & Travis 1992). If this is the case, then X-ray variability should be more pronounced for HBLs thanfor LBLs, in that the peak of the synchrotron emission liesnear the X-ray band for the former class. ACKNOWLEDGMENTS
This research was partially supported by Scientific ResearchFund of the Bulgarian Ministry of Education and Sciences(BIn - 13/09 and DO 02-85) and by Indo − Bulgaria bilateralscientific exchange project INT/Bulgaria/B − c (cid:13) , 1– ?? ptical IDV of HBLs Research and Technology – Hellas, and the Max-Planck-Institut f¨ur Extraterrestrische Physik.
REFERENCES
Abdo A. A., et al., 2010, ApJ, 715, 429Aharonian F., et al., 2003, A&A, 406, L9Albert J., et al., 2006, ApJ, 639, 761Albert J., et al., 2007, ApJ, 662, 892Bachev R., Strigachev A., Semkov E., 2005, MNRAS, 358,774Beckmann V., Wolter A., Celotti A., Costamante L., Ghis-ellini G., Maccacaro T., Tagliaferri G., 2002, A&A, 383,410B¨ottcher M., 2005, ApJ, 621, 176Camenzind M., Krockenberger M., 1992, A&A, 255, 59Carini M. T., Miller H. R., 1992, ApJ, 385, 146Catanese M., et al., 1998, ApJ, 501, 616Chakrabarti S. K., Wiita P. J., 1993, ApJ, 411, 602Dai B. Z., Xie G. Z., Li K. H., Zhou S. B., Liu W. W.,Jiang Z. J., 2001, AJ, 122, 2901Daniel M. K., et al., 2005, ApJ, 621, 181de Diego J. A., Dultzin-Hacyan D., Ram´ırez A., BenitezE., 1998, ApJ, 501, 69de Diego J. A., 2010, AJ, 139, 1269Doroshenko V. T., Sergeev S. G., Efimov Y. S., NazarovS., Pronik V. I., Sergeeva E. A., Sivtsov G. A., 2007, Ap,50, 40Elvis M., Plummer D., Schachter J., Fabbiano G., 1992,ApJS, 80, 257Falomo R., Kotilainen J. K., Treves A., 2002, ApJ, 569,L35Fan J.-H., Kurtanidze O. M., Nikolashvili M. G., GuptaA. C., Zhang J.-S., Yuan Y.-H., 2004, ChJAA, 4, 133Fiorucci M., Tosti G., Rizzi N., 1998, PASP, 110, 105Fossati G., Maraschi L., Celotti A., Comastri A., GhiselliniG., 1998, MNRAS, 299, 433Ghisellini G., et al., 1997, A&A, 327, 61Ghisellini G., Celotti A., Fossati G., Maraschi L., ComastriA., 1998, MNRAS, 301, 451Giommi P., Ansari S. G., Micol A., 1995, A&AS, 109, 267Giommi P., Padovani P., Perlman E., 2000, MNRAS, 317,743Gopal-Krishna, Wiita P. J., 1992, A&A, 259, 109Gopal-Krishna, Stalin C. S., Sagar R., Wiita P. J., 2003,ApJ, 586, L25Grube J., 2008, AIPC, 1085, 585Gupta A. C., Banerjee D. P. K., Ashok N. M., Joshi U. C.,2004, A&A, 422, 505Gupta A. C., et al., 2008, AJ, 136, 2359Heidt J., Wagner S. J., 1996, A&A, 305, 42Heidt J., Wagner S. J., 1998, A&A, 329, 853Heidt J., Nilsson K., Sillanp¨a¨a A., Takalo L. O., PursimoT., 1999, A&A, 341, 683Holder J., et al., 2003, ApJ, 583, L9Horns D., Konopelko A., 2002, IAUC, 7907, 2Hughes P. A., Aller H. D., Aller M. F., 1991, ApJ, 374, 57Jang M., Miller H. R., 1997, AJ, 114, 565Jannuzi B. T., Green R. F., French H., 1993, ApJ, 404, 100Jannuzi B. T., Smith P. S., Elston R., 1994, ApJ, 428, 130Krawczynski H., 2004, NewAR, 48, 367 Krawczynski H., et al., 2004, ApJ, 601, 151Kurtanidze O. M., Richter G. M., Nikolashvili M. G., 1999,in Blazar Monitoring towards the Third Millenniumn, eds.C. M. Raiteri, M. Villata and L.O. Takalo, p. 29Kurtanidze O. M., Nikolashvili M. G., 2002, in Blazar As-trophysics with BeppoSAX and Other Observatories, eds.P. Giommi, E. Massaro and G. Palumbo, p. 197Ma L., Xie G. Z., Yi T. F., Zhou S. B., Li K. H., ZhangX., Dai H., 2010, Ap&SS, 327, 35Mangalam A. V., Wiita P. J., 1993, ApJ, 406, 420Marcha M., 1994, Ph.D. Dissertation, University of Manch-esterMarscher A. P., Gear W. K., 1985, ApJ, 298, 114Marscher A. P., Gear W. K., Travis J. P., 1992, in Vari-ability of Blazars, eds. E. Valtaoja and M. Valtonen, Cam-bridge University Press, p. 85Marscher A. P., 1996, ASPC, 110, 248Miller H. R., Ferrara E. C., Daya A. B., Wilson J. W., FriedR. E., Noble J. C., Jang M., 1999, in Blazar Monitoringtowards the Third Millenniumn, eds. C.M. Raiteri, M.Villata and L.O. Takalo, p. 20Padovani P., Giommi P., 1995, ApJ, 444, 567Perlman E. S., et al., 1996, ApJS, 104, 251Qian S. J., Quirrenbach A., Witzel A., Krichbaum T. P.,Hummel C. A., Zensus J. A., 1991, A&A, 241, 15Rani B., et al., 2010, MNRAS, 404, 1992Rani B., Gupta A. C., Joshi U. C., Ganesh S., Wiita P. J.,2011, MNRAS, 324Romero G. E., 1995, Ap&SS, 234, 49Romero G. E., Cellone S. A., Combi J. A., 1999, A&AS,135, 477Sambruna R. M., Maraschi L., Urry C. M., 1996, ApJ, 463,444Schachter J. F., et al., 1993, ApJ, 412, 541Sikora M., Madejski G., 2001, AIPC, 558, 275Sol H., Pelletier G., Asseo E., 1989, MNRAS, 237, 411Stalin C. S., Gopal-Krishna, Sagar R., Wiita P. J., 2004,MNRAS, 350, 175Stetson P. B., 1987, PASP, 99, 191Stetson P. B., 1992, JRASC, 86, 71Stocke J. T., Liebert J., Schmidt G., Gioia I. M., MaccacaroT., Schild R. E., Maccagni D., Arp H. C., 1985, ApJ, 298,619Stocke J. T., Morris S. L., Gioia I. M., Maccacaro T., SchildR. E., Wolter A., 1989, LNP, 334, 242Tagliaferri G., et al., 2008, ApJ, 679, 1029Tonello N., Kranich D., HEGRA collaboration, 2003,ICRC, 5, 2615Urry C. M., Padovani P., 1995, PASP, 107, 803V´eron-Cetty M.-P., V´eron P., 2006, A&A, 455, 773Villata M., Raiteri C. M., Lanteri L., Sobrito G., CavalloneM., 1998, A&AS, 130, 305Villata M., Raiteri C. M., Popescu M. D., Sobrito G., DeFrancesco G., Lanteri L., Ostorero L., 2000, A&AS, 144,481Villata M., et al., 2007, A&A, 464, L5Wagner S. J., Witzel A., 1995, ARA&A, 33, 163Xie G. Z., Zhou S. B., Dai B. Z., Liang E. W., Li K. H.,Bai J. M., Xing S. Y., Liu W. W., 2002, MNRAS, 329,689 c (cid:13) , 1– ?? Gaur et al.
Table 1.
Details of telescopes and instrumentsSite: ARIES Nainital NAO Rozhen NAO Rozhen AO Belogradchik Skinakas Observatory, CreteTelescope: 1.04-m RC Cassegrain 2-m Ritchey-Chr´etien 50/70-cm Schmidt 60-cm Cassegrain 1.3-m Modified RCCCD model: Wright 2K CCD PI VersArray:1300B FLI PL160803 FLI PL09000 Andor DX436-BV-9CQChip size: 2048 × × × × × × µ m 20 × µ m 9 × µ m 12 × µ m 13 . × . µ mScale: 0.37 ′′ /pixel 0.258 ′′ /pixel 1.079 ′′ /pixel 0.330 ′′ /pixel a ′′ /pixelField: 13 ′ × ′ . ′ × . ′ . ′ × . ′ . ′ × . ′ . ′ × . ′ Gain: 10 e − /ADU 1.0 e − /ADU 1.0 e − /ADU 1.0 e − /ADU 2.687 e − /ADURead Out Noise: 5.3 e − rms 2.0 e − rms 9.0 e − rms 8.5 e − rms 8.14 e − rmsBinning used: 2 × × × × × ′′ to 2.8 ′′ ′′ to 3.5 ′′ ′′ to 4 ′′ ′′ to 3.5 ′′ ′′ to 2 ′′ a With a binning factor of 1 × (cid:13) , 1– ?? ptical IDV of HBLs Table 2.
Observation log of optical photometric observations of 1ES 1959+650Date of Telescope Data Pointsobservation Filters(yyyy mm dd) (B,V,R,I)2009 07 23 B 1,1,1,12009 07 26 D 56,2,56,22009 07 27 D 58,1,58,12009 07 28 D 57,1,57,12009 07 29 D 56,2,55,22009 07 30 D 76,1,37,12009 08 03 E 42,2,43,22009 08 05 E 102,2,102,22009 08 17 D 57,1,56,12009 08 20 D 50,2,49,22009 08 21 D 62,2,61,22009 08 27 E 90,2,90,22009 09 22 C 90,2,90,22009 10 03 E 65,2,67,22009 10 09 A 1,1,81,12009 10 10 A 1,1,1,12009 11 13 D 2,2,2,22009 11 14 D 2,2,2,22009 11 16 C 2,2,2,22009 11 17 C 2,2,2,22009 11 19 C 2,2,2,22009 11 20 C 2,2,2,22009 11 21 C 2,2,2,22009 11 25 B 2,2,2,22010 05 09 C 3,2,96,42010 05 13 C 2,2,2,22010 06 08 C 2,2,2,22010 06 10 C 2,2,2,22010 06 12 C 2,2,2,22010 06 19 E 61,2,61,22010 07 15 D 20,2.19,22010 07 16 D 32,2,32,22010 07 17 E 78,2,78,22010 07 17 C 2,2,80,22010 07 19 E 62,2,62,22010 07 21 E 72,2,72,22010 08 05 D 21,2,21,22010 08 06 C 2,2,2,22010 08 07 C 2,2,2,22010 08 08 C 2,2,2,22010 08 09 C 2,2,2,22010 08 10 D 15,1,15,12010 11 04 C 2,2,2,22010 11 05 C 2,2,2,2c (cid:13) , 1– ?? Gaur et al.
Table 3.
Observation log of optical photometric observations of 1ES 2344+514Date of Telescope Data Pointsobservations Filters(yyyy mm dd) (B,V,R,I)2009 07 23 B 1,1,1,12009 08 04 E 52,2,52,22009 08 18 D 45,8,41,82009 08 19 D 42,6,40,62009 08 21 C 1,1,1,12009 08 22 C 36,1,36,12009 08 25 C 44,2,44,22009 08 26 E 46,3,48,32009 08 28 C 45,2,46,22009 08 29 C 54,2,54,22009 09 18 E 42,2,42,22009 09 21 C 29,1,32,22009 10 03 E 32,2,32,22009 10 10 A 1,1,65,12009 11 13 D 2,2,2,22009 11 14 D 2,2,2,22009 11 16 C 2,2,2,22009 11 17 C 2,2,2,22009 11 19 C 2,2,2,22009 11 20 C 2,2,2,22009 11 21 C 2,2,2,22009 11 25 B 2,2,2,22010 01 10 A 1,1,49,12010 01 11 A 1,1,50,12010 01 20 A 1,1,30,12010 06 09 C 2,2,2,22010 06 11 C 2,2,2,22010 06 13 C 2,2,2,22010 07 18 C 2,2,114,22010 07 18 E 62,4,62,42010 07 20 E 86,2,86,22010 07 22 E 17,2,17,22010 08 06 C 2,2,2,22010 09 08 C 2,2,2,22010 09 09 C 2,2,2,22010 10 31 C 2,2,2,22010 11 01 C 2,2,2,22010 11 04 C 2,2,2,22010 11 05 C 2,2,2,22010 11 06 C 2,2,2,2A : 1.04 meter Sampuranand Telescope, ARIES, Nainital, IndiaB : 2-m Ritchey-Chretien Telescope at National Astronomical Observatory, Rozhen, BulgariaC : 50/70-cm Schmidt Telescope at National Astronomical Observatory, Rozhen, BulgariaD : 60-cm Cassegrain Telescope at Astronomical Observatory Belogradchik, BulgariaE : 1.3-m Skinakas Observatory, Crete, Greece c (cid:13) , 1– ?? ptical IDV of HBLs Table 4.
Standard stars in the blazar fieldsSource Standard B magnitude a V magnitude R magnitude I magnitude a ReferencesName star (error) (error) (error) (error)1ES 1959+650 1 13.37(0.02) 12.67(0.04) 12.29(0.02) 11.92(0.01) 1, 32 13.45(0.02) 12.86(0.02) 12.53(0.02) 12.22(0.01) 1, 33 14.93(0.02) 13.18(0.02) 12.27(0.02) 11.37(0.01) 1, 34 15.28(0.02) 14.53(0.03) 14.08(0.03) 13.62(0.01) 1, 35 15.60(0.03) 14.54(0.03) 14.00(0.02) 13.36(0.02) 1, 36 15.97(0.02) 15.20(0.03) 14.78(0.03) 14.37(0.01) 1, 37 16.01(0.02) 15.24(0.03) 14.79(0.03) 14.37(0.01) 1, 31ES 2344+514 C1 - 12.61(0.04) 12.25(0.04) 11.90(0.04) 2C2 - 14.62(0.06) 14.20(0.05) 13.84(0.04) 2C3 - 15.89(0.08) 15.40(0.08) 14.89(0.08) 21. Villata et al. 1998; 2. Fiorucci, M., Tosti, G., Rizzi N., et al. 1998. a
3. Doroshenko et al. 2007.c (cid:13) , 1– ?? Gaur et al.
Table 5.
Results of intra-day variability observations of 1ES 1959+650Date Band N C-test F-test χ test ANOVA Variable C , C F , F , F c (0 . , F c (0 . χ , χ , χ . , χ . F , F , F c (0 . , F c (0 . (cid:13) , 1– ?? ptical IDV of HBLs Table 6.
Results of intra-day variability observations of 1ES 2344+514Date Band N C-test F-test χ test ANOVA Variable C , C F , F , F c (0 . , F c (0 . χ , χ , χ . , χ . F , F , F c (0 . , F c (0 . (cid:13) , 1– ?? Gaur et al.
Table 7.
Results of short-term variability observationsSource Name Band N C-Test F Variable A (%) C , C F , F , F c (0 . , F c (0 . (cid:13) , 1– ?? ptical IDV of HBLs Figure 1.
The B (middle panel), R (lower panel) and (B-R)(upper panel) light curves of 1ES 1959+650. The x -axis is JD (2455000+)and the y -axis is the differential instrumental magnitude. Open circles (also in red colour) give the DLC of Blazar-Star1, filled circles(in blue colour), the DLC of Blazar-Star2 while star symbols (in black colour) represents DLC of comparison stars (st1-st2). Dates andtelescopes are given on top of each plot. The magnitudes of each band are adjusted with an arbitrary offset (for clarity) in each panel offigure.c (cid:13) , 1– ?? Gaur et al.
Figure 2.
As in Fig. 1 for additional nights for 1ES 1959+650. c (cid:13) , 1– ?? ptical IDV of HBLs Figure 3.
As in Fig. 1 for 1ES 2344+514.c (cid:13) , 1– ?? Gaur et al.
Figure 4.
As in Fig. 1 for 1 ES 2344+514. c (cid:13) , 1– ?? ptical IDV of HBLs Figure 5.
Short − term variability light curve of 1ES 1959+650. Starred, solid circle, open circle, triangle and square symbols representdata from the telescopes A, B, C, D and E respectively.c (cid:13) , 1– ?? Gaur et al.
Figure 6.
Short − term variability light curve of 1ES 2344+514. Starred, solid circle, open circle, triangle and square symbols representdata from the telescopes A, B, C, D and E, respectively. Y-axis is magnitude and X-axis is JD (2455000+) in each plot. In the B-banddata set, the upper LC is (blazar − Star1) and the lower one is the differential magnitude of (Star1 − Star2). In the remaining panels, theupper curves are the calibrated blazar magnitudes and lower ones are differential magnitudes of (Star1 − Star2) in V, R , I and (V-R).c (cid:13) , 1–, 1–