Introduction

FLAMES (Fibre Large Array MultiElement Spectrograph) is the VLT FibreFacility, installed and commissioned atthe Nasmyth A focus of UT2 (KueyenTelescope). FLAMES was built and as-sembled in about four years through aninternational collaboration between teninstitutes in six countries and three con-tinents. It had first light with the fibre linkto the red arm of UVES on April 1, andwith the GIRAFFE spectrograph on

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No. 110 – December 2002

Installation and Commissioning of FLAMES,the VLT Multifibre FacilityL. PASQUINI1, G. AVILA1, A. BLECHA2, C. CACCIARI3, V. CAYATTE4, M. COLLESS5, F. DAMIANI6, R. DE PROPRIS5, H. DEKKER1, P. DI MARCANTONIO7, T. FARRELL8, P. GILLINGHAM8, I. GUINOUARD4, F. HAMMER4, A. KAUFER9, V. HILL4, M. MARTEAUD4,A. MODIGLIANI1, G. MULAS10, P. NORTH2, D. POPOVIC8, E. ROSSETTI3, F. ROYER2, P. SANTIN7, R. SCHMUTZER9, G. SIMOND2, P. VOLA4, L. WALLER8, M. ZOCCALI1

1European Southern Observatory, Garching, Germany; 2Observatoire de Genève, Sauverny, Suisse; 3INAF-OABo, Bologna, Italy; 4Observatoire de Paris-Meudon, Meudon, France;5Research School of Astronomy and Astrophysics, Australian National University, Canberra, Australia;6INAF-OAPa, Palermo, Italy; 7INAF-OATs, Trieste, Italy; 8Anglo Australian Observatory, Sydney, Australia;9European Southern Observatory, Chile; 10INAF-OACa, Capoterra, Cagliari, Italy

Figure 1: The fibre positioner and GIRAFFEas seen during the GIRAFFE integration onthe Nasmyth platform. The positioner is look-ing towards the Nasmyth Focus, where the cor-rector is placed (only the corrector support isvisible). GIRAFFE is open, and all the opto-mechanics components are visible. The pic-ture was taken when the OzPoz enclosure hadnot yet been installed.

(www.eso.org/instruments/FLAMES/manuals).

The characteristics of the differentobserving modes are given in Table 1,which summarizes the observationalcapabilities of FLAMES.

UVES-Fibre

UVES is the high-resolution spectro-graph of the VLT UT2 (D’Odorico et al.2000). Each positioner plate has eightfibres connected to the UVES Red Arm.With an aperture on the sky of 1 arcsec,the fibres project to five UVES pixelsgiving a resolving power of ~ 47,000.Only the three standard UVES Red set-ups are offered, with central wave-lengths of 520, 580 and 860 nm, re-spectively.

GIRAFFE

GIRAFFE is a medium-high resolu-tion spectrograph (R = 6000–33,000)for the entire visible range (370–950nm). It is equipped with two echelles forlow and high resolution and uses inter-ference order sorting filters to select therequired spectral range within an order.The typical spectral coverage in oneexposure is 60–100 nm in low resolu-tion and 20–40 nm in high resolution.

The fibre system feeding GIRAFFEconsists of the following components(cf. Avila et al. 2002 for details):

• MEDUSA fibre slits, one per posi-tioner plate. Up to 132 separate objects(including sky fibres) are accessible inMEDUSA single fibre mode, each withan aperture of 1.2 arcsec on the sky.

• IFU slits. 15 deployable IntegralField Units (IFU) per plate, each con-sisting of an array of 20 square micro-lenses, for a total (almost rectangular)aperture of ~ 3 × 2 arcsec. For eachplate there are also 15 IFU dedicated tosky measurements.

• 1 ARGUS slit. A large integral unitconsisting of a rectangular array of 22by 14 micro-lenses, fixed at the centreof one positioner plate. Two scales areavailable: one with a sampling of 0.52arcsec/micro-lens and a total apertureof 12 by 7 arcsec, and one with a sam-

pling of 0.3 arcsec/micro-lens and a to-tal coverage of 6.6 by 4.2 arcsec. 15deployable ARGUS single sky fibresare also available.

GIRAFFE is operated with 30 fixedset-ups (22 high-resolution + 8 low-res-olution modes). For performance esti-mates (based on expected transmis-sion curves and performances) theuser is referred to the Exposure TimeCalculator (ETC), available at the ESOETC web page: www.eso.org/observ-ing/etc

The FLAMES observing software(OS) is coordinating the operations ofthe Telescope, the Positioner and theUVES and GIRAFFE spectrographs. Inaddition, it allows COMBINED observa-tions: that is the simultaneous acquisi-tion of UVES and GIRAFFE spectrawith the specific observing modes listedin Table 1.

LAYOUT

Figure 1 shows a view of the maincomponents of the FLAMES facility, thefibre Positioner and GIRAFFE, as seenon the telescope platform. The Cor-rector support is also visible, attachedat the Nasmyth rotator.

1. Corrector

The optical corrector is a big doubletof 880 mm free aperture. The functionof the corrector is to give an excellentimage quality over the whole 25 arcminFLAMES field of view and to provide apupil located at the centre of curvatureof the focal plate, to avoid vignetting foroff-axis fibres. When the whole opticaltrain is taken into account (including tel-escope optics and vignetting), the ef-fective transmission of the corrector de-pends on the observing wavelengthand on the distance of the object to thefield centre, varying from about 78% at370 nm to 90% in the visible, and 82%at 1 micron.

A system of lamps has been installedto illuminate the Nasmyth screen,which is located in the telescope cen-trepiece, about 2 m in front of the focalplane.

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July 3. We report here on the complexprocess of integration and commission-ing, and we compare the expected andobserved astronomical requirements.

Overview

FLAMES is the multi-object, interme-diate and high-resolution fibre facility ofthe VLT. Mounted at the Nasmyth Aplatform of UT2 it offers a large correct-ed field of view (25 arcmin diameter)and it consists of several components,developed by several consortia (Aus-tralis (AAO/ANU/UNSW), Paris-Meu-don (OPM), Geneva and Lausanne(OGL), Bologna, Cagliari, Palermo andTrieste (ITAL Consortium)) and ESO:

• An optical Corrector, providing ex-cellent image quality and telecentricityover the full field of view of 25 arcmindiameter (developed at ESO).

• A Fibre Positioner hosting twoplates. While one plate is observing,the other one is positioning the fibresfor the subsequent observations, there-fore limiting the dead time between ob-servations (developed at AAO).

• A link to the UVES spectrograph(Red Arm) via eight single object fibres /plate (developed at OPM in collabora-tion with ESO).

• A high and intermediate resolutionoptical spectrograph, GIRAFFE (devel-oped at OPM and ESO).

• Three types of fibre systems: MEDU-SA, IFU, ARGUS (developed at OPM).

• A coordinating observing software,that allows simultaneous UVES and GI-RAFFE observations (developed at ITALConsortium in collaboration with ESO).

• As for all VLT instruments, a fullData Reduction Software (DRS) pack-age is provided, integrated in the ESOData Flow Software (developed at OGLwith contributions from OPM (GI-RAFFE) and ITAL Consortium (UVES-Fibre)).

The FLAMES components are de-scribed in detail elsewhere (cf. e.g. Avi-la et al. 2002, Blecha et al. 2000, Gilling-ham et al. 2000, 2002, Jocou et al. 2000,Mulas et al. 2002, Pasquini et al. 2000,Royer et al. 2002), on the ESO web siteand in the FLAMES User Manual

Spectrograph Mode Nr. Objects Aperture R Coverage (nm)

UVES8 Red Arm (with sky) 8 1 47000 200, 400UVES7 Red Arm (with sky) Simult. Calibration 7 1 47000 200GIRAFFE MEDUSA 132 (with sky) 1.2 20500* λ/22*GIRAFFE MEDUSA 132 (with sky) 1.2 7000* λ/7*GIRAFFE IFU 15 (+15 sky) 2 × 3 33000* λ/22*GIRAFFE IFU 15 (+15 sky) 2 × 3 11000* λ/7*GIRAFFE ARGUS 1 11.5 × 7.3 33000* λ/22*

Or 6.6 × 4.2 GIRAFFE ARGUS 1 11.5 × 7.3 11000* λ/7*

Or 6.6 × 4.2

Table 1. Observational capabilities of FLAMES.

*The GIRAFFE resolving power (R) and wavelength coverage given here are only average values. They will vary for the different set-ups by ± ~ 20%. Apertureis in arcseconds.

2. Fibre Positioner (OzPoz)

The Fibre Positioner (OzPoz) is atthe core of the FLAMES facility. OzPozis a rather large and complex systemequipped with four plates, two of whichare currently in use (cf. Gillingham et al.2002). The Positioner can be subdivid-ed into the following subsystems:

• Plates: Two metallic disks, on whichthe magnetic buttons holding the fibresare attached. Each of the plates has ahole in the centre. In one plate this holehosts ARGUS. Each plate has a curva-ture of 3950 mm to match the curvatureof the corrector focal surface.

• Retractors: Mechanical systemsmaintaining the fibres in constant ten-sion. Each fibre is equipped with oneretractor. The retractors are externallythe same for all fibres.

• Exchanger: Main structure holdingthe plates. The exchanger can performtwo main movements: it can retract (orapproach) the Nasmyth adapter to en-gage the plate (or disengage it).Furthermore, it can rotate the upperpart in order to perform the plate ex-change.

• r-θ system and Gripper: This unitgrips and releases the magnetic but-tons at the positions reached via the r-θ (polar) robot. The gripper requires aback-illumination system, sending abrief light flash along each fibre fromthe spectrograph towards the plate. Acamera records this back-illuminationlight and performs an image analysis.The back-illumination light is used forseveral purposes: to reach the requiredpositioning accuracy and to detect if themagnetic button was properly liftedfrom the plate.

• OzPoz is equipped with a calibra-tion box which can direct the light eitherfrom a tungsten, or from a Th-Ar, orfrom a Ne lamp. In this way FF, Th-Arand Ne calibrations can be obtained forGIRAFFE and for UVES.

• Field Acquisition Bundles (FACBs):Four magnetic buttons on each plateare equipped with a system of 19 co-herent fibres each. This bundle is usedto obtain images of fiducial stars, oneper button. The four images are viewedby an ESO technical CCD; the imagecentroids are computed and the properoffsets are calculated to centre the fidu-cial stars into the bundles. Each FACBbundle has an effective diameter of 2.4arcseconds.

• Positioning Software: It is writtenaround the so-called delta-task pro-gram, developed initially for the 2dFsystem at AAO. This program allowscrossing between the fibres and it de-termines the button movements se-quence.

3. Buttons and Fibre Systems

FLAMES is equipped with differenttypes of fibres for UVES and for the dif-

ferent modes forGIRAFFE. It isworth mentioningthat, in addition tothe object or skyfibres coming fromthe Positioner,each GIRAFFE fi-bre slit has five fi-bres devoted to si-multaneous wave-length calibration,which provide si-multaneous cali-bration spectra foreach observationacquired with GI-RAFFE (Hammeret al. 1999).

The histogram distribution of thetransmission is given in Figure 2 for thedifferent fibre types. The FLAMES fi-bres have been proven very robust, sothat out of more than 1000 fibres, onlya few were broken (and replaced) at theend of the whole process.

3.1 UVES fibres

Each of the positioner plates hostseight 55-metre fibres to guide the lightto the UVES spectrograph located onthe opposite Nasmyth Platform. Fromthe UVES simultaneous calibrationbox, one additional 5-metre fibre reach-es the UVES-Fibre slit. The fibre cen-tres are separated by 1.7 fibre core,implying that there is some degree ofcontamination between adjacent fibres;the UVES-Fibre Data Reduction Soft-ware (Mulas et al. 2002) has been de-veloped to obtain a complete de-blend-ing of the spectra.

3.2 MEDUSA fibres

Each plate also hosts 132 MEDUSAfibres. The separation between each ofthe MEDUSA fibres in a sub-slit is 2.26× the fibre core; this ensures a fibre-to-fibre contamination below 0.5 per cent.

During our tests we have experi-enced some (about 2%) increasedstray light due to reflections from the fil-ter and the (polished) exit slit. We willattempt to reduce this with a stray lightmask that reduces the width of the re-flecting exit slit viewed from the side ofthe filter.

3.3 IFU fibres

Each Integral Field Unit (IFU) buttonis composed of twenty micro-lensesarranged in a rectangular shape. Themicro-lenses are 0.52″ squares. Themovable IFUs are a unique characteris-tic of GIRAFFE and can be placed allover the FLAMES field of view, with theexception of a small fraction of the cen-tre annulus. Each plate hosts fifteenIFUs plus fifteen Sky IFUs. The sepa-ration between the fibres’ centre is only

1.47 times the fibre diameter core,which implies that the contaminationbetween adjacent fibres is about 10 percent. With the chosen arrangement ofthe fibres this is acceptable since innormal seeing conditions a higher levelof contamination will be present at thefibre entrance level.

3.4 ARGUS fibres

The ARGUS system is a fixed arrayof 14 × 22 micro-lenses, similar to theIFUs, located in the middle of Plate 2.ARGUS is also equipped with a lens toswitch between a scale of 0.52″/micro-lens to a finer scale of 0.3″/pixel. AnADC ensures that the object images atdifferent wavelengths are maintained inthe same locations up to a Zenithal dis-tance of 60 degrees. In addition to theobject fibres, fifteen ARGUS sky fibresare present on the plate.

4. GIRAFFE

GIRAFFE is a fully dioptric spectro-graph. The fibres are arranged in 5 longslits (IFU 1/2, MEDUSA 1/2, ARGUS) atthe curved focal plane of the collimator.One of these fibre slits is placed in theworking position by a translation stage.The other slits are masked. After leav-ing the fibres at F/5, the light first pass-es through one order sorting filter be-

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Figure 2: Transmission distribution of all GIRAFFE fibres, as meas-ured in the laboratory.

Figure 3: Transmission of order sorting filterHR 17.

fore it is directed to the collimator by a45-degree mirror. The collimated beamis dispersed by one of the two echellegratings that are mounted back to backon the grating turret. After passingthrough the collimator again an inter-mediate spectrum is produced. Finally,a re-imaging system (“camera”) pro-duces a 2.5 times demagnified imageon the 2K × 4K CCD at F/2. A view ofGIRAFFE is given in Figure 1. GI-RAFFE has been conceived to mini-mize maintenance and night calibra-tions; special requirements have beenintroduced to reduce set-up shifts andto obtain accurate re-positioning, to beable to use wavelength calibrations andflats obtained in the afternoon. To beable to monitor instrument drifts, five fi-bres are always used to obtain simulta-neous calibration spectra.

A summary of the most relevantGIRAFFE characteristics is given inTable 2

GIRAFFE is operated with 30 fixedset-ups; 22 are used to cover the wholespectral range with the high resolutiongrating, 8 with the low resolution grat-ing. The different GIRAFFE sub-unitsare described in more detail in the fol-lowing sections.

4.1 Slit Unit

The slit unit is the most complex GI-RAFFE mechanical subsystem, be-cause it needs a very high stability andreproducibility. It has two linear move-ments: one to exchange the fibre slit,the other one is used for focusing andto press the non-used slits against abaffle that is also used for back-illumi-

nation. The slit unit is equipped with anumber of back-illumination LEDs,which are switched on before pick-upand after placing a fibre, to allow thegripper camera to view the fibre output.

4.2 Filters

After the fibre slit, an order-sorting in-terference filter is placed in the beam.These filters must have excellent trans-mission and steep bandpass edgesand out-of-band blocking, in order toavoid pollution from adjacent spectralorders. One example of filter transmis-sion is given in Figure 3. The transmis-sions of all filters, averaged over eachset-up, are given in Figure 4. This figureshows also the efficiencies of other GI-RAFFE components, as well as theoverall GIRAFFE efficiency.

4.3 Optics and gratings

The gratings are placed back to backon a turntable. The high-dispersiongrating is a MgF2/Ag coated, 204 × 408mm 63.6 degree echelle with a largegroove density (316 lines/mm). The370–950 nm spectral range is coveredby twenty-two set-ups in grating orders15 to 6. Silver was selected because ofits good VIS/NIR efficiency and the lowpolarization it produces on this grating.While the efficiency on one edge of thegrating is good through the used spec-trum, the efficiency degrades below500 nm as one approaches the otheredge. This is due to inhomogeneity inthe MgF2 layer. We are working to finda replacement up to specifications.

The low-resolution grating has 600

lines/mm coated with MgF2/Al. Its first-order blaze lies at 1.96 µm, so in therange 370–950 nm it works in orders5–2; eight set-ups are needed to coverthe full range. After being dispersed,the light passes again through the colli-mator, forms a real image at an inter-mediate focal plane and is finally im-aged by an F/2 reimaging camera with7 elements.

4.4 Spectral Format and Efficiency

In GIRAFFE the spectra are parallelin dispersion along the long side of thedetector (readout direction), while theshort side is along the slit direction.Spectra are slightly curved due to opti-cal distortion and lateral chromatism,with the central part closer to each oth-er than the edges. As is the case in anylong-slit spectrograph, the lines of con-stant wavelength are arranged on low-curvature arcs. An example of a Th-Aspectrum taken with MEDUSA is givenin Figure 5.

The resolving power and coverageare both a function of the grating angleand grating order: two or three differentgrating angles are required to fully cov-er a grating order, which causes the dif-ferences in resolution and spectral cov-erage between the different set-ups.The higher resolving power of the IFUand ARGUS modes (compared to theMEDUSA mode) is due to the smallersize of the fibres, which projects to~ 2.2 pixels instead of ~ 4.3 pixels ofMEDUSA , but the spectral coveragefor a given set-up is the same in theMEDUSA, IFU and ARGUS modes.

5. Integration andCommissioning Results

FLAMES components were installedbetween October 2001 and April 2002on UT2. The integration sequence hasroughly followed the light path: Correc-tor, UVES fibres, Positioner, GIRAFFE.GIRAFFE had first light in Europe inDecember 2001 and was re-assembledin Paranal in April 2002 without majorinconveniences, after the re-integrationof the positioner and in perfect timingwith the rest of the project.

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Fibre slit height 76.8 mmType Echelle + Order selection FiltersCollimator aperture F/5Collimator beam 18 cmH-R Echelle 204 × 408 mm, 316 lines/mm, 63.4º BlazeL-R Echelle 156 × 204 mm, 600 lines/mm, 33.7º BlazeReimager demagnification ratio 2.5Effective Camera Focal Length 360 mmDetector 2048 × 4096, 15 µm EEV CCDScale 0.20 arcsec/pixel

Table 2. Summary of GIRAFFE construction characteristics.

Figure 4: Component efficiencies per set-up and the resulting overall DQE (slit to detected photoelectrons) of GIRAFFE. All reported valuesare averages; note that grating efficiencies can vary by up to a factor of two within a set-up. The optics efficiency includes vignetting at thegrating, this is why the reported LR values are slightly lower than HR values at the same wavelengths.

During the first run at the end ofMarch 2002 the Positioner was inte-grated and tested and we could gatherthe first light with UVES-Fibres. Three10-minute spectra were obtained for 7stars (one fibre was dedicated to sky)belonging to the Globular clusterOmega Cen. Three of them were previ-ously observed with UVES in long slitmode (Pancino et al. 2002). A smallportion of the spectra around the lineHα line is shown in Figure 6. The starshad magnitude varying between V =11.5 and 12.9. After data reduction theequivalent widths of a few hundredlines were compared for the UVES –Fibre and UVES-long slit spectra of thesame star, finding an excellent agree-ment: EQW(Fibre) = EQW(Slit) + A withA in the range of 4–6 mÅ, with an rmsof 4–7 mÅ, which is largely due toequivalent widths measurement uncer-tainties.

During this March run, malfunctionswere experienced in the positioner ro-bot, due to the lower temperature at thetelescope compared to those prevailingduring the tests in Australia and in theAssembly Hall in Paranal. We tem-porarily fixed this using three heaters,

clearly not a suitable permanent solu-tion. This was corrected prior to a sec-ond commissioning period in June butnew mechanical problems, related tothe interface with the telescope rotator-

adapter were encountered. Theseproblems were cured in advance of theAugust commissioning run. On July 3,first GIRAFFE spectra could be ob-tained. Two sets of spectra of stars inan astrometric field were acquired,each of 10 minutes. The set-up usedwas LR4, Low Resolution centred at awavelength close to the V filter. 66MEDUSA fibres were positioned (somein sky positions), and the reduced spec-tra of some of them are shown in Figure7. The stars have R magnitudes be-tween 13 and 14. In all spectra theMagnesium triplet at 518 nm is clearlyvisible, as well as the 557 sky emissionline.

The stability and repeatability of GI-RAFFE on the telescope, measured bytaking a ThAr spectrum every 4 hoursover a period of 2 weeks is shown inFigures 8 and 9.

The tests have been very satisfac-tory, showing that the instrument ex-ceeds specifications.

A further commissioning took place inAugust 2002, enabling us to advancethe commissioning tests. ManyMEDUSA ‘rasters’ have been per-formed, to determine the telescope-plate geometry.

In these rasters an astrometric field isobserved with MEDUSA and the tele-scope is offset along a grid around thepointing position. For every star, thedistance between its nominal position(the centre of the grid) and the positionsof its maximum flux is then computedand used to correct the astrometricmodel. The results have been so farvery encouraging, and the found resid-uals are shown in Figure 10 , with anRMS of 0.15 arcseconds, which is verygood when considering that this num-ber includes all sources of errors, e.g.,the accuracy of the stellar astrometry.In Figure 11 we show the image of 4

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Figure 5: Th-A spectrum taken with a MEDUSA slit. The fibres in the centre of the slit havelines moved towards smaller pixels; e.g. fibres corresponding to the central fibres have slight-ly redder wavelength coverage than the ones at the CCD edges.

Figure 6: Reduced spectra of 7 giants in the globular cluster Omega Cen taken with UVES-fibre around the Hα line.

Figure 7: Reducedspectra of GIRAFFEfrom the first com-missioning run: thestars belong to anastrometric field.The Mg triplet andthe 557 nm skyemission line areclearly visible in allspectra.

stars from the Tycho catalogue wellcentred in the FACBs after the tele-scope field acquisition has been com-pleted.

Two recurrent questions when usingfibres are the capability of obtaininggood flat field corrections, and to sub-tract the sky. Sky subtraction tests areunder way, while flat field correction ca-pabilities have been evaluated. In Fig-ure 12 is shown the comparison be-tween the measured S/N ratio in a por-

tion of the spectraof several stars inthe SMC, acquiredwith the L8 set-up,a spectral rangeplagued by strongfringing. We search-ed for a ‘line free’region in the spec-tra and the S/N ra-tio has been com-puted in this small(typically less than2 nm) spectral re-gion. In ordinatethe Poisson noise

based on the number of detected elec-trons is given. The agreement betweenthe measured and Poisson S/N ratio isquite good.

In the past, several groups have re-ported problems in obtaining accurate

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Figure 8: the average position of ThAr lines in a time series that wastaken at 4-h intervals over a period of 2 weeks. Shifts are given inpixels relative to the first exposure. These exposures were taken inMedusa HR mode, with configuration changes between exposures.The jump on May 7 cannot be correlated to any known event like anearthquake.

Figure 9: Histogram of relative shifts after4, 12 and 24 hours, using the time series ofFigure 12. In HR mode, 0.1 pixel corre-sponds to about 300 m/sec. These plotsshow the intrinsic stability of Giraffe and willbe helpful to plan the frequency of calibra-tions. Note that we expect to obtain higherradial velocity precision by using the simul-taneous calibration fibres.

Figure 10: Residuals vectors after the subtraction of the position cen-troid from the object fibre, after a successful modelling of the plategeometry.

Figure 11: After a successful acquisition iscompleted, the four fiducial stars should bewell centred in the FACBs, as in the case ofthese 4 stars from the Tycho Catalogue. Theobservations were made at airmass 1.8which shows the good handling of atmos-pheric effects by the positioner SW. The fig-ure shows the Graphic User Interface at theUT2 Console.

flat fielding while using fibres, because,if spatial filters are present in the opticalpath, variable fringing pattern can beproduced in the spectra.

By dividing 3 FF taken with the fibrepositioner by the mean of 36 FF takenwith the Nasmyth screen, a S/N ratio of400 was obtained. This S/N ratio ob-tained was less than what expected byPoisson noise statistics, but it showsthat this problem is limited in theFLAMES system to very high S/N ob-servations. Note that the 3 positionerFF represent the standard calibrationsthat the observer will receive. Similarlywe have tested the stability of photo-metric response of the fibres to torsionand movements: by taking Nasmyth FFand after rotating by a full circle, thephotometric stability has been shown tobe better than 0.5 %. Finally, in order toperform an accurate sky subtraction,the relative fibre transmission needs tobe known accurately. By comparing fi-bre to fibre transmission by using 3 po-sitioner flats (standard calibrations)with that obtained by the average of the36 Nasmyth flats, we find that the twomeasurements agree to better than 3%rms. This is a good result, although notyet within the specifications (2%).

We have performed several compar-isons with the ESO Exposure TimeCalculator (ETC, www.eso.org/observ-ing/etc). In Figure 13 we show the re-sults of the comparison between theflux observed vs. the flux predicted asobtained by observing an astrometricfield (the astrometric solution in use inthis test had a residual rms of 0.24 arc-seconds). The Figure shows the num-ber of counts recorded vs. R magni-tude, while the curve shows the predic-tion by the ETC in the L8 set-up for aG0 star and the parameters (seeing,exposure time) appropriate for the ob-servation. As shown in the bottom pan-el, at any given magnitude, the numberof collected e– per object is represent-ed by a relatively broad distribution.This is due to several factors, such asthe variation in the fibre transmission,

the poor astrometry of some targets,and the difference in spectral typeamong the stars in the catalogue. On

the other hand, it is important that thepotential users become familiar withthe concept that for FLAMES the S/N

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Figure 12: S/N ratio as measured in cleanregions of spectra taken in the L8 set-up,compared with the Poisson S/N ratio ex-pected on the number of counts recorded.

Figure 13: Comparison between observed flux and predicted ETC flux for an astrometric fieldwith the L8 set-up. The histogram shows the ratio measured/expected. The ETC calculationsare based on a G0 spectral type.

Figure 14: Distribution of the difference between RV single measurements and RV mean for3 exposures taken in 1 night of 60 stars of the globular cluster NGC 6809, shown vs. the stel-lar magnitude. Points from the different observations are shown with different colours. In thelower panel, the left distribution refers to all the measurements, irrespective of stellar magni-tude. The right distribution includes only the bright (V < 14.3) subsample.

ratio predicted by the ETC will alwayshave to be intended in a statisticalsense.

A number of tests have been per-formed to characterize the GIRAFFEradial velocity (RV) capabilities, observ-ing the same objects several time in anight, as well as in different nights.Three not consecutive observations of15 minutes in the GIRAFFE H9 set-upwere obtained in the same night for 96giants at the centre of the GlobularCluster NGC6809. The stars were se-lected between magnitude 13 and 15and the observations made with a nonoptimal astrometric model. Out of thefull sample (96 stars) we analysed only60 objects giving a well determinedcross correlation peak.

Out of these 11 were further exclud-ed because their RV were either dis-cordant from the bulk of the cluster, orbecause they showed a very large RVdifference in the 3 measurements (4stars). Based on the 3 measurementsof the 49 stars left, the cluster velocityis found at 178.81 km/sec (note thatthis RV is computed with respect to ourdigital mask, no attempt has beenmade to transform it to an absoluteRV), with a dispersion of 4.057 km/sec.In Figure 14 the distribution of the dif-ferences of all the measurements withrespect to their mean values areshown. This figure clearly shows howthe Radial Velocity accuracy degradeswith increasing stellar magnitude (andtherefore decreasing fluxes); the threeobservation points are indicated withdifferent colours. When taking all thepoints, irrespective of the stellar magni-tudes, the σ of the distribution is of 113m/sec, while if only the bright end isconsidered, V < 14.3, the σ decreases

Red Giant Branch of the SmallMagellanic Cloud; the magnitudes varybetween V = 18.5 and 18.8, the expo-sure time was of 1 hour with an 0.9 arc-second seeing.

Spectra were acquired for a field ofgalaxies with the Giraffe IFUs, and theirreduction is not yet completed. We aimat making a release of FLAMES com-missioning data in the first quarter of2003.

6. Acknowledgements

The development of FLAMES, whichwas recommended by the ESO STC in1998, presented some remarkable as-pects, showing the feasibility of a com-plex, multi-institute facility in a shorttime. This could not have been possiblewithout appropriate structures butmostly without the dedicated work of avery large number of people, whosenames do not appear in the author list.We thank all the persons at ESO

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to 70 m/sec. These numbers are in-dicative of the accuracy reached in thetime scale of a few hours and normaloperations. Royer et al. (2002) havereached a short-term accuracy of 13m/sec with GIRAFFE, in a more stableenvironment, using much higher S/Nobservations of the solar light in theGarching Integration Hall.

We had the opportunity to observethe same NGC 6809 field also with GI-RAFFE in low resolution mode (L4) andin Figure 15 a fraction of the spectra ofthe same stars (4 giants belonging tothe cluster) observed in L4 and H9 at aresolving power of R = 6000 and R =26,000 respectively are shown. The in-tegration time in the HR9 was 2 timeslonger than in the LR4 .

We are acquiring spectra to verify thebehaviour of the instruments at the faintend of the magnitude range; Figure 16shows a collection of Hα spectra takenwith the high resolution grating of GI-RAFFE (set-up H14) of giants on the

Figure15: GIRAFFE Low (L4, black lines) and High (H9, red lines)resolution spectra of 4 giants belonging to the Globular Cluster NGC6809

Figure 16: Hα Spectraof 8 red giants in theSMC, no sky subtractionwas performed, and skyHα in emission is visibleat 6562.8 nm. The starshave V magnitude be-tween 18.5 and 18.8.The exposure time was1 hour, in the GIRAFFEH14 setup. Fluxes are inLog scale.

(Garching and Paranal) and at the dif-ferent institutions who have contributedto the realization of this ambitious proj-ect: J. Alonso, P. Ballester, J.-L. Be-ckers, P. Biereichel, B. Buzzoni, C. Ca-vadore, N. Cretton, C. Cumani, B. Dela-bre, A.V. Kesteren, H. Kotzlowski, J.-L.Lizon, R.S. Moreau , W. Nees, R. Pal-sa, E, Pozna, (ESO), P. Barriga, R.Castillo, J. Spyromilio AND MANYOTHERS at Paranal, C. Evans, R.Haynes, U. Klauser, B. Hingley, D.Mayfield, S. Miziarski, R. Muller, W.Saunders, G. Schafer, K. Shortridge,(AAO), L. Jocou and the fibre team, T.Melse and the mechanics workshopteam, S. Baratchart, L. Chemin, H. Flo-res, D. Horville, J.M., Huet, F. Rigaud,F. Sayède (GEPI, OPM), F.R. Ferraro,

P. Molaro, R. Pallavicini, I. Porceddu(ITAL Consortium).

ReferencesAvila, G. Guinouard, I., Jocou, L., Guillon, F.,

Balsamo, F. 2002, in Instrument Designand Performance for Optical/InfraredGround-Based Telescopes Proc. SPIE4841 (in press).

Blecha, A., North, P., Cayatte, V., Royer, F., Si-mond, G. 2000, Proc. SPIE Vol. 4008, p. 467.

D’Odorico, S. et al. 2000, Proc. SPIE Vol.4005, p. 121.

Gillingham, P., Miziarski, S., Klauser, U.2000, Proc. SPIE Vol. 4008, p. 914.

Gillingham P.R., Popovic D., Waller L.G.,Farrell T.J. 2002, in Instrument Designand Performance for Optical / InfraredGround-Based Telescopes, Proc. SPIE4841 (in press).

Hammer, F., Hill, V., Cayatte, V. 1999, Jour-nal des Astronomes Français, 60, 19.

Jocou, l. et al. 2000, Proc. SPIE Vol. 4008,p. 475.

Mulas, G., Modigliani, A., Porceddu, I.,Damiani, F. 2002, in Instrument Designand Performances for Optical/InfraredGround-Based Telescopes, Proc. SPIE4841 (in press).

Pancino, E. et al. 2002: ApJ 568, L101. Pasquini, L. et al. 2000, Proc. SPIE Vol.

4008, p. 129. Pasquini , L. et al. 2002, in Instrument

Design and Performances for Optical/Infrared Ground-Based Telescopes, Proc.SPIE 4841 (in press).

Royer, F., Blecha, A., North, P., Simond, G.,Baratchart, S., Cayatte, V., Chemin, L.,Palsa, R. 2002, in Astronomical DataAnalysis II, Proc SPIE 4847 (in press).

9

HARPS: ESO’s Coming Planet SearcherChasing Exoplanets with the La Silla 3.6-m Telescope

F. PEPE1, M. MAYOR1, G. RUPPRECHT3

email: Francesco.Pepe@obs.unige.ch; Michel.Mayor@obs.unige.ch; grupprec@eso.org

with the collaboration of the HARPS Team: G. AVILA3, P. BALLESTER3, J.-L. BECKERS3, W. BENZ5,J.-L. BERTAUX6, F. BOUCHY1, B. BUZZONI3, C. CAVADORE3, S. DEIRIES3, H. DEKKER3, B. DELABRE3,S. D’ODORICO 3, W. ECKERT 2, J. FISCHER5, M. FLEURY1, M. GEORGE 1, A. GILLIOTTE 2, D. GOJAK 2,3,J.-C. GUZMAN2, F. KOCH3, D. KOHLER4, H. KOTZLOWSKI, D. LACROIX3, J. LE MERRER10, J.-L. LIZON3,G. LO CURTO2, A. LONGINOTTI3, D. MEGEVAND1, L. PASQUINI3, P. PETITPAS4, M. PICHARD1, D. QUELOZ1,J. REYES3, P. RICHAUD4, J.-P. SIVAN 4, D. SOSNOWSKA1, R. SOTO 2, S. UDRY1, E. URETA2, A. VANKESTEREN 3, L. WEBER1, U. WEILENMANN 2, A. WICENEC 3, G. WIELAND 3 and advice by the InstrumentScience Team: J. CHRISTENSEN-DALSGAARD7, D. DRAVINS8, A. HATZES9, M. KÜRSTER9, F. PARESCE3,A. PENNY11

1Observatoire de Genève; 2ESO, La Silla; 3ESO, Garching; 4Observatoire de Haute Provence; 5Physikalisches Institut, Bern; 6Service d’Aéronomie du CNRS; 7Aarhus University, 8Lund Observatory,9Tautenburger Landessternwarte; 10Laboratoire d’Astrophysique de Marseille; 11Rutherford Appleton Laboratory

Introduction

An extensive review of past, presentand future research on extrasolar plan-ets is given in the article “ExtrasolarPlanets” by N. Santos et al. in the pres-ent issue of The Messenger. Here wewant to mention only that the search forextrasolar planets and the interpreta-tion of the scientific results haveevolved in recent years into one of themost exciting and dynamic researchtopics in modern astronomy.

As a consequence, ESO decided tohave a dedicated spectrograph built tobe installed at the La Silla 3.6-m tele-scope. An Announcement of Oppor-tunity was issued in 1998 and a con-sortium chosen for the realization of theproject which consists of Observatoirede Haute Provence (F), Physikalisches

Institut der Universität Bern (CH) andService d’Aéronomie du CNRS (F) un-der the leadership of the Observatoirede Genève (CH) where the PrincipalInvestigator and the Project Office arelocated. Further contributions comefrom ESO-La Silla and ESO Head-quarters (Garching). The agreementwas finally signed in August 2000.

According to this agreement, theHARPS Consortium bears the cost forthe spectrograph and all its compo-nents whereas ESO provides theCassegrain Fibre Adapter, the fibre link,the dedicated HARPS room in the tele-scope building and the complete detec-tor system. In return, the HARPSConsortium will be granted 100 HARPSobserving nights per year for a period of5 years after successful ProvisionalAcceptance in Chile. HARPS will of

course also be offered to the astronom-ical community like any other ESO in-strument.

The Radial Velocity Method:Error Sources

A description of the radial velocitymethod for the detection of extrasolarplanets in general and of the CORALIEspectrograph installed at the SwissLeonhard Euler Telescope at La Silla inparticular was given by Queloz & Mayor(2001); it describes the technique ap-plied also with HARPS and the spec-trograph which is its direct ancestor. Aninstrument like CORALIE achieves anaccuracy in radial velocity determina-tion of about 3 m/s, and with this accu-racy planets down to a minimum massof about one third the mass of Saturn

have been discovered. For HARPS anaccuracy of 1 m/s RMS is specified. Toachieve this goal, we have to analysethe limitations that are encountered bya spectrograph in order to overcomethem.

Photon noise

The contribution of photon noise tothe radial-velocity (RV) measurementerror can be expressed, according toHatzes & Cochran (1992), by thegeneric formula

σ RV ∝ S –0.5 ·λλ –0.5 ·R –1.5. (1)

The measurement precision is pro-portional to the square root of the flux Sand the wavelength range λλ of thespectrograph. It should be mentionedthat in practice it is not the wavelengthrange that determines the precision,but rather the number of spectral linesand their depth, since the informationon RV is contained in the stellar ab-sorption lines. Nevertheless, the formu-la tells us that we must maximize theamount of light entering the spectro-graph. This goal can be reached by in-creasing the spectral range, maximiz-ing the optical efficiency of the instru-ment, and, of course, by using thelargest telescope available.

The UVES-like optical design as wellas the choice of the optical componentsmake HARPS an efficient spectro-graph. However, this is not sufficient toget a precise RV spectrograph, since,

again according to formula (1), higherprecision is obtained when the spectralresolution R is increased. Unfortu-nately, telescope diameter D, spectralresolution R, and fibre diameter φ (pro-jected on the sky) are not independent,but linked by the formula

, (2)

where h is the diameter of the collimat-ed beam on the echelle grating and tanβ is the tangent of the echelle gratingangle. In other words, if we want to in-crease the spectral resolution we haveto increase the collimated beam diame-ter, or restrict the field of view φ andconsequently decrease the slit effi-ciency.

A trade-off therefore had to be madebetween spectrograph size, spectralresolution, and slit efficiency. Our solu-tion consists in building an instrumentusing the largest monolithic echellegrating available to date (a 837 × 208mm grating developed for UVES) andto balance spectral resolution againstslit efficiency. The compromise led to aspectral resolution of R = 93,000 whichin most cases is sufficient to resolve thestellar absorption lines. Going beyondwould not reduce the photon noise butonly increase slit losses. At the select-ed resolution the fibre covers 1 arcsecon the sky. The resulting slit losses areabout 50% on average. It should be re-called however that the goal was to op-timize the instrument for RV precision,which is proportional to R1.5 but only

proportional to thesquare root of S.Thus, a fibre twiceas large wouldcollect twice theamount of light, butwould also reducethe spectral resolu-tion by a factor of 2.The resulting RVprecision would bereduced by a fac-tor 2.

Instrumental errors

If we want to de-tect RV variationsof 1 ms–1 we mustbe able to detectchanges in the stel-lar absorption lineposition which are

smaller than 1/1000 of a CCD pixel. Atemperature variation by 1°C or thechange of atmospheric pressure by1 mbar would already produce effectsof the order of 100 ms–1! We havetherefore decided to act on differentfronts:

• Stabilize the spectrograph environ-ment. HARPS is installed in vacuum toremove effects caused by varying am-bient pressure. In addition, the temper-ature of the whole instrument is active-ly controlled.

• Simplify the opto-mechanical de-sign, avoid any moving or movablecomponents, design a robust spectro-graph.

• Use a spectral reference to trackresidual instrumental drifts.

Since it is impossible to avoid drifts of10 nm on the instrument, the spectralreference is mandatory. On HARPS wewill have the possibility of using either aThAr spectral lamp in simultaneous ref-erence mode, or an iodine absorptioncell in self-calibrating mode. Both tech-niques attain to date similar perform-ance at the level of about 3 ms–1. Themain mode of HARPS will be the ThArsimultaneous reference, since it allowsto cover a spectral range 3 times largerthan with the iodine cell, which in addi-tion absorbs about 50% of the light. Theefficiency is thus 6 times higher usingthe ThAr simultaneous referencemode.

Often, stability of the instrumentalprofile (IP) is associated with the RVprecision of an instrument. In practice,both spectral reference methods allowto avoid errors produced by internal IPvariations. It is however true that theThAr simultaneous reference techniquedoes not account for variations of theslit illumination. This technique workstherefore only if combined with a fibrefeed and with an image scrambler,whose sole function is to “scramble” thelight entering the fibre at the telescope.The goal is to avoid a change in spec-trograph illumination when the positionof the star on the fibre entrancechanges. We estimate that the residualeffects are of the order of 0.25 ms–1 andcan thus be neglected. It is worth men-tioning that good telescope guiding andhigh spectral resolution reduce thesekinds of error sources even further.

External errors

We should not forget that RV meas-urements may be influenced by errorsources related neither to the instru-ment nor to the telescope. Any of thosecan produce RV errors of the order ofseveral ms–1 if not correctly accountedfor. Without going into detail, we men-tion the following:

• Atmospheric absorption lines whichvary in relative position and intensitycompared to the stellar spectrum

• Atmospheric dispersion

10

Figure 1: Thefinished HARPSCassegrain FibreAdapter ready to beshipped from La Sillato Geneva forsystem tests.

• Sunlight reflected by the moon andsuperimposed on the stellar spectrum

• Stellar companions close to the tar-get star may contaminate its spectrum,depending on seeing conditions

• Stellar jitter, intrinsic variability ofthe star

HARPS: The Implementation

The La Silla 3.6-m telescope

As already mentioned, ESO decidedearly on to install HARPS at the La Silla3.6-m telescope. This telescope recent-ly underwent a major overhaul and wasretrofitted with a new control systembased on the VLT software and dataflow. These improvements, togetherwith the obvious fact that a 4-m-classtelescope is ideal in providing the largenumber of photons required for highsignal-to-noise ratio spectra, make thistelescope a superb complement to ournew instrument.

The HARPS Cassegrain FibreAdapter (HCFA)

To connect the fibre link (see below)to the telescope, the La Silla Engi-neering Department designed andmanufactured a new fibre adapter forthe Cassegrain focus. The HCFA fulfilsa number of functions: It allows the re-motely controlled exchange of the fi-bres for HARPS and the CES and pro-vides both fibre feeds with an atmos-pheric dispersion compensator (ADC)and the possibility to use the tele-scope’s guide camera for guiding onthe respective fibre entrance. ForHARPS, there is also a neutral densityfilter and a feed for the calibration fibrewhich carries the light from a separatecalibration unit (see below). This is acrucial part of the calibration concept ofHARPS (described in Queloz & Mayor2001). As a unique feature, HARPS willoffer the observer two options for pre-cise wavelength calibration: the defaultThorium-Argon method and the use ofthe Iodine absorption cell. This Iodinecell is also mounted in the HCFA andcan be moved in and out of the tele-scope beam under remote control.Figure 1 shows the finished HCFA.

Calibration unit

The HARPS calibration unit providesthe instrument with light for wavelengthand flatfield calibration. For this pur-pose it contains a set of hollow-cathodeThorium-Argon and halogen lampswhich can be remotely switched on andoff. A motorized exchange mechanismallows to position the calibration fibre infront of any desired lamp. The calibra-tion fibre pair connects the calibrationunit, which is located next to the air-conditioned HARPS enclosure in thecoudé west room, with the fibre adapterat the Cassegrain focus.

Fibre links, image scramblers

Strictly speaking HARPS is a distrib-uted system, and one of the most im-portant components connecting its var-ious parts is the observation fibre link.Its purpose is to feed the spectrographdown in the telescope building with (a)the star light collected at the Casse-grain focus and with (b) either the ThArspectrum for the simultaneous calibra-tion or with light from the night sky forbetter sky subtraction. For HARPS wechose two 70µm fibres (type FVP madeby Polymicro), corresponding to 1 arc-second on the sky, and put it in a“shower tube” made of sturdy steelmesh for mechanical protection. Thetotal length of the fibres from theCassegrain focus to the spectrographentrance (by way of the declinationbearing, telescope fork, northern tele-scope mount bearing and into thecoudé west room) is 38 metres.

To minimize focal ratio degradation,the light is coupled into the object andreference fibres by means of two microlens doublets per fibre. By projectingthe image on the fibre input end, the tel-escope pupil is at infinity. This designcombines an excellent image qualitywith easy, uncritical alignment.

A double image scrambler is locatedat the entrance of the object/referencefibres into the vacuum vessel of thespectrograph. In combination with thefibre feed, it serves to stabilize thespectrograph illumination: the objectmay move at the fibre entrance due toguiding errors or seeing, but the in-tensity distribution at the fibre exit, i.e.the spectrograph entrance, will notchange. In addition the scramblerserves as the feed-through for thefibres into the vacuum and it also

houses, on the atmosphere side, theexposure shutter.

The light is finally led to the spectro-graph entrance inside the vacuum ves-sel by means of two short (2-metre)pieces of fibre. The coupling to thespectrograph is again achieved by apair of doublet microlenses per fibre.

The second fibre link leads, as men-tioned before, from the calibration unitin the coudé west room up to theCassegrain focus. It is therefore alsoabout 38 metres long, consisting of apair of 300µm core diameter fibres(type FVP made by Polymicro).

Vacuum vessel with spectrograph

The vacuum vessel has the purposeof protecting the spectrograph properfrom temperature variations and fromthe effects of refractive index variationsof air. This vessel has a volume of ap-proximately 2 m3. It is evacuated bymeans of a turbo molecular pump be-fore the start of operations; we expectto repeat the regeneration of the vacu-um about once or twice per month.

Since the long-term stability of thespectrograph is of paramount impor-tance for the success of the exoplanetsearch, the vacuum should be brokenas seldom as possible. For this reasonthere are no moving functions insidethe vacuum except the focussingmechanism of the camera. This willhowever be adjusted and locked beforethe vessel is finally closed. Figure 2shows the closed vessel in the Genevaintegration hall.

The spectrograph itself is a cross dis-persed echelle spectrograph, very sim-ilar to UVES at the VLT. It is a whitepupil design with the grism cross dis-perser placed in the white pupil. The

11

Figure 2: The closed HARPS vacuum vessel in the Geneva integration laboratory.

echelle grating, a copy of the UVESmosaic, is operated in quasi-Littrow con-dition. An f/2.1 parabolic mirror servesas collimator and is used in triple pass.A dioptric camera images the cross-dis-persed spectra (one each from the ob-ject and reference fibres) side by sideonto a mosaic of two 2k × 4k EEVCCDs. 68 orders cover a spectral rangefrom 380–690 nm. At a spectral resolu-tion RS of 90,000, determined by the fi-bre diameter, a spectral resolution ele-ment is sampled by 4 pixels. All opticalcomponents are mounted on a stain-less steel optical bench. The optical pa-rameters are listed in Table 1.

Procurement of the optics turned outto be the most demanding of all. Thefirst company charged with the produc-tion of the cross disperser grism failedto deliver and a new contract had to benegotiated with another supplier. Thecollimator mirror also failed to material-ize, both within schedule and with ac-ceptable quality. The mirror which wasfinally delivered had to be re-polishedby another company to meet our re-quirements.

In order to improve the observing ef-ficiency by always applying the correctexposure time (according to the select-ed signal-to-noise ratio SNR) we fittedan exposure meter, again following theexample of UVES. Two photon coun-ters are used to separately measurethe light coming from the object and ref-

erence fibres which is reflected off thegap on the echelle between the twogratings comprising the mosaic.

Detector system

HARPS employs a mosaic of twoEEV type 44-82 CCDs (nicknamedJasmin and Linda). The spectral format

is thus 40962 pixels (15 µm square) ofwhich a field of 62.7 × 61.4 mm is actu-ally used at a sampling of 4 pixels perspectral element. The performance ofthe chips is summarized in Table 2.

As HARPS is a stationary instrument,its detectors are cooled by a continuous-flow cryostat (CFC) of the current ESOstandard design. A special feature of theHARPS cryostat is however the fact thatthe detector head (visible on the left inFig. 4) is mounted to the spectrographbench inside the vacuum vessel, with theactual CFC outside. Both are connect-ed by a stainless steel bellows whichprotects the detector high vacuum(10–6 mbar) from the mere “emptiness”(10–2 mbar) of the spectrograph vessel.The detector head window also servesas field lens of the camera optics.

As a consequence of this configura-tion, all cables connecting the detectorhead to the FIERA controller have topass through vacuum feed troughs.

Thermal enclosure

In order to keep the spectrographtemperature as constant as possible itwas decided to put the vacuum vesselin an additional thermal enclosure. Thisis a well insulated room on the coudéfloor of the 3.6-m telescope buildingwhich is itself already temperature sta-bilized. Based on preliminary measure-ments we expect to keep the tempera-ture variations of the spectrograph

12

Figure 3: The open HARPS vacuum vessel. The back side of the echelle grating is visible ontop of the optical bench as well as the rail system used for moving the large parts of the vessel.

System 2 fibres (1″ dia.), spectral range 380–690 nm, collimated beam diameter 208 mmEchelle R4, 31.6 gr/mm, blaze angle 75°, mosaic 2 ×1 on Zerodur monolith 840 × 214 × 125 mm, efficiency > 65% in the visibleCross disperser FK5 grism, 257.17 gr/mm blazed at 480 nm, 240 × 230 × 50mm, T = 73% (av)Collimator Zerodur, f = 1560 mm, used diameter 730 mm, triple passCamera All dioptric, 6 elements in 6 groups, f = 728 mm, f/3.3, T > 85%

Table 1 Summary of the HARPS spectrograph optics.

Figure 4: The figure shows the CCD detector system mounted below the spectrograph bench.Visible are also the field lens and a holder which still connects the head and the CFC. In thefinal configuration the CFC will be mounted outside the vacuum vessel.

bench inside the vacuum vessel (theonly place that really matters) below0.1K, ideally at 0.01K!

Performance

The instrument

Recently we have completed the in-tegration of HARPS at the GenevaObservatory. Tests of the completespectrograph have been carried out,and the spectroscopic characteristicsverified. The spectral format recordedcorresponds well to the calculated val-ues. All the orders – except order 115which falls between the two CCDs –could be localized and extracted usinga tungsten flat-field lamp (Fig. 5), whilethe wavelength calibration was doneusing the ThAr spectral lamp. In Figure6 we show a small segment of the ex-tracted and wavelength-calibrated ThArspectrum. About 3500 spectral linescovering a large intensity range couldbe detected and identified. From thespectrum shown, which was recordedduring an exposure of 7 s, we estimatethat the internal photon noise on the ra-dial velocity is only 10 cm s–1. This rep-resents a large gain compared toCORALIE, in part due to the higheroptical efficiency of the spectrograph,but mainly due to the higher spectralresolution. Indeed, the Gaussian fit of asingle ThAr line indicates a spectralresolution of R = 98,000, which is about5% higher than expected.

As mentioned above, HARPS isequipped with an iodine cell for the self-calibration mode. We have recorded aspectrum of our iodine cell by sendingthe white light of the tungsten lampthrough it. The extracted and wave-length-calibrated spectrum is shownin Figure 7. The intensity of the con-tinuum varies strongly across theechelle order because of the blaze re-

sponse of the grating. The angle of theechelle grating was not centred correct-ly at that time thus delivering an asym-metric blaze response. This was cor-rected later by tilting the echelle gratingphysically by 0.25°.

The total efficiency of the instrumenthas not yet been determined. Thismeasurement will be done duringCommissioning using a reference star.Nevertheless, the optical componentswere measured individually. Table 3summarizes the expected optical effi-ciency of HARPS derived from thesemeasurements. From this, and basedon data recorded with the CORALIE in-strument at La Silla (see Queloz &

Mayor (2001)), we have estimated theRV precision of HARPS as a function ofstellar magnitude (see their Fig. 14).The total RV efficiency of HARPS isthus about 75 times higher than that ofCORALIE. This extraordinary RV effi-ciency will offer exciting new possibili-ties in other research fields, for exam-ple in asteroseismology.

Operations

Often neglected, observational effi-ciency also plays an important role. Infact, searching for exoplanets requiresthat many measurements per objectare carried out. Telescope time musttherefore be used efficiently.

From the beginning, strong emphasiswas put on efficient operations scenar-ios to optimize the yield from an ob-serving run. Most important is certainlythe adoption of the VLT Data Flow Sys-tem, from the preparation of Observa-tion Blocks with P2PP through the useof BOB at the telescope to the imple-

13

DQE 82–85% peak, 70% minimumRead-out noise ~ 3e¯ at 50 kpx/sec (2 ports/chip)CTE 0.999999 (horizontal+vertical) at 166 kpx/secCosmetics Science grade (grade 1)Read-out modes 50 kpx/sec and 625 kpx/sec, one or two ports per chip

Table 2: Properties of the HARPS detector system.

Wavelength 380 nm 400 nm 450 nm 500 nm 550 nm 600 nm 650 nm 690 nm

Tel. + atm. 44% 47% 54% 57% 59% 59% 61% 63%“Slit” 46% 47% 48% 49% 50% 50% 51% 52%Instrument 8% 11% 15% 15% 17% 17% 15% 12%CCD 65% 78% 85% 85% 81% 79% 76% 72%

Total 1.0% 1.9% 3.2% 3.5% 4.0% 3.9% 3.4% 2.9%

Table 3: Optical efficiency of HARPS.

Figure 5: Spectral flat-field of HARPS usingthe tungsten lamp. Only one fibre (object) isilluminated for the localization of the orders.

mentation of a full online data reductionpipeline. A sophisticated software toolwill support the observer during the run:it will allow online calculation of the ex-posure time (based on the brightness ofthe star, the actual brightness of thenight sky, the seeing and the desired

SNR) and enable him to optimize dur-ing the night the short-term schedulingof observations of targets in a pre-pre-pared object catalogue.

All raw data are then stored in the lo-cal archive. In addition, observations inthe high-precision ThAr mode are auto-

matically reduced in near real time by apipeline running on a dedicated work-station: at the end of the night the ob-server will have available all final radialvelocities, ready for publication! All datacan be immediately written to DVD-ROM and are also archived in the cen-tral Garching science archive.

Outlook

The HARPS project is now close tocompletion. Following the initial An-nouncement of Opportunity in 1998 werethe project kick-off in February 2000,optics Final Design Review (FDR) inApril 2000, Preliminary Design Reviewin July 2000, FDR in March 2001, andsoftware FDR in July 2001. HARPS iscurrently undergoing extended systemtests at Geneva Observatory (see Pepeet al. 2002). As soon as these are fin-ished, we will perform PreliminaryAcceptance (Europe), scheduled forDecember 2002. After the transport toLa Silla and the subsequent re-integra-tion and functional tests we plan tohave First Light in February 2003, fol-lowed by two commissioning runs of 2weeks duration each. HARPS shouldthen be available to the ESO astro-nomical community from Period 72 on.At the specified performance, it will bea marvellous radial velocity machinewhich will among others serve tostrengthen the leading role Europeanastronomers play in the new and excit-ing field of extrasolar planet research.

Acknowledgements

Like many major instrument projectsthese days, HARPS would not have be-come a reality without the enthusiasticeffort of a large number of people inseveral institutes on different conti-nents. The main players in the HARPSTeam are listed at the beginning of thisarticle.

Industry played a further crucial rolein the project by providing key compo-nents. The major contractors are listedin Table 4.

The HARPS Consortium acknowl-edges financial support from variousfoundations and institutions in Switzer-land and France.

References

Hatzes, A., Cochran, W. in ProceedingsESO Workshop on High-ResolutionSpectroscopy with the VLT, M. Ulrich, ed.,ESO, (ESO Garching), 1992

Santos, N. et al.: this Messenger.Pepe, F. et al., “Performance Verification of

HARPS – First Laboratory Results”, inAstronomical Telescopes and Instrumen-tation 2002: Ground-based Telescopesand Instrumentation, Proc. SPIE Vol. 4841,(2002, in print).

Queloz, D., Mayor, M. 2001, From CORALIEto HARPS, The Messenger, 105, 1–7.

14

Figure 6: Segment of the extracted ThAr spectrum. In the small window one single line isshown. The FWHM of the superimposed Gaussian (dashed line) corresponds to a spectralresolution of R = 98,000.

Figure 7: Portion of the extracted and wavelength-calibrated iodine spectrum. The intensityvariation across the echelle order is caused by the blaze response of the grating.

Vacuum vessel APCO Technology (CH)Optical bench Deshors Ets. (F), Acrodur (F)Crossdisperser grism Cybernetix (F), Thermo RGL (USA)Echelle grating Thermo RGL (USA)Camera, collimator SESO (F), SAGEM (F)Coatings SAGEM (F), Gretag TFP (CH)Optical fibres SEDI (F)Iodine cell Hellma (D), Physikalisch-Technische Bundesanstalt (D)

Table 4. The main HARPS industrial contractors.

1. Introduction

In 2004, OmegaCAM will start op-erations on Paranal as the sole in-strument on the 2.6-m VLT SurveyTelescope. OmegaCAM is a huge opti-cal CCD imaging camera: its 16k × 16kCCD pixels cover the square degreefield of view of the VST almost entirely.The primary function of the VST and itsinstrument is to provide surveys in sup-port of VLT science, be it in the form oflarge homogeneous multi-colour imagingsurveys which form the basis for large-scale spectroscopic follow-up work, orin its ability to find rare or extreme as-tronomical objects for further study.

The designs of both VST andOmegaCAM try to take full advantageof natural good seeing, so it should alsobe a superb instrument for weak gravi-tational lensing surveys, or for monitor-ing projects designed to detect micro-

lensing or supernovae. In fact, applica-tions are manifold: one has only to lookat the exciting science that is now com-ing out of the Sloan Digitized SkySurvey to realize the potential of VST/OmegaCAM, which has a comparablefield of view to the SDSS camera, butwill operate continuously, with betterimage quality and higher throughput.

The scale of the instrument meansthat once operations start the challengeis not at all over: OmegaCAM will gen-erate of order 50 GByte of raw data pernight, year after year, and such a vol-ume of data can only be digested bymeans of a strict observing protocol(encoded in the Observation Blocks)combined with highly automated pro-cessing of the data. Exciting and chal-lenging times are ahead!

In this article, the OmegaCAM con-sortium presents the basic features anddesign of the instrument.

2. The VLT Survey Telescope

The VST (Arnaboldi et al. 1998), nowunder construction in Naples, is a 2.6-m modified Ritchey-Crétien telescopewhich will stand next to the four UT’s onParanal. It is specifically designed forwide-field imaging, and has been opti-mized for excellent image quality in nat-ural seeing. Thus, it will have active pri-mary and secondary mirrors, a re-tractable atmospheric dispersion cor-rector, a constant focal plane scale of0.21arcsec per 15 µm pixel over a 1.4degree diameter field, and a theoreticalPSF with 80% of its energy in a 2 × 2pixel area over the whole field.OmegaCAM will be the sole instrumenton the telescope, and will be mountedat the Cassegrain focus.

3. Overview of the Instrument

3.1 Detector system

The heart of OmegaCAM is the CCDmosaic (Fig. 1), being built at ESOheadquarters in Garching. It consists ofa ‘science array’ of 32 thinned, low-noise (5e–) 3-edge buttable 2 × 4kMarconi (now E2V) 44-82 devices, for atotal area of 16384 × 16384 15 µm pix-els (26 × 26 cm!). The science array fitssnugly into the fully corrected field ofview in the focal plane of the VST, andcovers an area of 1 × 1 degree at 0.21arcsec/pixel. Around this science arraylie four ‘auxiliary CCDs’, of the sameformat. Two of these are used for auto-guiding (on opposite sides of the field:the field is so large that also field rota-tion will be auto-guided), and the othertwo for on-line image analysis. For thispurpose the latter CCDs are deliberate-ly mounted out of focus (one 2 mm infront, one 2 mm behind the focal plane),and the resulting defocused imagescan be analysed on-line and used toinfer aberration coefficients such asdefocus, coma, or astigmatism every

15

OmegaCAM: the 16k ××16k CCD Camera for the VLT Survey Telescope K. KUIJKEN1,2, R. BENDER 3, E. CAPPELLARO 4, B. MUSCHIELOK 3, A. BARUFFOLO 5,E. CASCONE4, O. IWERT6, W. MITSCH 3, H. NICKLAS7, E.A. VALENTIJN 2, D. BAADE 6,K.G. BEGEMAN 2, A. BORTOLUSSI5, D. BOXHOORN 2, F. CHRISTEN 2,6, E.R. DEUL1,C. GEIMER6, L. GREGGIO5, R. HARKE7, R. HÄFNER3, G. HESS6, H.-J. HESS3, U. HOPP3,I. ILIJEVSKI3, G. KLINK8, H. KRAVCAR3, J. L. LIZON6, C. E. MAGAGNA5, PH. MÜLLER9,R. NIEMECZEK6, L. DE PIZZOL5, H. POSCHMANN8, K. REIF8, R. RENGELINK1, J. REYES 6,A. SILBER6, W. WELLEM7

1Leiden Observatory; 2NOVA/Kapteyn Astronomical Institute, Groningen; 3Universitäts-Sternwarte München; 4INAF - Osservatorio Astronomico di Capodimonte, Napoli; 5INAF – Osservatorio Astronomico di Padova; 6ESO, Garching; 7Universitäts-Sternwarte Göttingen; 8Sternwarte Bonn; 9Radioastronomisches Institut Bonn

Figure 1: Layout of CCDs in the focal plane. This arrangement minimizes the amount of deadspace between devices, given the constraints imposed by connecting the read-out ports. Theglobular cluster ω Cen is superimposed on the field, which covers a 1 × 1 degree area. Theauxiliary CCDs, shown in green and purple, are used for autoguiding and for online wavefrontanalysis.

Figure 2: General drawing of the OmegaCAM instrument. The view is dominated by the cryo-genic cooling system (with pre-amplifiers, vacuum equipment, etc. attached on the outside)and the detector head on top of it containing the CCD mosaic. Above the CCDs, obscuringthe curved dewar window, is the filter exchange system (shown in blue), which moves filtersbetween the two magazines and the beam. The exposure shutter is mounted at the top of theinstrument, just below the telescope flange.

minute. The whole detector system ismounted behind a large, curved dewarwindow (the final optical element in theVST design) and is cooled using a 40-lNitrogen cryostat. Readout of the fullmosaic takes 45 s, and is accomplishedby two FIERA controllers (a thirdFIERA takes care of the four guidingand image analysis CCDs).

The OmegaCAM detector team atESO is led by O. Iwert.

3.2 HardwareIn front of the dewar window is the

mechanical part of OmegaCAM: clos-est to the CCD window sits the filter ex-change mechanism, and above that theshutter. Both components have to fitinto a design space of a mere 16 cmbetween the dewar window and theVST’s Shack-Hartmann unit. The hous-ing provides the mechanical link be-tween the telescope flange and the de-tector/cryostat system.

Figure 2 gives a section view of thefinal design that foresees a cylindricalhousing with a spoke-like rib structureto support the axisymmetrical loads atthe Cassegrain focus. The housing canbe seen in Figure 3.

The filters are stored in two maga-zines which can move up and down,either side of the focal plane, throughlarge shafts in the housing. A linearstage slides filters into the beam, wherethey are locked into place by means ofmovable notches. High-precision filterpositioning ensures that intensity varia-tions in the flat fields due to optical im-perfections in the filters (dust grains,etc.) are less than 0.1%. The filter ex-change unit is built in such a way that itallows one filter to be pulled into thebeam while the previous one is pushed

the opening edge of the one blade andthe closing edge of the other blade withan identical time difference, even if theblades are still accelerating – this pro-vides an impact-free, high-accuracyphotometric shutter. Tests of the shutterconfirm that it meets the key technicalspecification: for an exposure time asshort as 1 second, deviations from ahomogeneous exposure are well below±0.2% over the whole field of view.

The OmegaCAM control electronicsare based on VME-based local ControlUnits (LCUs), with a higher level ofUnix-based workstations to managethe user interface, coordination, testingand maintenance. The LCU is a stand-alone VME crate equipped with aMotorola CPU board, an Ethernetboard, the real time operating systemVxWorks, as well as specialized con-trol and interface boards. All the con-trolled functions are standardized asmuch as possible, and the modular de-sign facilitates the maintenance andshould ensure efficient and reliable op-erations.

The mechanics are designed andbuilt by the German consortium part-ners in Göttingen (H. Nicklas – housing,filter exchanger) and Bonn (K. Reif –shutter). More details can be found inNicklas et al. (2002) and Reif et al.(2002). The control electronics weredesigned at INAF-Naples by E.Cascone, and are being assembled atMunich University Observatory underW. Mitsch.

3.3 Optics

The VST telescope will work in twoconfigurations, which can be selectedremotely. In the standard configuration,foreseen for work at small zenith dis-tances, a two-lens field corrector is

16

Figure 3: The1.5-m diameterhousing struc-ture during atest assembly ofthe main units:the two storagemagazines (tobe inserted intothe big shaft atleft and hiddenright), the filterexchange unitwith theexchange car-riage and a 420× 320 mm2

opaque,aluminium ‘filter’(at the bottomleft).

out, allowing efficient observing in spiteof the rather large distance the filtershave to travel.

The filters are large (in the languageof our latest ESO member state: asquare foot) and heavy: when fully load-ed with 12 filters, the instrument will con-tain 40 kg (90 lbs.) of filter glass alone!

The exposure shutter (Reif et al.2002) is one of the key units ofOmegaCAM (Fig. 4). It consists of twocarbon fibre blades which open andclose the light path. They are driven bymicro-stepper motors and movesmoothly on linear motion guides.These movements are controlled suchthat each individual CCD pixel ‘sees’

used. The second configuration re-places this corrector with one includingan Atmospheric Dispersion Corrector(ADC), consisting of one lens and twocounter-rotating prism pairs. The oper-ating wavelength ranges are 320–1014nm and 365–1014 nm for the two-lenscorrector and corrector + ADC respec-tively.

The only optical parts located in theinstrument are the filters, and the en-trance window to the cryostat, whichdoubles as a field lens.

The primary filter set of OmegaCAMwill be a set of Sloan u′, g′, r′, i′ and z′filters. In addition, there will be JohnsonB and V filters for stellar work and forcross-calibrating the photometric sys-tems, a Strömgren v filter, an Hα filterconsisting of 4 segments with redshiftsof up to 10,000 km/sec, and a seg-mented ugri filter for efficient photomet-ric monitoring of the sky.

The procurement of large format fil-ters of the required size turned into achallenging task. Only one manufactur-er (the French company SAGEM, for-merly REOSC, who figured the VLT pri-mary mirrors) could make an offer forproducing the primary set of filters with-out resorting to a segmented design,which would have created vignettingshadows on the detector array. Ratherthan using coloured glass – barelyavailable in the required size – the filterpassband is generated by means ofmultiple layer coating of up to 5 sur-faces in a sandwich of three plates. Theexpected throughputs of the Sloan fil-ters are very high (Fig. 5).

Filter procurement is coordinated byU. Hopp and B. Muschielok.

3.4 Control Software

All instrument functions (filter ex-change, shutter, detector readout, aswell as monitoring the instrument state)are controlled in software. The pro-gramming environment is defined andprovided by ESO through the releasesof the VLT Common Software whichhas to be used as the basis for designand development. The partitioning ofthe OmegaCAM Instrument Software(OmegaCAM INS) into software sub-systems also follows the VLT stan-dards. Nevertheless there were severalchallenges peculiar to OmegaCAM.

The Autoguiding Software and ImageAnalysis modules normally belong tothe Telescope Control Software. In thecase of OmegaCAM it was necessaryto move these functionalities to the INSbecause during normal operations theVST guiding arm will not be used, as itslightly vignets the science array. A newsoftware algorithm was developed toextract optical aberration coefficientsfrom the out-of-focus images recordedon the Image Analysis CCDs. On thedetector software side, particular atten-tion had to be paid to the coordination

of the readouts by the differentFIERA’s, and to the efficient storage ofthe data on disk.

The Instrument Software is beingproduced by the Italian part of the con-sortium, headed by A. Baruffolo (INAF-Padua), and is described in more detailin Baruffolo et al. (2002).

4. Calibration and DataReduction Software

The amount of data produced byOmegaCAM will be truly huge. We es-timate that there will be over 15 Tera-byte of raw data per year. This raw datavolume contains roughly 5 Terabyte ofcalibration data and 10 Terabyte of rawscience data. Data processing will thenproduce another 10 Terabyte of re-duced science data and may create,with about 100,000 astronomical ob-jects per OmegaCAM field, enormouscatalogues. To efficiently handle thisdata volume the data acquisition, cali-brations and the pipeline reductions arestrictly procedurized, a key aim being to

maintain the instrument, not individualdata sets, calibrated at all times. ESOwill operate the instrument in servicemode, optimizing the observing pro-gramme to ambient conditions, androutinely taking calibration data. Thuseach night the instrument’s overall re-sponsivity and also the transmission ofthe atmosphere will be monitored in theu′, g′, r′ and i′ bands irrespective of theschedule of science observations. Datareduction recipes, run in ESO’s DFS,will provide a continuous characteriza-tion of the behaviour of the instrumentin these key bands. When other filtersare used, the calibration plan foreseesa cross calibration of these filters ver-sus these key bands.

The basic technique to overcomeany gaps or artefacts in the CCD pixelsis to take more exposures of the samefield with slightly shifted field centre andto co-add the images off-line. We dis-tinguish the following observing modes:

• Dither has offsets matching themaximum gap between CCDs, ∼ 400pixels (5.6 mm). It will be operated with

17

Figure 4: The OmegaCAM exposure shutter. The aperture size is 370 × 292 mm, the short-est possible exposure time is smaller than 1 msec, the deviations of the effective exposuretime from pixel to pixel (homogeneity) are smaller than ±0.2% for a 100-msec exposure, theexposure time accuracy is about 0.3 msec. (The laptop computer gives an idea of thescale.)

Figure 5:Theoreticalthroughputcurves fromSAGEM for theSDSS filter set,and measuredquantumefficiency of oneof theOmegaCAMCCDs.

N (with 5 as the default value) pointingson the sky. Although this will nearly cov-er all the gaps in the focal plane andmaximizes the sky coverage, the con-text map of such data is complex. Anadvantage is that it will be relativelyeasy to couple the photometry amongthe individual CCDs.

• Jitter has offsets matching thesmallest gaps in CCDs ∼ 5 pixels. Thismode optimizes the homogeneity of thecontext map and will be used during ob-servations for which the wide gaps arenot critical, but which, for instance, re-quire a well-mapped smoothly varyingPSF.

• Stare allows re-observing one fixedpointing position multiple times. It is themain workhorse for monitoring the in-strument and allows detection of opticaltransients.

• SSO is the mode for observingSolar System Objects, which requiresnon-sidereal tracking.

For all these modes dedicated ob-serving templates are being developed.

An observing strategy employs oneor a combination of the basic observingmodes. It also defines a number of ad-ditional instructions for scheduling ofthe observations. We distinguish thefollowing strategies:

• Standard which consists of a singleobservation (observation block)

• Deep which does deep integrations,possibly taken at selected atmosphericconditions over several nights

• Freq which frequently visits (moni-tors) the same field on time scalesranging from minutes to months andhas overriding priority on the telescopeschedule

• Mosaic maps areas of the sky larg-er than 1 degree, which is essentiallyan item for the scheduling, as thepipeline has to produce uniform qualitydata anyway. The combination of vari-ous field centres into one image is notconsidered a standard pipeline task.

The observing modes and strategiesare fully integrated with the data reduc-tion software being developed by theOmegaCAM consortium. We distin-guish between a calibration pipelineproducing and qualifying calibrationfiles, and an image pipeline that appliesthe calibration files to raw data andtransforms them into astrometricallyand photometrically calibrated images.ESO users will be provided with theoutput of the image pipeline, run inGarching, on the data contained in asingle OB. The nominal photometricaccuracy of this pipeline will be ± 0.05mag, exceptionally ± 0.01 mag. Thenominal accuracy for the astrometry is± 0.1 arcsec rms over the entire field ofview.

As part of the contract, theOmegaCAM consortium will deliversoftware modules that ESO will inte-grate into the image pipeline. In addi-tion, a project has been set up amongEuropean wide-field imaging groups toprovide a ‘wide-field imaging surveysystem’ that will combine pipeline pro-cessing of image data with archivingand data mining tools. Further detailscan be found on http://www.astro-wise.org, and in Valentijn & Kuijken(2002).

The development of the analysissoftware is being done by a team basedin Groningen and Leiden, led by E.Valentijn.

5. Current Status

The OmegaCAM project is now wellinto in the manufacturing phase. Mostof the CCDs have been delivered andtested; most of the mechanics exist andare ready to be integrated; instrumentcontrol and data analysis software isbeing coded. Extensive tests in Europeare foreseen for the second half of2003, and the camera should see firstlight early in 2004. Exciting times!

Acknowledgements

The consortium was formed in 1998in response to an announcement of op-portunity from ESO, and comprises in-stitutes in the Netherlands (NOVA, inparticular the Kapteyn Institute Gronin-gen and Leiden Observatory), Ger-many (in particular University Observa-tories of Munich, Göttingen and Bonn)and Italy (INAF, in particular Padua andNaples observatories). The ESO Op-tical Detector Team provides the detec-tor system at cost to the consortium.OmegaCAM is headed by PI K. Kuijken(Groningen and Leiden University) andco-PI’s R. Bender (Munich USM/MPE)and Cappellaro (INAF Naples/Padua),and project management is done by B.Muschielok and R. Häfner (USM).

OmegaCAM is funded by grants fromthe Dutch Organization for Research inAstronomy (NOVA), the German Fed-eral Ministry of Education, Science, Re-search and Technology (grants 05AV9MG1/7, AV9WM2/5, 05 AV2MGA/6and 05 AV2WM1/2), and the ItalianConsorzio Nazionale per l’Astronomiae l’Astrofisica (CNAA) and Istituto Na-zionale di Astrofisica (INAF), in additionto manpower and materials provided bythe partner institutes.

ReferencesArnaboldi, M., Capaccioli, M., Mancini, D.,

Rafanelli, P., Scaranella, R., Sedmak, G.and Vettolani, G. P., 1998. The Mes-senger 93, 30.

Baruffolo, A., Bortolussi, A., De Pizzol, L.2002, Proc. SPIE 4848, in press.

Nicklas, H., Harke, R., Wellem, W., Reif, K.2002, Proc. SPIE 4836-34, in press.

Reif, K., Klink, G., Müller, Ph. and Posch-mann, H. 2002 in Scientific Detectors forAstronomy, Beletic, Amico eds., Astro-physics and Space Sciences Library(Kluwer: Dordrecht), in press.

Valentijn, E.A. & Kuijken, K. 2002 in Towardan International Virtual Observatory,Quinn, P., ed., ESO Astrophysics Sym-posia Series (Springer-Verlag), in press.

18

The VLTI – 20 Months after First Fringes A. GLINDEMANN, ESO

1. Introduction

In 2002, the second year of fringes atParanal, the VLTI has made substantialprogress. The highlight was the comple-tion of the combination in pairs of all fourUnit Telescopes on September 15/16and 16/17 using a total of five differentbaselines. Only the combination MELI-PAL – YEPUN could not be provideddue to the current configuration of delaylines in the interferometric tunnel.

Of equal importance was the start ofa total of 150 hours shared risk scienceoperations with the VLTI in October.

Forty proposals from the communitywere received representing about 10%of all proposals submitted to ESO forthe VLT observatory. A summary of thefirst semester with VLTI science opera-tions will be given at the end of this se-mester. For Period 71, the shared riskscience operations became a part ofthe ESO Call for Proposals with 25 pro-posals submitted for the VLTI. A num-ber of observation preparation tools havebeen developed in collaboration withthe Jean-Marie Mariotti Centre for Inter-ferometry (JMMC) in Grenoble. Two ofthem are now available on the web

(http://www.eso.org/observing/etc/ pre-view.html). In the course of the year, allscience data between First Fringes inMarch 2001 and September 2002 havebeen released through the archive re-sulting in first scientific results whichare described in [1]–[4]. A summary ofthe first results is given in [5]. In thecontext of science operations, the re-sults of the on-going observations ofcalibrator stars are reported in [6], incollaboration with the NOVA ESO VLTIExpertise Centre (NEVEC) in Leiden.

Amongst the runners-up for achieve-ments are the integrated optics beam

combiner IONIC for VINCI, the accept-ance of the science instrument MIDIin Heidelberg in September (currentlybeing integrated at Paranal), and threenew contracts to extend the VLTI infra-structure: Delay Lines 4, 5 and 6 wereordered at Fokker Space in Leiden, andthe installation of the required rail sys-tems in the Delay Line Tunnel started inAugust. The contract for the fourthAuxiliary Telescope was signed withAMOS in Liège at the beginning ofSeptember (see [7]). The first PRIMAcontract for the delivery of the fringesensor unit was signed with ALENIASpazio, Torino, in July, and the secondcontract for the star separator unit willfollow soon. The PRIMA laser metrolo-gy system will be an in-house develop-ment. These three subsystems will formthe first phase of PRIMA.

IONIC was provided by the Observa-toire de Grenoble. IONIC is a beam com-biner for VINCI that can be used insteadof VINCI’s fibre beam combiner MONA.In IONIC, the beams are guided throughtwo silicon fibres into an integrated op-tics (IO) component where the beamcombination takes place. At the exit ofthe IO component, fibers are attachedguiding the light to the infrared detector.Since the fibre connectors are identicalto those used in MONA, the optical in-tegration was just a matter of decon-necting and connecting a few fibres.Unlike the K-band beam combinerMONA, IONIC, having silicon fibres andsubstrates, works in the H-band. TheVLTI had first fringes with IONIC onJuly 18, producing an interferometrictransfer function above 80% with anaccuracy of 1% which is as good asMONA (see [8]).

Last but not least there is alsoprogress with the adaptive optics sys-tems for the VLTI. The Paranalization ofthe tip-tilt sensors STRAP on two UnitTelescopes is finished, and MACAO,the high-order adaptive optics system(see [9]), saw first light in the laborato-ry in Garching in August, closing theloop on an artificial star.

2. Four Eyes Are Better

Why should one attempt to use morethan two telescopes if the interferomet-ric instrument can only combine twobeams? First of all, this exercisedemonstrated the ability to provide mul-tiple beams in the VLTI laboratory with-out losing time for mirror realignment orfor other reconfigurations. In 2003, AM-BER, the near infrared science instru-ment, will see first fringes. Then we willhave an instrument combining up tothree telescopes simultaneously.

In addition, multiple baselines are thekey to efficient interferometric imaging.The result of the observations of thestar Achernar in September, shown inFigure 1, gives an impression of the an-gular resolution capability of the VLTI

with the Unit Telescopes. The intensitydistribution was reconstructed from themeasurements with individual baselinesby Fourier transforming the distributionof values displayed in Figure 2. Everyballoon in that figure represents the re-sult of a visibility measurement with onebaseline vector. It is important to prop-erly refer to baselines as baseline vec-tors in order to understand why the dis-tribution of the balloons is considerablydifferent from the distribution of UnitTelescopes at Paranal (see Fig. 3). Thevisibility function depends on the differ-ence of telescope coordinates (project-ed perpendicular to the viewing direc-tion of the star). For instance, two tele-scopes separated by 100 m along thenorth-south direction have exactly thesame baseline vector as two other tele-scopes a few kilometres away that arealso separated by 100 m along north-south.

The baseline vectors are displayed inthe ‘uv-plane’, where the measuredvisibility forms the function value at thecoordinate given by the length andthe orientation of the baseline vector.Since the visibilities are derived fromthe intensity, a real and positive quanti-ty, the visibility function is point sym-metric; one visibility measurement de-

livers two function values, at (u, v) andat (-u, -v).

The balloons representing the meas-ured visibilities in the uv-plane aresmall disks although a point would beexpected when observing with an indi-vidual baseline. The reason is that thetelescope is not a pinhole but an 8-maperture allowing for a small variationof baselines around the ’centre-to-cen-tre’ baseline. Inside VINCI, optical fi-bres are used as spatial filters. The sizeof the optical fibres is exactly an Airydisk of the telescope, i.e. the field ofview is one Airy disk. This spatial filter-ing is equivalent to performing an aver-age over the visibilities inside the diskrepresented by the balloons.

For imaging it is important to note thatone measures as many independentimage points inside the Airy disk as onehas baselines, i.e. as one has degreesof freedom. Fortunately, the number ofbaselines is not only determined by thenumber of telescopes but also by theduration of the observation. Due to thesidereal motion, the baseline projectedonto the viewing direction changes allthe time in length and orientation de-pending on the object coordinates.Some of the baselines that are dis-played in Figure 2 have a slight peanut

19

Figure 1: The reconstructed, two-dimensional interferometric point spread function (PSF) ofthe star Achernar observed in the K-band. The width of the central fringe is 3 by 15 milli-arc-sec indicating the angular resolution limit of the VLTI for the baseline distribution shown inFigure 2. On the largest scale, the image is enveloped by the Airy disk of a single 8-m UnitTelescope. Its first minimum at 57 milli-arcsec off the centre can be clearly seen. The imageprovides a dramatic illustration of the 20-fold increase in resolution of the VLTI over a single8.2-m telescope.

shape due to the observing time of afew hours. Several individual meas-urements were taken during this time.Therefore, the number of individualpoints in the image will be considerablylarger than six when observing with allfour Unit Telescopes.

The present observations of Acher-nar give an impression of the pointspread function of the VLTI becauseAchernar with an angular diameter of 2milli-arcsec is smaller than the resolu-tion limit of 3 milli-arcsec of the VLTI inthe K-band. Due to the moderate num-ber of baselines, the point spread func-tion in Figure 1 shows a mix of a mod-erate number of individual fringe pat-terns. A larger number of visibilitymeasurements would make the individ-ual fringe patterns disappear, and onlythe central fringe would remain as the‘point’ in the point spread function.

The particular distribution of baselinesin Figure 2 is rather extended in thenorth-east/south-west direction (thebaseline of ANTU-YEPUN) and rathernarrow in the perpendicular directionbecause the baseline MELIPAL-YEPUNis missing. Consequently, the Fouriertransform of this distribution is rathernarrow in the north-east/south-west di-rection and rather wide in the perpendi-cular direction. The terms narrow andwide refer to the width of the centralfringe. The central fringe in Figure 1 hasa width of about 3 milli-arcsec in the di-rection of the longest baseline and ofabout 15 milli-arcsec in the perpendicu-lar direction. Due to the processingtechnique, the Airy disk of an individual

8-m telescope shows up in the pointspread function in Figure 1. The firstdark ring of the Airy disk with a radiusof 57 milli-arcsec in the K-band is read-ily apparent. This illustrates the gain inangular resolution with the VLTI com-pared to an individual Unit Telescope.

However, it is still a long way to go be-fore the VLTI will do imaging of arbitraryobjects since this requires measuringthe phase of the visibility function in ad-dition to its modulus i.e. its contrast.AMBER in 2003 will be a first step, be-ing able to use three telescopes andperforming a phase reconstruction viathe closure phase technique. Eventu-ally, from 2005 onwards, PRIMA will al-low measuring the phase directly perbaseline and, in addition, observe veryfaint objects. The scientific objectives ofthe PRIMA facility are described in [10].

3. OutlookThe coming year will be extremely

busy, with integrations of new subsys-tems about every two months. FINITO,the fringe sensor unit, will be first, fol-lowed by two MACAO systems, by thefirst two Auxiliary Telescopes, by AM-BER and by the carriages with cat’seyes of the three new Delay Lines. Thisis a big burden for the integration teamssince the VLTI is not the only ‘tele-scope’ at Paranal.

In 2004, the remaining two MACAOsystems and the third and fourthAuxiliary Telescopes will arrive atParanal so that the VLTI will be able tocombine the light from eight differenttelescopes, four Unit Telescopes and

four Auxiliary Telescopes. With sixDelay Lines, a common focus of six tel-escopes can be used at any given mo-ment. To fully exploit the VLTI infra-structure, the first second-generationinstrument should clearly be a six-waybeam combiner, measuring 15 base-lines at once and 28 during a night,without reconfiguring the telescopes.

Amongst all these integration andcommissioning activities one must notforget that the success of an observa-tory depends solely on the scientificoutput. MIDI will make a start next year,and 2004 will be largely devoted to sci-ence operations with the VLTI.

AcknowledgementOver the last 20 months, the number

of people both at ESO and in the com-munity contributing to the VLTI has be-come too large to name here. I wouldlike to thank not only those with thehighest visibility but also the large num-ber of contributors forming the base ofthe success by ensuring that every littlepiece of this complex machine worksevery night.

References

[1] Kervella, P., et al. 2002, in Exotic Starsas Challenges to Evolution, IAU Sym-posium 187, in press.

[2] Wittkowski, M. and Hummel, Ch. 2002, inInterferometry for Optical Astronomy II,ed. Traub, W., Proc. SPIE 4838, in press.

[3] Richichi, A. and Wittkowski, M. 2002, inThe VLTI: Challenges for the Future, eds.Garcia P.J.V., Glindemann A., HenningT., Malbet F., JENAM Workshop, in press.

20

Figure 2: The set of baselines used for observingAchernar with the four Unit Telescopes. Each baseline isrepresented by two opposite, short arcs, symmetric aroundthe origin (centre) of the diagram. The colour-coded patternreflects the telescope pairs (ANTU-KUEYEN = magenta,ANTU-MELIPAL = red, ANTU-YEPUN = green, KUEYEN-MELIPAL = cyan, KUEYEN-YEPUN = blue).

Figure 3: The lines indicate the Unit Telescope and AT sta-tions that have been combined at Paranal to date. Thecolour scheme for the UTs corresponds to the baselinesshown in Figure 2.

N

[4] Segransan, D., et al. 2003, A&A Letters,in press.

[5] Paresce, F., et al. 2002a, in Inter-ferometry for Optical Astronomy II, ed.Traub, W., Proc. SPIE 4838, in press.

[6] Percheron, I., et al. 2002, in The VLTI:Challenges for the Future, eds. Garcia P.

J. V., Glindemann A., Henning T., MalbetF., JENAM Workshop, in press.

[7] Koehler, B., et al. 2002, The Messenger,this volume.

[8] Kern, P., et al. 2002, in Interferometry forOptical Astronomy II, ed. Traub, W.,Proc. SPIE 4838, in press.

[9] Arsenault, R., et al. 2002, in AdaptiveOptical System Technologies II, eds.Bonaccini, D., Wizinowich, P., Proc.SPIE 4839, in press.

[10] Paresce, F., et al. 2002b, in Interferom-etry for Optical Astronomy II, ed. Traub,W., Proc. SPIE 4838, in press.

21

The Auxiliary Telescopes for the VLTI: a Status ReportB. KOEHLER1, C. FLEBUS2

P. DIERICKX1, M. DIMMLER1, M. DUCHATEAU1, P. DUHOUX1, G. EHRENFELD3, E. GABRIEL2, P. GLOESENER2, V. HEINZ3, R. KARBAN1, M. KRAUS1, J.M. MORESMAU1,N. NINANE2, O. PIRNAY2, E. QUERTEMONT2, J. STRASSER, K. WIRENSTRAND1

1ESO, Garching, Germany; 2AMOS, Liège, Belgium; 3ESO, Paranal, Chile

1. Introduction

In June 1998, ESO signed a contractwith the company AMOS (Belgium) forthe supply of the Auxiliary TelescopeSystem (ATS) for the VLTI. The originalscope covered two movable AuxiliaryTelescopes (ATs), as shown in Figure 1,as well as their associated site equip-ment including rails and interface de-vices for each observing station. Anamendment was signed in September1999 for the supply of a third AT and,last September, a fourth AT was or-dered.

The contract with AMOS is based ontop-level performance requirementsand includes the design, manufacturingand testing of the complete system in-cluding all mirrors and cells, completemechanical structure, drives, encoders,small mechanisms and low-level elec-tronics. The main sub-contractors ofAMOS are FISBA (Switzerland) for thecoudé mirrors, PHASE (Italy) for themotors and CSEM (Switzerland) for theM2 hexapod mechanism.

ESO is in charge of the design anddevelopment of the telescope controlhardware and software, as well as ofthe two star sensors located at thecoudé focus.

After integration of the ESO controlhardware and software, the telescopeis fully tested in Europe at AMOS withthe possibility of sky observation pro-vided by an outside observing stationincluded in the dedicated test facility.After this, ESO transports the ATs toParanal, re-integrates them in the so-called Mirror Maintenance Building(MMB) and finally commissions themon the sky.

The project entered into manufactur-ing phase in mid-1999 and is nowreaching the end of a very extensivetesting and verification phase in Europeon AT#1 before its shipment to Paranalnext year.

This article recalls the rationale at theorigin of the ATS development, pro-vides a description of the design and fi-nally reports on the performance asmeasured in Europe so far.

2. Why Does the VLTI NeedAuxiliary Telescopes?

The VLTI is primarily intended tocombine coherently the four VLT 8-mUnit Telescopes (UTs). Obviously, theultimate VLTI sensitivity will indeed beobtained when combining the UTs.However, the arrayof 1.8-m diameterATs is a key ele-ment for the tech-nical and scientificcapability of the in-terferometer. Themain reasons arelisted below.

• It provides thebest imaging capa-bility of VLTI bycomplementing thearray of UTs. Itgives access to 30telescope stationsincreasing there-fore the number ofaccessible base-line vectors, a fun-damental parame-ter for high-fidelityimage reconstruc-tion.

• It gives accessto the longest VLTIbaseline of B = 200

m versus a maximum of 130 m be-tween UTs. This is needed to reach theultimate angular resolution of the VLTIthat scales as λ/B (i.e. 0.6 milli-arcsecin the visible and 2 milli-arcsec in the Kband).

• It enables full-time use of the VLTIfacilities, since the ATs are entirely ded-icated to interferometric observations.This is an important factor for the VLTIscientific productivity (and for the amor-tization of its development cost!).

• It is required by the NarrowAngle Astrometry mode of PRIMA that

Figure 1: The firstAuxiliary Telescopefor the VLTI duringfinal testing atAMOS in Liège(Belgium).

needs regular and long-term moni-toring out of reach with the highly sub-scribed UTs.

• It can be used for testing and com-missioning of other VLTI facilities (e.g.for the 2nd generation of instruments)without using precious UT time.

For these reasons, the ATs are ex-pected to become the ‘workhorses’ ofthe VLTI and to deliver a large amountof scientific results for those pro-grammes that do not fundamentallyneed the sensitivity provided by theUTs.

The typical magnitude limits for theVLTI in the K band (2.2 µm) and in theN band (10 µm) are expected to be re-spectively K = 11 and N = 4 for the UTsequipped with Adaptive Optics. For theAT equipped fast tip-tilt correction only,these become K = 9 and N = 1.

The design of the VLTI infrastructureforesees up to eight Auxiliary Tele-scopes and eight Delay Lines. The rea-sons for increasing the number of tele-

scopes in the array are essentiallylinked to the imaging capability. Indeed,the fourth AT has recently opened thepossibility for second-generation instru-ments to observe 6 baselines simulta-neously. This will drastically improvethe snapshot imaging capability, partic-ularly useful for variable objects. Evenwith a 3-beam instrument like the VLTIinstrument AMBER, the 6 baselinescan then be observed quasi simultane-ously by a fast switching of a mirror inthe VLTI optical train instead of relocat-ing the ATs, an operation that can bedone during the day only. Similarly, thenumber of phase closures available perAT configuration will increase from 1 to3. This represents a three-fold increasein observing efficiency in this powerfulimaging mode where the phase errorsintroduced by the atmosphere can becancelled out in order to retrieve thephase of the object, a critical parameterfor image reconstruction. A fourth ATwill also increase the limiting magnitude

by using the so-called ‘boot-strapping’technique where the fringes are trackedon the relatively short intermediatebaselines while the low-visibility fringesare integrated on the extreme baseline.In addition, the fourth AT will permit todedicate two ATs for the PRIMA as-trometry mode while two others are be-ing used by MIDI or AMBER. This mayprove a fundamental advantage duringthe challenging commissioning phaseof PRIMA by alleviating the possibleconflicts with the regular VLTI scienceobservation.

3. Specifications and DesignDescription

3.1 Basic conceptual choices andmain specifications

The ATS specifications originate fromits dedication to interferometric obser-vations down to visible wavelengths.This calls for an extreme mechanicalstability, at the level of a few nano-me-ters, together with the capability to eas-ily relocate the telescope.

The following conceptual choiceswere made as a result of early feasibil-ity studies.

• Primary mirror diameter 1.8 m, as acompromise between science capabilityand cost.

• Self-contained autonomous Tele-scope with built-in ‘Transporter’ en-abling relocation of the telescope onobserving stations connected by a railnetwork and built-in Enclosure ensuringsurvival in outdoor conditions (wind,earthquake).

• No active optics to limit control com-plexity already inherent in the interfer-ometer, but 5-axis M2 support for re-motely controlled alignment.

• Field acquisition and fast tip-tilt sen-sors at the coudé focus and provisionfor later implementation of AdaptiveOptics.

• Stiff Alt-Az mount with Altitude axis5 m high as a compromise betweenstiffness and ‘ground seeing’ effects.

• Pseudo-kinematic interface be-tween telescope and ground for highposition and angular repeatability sothat only fine remotely controlled align-ment is required after a relocation.

• Mechanical bearing for both axes tolimit maintenance and avoid hydraulicconnections at station or on-boardpumps.

• Direct drives for improved controlla-bility associated with high accuracy op-tical encoders and analogue tachome-ters.

Table 1 summarizes the main per-formance requirements specified forthe Auxiliary Telescope System. Themost stringent ones imposed by inter-ferometry are the excellent image qual-ity, the high tracking accuracy and theOptical Path Length (OPL) stability.Other constraints imposed by the site

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Figure 2: Overview of the Auxiliary Telescope System. The telescope is anchored on one ofthe 30 observing stations. It is surrounded by its built-in “Transporter” used to relocate the tel-escope, to support its enclosure and to house electronic cabinets and service equipment.Inside the station pit, the cylindrical structure contains the coudé sensors and the RelayOptics that sends a parallel beam to the Delay Line Tunnel trough underground light ducts.

Optical quality of the complete telescope < 110 nm r.m.s WFE Wavelength range: 0.4 to 25 µmTelescope position repeatability: ± 10 arcsec, ± 0.1 mm after any reloca-tionFirst telescope eigenfrequency: > 10 HzPointing accuracy: < 1 arcsec r.m.s. Tracking accuracy (blind): < 0.1 arcsec r.m.s.Tracking accuracy (Field Stab.): < 0.025 arcsec r.m.s. Optical Path Length Stability: < 30 nm rms over 10 msec Time to relocate to any station: < 3 hours with 2 operatorsTotal Mass: < 30 TonsDesign volume: H < 7.1m, W < 5.4 m & L < 5.8 m

Table 1: Main AT performance requirements.

infrastructure such as the total massand the design volume are also drivingfactors for the overall AT design.

3.2 Optical design

The AT optical design is similar tothat of the UT and includes 11 mirrors.The first three (M1-M3) constitute aRitchey-Chrétien telescope with a 1.8-m diameter primary mirror. A fast f/1.4primary mirror was selected to providea compact tube design and limit windbuffeting of particular concern for inter-ferometric application where nanome-ter-level stability is required in theOptical Path Length (OPL). One conse-quence of that particular design is thehigh sensitivity of the telescope to thesecondary mirror (M2) defocus & de-centre which places high demand onthe stiffness of the tube structure. TheNasmyth focus of the Ritchey-Chrétienis unused for normal operation. It is re-imaged by the 5-mirror coudé train (M4-M8) to form a coudé focus located un-derground. In addition, the coudé traincreates an image of the telescope pupilon M6 where a fast tip-tilt mirror (andpossibly later an adaptive optics mirror)is located. Just above the coudé focus,a dichroic mirror (M9) transmits the vis-ible to a CCD camera and an APD mod-ule (STRAP) for acquisition and fastguiding. The infrared light reflected byM9 is collimated by the Relay Optics(M10-M11) to form an 18mm parallelbeam sent towards the VLTI delaylines.

3.3 Telescope structure

An overview of the telescope designcan be seen in Figure 2. The main driv-ers were the maximization of the firsteigenfrequency, the minimization of theM1-M2 de-space and de-centre (typi-cally < 10 µm) under gravity, wind loadand thermal load as well as the mini-mization of the wind torque. This had tobe achieved while complying with theconstraints of mass, design volume, sta-tion interface and earthquake survival.

The Tube structure is a classicalSerrurier strut made of hollow tubesconnecting to a top-ring made out of fullrods for weight balance. The M2 spi-ders were optimized in order to limit thedesign volume, to reduce the torquedue to wind load on the top-ring and toincrease axial stiffness in view of OPLstability. A particular feature of the tubedesign is the solution adopted, afterseveral trade-offs, for the thermal com-pensation rendered necessary by theextreme sensitivity to M2 defocus. Thecompensation is provided by a properselection of material for the Serrurier(steel), top-ring (steel) and M2 spider(invar). When the temperature increas-es, the tube expands and the top-ringdiameter increases. This later creates astress in the M2 spiders that results ina downward motion of M2 compensat-ing the tube and M1 cell expansion.Several trade-offs were made on thecentrepiece that is a main contributor tostiffness and wind torque. The latterconstraint led to the final design based

on a complex strut structure with ap-propriate welded box/plates interfacesto the altitude shaft, Serrurier structureand M1 cell. On each side of the cen-trepiece, a semi-circular plate holds themagnets for the altitude motors andtachometers and a brake disk. For thesupport of M3, a spider was preferredover a more classical tower solution be-cause it provides better static and dy-namic results while easing the re-align-ment after an M1 removal for re-coating.

The tube is connected to the Fork onboth sides by double row angular con-tact ball bearings providing high stiff-ness, low run-out and low friction. TheFork is a welded box structure with in-ternal ribs highly optimized for stiffnessand mass. At its lower part, the Forkconnects to the Azimuth bearing, a keyelement to the overall telescope stiff-ness. The Azimuth bearing is a three-row roller bearing with very high radial,axial and tilt stiffness but, at the sametime, very low friction torque necessaryfor the pointing/tracking accuracy.

Below the Azimuth bearing, theGround Interface Structure (GIS) pro-vides the connection to the groundthrough the pseudo-kinematic anchor-ing and clamping devices. This part ofthe telescope was also intensely opti-mized by FEM due to its importance inthe first telescope eigenfrequency.

Inside the station pit the Relay OpticsStructure (ROS), a double-skin weldedbox cylindrical structure, houses theRelay Optics and the coudé sensors.When relocating an AT on a new sta-

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Figure 3: The 1.8-m primary mirror in its cell during integration intothe telescope at AMOS. The mirror is lightweighted to reduce itsmass by a factor about 2. A 54-points wiffle-tree axial support wasselected to comply with the stringent image quality requirements.

Figure 4: The 140-mm diameter secondary mirror is mounted on anhexapod to enable remotely controlled focus, centring and angularalignment. The hexapod, built by CSEM (Switzerland) is a very com-pact, very high accuracy mechanical device featuring a resolutionsmaller than 0.2 microns and a very small cross-coupling betweenfocus and tilt motions to avoid mispointing while re-focusing the tel-escope.

tion, this structure is lowered into thestation hole on soft pads, the telescopeis placed on top and a centring andclamping mechanism is activated be-tween the structure and the GroundInterface Structure. The position re-peatability is such that local alignmentis not required after relocation. Thisholds also for the repeatability of thetelescope anchoring system. The onlytelescope alignment necessary after re-location is performed remotely from theControl Room.

3.4 Drives, tachometer andencoders

The motors selected for both Altitudeand Azimuth axes are direct drive,brushless torque motor, similar in con-cept to those installed in the UT. Twocurved linear motor pads are used dia-metrically opposed for Azimuth and onepad is attached to each side of the forkfor Altitude. Each motor is equippedwith internal position sensors used inparticular for motor commutation.

Each axis is equipped with an AC-DCanalogue tacho-generator using thesame magnet strip and similar pads asthe motor. The power amplifier is basedon fast DSP technology and providesmotor current control, overspeed limita-tion and tacho demodulation.

The axis position information is givenby high accuracy Heidenhain opticalencoders. The high interpolation factorleads to the very high angular resolu-tion of 9 milli-arcsec on Altitude and 3milli-arcsec on Azimuth.

3.5 Primary mirror unit

The M1 unit is shown on Figure 3. Itconsists of the primary mirror, its cellwith axial and lateral supports and fi-nally the interface with the telescopetube.

The primary mirror is made out of a200-mm thick blank of Zerodur materialfrom SCHOTT (Germany). Due tomass constraints, the mirror is light-weighted by drilling hexagonal pocketson its rear face as can be seen onFigure 3. It is polished to a hyperbolicshape.

The mirror axial support is the resultof an intensive optimization by FiniteElement Modelling in order to minimizethe effect of gravity deflection. The se-lected solution compatible with thelightweight geometry and the low-levelerror budget is a 54-axial-point supportmounted on a three-stage wiffle-treesystem using flexible pivots.

For the lateral support, the retainedoptimal solution consists of 16 astaticlevers using low-friction bearing and

counter-weights. In addition, three radi-al fixed points are located inside the M1central hole.

The M1 cell consists mainly of awelded steel structure with a doubleskin for the bottom and lateral parts toimprove the cell stiffness. The structureis made of two parts to facilitate the mir-ror integration and maintenance. Thecell’s lower part interfaces with the axi-al support while the upper part inter-faces with the lateral support. The firsteigenfrequency of the M1 unit is 20 Hzand the mirror degradation in its cell un-der gravity is less than 11 nm RMS sur-face.

3.6 Secondary mirror unit

Due to the selected fast f/1.4 primarymirror, the secondary M2 mirror is arather small mirror 140-mm in diameter.However, its highly aspherical and fast(f/1.3) shape makes it a difficult mirrorto manufacture. It is also made out of ablank of Zerodur, 25 mm thick.

The M2 is mounted on a hexapodsupport (see Fig. 4) providing defocus,de-centring and tip-tilt motions withsub-micron and arcsecond level resolu-tion for remote alignment. Its design isprimarily driven by the very limitedspace available (in the shadow of M2)

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Figure 5: The first telescope on its test bench inside the dedicatedhall at AMOS. In addition to tests at component level such as opticalquality of each mirror, friction of the axis bearings, etc., the first tel-escope undergoes a thorough acceptance test programme includingmodal tests of the complete structure, dynamic response of the azi-muth and altitude axis, OPL stability measurement, verification of struc-tural deformation under gravity, pointing repeatability check and others.

Figure 6: The Transporter in its relocation configuration. The cylin-der in the foreground houses the Relay Optics and the coudésensors (CCD and STRAP). The light beam exits the telescope byone of the two black openings. During relocation, the cylinder is low-ered underground inside the station pit before the telescope isplaced on top.

and the requirement to limit cross cou-pling between focus and tilt motion.

3.7 Enclosure

The basic enclosure concept is toprovide good natural ventilation duringthe night to limit “dome and mirror see-ing” and therefore to leave the part ofthe telescope above the altitude axisexposed to the outside wind (see Fig.1). Following an analysis of the windload, thermal requirements and massand volume constraints, the selecteddesign consists of composite sphericalshells that open up with two hydrauliccylinders attached to each of the twoupper shells. When closed, the two up-per shells are clamped together bymeans of 6 electrical actuators and in-flatable seals are inflated in-betweeneach shell. The normal status of the en-closure is fully open or fully closed.However, it is possible to close totally orpartially any of the two sides in order toact as a wind shield in case of high windcoming from S-E or N-W, the two dom-inant wind directions.

3.8 Transporter and ServiceModules

The main function of the Transporteris to provide the handling means to re-locate the telescope without the needfor any additional external devices. Atypical relocation sequence can besummarized as follows: (i) unclamp theRelay Optics Structure from the tele-scope, (ii) unclamp the Telescope, (iii) liftthe Telescope, (iv) unclamp the Trans-porter, (v) lower down the Transporterwheel on the rails and lift the Trans-porter in driving position, (vi) drive theTransporter beside the station, (vii) at-tach the Relay Optics to its handling de-vice integrated on the Transporter andlift it, and (viii) drive to the next station.When a change of direction is needed(at one of the rail crossing), the Trans-porter is stopped at the crossing, low-ered and clamped on special interfaceplates. The wheels are lifted up and ro-tated by 90° before being lowered onthe perpendicular rails. Unclampingand lifting the Transporter enables todrive to the next crossing. At the desti-nation station, the reverse operationsare performed.

The Transporter (see Fig. 2) consistsof a frame structure made of hollow alu-minium rectangular profiles. It has beenoptimized with respect to mass mini-mization (aluminium), space con-straints and earthquake survival duringobservation and relocation (with tele-scope load). An inner cylindrical struc-ture supports the enclosure and theTelescope lifting devices. An outercylinder structure, covered with alu-minium panels and equipped with ac-cess doors, provides protection of theelectronics cabinets and the Service

Modules that are fitted in-between thetwo cylinders. The four drive wheelblocks, attached to the outer cylinder,include the electrically driven wheeland hydraulically driven lifting, rotationand blocking mechanisms. The tele-scope lifting device consists of 8 liftingactuators linked to cardan shafts whichare driven by two electric motors. TheRelay Optics handling device uses twoelectric linear actuators linked by afloating shaft.

The other functions of the Trans-porter are to provide air conditioning,cooling liquid, auxiliary power as wellsas hydraulic and pneumatic power forthe various mechanisms. It also pro-vides the space for the control elec-tronics cabinets.

The air conditioning system providestemperature control during the daybased on three temperature sensors lo-cated on the primary mirror, plus sixother sensors distributed on the tele-scope structure. The mirror tempera-ture is controlled to ±0.5°C of ambientair with the capacity to cool the mirrorby 2°C in about 3 hours. The coolingpower is injected with filtered, re-circu-lated air through a circular fabric ductwith constant flow and high induction toensure temperature homogeneity. Thecontrol is done through a PLC with anoptimal set point computed by the ESOAT Control System based on outsidetemperature measurement and fore-cast for the next night.

The liquid cooling module is provid-ing cooling power during observationfor temperature control of electroniccabinets and heat-dissipating equip-ment such as motors, coudé sensorsand small mechanisms. The remainingheat dissipated in the vicinity of the op-tical beam is maintained below 25Wand surface temperatures are kept atambient ±1.5 °C. In order to avoid thevibration of the chiller during observa-tion, the coolingenergy is stored in-side a tank of heatlatent nodules. Thecooling capacity is1 kW during 12 hand is re-generat-ed during the dayby running thechiller.

The auxiliary power module consistsof a set of batteries feeding a UPS toprovide uninterrupted power to the con-trol electronics as well as an electrogengroup connected to the UPS to provideelectric power during relocation whenthe telescope is unplugged from its sta-tion socket.

The hydraulic group provides therequired oil pressure and flow for thevarious hydraulic mechanisms. Suchmechanisms are used for the Trans-porter wheels 90° rotation, the Trans-porter lifting, the clamping of Trans-porter to ground, Telescope to groundand Relay Optics Structure to Telescope,as well as for the Enclosure opening/closing.

Finally, a pneumatic system is usedto inflate/deflate the Enclosure sealsand to activate the Telescope axesbrakes.

3.9 Station equipment and rails

The station equipment consists of aset of Telescope interfaces, Transporterinterfaces and a station lid. The tele-scope interface is made of a platewhose upper face is machined to pro-vide a perfect supporting plane for thetelescope anchoring counterparts.These counterparts are male elementsconsisting of a clamping finger placedinside a cone. They come into the cor-responding conical part of the tele-scope clamping device to form a pseu-do- kinematic interface. The station Lidis a composite hexagonal cover fittedwith a seal. It ensures good thermal in-sulation, air tightness and resists to a500-kg central load. It can be handledby two persons.

The rails are standard rolled steelprofile with end cut at 45° to ensuresmooth load transfer at the junctions. Ata crossing, a specific design enables totransfer the load from the central rolling

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Figure 7: The firstTransporter sees thesun light for the firsttime at AMOS (Liège,Belgium) during itsfunctional tests lastOctober. A dummystation will permittesting of the tele-scope on the sky be-fore shipment toParanal.

surface of the wheel to its two lateralflanges. At each track end, an end-stopis installed.

3.10 Control system

The design & implementation of thecontrol system is shared betweenAMOS and ESO. AMOS is responsiblefor the Transporter control including thecomplete relocation sequence and theService Modules while ESO is incharge of the complete Telescope Con-trol including control of the axes, the M2support, the field stabilization loop aswell as the small telescope mecha-nisms such as M10.

On the AMOS side, two Programma-ble Logical Controllers (PLC’s) performthe control tasks. The first one (Sauter)is dedicated to the Air Conditioning andLiquid Cooling Module while the sec-ond one (Siemens) controls the otherTransporter functions including thoseneeded during relocation. For the latter,the operation is done locally from aportable control panel (see Figure 7).Both communicate with the ESO LocalControl Units (LCU’s) for exchange ofstatus information and commandswhen needed (e.g. switch on/off airconditioning, set reference tempera-ture, open/close enclosure, etc.)

On the side of ESO, the control sys-tem adheres to the Hardware andSoftware standards adopted for theVLT. It is based on a distributed archi-tecture of workstations (Unix) connect-ed to the LCU’s (VxWorks) via a LAN(Ethernet). For each AT, the controlelectronics located on-board theTransporter is connected through thestation plug and the LAN optical fibresto the VLT or VLTI Control Room fromwhere all control is done during regularoperation (i.e. except for relocation).

The design of the Software makesextensive re-use of the TelescopeControl Software (TCS) implementedon the UT with a few modifications re-quired by the difference in hardware.

However, two significant new ap-proaches have been applied on the AT.On the Hardware side, a fully digitalcontrol scheme was selected in placeof the former analogue velocity con-trollers for the telescope axes. This wasone of the first applications of what hasbecome now a new VLT standard: thegeneric Tool for Advance Control (TAC).Its main advantage is the possibility toconfigure and optimize the control algo-rithm in a very flexible manner even atruntime. TAC is now used for severalapplications in the VLT/VLTI (OPDController, FINITO). On the Softwareside, the Unified Software Develop-ment Process with the Unified ModelingLanguage (UML) has been introducedat the beginning of the AT project.

4. Current Status

As far as optics are concerned, the fullset of 11 mirrors is complete for each ofthe first three ATs except the primarymirrors for AT#2 and AT#3 that are re-spectively under final figuring and ma-chining. The blank for AT#4 has beenordered and light-weighting will startbefore the end of the year. Each mirrorhas been successively polished withinits very tight allocated specification de-rived from the glob-al telescope imagequality given inTable 1. As an ex-ample, the phasemap of the first pri-mary mirror for AT#1is shown on Figure 9.The achieved wave-front error amountsto less than 24 nmRMS. This is an ex-

cellent achievement for the first mirrorof this size and short focal ratio manu-factured by AMOS.

Currently, the first AT is in its last test-ing phase at AMOS. The telescope, ful-ly integrated with optics, is anchored ona dedicated station (see Fig. 5) whereall the major performance tests havebeen completed, as described in thenext section. Meanwhile, the Trans-porter has also been fully integratedand the second telescope – partially in-tegrated – has been placed inside inorder to proceed with an extensivefunctional test phase during which allmotions required for the relocationhave been exercised and tuned (seeFig. 7). One of the main remaining ac-tivities is the testing of the relocationperformance, including the time to relo-cate (less than 3 hours) and the re-peatability of the telescope positioning(± 100 µm and ± 10 arcsec). The latterwas already successfully tested on abreadboard during the design phase.The other important remaining test con-cerns the performance of the thermalcontrol provided by the air conditioningduring the day and by the liquid coolingmodule during night observation. InJanuary, the formal technical accept-ance in Europe of the first AT will take

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Figure 8: Installation of the AT interface devices inside each of the 30 observing stations at Paranal. A high-precision alignment method us-ing theodolite and specific tools is used to reach a positioning accuracy below a millimetre and an arcminute.

Figure 9: The 1.8-mprimary mirror of AT#1after polishing. Thewavefront error is be-low 24 nanometersRMS

place. Subsequently, the system will behanded over to ESO for a 3-month testperiod in Europe where the ESOControl System will be tested and fine-tuned with a possibility for sky observa-tions. The first AT will then be packedand sent to Paranal where re-integra-tion will start mid-2003.

The second and third telescopes arecurrently in the assembly process. Mostof the sub-contracted equipment for thetelescopes such as the main drives, M2hexapods, bearings are integrated. TheTransporter for AT#2 is being assem-bled while the one for AT#3 is still undermanufacturing.

At Paranal, the preparatory activitiesare taking up momentum in order to beready to receive the first AT. The tele-scope interface devices have been in-stalled and carefully aligned in each ofthe 30 stations (see Fig. 8) and the railinstallation is about to start.

5. First Results from TelescopePerformance Tests

In order to verify system-level per-formance requirements before the finalcommissioning on site, an extensivetest program has been carried on AT#1at AMOS. Some of the most importantresults are reported below.

5.1 Verification of the telescopeaxis dynamic response

The dynamic transfer function be-tween the motor torque reference and

the signals from the tachometer or theposition encoder is a critical character-istic for an accurate, high-bandwidthcontrol of the telescope axes (under theresponsibility of ESO). For this reason,the dynamic responses of both altitudeand azimuth axes have been specifiedand measured. Figure 10 shows the re-sult of the measurement to which asimplified dynamic model of the tele-scope could be fitted in order to verifythat the control bandwidth and stabilitycriteria will be met when ESO plugs inits telescope control system. The resultsare well within the specifications of 8 Hzbandwidth for the velocity loop and 2Hz bandwidth for the position loop.

5.2 Verification of theOptical Path Length stability

One of the most unusual and difficultrequirements for the AT is the stabilityof the Optical Path Length (OPL) insidethe telescope that is required by inter-ferometric observations. The OPL vari-ation was a driving parameter of the tel-escope mechanical design where opti-mization processes were necessary tofulfil the specifications. The identifiedcontributors to the OPL variation are thevibrations resulting from the dynamicwind load, the normal micro-seismic ac-tivity and the Transporter and telescopeactive sources during operation. In or-der to verify the contribution of the tele-scope internal vibrations to the totalOPL stability budget, a dedicated OPLtest set-up was designed using a com-

bination of laser interferometer and ac-celerometer measurements. The M1-M2 OPL variation was measured bymeans of two high-sensitivity ac-celerometers while the OPL variationfrom M2 to M11 was measured bymeans of an HP Doppler interferome-ter. The results of the OPL test werecombined together with the other con-tributors to the OPL stability as comput-ed by analysis such as the wind loadand the micro-seismic disturbances.The consolidated budget appears to becompliant with the specifications.

5.3 Verification of non-repeatablepointing error

The pointing accuracy is particularlyimportant for a telescope dedicated tointerferometry. Indeed, the standard in-terferometric observing strategy con-sists in switching rapidly (every 1 to 3minutes) between the scientific objectand a calibrator of known diameter thatis used to calibrate the time varying in-strumental and atmospheric visibilitylosses. Poor pointing repeatability re-quiring full acquisition sequences cantherefore seriously affect the efficiencyof the observation. A test was dedicat-ed to measure the non-repeatable partof the AT pointing accuracy for fivepointing directions distributed over theoperational range. A collimator and fivefolding mirrors were installed above thetelescope. The collimated light, sent viaone folding mirror, is focused at thecoudé focus on a CCD. The image po-

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Figure 10: Azimuth axis dynamic response (measured and fittedmodel) used to verify the control bandwidths needed for accurate tele-scope tracking.

Figure 11: Overview of measured non-repeatable pointing error inarcsec on the sky for various angular motions of the altitude and az-imuth axes. The circles represent the requirement specifications forthe various angular ranges.

sition is recorded before and after thetelescope has been moved back andforth by a given angle (from 10 arcmin-utes, up to 60 or 360 degrees) in alti-tude and azimuth directions. Due to thevery small pointing errors to be meas-ured (0.1 to 1 arcseconds), particularattention had to be paid to the thermalenvironment and to the processingmethod in order to eliminate thermaldrifts of the measurement set-up.Figure 11 gives an overview of the ex-cellent results obtained, basically limit-ed by the residual measurement noiseof 0.13 arcsec.

6. Conclusion

This article provides an overviewof the current development of theAuxiliary Telescopes for the VLTI. Themain performance tests of the first tele-scope are completed. They haveshown that the system is on track tomeet the severe and specific require-ments originating from its dedicationto interferometry, in particular a high dy-namic performance. The functionaltests related to the relocation of the tel-escope have been successfully com-pleted on the first Transporter and the

related performance tests are cur-rently in progress. Early next year, thefirst AT will be ready for a short ESOtest period in Liège to tune the tele-scope control system developed inparallel at ESO. The year 2003 willsee the installation at Paranal of thefirst two ATs. This will considerablyboost the scientific productivity of VLTIthat will, at that time, be equippedwith its two first-generation instru-ments: MIDI and AMBER. AT#3 isexpected to be ready for scientificobservation in mid-2004 and AT#4 inearly 2005.

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Paranal Observatory – 2002ROBERTO GILMOZZI and JASON SPYROMILIO

We live in interesting times. Last yearwas our first with all four 8-m tele-scopes in operation. We had only oneinstrument per telescope but we werekept busy with UVES and ISAAC work-ing round the clock and commissioningand installations during bright timewhen FORS1 and FORS2 were takinga breather. Now at the end of 2002 welook at Paranal and a very different pic-ture appears in front of us.

We continued to operate with lowtechnical time losses. In the period fromApril to October 2002 the down timewas 3.6% on Antu, 4.4% on Kueyen,

2.9% on Melipal, and 3.1% on Yepun.Observing time losses due to adverseweather conditions represented 12% ofthe science time, somewhat more thanin previous years. We archived 57,810frames on ISAAC, 34,329 frames onUVES, 22,853 frames on FORS1 and37,847 frames on FORS2.

NAOS/CONICA has been broughtinto operation in a major effort by theconsortia that built it, ESO instrumenta-tion division and Paranal. All partiespulled out all the stops to start scienceoperations on the 1st of October 2002.NACO, as we now call the otherwise

unwieldly named instrument, is widelyrecognized as the most powerful adap-tive optics facility instrument in theworld. Already fabulous results haveappeared in press releases and moreare to come. The spectacular results onthe Galactic centre have demonstratedthe excellent capabilities of the instru-ment. UT4 now has two instrumentsand both dark and bright time are fullyexploited for science.

VIMOS appeared, somewhat over-weight and with all the challenges ofwhat in effect is four imaging, multiob-ject spectrographs mounted on a singleframe. Starting with two arms, it thenhad a leg added to it and another twoarms making the whole thing quite asight to be admired on the Nasmyth Bplatform of UT3. All of this under thecareful supervision of Paranal engi-neering and extensive testing and dis-cussions to ensure that nothing unto-ward occurred as we pushed the rota-tors beyond their specified and testedrange. Again with a lot of work from theconsortium and the support of the in-strumentation division and the everpresent Sandro D’Odorico, we havenow reached the stage of Paranaliza-tion of the instrument and we have afixed date with our customers on the 1stof April 2003. Already the PI (Olivier LeFevre) has claimed success with theexecution of the entire CFR Survey ex-ecuted in a single exposure.

In the mean time FLAMES arrived inparts and has slowly grown to occupythe totality of the Nasmyth A platform ofUT2. First OzPoz, the fibre positionerand then GIRAFFE the spectrograph totake the 130 spectra at medium resolu-tion simultaneously. OzPoz did puzzleus for a while. Mostly it puzzled the con-sortium that built it. After a few inter-ventions on the instrument to makesure it fit to the telescope properly andto make sure it could cope with the realFigure 1: The VST enclosure taking shape at the edge of the Paranal platform.

operational temperatures, it workedwell enough to give us confidence to of-fer the instrument to the communityalso on April 1st 2003. Paranalization isin full swing also for FLAMES and thevisits from Australia to fix it are rare.

FORS1 moved back to UT1 to makedark time available to VIMOS. NowUT1, UT2 and UT4 have dark andbright time instruments attached and ei-ther offered or just about to be. UT3with VIMOS is eagerly awaiting VISIRnext year.

On UT4 we added a laser cleanroom, kindly provided by telescope divi-sion in Garching. This is a whole newplatform under the Nasmyth B focus ofUT4 with a fancy room to house thelaser for the artificial guide star to be in-stalled at the VLT in 2003. We mounteda small telescope on the back on thesecondary of UT4 to check for flexuresand entertained ourselves and MartinCullum trying to work out where thething was pointing. In any case the flex-ures are low and we await the laser andthe fibre to carry the light plus thelaunch telescope during 2003.

This would be enough for most ob-servatories to be going on. However,Paranal is more than a home for four 8-m telescopes. The interferometer con-tinues to grow in capabilities and im-proves its operations. Last year westarted the installation and first runswith the tip-tilt boxes at the coudé foci.Stabilized beams were fed from UT1and UT3 into VINCI and excellentfringes were obtained. In 2002 the sys-tems were stabilized operationally withvery few nights of sky time. The othertwo coudé trains (UT2 and UT4) wereinstalled aligned and tested. The read-er should be aware that 11 mirrors ap-pear in the normal coudé train beforethe beam is sent on the interferometerdelay line tunnel. In September we ob-tained fringes with all four telescopesmarking a major milestone in the devel-opment of the interferometric capabili-ties of Paranal.

In November, MIDI, our first fullyfledged interferometric instrument ar-rived on site and is being integrated asthis article is being written. First lightshould arrive soon. AMBER, FINITOand full coudé adaptive optics on theUTs for VLTI are all in the plan for 2003.

Three additional delay lines are be-ing installed into the tunnels and vari-able curvature mirrors added to thesystem to better control the position ofthe pupil. The VLTI complex is reachinga maturity of operation at Paranal thatis comparable to that of the UTs. Weeagerly await the arrival of the first aux-iliary telescope next year and thepreparations on Paranal are ongoingwith the alignment of the docking sta-tions and the installation of the tracks.

In September of 2002, a set of boxesarrived on Paranal containing suspi-ciously Paranal blue metal pieces. At

the time of writing these pieces arestarting to be placed on top of the VSTconcrete pier for the enclosure of ournext telescope. The enclosure is ex-pected to be finished in early 2003 andthe telescope is to be installed soon af-ter. A few weeks before the enclosureerection begun we received the boxescontaining the VST 2.6-m primary andsecondary mirrors. Unfortunately theprimary mirror was destroyed duringtransport. Our colleagues at the Osser-vatorio di Capodimonte in Naples whoare building the telescope very quicklyordered a replacement mirror and if allgoes well we expect no significant de-lay to the project.

We welcomed the VISTA site super-visor to Paranal and we prepare for thestart of the works to accommodate the4-m infrared wide field telescope.

The smallest telescope on Paranalwas also commissioned. Called Mascot(ask J.G. Cuby) it is an all-sky camerabased on simple commercially avail-able components and with some inhouse ingenuity it makes its imagesavailable in real time on the UT con-soles. The operatorsstill go out to have alook.

Coating mirrors onParanal is a neverending story. How-ever, in 2002 we havereached a peak of ac-tivity. At the time of

writing we are removing the primaryand tertiary from UT1 and plan in thenext few days to have the fourth 8-mcoating this year. Moving these mirrorsaround and going through the cleaningand coating process, along with all thepreventative maintenance that is per-formed on the domes and telescopeswhile the glass is out, is a major engi-neering effort organized to the minuteand synchronized with great precision.The observatory is now undertaking themaintenance of the passive supportsfor the cell in an effort that is supportedby La Silla who are providing the work-shop space.

In January of 2002 we all moved intothe new Paranal residence. The longstay in the containers is finally over andwe now have a pleasant and comfort-able environment in which to live. Workon Paranal is by necessity associatedwith long absences from family andfriends and extreme weather condi-tions, in particular very low humidity.The new residence attempts, and webelieve succeeds, to reduce the hard-ships associated with work on site.

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Figure 2: The primary of UT3 coming out of the coating chamber with a reflectivity of 92%and a micro-roughness of 7 Ångströms. The best coating ever at Paranal.

Figure 3: All-sky imagefrom Mascot Camerawhich monitors thecloud cover on Paranalcontinuously during ob-servations.

Some of the containers are being re-tained as we continue to build onParanal (e.g. VST and VISTA) and of-ten have more people working on sitethan are expected in a steady state ofoperation. The recent MIDI commis-sioning required for the reactivation ofsome containers for ESO and consor-tium staff. Not all were happy with themove back to the not-so-good-old days.

The new Visitor’s Centre was con-cluded and is in use for the weekendvisitor programme, and as a startingpoint for VIP visitors.

The scale of the Paranal operation isoften difficult to appreciate. Ourselveson the mountain find it all fairly normal.To put it into context the ESO Paranal

casino serves approximately 100,000meals every year and we have 40,000overnight stays per year. Paranal is notconnected to the Chilean electricity gridand has to generate its own power(up to three megawatts). Water is ofcourse a necessary resource and wecan store up to 1 million litres in ourtanks. Keeping the observatory sup-plied requires a water tanker truck to ar-rive on site every 8 hours, every day ofthe year.

In 2002, Joerg Eschwey, who morethan any other individual has personi-fied the challenge and success of cre-ating a whole observatory and its infra-structure in the middle of the Atacamadesert, has moved on from the VLT to

the ALMA project. With this move asmall shake up of the observatory or-ganization was necessary and it hasbeen a pleasure to welcome FrankRuseler from the Santiago office ofESO to Paranal where he has takenover the logistics department.

The year ahead of us promises to beas exciting as this past one. VISIR,VST, OmegaCAM, MACAOs for theUTs, FINITO, AMBER, Auxiliary tele-scopes and the laser guide star all in-tend to arrive on Paranal. Times willcontinue to be interesting on Paranalfor the foreseeable future. Excitingnew facilities and capabilities continueto be added to this astronomy wonder-land.

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News from La Silla: Science Operations DepartmentO. HAINAUT

This year has seen a major restruc-turing of the internal workings of LaSilla observatory. While this is not im-mediately obvious to the visiting as-tronomer, it prepares the observatoryfor the future and decreasing staffinglevels. In particular, the engineering andtelescope teams have been reorgan-ized and now constitute two depart-ments:

(1) La Silla Engineering Department(LED). This is a merger of the previousMechanics, Electronics and Instrumen-tation teams. This department is re-sponsible for the maintenance of thetelescopes and projects taking place atLa Silla.

(2) La Silla Science Operations(SciOp). This team actually operates thetelescopes.

The Infrastructure Support Group(ISG), Software and Communications(SWC), Logistics, and Management de-partments all keep the same structureas before.

La Silla SciOp

This is a merger of the three formertelescope teams (NTT, 3.6-CAT andMedium-Sized Telescopes) which wereabandoned in order to optimize humanresources. SciOp currently consists of17 astronomers, 18 TIOs, and 2 opera-tion engineers who are allocated to dif-ferent “Instrument Forces”. The as-tronomers and TIOs work within a spe-cific instrument force to focus their ex-pertise on instruments of a particulartype. For example, the Infrared Instru-ment force consists of SOFI and TIM-MI2. All people working within this forcewill support both of these instruments.

Each force is led by one of theInstrument Scientists, and primary re-sponsibilities include implementing acoherent calibration plan, producingconsistent documentation, observingtemplates, etc., and following-up devel-opments and problems that may occurwith the instruments within the instru-ment force.

The La Silla web pages will be re-structured to reflect the organizationalchanges in the observatory, and in par-ticular to make comparisons betweeninstruments simpler. If you have aquery about a particular instrument,you should contact the correspondinginstrument force (see below). You canalso contact La Silla SciOp by sendingan email to lasilla@eso.org. This ac-count is continuously monitored by theSciOp Shift Leader (one of the as-tronomers on duty) who will then for-

ward your query to the correct person.This email address is the best way tocontact us in order to receive a fast reply.

The new web page of SciOp is avail-able at http://www.ls.eso.org/lasilla/sciops/ . Here you will find links to allthe instrument pages, as well as moreinformation on the new structure.

The RITZ!

For the visiting astronomer, most ofthe above will probably go largely un-noticed. The one big change for visitorsis the opening of the new control build-ing (nicknamed “RITZ”, for RemoteIntegrated Telescope Zentrum). Thisnew, central observing hub of La Silla islocated at the bottom of the NTT ac-cess ramps in front of the “Sarco-phagus” (Figure 1). It is 300 metressquare and will ultimately host the con-

Instrument Force Instruments Contact email

Imaging WFI 2.2-m ls-imaging@eso.orgSUSI2 NTT

Visible Spectro-Imagers EFOSC2 3.6-m ls-spectro@eso.orgEMMI NTT

Infrared TIMMI2 3.6-m ls-infrared@eso.orgSOFI NTTADONIS 3.6-m

High-Res. Spectroscopy FEROS 1.5-m/2.2-m ls-hires@eso.orgCES 3.6-mHARPS 3.6-mEMMI-Echelle NTT

Telescopes NTT, 3.6-m, 2.2-m ls-telescopes@eso.org

trol rooms of the NTT, 3.6-m and 2.2-mtelescopes in a common area, allowingus to operate the telescopes more effi-ciently.

The RITZ was built with the comfortof its inhabitants in mind. Toward thewest, it has many large windows over-looking the valley and ocean, with afantastic view of the sunset (and poten-tial green-flashes). On the other hand,its walls are almost completely blind to-ward the north-east since this is the di-rection from which the wind blows mostof the time. All of the control computersare located in separate rooms behindthe consoles, allowing us to keep thecomputers cold and the control roomtemperate. These computer rooms alsoprovide additional thermal insulation forthe main working area, and keep it freefrom the noise of the many worksta-tions. The floor plan of the main controlroom, as well as the materials chosenfor the walls and ceiling are ergonomi-cally optimized, ensuring that soundsare damped and glare from the win-dows is avoided.

Since October 18, the NTT has beenoperated fully remotely from the RITZ.We hope to move the 2.2-m controlroom this year, and the 3.6-m early nextyear.

News from the Instruments andTelescopes

From the start of Period 70 (start ofOctober), and as reported in previousMessenger articles, the two 1.5-m tele-scopes on La Silla are no longer offeredto the ESO community. The ESO 1.52-m will be operated until the end of 2002,with 100% of the time allocated toBrazilian observers, and the Danish1.54-m will continue to be used by theDanish community only.

A few months ago, the red CCD ofEMMI was upgraded to a mosaic of 2 ×1 MIT/LL 2k × 4k chips. The quantum

efficiency of the mosaic is significantlyhigher than that of the former TektronixCCD and the larger size of the CCDsleads to a wider field of view. In addi-tion, the pixel size (0.166″) gives a bet-ter sampling in good seeing imagesand permits narrower slits to be used.However, the most striking feature ofthe new detector is its read-out time:thanks to the FIERA controller, the 4kchips are read in only 18 seconds, withvery low read-out noise. More informa-tion can be found on the updated EMMIweb pages.

FEROS, the high-resolution, fibre-fed spectrograph that used to bemounted on the ESO 1.52-m telescope,has been transferred to the 2.2-m,where it is now installed in parallel withthe Wide Field Imager. The move tookplace in early October and the instru-ment is currently being commissionedin its new home. We expect that this will

result in an improvement of the spec-trograph’s performance by a factor of 2.Keep an eye on the FEROS pages forup-to-date information.

TIMMI2, the mid-IR spectro-imagermounted at the 3.6-m, is currently un-dergoing a major upgrade: the controlelectronics and software are being en-tirely replaced by an IRACE controllerand VLT-software. This will improve theperformance of the instrument, butmore notably, it will make its operationand maintenance much more straight-forward. Also, for the observer, themain difference will be that TIMMI2 willbe operated from the standard P2PP-BOB interfaces.

Finally, ADONIS, the adaptive opticssystem at the 3.6-m, was retired at theend of August after almost 10 years inservice. The field of high resolution im-aging is now carried by NAOS-CONICAon the VLT.

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Two unusual views of La Silla. The left-hand picture shows La Silla immediately after the snow storm of 28 August 2002. Some weeks lat-er the desert around La Silla was for a short time transformed into a sea of flowers. Photos by Peter Sinclaire.

Introduction

The widely accepted picture of stellarformation tells us that a planetary sys-tem is a simple by-product of the stellarformation process. When a cloud of gasand dust contracts to form a star, con-servation of angular momentum in-duces the formation of a flat diskaround the central newborn “sun”. By aprocess still not fully understood, thisdisk is believed to be the stage for theplanetary formation. According to thetraditional paradigm, dust particles andice grains in the disk are gathered toform the first planetary seeds (e.g.Pollack et al. 1996). In the “outer” re-gions of the disk, where ices can con-

densate, these “planetesimals” arethought to grow in a few million years.When such a “planetesimal” achievesenough mass (about 10 times the massof the Earth), its gravitational pull is suf-ficiently strong for it to start accretinggas in a runaway process that givesorigin to a giant gaseous planet similarto the outer planets in our own SolarSystem. Later on, in the inner part ofthe disk, where temperatures are toohigh and volatiles cannot condensate,silicate particles are gathered to formthe telluric planets like our Earth.

In the past decade, images taken bythe NASA/ESA Hubble Space Tele-scope (HST) have revealed a multitudeof such proto-planetary disks in the

Orion stellar nursery, showing thatdisks are indeed very common aroundyoung solar-type stars. This supportsthe idea that extra-solar planets shouldbe common. However, such systemshave escaped detection until very re-cently.

In fact, it was not until 1995, followingthe discovery of the planet orbiting thesolar-type star 51Peg May95, that thesearch for extra-solar planets had itsfirst success1. This long wait was main-ly due to the difficulty of detecting suchbodies. Planets are cold objects, andtheir visible spectrum results basicallyfrom reflected light of the parent star. Asa result, the planet/stellar luminosity ra-tio is of the order of 10–9. Seen from adistance of a few parsec, a planet is nomore than a small “undetectable”speckle embedded in the diffractionand/or aberration of the stellar image.

The detection of exoplanets has thusbeen based, up to now, upon “indirect”methods. In particular, all the planetarydiscoveries were only possible due tothe development of high-precision radi-al-velocity techniques. These methods,that measure the motion of a star alongthe line of sight (by measuring theDoppler shift of spectral lines), havenow achieved precisions of the order ofa few m s–1 (∆λ/λ ∼ 10–8). Such a highprecision is indeed needed to find aplanet: for example, Jupiter induces aperiodic perturbation with an amplitudeof only 13 m s–1 on the Sun!

In this article we will review the cur-rent status of planetary searches, pre-senting the major challenges that weare facing at this moment. We will thendiscuss how new and future generationinstruments and missions will help toanswer the most important questions.We will concentrate mostly on the re-sults we can expect from future radial-velocity campaigns with state-of-the-artinstruments like HARPS on the ESO3.6-m telescope (see article by Pepe etal. in this issue).

A Diversity of Planets

Today, about 100 extra-solar plane-tary systems have been unveiled

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REPORTS FROM OBSERVERS

Extra-Solar Planets N.C. SANTOS, M. MAYOR, D. QUELOZ, S. UDRY

Observatoire de Genève, Sauverny, Switzerland

Figure 1: Schematic orbital configurations for some of the newly found extra-solar planets inthree different scales. In the upper-left panel we represent the orbits of the shorter-periodcompanions. The Sun (yellow circle) is drawn to scale. This plot illustrates well both the prox-imity of these planets to their host stars, and the complete lack of planetary companions or-biting closer than a certain distance (see text for more details). The upper-right and lower pan-els illustrate the situation concerning longer orbital period planets. In these two plots, the or-bits of the Earth and Jupiter are also drawn for comparison (with the usual symbols). Thesethree panels clearly illustrate the huge variety of orbital parameters presented by the extra-solar planets.

1Before this discovery, only planets around apulsar had been detected (Wolszczan & Frail 1992);however, these are probably second-generationplanets. In this article we will concentrate on plan-ets around solar-type stars.

Figure 2: Distribu-tion of minimummasses for the cur-rently discoveredlow-mass compan-ions to solar-typestars. Although theradial-velocitymethod has a higher sensitivity tohigher-mass com-panions, theobserved distribu-tion rises verysteeply towards thelow-mass domain.From this mass upto the stellarregime, only a fewobjects weredetected; thisregion is usuallydenominated the“Brown-Dwarfdesert”. This gap inthe mass distribu-tion of low-masscompanions to so-lar-type stars supports the view that the physical mechanisms involved in the formation ofthese two populations (planets vs. stars) are very different.

around stars other than our Sun2.These discoveries, which include ∼ 10multi-planetary systems, have broughtto light the existence of planets with ahuge variety of characteristics, openingunexpected questions about theprocesses of giant planetary formation.

The diversity of the discovered extra-solar planets is well illustrated inFigure 1. Unexpectedly, they don’thave much in common with the giantplanets in our own Solar System.Contrarily to these latter, the “new”worlds present an enormous and unex-pected variety of masses and orbitalparameters (astronomers were basical-ly expecting to find “Jupiters” orbiting at∼5 A.U. or more from their host stars inquasi-circular trajectories). The majori-ty of the discovered planets were noteven supposed to exist according to thetraditional paradigm of giant planetaryformation (e.g. Pollack et al. 1996).Their masses vary from sub-Saturn toseveral times the mass of Jupiter.Some have orbits with semi-major axessmaller than the distance from Mercuryto the Sun, and except for the closestcompanions, they generally follow ec-centric trajectories, contrarily to thecase of the giant-planets in the SolarSystem.

These findings have put into questionthe former planetary formation para-digm. However, the relatively large num-ber of discovered planets is alreadypermitting us to undertake the first sta-tistical studies of the properties of theexoplanets, as well as of their host stars.This is bringing new constraints to themodels of planet formation and evolu-tion. Let us then see in more detail whatkind of problems and information thesenew discoveries have brought.

The Period Distribution

One of the most interesting problemsthat appeared after the first planetswere discovered has to do with theproximity to their host stars. The firstplanet discovered (around 51Peg –Mayor & Queloz 1995) is exactly thefirst such example: it orbits its star onceevery 4.2 days, corresponding to an or-bital radius of only 0.05 A.U. This ismuch less than the distance fromMercury to the Sun.

The problem that was raised with thefinding of these 51Peg-like planets(usually called “hot-Jupiters”) resides inthe fact that at such close distances thetemperatures are too high for ices tocondensate, and there does not seemto exist enough available material toform a Jupiter-mass planet. It is thusvery difficult to imagine that suchworlds could be formed so close to thecentral stellar furnace.

In order to explain the newly foundsystems, several mechanisms havethus been proposed. Current resultsshow that in situ formation is very un-likely, and we need to invoke inward mi-gration, either due to gravitational inter-action with the disk (Goldreich &Tremaine 1980) or with other compan-ions to explain the observed orbital pe-riods. In other words, the observedclose-in planets could simply havebeen formed far from their host star,and then migrated inwards.

But the migration mechanisms, thathave broken long-lasting ideas of “sta-bility” of the planetary systems, havesome problems to solve. According tothe models, the timescales of planetarymigration in a disk are particularlyshort. This means that more than wor-rying about how planets migrate after orduring their formation, we need to un-derstand how migration can be stopped(and/or slowed down)!

One particularly interesting cluecomes from the observation that thereis a clear pile-up of planetary compan-ions with periods around 3 days, ac-companied by a complete absence ofany system with a period shorter thanthis value. This result, which is in com-plete contrast with the period distribu-tion for stellar companions (we can finddouble stars with periods much shorterthan 3 days), means that somehow theprocess involved in the planetary mi-gration makes the planet “stop” at a dis-tance corresponding to this orbital peri-od. To explain this fact, several ideashave been presented. These invoke dif-ferent mechanisms like e.g. a magne-tospheric central cavity of the accretiondisk (once the planet gets into this cav-ity it will no longer strongly interact with

the disk and consequently stops the mi-gration), photo-evaporation, tidal inter-action with the host star, or Roche-lobeoverflow of the young inflated giantplanet (processes resulting in an in-crease of the orbital radius of theplanet, thus opposing the migrationtendency)3 .

In any case, even if these mecha-nisms are able to explain how the short-er period planets have stopped migra-tion, they do not explain how the longerperiod ones (like Jupiter itself) did notmigrate to distances closer to their hoststars. The key to this might have to dowith a combination of parameters, likethe disk masses, lifetimes, viscosities,and initial planetary masses and/ornumber of bodies formed, that will af-fect the final orbital configuration of aplanet. Some of these parameters are notwell known (maybe the planetary for-mation itself is controlling them), some-thing that complicates the discussion.

The Mass Distribution

Very important information is broughtto us by the analysis of the mass spec-trum of the planetary companions. Inparticular, a plot showing the mass dis-tribution of companions to solar-typestars (as shown in Figure 2), shows aclear discontinuity for the mass regimebetween about 30 and 50 times themass of Jupiter: there are basically nocompanions found to date having thosemasses. This result is even more strik-ing if we note that the radial-velocity

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2For a complete and updated list of the knownexoplanets, see e.g. table at http://obswww.unige.ch/exoplanets.

3Planets migrating more than this approximatelimit might “simply” also evaporate/transfer materi-al to the host star, and thus disappear or becometoo low-mass to be detected.

Figure 4: Eccentricityvs. orbital period (inlogarithmic scale) forthe discovered extra-solar planets (redpentagons), stellarbinaries (filled dots),and for the giantplanets in our SolarSystem (green stars).The earth is repre-sented by the usualsymbol. Althoughsome long-period ex-oplanets exist havinglow values for theeccentricity, most ofthe systems presentmuch higher eccen-tricities than thoseobserved for theSolar System giantplanets. Possible ex-planations invokemechanisms capableof pumping theeccentricity, likegravitational interactions between planets in a multi-planetary system, or with a distant stellar companion.

Figure 3: Distributionof exoplanet masses.The histogram repre-sents the observedminimum mass distri-bution, while the redline represents thestatistical true plane-tary mass distributionresulting from a de-convolution of the un-known orbital inclina-tions. The distributionreaches “zero” at amass of about 10times the mass ofJupiter, which probablycorresponds to theupper limit for the massof a giant planet. Thenature of the objectswith masses between10 and ∼17 times themass of Jupiter is stillan open question. Asin Jorissen et al.(2001).

technique is more sensitive to massivecompanions than to their lower masscounterparts.

This gap, usually called the “browndwarf desert”, separates the low mass“planetary” companions from their highmass “stellar” counterparts, and isprobably telling us something very im-portant about the physical processesinvolved in the formation of these twopopulations: stars, even the low massones, are thought to be formed as theresult of the gravitational collapse andfragmentation of a cloud of gas anddust. On the other hand, a planet formsin a circumstellar accretion disk.

More information is provided if weanalyse the shape of the distribution forthe planetary mass regime (Fig. 3).This distribution is observed to de-crease smoothly with increasing mass,reaching “zero” at about 10 Jupitermasses (Jorissen et al. 2001). This lim-it is clearly not related to the Deuterium-burning mass limit of ∼ 13 MJup, some-times considered as the limiting massfor a planet (this latter value is in fact anarbitrary limit used as a possible “defi-nition”, but it is not related to the plane-tary formation physics). As it was re-cently shown by several authors (e.g.Jorissen et al. 2001), this result is notan artefact of projection effects (the un-known orbital inclination implies that wecan only derive a minimum mass for thecompanion from the radial-velocitymeasurements), but a real upper limitfor the mass of the planetary compan-ions discovered so far.

The Period-Mass Relation

Recent results have also unveiledsome interesting correlations betweenthe planetary mass and its orbitalperiod. In fact, there seems to exist apaucity of high-mass planetary com-

panions (M > 2 MJup) orbiting in shortperiod (lower than ∼ 40-days) trajecto-ries. Similarly, there seems to be a lackof long-period and very low-massplanets.

These results are further helping usto understand the mechanisms of plan-et migration, since they are compatiblewith the current ideas about planetaryorbital evolution (either due to an inter-action with the disk or with other com-panions), that suggest that the higherthe mass of a planet, the more slowly(and less) it will migrate.

The Eccentricity

One of the most enigmatic resultsfound to date has to do with the analy-

sis of the orbital eccentricities of theplanetary-mass companions. Accord-ing to the traditional paradigm of plane-tary formation, a planet (formed in adisk) should keep a relatively circular(low eccentricity) trajectory. Currentmodels shown indeed that the interac-tion (and migration) of a low mass com-panion within a gas disk has the effectof damping its eccentricity (Goldreich &Tremaine 1980). However, opposite toexpectations, if we look at the eccen-tricities of the planetary companions wecan see that they are spread throughvalues that go from nearly zero to morethan 0.9 for the planet orbiting the starHD80606!

In Figure 4 we plot the eccentricity asa function of orbital period for the plan-etary companions to solar-type stars,as well as for the stellar companions.First of all, it is important to describe thegeneral tendencies observed in theplot. The low eccentricity observed forshort-period binaries is the result of awell known effect: the proximity to theirprimary stars induces tidal interactionsthat have the effect of damping the ec-centricity. Since the tidal effect de-creases very fast with distance, abovea given orbital period (about 10 days fordwarf star binaries), tidal circularizationis no longer effective, and all compan-ions having periods longer than a givenvalue simply keep their “initial” orbitaleccentricity.

While both distributions show the sig-nature of tidal effects on the eccentrici-ty, a first glimpse also tells us that thereis no clear difference between the twogroups of points: stars and planetshave a similar distribution in this dia-gram. This poses the problem of un-derstanding how planetary companionsformed in a disk can have the same ec-

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centricity distribution as their stellarmass counterparts. And how can thisbe fit into the picture of a planet forming(and migrating) in a disk?

The explanation for these facts maybe other processes capable of excitingthe eccentricity of the planetary orbits.These include the interaction betweenplanets in a multiple system or betweenthe planet and a disk of planetesimals,the simultaneous migration of variousplanets in a disk, or the influence of adistant stellar companion. All or at leastsome of these physical processesmight play an important role in definingthe “final” orbital configuration.

Although still not clear, a close inspec-tion of Figure 4 permits us to find a fewdifferences between the eccentricitiesof the stellar and planetary compan-ions. In particular, for periods in therange of 10 to 30 days (clearly outsidethe circularization period by tidal inter-action with the star), there are a fewplanet hosts having very low eccentric-ity, while no stellar binaries are presentin this region. The same and evenstronger trend is seen for longer peri-ods, suggesting the presence of agroup of planetary companions with or-bital characteristics more similar tothose of the planets in the Solar Sys-tem (with low eccentricity and longperiod). On the other hand, for thevery short period systems, we can seesome planetary companions with ec-centricities higher than those found for“stars”. This features may be telling usthat different formation and evolutionprocesses took place: for example, theformer group may be seen as a sign forformation in a disk, and the latter oneas an evidence of the influence of alonger period companion on the eccen-tricity.

Clues from the Planet Hosts:the Stellar Metallicity

Up to now we have been reviewingthe results and conclusions obtained di-rectly from the study of the orbital prop-erties and masses of the discoveredplanets. But another fact that is helpingastronomers understand the mecha-nisms of planetary formation has to dowith the planet host stars themselves.Indeed, they were found to be particu-larly metal-rich, i.e. they have, on aver-age, a higher metal content than thestars without detected planetary com-panions (see Santos et al. 2001 for themost recent results) – see Figure 5.

A possible and likely interpretation ofthis may be that the higher the metallic-ity of the cloud that gives origin to thestar/planetary system (and thus thedust content of the disk), the faster aplanetesimal can grow, and the higherthe probability that a giant planet isformed before the proto-planetary diskdissipates. In other words, the metallic-ity seems to be playing a key role in the

formation of the currently discoveredextra-solar planetary systems.

Boss (1997) has suggested that be-sides the core accretion scenario (seeintroduction), giant planets might alsobe formed as a result of disk instabilityprocesses: by the formation and con-densation of clumps of gas and dust inthe protoplanetary disk. This processis, however, not very dependent on themetallicity. In other words, if the disk-in-stability models were the most impor-tant mechanism involved in the forma-tion of giant planets, we should not ex-pect to see a strong dependence on therate of planet detection as a function ofthe metallicity. The huge dependenceobserved is thus probably a sign thatthe core accretion scenario is the im-portant mechanism involved in the for-mation of giant planets.

It is, however, important to stress thatit is not precisely known how the metal-licity is influencing the planetary forma-tion and/or evolution; for example, themasses of the disks themselves, whichcan be crucial to determine the efficien-cy of planetary formation, are notknown observationally with enoughprecision.

A Case of Stellar “Cannibalism”

Recent observations also suggestthat planets might be engulfed by theirparent stars, whether as the result oforbital migration, or e.g. of gravitationalinteractions with other planets or stellarcompanions. Probably the clearest evi-dence of such an event comes from thedetection of the lithium isotope 6Li in theatmosphere of the planet-host star HD82943 (Israelian et al. 2001). This frag-ile isotope is easily destroyed (at only1.6 million degrees, through (p,α) reac-tions) during the early evolutionarystages of star formation. At this stage,the proto-star is completely convective,and the relatively cool material at the

surface is still deeply mixed with the hotstellar interior (this is not the case whenthe star reaches its “adulthood”). 6Li isthus not supposed to exist in stars likeHD 82943, and the simplest and mostconvincing way to explain its presenceis to consider that planet(s), or at leastplanetary material, have fallen into HD82943 sometime during its lifetime.

The most recent and detailed analy-sis seems to clearly confirm the pres-ence of this isotope. The question is thenturned to know whether this case is iso-lated or if it represents a frequent out-come of the planetary formation process.How much can this process increasethe observed metallicity of the planethosts? Current results suggest that atleast the degree of stellar “pollution” isnot incredibly high (Santos et al. 2001).

Black Sheep

When measuring the spectrum of astar we are obtaining the integratedlight of the whole stellar disk, and gath-ering photons coming from differentpoints in the stellar surface. Each indi-vidual point has its own spectrum, witha different Doppler shift that is a func-tion of the velocity field in that specificregion of the stellar photosphere. As aconsequence, any phenomenon capa-ble of changing the velocity field of agiven region in the stellar surface willchange the global spectrum Dopplershift, and consequently the measuredradial-velocity.

This result has an important impactwhen dealing with radial-velocity meas-urements: the radial-velocity techniqueis not sensitive only to the motion of astar around the centre of mass of astar/planet system, but also to eventualvariations in the structure of the stellarsurface.

In fact, phenomena such as stellarpulsation, inhomogeneous convectionor the presence of dark spots (e.g. Saar

35

Figure 5: Left: Metallicity ([Fe/H]) distributions for stars with planets (yellow histogram) com-pared with the same distribution for field dwarfs in the solar neighbourhood (open red his-togram). In this panel, both distributions are normalized by the total number of points. The[Fe/H] scale is logarithmic, and the Sun has by definition [Fe/H] = 0. Most planet hosts aremore metal-rich than our Sun; Right: the percentage of stars that have been found to harboura planet for each metallicity bin, plotted as a function of the metallicity. This plot clearly showsthat the probability of finding a giant planet increases with the metallicity of the star. As inSantos et al. (2001).

& Donahue 1997) are expected to in-duce radial velocity variations. Further-more, contamination from the spectrumof other stellar companions can also in-duce spurious radial-velocity signals.These cases can prevent us from find-ing planets (if the perturbation is largerthan the orbital radial-velocity varia-tion), but perhaps more importantly,they might give us false candidates ifthey produce a periodic signal (e.g. arotating stellar spot). The radial-veloci-ty technique is thus most efficient forlow chromospheric activity single “old”dwarfs.

A few good examples of spurious pe-riodic radial-velocity variations can befound in the literature. The first to bedescribed was the periodic signal ob-served for the dwarf HD 166435, thatwas shown to be due to the presence ofa dark spot rather than to the gravita-tional influence of a planetary compan-ion.

At the current (and increasing) levelof precision obtained using the radial-velocity techniques, such kinds of ex-amples might become more and morecommon. It is thus very important to de-velop ways to disentangle e.g. activity-related phenomena from real planetarycandidates. Such methods might bebased on the study of the shape of thespectral lines (usually called the “bisec-tors”), the photometric variability of thestar, and/or of chromospheric activityindexes. This is very important for proj-ects like HARPS, for which the barrierof the 1 ms–1 will be achieved.

Guidelines for the Future

The study of extra-solar planetarysystems is just beginning. After only 7years, we can say that at least 5% ofsolar-type dwarfs have giant planetarycompanions in relatively short periodorbits (≤ 3 years). However, and as wehave seen in the previous sections,many interesting but troubling problemsstill await a solution. In fact, the newly-found planets have clearly disturbedthe long-standing theories of giant plan-etary formation and evolution. The def-

inition of a planet itself is currently un-der debate.

To solve this problem, there are highexpectations from new or soon-to-beavailable instruments (see the paper onthe HARPS spectrometer in this issue).The incredible precision gain achievedby these radial-velocity “machines” willbe crucial in this aspect. But what areexactly the important lines of study tofollow? And what kind of answers canwe expect in the next few years?

Increasing the Statistics

As we have seen above, the ob-served correlations between the orbitalparameters of the newly found planetsare giving astronomers a completelydifferent view on the formation and evo-lution of planetary systems. We nolonger have the Sun as the only exam-ple, and today we have to deal with thepeculiar characteristics of the “new” ex-tra-solar planets: a huge variety of peri-ods, eccentricities, masses.

To clear up current uncertainties weneed more data. Only a large and sta-tistically uniform set of data may enableus to clarify the current situation. In thissense, the hundreds of new planets ex-pected to be discovered in the next fewyears will have a very important out-come.

Lower and Lower Masses

One clear result of the increase inprecision of the current and future radi-al-velocity surveys is the ability to findlower-mass planets.

As we have seen above, the massfunction for planetary companionsaround solar-type stars is rising towardthe low-mass regime. But given that theradial-velocity technique is more sensi-tive to the more massive companions,the lower mass bins in the plot ofFigure 2 are definitely not well repre-sented. How then does this function be-have for the lower-mass regime? Whatis the minimum mass of a giant planet?How does this depend on the orbital pe-riod?

The increase in long-term precision,and the continuation of the currenthigh-precision radial-velocity pro-grammes, will also give us the opportu-nity of finding more and more long-pe-riod planets. Current models predictthat the planetary formation and evolu-tion processes should produce morelong-period planets than their short pe-riod counterparts. In fact, the currentsurveys are just starting to discoverexo-planets with periods comparable tothe ones found for the Solar System gi-ants. This is an essential goal in orderto improve the statistics of these ob-jects, and to check if the Solar Systemis anomalous or common.

Of course, other problems will ariseand become more important as a con-sequence of the dramatic increase inprecision. In particular, the intrinsic stel-lar radial-velocity “jitter” produced bye.g. chromospheric activity related phe-nomena (e.g. Saar & Donahue 1997)might impose serious (and still un-known) limits on the final precision, andon the consequent ability to find verylow-mass (and long-period) systems.This is particularly true for the young-est stars. Some effort should thusbe put into the development of diag-nostics capable of confirming the plan-etary nature of the radial-velocity sig-nal. Furthermore, we might even imag-ine that the spurious radial-velocityvariations caused by activity might bemodelled and corrected, leaving onlythe real gravitational effects on thesignal.

Planets Around M Dwarfs

Although more than 100 exoplanetcandidates are now in the lists, only twoplanetary companions around M dwarfshave been detected (both in the solelysystem Gl 876). The very low numberof M-dwarf planetary companions canin fact be largely explained by observa-tional biases: the very low mass dwarfsare faint and it is difficult to obtain ac-curate radial velocities for them.

However, to constrain the variousscenarios of planetary-system forma-tion and evolution, it is now crucial toobtain better statistics for planetsaround the most numerous stars in ourgalaxy. M dwarfs compose 80% of themain sequence stars. How many ofthem have planets? How these planetsdiffer from those orbiting the moremassive G dwarfs is totally unknown.These questions await the future capa-bilities of instruments like HARPS.

The Chemical Link

Planet host stars seem to be, on av-erage, particularly metal-rich. This in-teresting result probably reflects the im-portance the quantity of available rockymaterial in the disk has on the forma-tion of giant planets.

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Figure 6: Radial-ve-locity measurementsand best 2-planetKeplerian solution forHD 82943. The radi-al-velocity measure-ments show that thisstar has a system oftwo resonant plan-ets, orbiting it withperiods of ∼ 220 and440 days, respec-tively. This same starwas found to have6Li in its atmosphere,a signature of possi-ble planet engulf-ment (Israelian et al.2001).

As discussed above, this link mighthold the key to understanding how gi-ant planets are formed. The two com-peting theories (core accretion and diskinstability) should have different sensi-tivities to the metallicity. If most planetswere formed by the latter of these twoprocesses we should not expect anyspecial metallicity sensitivity.

Although the metallicity trend isclearly seen (Santos et al. 2001), thereare nevertheless a few planet hosts thatare particularly metal poor (see e.g. caseof HD 6434). How can this be explained?One elegant way of solving this puzzleis to look for the frequency of planetsaround metal-poor halo dwarfs. If, asexpected, no planetary mass objectsare found around such a sample ofstars, then the disk instability model isclearly put into question. However, ifsome are found, giant planets might beformed by different processes.

There are also some traces of stellar“pollution” among the planet host stars(Israelian et al. 2001). This opens thequestion of understanding how muchplanetary material might fall into theconvective envelope as a consequenceof the planetary formation process it-self. How much can this change themetallicity of the star? If important, thiscould have consequences even forstudies of the chemical evolution of thegalaxy.

Although current results seem to re-fute any strong generalized stellar pol-lution among planet hosts, it is impor-tant to cross-check. One interestingway of addressing this problem mightthen be the study and comparison ofthe chemical contents of stars in visualbinary systems composed of similarspectral type solar-type stars, one (orboth) of which having planetary-masscompanions. Strong differences foundwould be interpreted as a sign that stel-lar pollution is quite common. However,only one clear case has been studied todate: the double visual system 16 CygA and B, where the B component isknown to harbour a planet. This casedoes not show any clear difference inthe iron abundance, while a curiouslithium abundance difference is found.

Transit Candidates

Another important result of instru-ments like HARPS will be their ability tofollow up planetary-like transit signa-tures detected by photometry.

There are currently more than 20groups around the world trying to lookfor signatures of the presence of plan-ets around field dwarfs by looking forthe brightness dimming as a putativeplanet crosses its disk. In spite of theefforts, only very few results have beenannounced, and none of these wasconfirmed to have a planetary origin.The only clear case of a real planetarytransit detected so far was found for the

star HD 209458, a dwarf that was pre-viously discovered to host a very-shortperiod planetary companion (Charbon-neau et al. 2000).

This detection, and the subsequentrelated studies, have had an enormousimpact for the understanding of thesesystems. For example, it was possibleto estimate the mean density of theplanet, and to prove that it is orbiting inthe same direction and plane as thestar’s equator.

In the near future we can expect thatother such events might be brought tolight. In particular, many hundreds ofphotometric transit candidates are ex-pected from space missions likeCOROT, Kepler or Eddington. How-ever, based only on photometry wecannot determine whether the ob-served transiting body is a planet orsimply a low-mass star (since the effectis of similar magnitude because of thelarge degeneracy of the radius of theseobjects). The follow-up of the photo-metric observations by radial-velocitysurveys is thus essential and will permitus to obtain the real mass of the planet.

With such data we can hope to deriveempirical relations between variablessuch as the planet’s mean density andits distance from the star, its mass, andthe stellar metal abundance. In thissense it is important to say that the veryhigh precision of HARPS, together withthe relatively large aperture of the ESO3.6-m telescope, will play an importantrole, since it will give the opportunity toobtain masses (or at least meaningfulupper limits) for the least massive planetsdetected by the photometric missions.

Multiple Systems: DynamicalInteractions

Among the many planets that are ex-pected to be found, some will surely be-long to multi-planetary systems. Today,only about 10 such cases are known,but many stars that are already knownto harbour a planet also show system-atic trends in radial velocity, indicatingthat at least a second companion ispresent in the system. While for the ma-jority of cases this tendency might besimply due to the presence of low-massstellar companions, in some othersthey might be the telltale signatures ofa multi-planetary system. The gain inprecision with instruments like HARPSwill definitely permit us to search the al-ready known planet hosts for otherplanet-mass companions, and to in-crease the number of known multiplesystems.

There is in fact much interesting in-formation that can come from thesecases. Current results have shown thatplanets in multiple systems come fre-quently in resonant orbits (see e.g.HD82943 – Figure 6). This is telling usa lot about the formation and migrationof the exoplanets.

On the other hand, the strong inter-action between planets in such sys-tems will be reflected as an observableevolution of their orbital parameters. Adynamical analysis of this will give usthe opportunity to obtain information onthe masses and relative orbital inclina-tions for the companions.

Planets in Binaries

To date, several planets have beendiscovered in known multiple stellarsystems. Moreover, a fraction of starsknown to host planets exhibit a drift inthe γ-velocity indicating the presence ofan additional distant companion.

These observations show that de-spite the gravitational perturbation ofthe stellar companion, planets may formand survive around stars in multiple sys-tems. The properties of such planetshold important clues on the mecha-nisms of planetary formation. For ex-ample, according to the standard coreaccretion model of planetary formation,a giant planet is formed by the accre-tion of gas around a ∼ 10 earth masscore of rocky material. This is supposedto take place at distances comparableto the Jupiter–Sun separation (∼5 A.U.).Opposing this model, Boss (1997) hasproposed that giant gaseous planetsmight also be formed from the conden-sation of clumps resulting from gravita-tional instabilities in the disk. How canwe distinguish these two scenarios?

One of the keys may come from thestudy of planets in binaries. The pres-ence of a stellar companion possiblyplays an important role in the formationof planets. It has been shown, for ex-ample, that a stellar companion cantruncate the proto-planetary disk at aradius that depends mainly on the dis-tance between the two stars. If so, andconsidering the core accretion sce-nario, we should not be able to seeplanets around stars members of bina-ry systems that are closer than a givenlimit. How close can a star have a com-panion and still have planets? The an-swer to this question is very importantto understand how giant planets areformed.

Concluding Remarks

As the planet search programmesare on their way, many more planetarycompanions are expected to be discov-ered in the next few years. Many hopesare now placed on instruments likeHARPS, that will provide radial veloci-ties of stars with a precision of 1 ms–1

or better. This will give us the opportu-nity to dramatically improve the sam-ples.

Other major contributions will comefrom future space missions likeCOROT, Eddington, or Kepler, whichwill unveil thousands of short-periodplanets around stars in the solar neigh-

37

bourhood. And, of course, the use ofhigh-precision astrometric measure-ments with instruments like the VLTI orthe Keck interferometer will survey“nearby” stars for long-period systems.Altogether, these coming observationalfacilities will definitely help us to con-struct a new and more complete view ofhow planetary systems are born andhow they evolve.

References

Boss A.P., 1997, Science 276, 1836. Charbonneau D., Brown T.M., Latham D.W.,

Mayor M., 2000, ApJ 529, L45. Goldreich P., Tremaine S., 1980, ApJ 241,

425. Israelian G., Santos N.C., Mayor M., Rebolo

R., 2001, Nature 411, 163. Jorissen A., Mayor M., Udry S., 2001, A&A

379, 992.

Mayor M., Queloz D., 1995, Nature 378,355.

Pollack J.B., Hubickyj O., Bodenheimer P.,Lissauer J.J., Podolak M., Greenzweig Y.,1996, Icarus 124, 62.

Saar S.H., Donahue R.A., 1997, ApJ 485,319.

Santos N.C., Israelian G., Mayor M., 2001,A&A 373, 1019.

Wolszczan A., Frail D. A., 1992, Nature 355,145.

38

FIRES: Ultradeep Near-Infrared Imaging withISAAC of the Hubble Deep Field SouthI. LABBÉ1, M. FRANX1, E. DADDI3, G. RUDNICK2, P.G. VAN DOKKUM4, A. MOORWOOD3, N.M. FÖRSTER SCHREIBER1, H.-W. RIX5, P. VAN DER WERF1, H. RÖTTGERING1, L. VAN STARKENBURG1, A. VAN DE WEL1, I. TRUJILLO5, andK. KUIJKEN1

1Leiden Observatory, Leiden, The Netherlands; 2MPA, Garching, Germany; 3ESO, Garching, Germany; 4Caltech, Pasadena (CA), USA; 5MPIA, Heidelberg, Germany

1. Introduction

Between October 1999 and October2000 an undistinguished high-galacticlatitude patch of sky, the Hubble DeepField South (HDF-S), was observedwith the VLT for more than 100 hoursunder the best seeing conditions. Usingthe near-infrared (NIR) imaging mode

of the Infrared Spectrometer and ArrayCamera (ISAAC, Moorwood 1997), weobtained ultradeep images in the Js(1.24 µm), H (1.65 µm) and Ks (2.16µm) bands. The combined power of an8-metre-class telescope and the high-quality wide-field imaging capabilities ofISAAC resulted in the deepest ground-based NIR observations to date, and

the deepest Ks-band in any field. Thefirst results are spectacular, demon-strating the necessity of this deep NIRimaging, and having direct conse-quences for our understanding ofgalaxy formation.

The rest-frame optical light emittedby galaxies beyond z ∼1 shifts into thenear-infrared. Thus, if we want to com-pare 1 < z < 4 galaxies to their present-day counterparts at similar intrinsicwavelengths – in order to understandtheir ancestral relation – it is essentialto use NIR data to access the rest-frame optical. Here, long-lived starsmay dominate the total light of thegalaxy and the complicating effects ofactive star formation and dust obscura-tion are less important than in the rest-frame ultraviolet. This therefore pro-vides a better indicator of the amount ofstellar mass that has formed. Com-pared to the selection of high-redshiftgalaxies by their rest-frame UV light,such as in surveys of Lyman BreakGalaxies (LBGs, Steidel et al. 1996a,b),

Figure 1: Three-colour composite image ofthe ISAAC field on top of the WFPC2 main-field, outlined in white, and parts of threeWFPC2 flanking fields. The field of view isapproximately 2.5 × 2.5 arcminutes andNorth is up. The images are registered andsmoothed to a common seeing of FWHM ≈0.46″, coding WFPC2 I814 in blue, ISAAC Jsin green and ISAAC Ks in red. There is astriking variety in optical-to-infrared colours,especially for fainter objects. A number ofred sources have photometric redshifts z > 2and are candidates for relatively massive,evolved galaxies. These galaxies would notbe selected by the U-dropout technique be-cause they are too faint in the observer’s op-tical.

selection in the NIR Ks-band gives amore complete census of the galaxiesthat contribute to the stellar mass con-tent of the early universe. And, bystudying these systems over a sub-stantial redshift range, we can directlysee how they were assembled as welook further back in time. Very deepoptical-to-infrared data in many filtersare required not only to access therest-frame optical light of galaxies andconstrain their stellar composition, butalso to determine the redshifts offaint distant galaxies from their broad-band photometry alone. The majori-ty of the galaxies detected are toofaint to be observed spectroscopically,even with powerful telescopes likethe VLT.

Here, we discuss the full NIR data setof the HDF-S. The observations, reduc-tion techniques, the catalogue ofsources, and the photometric redshiftsare described in detail in Labbé et al.(2002). Throughout, we will assume aflat Λ-dominated cosmology (ΩM = 0.3,Λ = 0.7, H0 = 100h km s–1 Mpc–1) andall magnitudes are expressed in theJohnson photometric system unlessexplicitly indicated by the subscript ABfor the AB magnitude system.

2. Observations

The observations of the HDF-S arepart of the Faint InfraRed ExtragalacticSurvey (FIRES, Franx et al. 2000), alarge public programme carried out atthe VLT consisting of very deep NIRISAAC observations of two selectedfields with existing deep optical WFPC2imaging. The second and somewhatshallower field around the z ≈ 0.83 clus-ter MS1054-03 (Förster Schreiber et al.in preparation) was observed for 80hours over an area four times largerthan the HDF-S. The full NIR data setof the HDF-S consists of 33.6 hours ofimaging in Js, 32.3 hours in H, and 35.6hours in Ks with ISAAC, which has a2.5′ × 2.5′ field of view. The effectiveseeing in the reduced images is ap-proximately 0.46″ in all bands, close tothe best seeing that can be reasonablyobtained from Paranal. We reach a depthof 25.9 in Js, 24.8 in H, and 24.4 in Ks(total magnitude for point sources, 3σ).Complemented by ultradeep HST/WFPC2 imaging (version 2, Casertanoet al. 2000) in the optical filters U300,B450, V606, and I814 bands (the subscriptindicating the effective wavelength), weassembled a Ks-selected cataloguecontaining 833 sources, of which 624have seven-band photometry, covering0.3–2.2 µm. We determined the photo-metric redshifts of all extragalacticsources using a method detailed inRudnick et al. (2001, 2002a in prepara-tion). Comparison with spectroscopicredshifts available for 49 sources in thefield shows excellent agreement: |zspec– zphot | / (1+zspec) ≈ 0.08.

The colour image shown in Figure 1combines HST I-band data with ISAACdata in the Js and Ks bands. Mostsources visible are very faint distantgalaxies and a rich variety in their opti-cal-to-infrared colours is readily appar-ent. Some high-redshift galaxies, visi-ble as red sources, are hardly detectedor even absent in the optical I-band butare quite bright in the NIR Ks.

3. Evolved Galaxies at HighRedshift

Remarkably, the HDF-S containsmany high-redshift sources that are rel-atively bright in the infrared and ex-tremely faint in the optical, whereas theHDF-N contains far fewer such galax-ies. This is clearly visible in Figure 2,which shows V606 – H colours versusthe infrared H band magnitudes of NIRselected galaxies between redshifts 2 <z < 3.5. The filters are chosen to matchthose used in Figure 1 in Papovich,Dickinson & Ferguson (2001), allowinga direct comparison with the HDF-N.For a similar survey area and similarlimiting depth, we find 7 galaxies redderthan V606,AB – HAB > 3 and brighter thanHAB < 25 (on the AB magnitude sys-tem), compared to one in the NIR sam-ple of HDF-N. These systems are prob-able candidates for relatively massive

and evolved galaxies with comparative-ly low star-formation rates, and modelsof galaxy formation must allow for theirformation in sufficient numbers at z ∼ 3.It appears that in recent hydrody-namical simulations (Weinberg, Hern-quist & Katz 2002) high-redshift galax-ies are converting gas into starscontinuously and at a fairly high rate,keeping them blue or “unevolved”.While this is qualitatively consistentwith the lack of red galaxies in the HDF-N, the situation in the HDF-S is signifi-cantly different.

All high-redshift galaxies with veryred V606 – H colours also have very redinfrared Js – Ks colours. Redshiftingmodel spectra of galaxies shows thatsuch red colours are produced byevolved systems at redshifts z > 2. Inour Ks-band selected catalogue we find13 relatively bright galaxies (Ktot < 22)that have red Js – Ks > 2.3 colours(Franx et al. in preparation), comparedto 37 U-dropouts to the same flux limit.The photometric redshifts of thesegalaxies, determined from the optical-to-infrared spectral energy distributions(SEDs), are between 2 < zphot < 4. Yet,most of the red galaxies are so faint inthe observers optical that they would bemissed by standard optical selectioncriteria, such as the U-dropout method.In principle, the red colours can be

39

V606

– H

Figure 2: V606 – H versus H colour-magnitude diagram (on the AB system) for galaxies in theHDFS Ks-selected catalogue with 1.95 < zphot < 3.5. Filled symbols indicate galaxies withspectroscopy and the data points are grouped into two redshift ranges: 1.95 ≤ z < 2.7 (red)and 2.7 ≤ z ≤ 3.5 (blue). We find 7 galaxies redder than V606,AB – HAB > 3 and brighter thanHAB < 25, compared to only one in the HDF-N. Galaxies with S/N < 2 for the V606 measure-ment (triangles) are plotted at the 2σ confidence limit in V606, indicating a lower limit on theV606 – H colour. Galaxies having red (Js – Ks)J > 2.3 colours (open squares) are also shown.The number of candidates for red, evolved galaxies is much higher than in the HDF-N for asimilar survey area, as shown in an identical plot in Figure 1 of Papovich, Dickinson &Ferguson (2001). The transformation of the V606 – Hs colour from the AB system to theJohnson magnitude system is (V606 – H)J = (V606 – H)AB + 1.26.

Figure 4: The angu-lar two-point correla-tion function forgalaxies with red Js– Ks > 1.7 coloursbetween redshifts of2 and 4, with thebest fitting power-law(solid line) and 1-sigma errors (dottedlines). A positive andlarge clustering sig-nal is detected, sug-gesting that thesered FIRES galaxiesare stronglyclustered in realspace, with a corre-lation length r0 = 8.3± 1.2 h–1 Mpccomoving. This isthe largest level ofclustering ever foundfor distant, z > 2galaxy populations.

caused by dust, by contribution ofprominent emission lines falling in theKs-band, or by the redshifted Balmer or4000 Å break, which indicates the pres-ence of evolved stars. Although emis-sion lines have been detected in spec-troscopically confirmed Js – Ks > 2.3galaxies in the FIRES MS1054-03 fieldby van Dokkum et al. (in preparation),their contribution to the broadband NIRphotometry is estimated to be small. Inmany cases the spectral energy distri-butions show a clear break between Jsand Ks, which is more easily explainedas an aging effect than as a result ofdust reddening. These galaxies con-tribute significantly to the rest-frame op-tical luminosity density (Rudnick et al.2002a in preparation) and from theirred rest-frame optical colours – imply-ing high mass-to-light ratios – we esti-mate that they contain a substantialfraction of the total stellar mass presentin all galaxies at z ∼ 3.

The morphologies of the red galaxiesare generally very compact at all wave-lengths, as can be seen in Figure 3. Anotable exception is the galaxy inFigure 3c which is clearly extended andhas an exponentially decreasing sur-face brightness profile out to 1″ radiusin the H-band. This galaxy seems tohost blue knots which could form a spi-ral arm in the optical WFPC2 images,while it has an extended disk-like ap-pearance with a prominent bulge in thecentre in the ISAAC H and Ks images;possibly these are actively star-formingsites embedded in a larger mature hostgalaxy.

4. Strong Clustering of FaintRed Galaxies at 2 < zphot < 4

Presumably, high-redshift galaxieswith red Js – Ks colours are among theoldest and most massive galaxies attheir cosmic epoch and they haveformed at the highest density peaks inthe matter distribution at significantlyearlier times. If so, this populationshould be more clustered than lessmassive and less evolved (bluer) ob-jects at similar redshifts. Using theFIRES data of the HDF-S, Daddi et al.(2002) have studied the clustering be-

haviour of Ks-selected galaxies at 2 <zphot < 4, finding that the amount ofclustering depends strongly on the Js –Ks colour of galaxies , with red galaxiesmore clustered than blue galaxies, verysimilar to what is observed in the opti-cal in the local universe. Dividing thesample at Js – Ks = 1.7 in two subsam-ples reveals that the galaxies with Js –Ks > 1.7 colours have the largest-everlevel of clustering measured for z > 2galaxies (see Fig. 4). The derived cor-relation length for the faint red galaxiesis r0 = 8.3 ± 1.2 Mpc (comoving): a fac-tor of 3–4 higher than that of Lyman

40

Figure 3: Morphology of bright red Js – Ks > 2.3 galaxies in the HDF-S. The left grayscale panels show the averaged WFPC2 V+ I data, andright-hand panels show our VLT/ISAAC Js, H, Ks data. The intensity is proportional to Fλ, with arbitrary normalization for each galaxy. Thecolour images show a combination of I, Js, Ks data, after matching the image quality to that of the Ks-band. The right-hand column shows thespectral energy distributions. Many of these galaxies are small and show prominent breaks in the infrared (rest-frame optical). Note that thegalaxy in the top row is barely visible even in Js. The SED of the object in the middle is consistent with a strong Balmer/4000Å break at z ∼2.5. The galaxy in the bottom row is very extended in the rest-frame optical. It shows faint emission in the H-band out to 1″, consistent withan exponential profile.

Break galaxies over similar redshiftranges and with similar number densi-ties. The overall properties of thisred J – K > 1.7 strongly clustered pop-ulation suggests that they are the pro-genitors of present-day massive ellipti-cals and extremely red objects (EROs)at z ∼ 1.5.

5. Large Disk-Like Galaxies atHigh Redshift

The high-redshift galaxies in HDF-Sshow a large variation in both mor-phologies and spectral energy distribu-tions. Spectacularly, we find some z > 2galaxies that are very large in the rest-frame optical and show profound differ-ences between the intrinsic ultravioletand optical morphologies (Labbé et al.in preparation). Three of the galaxies,shown in Figure 5, are LBGs (selectedby applying the U-dropout criteria ofMadau et al. 1996), while one is thepreviously described red galaxy inFigure 3c. Their spectral energy distri-butions show a pronounced break inthe rest-frame optical, identified as theage-sensitive Balmer/4000 Å break.They are amongst the brightest andreddest in the rest-frame optical andthey are probably the most massivegalaxies at z ∼ 2–3.

The I-band morphologies (probingrest-frame 2000–2700 Å) are complexwith a patchy distribution that is notconcentrated towards the centre.Conversely, the NIR light is compactand peaks right in the centre of the I-band structure, surrounded by a diffuseand extended disk-like distribution. Themost straightforward explanation is that

these systems are very luminous giantspiral galaxies, with patchy star forma-tion residing in a diffuse disk with a redbulge in the centre. It is well possiblethat we have discovered the predeces-sors of the large disk galaxies we seein the lower-redshift universe. Mor-phological studies in the rest-frame UVlight of LBGs have always emphasizedthe compact and small sizes of the gal-axies (Giavalisco et al. 1996, Lowen-thal et al. 1997), as indeed is true forthe bulk of our sample. However, galax-ies that are large in the rest-frame opticalwith classical spiral morphologies havenever been seen before, in particular notin the HDF-N (Dickinson et al. 2000)even though such structures wouldhave been well recognizable in theirdeep WPFC2 and NICMOS imaging.

Follow-up VLT/FORS spectroscopy(Rudnick et al. 2002b in preparation)confirmed the redshift of the galaxy inFigure 5c, a disk-like U-dropout at zspec= 2.793 ± 0.003, implying a scale lengthof 10 h–1kpc. Vanzella et al. (2002) foundan identical redshift for this galaxy andfor another, shown in Figure 5a, theyfound zspec = 2.027. Theories of diskformation generally predict that largedisk galaxies have been assembled re-cently (z < 1), and that high-redshift z >2 disks are small and dense, with typi-cal sizes of ∼ few kpc. Producing largedisk-like galaxies with reasonable likeli-hood may therefore pose a challenge tothe current hierarchical models.

6. Summary

These first results on the HDF-Sdemonstrate the necessity of extending

deep optical observations to near-IRwavelengths for a more complete cen-sus of galaxies in the early universe. Inparticular, the deepest-ever Ks-banddata have proven to be invaluable,probing well into the rest-frame opticalat 2 < z < 4, where long-lived stars maydominate the light of galaxies.

Provided with a new window on theearly universe we find previously undis-covered populations of galaxies, withpossible far-reaching consequences forour understanding of galaxy formation.We find that the HDF-S has many moregalaxies at z ∼ 3 with very red V606 – Hcolours than the HDF-N; these are can-didates for relatively massive evolvedsystems. Closely connected, we find asubstantial population of red Js – Ks >2.3 galaxies at z > 2, many of which arebarely detectable even in the deepestoptical images, and would be missedby optical colour selection techniquessuch as the U-dropout method. Yet,these galaxies probably contribute sub-stantially to the total stellar mass pres-ent in the early universe at z ∼ 3. Wedetect strong clustering (r0 = 8.3 ± 1.2Mpc, comoving) of galaxies with Js – Ks>1.7 colours at 2 < zzphot < 4. The strongclustering and their red Js – Ks colours,suggest that there is a direct evolution-ary link between these systems, EROsat z ∼1.5, and massive elliptical galax-ies in the local universe. We also findhigh-redshift systems that are verylarge in the rest-frame optical withmorphologies similar to those ofnearby giant spiral galaxies, containinga red bulge surrounded by a diffusedisk and with scattered patches of star-formation. While these are only a few

41

Figure 5: Same as Figure 3 for the three largest Lyman-break galaxies in the HDF-S. The K-band images (rest-frame optical) are more cen-trally concentrated than the WFPC2 images; this is a real change in light distribution and not caused by the decreased image quality in theNIR. The redshifts of galaxies (a) and (c) have been confirmed spectroscopically; they have z = 2.027 and z = 2.793 respectively, implying ascale length of 10h–1kpc for the latter.

examples, it is currently not clearwhether galaxy formation models canproduce objects of comparable coloursand sizes in sufficient numbers to beconsistent with our observations.

Lastly, the substantial differencesbetween HDF-N and HDF-S demon-strate that results based on such smallfields can be seriously affected bycosmic variance. The results of thesecond field in the FIRES survey, themuch larger MS1054-03 field, shoulddecide which one of the fields is atypi-cal. We are pursuing extended follow-up programmes to obtain more spec-troscopic confirmation of the above re-sults, allowing us to fully investigatethe nature of these galaxies and theclues they provide for models of galaxyformation. Updates on the FIRES pro-gramme and access to the reduced im-ages and catalogues can be found atour website http://www.strw.leidenuniv.nl/~fires.

Acknowledgements

We wish to thank the ESO staff foroperation of the VLT and ISAAC andtheir kind assistance. We are gratefulfor their enormous efforts in obtainingthese data in service mode and makingthem available to us.

References

Moorwood, A.F.M., “ISAAC: a 1-5 µm imag-er/spectrometer for the VLT”, in Opticaltelescopes of today and tomorrow. A.L.Ardeberg ed., proc. SPIE 2871, pp.1146–1151, 1997.

Steidel, C.C., Giavalisco, M., Dickinson, M., &Adelberger, K.L., AJ 112, pp. 352-358, 1996.

Steidel, C.C., Giavalisco, M., Pettini, M.,Dickinson, M., & Adelberger, K.L., ApJ462, L17, 1996.

Rudnick, G. et al., AJ 122, pp. 2205–2221,2001.

Labbé, I.F.L. et al., AJ accepted, 2002.Förster Schreiber, N.M. et al., in preparation,

2002.

Oke, J.B. 1971, ApJ, 170, 193. Franx, M. et al., “FIRES at the VLT: the Faint

InfraRed Extragalactic Survey”, TheMessenger 99, pp. 20–22, 2000.

Casertano, S. et al., AJ 120, pp. 2747–2824,2000.

Papovich, C., Dickinson, M., & Ferguson,H.C., ApJ 559, pp. 620–653, 2001.

Weinberg, D.H., Hernquist, L., & Katz, N.,2002, ApJ, 571, 15.

Daddi, E. et al., ApJ, submitted, 2002.Madau, P., Ferguson, H.C., Dickinson, M.E.,

Giavalisco, M., Steidel, C.C., & Fruchter,A., MNRAS 283, pp. 1388–1404, 1996.

Giavalisco, M., Steidel, C. C., & Macchetto,F.D., ApJ 470, 189, 1996.

Lowenthal, J.D. et al., ApJ 481, pp.673–688, 1997.

Rudnick, G. et al., in preparation, 2002a.Rudnick, G. et al., in preparation, 2002b.Dickinson, M., “The first galaxies: structure

and stellar populations” Philos. Trans. R.Soc. London A 358, p. 2001, 2000.

Dickinson, M. et al., ApJ 531, pp. 624–634,2000.

Vanzella, E. et al., A&A accepted, 2002.

42

OTHER ASTRONOMICAL NEWS

Summary of the Workshop on

EXTRAGALACTIC GLOBULAR CLUSTER SYSTEMShosted by the European Southern Observatory in Garching on August 27–30, 2002

By M. KISSLER-PATIG

Extragalactic Globular Clusters,a Booming Field of Research

Globular cluster systems were estab-lished in the last decade as powerful

tools for the study of galaxy formationand evolution. For this purpose they areused in nearby galaxies with as muchsuccess as the diffuse stellar popula-tions and complement the latter studies

by being superior in several practicalaspects. For instance, globular clustersare better chronometers than the dif-fuse stellar population since each glob-ular cluster can be identified as a sin-gle-age population. The age determina-tion of the major globular cluster sub-populations allows one to preciselydate the star formation events in thehost galaxy.

Further, the study of globular clustersystems is the easiest way to detectmultiple old or intermediate-age sub-populations within a galaxy. These inturn can trace multiple major star-for-mation episodes at early times, invisi-ble in the studies of the diffuse galaxylight in which all populations are mixed.

Another advantage is that globularclusters can be traced far out in thehalo, probing stellar populations andkinematics at several effective radii ofthe host galaxy – a range inaccessibleto diffuse stellar light studies. The glob-ular clusters can be used for dynamicalstudies at outer radii, shedding light onthe assembly histories of the systemsand the host galaxies.

Finally, globular clusters are verygood calibrators at all wavelength forsingle stellar population models, an es-

Figure 1: NGC 6946 – a nearby spiral with a large number of very luminous young stellar clus-ters. Taken from Larsen’s contribution. The observations were made with the ALFOSC in-strument on the Nordic Optical Telescope on La Palma. ALFOSC is a twin instrument ofDFOSC on the Danish 1.54-m and mounted on the NOT; it has a field size of about 6 × 6 arc-min. The colour image was generated from a mosaic of 3 pointings with red channel = V + I+ Hα; green channel = V; blue channel = B.

sential tool for all galaxy formation andevolution studies.

In summary, globular clusters are avery strong complement to diffusegalaxy-light studies. They are verygood tracers of the major star-formationepisodes of galaxies, including extremeconditions such as star-burst during vi-olent merger events. The study of thesesystems allows us to learn not onlyabout the formation of the globular clus-ters and their systems, but also aboutthe formation and evolution of galaxies.

The ESO Workshop, a Follow-upon the Pucon IAU Symposium

Last year’s IAU Symposium 207 on“Extragalactic Star Clusters” in Pucon(the first IAU symposium to be held inChile) highlighted preliminary resultsfrom some of the currently ongoing, ex-tensive programmes to study globularcluster systems. The size of the meet-ing did not lend itself to interactive dis-cussions. However, the talks presentedat this Symposium underlined the dra-matic progress currently underway inthis field.

This motivated us to follow it up ayear and a half later with a workshop ona more focused aspect, namely “Extra-galactic Globular Cluster Systems”.The meeting was held at the headquar-ters of ESO in Garching, profiting fromthe freshly renovated Auditorium. About60 participants enjoyed 4 days of livelydiscussion and excellent presentations.The low mean age of the researchers inthis dynamic field spoke for the timeli-ness of the research. Participants whopioneered the research a couple ofdecades ago (such as W.E. Harris, andD. Hanes) expressed several times thejoy of seeing the field blossom, even ifit implied that they are now regarded as“dinosaurs” by the young graduatesand post-graduates.

The workshop programme wasplanned such that no less than 2 hoursfor lunch break and another one hour (e.g. the presentation of Enrico Ves-

perini).Young star clusters are, however, ex-

tensively observed and studied as awhole session demonstrated. The maincurrent result is that these young clus-ters get better and better connectedwith their old counter-parts and that lit-tle doubt remains that a vast number ofthe observed massive young clusterswill evolve into objects similar to theold globular clusters observed today,although their initial properties mightcontain surprises (e.g. S. Mengel’s re-view).

Early- and late-type giant galaxies re-main the favourites for studying clustersystems as a whole. Spectroscopicstudies with 8- to 10-m telescopes pro-duce impressive results both on thekinematics of the outer regions ofgalaxies (e.g. the many hundred veloc-ities collected in the outskirts of NGC

43

Figure 2: Deep mosaic of the Centaurus galaxy clusterobtained with FORS1 on the VLT (V and I filters, seeing0.6″, reaching V = 26 mag). Taken from M. Hilker's con-tribution.

and a half for coffee breaks were avail-able each day. With 10–15 minutes dis-cussion after each talk, this workshopwas a very interactive event. The shortweek culminated in a final discussion(served with wine and cheese) that wenton for several hours and only finishedas the beverage supply finally ran out.

Summary and Highlights of thePresentations

The author of this summary apolo-gizes for the injustice of not citing all in-teresting contributions but only somehighlights of the meeting.

At the beginning of the week, wewere reminded by Keith Ashman thatglobular cluster formation and destruc-tion are still not well understood, al-though numerical simulations are mak-ing steady progress in explaining theevolution of globular cluster systems

Figure 3: The smoothed density map of globular cluster candidates shows clear over-densi-ties outside the isophotes of the major galaxies. Centaurus must host a significant populationof free floating, intra-galaxy globular clusters.

1399 by Tom Richtler and his group), aswell as on the chemistry of the old clus-ters (e.g. Puzia et al.). A number ofstudies also reported on large interme-diate age populations of globular clus-ters in early-type galaxies, detectedboth by a combination of optical andnear-infrared photometry as well as byspectroscopy (e.g. Hempel et al.,Goudfrooij et al.). Interestingly, theglobular clusters in dwarf galaxies look“almost but not quite unlike” their count-er-parts in giant galaxies as JenniferLotz showed us. Despite the similarity,populations in dwarfs can extend tovery low metallicity.

Interestingly (even if not unexpected-ly), globular clusters are found in largenumbers in the intergalactic mediumof galaxy clusters. Michael Hilker andcollaborators show clear evidence forglobular clusters floating through theHydra I and Centaurus galaxy clus-ters.

Finally, globular clusters, as the old-est objects known in the sky, are nowbeing put in a cosmological context.Mike Beasley investigated the con-straints that they put on semi-analyticmodels in the hierarchical clusteringscenario of galaxy formation. MichaelSantos presented a model in which the

old metal-poor globular clusters formedprior to re-ionization.

Clearly, extragalactic globular clus-ters currently dominate the study ofstellar populations in nearby galaxies.Their properties set hard constraints ongalaxy formation and evolution models,but also shed new light on the star for-mation history of the universe. This re-search area will clearly produce a numberof astonishing results in the next 5 years.

The organizers would like to warmlythank Christina Stoffer who perfectlycontrolled the logistics of the confer-ence, as well as Pam Bristow for help-ing in the proceedings editing.

44

The VLTI: Challenges for the Future WORKSHOP AT JENAM 2002 IN PORTO

A. GLINDEMANN, ESO

On the verge of becoming a majorscience facility, the VLTI was the sub-ject of one of the workshops at thisyear’s Joint European and NationalAstronomical Meeting in Porto (Portu-gal) in September. The two and a halfday workshop had the aim of introduc-ing interferometry and the VLTI to thenon-expert. About 65 participants,many of them young and from outsidethe interferometric community (the twoessentials for success) showed thelarge interest in interferometry.

After an introduction to interferome-try, to the VLTI and its instruments, andto the two other large interferometers –

the Keck Interferometer and the LargeBinocular Telescope – on the first day, thesecond day was devoted to science. In16 presentations, stellar astrophysics wasaddressed including accretion disks andoutflows of young stellar objects, surfacestructure and circumstellar envelopes ofMira stars, diameter of Cepheids, andmass loss of Wolf-Rayet stars. The pos-sibilities of observing nuclear regions ofAGNs with the VLTI were also explored.The emphasis of the presentations wason the possibilities with the VLTI. Someof the presentations contained resultsof interferometric observations withsmaller interferometers. The report on

Mira was based on new observationsdone with the VLTI and made publicthrough the regular VLTI releases.

On the last day, the discussion on thefuture of the VLTI and on projects forsecond-generation instrumentation con-cluded the workshop.

This workshop was a success in at-tracting the non-interferometrists in theastronomical community and in sum-ming up the scientific topics to be tack-led in the near future. It became evidentthat it is now important to finally makethe observations and produce the re-sults that have been discussed over thelast years.

Summary of a Meeting on Science Operations withALMA, held on Friday, 8 November 2002P. SHAVER (ESO) and E. VAN DISHOECK (Leiden)

With the recent approvals by theESO Council and the US NationalScience Board for the construction ofthe Atacama Large Millimeter Array(ALMA), it was thought timely to updatethe European astronomical communityon the project and to solicit input on theplans for science operations and usersupport. To this end, a one-day meetingwas held at ESO Headquarters inGarching on Friday, 8 November. It wasvery well attended; the auditorium wasfilled to capacity with some 100 partici-pants from all over Europe.

The meeting began with an overviewof the project and its current status byS. Guilloteau. Three very stimulatingtalks reviewed some of the major sci-ence drivers for ALMA: the high-redshift

Universe (S. Lilly), star and planet for-mation (A. Natta), and late-type stars(H. Olofsson).

The meeting then moved on to con-sider operations: concepts and plansfor the operations phase were outlinedby E. van Dishoeck, D. Silva talkedabout the relevant operational lessonsfrom the VLT, and R. Lucas discussedthe ALMA data reduction software andobserving tools. After lunch P. Cox dis-cussed the coordination of the Europeanastronomical community and prepara-tion for ALMA, including the opportuni-ties within the EU Framework 6 pro-gramme. The discussion session thenstarted with seven short contributionson a variety of topics, which led on to avery stimulating and useful open dis-

cussion chaired by J. Richer and E. vanDishoeck. Topics of discussion includ-ed the role and nature of a possible Eu-ropean Regional Support Centre (RSC)which could assist users in the obser-vation preparation and data analysisprocesses, and ideas and priorities con-cerning software, receiver bands, sur-veys, future enhancements, and prepa-ration for ALMA science. In view of thesuccess of this meeting, there will prob-ably be more such meetings for thecommunity as the project evolves.

The viewgraphs from the invited talksgiven at the meeting are posted on theALMA website, http://www.eso.org/proj-ects/alma/doclib/talks/, where more de-tails on the ALMA project can also befound.

ESO has always been a forward-looking organization, focussed on pres-ent and future programmes and proj-ects. In its 40th year, this is as true asever, as ESO keeps its eyes on imple-menting the VLTI, initiating the devel-opment of the second-generation in-strumentation for the VLT and, ofcourse, embarking on Phase II of theALMA project, for which the ESOCouncil gave the green light at its mid-year meeting in London. Still, a 40th an-niversary warrants a look at thedecades passed and a reflection on therole and achievements of the organiza-tion during a period of time in which ourscience has progressed tremendously.That ESO has achieved its goals isundisputed, providing evidence of thewisdom of those astronomers whogathered in Leiden in early 1954 to for-mulate the need for this organizationand justifying the trust displayed by thefive countries, which were the originalsignatories to the 1962 convention thatformed the legal base for ESO.

There were many milestones onESO’s road to success. Not surprising-ly, many of them were recalled – andretold – during the various celebrationsin September and October this year,described in the Anniversary issue ofThe Messenger, as well as, undoubted-ly, in numerous conversations amongmembers of the ESO ‘family’, be it for-mer or present staff, members of com-mittees or Council – or users in the sci-entific community.

Rather than organizing a single,high-profile central celebration event,ESO opted for a series of events, be-ginning with the official visit on 9

September by Dr. Uwe Thomas, Secre-tary of State for education and researchof the German Federal Government.Speeches by Dr. Thomas, Prof. RalfBender (LMU), Prof. Wolfgang Hille-brandt (MPA), ESO’s Director Generaland Dr. Arno Freytag, President of theESO Council, were followed by ascreening of a (pre-premiere) version ofthe anniversary film ‘Europe Reachesfor the Stars – Forty Years ESO’. Also,the anniversary book ‘GeheimnisvollesUniversum – Europas Astronomenentschleiern das Weltall’ by Germanscience journalist Dirk Lorenzen wasshown for the first time with hand-

bound copies presented to Dr. Thomas,Dr. Freytag, Mr. Manfred Solbrig, theMayor of Garching, and to ESO’sDirector General.

On 14 October, a joint celebration bythe Garching staff and the ESO Counciltook place, with speeches by ESO’sCouncil President, the Director Generaland Claudio Cumani, Chairman of theESO International Staff Association.The event was ‘broadcast’ via ESO’svideo-conference system. The anniver-sary film was shown to the entire staff,both in Garching and at simultaneousanniversary parties in Santiago, at LaSilla and Paranal.

As mentioned in the last issue of TheMessenger, ESO’s anniversary is alsoa public event. Perhaps most visible isthe new planetarium show, producedby the Association des planétariums delangue française (APLF) in collabora-tion with ESO. It is entitled ‘Mys-teries of the Southern Sky’ and will be-come available in French, German,English, Dutch, Spanish and possiblyother languages. Until now, about 40planetaria in France and seven othercountries have signed up and will stagethe show in the coming weeks andmonths.

The formal premiere took place on 11October at the Cité de l’Espace inToulouse. The premiere event for an in-vited audience encompassed scientifictalks by Drs. Françoise Combes andMarc Lachieze-Rey and an evening re-ception accompanied by live LatinAmerican tunes.

Meanwhile, the ESO anniversary filmis available on a three-language DVD(English, French, German), as well as

45

Celebrating ESO’s 40th AnniversaryC. MADSEN, ESO

The newly refurbished auditorium in Garching was packed during the anniversary celebrationon 14 October.

Dr. Uwe Thomas addressing an invited audience during his visit to the ESO Headquarters on9 September.

in single-language versions on VHStapes. Among other places, the DVDwill be widely available through thoseplanetaria staging the ESO show.

Finally, the CD-ROM (in English) en-titled ‘3D Atlas of the Universe’, is now

scheduled for early December. TheCD-ROM, which is produced in a col-laboration between ESO, USM (Mu-nich) and Planetary Visions (UniversityCollege, London), includes video clips,photos and texts about ESO as well as

some of the most spectacular astro-nomical images based on observationsat La Silla and Paranal combined withadvanced 3D simulation software. It willbe commercially available all overEurope.

46

ESO’s Director General addressing the invited audience at the pre-miere of the planetarium show in Toulouse.

In spite of the decidedly unpleasant weather in Garching on 14October, good catering contributed to the convivial atmosphere in theparty tent, as about 250 people gathered for the dinner for staff andCouncil.

Breaking the Ground for the EuropeanResearch Area – The Conference ‘EuropeanResearch 2002’C. MADSEN, ESO

In the days November 11–13, about9000 scientists, science administratorsand policy makers gathered in Brusselsto attend the Launch Conference forthe ‘6th Framework Programme of theEuropean Community for research,technological development and demon-

stration activities’ – or for short, ‘FP-6’.While most participants came from themember states of the European Union,candidate countries and associatedstates, the meeting was in fact attend-ed by people from 65 countries,demonstrating the wide scope and the

importance of the process set in motionto create the European Research Area.Some 50 TV teams and 230 journalistsfrom the print media covered the event,which El País, the leading Spanishnewspaper, described as ‘The ScienceSummit in Brussels’. The strong mediainterest bears witness to the fact thatscience and technology (and with them,also education) are playing an increas-ingly important and visible role in thepublic sphere and that the organizationand execution of research, as well asthe exploitation of scientific results, areassuming importance in the main-stream political debate.

Winds of Change

The background for the strong partic-ipation to the meeting is certainly to befound in the notion of a change of thepolitical climate in Europe with respectto science and technology. The Lisbondeclaration by the European Council inMarch 2000, stated that the EU should‘become the most competitive and dy-namic knowledge-based economy inthe world, capable of sustainable eco-nomic growth with more and better jobsand greater social cohesion’. It was re-The ESO stand formed a part of the EIROforum area.

inforced by the decision at theBarcelona summit in March 2002 to aimfor a 50% increase in research spend-ing in Europe, to reach the 3 % level ofthe GDP by the year 2010. One of themost important elements in the effort torevive European competitiveness inR&D is the concept of the EuropeanResearch Area (ERA) and the 6th

Framework Programme is explicitly de-scribed as a tool to help facilitate theformation of the ERA.

Clearly, creating the ERA is a com-plex process with a diversity of actorsand stakeholders. The conference pro-vided a forum for discussions among allgroups involved, from science ministersto the scientists themselves, whoturned up in huge numbers.

The subjects under discussion in-cluded ‘Research and Innovation’, ‘Hu-man resources and mobility’, ‘Euro-pean research in a global context’,‘Infrastructures for the European Re-search Area’ as well as ‘Science andSociety’. In parallel, ‘how-to’ seminarsgave practical advice for scientistsplanning to seek project funding underthe new Framework Programme.

Furthermore, the ‘Participants’ forum’comprised more than 80 workshopsand events covering a wide spectrum oftopics – from fusion research, ge-nomics, ‘e-science’ and grid technolo-gies to intellectual property rights andscience communication.

EIROforum and the ERA

One of the first tangible results of theERA process was the collaboration be-tween the European IntergovernmentalResearch Organizations (CERN, EFDA-JET, EMBL, ESA, ESO, ESRF, ILL), lead-ing to the formation of the EIROforum.EIROforum has been active for a whileand already achieved significant, prac-tical results and the collaboration wasformally sealed at the conference inBrussels. On Tuesday, November 12,the Directors General of the organiza-tions signed the EIROforum Charter inthe presence of Philippe Busquin, Euro-pean Commissioner for Research, and nu-merous journalists. ‘The establishmentof EIROforum is a concrete example of

the dynamic created by the EuropeanResearch Area. Europe has unques-tioned excellence in science. By work-ing together, Europe’s leading researchorganizations can make that more visi-ble on the European and world stage,’said Busquin in a press statement.

The signing of the charter was alsomarked by the launch of the EIROforumweb site (www.eiroforum.org) and abrochure describing the aims and thespecific elements of the collaboration.

EIROforum at the ConferenceIndeed, EIROforum maintained a high

visibility at the conference. Representa-tives of the member organizations tookpart in several round-table discussionsand workshops, and the EIROforum ex-hibition attracted many visitors and pro-vided inspiring surroundings for manydiscussions. The EIROforum stand cov-ered 400 square metres with an open-plan architecture, providing individualspace for each of the seven memberorganizations and a centrally locatedcommon area for the EIROforum itself.The press event in connection with thesigning of the EIROforum Charter tookplace directly on the stand.

The ESO StandFeaturing prominently on the 36

square metres ESO stand was a scale

model of the 100-m OWL telescope. Itcaught the attention of many visitorsand was also shown in ‘The SixthSense’, the daily conference news-paper. Apart from the OWL project,subjects covered on the ESO standwere the VLT/VLTI and the ALMAproject as well as those programmesfor which the European Commissionhas provided major financial support,e.g. for Adaptive Optics, the Astro-physical Virtual Observatory (AVO) andthe educational outreach programmes,notably ‘Life in the Universe’ and‘Physics on Stage’. Clearly, exposingESO and its activities and futureprojects to such a huge, but select au-dience, was important, and the sevenESO staff that took turns on the standhad many interesting exchanges withthe visitors.

‘European Research 2002’ sent aclear signal that the European re-search landscape is undergoing a ma-jor transformation, leading Le Figaro,the French daily, to talk about ‘a newspirit for European scientists’. In twoyears time, when FP-6 reaches itshalf-term, the follow-up meeting, an-nounced by Commissioner Busquin,will surely be able to determine towhat extent this new spirit has pro-duced tangible progress for Europeanscience.

47

The EIROforum stand was packed with attendants as the EIROforum charter was signed.

A Nobel Prize for Riccardo GiacconiRiccardo Giacconi has been awarded the Nobel Prize„for pioneering contributions to astrophysics, which haveled to the discovery of cosmic X-ray sources”. In thewords of The Royal Swedish Academy of Sciences, „hedetected for the first time a source of X-rays outside oursolar system and he was the first to prove that the uni-verse contains background radiation of X-ray light. Healso detected sources of X-rays that most astronomersnow consider to contain black holes. Giacconi construct-ed the first X-ray telescopes, which have provided uswith completely new – and sharp – images of the uni-

verse. His contributions laid the foundations of X-ray astronomy.”(http://www.nobel.se/). Over the last two decades he has also been a leader inother fields of astronomy, and served as ESO’s Director General over the years1993–1999. ESO and its community send Riccardo their congratulations onthis great honour.

For the last twoyears, I have workedas a fellow at theSantiago offices of theEuropean SouthernObservatory (ESO)with support responsi-bilities at the Very

Large Telescope (VLT) on Cerro Para-nal in the north of Chile. This is es-

PhD at the University of Cambridge inthe summer of 2000.

48

Agreement Between the Government of theRepublic of Chile and ESO for Establishing a NewCentre for Observation in Chile – ALMAReproduced from ESO Press Release 18/02 (23 October 2002)

On October 21, 2002, the Minister ofForeign Affairs of the Republic of Chile,Mrs. María Soledad Alvear, and theESO Director General, Dr. CatherineCesarsky, signed an Agreement thatauthorizes ESO to establish a newcentre for astronomical observation inChile.

This new centre for astronomical ob-servation will be for the Atacama LargeMillimeter Array (ALMA), the largestground-based astronomical project forthe next decades.

On this occasion, Minister Alvearstated that “we want to have ALMAworking as soon as possible, which willconstitute a pride not only to Chileanscientists but for the whole country andin particular, for the community of theAntofagasta Region”.

ESO Director General Cesarsky saidthat “signing this agreement betweenthe Government of Chile and ESOis a historical step in the astronomical

collaboration between Chile and ESOand it will allow Chile to host, once

again, a project ofworldwide interestand impact”.

ALMA is a jointproject on equalbasis betweenESO and AUI (As-sociated Universi-ties, Inc.). Theseorganizations rep-resent the scien-tific interests ofEurope on oneside and theUnited States withCanada on theother side. Chile-an astronomersare closely in-volved with the

project, and 10 % of the observingtime will be reserved for Chilean sci-ence.

ALMA will be built in the Andes, onthe Plateau of Chajnantor, 5000 metresabove sea level and 60 km east of thetown of San Pedro de Atacama. The ar-ray will be comprised of 64 antennaswith unprecedented sensitivity and an-gular resolution that will allow studyingthe origin of galaxies, stars and planets,opening new horizons for astronomy,and being able to observe galaxiesacross the universe where stars are be-ing formed.

The agreement now signed betweenESO and the Government of theRepublic of Chile recognizes the inter-est that the ALMA Project has for Chile,as it will deepen and strengthen coop-eration in scientific and technologicalmatters between the parties.

Signing of the new Agreement that authorizes ESO to establish a new centre for astro-nomical observation in Chile. General view.

Chilean Minister of Foreign Affairs of the Republic of Chile, Mrs.María Soledad Alvear (seated, right) and the ESO Director General,Dr. Catherine Cesarsky (seated, left).

Eighty Nights Up a MountainBy SARA ELLISON

sentially a postdoctoral position; I start-ed here directly after I completed my

A version of this article first appeared in August 2002 in the onlinepublication Next Wave (www.nextwave.org), published by ScienceMagazine on the occasion of the UK joining ESO.

As Sara Ellison describes, she has been a Paranal Fellow for the pasttwo years and is currently in her third year at the Pontificia UniversidadCatólica de Chile in Santiago.

49

first Chilean telescope site. Contrac-tually, we spend 80 nights a year (usu-ally in bite-size shifts of 1 week) on themountain doing support work, plus upto a month doing duty-related work inSantiago. The rest of the time is yourown to pursue whatever scientific re-search takes your fancy. The advan-tage in this set-up is that your duty re-sponsibilities can be quite easily com-partmentalized: When you’re on themountain you do support and whenyou’re in Santiago you do science.

For every week you spend on themountain, you officially get the follow-ing week off as ‘compensation time’.However, most people tend to only takea few days off in order to maximize thetime they can spend on their science,as research output is always an impor-tant factor in securing the next job. ForChile fellows, the contracts are nomi-nally four years (a recent – and wel-come – change from the old style 3-year contracts); the first three years arespent divided between research andsupport whilst there are various optionsavailable for the final year. You mightcontinue working at ESO Chile, but withonly 40 nights of support work, or youmight go back to the headquarters inGermany and do a similar duty load inone of the divisions there. Alternatively,you can elect to go to one of theChilean universities where you will do100% research and have the benefitsof applying for telescope time as aChilean. Due to the relatively smallnumber of Chilean astronomers com-pared to the large number of telescopehours they are entitled to, this offers agreat opportunity to collect lots of data.A final option that has just been ap-proved is that fellows may return toa university department that is willingto ‘adopt’ them in any ESO membercountry with their salary still paid byESO. This is a particularly attractive op-tion for many as it gives fellows whohave been away from their home coun-tries for three years the opportunity toadvertise themselves to universitiesthat they see as potential employerswhen their fellowship ends. This latteroption also offers the potential to gainsome teaching or supervisory experi-ence, opportunities that are limited atESO itself.

The flexibility, diversity, and fourguaranteed years are just a few of theadvantages of the ESO fellowshipscheme. Other professional plusesinclude the opportunity to meet thelarge number of European astronomersthat pass through to conduct their ob-servations, providing ample possibili-ties for forming new collaborations andmaking yourself known within the com-munity. In addition, it is generally at-tractive for future employers to seethat you have had experience with

instrumentation and observatory work,which is a good complement to inde-pendent research. ESO fellowships arealso very well regarded and presti-gious, so it is certainly a good positionto have held.

On the practical side, ESO fellowsare well paid, with allowances for beingoverseas and having to spend time onthe mountain, and the salary is essen-tially tax-free. There are also salarybenefits if you are married and havechildren (who additionally receive aneducation allowance). Taking into ac-count the lower cost of living in Chile,you can certainly live very well on thesalary you take home, and there arelots of opportunities to live well in Chile!The lifestyle here is definitely an attrac-tion – Santiago has the beach, moun-tains, and vineyards on its doorstepand all of the trappings of a big citysuch as cinemas, restaurants, and the-atres.

Of course, every job has its associat-ed drawbacks, and fellowships are noexception. The support work is very de-manding, requires a lot of commitment,and involves a lot of travelling (it’s 6hours door-to-door to Paranal) andnight work. You also have to spend asignificant fraction of your time awayfrom your family, and spouses do nothave work status. Also, the supportwork obviously means you can spendless time on science, although this issomething you have presumably ac-cepted if you are considering an obser-vatory job. For me, I have actuallyfound that my productivity has not seri-ously suffered and that time away at thetelescope is a refreshing change fromthe office. The key to success is defi-nitely effective short-term time man-agement; it is important to be able toplan your work on time scales of a fewweeks between each telescope shift.This means breaking down large tasksinto smaller jobs and setting goals anddeadlines on a weekly basis. Focusingin this way has allowed me to be muchmore efficient with my time when I’mnot on the mountain, a technique thatwill be useful even when I’m no longerworking at ESO.

My contract with ESO will come toan end in September 2003. However, Ialready have a tenure-track job lined upat the University of Victoria in Canada,allowing me to satisfy my ‘wanderlust’still further and get to know anothernew country. Apart from my researchprofile, the experience and skills I havedeveloped at ESO are definitely part ofwhat secured this position, as Canadaincreases its access to large tele-scopes and its involvement in futureprojects such as NGST and ALMA. I’llcertainly be sad to leave Chile though,although I’m sure it will only be hastaluego and not adios.

There are many possibilities open toa completing graduate student and itcan be quite daunting filtering throughthe various options, academic and oth-erwise. For me, the decision was basedon a few fundamental academic andpersonal priorities. I wanted a destina-tion that was going to provide an inter-esting and fulfilling lifestyle, preferablywith the chance to learn a new lan-guage. I was also looking for a workingenvironment that would offer diverseopportunities for professional develop-ment. Working in an observatory wastherefore very appealing, and at thattime (as is still the case a few years lat-er) there were plenty of positions as ob-servatories expanded their horizons tothe new generation of large telescopes.In addition, as an undergraduate, I hadspent five consecutive summer vaca-tions working at various observatoriesin locations including Australia, Hawaii,and the Canary Islands, and undertak-ing support work and dabbling in instru-mentation seemed a good complementto my scientific research.

Therefore, many of my postdoc ap-plications were aimed at observatories,and I was particularly excited by thepossibility of moving to Chile. The fel-lowship application consists of astraightforward form (no long researchproposal like most places!) followed byan interview at the ESO headquartersnear Munich in Germany with a videolink to Chile. Applicants are selected onthe combined criteria of a strong scien-tific record and suitability for supportwork (particularly important for Chile-bound fellows). Priority for posts goesto nationals of ESO member countries,and at that time the UK was not a mem-ber (it joined this July). However, be-cause Chilean fellows have to fulfillvery specialist duties at the telescopes,there are fewer suitable applicants (andusually fewer people who want to moveto Chile) and this criterion is thereforerelaxed. Because the ESO fellowshipwas my first choice amongst the vari-ous positions I applied for, it was aneasy decision when, just before thestart of the Christmas holidays, I wasoffered a position to start the followingautumn.

Many people expect that moving toChile may involve a large culture shock.Although it is certainly not the same asEurope, the facilities and infrastructurein Santiago are very well developedand settling in was not a problem at all.ESO has been dealing with Europeansmoving to Chile for many years and hasa well-oiled system that handles every-thing from moving your household andgetting your local ID and visa proc-essed to helping you find somewhere tolive and opening a bank account. Chilefellows are allocated duty stations at ei-ther Paranal (like me) or La Silla, ESO’s

50

GEORGIEVA, Radostina (BG), AssociateHAUPT, Christoph (D), System EngineerKOLB, Johann (F), StudentLARSEN, Soeren (DK), FellowLOMBARDI, Marco (I), FellowPANGOLE, Eric (F), System EngineerPIGNATA, Giuliano (I), StudentREJKUBA, Marina (HR), FellowRUDOLF, Hans (D), System EngineerSOENKE, Christian (D), AssociateSOMMER, Heiko (D), ALMA Software DeveloperVERDOES KLEIJN, Gijsbert (NL), FellowVERINAUD, Christophe (F), AssociateVICENTE, Silvia (PL), Student

CHILEANDERSSON, Andreas (S), Associate SESTCLARKE, Fraser (UK), FellowGANDHI, Poshak (IND), FellowHUELAMO, Nuria (E), FellowMEDVES, Giuseppe (I), Paid Associate ALMANAKOS, Theodoros (GR), StudentO’BRIEN, Kieran (UK), Operations Staff AstronomerVREESWIJK, Paul (NL), Fellow

DEPARTURES

EUROPEARNOUTS, Stéphane (F), FellowBACHER, Arntraud (A), AssociateBERGOND, Gilles (F), AssociateCAVADORE, Cyril (F), CCD Detector SpecialistCORREIA, Serge (F), AssociateEULER, Christa (D), Administrative AssistantHOSE, Jürgen (D), AssociateTOPLAK, Nenad (HR), AssociateZIEBELL, Manfred (D), Head of Technical Division

ANNOUNCEMENTS

FIRST ANNOUNCEMENT

ESO WORKSHOP ON

High Resolution Infrared Spectroscopy in AstronomyESO-Headquarters, Garching near Munich, November 18–21, 2003

Infrared spectroscopy at a resolution of a few km/s offers a unique tool to study rotational-vibrational transitions of many abundant mole-cules as well as important atomic lines in a multitude of interesting astrophysical environments. Applications include the possible directdetection of extrasolar planets, measurements of the abundances and magnetic fields of stars, studies of ISM chemistry and the kine-matics of stars and gas in galactic centres.The ESO VLT will shortly be equipped with two unique spectrometers which not only offer this high spectral resolution but also spatial res-olutions of ~ 0.2″:

• CRIRES, an adaptive optics fed 1–5 µm spectrograph yielding a resolution of 3 km/s• VISIR, which includes a high-resolution mode yielding a resolution of < 8 km/s between 8 and 13 µm

The aims of this workshop are to:

• present the latest status of the high-resolution infrared spectroscopic capabilities of the VLT and other observatories• bring together the community interested in the application of high resolution infrared spectroscopy and to foster new

collaborations• provide ESO with feedback in the phase when operating and data reduction software is being defined and coded

For space reasons the number of participants is limited to 110.

Scientific Organizing Committee:Bengt Gustafsson (Uppsala, chair), Catherine de Bergh (Meudon), Ewine van Dishoeck (Leiden), Artie Hatzes (Tautenburg), Ken Hinkle(Tucson), Ulli Käufl (ESO), Alan Moorwood (ESO, co-chair)

More details, preliminary pogramme and registration form can be retrieved from http://www.eso.org/gen-fac/meetings/ekstasy2003or contactekstasy2003@eso.org for further information

Symposium Secretary: Christina StofferEuropean Southern Observatory, Karl-Schwarzschild-Str. 2,D-85748 Garching bei Muenchen, GermanyFax: +49 89 3200 6480; E-mail ekstasy2003@eso.org

ESO Workshop on LargeProgrammes and Public

SurveysDates: 19–21 May 2003

Location: Garching, Germany

An evaluation of the scientific success of the first completed LargeProgrammes at the VLT and public surveys should take place be-fore the survey telescopes VST and VISTA start operating. To as-sess the scientific return of the Large Programmes, ESO is or-ganizing a three-day workshop in Garching from 19 to 21 May2003. To ESO it will offer an opportunity to hear directly from major usersof the facilities of their experiences, problems and wishes. The re-lation of surveys to other forms of observational programmesshould be discussed.

PERSONNEL MOVEMENTSInternational Staff (1 October – 31 December 2002)

ARRIVALS

EUROPE

BASTIAN, Nathan (USA), StudentDOLENSKY, Markus (D), System Engineer

51

CHILE

DAHLEM, Michael (D), Operations Staff AstronomerDELSANTI, Audrey (F), StudentFAURE, Cécile (F), StudentZALTRON, Paolo (I), Associate

Local Staff (1 September – 30 November 2002)

ARRIVALSCANIGUANTE BARCO, Luis Alejandro, Electronic EngineerHENRIQUE LEON, Juan Pablo, Mechanical Technician

duty in the Visiting Astronomers’ Section then headed by Dr. A.B.Muller.

During the twenty-six years that Christa worked in this Section,apart from being the key person whose diligence was much appreci-ated by all the astronomers having to travel to Chile, she has –among many other things – very efficiently handled over 20,000 ob-serving proposals, and organized about fifty OPC meetings, includ-ing the corresponding “exclusive” OPC dinners. Today, her name isfamiliar to most European astronomers as well as to many othersoverseas, as demonstrated by the numerous expressions of heartfeltgratitude received for the farewell booklet that was presented to heron November 28, on the occasion of the OPC dinner, which was heldin her honour.

Friends, colleagues and astronomers will certainly miss ChristaEuler. We wish her all the best for the years to come.

Jacques Breysacher

LEIVA BECERRA, Alfredo, Instrumentation EngineerROA FIGUEROA, Luis, Mechanical TechnicianROJAS CHANAMPA, Chester, Instrument TechnicianTORRES ZAMORANO, Manuel, Maintenance Electronic

Technician

DEPARTURESALVAREZ CASTELLON, Pedro,

Precision MechanicCASTRO CASTRO, Wilson, MechanicMIRANDA MANRIQUEZ, Marcela,

Administrative Secretary

Manfred Ziebell RetiresIt’s hard to believe, but by the end

of November, Manfred’s soundingfootsteps will no longer be heard atESO Garching.

Manfred initiated his engineeringcareer at ESO when the 3.6-m tele-scope was in the design phase andhad to be brought into a reality. Newservo controls, positioning systems,communication protocols were his lot;a time when ESO struggled to reachits autonomous design maturity.Manfred proved to be a key figure forthe future technological challenges.

Faced with a jungle of technical ini-tiatives, he managed to develop elec-

tronic standards that merged successfully into the VLT design.Manfred assembled design and support teams for the telescope

and the instrumentation controls. Convincing rather than imposingwas his style, a style that was very much appreciated and highly re-spected by his colleagues.

Service-minded, he became a paradigm of support endeavours,by solving technical problems, contractual issues and human rela-tions. He was ready to tackle any issue, including Hausmeister andsafety functions.

Problems? Ask Manfred! Manfred was on the phone with Chile or… in Chile!The VLT encoding units failed? Manfred was already negotiating

alternatives!As Head of the Technical Division, he established a support cul-

ture that placed ESO on the trail to success.Manfred’s hurried and sounding steps will be missed at ESO.“Un hombre de primeras aguas, ahora otras aguas lo esperan”

Daniel Hofstadt

Thirty-Seven Years of Servicewith ESO!

On December 1st, 2002, after thir-ty-seven years of service, first in Chileand then in Garching, Ms. ChristaEuler will leave ESO to enjoy a well-deserved retirement. Among the cur-rent staff, she is probably the only per-son who started her career at ESOjust four years after the Organizationwas founded.

Christa joined ESO Chile on April1st, 1966, at the time when Prof.Heckmann was Director General. Inthese early days she was responsiblefor all secretarial work in the Santiago

office, which was first located in the Guesthouse before moving to theVitacura building. Three years later, upon the arrival of Prof. B.E.Westerlund, she took over the post of secretary to the ESO Directorin Chile. With the exception of a ten-month period spent in thePersonnel Department in Hamburg, she held this position until mid-1976. Her definitive move to ESO Europe took place in September1976. At the newly installed Headquarters in Garching, she took up

New ESO Proceedings

The Proceedings of the Topical Meeting, held in Venice from 7 to 10 May 2001

Beyond Conventional Adaptive Optics

(ESO Conference and Workshop Proceedings No. 58)

have now been published. The price for the 490-page volume,edited by E. Vernet, R. Ragazzoni, S. Esposito and N. Hubin, is8 50.– (prepayment required).

Payments should be made to the ESO bank account 2102002with Commerzbank München (BLZ 700 400 41, S.W.I.F.T. Code:COBADEFF 700) or by cheque, addressed to the attention of

ESO, Financial Services, Karl-Schwarzschild-Str. 2D-85748 Garching bei München, Germany

ESO Workshop Proceedings StillAvailable

Many ESO Conference and Workshop Proceedings are still avail-able and may be ordered at the European Southern Observatory.Some of the more recent ones are listed below.

No. Title Price

45 ESO/EIPC Workshop “Structure, Dynamics and Chemical Evolution of Elliptical Galaxies”, 1993 7 45.–

46 Second ESO/CTIO Workshop on “Mass Loss on the AGB and Beyond”, 1993 7 35.–

47 5th ESO/ST-ECF Data Analysis Workshop, 1993 7 15.–

48 ICO-16 Satellite Conference on “Active and Adaptive Optics”, 1994 7 45.–

49 ESO/OHP Workshop on “Dwarf Galaxies”, 1994 7 45.–

50 ESO/OAT Workshop “Handling and Archiving Data from Ground-based Telescopes”, 1994 7 20.–

51 Third CTIO/ESO Workshop on “The Local Group: Comparative and Global Properties”, 1995 7 25.—

52 European SL-9/Jupiter Workshop, 1995 7 40.—

52

ESO, the European Southern Observa-tory, was created in 1962 to “... establishand operate an astronomical observato-ry in the southern hemisphere, equippedwith powerful instruments, with the aim offurthering and organising collaborationin astronomy...” It is supported by tencountries: Belgium, Denmark, France,Germany, Italy, the Netherlands, Portu-gal, Sweden, Switzerland and the UnitedKingdom. ESO operates at two sites inthe Atacama desert region of Chile. Thenew Very Large Telescope (VLT), thelargest in the world, is located onParanal, a 2,600 m high mountain ap-proximately 130 km south of Antofa-gasta, in the driest part of the Atacamadesert where the conditions are excellentfor astronomical observations. The VLTconsists of four 8.2-metre diameter tele-scopes. These telescopes can be usedseparately, or in combination as a giantinterferometer (VLTI). At La Silla, 600 kmnorth of Santiago de Chile at 2,400 maltitude, ESO operates several opticaltelescopes with diameters up to 3.6 mand a submillimetre radio telescope(SEST). Over 1300 proposals are madeeach year for the use of the ESO tele-scopes. The ESO headquarters are lo-cated in Garching, near Munich, Ger-many. This is the scientific, technical andadministrative centre of ESO where tech-nical development programmes are car-ried out to provide the Paranal and LaSilla observatories with the most ad-vanced instruments. There are also ex-tensive astronomical data facilities. ESOemploys about 320 international staffmembers, Fellows and Associates inEurope and Chile, and about 160 localstaff members in Chile.

The ESO MESSENGER is publishedfour times a year: normally in March,June, September and December. ESOalso publishes Conference Proceedings,Preprints, Technical Notes and other ma-terial connected to its activities. PressReleases inform the media about partic-ular events. For further information, con-tact the ESO Education and PublicRelations Department at the followingaddress:

EUROPEAN SOUTHERN OBSERVATORYKarl-Schwarzschild-Str. 2 D-85748 Garching bei München Germany Tel. (089) 320 06-0 Telefax (089) 3202362ips@eso.org (internet) URL: http://www.eso.orghttp://www.eso.org/gen-fac/pubs/messenger/

The ESO Messenger:Editor: Peter ShaverTechnical editor: Kurt Kjär

Printed by Universitätsdruckerei WOLF & SOHNHeidemannstr. 166D-80939 MünchenGermany

ISSN 0722-6691

ContentsTELESCOPES AND INSTRUMENTATION

L. Pasquini et al.: Installation and Commissioning of FLAMES, the VLTMultifibre Facility . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 1

F. Pepe, M. Mayor, G. Rupprecht: HARPS: ESO’s Coming Planet Searcher.Chasing Exoplanets with the La Silla 3.6-m Telescope . . . . . . . . . . . . . . . 9

K. Kuijken et al.: OmegaCAM: the 16k ×16k CCD Camera for the VLTSurvey Telescope . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15

A. Glindemann: The VLTI – 20 Months after First Fringes . . . . . . . . . . . . . . . 18B. Koehler, C. Flebus, P. Dierickx, M. Dimmler, M. Duchateau, P. Duhoux,

G. Ehrenfeld, E. Gabriel, P. Gloesener, V. Heinz, R. Karban, M. Kraus, J.M. Moresmau, N. Ninane, O. Pirnay, E. Quertemont, J. Strasser,K. Wirenstrand: The Auxiliary Telescopes for the VLTI: a Status Report . . 21

Roberto Gilmozzi and Jason Spyromilio: Paranal Observatory – 2002 . . . . . . 28O. Hainaut: News from La Silla: Science Operations Department . . . . . . . . . . 30Two unusual views of La Silla . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 31

REPORTS FROM OBSERVERS

N.C. Santos, M. Mayor, D. Queloz, S. Udry: Extra-Solar Planets . . . . . . . . . . 32I. Labbé, M. Franx, E. Daddi, G. Rudnick, P.G. van Dokkum,

A. Moorwood, N.M. Förster Schreiber, H.-W. Rix, P. van der Werf, H. Röttgering, L. van Starkenburg, A. van de Wel, I. Trujillo, and K. Kuijken: FIRES: Ultradeep Near-Infrared Imaging with ISAAC of theHubble Deep Field South . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 38

OTHER ASTRONOMICAL NEWS

M. Kissler-Patig: Summary of the Workshop on Extragalactic Globular Cluster Systems hosted by the European Southern Observatory in Garching on August 27–30, 2002 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 42

A. Glindemann: The VLTI: Challenges for the Future. Workshop at Jenam 2002 in Porto . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 44

P. Shaver and E. Van Dishoek: Summary of a Meeting on Science Operations with ALMA, held on Friday, 8 November 2002 . . . . . . . . . . . . . 44

C. Madsen: Celebrating ESO’s 40th Anniversary . . . . . . . . . . . . . . . . . . . . . . 45C. Madsen: Breaking the Ground for the European Research Area – The

Conference ‘European Research 2002’ . . . . . . . . . . . . . . . . . . . . . . . . . . 46A Nobel Prize for Riccardo Giacconi . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 47Agreement Between the Government of the Republic of Chile and ESO for

Establishing a New Centre for Observation in Chile – ALMA . . . . . . . . . . . 48Sara Ellison: Eighty Nights Up a Mountain . . . . . . . . . . . . . . . . . . . . . . . . . . . 48

ANNOUNCEMENTS

First Announcement of an ESO Workshop on High-Resolution Spectroscopy in Astronomy . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 50

ESO Workshop on Large Programmes and Public Surveys . . . . . . . . . . . . . . 50Personnel Movements . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 50Daniel Hofstadt: Manfred Ziebell Retires . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51Jacques Breysacher: Christa Euler – Thirty-Seven Years of Service with ESO! 51New ESO Proceedings . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 51ESO Workshop Proceedings Still Available . . . . . . . . . . . . . . . . . . . . . . . . . . 51

No. Title Price

53 ESO/ST-ECF Workshop on “Calibrating and understanding HST and ESO instruments”, Garching, Germany. P. Benvenuti (ed.) 7 30.–

54 Topical Meeting on “Adaptive Optics”, October 2–6, 1995, Garching, Germany. M. Cullum (ed.) 7 40.–

55 NICMOS and the VLT. A New Era of High Resolution Near Infrared Imaging and Spectroscopy. Pula, Sardinia, Italy, May 26–27, 1998 7 40.–

56 ESO/OSA Topical Meeting on “Astronomy with Adaptive Optics – Present Results and Future Programs”. Sonthofen, Germany, September 7–11, 1999. D. Bonaccini (ed.) 7100.–

57 Bäckaskog Workshop on “Extremely Large Telescopes”. Bäckaskog, Sweden, June 1–2, 1999. T. Andersen, A. Ardeberg, R. Gilmozzi (eds.) 7 30.–