The Australia Telescope 20 GHz (AT20G) Survey: The Bright Source Sample
Marcella Massardi, Ronald D. Ekers, Tara Murphy, Roberto Ricci, Elaine M. Sadler, Sarah Burke, Gianfranco De Zotti, Philip G. Edwards, Paul J. Hancock, Carole A. Jackson, Michael J. Kesteven, Elizabeth Mahony, Christopher J. Phillips, Lister Staveley--Smith, Ravi Subrahmanyan, Mark A. Walker, Warwick E. Wilson
aa r X i v : . [ a s t r o - ph ] F e b Mon. Not. R. Astron. Soc. , 1–18 (2002) Printed 27 October 2018 (MN L A TEX style file v2.2)
The Australia Telescope 20 GHz (AT20G) Survey:The Bright Source Sample
Marcella Massardi , ⋆ , Ronald D. Ekers , Tara Murphy , , Roberto Ricci , ElaineM. Sadler , Sarah Burke , Gianfranco De Zotti , , Philip G. Edwards , Paul J.Hancock , Carole A. Jackson , Michael J. Kesteven , Elizabeth Mahony , Christo-pher J. Phillips , Lister Staveley–Smith , Ravi Subrahmanyan , Mark A. Walker ,and Warwick E. Wilson SISSA/ISAS, Via Beirut 2–4, I-34014 Trieste, Italy Australia Telescope National Facility, CSIRO, P.O. Box 76, Epping, NSW 1710, Australia School of Physics, University of Sydney, NSW 2006, Australia School of Information Technologies, University of Sydney, NSW 2006, Australia Department of Physics and Astronomy, University of Calgary, 2500 University Drive NW Calgary, AB, Canada Swinburne University of Technology, P.O. Box 218, Hawthorn, Vic 3122, Australia INAF, Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, I-35122 Padova, Italy School of Physics,University of Western Australia, 35 Stirling Highway Crawley, WA 6009, Australia Raman Research Institute, Sadashivanagar, Bangalore 560080, India Manly Astrophysics Workshop Pty Ltd, 3/22 Cliff St., Manly 2095, Australia
ABSTRACT
The Australia Telescope 20 GHz (AT20G) Survey is a blind survey of the whole South-ern sky at 20 GHz (with follow-up observations at 4.8 and 8.6 GHz) carried out withthe Australia Telescope Compact Array (ATCA) from 2004 to 2007.The Bright Source Sample (BSS) is a complete flux-limited sub-sample of theAT20G Survey catalogue comprising 320 extragalactic ( | b | > . ◦ ) radio sources southof δ = − ◦ with S > .
50 Jy. Of these, 218 have near simultaneous observationsat 8 and 5 GHz.In this paper we present an analysis of radio spectral properties in total inten-sity and polarisation, size, optical identifications and redshift distribution of the BSSsources. The analysis of the spectral behaviour shows spectral curvature in mostsources with spectral steepening that increases at higher frequencies (the median spec-tral index α , assuming S ∝ ν α , decreases from α . . = 0 .
11 between 4.8 and 8.6 GHzto α . = − .
16 between 8.6 and 20 GHz), even if the sample is dominated by flatspectra sources (85 per cent of the sample has α . > − . Key words: surveys – radio continuum: general – galaxies: active – cosmic microwavebackground .
Our knowledge of the high frequency radio-source popula-tion is poor. Important advances were made recently by the15 GHz surveys with the Ryle telescope (Taylor et al. 2001; ⋆ E-mail: [email protected]
Waldram et al. 2003) covering 520 deg to a flux density limitof 25 mJy and going down to 10 mJy in small areas. TheWilkinson Microwave Anisotropy Probe (WMAP) satellitehas surveyed the whole sky at 23, 33, 41, 61 and 94 GHz tocompleteness limits of > ∼ c (cid:13) Massardi et al. formation in the flux densities range between 200 mJy and1 Jy.High frequency surveys are very time-consuming. Fortelescopes with diffraction limited fields of view the numberof pointings necessary to cover a given area scales as ν .For a given receiver noise, the time per pointing to reachthe flux density level S scales as S − so that, for a typicaloptically thin synchrotron spectrum ( S ∝ ν − . ), the surveytime scales as ν +3 . : a 20 GHz survey takes more than 110times longer than a 5 GHz survey with the same aperturecovering the same area of sky to the same flux density level.However, Cosmic Microwave Background (CMB) stud-ies, boosted by the on-going NASA WMAP mission and bythe forthcoming ESA Planck mission, require an accuratecharacterization of the high frequency properties of fore-ground radio-sources both in total intensity and in polarisa-tion. Radio sources are the dominant contaminant of small-scale CMB anisotropies at mm wavelengths: their Poissoncontribution to temperature fluctuations is inversely propor-tional to the angular scale, i.e. linearly proportional to themultipole number l , while the power spectra of the CMB andof Galactic emissions decline at large l . As a result, Poissonfluctuations dominate for l > ∼ > ∼
200 mJy at all frequencies) over a largearea, complemented with a good characterization of sourceproperties, can be used to correct the contaminating effect.It may be necessary for the high frequency surveys togo even fainter than the rms in the CMB observations. Thisis because fluctuations in numbers of weak sources (particu-larly if there is clustering) could affect power spectrum esti-mation: it depends on the source counts at weak flux densitylevels and the contribution to the angular power spectrumfrom sources at different flux density intervals.Furthermore, forthcoming telescopes in the South-ern hemisphere, like the Atacama Large Millimeter Array(ALMA), that will operate at frequencies above 90 GHz, re-quire suitable calibrators which can be readily selected usinglarge area high frequency surveys (Sadler et al. 2006).Optical surveys for transients, which would be severelycontaminated by Galactic novae, could use high frequencyradio surveys as templates to identify potential transientsassociated with AGNs (Rau et al. 2007).A Pilot Survey (Ricci et al. 2004; Sadler et al. 2006)at 18.5 GHz was carried out in 2002 and 2003 with theAustralia Telescope Compact Array (ATCA ). It detected173 objects in the declination range − ◦ to − ◦ down to100 mJy.The Pilot project characterised the high-frequency ra-dio source population and allowed us to optimise the obser-vational techniques for the full Australia Telescope 20 GHz(AT20G) Survey. The full survey will cover the whole South-ern sky to a flux density limit of ≃
50 mJy. It began in 2004and will be completed in 2007. To date we have completedthe survey from the South Pole to declination δ = − ◦ .More than 4400 sources were detected, down to a flux den-sity of 50 mJy. We expect another 1500 sources in the dec-lination range − ◦ < δ < ◦ for which the analysis is stillon-going. This paper presents the analysis of the brightest( S > .
50 Jy) extragalactic ( | b | > . ◦ ) sources in theAT20G Survey based on the observations in the declinationrange − ◦ < δ < − ◦ surveyed between 2004 and 2007.In § § §
4, in § § The first phase of our observations is to make a set of blindscans. More details on the survey mode observations, themapmaking, the source detection in maps and the complete-ness and reliability of the survey will be provided in forth-coming papers on the full AT20G sample. Here we presentonly a general description.We have exploited the ATCA fast scanning capabili-ties (15 degrees min − in declination at the meridian) andthe 8 GHz bandwidth of the wideband analogue correla-tor originally developed as part of the collaboration for theTaiwanese CMB experiment AMiBA (Lo et al. 2001) andnow applied to 3 of the six 22 m dishes of the ATCA. Thelag-correlator measures 16 visibilities as a function of differ-ential delay for each of the three antenna pairs used. Thiswideband analogue correlator has no mechanism to allow forgeometrical delay as a function of the position in the sky,so the scan has to be performed along the meridian corre-sponding to zero delay for the EW configuration used. Thereis no fringe stopping.The scanning strategy consists of sweeping sky regions10 ◦ or 15 ◦ wide in declination, using a whole Earth rotationto cover all the right ascensions in a zig-zag pattern. Eachdeclination strip requires several days to be completely cov-ered by moving the scanning path half a beam apart fromday to day. Along the scan a sample is collected every 54 ms(3 samples per beam), enough to reach a rms noise of 12 mJy.With this exceptional continuum sensitivity, along with pre-cise and high speed telescope scanning capability, we canscan large areas of the sky, despite the small ( ∼ . ≃
10 mJy.The initial calibration is based on a transit observa-tion of a known calibrator observed every 24 hours between c (cid:13) , 1–18 he AT20G Bright Source Sample Table 1.
Follow-up observations to confirm candidate sources at 20 GHz (flagged as C ), to observe them at 5 and 8 GHz ( O ) or torepeat previous bad quality observations ( R ). ( M ) refers to the observation in which we observed the very extended sources in mosaicmode. Epoch Declination
Central Array Configuration
Beamsizeref. range
Frequencies(MHz) (shortest spacing [m]) [arcsec]
Dates Reason − ◦ , − ◦ . × . − ◦ , − ◦ . × . . × .
13 04 Nov - 08 Nov 2004 O2 − ◦ , − ◦ . × . − ◦ , − ◦ . × . . × . − ◦ , − ◦ − ◦ , − ◦ − ◦ , − ◦ . × . − ◦ , − ◦ . × . . × . − ◦ , − ◦ − ◦ , − ◦ − ◦ , − ◦ . × . scans. All the sources detected in the scans that have knownflux densities and positions (about 10 for each scan) arethen used in a bootstrap process to refine the scan calibra-tion. From this we produced an initial list of positions andflux densities for candidate sources brighter than 5 σ (about50 mJy). Each of the candidate sources selected in the first phasehas been re-observed to confirm they are genuine sourcesand to get accurate positions, flux densities and polarisationinformation.Note that this procedure will exclude any fast (withinfew weeks) transient sources, if they exist. We intend tocheck for such objects in a future analysis. The follow-uphas been performed with an hybrid array configuration (i.e.,with some of the baselines on the NS direction) with the nor-mal ATCA digital correlator with two 128 MHz bands cen-tered at 18752 MHz and 21056 MHz and two polarisations.The combination of the two close bands could be consid-ered as a single 256 MHz wide band centered at 19904 MHz,which is the reference frequency for our ‘20 GHz’ observa-tions.The follow-up observations exploit the fast mosaic ca-pabilities of the ATCA to reduce the slewing time betweenpointings. In our observing strategy each mosaic point isa pointing on a candidate source. The same source has tobe observed more than once to improve the visibility planecoverage. The sources discussed in this paper have been ob-served at least twice and in some cases up to 8 times at differ-ent hour angles. Up to 500 candidates could be followed-upin a day. A set of secondary calibrator sources are regularlyobserved between blocks of candidate sources.Within a couple of weeks, we observed the confirmedsources with an East-West extended array configurationwith two 128 MHz bands centered at 4800 MHz and8640 MHz to study their radio spectral properties. Thoseare the frequencies to which we will refer in the following as‘5’ and ‘8’ GHz. In Table 1 we have summarised the arrayconfigurations used to observe sources or to replace previousbad quality data in the various sky regions. The simultaneityof observations at different frequencies is necessary to study the spectral properties of the sources, avoiding errors due tothe source variability.The primary beam FWHM is 2 .
4, 5 .
5, and 9 . We have developed a fully automated custom analysispipeline to edit, calibrate, and reduce the data for all thefollow-up observations (Fig. 1). This procedure has been de-veloped to ensure consistent data quality in the final cat-alogued data. The software was built using the scriptinglanguage Python, and the underlying data reduction wasdone with the aperture synthesis reduction package Miriad(Sault, Teuben & Wright 1995).After an initial manual inspection of the data to flagbad data, the pipeline generates the calibration solutions.Once source flux densities are calibrated, a set of processes isapplied to determine positions, peak flux densities, extend-edness, integrated flux densities, polarisation properties andto generate images. The final result is a list of confirmedsources with all the available information and images foreach epoch and for each frequency. From this list we haveselected the Bright Source Sample that we will analyze inthe following sections. c (cid:13) , 1–18 Massardi et al.
In the rest of this section we describe the details of thedata reduction.
An initial inspection of the correlator output is necessaryin order to identify interference or any problems in the dataacquisition that mean the data should be excluded from fur-ther analysis.Weather conditions can seriously affect the quality ofthe data. Attenuation of the signal by atmospheric watervapour can decrease the sensitivity of the observations, andatmospheric turbulence can produce phase fluctuations thatmay produce visibility amplitude decorrelation. Data col-lected in periods of bad weather have to be removed. Inparticular, calibrator data must be of high quality otherwiseit introduces errors in the calibration solutions that affectthe whole dataset (Thompson, Moran & Swenson 2001).A seeing monitoring system is run at the ATCA sitesimultaneously with the main array. Two 40 cm dishes on a240 m baseline monitor the differential phase variations ina geostationary satellite signal caused by tropospheric wa-ter vapour fluctuations. These fluctuations can be used toestimate the decorrelation in the interferometric data (Mid-delberg, Sault & Kesteven 2006). In addition, the absorp-tion due to atmospheric water vapour is estimated for eachmain antenna receiver by measuring the system temperature( T sys ) changes due to tropospheric emission.We used the seeing monitor data (to measure amplitudedecorrelation) in conjunction with T sys (to estimate tropo-spheric opacity) to develop semi-automatic flagging criteria.Specifically, we discarded data from all the periods in whichthere was decorrelation greater than 10 per cent. This im-proves the uniformity and data quality across all our observ-ing epochs.Flux density measurements for unresolved targetsources suffering from significant decorrelation could still berecovered using triple product techniques (see below), butimaging for these sources was not possible. Calibrators withsignificant decorrelation were excluded, and the blocks oftarget sources associated with those calibrator observationswere also excluded. Very occasionally, bad weather requiredlarge blocks of data to be edited out and hence a small num-ber of sources do not have near-simultaneous data at thelower frequencies (5 and 8 GHz). Primary flux calibration and bandpass calibration were car-ried out in the standard way using PKS B1934 −
63 as theprimary and PKS B1921 −
293 as the bandpass calibrator.For the secondary flux calibration we follow a non-standard procedure, which we describe here. Our follow-upobserving schedule follows the pattern: • a nearby secondary calibrator is observed for ∼ • a block of ∼
20 target sources are observed for ∼ • the secondary calibrator is re-observed for ∼ Manual inspection and flagging
Raw data
RPFITS files
Load data Apply flags
Calibrate bandpass and primary
Calibrate secondaries (scan based)
Plot tsys seemon List of flagged regions
Plots primary
Determine secondary fluxesCalibrate target sources
Scalar secondary fluxes
Further processing
List of flagged regions List of sources to repeatScalar secondary fluxes
Plots secs
Plots targets
Quick image target sources Determine primary beam correctionDetermine target fluxes
Images targetsScalar target fluxes
Identify extended sources
Extended source fluxes
Determine polarisation properties
Figure 1.
Diagram of the analysis pipeline process. we typically observed around ∼
50 secondary calibratorsduring one epoch of our observations. To calculate an accu-rate flux density for each secondary calibrator we calculatethe mean of the individual snapshot flux densities across thewhole run excluding only a snapshot which has a flux den-sity more than 2 standard deviations away from the mean.The rest of the snapshots are averaged to calculate the fluxdensity for that secondary. Finally, each target source is cal- c (cid:13) , 1–18 he AT20G Bright Source Sample ibrated using the secondary calibrator associated with itsobserving block. For each target source we calculate the po-sition, flux density, primary beam corrections and Stokesparameters.Full polarisation data ( I , Q , U and V Stokes parame-ters) are determined for all of the target sources. These arecalculated in the pipeline using a polarisation specific pro-cess. Firstly, a correction is applied for the time dependentphase difference (automatically monitored in real time atthe telescope) between the orthogonal, linear antenna feeds(which are referred to as x and y). After making this correc-tion, a small residual xy-phase signal still remains. Becausewe have insufficient secondary calibrator data to accuratelydetermine all the free parameters involved in instrumentalpolarisation corrections (e.g., leakages, residual xy-phase dif-ference, and time-dependant gains), leakage terms were cal-culated using the primary calibrator, PKS B1934 − Q and U Stokes parameters of the calibrators. The polarisation cali-bration was then applied to the target sources.
If a source is extended more than few arcsec (depending onthe array configuration) we will underestimate its total fluxdensity using either the image peak or the triple correlation.We could use the shortest spacing or integrate the imageover a larger area to recover the total flux density for anextended source, but this does not optimise the sensitivityfor a point source. Hence we need an automatic procedurecapable of distinguishing point-like sources from extendedsources. To do this we exploited the properties of the ob-served phase closure. The phase closure calculated on threeantennas (a baseline closure triangle) is the vector combi-nation of the phase of the correlated signal between eachcouple of antennas:Φ cl = Φ , + Φ , − Φ , . (1)It is null for a point source. It is also null for any flux densitydistribution that is an autocorrelation function such as asymmetrical Gaussian, but this is unlikely to occur for thesample of extragalactic objects we are considering in thispaper.In an array with more than three antennas the rootmean square (rms) of the phase closure can be calculated forall the possible combinations of three antennas in the array.Analogously to the three antenna case, it is expected to benull for a point source: the phase closure rms is differentfrom 0 if the source is extended or if there is more thanone source in the beam area. Receiver noise will contributeto the phase closure errors but the phase closure rms doesnot depend on antenna based instrumental and atmosphericphase effects or on the position of the source in the field.For each source we compare the observed phase closureto the predicted phase closure due to receiver noise. This is determined by Monte Carlo simulations of our observationsfor point sources with receiver noise added.Then we have defined the extendedness parameter as theratio of the predicted phase closure rms due to noise and theobserved value. The discrimination between point-like andextended sources is for the extendedness parameter equal to3, a good trade off, minimizing the wrong assignments tothe two classes. An incorrect assignment will result in a fluxdensity error of at most 20 per cent passing from one class toanother. The largest errors are made for faint objects (wellbelow 0 .
50 mJy).With the 214 m array the threshold means that a sourceis extended if it has significant flux density ( >
10 per cent)at 20 GHz on scales larger than 6 arcsec.The same criterion could be applied to all the frequen-cies, but, in the following, we consider that a source is ex-tended if its extendedness parameter is larger than 3 at20 GHz. A more refined method will be required to cor-rect for confusion due to faint sources especially at 5 GHz,but this correction is negligible for the present sample.
Source positions have been measured on the source centroidof the cleaned and restored images. Formal positional errorsin right ascension and declination have been obtained byquadratically adding a calibration term ( σ cal ) and a noiseterm ( σ n ). We have statistically determined the calibrationterm by cross-matching the Bright Source Sample (233 ob-servations in different epochs) with the International Coor-dinate Reference Frame catalogue (ICRF, Ma et al. 1998).The VLBI-measured positions in the ICRF catalogue areaccurate to − arcsec, so any discrepancy between thepositions of our target sources and the ICRF positions canbe attributed to positional errors in our sample. The rms po-sitional error is 0.5 arcsec in right ascension and declinationwith small variations due to changing weather conditions.For the Bright Source Sample the noise term is always neg-ligible. We have obtained the flux densities for bright point-likesources using the triple product method implemented in theMiriad task CALRED. The amplitude of triple product isthe geometric average of the visibility amplitudes in a base-line closure triangle A TP = p A , · A , · A , (2)and its phase is the phase closure (eq. 1).This way of measuring flux densities is particularly wellsuited for strong and point-like sources and it is able torecover the flux density lost in imaging because of phasedecorrelation. We have derived formal flux density errors, byadding quadratically a calibration term (gain error, σ gain )and a noise term ( σ n ). The gain error is a multiplicativeterm (i.e., it is proportional to the source flux density) andis a measure of the gain stability over time. We estimated σ gain for each observational epoch and frequency from thescatter in the visibility amplitudes of the calibrators in eachobserving run. Such average values for the gain errors were c (cid:13) , 1–18 Massardi et al. found to be of the order of a few per cent. The noise term isan additive term strictly related to the interferometer noisewhich is proportional to the system temperature. Since nosource has significant Stokes V , the rms noise levels in the V images have no gain error and are used as an estimate ofthe σ n value for each target source.For sources that have been defined as extended at20 GHz, integrated flux densities at 5, 8 and 20 GHz havebeen estimated from the amplitude of the signal measuredby the shortest baseline. Any source extended at 20 GHzis assumed to be extended at 5 and 8 GHz. Sources whichare extended at 5 or 8 GHz but core-dominated at 20 GHzwon’t be considered as extended according to this proce-dure. In this case we are assuming a dominant point sourceand the flux densities at all the frequencies will be for thecore and not the total source. The shortest baseline usedin the follow-up (see Table 1) is 60 or 80 m so we still un-derestimate flux densities for sources larger than 20 arcsec.For extended sources the error is increased by the squareroot of the number of baselines n base (normally 10 for our5-antenna follow-up arrays) to correct for the fact that theflux densities for these sources are estimated using only one(the shortest) baseline instead of n base . Images in Stokes U , Q and V are calculated for all the targetsources using the calibration procedure described in § V at our sensitivity the V image is used to estimate the noise error. If a source isdetected, P , the polarised flux, is calculated in the usualway P = p Q + U with no noise debias factor. For theintensity ( I ) we were able to avoid the effect of phase decor-relation by using the triple product but we don’t have anequivalent measure for U and Q . However, the troposphericphase decorrelation affects Stokes parameters Q and U inexactly the same way as Stokes I , so that we can use thetriple product amplitude, I tp , and the restored Stokes I im-age peak, I map , to calculate the factor by which the fluxdensity is reduced due to decorrelation χ = I map /I tp . Thenthe corrected polarised flux is P/χ .The error on P is P ERR = √ σ V /χ , where σ V is thenoise error from the V image: that is the propagation of theerror on P assuming that both the errors on U and Q areequal to σ V .For the non-detections ( P < σ V ) we calculate an upperlimit on P setting U and Q to 3 σ V and calculating the valueof P as above.To avoid bias in P , it is always measured at the posi-tion of the peak in I for point sources. For extended sourceswe need to integrate the polarisation vectors over the sourcewhich is the same as the integrated value of U and Q . Thishas been done for the extended sources that have been ob-served in the mosaic mode, but at this time we have notdetermined the integrated polarisation for the slightly ex-tended sources.Unfortunately, an instrumental phase problem hasspoiled the phase measurements for May 2007 observations:for this epoch flux densities could be recovered with triplecorrelation techniques, but the polarisation information hadto be flagged. Figure 2.
Equal area projection of the Southern sky in equato-rial coordinates, showing the BSS sources. The symbols size is afunction of the flux density at 20 GHz, as in the inset. The dottedlines indicate the regions of Galactic latitude b = ± ◦ and theGalactic plane. From the confirmed sources observed in the period 2004-2007 we selected those which have good quality data at20 GHz and flux densities above 0.50 Jy and Galactic lati-tude | b | > . ◦ . Some sources were observed at more thanone epoch, in which case the flux density selection thresh-old has been applied to the measurements with the highestquality and the smallest primary beam correction. To avoidany selection bias caused by variability, sources were onlyincluded if they were above the threshold for the epochswith the highest quality observation. This is necessary toavoid any bias caused by variability. The final distribution incoordinates, both equatorial and Galactic, is homogeneous(Fig. 2). The median errors in flux density estimation is4.8 per cent at 20 GHz, and 2.5 and 1.5 per cent respec-tively at 8 and 5 GHz.A small number of very extended sources are known tohave 20 GHz flux density above our 0.50 Jy cut. These arediscussed in § Tables 2 and 3 catalogue the 320 sources in the sample. Ta-ble 2 lists positions, flux densities, identifications with otheroptical or radio catalogues, and redshifts. Table 3 lists the in-formation about polarisation (polarised flux densities, frac-tions and angle of polarisation). For the full sample thesource names reflect the source J2000 equatorial coordinatesas ‘AT20G JHHMMSS-DDMMSS’. For sake of simplicity in c (cid:13) , 1–18 he AT20G Bright Source Sample this paper we will refer to the sources according to theirsequential number as listed in the first column of Table 2.The content of the columns are as follows for Table 2. • (1) Sequential number. An asterisk (‘*’) following thenumber indicates that the source is listed in the Appendix Aor has been commented on in the text. • (2–3) Right ascension and declination (J2000). The av-erage error in right ascension and declination is 0.5 arcsec(see § • (4–5) Flux density at 20 GHz and its error in Jy. • (6–7) Flux density at 8 GHz and its error in Jy. • (8–9) Flux density at 5 GHz and its error in Jy. When-ever available we give the results of 5 and 8 GHz observationsalmost simultaneous to the 20 GHz ones, otherwise we re-fer to the best observations available for the source at eachfrequency. • (10–11) Flux density at 1.4 GHz and its error fromNVSS (Condon et al. 1998). • (12–13) Flux density at 0.843 GHz and its error fromSUMSS (version 2.0). • (14–15) Redshift and its reference, obtained as dis-cussed in § • (16) Optical B magnitude for sources with SuperCOS-MOS counterparts. • (17) SuperCOSMOS identifications: ‘G’ for galaxies, ‘Q’for QSOs. A blank space indicates that no identification waspossible (see § • (18) Flags column where we collected some flags forsource properties in the following order:– the epoch of the 20 GHz observations: numbers referto the epoch reference number in Table 1;– spectral shape: ‘F’ for flat, ‘I’ for inverted,‘P’ forpeaked, ‘S’ for steep,‘U’ for upturning, as in Table 4;– galactic position: a ‘G’ indicates that the source iswithin 10 ◦ from the galactic plane;– epoch of observation at 8 and 5 GHz respectively, incase of not simultaneous observations (numbers refer tothe epoch reference number in Table 1): in such cases wehave listed the flux densities measured in the best obser-vation available.– extendedness: ‘E’ if the source is extended at 20 GHz,‘M’ if it has been observed in the mosaic mode. The fluxdensity for the ‘M’ sources corresponds to the integratedflux density of the source in the mosaic area;– a flag ‘C’ means that the source is listed in the ATcalibrator manual • (19) Alternative name from other well known catalogues(PMN, PKS) at radio frequency. • (20) Identification number in the WMAP 1-yr catalogue(Bennett et al. 2003).In Table 3 we collected the following columns • (1) Sequential number as in Table 2. • (2–3) Right ascension and declination (J2000). • (4–5) Integrated polarised flux in Jy and its error at20 GHz. • (6) Fractional polarisation at 20 GHz (per cent). • (7) Polarisation angle at 20 GHz in degrees. Figure 3.
Differential source counts at 20 GHz, with their Pois-son errors, normalised to Euclidean counts. The statistics is verypoor above ≃ • (8–9) Integrated polarised flux in Jy and its error at8 GHz. • (10) Fractional polarisation at 8 GHz (per cent). • (11) Polarisation angle at 8 GHz in degrees. • (12–13) Integrated polarised flux in Jy and its error at5 GHz. • (14) Fractional polarisation at 5 GHz (per cent). • (15) Polarisation angle at 5 GHz in degrees. The differential source counts for the present sample (Fig. 3)are in good agreement with the 9C counts at 15 GHz (Wal-dram et al. 2003) and with the WMAP counts at 23 GHz(Hinshaw et al. 2007; L´opez-Caniego et al. 2007), as well aswith the predictions of the model by De Zotti et al. (2005).However, we must beware of resolution effects. The sourcedetection technique is optimised for point-sources, and therewill be some bias against extended sources with angularsizes larger than about 30 arcsec. An outstanding case isFornax A, one of the brightest sources in the Southern sky,which was missed by our survey because its compact nu-cleus (and any other compact component) is fainter thanour blind scan detection limit (as was expected accordingto previous observations, e.g., Morganti et al. 1997) and itslobes are completely resolved by the 30-m baseline used forthe blind scan.By the same token, although no other bright sourceappears to have been completely missed by the AT20G Sur-vey, the flux densities of the most extended objects may fallbelow our threshold because they are underestimated. Toovercome this problem we have searched low-frequency cat-alogues for bright and extended sources, expected to haveintegrated 20 GHz flux densities above our 0.50 Jy thresholdbut missed by our selection (see § c (cid:13) , 1–18 Massardi et al. et al. in preparation). Except for Fornax A, all the knownbright extended sources in our area have been mosaiced.Another source of uncertainty in the sample selectionis variability, making sources move in or out of a given fluxdensity bin, depending on the epoch of observations. Sincewe have been gathering flux density measurements made atdifferent times we do not have a uniform view of the sur-veyed sky region. Only 30 BSS sources have more than oneobservation at 20 GHz in the 2002–2007 period (consideringalso the Pilot Survey observations), too small a sample fora meaningful analysis of variability. However, Sadler et al.(2006) found, at 20 GHz and on timescales of a few years,a median debiased variability index, that takes into accountthe uncertainties in individual flux density measurements,of 6.9 per cent, uncorrelated with the flux density, with onlya few sources more variable than 30 per cent. Also, a goodfraction (201 sources corresponding to the 63 per cent of thesample) of our sources are ATCA calibrators and have there-fore been observed repeatedly. Again, the variability turnsout to be relatively modest. Thus, variability does not affectsource counts, since it implies, on average, an equal numberof sources to change to lower or higher values of flux density.A much better assessment of variability will be provided bythe analysis of the full AT20G data. Since we selected theobservation to which we applied the selection threshold onthe basis of its quality and not on the basis of the flux den-sity itself (i.e. the best observation is not necessarily thatwith the higher value of flux density) we avoid any bias to-wards higher flux density values that could affect the sourcecounts.
The spectral index between the frequency ν and ν is de-fined as α = log( S /S )log( ν /ν ) (3)i.e. S ν ∝ ν α . Figure 4 shows the so called colour-colourradio plot (Kesteven et al. 1977): it is the comparison ofspectral indices at low and high frequencies. Only the al-most simultaneous data have been used in this analysis: thesub-sample consists of 218 sources. The diagram shows thevariety of spectral behaviours, with a relatively small num-ber of power-law spectra. Most of the points lie below thediagonal in Fig. 4, which implies that most sources steepenwith increasing frequency. The median of the difference ofthe spectral indices α − α is − α to extrapolate from 8 to 20 GHz could result, on average, ina 36 per cent error in the flux density estimation. Thus, sim-ple extrapolations in frequency using low-frequency spectralindices are highly unreliable.In Table 4 we have classified the spectra shapes on thebasis of the spectral indices between 5 and 8 GHz and be-tween 8 and 20 GHz. Examples of spectra in total intensityand polarisation are plotted in Fig. 5, where the NVSS andSUMSS measurements at 1.4 and 0.843 GHz are also shown.Table 4 also gives the fractions of ‘steep’- and ‘flat’-spectrum Table 4.
Distribution of spectral behaviour for the 218 BSSsources with almost simultaneous 5, 8 and 20 GHz data. The ab-breviations in the parenthesis in the second column refer to theclassification used to flag the sources according to their spectralbehaviour in Table 2. In the third column there are the numbersof object for each spectral class including a separate ‘very flat’source class. No selection has been applied for flat sources to getthe numbers in the last column. See the text for details.
No. (%)
No. (%)
Spectrum incl. flat class excl. flat class α > α > (17.9) (26.6) α > α < (23.4) (37.6) α < α > (0.9) (4.1) α < α < (20.2) (31.7) − . < α < . − . < α < . (37.6) α < − . (15.6) α > − . (84.4) α < − . (8.3) α > − . (91.7) sources, based on the commonly used classification (spectralindices smaller or larger than − . − . < α < − .
5. This will be discussed in moredetails in forthcoming papers on the whole Survey sample.As expected, the 20 GHz sample is dominated by flat-spectrum sources. A significant trend towards a steepeningof spectral indices at higher frequencies can be noted (seeFig. 6). The median spectral index between 5 and 8 GHz is0 .
11 and the fraction of ‘steep’-spectrum sources is ≃ − .
16 and the fraction of ‘steep’-spectrumsources almost doubles to ≃ . S lim , >
150 mJy)selected sample of the AT20G Survey (Sadler et al. 2007).The presence of spectral curvature provides valuable in-formation about the physical conditions in a radio source.Two mechanisms which generate spectral curvature are theenergy losses by synchrotron radiation causing steepening ofthe spectrum at high frequencies (e.g., Pacholczyk 1970) andoptical depth effects in compact sources at lower frequencieswhich may be due to either free-free opacity or synchrotronself absorption. The clear evidence for spectral steepeningof integrated flux density in the majority of the sources inthe 20 GHz Bright Source Sample (Fig 4) and the increasedspectral steepening observed at higher frequencies (Fig. 7)is in stark contrast to the lack of spectral steepening in theintegrated flux density for radio sources in low frequencysurveys (e.g. Laing and Peacock 1980). Spectral steepeningin the resolved structure in radio source lobes is commonlyseen and successfully modelled by a combination of energylosses and continual reacceleration in the lobes (e.g., Jaffe& Perola 1973, Subrahmanyan et al. 2006). c (cid:13) , 1–18 he AT20G Bright Source Sample Figure 5.
Some spectra as example of the large variety of spectral behaviour in total intensity (squares) and polarisation (diamonds) fora set of point sources. We selected examples of inverted, flat, peaked and steep total intensity behaviour similar (top panels) and different(bottom panels) to the polarisation behaviour. The triangles correspond to the fraction of polarisation. The low frequency values referto data from SUMSS (0.843 GHz) and NVSS (1.4 GHz) catalogues in total intensity (small squares) and, where available, polarisation(small diamonds).
Figure 4.
Colour-colour radio plot for the 218 sources with nearsimultaneous observations: the comparison of the spectral be-haviour in two ranges of frequencies shows the distribution of thespectral shapes in the whole sample. Power-law spectra sourceslie on the dashed diagonal line. A general steepening of the spec-tra from low (5 to 8 GHz) frequency to high (8 to 20 GHz) isclearly shown by the large number of sources with α < α . Figure 6.
Distributions of spectral indices α (upper panel), α (central panel), and α (bottom panel). Data at ∼ − (cid:13) , 1–18 Massardi et al.
Figure 7.
Plot of the median spectral indices as they have beencalculated for each frequency range for the BSS (solid lines), com-pared with NVSS to get the value between 1.4 and 4.85 GHz, forthe NEWPS catalogue (dashed lines) and for the observations at95 GHz of a flux limited sample of the AT20G Survey (dottedline).
The class of flat and inverted spectrum objects whichdominates the high frequency AT20G sample is quite differ-ent. The objects are small and almost certainly in a youngerevolutionary phase which includes the ‘Gigahertz PeakedSpectrum’ (GPS) sources (e.g., O’Dea 1998, Tinti & DeZotti 2006).The spectral steepening of sources in the lower leftquadrant in Fig. 4 could be due to synchrotron aging whichwould be much more rapid in the compact radio sourcesbecause the magnetic fields are higher.The BSS sample contains 64 objects (29.4 per cent ofthe 218 objects with simultaneous observations at 5, 8 and20 GHz) with α > α and α > .
3, i.e. peaking above5 GHz. Tinti et al. (2005) argued that a large fraction ofsources showing spectral peaks at several GHz are not trulyyoung (GPS) sources but blazars where a flaring, stronglyself-absorbed synchrotron component, probably originatedat the base of the relativistic jet, transiently dominates theemission spectrum. Although our evidence for variability(see § The comparison of the extendedness parameters at differentfrequencies (Fig. 8) for the BSS sources confirms the expec-tation that the extended, steep-spectrum radio lobes are lessand less prominent at higher frequencies.In Fig. 8 we can see three clear effects. There are point
Figure 8. sources spread into a circular patch by noise (a), a group ofsources extended at 5 GHz but still dominated by a pointcore at 20 GHz (b) and a group of sources extended atboth 5 and 20 GHz (c) which have a steeper spectral in-dex for the extended component. The solid line correspondsto α = − . − § > .
50 Jy, (that happens in 7 cases that are flagged withan ‘M’ in Table 2 and in Table 5) but present in the initialBSS selection only if their core component has flux densityabove 0 .
50 Jy (see Table 5). All of these have been observedwith the mosaic mode. For these sources we have integratedflux densities at 20 GHz but no flux densities at lower fre-quencies. Therefore we could not determine the extended-ness parameter at low frequencies or the spectral indices forthem, that are thus missing in Fig. 8 and 9.A summary of the properties of the extended sources inthe BSS is in Table 5. A few more sources that lie at theedge of our classification have been listed and commentedin the Appendix A. c (cid:13) , 1–18 he AT20G Bright Source Sample Table 5.
Table of extended sources in the BSS. The first column lists the sequential number of the sources as in Table 2. An ‘M’ indicatesthat they have been observed in mosaic mode. The 20 GHz flux densities in column 4 refer to the core region whereas those in column5 are the integrated flux densities. For 3 sources observed in mosaic mode, we believe we have acquired the flux density values only forsubregions, so we consider them as lower limits of the total integrated flux densities. P.A. is the position angle (in degrees) of the majoraxis of the source.Seq. RA δ S GHz S GHz S . GHz S . GHz P GHz z Size P.A.
Alternative core [ Jy ] [ Jy ] [ Jy ] [ Jy ] [ Jy ] [arcmin] [ ◦ ] Name
20M 01:33:57.6 -36:29:34.9 0.041 > .. ... .. ... .. ... ... 6.1 79 PKS 0131-36
52 04:08:48.75 -75:07:20.1 ... 0.86 ... 0.693 0.1 45
PKS 0410-75
69M 05:19:49.7 -45:46:44.2 1.33 ... .. ... .. Pictor A
71 05:22:57.94 -36:27:30.4 ... 3.91 ... 0.0553 0.5 55
PKS 0521-36
92 06:35:46.33 -75:16:16.9 ... 3.24 ... 0.653 0.2 90
PKS 0637-75
100 07:43:31.60 -67:26:25.7 ... 1.22 ... 1.51 0.2 14
PKS 0743-67
118 09:19:44.06 -53:40:05.1 ... 0.94 ... ... 0.3 40
PMN J0919-5340
157 12:05:33.37 -26:34:04.9 ... 0.84 ... 0.789 0.4 .
PKS 1203-26 > .. ... .. ... .. ... 0.00183 10.9 34 Centaurus A > .. ... .. ... .. ... 0.01254 31 53 PKS 1333-33
185 13:46:48.95 -60:24:29.0 ... 5.30 ... ... 0.1 .
PKS 1343-60 ... .. ... .. PKS 1610-60 ... .. ... .. PKS 2153-69
310 23:33:55.28 -23:43:40.8 ... 0.82 ... 0.0477 21 -43
PKS 2331-240 ... .. ... .. PKS 2356-61
Note: sources number 20, 69, 182, 284, 310 and 319 are characterized by a core and double lobes; 71, 92 and 100 have a core and a jet;179 is the inner double lobe of the giant radio galaxy Centaurus A with total extent of 5 degrees; 216 is a wide angle tail source; 310 isthe core region of a highly-extended radio galaxy: it is difficult to determine the correct size without a mosaic observation. Referencesfor redshift are given in Table 2. Useful references for the single sources are as follows. 20: Ekers et al. (1978); 69: Perley, Roser, &Meisenheimer (1997); 71: Birkinshaw, Worrall, & Hardcastle (2002);92: Schwartz et al. (2000); 182: Killeen, Bicknell, & Ekers (1986);284: Fosbury et al. (1998).
Figure 9.
The spectral indices between 8 and 20 GHz versus the20 GHz extendedness parameter.
All the follow-up measurements include polarisation. Oncethe low quality data have been removed from the sample, wetake, as ‘detections’, measurements of integrated polarisedflux at least 3 times higher than their errors (see § Table 6.
Matrix of the number of objects according to the com-bination of the total intensity and polarisation spectral behaviourfor the sources with almost simultaneous total intensity and po-larisation detection at 5,8 and 20 GHz. On the rows there are thespectra shapes in polarisation, on the columns the spectra shapesin total intensity. The spectral types are defined in Table 4. S → U I F P S
P ol. ↓ U I F P S Table 7.
The same as Table 6, but on the rows there are thespectral shape of the fractional polarisation. The spectral typesare defined in Table 4. S → U I F P S m [%] ↓ U I F P S per cent at 8 GHz and 1.7 per cent at 5 GHz (see Fig. 12).A similar trend was found by Burke et al. (in prep.) for thesub-sample observed in October 2006 during the observationrun dedicated to high sensitivity polarisation observations. c (cid:13) , 1–18 Massardi et al.
A detailed analysis of polarisation data will be presented inthat paper.As can be seen from Fig. 5 the spectra for polarisedflux density are very diverse and show little correlation withtotal flux density. This makes it even more difficult to predicthigh frequency polarisation properties from low frequenciesobservations than it is to predict I .There is no clear relation between the spectral proper-ties of the sources and their polarised flux, nor is there anyunique trend in the spectral behaviour of the total intensityand the polarised emission. The spectral shape in the polar-isation is often quite different from the spectral shape in thetotal intensity.The matrices of spectra in Tables 6 and 7 are complexbut not random. The diagonal cells dominate indicating thatnearly 50 per cent of the sources have polarised spectra sim-ilar to those in I . However the flat and peaked spectrumsources stand out with an excess of rising polarisation spec-tra. For sources with peaked spectra the polarised fractiongenerally decreases below the turnover frequency; an exam-ple of this behaviour in Fig. 5 (third panel, bottom row).This is not surprising as a polarisation mode which has ahigh emission coefficient should also have a high absorptioncoefficient, so in moving from optically thin (high frequen-cies) to optically thick (low frequencies) conditions we ex-pect that the ratio of the intensities in the two modes will de-crease. There are, however, other reasons why the polarisedfraction might decrease at lower frequencies, including: • depolarisation due to Faraday rotation intrinsic to thesources; • superposition of multiple components with different po-larised spectra; • depolarisation due to spatial variations in Faraday Ro-tation across the source; • bandwidth depolarisation due to very high levels ofFaraday Rotation.Figures 10 and 11 plot the polarised flux and fractionalpolarization as a function of flux density. The data from ourpilot observations (Sadler et al 2006) suggested a marginaltrend for weaker sources to have higher fractional polariza-tion. Although this seems to be present in Figure 11 themedian fractional polarization as a function of flux densityhas no trend and indicates that the apparent effect is dueto the increased density of points at lower flux levels. Thesources with a peak in the spectrum above 5 GHz have lowerfractional polarization at 20 GHz but this effect is not verypronounced. Fig. 12 shows the distribution of fractional po-larisation at 5, 8 and 20 GHz. Due to the lack of deep large area surveys at frequenciesabove 15 GHz the comparison of our results has to be donewith low frequency catalogues. Because of variability be-tween catalogue epochs a direct comparison can only providehints on the spectral behaviour as discussed in the previoussection. The results of the cross-correlation with NVSS at
Figure 10.
Integrated polarised flux as a function of total inten-sity 20 GHz flux. The bright source at P = 1 . Figure 11.
Fractional polarisation as a function of total intensity20 GHz flux density. The dashed lines shows the median fractionalpolarization by bins (the dotted lines indicates the bin ranges) offlux density for the full sample. Filled simbols refer to objectswith α > α and α > . c (cid:13) , 1–18 he AT20G Bright Source Sample Figure 12.
Distribution of fractional polarisation at 5, 8 and20 GHz. Dashed lines are the median values.
Figure 13.
Comparison of 5 GHz flux densities with the Parkesquarter Jy sample. sources that fall below ∼ . The contamination due to point sources is a crucial limi-tation to the CMB power spectrum determination on thesmaller angular scales (less than ∼
30 arcmin).The best frequency region to study the CMB is around
Figure 14.
Comparison of 5 GHz flux densities with the PMNcatalogue.
70 GHz, where the effects of foregrounds emissions is at aminimum, but in any case it is necessary to enlarge the fre-quency range as much as possible to try to single out all theforeground components and improve the component sepa-ration techniques. However, the efficiency of these methodsrelies on a good knowledge of the source populations to im-prove the capabilities of blind detection methods for ‘point’sources (L´opez-Caniego et al. 2007).The variety of source spectral behaviours implies that,as mentioned, it is extremely difficult to make reliable fluxdensity extrapolations from low to high frequency. Variabil-ity and confusion effects complicate the situation even more.Also, the forthcoming Planck mission will be stronglyconfusion limited. According to L´opez-Caniego et al. (2006),the 5 σ detection limits range from ≃
520 mJy at 30 GHz to ≃
180 mJy at 100 GHz, while the rms noise levels are farlower (from ≃
19 mJy at 30 GHz , Valenziano et al. 2007, to ≃
14 mJy at 100 GHz; Lamarre et al. 2003): this means thatthere is a lot of astrophysical information in Planck mapsbelow the confusion limit, that can be to some extent ex-tracted, e.g. using stacking techniques, thanks to the AT20GSurvey and follow-up observations at higher frequencies.As a test of high frequency predictions from low fre-quency samples we selected a sample from the PMN cata-logue with declination below − ◦ and | b | > ◦ and cross-matched it with SUMSS to obtain the low frequency spectralbehaviour. Then we divided it into sub-samples accordingto different limits in flux density at 5 and 1 GHz and/oraccording to different spectral indices at those frequencies.Finally, for each sub-sample we considered how many PMNsources are in the sample and how many of them have acounterpart in the BSS. In fact, the efficiency of the detec-tion depends on the ratio between what is present at theselection frequency and what is effectively found at the de-tection frequency ( detection rate ), and on the completenessof the sample obtained at the detection frequency.There are 154 BSS sources with declination below − ◦ and | b | > ◦ and 152 have a PMN counterpart. However, 35PMN counterparts have flux density at 5 GHz below 0.50 Jy,so that a low frequency selection threshold at 500 mJy wouldhave lost them. Selecting only inverted sources ( α . . > c (cid:13) , 1–18 Massardi et al.
Figure 15.
Comparison of the BSS 20 GHz flux densities withthe NEWPS catalogue at 23 GHz. α . . > . α . . > .
5) results in a low detection rate(3 .
6, 3 .
2, 3 . .
2, 22 .
2, 9 . − ◦ and | b | > ◦ have flux density above 500 mJy and nocounterpart in the BSS. Combining spectral and flux densitylimits or adding further selection criteria at 1 GHz improvesthe detection rate but at the cost of a very low completenessof the high frequency sample. Thus, it is clear that low fre-quency catalogues could provide positions for constrainedsearch techniques (cf. L´opez-Caniego et al. 2007), but areinadequate to forecast the high frequency population.The comparison of flux densities with WMAP map-based catalogues shows a good agreement in general (weused the NEWPS catalogue as in Gonz´alez-Nuevo et al. 2007in Fig. 15). The epochs of observations partially overlap, butsince the WMAP maps have been averaged over three years,transient phenomena have been smoothed out.Furthermore, both space- and ground-based missionsrequire a set of carefully selected sources to work as calibra-tors for pointing, total intensity flux densities and polarisa-tion angle.The Bright Source Sample we have discussed is well-suited for CMB studies in the next years. It collects a sampleof the brightest sources in the Southern sky, that, thanks tothe low variability observed, will be observable or detectableby any detection method. The observational frequency is, sofar, the closest to the region of the spectrum of interest forCMB studies.The Bright Source Sample provides a direct test ofsource detection algorithms, quantifying the completeness,the fraction of spurious detections, the effective beam size(and therefore the flux calibration) and the possible pres-ence of biases in flux density estimates. It also provides arich list of candidate flux density and pointing calibratorsover a large fraction (37 per cent) of the sky. Finding suitable polarisation calibrators for CMB ex-periments is much more complicate. For example, the largelow frequency beams of the Planck satellite (33 arcmin beamat 30 GHz) dilute the polarised signals by summing over dif-ferently oriented polarisation vectors. Thus, finding sourceswith large enough polarised flux density within such beamsis very hard (Figs. 10 and 11). A more extensive discussionabout such calibration will appear in Burke et al. (in prepa-ration). To make optical identifications for objects in the BrightSource Sample, we searched the SuperCOSMOS catalogue(Hambly et al. 2001) near the positions of all sources. Ob-jects within 10 ◦ of the Galactic plane (flagged with a ‘G’ inTable 2) were excluded from the analysis because the pres-ence of foreground stars and Galactic dust extinction makesoptical identifications incomplete in this region. This cutoffin Galactic latitude excluded 69 of the 320 BSS sources. Twoother sources were also excluded from the optical analysis:sources number 57 and 160 (according to the sequential nu-meration in Table 2) lie so close to bright foreground starsthat no optical identification is possible from the DSS im-ages. Source number 73 lies within the boundaries of theLarge Magellanic Cloud and its identification is uncertain.An optical object was accepted as the correct ID if itwas brighter than B J =22 mag and lay within 2.5 arcsec ofthe radio position. Monte Carlo tests imply that at least 97per cent of such objects are likely to be genuine associations(Sadler et al. 2006).We found a DSS identification for 238 of 249 sources,with 235 of the optical IDs having B J . B J magnitude is 18.6 for QSOsand 17.7 for galaxies (see Fig. 16).We have also checked in the NASA ExtragalacticDatabase (NED ) for optical identifications in order to dis-tinguish between Galactic and extragalactic objects: noneof the sources in the BSS which have a clear identificationare Galactic objects (i.e. HII regions, planetary nebulae orSNRs). After completing the optical identifications, we checked(NED) to search for published redshifts. A listed redshiftwas accepted only if it could be traced back to its origi-nal source and appeared to be reliable. 177 of the 249 BSSobjects (71 per cent) had a reliable published redshift, in-cluding three of the sources which are blank fields on theDSS (these objects were identified in deeper optical imagesby other authors).The 72 objects without a published redshift includeseven objects (sources number 10, 19, 30, 42, 85, 221 and 278 (cid:13) , 1–18 he AT20G Bright Source Sample as listed in Table 2) which have a redshift listed in NED. Inthese cases, we were either unable to trace back to its orig-inal source, or considered to be unreliable for other reasons(PKS 0332 − Redshifts for two BSS objects (sources number 33, a QSO at z =0.466 QSO, and 313, a QSO at z =0.626 based on a singlebroad emission line identified as MgII) were obtained froma pre-release version of the final redshift catalogue from the6dF Galaxy Survey (Jones et al. 2004, 2007 in preparation).The redshift for source 138 has been measured with the ESO3.6 m telescope by PGE and his collaborators (Edwards etal. in preparation).Optical spectra of nine other BSS sources were obtainedat the ANU 2.3 m telescope in April and June 2007 by R.W.Hunstead and two of the authors (PH and EM). Redshiftswere measured for seven of these objects (sources number 77,98, 140, 162, 166, 208 and 246). The spectra of two otherobjects (number 68 and 78) showed a featureless optical con-tinuum from which no redshift could be measured.Among the 186 objects with redshifts, 144 are QSOsand 36 are galaxies. The median redshift is 1.20 for the QSOsand 0.13 for the galaxies(Fig. 17). No correlation is observedbetween redshift and total 20 GHz flux density or polarisedflux. As noted by Sadler et al. (2006) there is a correla-tion between redshift and optical magnitude for galaxies inthe AT20G sample, but this does not apply to the AT20Gquasars (see Fig. 18). Six BSS objects with good–quality optical spectra (ei-ther from the published literature or from unpublished6dF/2.3 m data), have no measured redshift because thespectra are featureless. Such objects generally fall into theBL Lac class, though it is possible that some of them fall inthe ‘redshift desert’ at z ∼ . − . The correlation between the difference of the spectral indicesat high and low frequencies ( α − α ) with redshift (Fig. 19)shows a clear curvature in the spectra. Note that the 20 GHzflux densities from the higher redshift objects correspond toflux densities from the higher frequencies in the rest frame(scaling as 1 + z ). Since the median redshift of the QSO inthe sample is 1.20, the steepening is occurring at frequency ν >
50 GHz in the rest frame, and grows steeper to above ν >
70 GHz in the rest frame for the objects at z ∼ . Figure 16.
B-magnitude distribution.
Figure 17.
Redshift distribution. The model by De Zotti et al.(2005) has been overlapped for comparison. complete redshift information, since there may be selectioneffects in the sub-samples with existing redshift information.
We have presented a complete sample of 320 sources se-lected within the AT20G Survey catalogue as those havingflux density S > .
50 Jy, | b | > . ◦ . Almost simulta-neous 5 and 8 GHz observations have been used for spectralbehaviour analysis.Information on polarisation are available at all the fre-quencies. We found that the median fractional polarisationis increasing with frequency.Neither the high frequency total intensity nor the po-larisation behaviour can be estimated from low frequencyinformation. We examined a set of issues that support thisstatement: • the colour-colour plots show a broad range of spectralshape: most sources spectra are not power-law so do notallow easily extrapolation from one frequency to the other; c (cid:13) , 1–18 Massardi et al.
Figure 18.
B-magnitude versus redshift for galaxies and QSO.
Figure 19.
Plot of the difference between spectral indices α and α with redshift. • the comparison with low frequency selected samplesshowed that by increasing the constraints on the low fre-quency sample the number of low frequency objects recov-ered also at high frequency increased, but that the complete-ness of the predicted high frequency sample gets poorer. Itis necessary to fine tune the conditions on the low frequencysample to obtain a good trade-off between completeness andidentification rate, but there is no way to select a low fre-quency sample that guarantees that all the sources will con-stitute a complete high frequency sample; • the polarisation spectral shape does not agree in all thecases with the total intensity: the lack of knowledge on polar-isation properties, together with unpredictable polarisationspectral behaviour make any forecast extremely difficult.It is clear that actual high frequency samples are betterthan trying to predict them from lower frequencies.So, the Bright Source Sample constitutes an unprece-dented collection of information at 20 GHz, that will turnto be of importance by itself and for any future observationsat high radio frequencies.The whole AT20G Survey, in fact, will improve the ra- diosource population knowledge to much lower flux densi-ties. This amount of information will be of crucial interestfor the next generation telescope, to provide good sampleof calibrators, and for the CMB targeted missions, as a testfor point source detection techniques, as a help in point-source removal in any component separation exercise and asa list of candidate pointing, flux and possibly polarisationcalibrators.Even with the relatively superficial analysis presentedhere, we find interesting new physical effects from this sam-ple: • spectral steepening is common in this class of object; • the spectral steepening correlates with redshift, possi-bly due to changing rest frame frequency; • sources with spectral peaks in the GHz range are com-mon in this sample and have high depolarisation on the lowfrequency side of the peak. ACKNOWLEDGEMENTS
MM and GDZ acknowledge financial support from ASI (con-tract Planck LFI Activity of Phase E2) and MUR.We gratefully thank the staff at the Australia TelescopeCompact Array site, Narrabri (NSW), for the valuable sup-port. The Australia Telescope Compact Array is part of theAustralia Telescope which is funded by the Commonwealthof Australia for operation as a National Facility managed byCSIRO.This research has made use of the NASA/IPAC Ex-tragalactic Database (NED) which is operated by the JetPropulsion Laboratory, California Institute of Technology,under contract with the National Aeronautics and SpaceAdministration.This research has made use of data obtained from theSuperCOSMOS Science Archive, prepared and hosted bythe Wide Field Astronomy Unit, Institute for Astronomy,University of Edinburgh, which is funded by the UK ParticlePhysics and Astronomy Research Council.We thank the referee for his useful comments and cor-rections.
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APPENDIX A: INDIVIDUAL SOURCES NOTES
Table 2 source 61 : PKS 0454 −
81 appears in the scanmaps, but the follow-up data were degraded by bad weatherand we didn’t have the opportunity to re-observe it. Forthis source we obtained a flux density measurement from itsobservations as a secondary calibrator in October 2006.
Table 2 source 92 (PKS 0637-752) is a quasar with anasymmetric jet seen in radio and Xray images (Schwartzet al. 2000). The tabulated flux density is dominated by thecore with about 10% in the 15 arcsec jet. It is one of thelargest (100 kpc) and most luminous jets known with prop-erties similar to 3C273.
Table 2 source 109 (PMN J0835 − α = +0 .
88, buthas no obvious optical counterpart. Although the Galacticlatitude is relatively low ( b = 11 ◦ ), the optical extinction isonly 1.1 mag in the B band. The lack of optical ID suggeststhis could be a distant radio galaxy rather than a QSO. Table 2 source 151 (PKS 1143 − Table 2 source 211 (PKS 1548 −
79) is a relatively nearby( z = 0 .
15) galaxy with an unresolved radio source which hasa steep spectrum in our 5, 8 and 20 GHz data. The galaxyhas strong optical emission lines, and has been studied indetail by Tadhunter et al. (2001).
Table 2 source 221 appears to be one component of asource (PKS 1622 −
29) which is double (component sepa-ration ∼ . Table 2 source 258
The AT20G source (corresponding toPKS 1932 −
46) is flagged as extended, and the image ap-pears to show a compact double. The source is a 30 arcsecdouble at 5 GHz (Duncan & Sproats 1992). The optical po-sition given in NED is associated with a z = 0 .
231 galaxyat (J2000) 19:35:56.5 − Table 2 source 273 (PKS 2052 −
47) is a z = 1 . c (cid:13) , 1–18 Massardi et al.
The calibrator data suggest that our AT20G observation ofthis object in October 2004 took place during the declin-ing stage of a flaring phase, during which the flux densityof the source changed rapidly. This fast change in flux andpolarisation properties is clearly visible in our data, withthe 20 GHz flux density decreasing by a factor of 2.5 in twodays. This makes it difficult to give a reliable value for theflux density and fractional polarisation of this source.
Table 2 source 292 (PKS 2227 − Table 2 source 310 , flagged as extended, appears to bethe core of a well known and highly-extended radio galaxyPKS 2331 − Table 2 source 319 , (PKS 2356 − A TEX file preparedby the author. c (cid:13) , 1–18 able 2. The AT20G Bright Source Sample.Seq. RA δ S S . S . S . S . z z Ref. B J mag
Opt. Flags Alternative WMAP Jy ] [ Jy ] [ Jy ] [ Jy ] [ Jy ] ID name ID1 00:04:35.65 -47:36:19.1 0.87 ... ... ... ... La01 ... ...
Wr83 ... ... ... ... ... ...
Ta93 ... ... ... ...
Ja84 ... ...
Pe76 ... ... ... ... ... ...
Wh88 ... ... ... ... ... ... ... ... ... ...
Wi83 ... ... ... ... ... ...
Wh88 ... ... ... ... ... ...
Wi83 ... ... ...
St93 ... ... ... ... ... ...
Wr83 ... ... ... ... > ... ... ... ... ... ... ... ... ... RC3
Io96 ... ...
Wr83 ... ...
Wr83 ... . 4S....C PKS 0142-278 .25 01:53:10.19 -33:10:26.7 0.54
Wr77 ... ...
Wr83 ... ...
Wi00
Ja02
Dr97 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
Ar67 ... ...
Ma95 ... ...
Wi00 ... ... ... ... ... ... ... ... ... ...
Ja78 ... ...
Wi83 ... ...
Ba95 ... ... ... ... ... ... ... ...
FCSS ... ...
Sb05 ... ...
Wh88 able 2.
Continue.Seq. RA δ S S . S . S . S . z z Ref. B J mag
Opt. Flags Alternative WMAP Jy ] [ Jy ] [ Jy ] [ Jy ] [ Jy ] ID name ID46 03:48:39.28 -16:10:17.2 0.94 ... ... ... ... ... ... El01 ... ... ... ...
Pe76
St94 ... ...
Dr97 ... ...
Ta93 ... ...
Hu78
Wr77
He04 ... ...
Wr83
Wi83 ... . 1F..... PKS 0435-300 .58 04:39:00.83 -45:22:22.6 0.70 ... ... ... ... ... ...
Mo78 ... ... ... .. ... ... ... ... ... ... ... ... St94 ... ...
Wi83 ... ...
Su04 ... ...
St89 ... ...
Wr77 ... ...
St89 ... ...
St93 ... ... ... ... ... ... ... ... ... ... ... ...
RC3 ... ...
Wr79 ... ...
Ke85 ... ... ... ... ... ... ... ... ... ...
Ca00 ... ... ... ...
Pe76 ... ...
Os94 ... ...
SSO ... ... ... ... ... ... ... ...
Ja02 ... ... ... ... ... ...
Pe98 ... ... ... ... ... ... ... ...
Ja02 ... ...
Wr79 ... ...
Hu78 ... ... ... ... ... ... ... ... ... ...
Pi98 ... ...
Qu95 ... ... ...
Sb06 ... ... ... ... ... ...
Ja84 ... ... ... ... ... ... ... ... able 2.
Continue.Seq. RA δ S S . S . S . S . z z Ref. B J mag
Opt. Flags Alternative WMAP Jy ] [ Jy ] [ Jy ] [ Jy ] [ Jy ] ID name ID92* 06:35:46.33 -75:16:16.9 3.24 ... ... Hu78 ... ...
Ho03 ... ... ... ... ... ... ... ... ... ... ... ... ... ...
Ja02 ... ... ... ... ... ...
SSO ... ... ... ... ... ...
Al94 ... ... ... ... ... ... ... ... ... . 1.G...C PKS 0745-330 .102 07:48:03.09 -16:39:50.3 0.77 ... ... ... ... ... . 4PG...C PMN J0748-1639 .103 07:56:50.65 -15:42:04.7 0.73 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... . 2I..... PMN J0835-5953 .110 08:36:39.21 -20:16:58.9 2.68 ... ...
Fr83 ... ... ... ... ... . 1FG.... PKS 0835-339 .112 08:45:02.47 -54:58:08.8 0.92 ... ... ... ... ... ... ... ... ... ...
Wh88 ... ...
St93 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
Wr79 ... ... ... ... ... ... ... ... ... ...
Pe79 ... ... ... ... ... ... ... . 2UG.... PMN J0958-5757 .124 10:01:59.89 -44:38:00.2 0.81 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
Ja02 ... ... ... ... ... ... ... ... ... ...
Dr97 ... ... ... ... ... ...
Wr79
Ja84 ... ...
St89 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... able 2.
Continue.Seq. RA δ S S . S . S . S . z z Ref. B J mag
Opt. Flags Alternative WMAP Jy ] [ Jy ] [ Jy ] [ Jy ] [ Jy ] ID name ID136 10:48:06.58 -19:09:35.3 1.24 ... ... St93 ... ... ... ... ... ... ... ... ... ... ... ...
ESO ... ... ... ... ... ... ... ... ... ... ... ...
SSO ... ... ... ... ... ...
Dr97 ... ... ... ... ... ...
Pe79 ... ... ... ... ... ... ... ... ... ... ... ... ... ...
Ja84 ... ...
Ja02 ... ...
Dr97 ... ... ... ... ... ... ... ... ... ... ... ... ... . 3SG...C PKS 1133-681 .151* 11:45:53.58 -69:54:04.1 0.86 ... ... ... ... ... ...
Dr97
St89 ... ... ... ... ... ... ... ... ... ... ... . 3P....C PKS 1150-834 .156 11:54:21.79 -35:05:29.2 0.89 ... ...
Ta93 ... ...
Wr79 ... ...
Ja02 ... ...
Ja02 ... ... ... ... ... . 4F....C PKS 1213-17 173161 12:18:06.26 -46:00:30.3 0.70 ... ...
Mu84 ... ... ... ... ... ...
SSO ... ... ... ... ... ...
Sa76 ... ...
Od91 ... ...
SSO ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
Ja84 ... ...
Ja02 ... ...
De84 ... ... ... ... ... ... ... ... ... ... ... . 5.G.... PMN J1303-5540 .174 13:05:27.47 -49:28:04.8 0.73 ... ... ... ... ... ...
RC3 ... ... ... ... ... ... ... ... ... ... ... . 5.G.... PMN J1315-5334 .176 13:16:08.09 -33:38:58.9 1.50 ... ...
Ja82 ... ... ... ... ... ...
RC3 ... ...
Dr97 > ... ... ... ... ... ... ... RC3 ... ... ... ... ... . 4IG...C PMN J1326-5256 . able 2.
Continue.Seq. RA δ S S . S . S . S . z z Ref. B J mag
Opt. Flags Alternative WMAP Jy ] [ Jy ] [ Jy ] [ Jy ] [ Jy ] ID name ID181 13:29:01.13 -56:08:02.6 0.93 ... ... ... ... ... ... ... ... ... ... ... . 5.G...C PMN J1329-5608 .182* 13:36:39.0 -33:57:58.2 > ... ... ... ... ... ... ... ... ... R3C ... ... ... ... ... ... ... ...
Ja02 ... ... ... ... ... ... ... . 5.G.2E. PKS 1343-60 .186 13:57:11.27 -15:27:29.5 0.55 ... ...
Wi83 ... ...
Dr97 ... ... ... ...
Ja02 ... ... ... ... ... ... ... ... ... . 5.G.... PMN J1419-5155 .190 14:24:32.24 -49:13:49.3 2.64 ... ... ... ... ... . 5S....C PKS B1421-490 .191 14:24:55.56 -68:07:57.8 1.50 ... ... ... ... ... ... ... . 5.G44.. PKS 1420-679 .192 14:27:41.31 -33:05:31.9 0.80 ... ... ... ... ... . 5I....C PKS 1424-328 193193 14:27:56.30 -42:06:18.9 2.75 ... ...
Wh88 ... ...
Dr97 ... . 5S....C PKS 1430-155 .195 14:38:09.46 -22:04:54.6 0.80 ... ...
Dr97 ... ... ... ...
Ja02 ... ... ... ...
Ja02 ... ...
Hu78 ... ... ... ... ... ... ... ...
Dr97 ... ... ... ... ... . 5FG...C PMN J1514-4748 .203 15:17:41.76 -24:22:20.3 3.45 ... ...
Ro77 ... ...
He04 ... ... ... ... ... ... ... ... ... ... ... . 5.G.... PMN J1534-5351 .206 15:34:54.68 -35:26:23.8 0.73 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... . 5.G.... PMN J1535-4730 .208* 15:46:44.51 -68:37:28.9 0.51 ... ... ... ... ... ... ... ...
SSO ... ... ... ... ... ...
Ja84 ... ...
Ta01 ... ...
Wr79 ... ... ... ... ... ... ... ... ... ... ... . 1.G.... PMN J1600-4649 .214 16:03:50.67 -49:04:05.1 0.55 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
R3C ... ... ... ... ... ...
Hu80 ... ...
Th90 ... ... ... ... ... ... di94 ... . 4.....C PKS 1622-253 .221* 16:26:06.04 -29:51:26.6 1.79 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... . 1.G44.. PMN J1636-4101 .223 16:47:37.79 -64:38:01.0 0.75 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... able 2.
Continue.Seq. RA δ S S . S . S . S . z z Ref. B J mag
Opt. Flags Alternative WMAP Jy ] [ Jy ] [ Jy ] [ Jy ] [ Jy ] ID name ID227 17:00:53.27 -26:10:52.6 0.51 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... . 5.G22.C PMN J1701-5621 .229 17:03:36.34 -62:12:38.2 1.05 ... ... ... ... ... ... ... ... ... . 1SG.... NVSS J170918-352521 .231 17:09:34.40 -17:28:52.7 0.55 ... ... ... ... ... ... ... ... ... . 4.....C PMN J1709-1728 .232 17:13:10.02 -34:18:27.7 1.03 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... RC3 ... ... ... ... ... . 1PG...C PMN J1733-3722 .238 17:33:40.43 -79:35:55.7 1.11 ... ... ... ... ... ... ... ... ... ... ... ... ... . 2S....C PKS 1740-517 .240 18:02:42.66 -39:40:07.8 1.41 ... ... ... ... ... ... ... . 5.G...C PMN J1802-3940 .241 18:03:23.56 -65:07:36.8 1.24 ... ... ... ... ... ... ... ... ... ... ... ...
Da79 ... ... ... ... ... ... ... ... ... ... ... ... ... ...
SSO ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... . 4.G...C PKS 1830-211 .249 18:34:27.29 -58:56:36.7 1.43 ... ... ... ... ... ...
Ja84 ... ... ... ... ... ...
Ha03 ... ...
Go98 ... ...
Ha03 ... ... ... ...
Wi83 ... ... ... ... ... ...
Ja78 ... ... ... ... ... ...
Ta93 ... . 1...... PKS 1932-46 .259 19:37:16.22 -39:58:01.6 1.76
Dr97 ... ...
Ta93 ... ...
Ja84 ... ...
Os94 ... ...
Gr83 ... ...
Ja82
Br75 ... ...
Br75 ... ...
Os94 ... ...
Hu78 ... ... ... . 1..44.C PKS 2002-375 .270 20:09:25.45 -48:49:53.9 0.88 ... ... ... ... ... ...
Fa87 able 2.
Continue.Seq. RA δ S S . S . S . S . z z Ref. B J mag
Opt. Flags Alternative WMAP Jy ] [ Jy ] [ Jy ] [ Jy ] [ Jy ] ID name ID271 20:11:15.70 -15:46:40.2 2.10 ... ... Pe79 ... ...
Br75 ... ... ... ... ... ...
Ja84
Ja02 ... ... ... ... ... ... ... ... ... ...
Wh88 ... ... ... ... ... ... ... ...
Os94 ... ... ... ... ... ... ... ... ... ...
Wi86 ... ... ... ... ... ...
RC3 ... ... ... ... ... ...
Wh88 ... ... ... ... ... ...
Ja78 ... ...
Mo82 ... ...
Wi83 ... ... ... ... ... ...
Wr83 ... ... ... ... ... ... ... ... ... ... ... ...
Ja84 ... ... ... ... ... . 2P..... PKS 2236-572 201295 22:43:26.47 -25:44:31.4 0.65 ... ... ... ... ... ...
St93 ... ... ... ...
Wi83
Pe79 ... ... ... ... ... ... ... ... ... ... ... ...
RC3 ... ... ... ... ... ...
Pe76 ... ...
Hu78 ... ...
Wi86
Dr97 ... ... ... ... ... ...
Pe72 ... ... ... ... ... ...
Wr83
Ja78 ... ...
Wi76 ... ... ... ... ... ... ... ... ... . 2S....C PKS 2333-528 195313* 23:45:12.47 -15:55:08.0 0.93 ... ... ... ...
Wr83 ... ...
Pe79 able 2.
Continue.Seq. RA δ S S . S . S . S . z z Ref. B J mag
Opt. Flags Alternative WMAP Jy ] [ Jy ] [ Jy ] [ Jy ] [ Jy ] ID name ID316 23:56:00.67 -68:20:03.6 0.84 ... ... Pe72 ... ...
Ja84 ... ... ... ... ... ... ... ... ... ... ... ... ... ... ... ...
Lo96 able 3.
The AT20G Bright Source Sample: polarization data.Seq. δ P GHz [ Jy ] m GHz [%] θ GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ]1 00:04:35.65 -47:36:19.1 0.017 < ... ... ... 0.038 < ... ... ... < ... ... ...7 00:38:14.72 -24:59:01.9 0.017 < ... ... ... 0.008 ... ... ... ... ... ... ...10 00:58:46.64 -56:59:11.4 0.012 < ... ... ... 0.006 < ... ... ...13 01:06:45.11 -40:34:19.5 0.055 ... ... ... ... ... ... ...14 01:17:48.81 -21:11:07.4 0.013 < ... ... ... 0.005 ... ... ... ... ... ... ... ... ... ... ...21 01:34:32.14 -38:43:33.7 < ... ... ... 0.010 < ... ... ... 0.004 < ... ... ... 0.009 < ... ... ... ... ... ... ... ... ... ... ...32 02:36:31.11 -29:53:55.1 0.011 < ... ... ...33 02:36:53.27 -61:36:15.2 0.026 ... ... ... ... ... ... ...34 02:40:08.13 -23:09:15.8 0.030 < ... ... ... 0.013 ... ... ... ... ... ... ...38 03:09:56.12 -60:58:39.0 0.008 ... ... ... ... ... ... ...39 03:11:55.33 -76:51:51.2 0.026 ... ... ... 0.015 ... ... ... ... ... ... ...43 03:36:54.12 -36:16:06.0 0.005 able 3. Continue.Seq. δ P GHz [ Jy ] m GHz [%] θ GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ]46 03:48:39.28 -16:10:17.2 0.043 ... ... ...48 03:52:11.00 -25:14:50.2 0.013 < ... ... ... 0.007 ... ... ... 0.029 ... ... ... ... ... ... ... ... ... ... ...53 04:16:36.61 -18:51:08.9 0.019 < ... ... ... < ... ... ...57 04:37:36.56 -29:54:03.9 0.019 ... ... ... ... ... ... ... ... ... ... ...62 04:53:14.64 -28:07:37.4 0.047 ... ... ... 0.086 ... ... ... ... ... ... ...66 05:13:49.10 -21:59:17.4 0.077 ... ... ... ... ... ... ...69 05:19:49.7 -45:46:44.2 1.400 ... ... ... ... ... ... ...70 05:22:34.40 -61:07:57.0 0.018 ... ... ... ... ... ... ... ... ... ... ...72 05:25:06.48 -23:38:11.1 0.004 < ... ... ... < ... ... ...73 05:29:30.02 -72:45:28.2 0.020 ... ... ... ... ... ... ...74 05:36:28.45 -34:01:10.8 0.016 ... ... ... ... ... ... ...81 06:04:25.13 -42:25:30.1 0.042 ... ... ... ... ... ... ...82 06:08:59.76 -22:20:21.3 0.009 < ... ... ...83 06:09:41.03 -15:42:41.6 0.221 < ... ... ... 0.011 ... ... ... ... ... ... ...87 06:27:06.73 -35:29:16.1 0.012 ... ... ... ... ... ... ...90 06:33:26.76 -22:23:22.6 0.020 < ... ... ... < ... ... ...91 06:34:58.99 -23:35:12.6 0.046 able 3. Continue.Seq. δ P GHz [ Jy ] m GHz [%] θ GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ]92 06:35:46.33 -75:16:16.9 ... ... ... ... ... ... ... ... ... ... ... ...93 06:48:14.18 -30:44:19.3 0.020 ... ... ... ... ... ... ...97 07:31:06.67 -23:41:47.8 0.031 ... ... ... 0.022 < ... ... ...100 07:43:31.60 -67:26:25.8 ... ... ... ... ... ... ... ... ... ... ... ...101 07:47:19.72 -33:10:46.6 0.026 ... ... ... ... ... ... ...102 07:48:03.09 -16:39:50.3 < ... ... ... < ... ... ... < ... ... ...103 07:56:50.65 -15:42:04.7 0.013 < ... ... ...106 08:16:40.41 -24:21:05.8 0.020 < ... ... ... 0.012 < ... ... ... 0.031 < ... ... ... < ... ... ... < ... ... ...110 08:36:39.21 -20:16:58.9 0.026 ... ... ... 0.015 ... ... ... 0.016 ... ... ... ... ... ... ... ... ... ... ...119 09:20:43.25 -29:56:30.6 0.008 < ... ... ... 0.003 ... ... ... ... ... ... ...122 09:27:51.90 -20:34:50.4 < ... ... ... 0.018 < ... ... ... < ... ... ... < ... ... ...124 10:01:59.89 -44:38:00.2 < ... ... ... 0.006 ... ... ... ... ... ... ...127 10:14:50.33 -45:08:41.2 0.015 ... ... ... ... ... ... ...128 10:18:28.76 -31:23:53.3 0.019 ... ... ...129 10:23:43.47 -66:46:47.8 0.015 < ... ... ... < ... ... ... 0.003 able 3. Continue.Seq. δ P GHz [ Jy ] m GHz [%] θ GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ]136 10:48:06.58 -19:09:35.3 0.024 ... ... ... ... ... ... ... ... ... ... ...139 11:01:54.42 -63:25:22.6 0.012 ... ... ... ... ... ... ...141 11:03:52.17 -53:57:00.8 0.007 ... ... ... ... ... ... ...144 11:07:12.85 -68:20:50.6 < ... ... ... 0.008 < ... ... ... 0.009 < ... ... ... < ... ... ...150 11:36:02.21 -68:27:05.4 0.028 < ... ... ... 0.002 < ... ... ...157 12:05:33.37 -26:34:04.9 ... ... ... ... ... ... ... ... ... ... ... ...158 12:09:02.64 -24:06:19.8 0.020 < ... ... ... < ... ... ...160 12:15:46.88 -17:31:45.3 0.033 ... ... ...162 12:27:26.74 -44:36:39.8 < ... ... ... ... ... ... ... ... ... ... ...163 12:45:53.72 -16:16:44.5 0.074 ... ... ... ... ... ... ... ... ... ... ...165 12:48:23.88 -19:59:18.4 ... ... ... ... ... ... ... ... ... ... ... ...166 12:48:28.53 -45:59:47.8 0.020 ... ... ... ... ... ... ... ... ... ... ...168 12:54:37.24 -20:00:56.2 ... ... ... ... ... ... ... ... ... ... ... ...169 12:54:59.80 -71:38:18.4 ... ... ... ... ... ... ... ... ... ... ... ...170 12:57:59.20 -31:55:15.2 0.067 ... ... ... 0.026 ... ... ... ... ... ... ... ... ... ... ...174 13:05:27.47 -49:28:04.8 ... ... ... ... ... ... ... ... ... ... ... ...175 13:15:04.24 -53:34:36.0 ... ... ... ... ... ... ... ... ... ... ... ...176 13:16:08.09 -33:38:58.9 0.068 ... ... ... ... ... ... ... ... ... ... ...178 13:21:14.02 -26:36:10.2 ... ... ... ... ... ... ... ... ... ... ... ...179 13:25:27.7 -43:01:07.0 ... ... ... ... ... ... ... ... ... ... ... ...180 13:26:49.58 -52:55:35.6 0.025 able 3. Continue.Seq. δ P GHz [ Jy ] m GHz [%] θ GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ]181 13:29:01.13 -56:08:02.6 ... ... ... ... ... ... ... ... ... ... ... ...182 13:36:39.0 -33:57:58.2 ... ... ... ... ... ... ... ... ... ... ... ...183 13:37:52.37 -65:09:25.4 ... ... ... ... ... ... ... ... ... ... ... ...184 13:42:04.79 -20:51:29.0 0.042 < ... ... ...185 13:46:48.95 -60:24:29.1 ... ... ... ... ... ... ... ... ... ... ... ...186 13:57:11.27 -15:27:29.5 0.007 < ... ... ... 0.013 ... ... ... ... ... ... ... ... ... ... ...189 14:19:35.22 -51:54:58.8 ... ... ... ... ... ... ... ... ... ... ... ...190 14:24:32.24 -49:13:49.3 ... ... ... ... ... ... ... ... ... ... ... ...191 14:24:55.56 -68:07:57.8 ... ... ... ... ... ... ... ... ... ... ... ...192 14:27:41.31 -33:05:31.9 ... ... ... ... ... ... ... ... ... ... ... ...193 14:27:56.30 -42:06:18.9 ... ... ... ... ... ... ... ... ... ... ... ...194 14:33:21.45 -15:48:45.0 ... ... ... ... ... ... ... ... ... ... ... ...195 14:38:09.46 -22:04:54.6 0.009 ... ... ... ... ... ... ...196 14:54:27.46 -37:47:33.1 ... ... ... ... ... ... ... ... ... ... ... ...197 14:54:32.92 -40:12:32.6 ... ... ... ... ... ... ... ... ... ... ... ...198 14:57:26.70 -35:39:10.8 ... ... ... ... ... ... ... ... ... ... ... ...199 15:07:04.78 -16:52:30.3 ... ... ... ... ... ... ... ... ... ... ... ...200 15:08:38.98 -49:53:01.8 ... ... ... ... ... ... ... ... ... ... ... ...201 15:13:57.01 -21:14:57.7 0.039 ... ... ... < ... ... ...202 15:14:40.05 -47:48:28.7 ... ... ... ... ... ... ... ... ... ... ... ...203 15:17:41.76 -24:22:20.3 0.137 ... ... ... ... ... ... ... ... ... ... ...206 15:34:54.68 -35:26:23.8 ... ... ... ... ... ... ... ... ... ... ... ...207 15:35:52.23 -47:30:21.8 ... ... ... ... ... ... ... ... ... ... ... ...208 15:46:44.51 -68:37:28.9 0.013 ... ... ... ... ... ... ...209 15:50:58.66 -82:58:06.9 ... ... ... ... ... ... ... ... ... ... ... ...210 15:54:02.44 -27:04:39.8 0.016 < ... ... ...211 15:56:58.72 -79:14:04.9 < ... ... ... 0.009 < ... ... ...212 15:59:41.26 -24:42:40.2 0.032 < ... ... ... 0.003 ... ... ... ... ... ... ...214 16:03:50.67 -49:04:05.1 < ... ... ... ... ... ... ... ... ... ... ...215 16:04:31.09 -44:41:31.3 0.060 ... ... ... ... ... ... ...217 16:17:17.96 -58:48:06.1 0.073 < ... ... ... ... ... ... ... ... ... ... ...220 16:25:46.99 -25:27:39.3 0.060 ... ... ... ... ... ... ...221 16:26:06.04 -29:51:26.6 0.098 ... ... ... ... ... ... ...222 16:36:55.26 -41:02:00.7 0.006 ... ... ... ... ... ... ...223 16:47:37.79 -64:38:01.0 0.131 ... ... ... ... ... ... ...224 16:48:42.34 -33:01:47.7 < ... ... ... 0.012 ... ... ... ... ... ... ... ... ... ... ... able 3. Continue.Seq. δ P GHz [ Jy ] m GHz [%] θ GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ]227 17:00:53.27 -26:10:52.6 0.009 ... ... ... ... ... ... ...228 17:01:44.84 -56:21:55.6 ... ... ... ... ... ... ... ... ... ... ... ...229 17:03:36.34 -62:12:38.2 0.005 ... ... ...231 17:09:34.40 -17:28:52.7 0.107 ... ... ... ... ... ... ...232 17:13:10.02 -34:18:27.7 0.075 ... ... ... ... ... ... ...234 17:13:50.99 -32:26:08.9 0.025 < ... ... ... ... ... ... ... ... ... ... ...237 17:33:15.32 -37:22:30.6 < ... ... ... 0.064 ... ... ... ... ... ... ... ... ... ... ...239 17:44:25.25 -51:44:45.2 < ... ... ... < ... ... ... < ... ... ...240 18:02:42.66 -39:40:07.8 ... ... ... ... ... ... ... ... ... ... ... ...241 18:03:23.56 -65:07:36.8 0.003 ... ... ... ... ... ... ...242 18:09:57.79 -45:52:41.2 < ... ... ... ... ... ... ... ... ... ... ...243 18:19:34.98 -63:45:48.2 ... ... ... ... ... ... ... ... ... ... ... ...244 18:19:45.29 -55:21:21.8 0.031 ... ... ... ... ... ... ... ... ... ... ...246 18:29:12.57 -58:13:54.9 < ... ... ... 0.022 ... ... ... ... ... ... ...248 18:33:39.95 -21:03:41.2 0.110 ... ... ... ... ... ... ...249 18:34:27.29 -58:56:36.7 0.118 ... ... ... ... ... ... ...252 19:11:09.71 -20:06:55.7 ... ... ... ... ... ... ... ... ... ... ... ...253 19:12:40.38 -80:10:07.0 0.011 ... ... ... ... ... ... ...254 19:23:32.27 -21:04:33.4 0.079 ... ... ... 1.329 ... ... ... ... ... ... ...257 19:32:44.95 -45:36:38.3 0.017 ... ... ... ... ... ... ...258 19:35:57.62 -46:20:43.8 0.017 ... ... ... ... ... ... ...259 19:37:16.22 -39:58:01.6 0.083 ... ... ... ... ... ... ...260 19:39:24.83 -63:42:45.4 < ... ... ... < ... ... ... < ... ... ...261 19:39:26.74 -15:25:43.3 ... ... ... ... ... ... ... ... ... ... ... ...262 19:40:25.74 -69:07:58.0 0.014 < ... ... ... < ... ... ... < ... ... ...264 19:56:59.41 -32:25:46.0 0.007 ... ... ... ... ... ... ...266 20:00:57.08 -17:48:57.4 ... ... ... ... ... ... ... ... ... ... ... ...267 20:03:24.04 -32:51:42.4 < ... ... ... ... ... ... ... ... ... ... ...268 20:05:17.30 -18:22:03.1 ... ... ... ... ... ... ... ... ... ... ... ...269 20:05:55.03 -37:23:39.7 < ... ... ... ... ... ... ... ... ... ... ...270 20:09:25.45 -48:49:53.9 0.010 ... ... ... ... ... ... ... able 3. Continue.Seq. δ P GHz [ Jy ] m GHz [%] θ GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ]271 20:11:15.70 -15:46:40.2 ... ... ... ... ... ... ... ... ... ... ... ...272 20:24:35.58 -32:53:35.6 ... ... ... ... ... ... ... ... ... ... ... ...273 20:56:16.40 -47:14:47.9 0.009 ... ... ... ... ... ... ...274 20:57:41.64 -37:34:02.3 0.035 ... ... ... ... ... ... ...275 21:05:44.98 -78:25:35.0 0.011 ... ... ... ... ... ... ...277 21:21:13.19 -37:03:08.9 0.010 < ... ... ...278 21:26:30.69 -46:05:48.2 < ... ... ... < ... ... ... 0.028 ... ... ... ... ... ... ... ... ... ... ...280 21:42:31.00 -24:44:39.1 ... ... ... ... ... ... ... ... ... ... ... ...281 21:46:29.73 -77:55:54.9 0.046 ... ... ... ... ... ... ... ... ... ... ...283 21:51:55.58 -30:27:53.8 0.023 ... ... ... ... ... ... ...284 21:57:06.07 -69:41:23.2 0.087 ... ... ... ... ... ... ...285 21:58:06.28 -15:01:09.3 ... ... ... ... ... ... ... ... ... ... ... ...286 22:00:54.69 -55:20:08.2 < ... ... ... < ... ... ... 0.005 ... ... ... ... ... ... ...288 22:06:10.60 -18:35:39.5 0.235 ... ... ... ... ... ... ...289 22:07:43.82 -53:46:34.1 0.060 ... ... ... ... ... ... ...291 22:29:00.22 -69:10:29.7 0.003 ... ... ... ... ... ... ...292 22:30:40.34 -39:42:51.8 ... ... ... ... ... ... ... ... ... ... ... ...293 22:35:13.28 -48:35:58.7 0.018 ... ... ... 0.007 ... ... ... ... ... ... ...296 22:46:16.77 -56:07:46.1 0.068 ... ... ... ... ... ... ...300 22:56:41.26 -20:11:41.3 0.031 ... ... ... ... ... ... ...301 22:57:10.50 -36:27:44.6 ... ... ... ... ... ... ... ... ... ... ... ...302 22:58:05.97 -27:58:21.7 0.029 ... ... ... ... ... ... ...303 23:03:03.02 -18:41:26.1 ... ... ... ... ... ... ... ... ... ... ... ...304 23:03:43.46 -68:07:37.7 0.030 ... ... ... ... ... ... ...305 23:14:48.56 -31:38:38.6 < ... ... ... 0.015 ... ... ... ... ... ... ... ... ... ... ...309 23:31:59.43 -38:11:47.4 0.021 ... ... ... ... ... ... ... ... ... ... ...311 23:34:44.88 -52:51:19.4 0.008 ... ... ... ... ... ... ... ... ... ... ...314 23:48:02.63 -16:31:12.0 0.021 ... ... ... 0.054 ... ... ... ... ... ... ... ... ... ... ... able 3. Continue.Seq. δ P GHz [ Jy ] m GHz [%] θ GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ] P . GHz [ Jy ] m . GHz [%] θ . GHz [ deg ]316 23:56:00.67 -68:20:03.6 < ... ... ... ... ... ... ... ... ... ... ...317 23:57:53.41 -53:11:12.5 < ... ... ... 0.027 ... ... ... ... ... ... ...319 23:59:04.7 -60:55:01.1 0.053 ... ... ... ... ... ... ...320 23:59:35.41 -31:33:44.9 < ... ... ... 0.0060.001