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First Results from the Very Small Array K. Grainge

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First Results from the Very Small Array K. Grainge FIRST RESULTS FROM THE VERY SMALL ARRAY K. GRAINGE Cavendish Astrophysics, Cavendish Laboratory, Madingley Road, Cambridge, England. The Very Small Array (VSA) is an interferometric array designed to measure the primordial fluctuations in the Cosmic Microwave Background (CMB). Interferometers are largely free of several of the sources of systematic error from which balloon-bourne bolometer experiments suffer. In addition, the VSA has two key differences from other CMB interferometers: in­ dividually tracking elements and a dedicated source subtraction system. In the first year of operation, the VSA has observed the CMB in a total area 101 square degrees in three regions of sky at 34 GHz on angular scales of 3° .6-0°.4 (angular multipoles (=150-900). Primor­ dial anisotropies are detected at high significance in all three regions, and the signals seen in overlapping paintings agree well with each other. 1 Introduction The Very Small Array is a heterodyne interferometerarray built by a collaboration consisting of the Cavendish Astrophysics group in Cambridge, the Jodrell Bank Observatory in Manchester and lnstituto de Astroflsica de Canariasin Tenerife. The VSA is designed to detect the primor­ dial CMB anisotropies which are an image of the temperature fluctuations in the universe at the epoch of recombination, z ::::; llOO. In this paper we give an overview of the VSA telescope and describe its key differences from other CMB experiments. We present the results from the first year of observations with the VSA in its compact configuration. These data show clear detections of primordial anisotropies in the CMB on angular scales of 3° .6-0° .4. The Instrument 2 The VSA is a 14-element interferometer operating at a frequency between 26 and 36 GHz with a receiver bandwidth of 1.5 GHz. The telescope is sited at the Teide Observatory, Tenerife at "' 1: Figure The VSA main array inside its ground screen. The small, compact array antennas are fitted. an altitude of 2400m. We have found this to be an excellent site for 30 GHz observations with the atmosphere contributing approximately 5 K to the total system temperature of 30 K, and less than 53 of the observing time being lost due to bad weather. The antennas used for the observations reported here have illuminated apertures of 143 mm diameter, giving a primary beam of 4.6° FWHM and are designed for use in a "compact array" , for observations up to a maximum spherical harmonic mnltipole of about 900. For subsequent observations at higher C = C the antennashave been modified by increasingthe length of the horn and fittinga larger mirror, giving an aperture of 322 mm and a beam of 2.0° FWHM; observations are currently underway with this "extended array" . The telescope is located inside a metal enclosure to reduce pick up from the ground (Figure 1) . 3 Comparison with other CMB experiments Many different experimental approaches have been used to measure features in the CMB, mostly using switched- or swept-beam systems with total power detectors. Such systems can be made with very high sensitivity, particularly when using broad-band bolometric detectors, and progress in the field hasmostly been due to improved techniques for suppressing unwanted signals such as atmospheric emission and differential emission from the telescope optics. An alternative approach is to use interferometers. These provide excellent rejection of systematics, since only signals entering both antennas with the correct path and phase modulation are detected while signals such as ground radiation and atmospheric emission are strongly suppressed 1 . They can also be made relatively insensitive to receiver gain fluctuations, have an envelope beam which can be determined to high accuracy and, since they are generally mounted on the ground, have small pointing errors. Interferometric systems also offer the opportunity to target a specific range of angular scales on the sky, determined by the spacing of the elements of the array, and are well-suited to measuring the power spectrum of the CMB since they directly sample the Fourier modes on the sky which can then be converted to a power spectrum. Their main disadvantages are the restricted bandwidth compared to bolometers, due to the need for coherent receivers, and the relative complexity and expense of the correlator. In addition to the VSA, two other interferometers have recently published measurements of the CMB, the Cosmic Background Imager 3 and the Degree Angular Scale Interferometer 2• The VSA hastwo key differences from these telescope, a dedicated source subtraction baseline, and individually tracking elements which allow simple Fourier filtering of our data. 3. 1 Source Subtraction Strategy The most important foreground which contaminates the VSA observations is from radio-galaxies and quasars. The strategy that we have adopted to remove their contribution is to survey the VSA fields with the Ryle Telescope at 15 GHz. We then monitor the sources found with a single baseline interferometer comprising a pair of 3.7-m antennas working at the VSA's observing frequency of 34 GHz. It is not possible to use existing low-frequency surveys at 1.4 GHz for the source finding because this strategy would miss a significant number of inverted spectrum radio­ sources. The source subtraction antennas are fitted with identical front ends and IF electronics to the main VSA array, and are also sited at the Teide Observatory. We therefore believe that cross-calibration between the source subtraction system and the VSA itself can be made very accurate. 3.2 Filtering based on AstronomicalFr inge Rate Because the VSA antennas are individually tracking, the relative pa.th length to the antennas changes during the course of an observation, resulting in the correlated astronomical signal being phase modulated at a known rate. We can therefore filter any unwanted signals using a Fou rier filter on the time-ordered data. Figure 2 shows an example of the application of this technique to an observation of Jupiter in the presence of contaminating signals from the Sun. Figure 2a is an image produced by simply mapping the raw data. Jupiter, whose flux density is approximately 90 Jy, is almost entirely obscured by synthesised beam grating rings from the Sun, whose flux density is approximately 10 MJy. However, since the Sun is 11° from Jupiter it has a different astronomical fri nge rate. We therefore fringe rotate the data to the Sun, apply a high pass Hanning (cosine-edged) filter in the Fo urier plane to remove the signal from the Sun, and then fringe rotate back to the position of Jupiter. It is then necessary to flag out any data where the fringe rates of Jupiter and the Sun are sufficiently similar that the filtering will have affected the Jupiter signal. Mapping this filtered data now gives Figure 2b. All signals from the Sun have been removed and a point source at the correct position and flux density for Jupiter are visible. The noise on the map is consistent with the thermal noise predicted from theory. In practice we never observe a field within 40° of the Sun, beyond which the primary beam of the telescope provides a further attenuation of approximately 30 dB. 4 Observing Programme and Results We have made deep mosa.iced observations of eight fields in three evenly spaced regions of sky. Each mosa.iced field was observed for "' 400 hours, reaching a thermal noise of approximately 30 mJy. Mosa.icing in this way enables us to increase the £-resolution of our measurements whilst also reducing sample variance. Although the CMB power spectrum is calculated directly from the complex visibility data obtained from our observations (see Taylor, these proceedings) , image-plane analysis provides valuable consistency checks on the data. Figure 3 shows a map of the VSA3A field. CMB fluctuations a.re clearly visible within the primary beam. In addition, one bright source is clearly detected with a primary beam corrected flux density of 214 ± 36 mJy. This is in good agreement with that determined by the source subtracter antennas of 238 mJy. The rightha.nd map shows an 'a.utosubtracted' image where we have ta.ken the time-ordered visibility data and reversed the sign of alternate visibilities. On the time sea.le between adj acent visibility points (64 seconds) , both astronomical and spurious signals a.re effectively coherent, and a.re thus cancelled out. In contrast, the noise on adj acent visibilities is completely uncorrelated, and the rms noise on the map is unaffe cted. This technique therefore provides a robust estimate of our thermal noise. Since the rms power in the a.utosubtracted map is consistent with that far from the centre of our standard field map, we conclude that there Figure 2: VSA observation of Jupiter (flux density � 90 Jy) with the Sun (flux density � 10 MJy) at 11° clistance. Note that the FWHM of the primary beam is 4.6°. The two plots show the map (a) before and (b) after filtering, with contour levels of 20 Jy and 5 Jy respectively. The filtered image shows no significant contamination from the Sun and maps Jupiter as a point source with the preclicted flux density. -100 100 -100 47 46 46 45 44 " " z 43 z 43 0I 144 !i 42 � .. g 41 �:!I 41 .. 40 39 38 38 15 50 40 30 20 10 1550 40 30 20 10 RIGHT ASCENSION {J2000} RIGHT ASCENSION (J2000} Figure 3: Maps VSA field (VSA3A) at resolution. before source subtraction. of a sn.mple full Left: the map Positions of sources to be subtracted are marked with crosses, and FWHM of the primary beam is inclicated the by a circle. CMB fluctuations are clearly visible within the envelope and a bright point source is detected beam, to the north-west of the pointing centre.
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