No 15 - December 1994

Future VLT Instruments:

Scientific Drivers and

Concept Definitions

Edited by Sandro O'Odorlco

Published by EUROPEAN SOUTHERN OBSERVATORY Karl-Schwarzschild-StraBe 2, D-85748 Garching bei Miinchen Gennany

© Copyright ESO 1995

FOREWORD

In May 1993 the Scientific Technical Committee of ESO established a Working Group with the assignment "to propose a set of main scientific priorities to serve as guidelines for future discussions on the VLT and its instrumentation". The final membership of the WG included K.S. de Boer, S. D'Odorico, B. Fort, R-P. Kudritzki, B. Marano, L. Vigroux (Chairman), J. B. Walsh and J. Wampler. A preliminary report from the WG was presented to the STC in December 93. In the subsequent months the WG did involve a number of scientists in the preparation of the scientific cases and the concept definitions for 8 instruments which were identified as potentially important for the scientific use of the VLT in addition to the six now under construction (ISAAC, FORS 1, FORS2, CONICA, UVES and FUEGOS). Their work was presented in a preliminary form at the ESO Workshop "Science with the VL T" held in Garching in June 94.

ESO gratefully acknowledges these efforts to the benefit of the scientific success of the VL T. The definition studies for 7 of the new instruments are collected in this report as a reference for the discussion on future VL T instrumentation. For the eighth instrument, the Mid Infrared Imager Spectrometer, a Phase A study· by a consortium of institutes with P.O. Lagage as P.I. has been completed in October 94

and documentation is available separately. Additional discussions on the scientific

use of the VLT with these instruments can be found in the Proceedings of the June 94 Workshop edited by J. Walsh and J. Danziger.

Sandro D'Odorico, ESO

3

TABLE OF CONTENTS

Foreword S. D 'Odorico .......................................................................................... ....... ..... .......... 3

FRISPI: A Fast Recording Imager and Spectro-Photometric Instrument for the VL T H. Barwig, K.-H. Mantel and S. Kiesewetter ................................................................ 7

NIRMOS: A Wide Field Near-IR Multislit Imaging-Spectrograph for the VLT O. Le Fevre, P. Felenbok, F. Hammer, L. Tresse, B. Delabre, P. Vettolani, Y. Mellier, J.P. Picat, S. J. Lilly ........................................................ : ......................... 17

VHRS: The Very High Resolution Spectrograph for the VLT P. Magain, H. Dekker, B. Delabre .............................................................................. 37

WFIS: A Wide Field Visual Multislit Imaging Spectrograph for the VL T G. Vettolani, B. Delabre, O. Le Flvre, F. Hammer, G. Zamorani ............................... 53

The VLT Wide-Field Direct Visual Camera (WFDVC) E. J. Wampler ............................................................................................................. 63

Preliminary Design Document for the Visible High Angular Resolution Camera (VHARC) G. Weigelt ................................................................................................................... 77

CRIRES: the VL T High-Resolution Infrared Spectrometer G. Wiedemann, B. Delabre and A. Moorwood ............................................................. 87

FRISPI: A FAST RECORDING IMAGER AND SPECTRO-PHOTOMETRIC INSTRUMENT

FOR THE VLT

H. Barwig, K.-H. Mantel, S. Kiesewetter Universitats-Sternwarte, Scheinerstrafie 1, D-81679 Miinchen, FRG

1 The need for high-speed observations with the VLT Among the many instruments proposed for the VLT, observational facilities for adequately exploring compact stellar objects and related phenomena like white dwarfs, neutron stars and black holes are missing. .

These targets predominate the observational parameter domain of short time scales due to their small geometrical extent. Furthermore, in their environment very high energy densities and enormous magnetic fields are present. Maximum radiation is released at short wavelengths (UV - X-rays), which makes these objects appear relatively faint in the optical. Though of fairly low absolute visual luminosity, these objects play an important role in several fields of astrophysics. .

Further fundamental insights in the physics of these objects require spectrally resolved observations with integration times in the range between milliseconds and several seconds with high SIN and can therefore only be performed with an appropriate instrument at the largest telescopes.

2 Scientific programs and observational requirements for exploring short-time-scale phenomena

The most challenging scientific programs in the field of compact object research are given in Table 1. For each project the adequate combination respectively of time and spectral resolution is indicated.

Spectro-photometry of cataclysmic variables (eV) Spectral eclipse mapping of accretion disks in ev s Flickering and light-echo mapping in accretion disks Magnetic-field tomography in eclipsing AM H:er stars Determination of white dwarf masses in faint eclipsing evs Search for millisecond-pulsars in globular clusters Phase-resolved spectro-photometry of pulsars Doppler tomography of accretion flow in interm. polars System parameter determination of close detached binaries Light-echo mapping in X-ray binaries Line profiles of pulsating white dwarfs Search for optical counterparts of gamma-ray bursts

Table 1

7

time res.

1 s 1 s

0.1 s 0.1 s

1-10 s 0.1-10 ms 0.1-10 ms

1-10 s 1-10 s

0.1 s 1-10 s

0.1-10 s

spect. res.

5-100 1000 100 100 10 1

100 1000

10 10

~ 103

1

These scientific programs have been proposed by the authors and the following col­leagues:

• K. Beuermann, Gottingen, FRG

• D. Dravins, Lund, Sweden

• C. Gouiffes, Saclay, France

• R. Hifner, Miinchen, FRG

• F. Hessman,· Gottingen, FRG

• K. Home, St. Andrews, UK

• C: Motch, Strassbourg, France

• M. Pakull, Strassbourg, France

• J. van Paradijs, Amsterdam, NL

• H. Pedersen, Kopenhagen, DK

• A. Schwope, Potsdam, FRG

• K. Simon, Miinchen, FRG

• H.-C. Thomas, Garching, FRG

• F. Verbunt, Amsterdam, NL

The following section presents some topics of research work done by the authors at the Universitiits-Sternwarte Miinchen in the field of accretion physics.

3 Investigation of accretion physics as an example for high-speed spectro~photometry at the VLT

In the resent years different new numerical methods have been developed for the analysis of observations of compact accreting objects. These methods require high SIN ratios combined with medium spectral and high time resolution.

3.1 Eclipse mapping of accretion disks in close binaries

3.1.1 Broad band observations

Among compact targets, close binary systems e.g. cataclysmic variables (CV) are of enormous scientific interest. CV s are interacting binaries with a white dwarf primary and a red dwarf secondary which fills its Roche lobe. Orbital periods range between 80 minutes and several hours. A gas stream flowing through the inner Lagrangian point transfers matter from the secondary to an accretion disk around the primary. The transfer

8

20 P arcs

+

Pigure 1: Ecl ipse maps of an accretion disk reconstructed from four different. eclipse light curves of a dwarf nova. On t.he right. side, the Roche lobe filling secondary is indicated to show the relative system dimensions. (Note the high spatial resolution!)

of angular momentum outwards allows mattcr to move inwards towards the while dwarf, but the physica1 mechanism of this accretion-disk viscosity is noL yet understood. An important group among the CV systems arc dwarf novae (ON) which show quasi-periodic outbursts.

The process of mass accretion in a disk turned out to be of outstanding significance for many astrophysical processes. It is likewise found in disks of prolosLars and AGNs and is also relevant, c.g., for the evolution of galaxies and planetary systems as well. CVs provide ideal laboratories for studies of mass accretion mechanisms and for testing related theories since the components in these labs are fairly well known through observation of their spectra and their photometric signatures. Under these "calibrated" conditions the disk itself can be studied.

Most convenient for these investigations are ONs with high orbital inclination where the ecl ipse by the secondary provides scans across the accretion disk during quiescence and outburst. The brightes t of these objeds have apparent magnitudes of about 15th magnitude. Time and specLralJy resolved observations of eclipse phenomena using large telescopes like the VLT, combined with high-speed spectro-pholometric instruments pro­vide the input data for the famous eclipse-mapping method first introduced by Horne (1985).

As an example, fi rst results (Bobinger 1994) from the application of the eclipse­mapping technique to photometric data obtained in the V-band of a double eclipsing ON are presented in Pig.l . It shows two-dimensional disk maps reconstructed from one­dimensional ecl ipse light curves which had been observed with a 2-111 telescope during 4

9

different nights. Spatial brightness structures of the accretion disk are clearly visible at the micro-are-second level. Simulations of eclipses in ev systems assuming a steady state disk with two artificially shaped hot spots clearly proved that the quality of the eclipse-mapping reconstruction strongly depends on the SIN of the input data (Spruit 1994). Increasing the SIN by averaging over several orbital periods smears out spatial structures. In order to get information on the actual processes of accretion and to achieve spatial resolution of sub-micro-arc-second struct1ire8, single eclipse events have to be an­alyzed, a task which cannot be performed without the light-collecting power m an 8-m telescope.

3.1.2 Narrow-band observations

The need for large telescope apertures becomes even more stringent for applications in­volving the spectrally-resolved eclipse-mapping technique, a straight-forward extension of the original maximum-entropy eclipse-.mapping method. Spectra from isolated parts of the accretion disk can be synthesized from many narrow-band images of low resolution (Rutten et al. 1994). From comparison with theoretical disk spectra, the physical condi­tions in the accretion disk are derived. Also the spectra of the hot spot and the secondary, respectively, can be extracted. This technique requires accurate spectro-photometric data of high SIN and sufficient phase resolution to slice up the accretion disk.

3.2 Doppler tomography and spectral eclipse mapping

Doppler tomography, developed by Marsh & Horne (1988), provides another powerful tool for exploring accretion phenomena. It allows to map the emission regions associated with the accretion flows in evs. For this purpose two-dimensional data consisting of the velocity profiles as a function of binary phase are used. A combination of the Doppler tomography method with the spectral eclipse-mapping technique allows one to study the velocity fields in accretion disks and hence the processes relevant for understanding the outburst mechanism. Precise time-resolved high-resolution spectro-photometry of single eclipses cannot be successfully performed with anything less than an 8-m class telescope.

3.3 Investigation of ftickering

Flickering seems to be a quite common signature of the accretion process and therefore could probably be related to the viscosity mechanism. This association also seems most important for studying the pre-outburst phase of evs. The flickering time scales range from a few seconds, possibly indicative of Keplerian motion near the white dwarf surface, to a fraction of the orbital period. Observations of the phase-dependent disappearance and reappearance of flickering activity observed in eclipsing ev systems indicate that this phenomenon is not exclusively confined to the hot spot region. In order to localize the site of origin, the eclipse-mapping technique can be applied to data which show the rms of the residUals obtained by subtracting a smoothed light curve from the original eclipse light curves averaged over many orbital cycles (Horne and Stiening, 1985). Knowing the site of flickering, high-speed spectro-photometry of the individual flickering events could provide information on the flickering process. Furth~rmore, one may search for light echoes using cross-correlation techniques, a first step towards a possible echo mapping of ev systems.

10

This task can be achieved using spectrally resolved photometry of single flickering events with integration times in the sub-second range.

4 General instrumental requirements In order to perform the above mentioned scientific programs the following observing modes are applied:

1. High speed multi-aperture spectro-photometry

2. High speed direct imaging photometry

3. High speed low/medium resolution spectroscopy (long- or multi-slit)

In principle a device for accurate high-speed aperture photometry for the VLT should meet the following specifications:

Measuring mode Simultaneous measurements of program target, sky background and of at least two nearby comparison stars are mandatory for each observing mode. This is most important in order to achieve high photometric accuracy and to allow operation even during not perfect photometric conditions, especially in target of opportunity mode. The spectral behaviour of the program target must be monitored in order to allow temperature determinations (e.g. of 2-dim. structures applying e·:.lipse mapping methods). This requires the application of spectro-photometry rather than successive filter measurements.

In direct imaging mode obscuration of bright· objects in the field of view is required in order to optimize the dynamic range of a photon counting detector.

Spectral range According to the scientific aim and the faintness of the targets the spectral resolution range from integral light (UV-IR) over those characteristic of broad-band photometric systems (e.g. Johnson UBVRI, Stroemgren, Geneva) to low (R~50) and high (R=5 104 ) resolution spectro-photometry. In order to investigate photometric data according to different scientific aspects it shall be possible to synthesize different standard photometric systems by software spectral binning after observation.

Target selection The selection of the targets in the focal plane of the telescope shall not be restricted by the entrance slit position of a spectrograph. Seeing and image movements must not influence the illumination of the spectrograph's entrance slit. This implies that the photometric apertures and the entrance slits of the spectro­graph must be decoupled. The only known solution for this requirement is the use of optical fibres between the Cassegrain focal plane and the spectrograph. Further­more, the adaption to different seeing conditions as well as utmost suppression of sky background requires exchangeable diaphragm sizes.

Field size For aperture photometry of objects with precisely known coordinates, a field of 1-2 arcmin at the VLT seems to be sufficient for the selection of comparison stars of suitable brightness.

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For imaging photometry (e.g. for extended objects, search for variables with large positional error boxes) the available field size should be considerably larger, preserv­ing the spatial resolution during good seeing conditions. A field of view of 5' x 5' should be achieved.

Time resolution Integration times between 100 fJsec and a few 10 seconds is required for the above mentioned scientific programs. Note, that the shortest exposure times are also limited by the dynamical range of photon-counting detectors (~ 106ctsls for a MAMA detector). In principle, the integration time shall match the time constant of the phenomenon to be observed, i.e. for a given object brightness only the spectral resolution can be adjusted for a required SIN.

Detectors For imaging photometers a 2-dim. photon-counting device is indispensable. In principle for fibre-optic aperture photometers photo-multipliers can be used to measure broad-band regions of a spectrum of very low resolution (e.g. the Multi­Channel-Multi-Colour-Photometer MCCP of the Universitats-Sternwarte Miinchen, Barwig et al. 1987). If the star light is split by a higher resolution spectrograph, position-sensitive photon-counting detectors (PCA, MAMA) have to be used.

Photometric accuracy Besides the errors due to photon statistics, systematic errors arise from insufficient extinction corrections and from instrumental instabilities. An overall accuracy of the order of a few promilles shall be reached.

Software (on-line facilities) Preliminary on-line data reduction and monitoring is mandatory. For special purposes (observations of faint objects through narrow fil­ters, yielding low count rates) the photon-time-tagged mode may be used with advantage (as is often done in space). The software shall provide for event-triggered acquisition, for on-line ,period folding, and for evaluation of statistical properties. It is also important to be able to quickly study the data in respect to (possibly spurious) frequencies.

Mounting at the VLT Both photometer devices shall be permanently mounted, cooled and electronically connected so that they can be taken into use at seconds' no­tice ("standby-mode"). Argument: other telescopes, being ground-based or space­borne, may provide discoveries of time-variable objects which merit immediate ob­servations. Such cases are likely to occur with increasing frequency as world-wide electronic communication becomes commonplace. For example data from the detec­tion of a Gamma-Ray-~urst (GRB) event by BATSE/GRO can reach La Silla and Paranal via satellite about 6 seconds after the trigger event. Since more than half of all GRBs last longer than this, there is still time to do strictly simultaneous ob­servations from ground, in particular, if the telescope has a reasonably fast slewing capability.

5 Technical realization

As an instrument incorporating all these requirements, the multi-object spectro-photome­ter FRISPI (Fast Recording Imager and Spectro-Photometric Instrument) is proposed for

12

the VLT (Fig.2). FRISPI is based on experiences obtained during ten years of photometer development at the Universitats-Sternwarte Miinchen (including MCCP, Barwig et al 1987, MEKASPEK, Mantel et al. 1988, 1993, Mantel & Barwig 1994).

It will be both a high-speed aperture spectro-photometer and a high-speed focal re­ducer.

5.1 Aperture photometer mode

In this mode FRISPI will be a fibre-fed spectro-photometer which allows to measure in the optical wavelength range (AA 3400 - 9 000 A) 4 sources simultaneously with selectable spectral resolution (R~ 1000).

Four fibre bundles can be positioned by computer control in the focal plane of the telescope. The fibre positioning device has to match the focal field curvature of the telescope.

An off-axis target-selector camera allows for object identification and precise fibre positioning.

A Fabry lens at the entrance of the fibre bundles projects the pupil of the telescope onto the circular fibre bundle with a diameter of 800 IJm. The entrance diaphragms are exchangeable with a maximum size of 2mm. The atmospheric dispersion must be compensated by an ADC in front of the instrument.

The rectangular outcome of the fibre bundles form the entrance slit of a grism spec­trograph. The grisms are exchangeable for different spectral resolutions and different spectral ranges.

A rectangular diaphragm (2.5 x 0.2 mm2) for the fibre bundle outcome and a f-number of 3 for the spectrograph entrance cone yield an imaging ratio of 3: 1 for the spectrograph optics. The f-number of the camera optic will be of the order of 1(!).

The spectra of the four sources are imaged onto the photo-cathode of a 2k x 2k PCA detector and are recorded by an image-acquisition computer. This computer registers the incoming photon events either in full images, in time-tagged mode or binned for objects and spectral resolution elements.

The overall spectral resolution of AI ~A ~ 1 000 allows to bin data by means of software into synthetic overlapping filter curves which are optimized for the scientific aim and which meet the requirements of the Sampling Theorem for adequate description of colour effects. In this way correct treatment of atmospheric extinction effects' as well as exact transformations to any other broad-band photometric system can be achieved.

5.2 Focal reducer mode

The focal reducer mode is realized by exchanging the multi-fibre collimator unit with a field collimator. This allows to use the high-speed capabilities of the PCA for direct imaging as well as for long-slit spectroscopy.

A mask and slit unit either provides up to several slits for spectroscopy or can be used for masking bright targets that may overload the detector.

The imaging optics for the focal reducer mode of FRISPI will be similar to the standard-resolution optics of the FORS instrument (Seifert et al. 1994) with a cam­era focal length of 280 mm and a collimator focal length of about 1250 mm. This yields a 3'.5(5'.6) field on the sky for a 25 (40) mm photo-cathode.

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Field

Viewer

(CCO)

Multi Fibre Unit

Crossdisperser

t-----..... Filter

Filter

Mask and

Slit Unit

• c:::::::: ::::::>

Field Collimator

-<: ::> -

Figure 2: Layout of the high-speed spectro-photometer proposed for the VLT.

14

Different filters can be switched into the light-path for observations in selected (broad­band) spectral ranges.

The photo-cathode of the PCA shall be selected for highest quantum efficiency in an extended wavelength range (AA 3400 - 9000 A).

A calibration unit for both photometer modes will be provided as part of the instru­ment.

5.3 Summarized instrumental characteristics 1. Fibre-optic Aperture Photometer

Focal field N umber of channels Positioning of apertures Entrance diaphragm Spectral range Spectral dispersion Photometric system Detector Detector size Pixel size Photo-cathode Quantum efficiency Overall dynamic range Pixel dynamic Dark current

2. Direct Imaging Photometer

Field size Field reduction

180 mm x 180 mm ~4

remote controlled =::; 4"(=::; 2mm) 3400 - 9000 A R =::; 1 300 (TBD) synthesized by software PCA (MAMA) 2k x 2k 25 p.m Trialkali =::; 20% 2000000 counts/second 400 counts/seconds/pixel 20 counts/hour/pixel

5' x 5' adjustable

3. Low/medium resolution spectrograph

Spectral resolution Number of slits Slit width

Time resolution for all modes

time tagged mode full image binning mode

R =::; 1300 (TBD) 2::4 adjustable

=::; 0.1 p.s ~ 4 s ~ 1 ms

15

6 Future prospect

At present the status of the instrument is in a preliminary concept phase. The authors have a vital interest in a high-speed spectra-photometer available for observations with the VLT. In course of the in-house development of the MCCP and MEKASPEK instruments a lot of experience has been gained. The authors would like to contribute their know-how in the field of high-speed spectra-photometric instrumentation to the development of an appropriate instrument for the VLT.

References

Barwig, H., Schoembs, R., Buckenmayer, C. (1987), A&A, 175, 327. Barwig, H. (1994), in Instruments for the ESO VLT, ed. Moorwood, A.F.M., 46. Bobinger, A. (1994), private communication Horne, K. (1985), Mon. Not. R. Soc., 213, 129. Horne, K., Stiening, R.F. (1985), Mon. Not. R. Soc., 216, 933. Mantel, K.-H., Barwig, H., Kiesewetter, S. (1988), in Proc. of New Directions in

Spectrophotometry, eds. Philip, A.G.D., Hayes, D.S., Adelman, S.J., L. Davis press, N.Y., USA, 283. Mantel, K.-H., Barwig, H., Kiesewetter, S. (1993), in Proc. of IAU CoIl., 136, 172. Mantel, K.-H., Barwig, H. (1994), in NATO ASI Series, eds. C. Sterken and M. de Groot,

329. Marsh, T.R. and Horne, K. (1988), Mon. Not. R. Soc., 235, 269. Rutten, R.G.M., Dhillon, V.S., Horne, K., Kuulkers, E. (1994), A&A, 283, 441. Seifert, W., Mitsch, W., Nicklas, H., Rupprecht, G. (1994), SPIE 2198, 213. Spruit, H.C. (1994), A&A, ~n press.

16

NIRMOS: A WIDE FIELD NEAR-IR MULTI SLIT IMAGING-SPECTROGRAPH FOR THE VLT

O. Le Fevrel, P. Felenbokl, F. Hammerl; L. Tressel, B. Delabre2, P. Vettolani3,

Y. Mellier4, J.P. Picat4, S.J. Lilly5

I DAEC, Observatoire de Paris-Meudon, 2 ESO Garching, 3 Istituto di Radioastronomia, CNR, Bologna, 4 Observatoire Midi-Pyrenees, .

5 Astronomy Department, University of Toronto

1 Introduction The ",1 magnitude gain offered by the VLT over 3.5m telescopes and the projected density of interesting faint objects in the near-IR fully justifies the need for an optimized near-IR spectrograph coupled to the large multiplexing gain offered by multi-slit spectroscopy.

One of the main science motivation is the study of very distant galaxies at epochs when the universe was 50% of the present age and younger, which is almost out of the reach of 4m telescopes and becomes fully feasible with the VLT. This will tell us much about the evolution of the stellar populations in the galaxies, and possibly their date of birth, the role of their environment, as well as the evolution in the large scale distribution of galaxies since these early epochs.

Present multi-object spectrographs working in the visible are limited to redshifts z",1.2 because above thisredshift [011]3727 A is redshifted out of the wavelength range and no strong lines are available in the UV until Lya comes in for z > 2. A near-IR spectrograph would therefore be most adapted to these types of studies as it would allow for z'" 1.5 access to spectral features including the well known region from 3727 A to 4340A at rest, and permit not only accurate redshift measurements but also to explore the spectral energy distribution in a well studied rest frame spectral domain. The density of sources is high enough for many programs to make use of the spectacular multiplexing gain offered by multi-slit spectroscopy. Multiplexing has indeed become a critical parameter in the efficiency of astronomical observations: 19 slits on the VLT is strictly equivalent to 94 slits on a 3.6m telescope for e.g. the observations of faint galaxies. It is of the outmost importance to maximize multiplexing as we are investing into collecting more photons with the VLT.

We are therefore proposing a near-IR multi-slit imaging spectrograph working in a 10'xlO' field split into 4 quadrants, from 0.8 to 1.8j.tm, possibly partially extended to 2.3 j.tm, with an average number of slits around 120. We argue that this instrument would have an outstanding impact as a survey spectrograph as well as a wide field IR imager and therefore fills a gap in the current VLT instrumentation. Such an instrument concept has not yet been publicly discussed by any other large telescope project, but will probably sooner or later as the large IR arrays are becoming available. N ow is therefore a good time to explore this prospective area, and give VLT users access to an outstanding science window.

We present below an overview of the science objectives, the instrument basic specifica­tions, and its expected performances. Operational constraints and some budget estimates are also presented.

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2 Scientific objectives In the paragraphs below, the aim has been to o~tline a few outstanding projects. This is by no means an exhaustive list. A near-IR, multi-object spectrograph. would expand on the visible range more that 1S0%, and increase the field size an order of magnitude compared to conventional near-IR spectrographs, and as such, will appeal to the many scientific programs that require near-IR low resolution spectroscopy.

2.1 Evolution of galaxies and structures at z > 1 One of the main drivers to the construction of the ESO /VLT, and other very large tele­scopes, has been the wide interest in having a better understanding of .the distant uni­verse. The main questions are (i) when do galaxies form, (ii) what is the time evolution of galaxies; (iii) what is the distribution of galaxies; (iv) what is the evolution of the large structures, when do they form and in which order. These questions all pertain to the same underlying quest: what is the structure of our universe, what has been and what will be the evolution of the universe and its major constituents.

A lot has been learned with 4m telescopes. The evolution of galaxies has been ex­plored in a more systematic way up to redshifts unity (Fig. 1a; Le Fevre et al., 1994a; Lilly et al., 1994; Colless, 1994), the distribution of galaxies in large scale structures has been demonstrated at low redshifts, and clustering of galaxies has been identified up to redshifts unity, when the universe was about half of its present age. Survey spectrographs such as MOS-SIS at CFHT (Le Fevre et al., 1994b) have been instrumental in these stud­ies. However, it is fair to say that very little is known of the content of the universe and the distribution of matter at redshifts above one, except for the very rare events like QSOs and radio-galaxies. The current limitations of 4m telescopes are not only that the spectral signatures of galaxies with z > 1 fallout of the visible band (0.4 to 0.S5 I'm) and enter the near IR, but more critically, that the collected flux is insufficient to be able to observe normal galaxies in the field or clusters at z > 1 in the near-IR. A near-IR spectrograph on one of the VLT Sm will overcome these limitations. Indeed, such critical spectral features as the [011]3727 A, the CaH&K, 4000A break, H, G band + H are above 0.S5 I'm for redshifts of 1.2S, 1.14, 1.12, 1.07, 0.96 resp. Moreover, one can expect a S/N=10 for a J=22.5 galaxy observed for Sh, this corresponds to a galaxy with a luminosity 0.5L. at z=1.5. From the current redshift distribution in the most complete survey of more than 700 1< 22 galaxies (CFRS survey at CFHT), one can predict that at J < 22 or H = 21, > 30% of the galaxies will be at z > 1 (see figure 1 b ).

One can identify 4 major themes:

• The evolution of "field" galaxies for z > 1: At a redshift of one, the observed evolution in the field population of galaxies is 1.2 magnitudes (survey CFRS: Le Fevre et al., 1994a; Lilly et al., 1994). What is the history of galaxy evolution

, which has led to this brightening? What is the star formation history in galaxies at redshifts z > 1. Answering these questions from the observations of z > 1 galaxies will only be possible from the IR, and will help to set strong constraints on the epoch of galaxy formation. The VLT will allow to obtain S/N=lO on J=22 or H=21 in 4h integrations (R=200). Since at these 'magnitudes, there are 40000 field galaxies per square degree (Broadhurst et al., 1992), a near-infra-red spectrograph will see about

18

300 galaxies in a field of only 5'x5', of which 30% are expected to be at z > 1. Five nights of observations would provide deep IR images and counts for galaxies down to H=24 in twenty fields 4x5'x5' each and allow target selection for the spectroscopy. With 120 slits per observations in a 4x5.1 'x5.1' field, a ",40 nights allocation would allow to obtain spectra of 10000 galaxies, of which ",3000 will be at z larger than 1 and the rest between 0:5 z :51. This is an outstanding prospect for cosmology.

• The evolution of the clustering of galaxies for z > 1: The latest deep redshift surveys have uncovered structures of galaxies up to red shift unity (Le Fevre et al., 1994c). The form of the clustering at these early epochs will strongly constrain our ideas on the evolution of the universe, and the formation and evolution of clustering. To cover an area of 100 x 10 Mpc2 at z=1.5 (2.5°x15'j Ho. 50, qo=O), one would need to obtain redshifts for 25000 galaxies at H=21 (R=400).' A 4 quadrant 4x5'x5' field spectrograph would allow much progress in this field in less than 100 nights, the size of large key-programs conducted on La Silla.

• The evolution of galaxies in clusters: Rich clusters of galaxies provide one of the most valuable laboratories for studying the evolution of galaxies, in part because they are recognizable out to large redshifts and provide large collections of galaxies at a common distance for convenient observation. The evolution of clusters them­selves has profound cosmological significance: as the most massive gravitationally bound systems in the universe, cluster properties and their evolution are highly sensitive to the physics of cosmic structure formation and the values of the funda­mental cosmological parameters. Our current knowledge of cluster galaxy evolution is mostly based on observations at z<0.5, and only a few clusters at z > 0.5 because of the difficulty in measuring redshifts at these great distances with 4m telescopes. Deeper surveys from future X-ray satellites (e~g. ESA/XMM) will allow to define coherent samples of distant clusters. The VLT will allow to probe clusters of galax­ies, or their progenitors, 3 to 4 magnitudes below the brightest galaxy member at z'" 1-1.5, therefore opening to spectroscopic measurements several tens of galaxies. This would set strong constraints on rich cluster evolution (Peebles et al., 1989). A 4 quadrant spectrograph in 4x5'x5' would allow to obtain on order 50 clusters galaxies spectra in 8h and to explore up to 6Mpc at z=1.5 (Ho=50, qo=O). This will allow to (i) quantify recent and on-going star formation in blue cluster galaxies at high redshift (i.e. extrapolate Butcher-Oemler results to z > 0.8)j (ii) measure spectral indices (e.g. 4000A break strengths, UV fluxes) for red cluster galaxies to probe the earliest epochs of cluster galaxy formationj (iii) Evaluate the dynamical state of rich clusters at very large redshifts via estImates of the cluster velocity dispersions.

• Search for Distant Quasars: The decrease of the QSO luminosity function at z > 3.5, or when the universe had less than 15% of its present age, is of great significance in cosmology. The highest redshift known is now 4.92, but the detection of high-z QSOs is potentially subjected to severe selection effects (the slitless detection of Lya is affected by the OH sky emissions beyond 7200Aj high-z QSOs fall in the stellar locus of visible multicolor diagrams). Ultra high-z QSOs (z > 5) will require IR observations since beyond z=4.8 elv is above 9000A, Lya can be observed in the visible only for z < 6, and the flux below Lya (redshifted above R band) is greatly

19

diminished by Lya forest absorbants. Recently, a previously unnoticed population of red QSOs has been identified (Webster et al., in prep.). QSO candidates can be selected in visible-IR multi color diagrams with 0.5 to 5 bright QSOs (absolute magnitude at rest -28 < MB < -26) predicted per square degree from z=5 to z=7. Such objects would be as bright as J=19-20 (a=-0.5) and can .be imaged in 120sec per filter with SjN",lO. The wide field imaging capability of NIRMOS would allow to obtain imaging of 1 square degree in 2.4 hours or 10 nights to cover a 30 square degree area, and potentially identify 15 to 150 ultra high-z QSOs candidates. The spectroscopic follow-'ap would require ",30min per each candidate to get S j N > 10 (R=1500) and identify these J=I9-20 objects (emission lines). Detection and studies of such objects when the universe at 15% of its present age is a challenge that the VLT can tackle only with a wide field, near-IR imaging-spectrograph.

2.2 Gravitational lensing and the distribution of dark matter

Gravitational lensing is a unique means to investigate the mass distribution in galaxies and galaxy clusters. A large fraction of multiple QSOs have no identified lens at visible wavelengths, which is likely related to the fact that these lenses are too faint and redshifted at z > 1 to be detected in the visible. Several arcs have now been found in very distant clusters (z > 0.6; Luppino, 1994, in preparation) and are most probably distorted images of z > 1 galaxies. Imagery and spectroscopy of lensing galaxies, arcs and cluster galaxies in the field would provide the first view of the dark matter distribution from z=l to z=2. Lensing galaxies responsible for multiple QSOs are generally L > L. galaxies, and then can be spectroscopically identified from z=1 to z=2 in less than 4h (J < 22 and H < 21; R=200). The multi-slit capability of an IR spectrograph would allow to study in detail the redshift distribution of other objects acting as lenses in the field, as well as arcs and very distant field galaxies magnified by the cluster gravitational field.

2.3 Objects identified with IRAS or ISO

Some of the objects that have been studied by IRAS or will be by ISO, will have red spectral energy distributions that will require near-IR spectroscopy to better understand the physical phenomena at play.

The ISO guaranteed time programs include e.g. deep survey programs that are ex­pected to uncover the evolutionary properties of starburst and active galactic nuclei at redshifts is excess of one. One also expects to identify primeval galaxies at very high redshifts in their first burst of star formation.

These surveys are expected to cover", 0.1502 at large sensitivity and '" 102 at moderate sensitivity and several hundred sources are expected to be identified. NIRMOS will then be most appropriate to follow-up the source identification by ISO, expected to be ",60 in the NIRMOS field~ to obtain the redshift and the spectral line information that will be critically needed to interpret ISO data.

2.4 Studies of sub-stellar objects in nearby star clusters

During the initial phases of their contraction, sub-stellar "brown dwarf" objects are as bright as main sequence stars. In the near IR spectral domain, numerous temperature and

20

120 eFRs: 616 galaxies

100

80 104 JDCOmp.

60

40

20

o o 0.5 1 1.5 Redshift

Figure la: Redshift distribution of >600 galaxies observed for the· Canada-France Redshift survey. The sample is magnitude limited (I~22), and ",Sh integrations were needed on the CFHT to obtain sufficient SIN to measure redshifts up t'o 85% completeness. This sample represents the current limit of 4m class telescopes in the measurement of faint galaxies redshifts obtained with state of the art multi-slit spectrographs (MOS-SIS with "'SO slits per integration). Only the Sm class telescopes will be able to break the z'" 1 barrier to study galaxy evolution at earlier epochs.

21

Recl8hIR dldrlbuUon (evol - U m .. for .>1.3)

- - H<21

'--1<22 .... :;: . .:,:., .. " "

10

10

······1<23

-1<22CFRS

o

:'1 ".' :'1 ' ......... ""

.:, I " .... ' .:/ I \,~, .:, '\ '

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'. '.

, , " , ,

" , , eo 0 .................. -0 ••

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°0~~~~~70~~~~~~~1~~~~--·1~--------~2

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Figure 1b: Predictions for the redshift distribution of galaxies selected brighter than H=21 or J=22, compared to the distribution expected for 1<23. (a) The amount of evolution used for the predictions is the same as what is observed for the CFHT/CFRS survey to z=l (the prediction for 1=22 is very close to the observations of this survey), with 1.3 mag of brightening at z=l; a very conservative amount of brightening equal to what is observed at z=l has been set for z>1.3 (no additional evolution above z=i). This gives a pessimistic view on the fraction of galaxies expected at z>l. (b) The amount of evolution observed in the CFRS survey has been fitted by a low order polynomial to reproduce the observed 1.3 mag of evolution at z=l, this would give a luminosity brightening of 2.1 mag by z=2. This clearly shows that the selection of galaxies in the near-IR allows to identify in the most conservative scenario more than 30% of the galaxies at redshifts larger than 1.

surface gravity sensitive transitions can be used to distinguish between sub-stellar objects and foreground dwarfs. Moreover, most of the opacity is coming from the water bands which dominate the near-IR spectral energy distribution. ·These features are of critical importance to compute accurate luminosities and effective tempera.tures. Young star clusters provides a very promising environment for brown dwarfs searches, recent surveys in found 1 K=15 brown dwarf candidate per 5 square arcmin (Jones et al., 1994; Comeron et al., 1993). To H=21, brown dwarf candidates could be observed out to distances of 2000 parsecs, with more than a thousand objects per 5'x5' field. 3 nights would allow 15 observations each of 3h and 50 slits, hence more than 700 such candidates.

3 Instrument specifications There are several main drivers to the specifications of a near-IR multi-object imaging­spectrograph to fulfill the above science goals:

• Need for efficient observing modes both in wide field imaging and wide field multi­slit spectroscopy. The possibility to switch from one mode to the other in time scales around a minute is essential. • The density of faint field galaxies, cluster galaxies, sub-stellar objects, stars in clusters or in the Magellanic clouds is high enough to allow large gains from multi-object spec­troscopy techniques as presently used in the visible. • With high transmission, one expects in 4h to get S/N=10 for J=22 or H=21. Spec­troscopy of L. galaxies at z > 1 will therefore become possible in the near-IR (see Fig. 2). • A large field allows (i) more objects to be observed at once, and (ii) to place more slits if the detector area is increased. The large unvignetted field at the Nasmyth focus of one of the VLT units is most appropriate. • The wavelength range from 0.8 to 2JLm can be observed without a fully cooled instru­ment. Flexibility in aperture mask installation is critical for the efficiency of observations, and should remain a priority. Simple cryogenic solutions keeping this in mind are to be studied to cool the mask environment and possibly lower the thermal background by 2/3. • Multi-spectroscopy is a technique which works most efficiently with slits, because of their high transmission, and excellent sky subtraction. A direct imaging capability to identify targets and prepare slit masks, coupled to the ability to switch from the imaging mode to spectroscopic modes, is essential. • Good sampling of the observed objects and the slit width is necessary to obtain good sky correction. An excellent compromise is 0.3" /pixels. which allows excellent sampling of slits 1" in width, and also of the site seeing. • Low to medium resolutions are adequate to measure velocities with accuracies on order 100 km/s; these can be achieved with standard gratings. Resolutions between 200 and 3500 could be obtained for a 1 arcsec slit.

4 Design solution A 4 quadrant spectrograph with 4 channels each covering a 5'x5' field would answer all of the above design goals and be extremely competitive. A design such as WFIS (Delabre et al., 1994), a wide field multi-object spectrograph for the visible region, is also adequate

23

for the near-IR, and such a solution is proposed here. The specifications are: • Imaging spectrograph located at the Nasmyth focus. Main modes of operation: wide field imaging, long slit spectroscopy, multi-slit spectroscopy. • wavelength range = 0.8 to 1.8 pm (possible extension to 2.3 p) • Field 4x5.1'x5.1' (limited by detectors, larger if2048x2048 arrays become available). • Detector: 4 Rockwell HgCdTe arrays, each 1024x1024 pixels, 18.5 pm/pix • pixel scale = 0.3 arcsec • Resolving power = 200 to 3500 • instrument image quality = better than 0.2" anywhere • Multi-slit masks at entrance focal plane • Flexure: less than 18.5 pm for a 4h motion in H.A.

There is a potential technical solution to expand the wavelength range of such a spectrograph to the K band up to 2.3 pm. Thermoelectric coolers can potentially be used to cool the slit environment to temperatures around -45°C, therefore reducing the thermal background seen by the detector arrays by more than 2/3 to allow for K band spectroscopy. We strongly emphasize the need for a detailed feasibility study on this possibility.

5 Instrument performance

5.1 Direct imaging The limiting magnitude (S/N=3) in Ih is expected to be J=25, H=24, K'=22 (assuming the sky brightness to be H=14.4, J=15.6). With thermo-electric cooling of the slits, the limiting magnitude in K' would be K'=23.

5.2 Spectroscopy The limiting magnitude (S/N=lO) in 4h is expected to be J=22, H=21, K'=19. for R=200 (assuming the sky brightness to be H=14.4, J=15.6), see Fig. 2. With thermo-electric cooling of the slits, the limiting magnitude in K would be K=20.

The average number of slits would be on order 120 simultaneous slits, or about 30 slits 10 arcsec in length per each of the 5.1'x5.1' fields. At the lowest resolution R=200, one might be able to fit two "stages" of spectra on one of the detectors, as is done successfully in the visible, therefore allowing a maximum of 240 slits.

Sky correction can be done in a classical way (each slit treated as a long slit). Recent observations in the near-IR show that sky correction can be achieved to levels of less than 0.5% of the sky brightness in J and H (Goodrich et al., 1994).

6 Technical concept

6.1 Optical design An optical solution for WFIS, a wide field visible 4 quadrant spectrograph, has been presented by Delabre, D'Odorico, Vettolani at the SPIE conference in March 1994. The

24

6.3.2 Detectors

Each detector will need to be aligned with respect to the focal plane multi-slits masks to a precision of 0.5 pixels over 1024 or to better than 0.0280

• Distortion within the instrument, if any, will need to be mapped for each quadrant via the use of focal plane masks with square reference grids.

6.3.3 Data processing

The number of photons per pixel element is indicated in Table 2 for various modes of observations, together with the maximum exposure time possible with the Rockwell array (number of electrons per pixel = 2/3 full well of 105e-).

Mode e- /pix Max. expo time (sec.) J band imaging 1.5xl04 4.5 H band imaging 2.1x104 3.2 J band spectro (R=400) 30 2230 H band spectro (R=400) 42 1600

Table 2: Max. exposure times

The operational constraints will then be for direct imaging:

• Fast readou.t time for the full array (4 quadrant each). Each quadrant of 512x512 pixels can be read in about 1.5 sec. in imaging mode (based upon the current readout rate of the Redeye/256x256 NICMOS cameras at CFHT). • narrow filters for imaging: a O.IJL filter in the J or H band will increase the maximum exposure time by a factor ",3 • Real time produ.ction of averaged images: would greatly facilitate the identification of sky targets prior to spectroscopy. Need a fast processor to average several tens of 4xl024x1024 images in a few minutes elapsed time.

The operational constraints are much less stringent for spectroscopy: one would need to be able to evaluate the data quality in a timely fashion on spectra roughly sky corrected and wavelength calibrated. This is currently achieved in time scales on order 1 min with present data hardware/software.

6.3.4 Multi-slit masks

The same slit cutting machine could be used for NIRMOS and a large field visible multi­slit spectrograph such as WFIS. Since it might be impossible to install the slit cutting machine in one of the 8m domes, provision should be taken to 'allow for several (3) masks to be loaded at one time. .

The flexibility to install slit masks should remain at any time. Positioning on the sky will be done via 2 reference holes cut into the aperture masks, at the actual focal plane location of suitably bright stars (J=18 'or above would be adequate for short positioning exposures) .

25

solution for the near-IR spectrograph presented here is of an almost identical concept (see attached document by Delabre). The spectral domain is from 0:8 to 1.8 p.m, the total field covered is 10.2'xl0.2' with spot diagrams indicating an intrinsic instrument image quality better than 0.2 arcsec in imaging as well as spectroscopy mode.

We note that the ,similarity of the optical concepts of WFIS and NIRMOS could lead to a substantial reduction in development time and expenses as the mechanical frames of the two spectrographs could be made nearly identical.

6.2 Detector specifications

The detector considered here is the 1024xl024, 18.5p.m pixels array presently offered by Rockwell International. Rockwell has delivered the first arrays this summer to the university of Hawaii, and the measured performances are excellent. Outstanding images have been obtained with the full array, e.g. on the Shoemaker-Levy 9 comet impact on Jupiter (Fig. 3; Kozlowski et al., 1994). The detector specifications as measured on this first array (Kozlowski et al., 1994) are given in Table 1.

Format (pixels) 1024x1024 Pixel pitch (p.m) 18.5 Fill factor >95% Quantum efficiency (77K) @0.8 p.m 50% @1.2/-tm 60% @2.35 /-tm 65% Long wavelength cutoff 2.5J.tm Read noise (e- @ 77 K) 8.6 Dark current (e- Isec @ 77 K) <0.1 Well capacity (e-) 105

Yield (working pixels) >99% Operating temperature >77K Number of outputs 4

Table 1: Rockwelll024xl024 HgCdTe array performances

We feel that this array is presently suitable for use in NIRMOS. Moreover, we have in mind that 2048x2048 HgCdTe arrays are in the plans of Rockwell, and the design of NIRMOS will have provisions for use of the larger arrays if they become available: the field size could be increased to cover in excess of 15'xI5'.

6.3, Operational requirements

6.3.1 Instrument configuration

The instrument hardware and software control should allow to switch from the imaging mode to the spectroscopy mode, without interfering with telescope guiding. The time scale for instrument configuration changes should be on order 1 min.

26

:z "-rJl

:z ~

SIN per resol. element. 1 VLT 8m

104

R-2OO R.O.N.-IO e-

1000 Dark-O.2 e -Iseo

" Alm.+Te1.+NIRMOS-12X " object on 4xO.3 pixels .....

...... ......

Wac sky - 14.4 lOa ...... ......

....

10

1 15 20

H magnitude

104

R-2oo R.O.N.-l0 e-

1000 '- Dark-O.2 e-Iseo 4m (1h) " - Alm.+Tel.+NIRWOS-12X .....

..... object on 4xO.3 pixels " ..... 100 .... Wa, sky ~ 15.6

.... ..... ", ...

10 '.

. ...

1 15 20

J magnitude

Figure 2: Signal to noise ratio for spectroscopy in the low resolution mode of NIRMOS. Param­eters used for the computation are identified in the figure. The full line is for one of the VLT 8m, the dashed line is for a 4m telescope.

27

Figure 3: Jupiter after the Shoemaker-Levy 9 comet collisions. F'irst. image obtained on the Ull 2.2m telescope with the 1024x1024, IlgCdTe array developed by Rockwell and the Uni versity of Hawaii (Kozlowski e t aI., 1994)

28

Slit masks with approx. 120 slits would be preferably cut in time scales on order 15 min. The accuracy of slit cutting is to be better than 10 p.m for high quality slits, which is well within the capability of conventional high-precision milling machines (Laser cutting

. allows 1p.m accuracy, which would be an overkill in this case).

7 Cost and manpower We mention here only the cost of the detector arrays. An estimate of cost and manpower for the full instrument should await a more detailed technical review based upon the specifications defined in the present document.

• The cost of 4 Rockwell1024x1024 science grade arrays and 2 engineering grade ar-' rays will be approximately (direct quote from Rockwell to us) = 0.8 MDM.

8 Comparison with existing or planned instruments To our knowledge, there are only two other instruments which have been identified for multi-slit spectroscopy in the near-IR: the OSIS spectrograph upgrade of the SIS spectro­graph at CFHT and the optional upgrade to the MOS spectrograph presently designed for the GEMINI project to cover up to 1.8p.m. The fields of both of these instruments is less than 5'x5', which will leave NIRMOS unsurpassed for survey work. Note that O. Le Fevre is co-P.I. on the OSIS upgrade at CFHT, and that the experience gained for this instrument will benefit NIRMOS (OSIS will be ready for CFHT by the end of 1995).

Acknowledgments We thank S. D'Odorico for his in-depth comments on this proposal as it was being pre­pared.

References Broadhurst, Ellis, Glazebrok, 1992, Nature, 355, 55. Colless, proc. of the 35th Herstmonceux Conf. "wide field spectroscopy and the distant

universe", Cambridge, 4-8 July, 1994. . Comeron, Riecke, Burrows, Riecke, 1993, ApJ, 416, 185. Delabre, D'Odorico, Vettolani, 1994, SPIE symposium "Instrumentation for the 21st

Century" . Goodrich, Veilleux, Hill, 1994, ApJ, 422, 521. Gunn, 1989, in "clusters of galaxies", STScI series 4, p.341. Jones, Miller, Glazebrook, MNRAS, 270, L47. Kozlowski, L.J., Vural, K., et al., 1994, proc. of the International Symposium on Optics, Imaging and Instrumentation, SPIE, 26-29 July 1994, San Diego, California (in press) Le Fevre, Crampton, Hammer, Lilly, Tresse, 1994a, these procedings Le Fevre, Crampton, Felenbok, Monnet, 1994b, A&A, April 20.

29

Le Pevre, Cra.mpton, Hammer, Lilly, Tresse, 1994c, ApJ , 423, L89. Lilly, Le Fevre, Crampton, Hammer, Tresse, proc. of tbe 35th Herstrnonceux Conf. "wide

field spectroscopy and t.he dist.ant. universe", Cambridge, 4-8 July, 1994. Luppino, 1994, pr ivate communication. Peebles, Daly, J uszkiewics, 1989, ApJ , 347, 563.

30

APPENDIX'

NIRMOS • Preliminary optical design of a large field multi slit Imaging spectrograph for the range 0.8 to 1.8 J1m

B.Delabre - ESO Garching· 13-04-94

The concept of the Wide field imaging spectrograph (see paper Optical design of a wide field imaging spectrograph for the VLT , Delabre, D'Odorico, Vettolani. SPIE Conference Hawaii 1994) has been adapted to the near infrared (range 0.8 to 1.8 f.UIl).

optical characteristics :

POV (arc:min) S.l x S.1 arcmin (4 times) POV(mm) 183 x 183 nun{4 times) ~il diameter 131mm Final P/ntio 1.60 Pinal scale 0.3 _~~!;.....I (18.~) Pcamera 210 nun Detector 1024 x 1024 (18.S JUll) RockweU HgCdTe Wavelength range 0.8 to 1.8 JUll Spectral coverage with 300 DID assuming slit in the middle of the field of view 300 mm-l grating (k=I) R. 2800 for 0.7S arc:sec slit at 1.6J.UD

Optical quality :

The average optical quality is 10 f.UIl nns polychromatic spot diameter (0.16 arcsec) in imaging mode. The quality of the spectroscopic mode is around 18 JUn (0.3 arcsec).

Field 1.00 2.20 3.SS 4.80 6.10 arcmio

6.10 12 9 8 9 12 4.80 9 9 8 8 9 3.SS 10 10 8 8 9 2.20 9 10 10 9 9 1.00 10 9 10 10 12

Rms spot diameter (polychromatic) for direct imaging mode in J1l1l

31

graling 300 mm- l , 1st order . used al i = 38° and j' = 7° R= 2800 fo r 0.75 arcsec slit

dispersion 1495 om 1570 om 1645 run 1720 run 1795 nm slit -9.533 mm -4.750 mm 0.027 mm 4.794mm 9.550 mm 2.6 arcmin 20 16 14 16 1.3 arcmin 17 16 II 17 axis 18 15 8 18

-1.3 acemia 18 16 10 18 -2.6 acemla 20 16 14 17

Rms Spot dianleler for spectroscopic mode in J.Lm

graling 300 mm- l, I SI o rder. used at i = 31.5° and j' = 13.5° R= 1600 for 0.75 arcsec slit

dispersion 8 15 nm 887 run 960 run 1036 om

16 18 18 18 17

lJI2nm slit -9.588 mm -4.912 mm -0. 185 mOl 4.7 18 mm 9.597 mm 2.6 acemin 22 16 15 15 18 1.3 acemia IS 14 12 14 IS ax~ 14 12 II 14 16

-1.3 arcmin IS 14 13 IS 16 -2.6 aremin 22 16 IS 16 21 Rms spot diameter fo r spectroscopic mode in J.Lm

TIle 4 idenlicaJ instru ment channels are amUlged over the modified Nasmyth focus as shown in Ihe foUowing figure:

Remarks:

5. 1 arcmin

EJEJ !S I=m;n 20rc mln r - :-..

EJ'E] ,-

2 0rcmrn

II is not possible to extend the wavelength range to 2.31J.1ll with the materials used in this preliminary design. Class ical infrared material s like Baf2 or Lif will be required. As no indices of refraction are given LUlIil 1.8 !.tm , CodeY exlrapolalion has been used. This extrapolation method may produce significanl erro r. A detector with larger puc! will simplify rhe design of the camera.

32

Efficiency:

The material used in this preliminary design have some absoIption specially around 1.8 J1m:

L wavelenath 800 om 1000 om 1200 DID 1500mn I 1800 DID I I absOlplioo S% S% S% 10% I 15% I

The following table gives the efficiency of the system with single layer MgFl coating: (better coatings can certainly be obtained)

I wavelength I 800 DID I 1000 mn I 1200 mn I 1500 mn 1800 om ItraDsm. I 62% I 76% I 80% I 72% 61%

List of figures:

1- Layout of a single channel 2- Spot diagrams for direct imaging 3- Spot diagrams for spectroscopic mode

33

09:54:21

Folding mirror

Collimator

(L1F20)

aspheric

(KZFSN2 - FK54)

Camera F/l.60 all surfaces are spherical (FK5-LAF20-FK54-LAF20-KZFSN5-FK54 )

F/15.45

Adaptor flange

Focal plane adaptation lens (SSK3-SF57)

Flat mirror or grating

250.00 MM

13-Apr-94

1.00 arcmin 2.20 arcmin 3.SS arcmin 4.80 arcmin

. ~ .V!9

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WFIS - Near infrared version (0.8 to 1.8 J.1m) Polychromatic spot diagrams (direct imaging) Field values are referenced to the telescope optical axis

35

.~=-.~ .. :: .. " "'. ,. , .

" , ..

6.10 arcmin

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~ ~ .. '.

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.

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Dispersion

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-

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@

WFIS - Near inrrared version (O.S to I.Sl1m) Spectroscopic mode with 300 mm-I grating used at i :;; 38° and i' :;; 7°

36

1720 nm 4.794 nun

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m

VHRS : THE VERY HIGH RESOLUTION

SPECTROGRAPH FOR THE VLT

P. Magain1 , H. Dekker2, B. Delabre2

1 Astrophysical Institute, University of Liege 2 ESO Garching

1 Introduction

One of the best opportunities provided by the large collecting area of the VLT is the ability to obtain spectra at very high spectral resolution, R> 3 X 106• Up to.now, only the UHRF at the AAT allows to explore this range of spectral resolution. However, its low efficiency limits the possibilities to the very brightest stars, therefore narrowing drastically the kind of science that can be carried out.

With an 8m telescope and an efficient instrument, very high resolution observations could be carried out for objects as faint as V = 17, opening the extragalactic universe to this kind of observations, which were traditionally reserved for studies of stars or interstellar matter.

A point which also merits consideration, in view of maximizing the scientific productivity of the coming generation of 8-10m class telescopes, is the complementarity in the instru­mentation between the different telescopes. Indeed, it would not appear very wise to equip all these telescopes with the same kind of instruments, while leaving whole important areas totally unexploited.

According to our knowledge of current projects, the very high resolution spectrograph (VHRS) of the VLT would be the only instrument allowing to explore this spectral resO­lution range. It would thus provide the European astronomical community with an unique opportunity to take the lead in an important area of science, an area, moreover, in which it has already developed considerable expertise.

After the ESO Worskshop on Science with the VLT, it appeared quite clear that the majority of scientific programs concerned require a resolving power either approaching R = 300000 or around R = 600000, these two resolution ranges separating roughly the stellar programs from the interstellar/planetary ones. Other important considerations are the instrumental stability, a clean and well known instrumental profile, and the ability to reach very high signal-to-noise ratios.

In the following section, we list a number of scientific programs which could be carried out with such a very high resolution spectrograph, while the last section presents a preliminary conceptual design for such an instrument.

37

2 Science objectives

2.1 Stars

2.1.1 Dynamics of stellar atmospheres

Numerical hydrodynamical simulations of stellar convection patterns are reaching a remark­able level of refinement for the Sun and indicate observable diagnostics in line profiles for the characteristic features of convection patterns in other stars (Dravins & Nordlund, A&A 228, 184, 1990).

Very high spectral resolution and SIN ratio are needed, as well as a clean and well-known instrumental profile, so that data now exist only for a few of the very brightest stars. Further studies of fainter stars (in particular metal-poor stars) would be of great significance, not only for a better understanding of convection and other dynamic phenomena in stellar atmo­spheres of different types, but also for more precise determinations of chemical abundances and other fundamental stellar parameters.

Very detailed line profiles will also lead to a better determination and modelling of the mass loss from cool AGB stars. This in tum would considerably improve our understanding of nucleosynthesis and galactic evolution. High resolution and SIN are essential, and the VLT as well, since most samples of the key stellar populations (e.g. halo and bulge) are distant and faint.

2.1.2 The first stars

One of the main advantages of the large collecting area of the VLT is that it will allow to reach fainter objects. As the number of stars of a given metallicity decreases strongly with the metallicity, going fainter also means going more metal-poor and, thus, closer to the primordial nucleosynthesis. One can also hope to identify objects whose atmospheric composition would be the result of a single post-Big Bang nucleosynthesis event (supermassive star, quasar, ... 1). .

However, the spectral lines of these extremely metal-poor objects become so faint that high resolution (R > 100000) and very high SIN (> 500) are required to be able to measure the line equivalent widths with acceptable accuracy. Access to the UV spectral range is also mandatory.

2.1.3 Abundance diagnostics of stellar evolution

As discussed most recently, e.g. by Charbonnel (A&A 1994, in press), isotope ratios of. several key elements like C, N., 0, AI, ... , in addition to the well-known examples of Li, Be, and B, can serve as detailed diagnostics of stellar model calculations in regions of the

38

HR diagram where other empirical information, e.g. from membership of binary systems or clusters, is missing.

As these isotope ratios are particularly sensitive indicators of the degrees of mixing and mass loss in the envelopes of giant stars, the information they provide on nucleosynthesis and enrichment of the interstellar medium is particularly relevant to models of the chemical evolution of the Milky Way Galaxy (see e.g. Edvardsson et al., A&A 275, 101, 1993).

Since this has now been shown to be different in different parts of the disk, particularly towards the inner disk and bulge (see also McWilliam & Rich, ApJS 91, 749, 1994), reaching distant and faint stars with the VLT is essential.

2.1.4 The 6Lij1Li isotopic ratio in the oldest stars

Recently the 6Li isotope has been detected in two Pop. II stars (Smith, Lambert & Nissen, ApJ 408, 262, 1993, Hobbs and Thorburn, ApJ 428, L25, 1994), with far reaching con­sequences for our understanding of Big Bang nucleosynthesis and early galactic evolution. The isotopic ratio is derived from very accurate observations of the profile of the lithium resonance line at 6707 A.

With the ESO 3.6m telescope and the fiber link to the CES the lithium line has recently been studied at a resolution of 110000 and SIN = 400 for stars of magnitude V=9.0 using total exposure times of the order of 10 hours. However, this leaves us with only a couple of Pop. II stars at the turnoff point, where 6Li probably survives depletion at the bottom of the convection zone. The VLT and a dedicated very high resolution spectrometer is needed to reach a reasonable sample of very metal deficient stars.

2.1.5 Heavy elements isotopes

The elements heavier than the iron peak are synthesised by neutron capture, either through the s-process (which is though to operate mainly in rather low mass AGB stars) or through the r-process (which might take place in supernova events). While the most abundant of these elements in solar-system material are mainly produced by the s-process, Truran (1981, A&A 97, 391) suggested that, in metal-poor stars, this is the r-process which is mainly responsible for the production of even such traditional s-process elements as barium.

Given the much shorter timescale of the r-process, this suggestion appeared quite rea­sonable and was apparently confirmed by subsequent spectroscopic analyses (e.g. Gilroy et al., ApJ 327, 298, 1988). However, this test is rather indirect and the spectroscopic analyses subject to large uncertainties.

A much better test would be the determination of the isotopic composition. In the case of barium, the even isotopes are mostly produced by the s-process while the r-process contributes significantly to the odd isotopes. Unlike the light elements, the isotopic shifts of the lines of these heavy elements are completely negligible. However, many of them show

39

hyperfine structure, which strongly depends on the parity of the nucleus.

This effect was used by Magain and Zhao (1993, A&A 268, L27) and Magain (1994, in preparation) to determine the fraction of odd isotopes of barium in the metal-poor star HD140283. This analysis, based on more than 30 hours of exposure time with the CES at R = 100000, shows that, at least for this star, the s-process alone accounts quite well for the isotopic mix.

To rea.ch a good accuracy for a significant number of very metal-poor stars, the VLT, with R '" 200000 and S / N > 500, is necessary.

2.1.6 Thorium abundances in F and G dwarfs

As shown by Butcher (Nature 328, 127, 1987) and Morell et al. (A&A 259, 543, 1992), the abundance of the radioactive element 232Th is particularly interesting for studies of the nucleosynthesis of r-process elements and determinations of the age of the Galaxy.

However, thorium abundances can only be determined from a faint Thn line at 4019.13A, which is blended by other lines. Thus very high resolution and SIN observations of the profile of the 4019A feature are needed.

With the CES it has not been possible to derive thorium abundances for the metal-poor F and G stars. Hence, the VLT and a dedicated high resolution spectrometer is needed to study the nucleosynthesis of thorium in the early Galaxy and to get an independent measure of the ages of the halo stars.

2.2 (Proto-) planetary systems,

Observations at high spectral resolution, high SIN and high temporal resolution are needed to better understand the protoplanetary disks, the prototype of which is found around the star f3 Pictoris (Lagrange-Henri et al., A&A 264, 637, 1992).

Spectral observations at very high resolution and very high stability also open the possi­bility to detect the presence of planets around stars by measuring the periodic radial velocity shifts in the stellar spectrum induced by the planets orbiting around it.

The amplitude of the radial velocity variations of the Sun due to the presence of Jupiter amounts to 13 m/s. On the other hand, the perturbation caused by the Earthis only 0.1 m/s. If this method appears feasible in the case of the giant planets, it seems to be out of the reach of terrestrial planets.

At a resolving power of 300000, a precision of 5 mls should be attained for a single absorption line if the S / N in the continuum reaches 500, which is quite feasible. Considering more lines allows to increase the precision as the square root of the number of lines conside~ed.

40

On the other hand, simulations indicate that, in this range of R, the single line precision is roughly proportional to R. For a fixed detector length, the precision will thus (as the number of lines available decrease with R) increase as y'R, at a given S / N. At R f"oJ 600000, it should thus be possible to a.pproach an accuracy of 1 m/s. This assumes that stability is not a problem. Anyhow, for some stars, interstellar lines superimposed on the stellar spectrum, if separated from the stellar lines, could provide a reference which is independent of the stability of the spectrograph.

Moreover, working at very high resolution has the essential advantage that the full line profile can be monitored, thus allowing to disentangle line shifts due to photospheric effects (oscillations, turbulence, ... ) from those caused by the presence of a planet.

2.3 Interstellar medium

The study of the interstellar medium is obviously one of the fields which would benefit most from a very high resolution spectroscopic facility. Indeed, typical interstellar absorption lines are very narrow and show several components which are completely blended at usual resolutions. With R '" 600000, it becomes possible to measure the individual components, and to derive important physical parameters of the clouds, such as the temperature and the importance of turbulent motions (e.g. Ferlet, The Messenger 73, 25, 1993).

The use of a large telescope would also allow to observe fainter objects, such as stars deeply embedded in cloud cores. The properties of these dense clouds could then be compared to those of the more commonly studied diffuse clouds.

Observations at very high resolution and S / N would also allow the isotopic ratios of lithium and beryllium to be measured in a variety of interstellar clouds, thus providing important constraints on the cosmological scenarios as well as on galactic chemical evolution models.

2.4 Extragalactic

A very important application of a very high resolution spectrograph fed by a large telescope lies in the detailed study of the narrow absorption lines in quasar spectra.

A very exciting program is the measurement of the temperature of the cosmic microwave background TOMB as a function of redshift. Locally, interstellar lines from the lowest rota­tionallevels of the eN molecule can be used to derive a precise value of TOMB (e.g. Palazzi et al. ApJ 398, 53, 1992). At high redshift, neutral carbon lines can be used for the same purpose. Although the required resolution is somewhat lower (Songaila et aI. Nature 371, 43, 1994), very high resolution would provide a much better separation of the possible blends and, thus, more reliable results.

Another key programme would be the determination of the deuterium-to-hydrogen ratio D /H in a variety of sites. Indeed, the strong - and uncertain - variation of this ratio with

41

time calls for observations in media as primordial as possible, if one wishes to derive secure values for the cosmological D IH ratio, which is a key test of cosmological scenarios. On the other hand, measurements of this ratio in clouds of different metallicities would allow to study the galactic evolution of deuterium. Here again, high resolution is essential for a reliable separation of the deuterium lines from the crowd of hydrogen lines in the Lyman-a forest.

3 Conceptual design

3.1 Introduction

To reach R = 600000 with a slit of 0.6" - equal to the median size of the seeing disc - on an 8m telescope requires a grating with a length of 7m. Since this is not feasible, narrower slits must be used and the question should be asked how one is going to avoid losing too much light on the slit jaws. Available solutions are image slicers and adaptive optics systems. At the present moment, no adaptive optics system for the blue/visual/NIR region has been demonstrated, nor does any group plan such a system for an 8m telescope. Plans do exist for an adaptive optics system working in the 0.5 - 1 /-Lm range feeding a simple high resolution spectrograph on a 1.5m telescope (R. Angel at the conference on High Resolution Spectroscopy with Very Large Telescopes, Tucson, Oct. 1994). For this study, we have restricted ourselves to image slicers, being the only viable solution for the near- and mid-term.

Adaptive optics systems might become available during the lifetime of the instrument. This will not impact so much the efficiency since slicers are quite efficient. The main impact of an AO system on the type of science that can be done with the VHRS is that it reduces the vertical extent of the spectrum on the CCD. This way it will improve the sensitivity at low SIN (where detector noise is limiting due to the large spread of light in the vertical direction) and it will give room for increasing the spectral coverage and obtaining spatial information in a crossdispersed echelle format. In a further study, it should be verified to which extent the VHRS design could or should anticipate feeding by a future AO system.

This preliminary conceptual design is an "existence proof" that an instrument with R = 600000 is feasible given the space and weight constraints at the N asmyth or other foci of the VLT. The grating size directly affects the amount of slicing and the performance in detector-noise limited SIN regime. Thus the concept also puts in evidence the tradeoff that must be made between grating size and cost, level of slicing, SIN requirements and detector noise properties, which are as important as the spectral resolution in determining the characteristics of the instrument.

3.2 First-order design parameters

Spectral format

42

Similar to the CES, we have opted for a predispersed echelle and a linear spectral format. The linear format helps in the data extraction process and allows good control over the instrument line spread function after extraction. It is also the only format compatible with the use of image slicers. At R = 300000, the 0.28" slit is adequately sampled by 3 pixels of 0.092". If the slit width is reduced to 0.14" to reach R = 600000, it is only sampled by 1.5 pixels which is insufficient for good reconstruction of the line profiles. A possible solutio~ would consist in applying a tilt of one pixel (top to bottom) to the slit (which is 69 pixels high at R = 600000) which would allow sub-pixel sampling of the line profile.

A few spectra can be accommodated within the vertical dimension of the chip. These could be sky spectra, spectra from different objects or fibers, or spectra from the same object taken by different unit telescopes. A very attractive possibility, to be investigated, would be to take spectra of the same object in different orders, introducing a multi-wavelength capability.

Table 1. First-order design parameters

Spectrograph type Predispersed echelle (CES-like) with linear data format and little or no tilt and distortion, etc., of spectral lines

Slit Image slicer, entrance aperture'" 1.4 x 1.6" with 5 slices: exit slit 0.28 x 8", R = 300000, sampling 3 x 34 pixels

with 10 slices: exit slit 0.14 x 16", R = 600000, sampling 1.5 x 69 pix.

Grating R4 echelle, size 1.70 x 0.2 m (elliptical beam 40 x 20 cm, 4 x 1 mosaic)

Camera F / 4.2 in dispersion direction F /1.7 in slit direction (anamorphic) field size'" 15 x 30(60) mm

Detector min. 1 K x 2 K CeD with 15JLm pixels Wavelength range 0.3 - 1.1JLm. Extension to 1.8JLm feasible Resolution Wavelength bin 6.7 mA at 6000A ~~ 0.23% of central wavelength, e.g. 13.7 Aat 6000A sampling single-pixel resolution 900000

Scale 0.23" Ipixel in Y Detection efficiency > 8% in V (top of the atmosphere to detected photons) Sensitivity I V = 17: SIN", 20 in 2 hrs limiting magnitude V = 12: SIN", 200 in 2 hrs

V = 7: SIN", 200 in 72 sec when shot noise limited

Location

For this photon-starved, detector limited application, the need to have as few reflec­tions as possible between the slit and Ml is pressing. Placing it at the Nasmyth focus

43

might be the best compromise between stability, space and weight requirements and opti­cal efficiency. Because of the compact optical arrangement, also other locations could be envisaged (Cassegrain, fiber-coupled Cassegrain, fiber coupled Nasmyth, Coude, combined Coude) without the need to substantially change the instrument. The final choice of location should take into account vibrations, thermal and gravitational stability; factors that could not be included at this point. The sensitivity calculations have been carried out assuming a location at N asmyth and no fiber feed. The layout is shown in Figure 1. The instrument "footprint" at Nasmyth is seen in Fig. 2.

The design parameters are summarized in Table 1. In the next section, the choices for the various components will be discussed.

3.3 Instrument description

The concept for the VHRS consists of a predisperser and a main spectrograph and is similar to that of the CES. We have investigated the 300 - 1100 nm range but, through the use of Silica lens optics and mirrors, it is not limited to this range.

The echelle is a mosaic of 4 R4 segments, each 40 x 20 cm in size, with a total length of f'V 1. 70 m. The echelle is used in-plane, with a very small angle of about 2 degrees be­tween incident and diffracted beams. This ensures the best possible geometry of spectral lines which is necessary for determining accurate line profiles while summing a fair number of pixels in the Y direction. This large size of 1. 7 m is the maximum we think might fit the budget and also space requirements; it is a factor 2 larger than the UVES monolithic echelles and a factor 1.4 larger than the HIRES echelle, which is an assembled mosaic. The groove density of the echelle is freely selectable but is likely to be set ~y the red and blue UVES rulings that are on order: 31.6 g/mrn (already ruled) or 41.6 g/mm. With a 40 cm collimated beam, the resolution-slit product is 84000, which means that this grating will produce R = 600000 with a 0.14" x 16" slicer slit (assumed slicer entrance aperture about 2 seeing disks; 1.5 x 1.5"); which implies a still moderate ten-fold slicing. With this level of slicing, requirements on the detector properties in terms of chip size, readout noise, cosmic rays and dark current are manageable even at low SIN levels as will be shown below. For instance with a 4 x smaller echelle, the slit width must be reduced to 0;035 arcsec to reach R = 600000 - diffraction effects will play a role at this level - and the slicer slit height increases by a factor of 16 to 1100 pixels at which point the detector noise will dominate shot noise to high SIN levels.

The ceD is a standard thinned 2 K x 2 K device with 15 /-1m pixels ,of which presently only 1 . K is used in the Y direction due to camera field limitations.

The image slicer is located in the telescope focal plane. The function of this device is to reformat a roughly rectangular entrance aperture into a slit. A Bowen-Walraven type slicer is envisaged. Preliminary determination of its parameters calls for .5 to 10 slices of a

44

'" 1.4 x 1.6 aperture, producing a slit that is 0.28 to 0.14 arcsec wide and 8 to 16 arcsec high. The output slit of the B-W slicer is tilted (45 degrees in glass, '" 35 degrees when viewed from air). The predisperser slit and the detector sh~uld be tilted by 12 degrees in order to maintain the focus along the slicer slit.

B-W slicers are studied at UCL and at ESO in order to provide longer slicer slits, at the same time avoiding the need for tilting the detector. UCL have installed a modified slicer at the UHRF at the AAT (see F. Diego et a1. in Instrumentation in Astronomy VIII, SPIE Vol. 2198).

The predisperser has two functions: - wavelength selection - creation of an anamorphic beam with major and minor axes 40 and 20 cm to reduce the number of segments in the mosaic from 2 x 4 (40 cm circular beam) to 1 x 4. This will reduce cost significantly, also that of the collimator/camera.

All predisperser optics are fused Silica. The F /15 Nasmyth beam is collimated by an as-pheric singlet. Adaptation of the instrument to different focal ratios, for instance at different focal stations of the telescope, is easily done by modification of the slicer and predisperser first lens. In its parallel beam, two prisms create simultaneously the dispersion required for order selection and an anamorphic effect of a factor 2. The beam is now elliptical. A second aspheric singlet focuses the spectrum on the predisperser slit. Since the singlets are not achromatic, they are focused as a function of wavelength. The slit width and position is set to admit only part of one order to the main spectrograph.

The main spectrograph unit has a 3-mirror collimator/camera that is used in double pass. The focal length of this system is 166 cm so ~ with an elliptical beam of 40 x 20 cm - its working speed is F /4.16 in the dispersion. In order to be able to manufacture it, the very fast M1 mirror is a concave sphere. M2 is an on axis convex parabola and M3 an off axis concave hyperbola. A flat folding mirror (M4) is inserted to get a better location of the image plane. Presently, the field in Y is limited to '" 15 mm to be able to keep M1 spherical and to reduce the off axis distance of the system. The preliminary image quality of a predesign of the whole system (predisperser and collimator in double pass, 31.6 g/mm echelle) is shown in Figure 3. The wavelength selected is 400 nm but the image quality is largely independent of wavelength.

The concept of a cylindrical lens has been investigated (Fig. 4). It consists of a doublet (actually, two cylindrical sapphire lens rods) that are placed some 4 mm in front of the detector. This lens combination compresses the height of the spectrum a factor 5, reducing it from 345 pixels (R = 600000) to 69, which is important to reduce the influence of detector noise. The Fino in the slit direction speeds up from 8.3 to 1.67.

45

3.4 Performance

The following table shows how we arrive at our estimate of detection efficiency of 8%. Gains in efficiency can be made by going to a two-arm layout which permits using higher efficiency coatings with a smaller bandpass.

Table 2. Detection efficiency estimate (N asmyth location, at blaze peak in the visible, from top of the atmosphere to detected photoelectrons):

1 airmass 0.89 3 telescope reflections (0.853 ) 0.61 central obstruction (1.2 m) 0.98

Image slicer and couplingl 0.70 Predisperser 0.80 Main spectrograph 0.70 Echelle (at blaze) 0.65 CCD QE 0.60

Total 0.08

lThis includes optical transmission losses in the slicer and geometrical slit losses at the slicer entrance.

A first estimate of the limiting magnitudes is established as follows. The photon rate for an 17th magnitude V star is 1.66 photons/m2 /s/ A at the top of the atmosphere. This star will produce 320 detected photoelectrons per 6.7 mA wavelength bin in 2 hours; or a shot-noise limited max. possible SIN of 18. Similarly, a 12th magnitude star will give 32000 photoelectrons or S / N = 180 in 2 h. Time-resolved spectroscopy with exposure times of 72 seconds can be done to V = 7 at S / N = 180.

These values represent the theoretical maximum, regardless of the size of the grating. To which extent one can reach these values depends on the detector noise properties since these determine how much one can afford to slice before detector noise starts to dominate. In the case of the V = 17 star, the detector noise variance must be less than 320, say 160. Splitting this up equally between read noise and dark noise, the variance of each source of noise that can be tolerated is 80.

Read noise can be reduced by binning, but the cosmic ray rate is going to determine how much binning one can perform in the vertical direction. In current devices it is about 2 events/h/mm2• There is little chance that the cosmi<; ray rate can be substantially reduced by future developments. If one allows a maximum of 1 percent of pixels that are affected by cosmic ray hits, more than 8-fold binning should not be used. The spectrum is 70 pixels high which is equivalent to 9 superpixels. The constraint on read noise is thus:

46

Fig. I. Optical layout (schema.tic)

Fig. 2. Instrument footprint at Nasmyth

47

",."'- --- _75_ . ..,- .1 .. -.aD_ ·16.11- .... - IUI_ ~-

".::: ~:." 1'-::'-

US_ ::.::~.;" I~ :. .~ . 1;:" ,,' ".

w '~., . .,. , .. ,

" ···.:·1.:: .. .".,' .. ~

I~ r' .:~

~ ;" -, .. ,

.' r, I ... ~ 15 .... I

I " .. , ·US_ :::;:!i :~ I~'" I:'.' ~.~ .j .

Fig. 3. Spot diagrams showing image quality along a 60 mm field.

IU_ ... -VHRS • AMmoIphic fIclcllcnl

1.00 ...

FIg. 4. Cylindrical doublet for vertical image compression that reduces the importance of detector noise. The compression factor is S.

48

read noise < 3 e-

The maximum dark current that can be tolerated from 70 pixels in 2 h is 80 e-, which means

dark current < 0.6 e- /pix/h

Although these values have been experimentally demonstrated on individual devices, we are not aware of any chip world-wide in operation with these properties (as well as good QE, large size, etc.). The best large devices currently available have about 4 e- read noise and 2 e- /pix/h dark current, which in this example would lead to a largely dark current noise limited SIN of 12 instead of 18.

The calculation shows that the read- and dark noise properties of future CCDs are going to have a decisive impact on the limiting performance in the low S /N regime. There is room for improvement in these areas. Use of multiple readouts will allow to reduce the read noise in proportion to the square root of the number of reads. Since the VHRS detector is relatively small and will often be used in binned mode, the increase in readout time can be tolerated. Devices that ESO currently has. on order will be able to be read out in MPP mode, which has potential to reduce dark current by a factor of 10 - 20.

Table 3 shows the spectral coverage of the proposed design at a number of wavelengths. With a single-pixel resolution of 900000 and a 2 K chip, the wavelength coverage (AA) will be 1/450 of the selected wavelength, i.e. 11.1A at 5000A. With a 31.6 g/mm echelle as pro­posed, the free spectral range (FSR) is 40.6A at this wavelength so a 2 K spectrum covers only 27% of a full order. The optics permit to extend the field in the dispersion direction as is evident from Fig. 3.

Table 3. Wavelength/order coverage

A AA FSR AA/FSR (A) (A) (A) (%)

3000 6.7 14.6 46 5000 11.1 40.6 27 7000 15.6 79.5 20 9000 20 131.4 15

3.5 Critical areas

The most critical element is clearly the echelle. We propose to build the VHRS mosaic as a replicated mosaic (simultaneous replication of a mosaic of submasters on a common

49

MIRRORS COATINGS Inc=5 Deg

v _

" x , "II ; - I I

I C , 0- I I -- I

I X I I • I -II I

... '" I I I • ,

It: I -'" I II I •

I "".".I.nlJ~h (nm)

"

~."~~r------r------~-----,------~------.----588.88 1588.88

Fig. 5. Efficiency of enhanced Silver coating in the visible (top) and NIR (b9ttom)

50

substrate) as it haS been developed in collaboration by Milton Roy and ESO for the EMMI and UVES ~4 echelles. This technology has the important advantage of an instrument LSF which will be stable. The EMMI echelle mosaic has demonstrated R in excess of 700 000. The UVES mosaics are specified to reach R = 500000. Their delivery is expected in 1995 and 1996. An alternative could be a mechanical mosaic (HIRES at KECK) or a mixture of both replicated and assembled mosaics. Note that the theoretical diffraction limited resolu­tion of a 1.7 m echelle is 6400000 at 0.5 I'm, more than 10 times the required resQlution. Hence, the mosaics need not be phased.

With the very big grating that we chose for this "existence proof", and the use of the cylindrical field lens, the CCD properties are not very critical and shot-noise limited per­formance will be possible beyond a S /N of 10 - 20. Improvements in CCD noise properties beyond the present state of the art might a.llow heavier slicing, hence a sma.ller, less costly instrument.

The coatings are an area. of concern, especia.lly the reflective coating to be used in the collimator/camera (7 reflections). Silver is pretty good beyond 4500A, but one loses the UV. Aluminium is not efficient enough. Exchanging mirrors as in the CES or a two-arm' solution as in UVES will complicate the design and increase cost, even more so if a third arm for the 1 - 2 I'm range shoul~ be added. An overcoated Silver coating, enhanced for the blue-UV might be the best option (Fig 5), this coating was developed fot EMMI. Its performance, especia.lly in the UV -blue, is less than optimised coatings could deliver. The process will need further industrial development for the typica.lly 60 cm diagonal mirrors.

Large Silica prisms need to be temperature controlled to better than 0.1 °c to preserve the image quality of the transmitted wavefront and require a stable thermal environment. At this high resolution, effects like air turbulence in the instrument, vibration and mechanical flexure due to thermal or gravitation influences become a source of concern. Presently, little can be said about these effects, except that they will be less important if one manages to keep the instrument compact.

3.6 Conclusion

In this type of linear-format image slicer spectrograph where one is not constrained by space in the spatial direction or by interorder space requirements, spectrograph size and cost are determined by the quality of the detector and by the SIN at which one wants to be photon noise limited. For detector parameters that represent the current state of the art, this design is photon noise limited down to a SIN of f"J20. Further study and optimisation must take into account the likely properties of detectors and SIN requirements since these have a large impact on the optomechanical concept. It will be necessary to perform a careful tradeoff between the available options.

51

WFIS: A WIDE FIELD VISUAL MULTI SLIT IMAGING-SPECTROGRAPH FOR THE VLT

G. Vettolanil, F. Delabre2 , O. Le Fevre3, F. Hammers, G. Zamorani4

1 Istituto di Radioastronomia, Bologna, 2 ESO, Garching, 3 DAEC, Observatoire de Paris-Meudon, 4 Osservatorio Astronomico, Bologna

1. Driving Science: why a Survey Instrument to Study the Properties of Large Scale Structure and Galaxy Evolution?

After the realization, in the late seventies, of the existence of a large scale organization of galaxies in structures, the efforts of a few groups in the USA has led to a systematic mapping of the local, i.e. up to 10,000 Km 8-1 , Universe [CfA1 and CfA2 (Geller and Huchra 1989), LDSS (Da Costa et al 1991) and Arecibo surveys]. These surveys (see Giovanelli and Haynes 1991, for a review) have convincingly demonstrated some important properties of the galaxy distribution, such as:

a. Galaxies are clustered on scales less than lOh-1 Mpc in dynamical systems of different richness (from poor groups to rich clusters).

b. Large, underdense regions (voids) have been detected with maximum sizes of the order of 50h-1 Mpc.

c. The "field" galaxies around the voids form structures which are connected and have the appearance of sheets with sizes up to 100h-1 Mpc. At intermediate depth (z ~ 0.2), redshift surveys have been confined to smaller areas

(few tens of square degree) but they have already confirmed the validity of the picture of the large scale distribution of galaxies derived from shallower surveys. Furthermore they have shown the beginning of interplay between large scale structure and galaxy evolution (ESP at ESO (Vettolani et al 1994), Las Campanas (Oemler et al1993) and Century (Geller 1994, private communication) from USA groups, surveys which are all in progress). Much fainter redshift surveys at the Anglo-Australian, Kitt Peak and CFHT telescopes have been limited to some hundreds of galaxies up to bj ~ 22 - 23.5 in small pencil beams and probing galaxy evolution to z = 0.3 - 0.6 (see Koo and Kron 1992). Despite the efforts of so many groups and the large investment of telescope time, many possible issues about the properties of large scale structure remain completely open and largely unsolved, for example the determination of the typical dimensions of structures and voids, their relative density contrast, their topology and statistical properties.

In order to det~rmine these properties, one has to study statistically representative samples of large numbers of galaxies covering large areas and volumes.

For surveys of galaxies with bj ~ 20 - 21 a fundamental role is played by the 4 meters class telescopes, if properly equipped with efficient multifiber spectrographs with a large field of view and hundreds of fibers. The best example of this instrumentation is the 2dF Spectrograph at the Anglo-Australian Telescope, with 400 fibers over a two degrees field of view (Taylor 1994).

53

However, for fainter galaxies we have reached the point at which the 4 m telescopes are totally limited by the necessity of collecting these huge amounts of data, requiring exposures of the order of many hours, to seriously investigate the Universe to the depth required to sample significant volumes and look-back times. Surveys of fainter galaxies are the domain of the very large telescopes as the gain they provide in the number of photons hitting the detectors is just enough to make surveys of faint galaxies feasible in a reasonable amount of time. It is, therefore, easily predictable that the large telescopes will have a strong impact on the study of the formation and growth of the diverse structures, as well as constraining galaxy evolution.

In the following we briefly describe the possible role of large telescopes in two nowadays very challenging problems: the nature of faint blue galaxies responsible for the galaxy counts excess and the study of th~ evolution of the galaxy clustering.

1.1 Counts of Faint Blue Galaxies Counts of galaxies in the blue band as a function of magnitude show a large excess ( ~

a factor 5 at bj = 24) over what is predicted extrapolating the number of galaxies per unit volume as measured locally (through the galaxy luminosity function) to distant volumes and large lookback times. These predictions are hampered by our poor knowledge of the galaxy luminosity function for different morphological types, the appropriate cosmological K-corrections at large redshifts and the amount of the evolutionary corrections which take into account the evolution of the stellar population in galaxies. These simple passively evolving models totally fail to predict the counts at faint magnitudes. Conversely, K band counts, Le. counts of galaxies selected on the basis of the luminosity in the older stellar population, show almost no excess over prediction.

One could explain the blue counts if these blue faint galaxies were distant galaxies with young stellar populations (hence the blue color) or a nearby very low luminosity population. However, the fact that their redshift distribution is similar to the predicted distribution from the same model which fails to predict the counts and that does not show a tail at high or very low redshifts, req~ires a more complex scenario.

Many solutions have been put forward, as for example: a.· non conservation of the galaxy numbers due to mergers; b. a population of dwarf galaxies disappearing at short lookback times; c. non zero cosmological constant;

What we do know is that some physical process involving a substantial fraction of the galaxy population was acting at redshifts around 0.5 and beyond. Identifying these process (or processes) is really challenging. Observing deeper samples (hence larger redshifts and lookback times) would permit to test some of the hypotheses which have been put forward which predict f.L different shapes of the redshift distributions in faint samples.

1.2 Correlation Function The modeling of the formation and growth of the diverse structures observed in the

Universe can be strongly constrained by the measure of the galaxy clustering as quantified by the two-point spatial correlation function. This function e(r) measures the joint prob­ability of finding galaxies in the volume elements 8Vl and 8V2 separated by the distance r.

54

At small scales e(r) evolves in time under the action of gravity while the signal on large scales strongly reflects the form of the initial spectrum of the perturbations in the Universe.

As its measure requires the knowledge of the galaxy distribution in three dimensions, we have very few information on the evolution of its shape and amplitude with cosmic time, due to the essentially mono dimensional nature of present deep surveys. In principle, the spatial correlation function can be indirectly determined from the angular correlation function. The procedure is quite similar to the counts modeling, in the sense that it requires to integrate along the line of sight to all galaxies in the sample contained within the appropriate volume of space. Again, this procedure requires the knowledge of the luminosity function, its dependence on the morphological types , the amount of luminosity evolution that galaxies have undergone with cosmic time etc. Given the difficulties in the counts modeling, it is clear that any determination of e(r) and of its evolution from the angular function suffers from large uncertainties.

The determination of the spatial clustering properties at large look-back times would also enormously help to solve the puzzle of the faint blue galaxies excess. The blue galaxies are vigorous star forming, therefore, once separated from the more normal population through for instance the width of the OIl line, their correlation function should clearly show if they are a not-clustered new population of which there is no trace in the near samples, or if they have similar clustering properties to the more normal galaxies, which would favour a merger hypothesis.

l.From these brief considerations it immediately follows that the main requirements of an instrument dedicated to studies of observational cosmology are those of combining the large collecting area of an 8 meter telescope with the flexibility of a multislit spectrograph to be able to gather, in a reasonable amount of observing time, large statistical samples of objects representative of a large fraction of the Universe volume.

2. Instrument Specification and Performance The requirement of a large field of view, with an as large as possible number of slits,

calls for the Nasmyth Focus of one of the VLT telescopes, with its unvignetted field of view of 20 arcmin in diameter.

Imaging is an obvious requirement to select objects to be observed spectroscopically, but will have also, by itself, important scientific outcomes.

As will be made clear from the examples in the next sections, there is no clear reason to require resolutions R larger than 2000 if the instrument is going to work as a survey instrument of cosmological interest. The constraints on resolution are posed essentially by the required accuracy in velocity measurement, which is of the order of 50 kmj s for the most accurate measurements of radial velocities of galaxies in clusters.

Spectral resolutions between 250 and 2000 seem quite well adapted to most of the programs, with the lowest value well tailored for quasar surveys or very faint galaxies and the highest value for accurate redshift measurements in clusters of galaxies.

The number of available slits is a crucial point for the evaluation of the instrument capabilities and for its comparison with similar existing or planned instruments.

55

TABLE 1

(1) (2) (3) (4)

mag N Galaxies B-R N stars

(bj) per square arcmm

19.5 0.014

20.0 0.036 1.53 0.19

20.5 0.08

21.0 0.16 1.40 0.32

21.5 0.34 22.0 0.58 1.43 0.50

22.5 1.05

23.0 1.89 1.23 0.78

23.5 3.48 24.0 6.12 1.19 0.90

24.5 10.57

Table 1 gives the relevant numbers of objects (galaxies or stars) in magnitude bins (cumu­lative distribution) adapted from Jones et al (1991). Basically, one should try to match the number of slits to the density of objects at some interesting magnitude. If we assume that bJ '" 23 is such a magnitude (see f.i. the example on large scale structure in the next section) the required slit density is of the order of 1.8 per square arcmin. The high density of objects per square arcmin calls for a mask system.

A fundamental constraint on slits, however, is given by their length. It is absolutely necessary to sample as well as possible the sky outside the object in order to get the best possible sky subtraction. Sky subtraction errors are in fact suppressive, in the sense that, for sky limited observations, the maximum attainable signal to noise of the spectrum is essentially defined by the error in the sky subtraction and, beyond a certain limit, it does not increase any longer with the exposure time. This forbids the use of too small slits. An acceptable length for the slits length is between 10 and 15 arcsec.

3. WFIS In the .following we describe an imager-spectrograph which fulfills most .of the requirements dictated by the science. The requirement of the largest possible field of view, calls, as already mentioned, for the Nasmyth focus where, however, a single .channel instrument cannot be accommodated with full use of the field. Therefore, we have developed a concept

.56

based on 4 identical channels, slightly separated, each covering a field of 7 x 7 arcminutes. Being each channel far from the optical axis , the N asmyth field curvature creates an asymmetry which limits considerably the image quality. Moreover for the spectroscopic mode it is necessary to have the multi-mask aperture on a flat surface perpendicular to the optical axis of each channel. These requirements are matched with a corrector lens between the adaptor flange and the N asmyth focus.

The number of optical elements in both the collimator and the camera has been kept to a minimum. In the collimator 3 lenses in two groups form an image of the pupil 300 mm in front of the second lens with a parallel beam of 135 mm diameter. The camera has a focal length of 250 mm and a FIN of 1.87. The field which can be covered is 70 mm diameter which is perfectly adapted for a rectangular detector of 4096 x 2048, with 15 microns pixels.

Detectors of this size could be either a monolithic device (as envisaged for Keck2 in­struments) or a mosaic of two 2048 x 2048 devices as presently foreseen for other VLT instruments. Only the central half of the detector area is used in imaging, while the entire area is reserved for spectroscopy.

Table 2 summarizes the optical characteristics of the instrument. More details on the optical design can be found in Delabre et al (1994).

The splitting of the field into four separate channels makes simpler the realization of the slit masks. Due to the partial correction of the field curvature provided by the focal plane adaptation lens, a flat plate of 252 x 252mm2 can be inserted after preparation of a slitlets pattern according to the distribution of targets in the field. This will be derived from direct images of the same field obtained from direct imaging from this same instrument. Assuming slitlets of 10-15 arcsec length, a total number of 100-120 spectra, or more depending on field geometry, overlap etc, can be obtained in a single exposure in the four CCDs, corresponding to a slit density of 0.4-0.6 per square arcminute. Therefore, complete sampling of all galaxies brighter than bj = 23 would be obtained with 3 to 4 exposures. With the parameters of the spectrograph described above, and an efficiency ~ 40%, exposure times of the order of 1 hour will provide SIN = 10 in the continuum. Such a signal to noise is more than adequate to obtain reliable redshift determinations.

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TABLE 2

(1)

FOV (arcmin)

FOV (mm)

Pupil Diameter

Final f/ratio

Final Scale

F camera

detector

wavelength range

Spectral coverage

Spectral coverage

4 A few possible scientific highlights

(2)

7.0 x 7.0 (4 times)

256 x 256 (4 times)

135mm

1.87

0.205 arcsec per pixel

252mm

Four 4096 x 2048, 15 microns pixel CCD

380 1000 nm

800nm with slit in the center

300 l/mm grating (R=780)

400nm with slit in the center

600 l/mm grating (R=1560)

We present here some examples of scientific goals which can be reached by using such an instrument. This is certainly not an exhaustive list of objectives: a large number of studies other than cosmological ones can also be accomplished, including stellar cluster dynamics, extragalactic globular clusters etc.

a) Large Scale Structure

We start describing a program which, for its size, is hard daring to propose, but we feel interesting to consider for two main reasons: first, because it directly compares to the most ambitious program ever conceived in cosmological research (namely the SLOAN survey), second, because it gives directly the flavor of how easily we will be able to accomplish programs with scientific objectives which are not thinkable without such an instrument on the VLT.

At the beginning of next century the SLOAN survey will be accomplished with a fully dedicated telescope. After a 5 years photometric survey, all galaxies (plus stars) brighter than the 19th magnitude in the northern sky will be observed spectroscopically, in a 5 years, or more, planned period. This survey will have a typical depth (at the maximum of the selection function) of '" 240h-1 Mpc, corresponding to an explored volume of ~ 8.3 x 106 Mpc3 , over 1.8 steradians. A redshift survey at bj = 23 has a depth (at the maximum of selection function) of ~ 900h-1 Mpc, depending on qo. Therefore it is easy to

58

show that '" 1800 fields (15 x 15arcmin) would give the same explored volume. These can be arranged in a way as to maximize the information on large scale structure transverse to the line of sight. Qbserving these fields with a galaxy sampling of one over three, this would call for some 200 observing nights: a large but not unfeasible program!

A program like this would allow to determine at z ~ 0.3 the properties of the large scale structure with the same statistical power as the SLOAN survey at z ~ 0.1. A comparison of the results of these two surveys would immediately give us interesting and possibly exciting results on the evolution with cosmic time of the large scale structure, on which nothing is known (except from N body models). Furthermore, a survey like this would have the enormous advantage, over the SLOAN survey, of tackling directly the problem of galaxy evolution (on which SLOAN tells nothing) over a lookback time 30% of the age of the Universe.

Now we turn to the problem of galaxy evolution with a simple realistic example.

b) Galaxy Evolution As mentioned, faint galaxies, while being overabundant over non-evolutionary models prediction, show a redshift distribution which is displaced towards low redshifts, consistent with such models. Vice versa, evolutionary models, which fit the counts, predict a redshift distribution which is displaced toward higher redshifts than expected. Large spectroscopic samples of faint galaxies are needed to constrain models which can fit at the same time both the counts and the redshift distribution. There are '" 1000 galaxies with bj ~ 24 in the field of view of WFIS and 10 - 15% of these are expected to be galaxies having 1 ~ z ~ 1.5 for which evolution could become dramatic, according to present modeling. All these galaxies can be observed with 10-12 exposures of 3 hours each, corresponding to 4-5 observing nights. A program like this over a few fields would provide a fundamental statistical sample to clarify the processes of the evolution of stellar populations in distant galaxies, in a reasonable amount of observing time. c) Galaxy clusters Clusters of galaxies are the more massive bound structures which ever formed in the Uni­verse, possibly through the hierarchical growth of subunits. A comprehensive dynamical study of a galaxy cluster, sampling its luminosity function well below M., requires a few hundred radial velocities of galaxies with accuracies of the order of 50km 8-1 . This has been accomplished, or is in progress, for several nearby clusters with fiber spectrographs at 4 meters telescopes. At intermediate redshifts (0.3 - 0.5), extensive studies of a few galaxy clusters have been mostly motivated by the study of the ~utcher-Oemler effect, i.e. an increasing fraction of blue, starburst and/or poststarburst galaxies, with respect to nearby clusters. More recently, comprehensive dynamical studies of a few clusters have been led by the discovery of arcs and arclets due to the presence of a rich cluster in the line of sight of a background faint galaxy.

These studies are at the borderline of 4 meters class telescopes, requiring very long inte­grations per object (or per field in the case of multi slit spectrographs), and therefore the number of clusters with useful dynamical information is quite limited.

At z ~ 0.3 - 0.4 a 15 arcmin field corresponds to ~ 4 - 5 h-1 Mpc, and is extremely well suited to study the whole velocity field of a rich cluster and its surroundings. Assuming

59

that 80 galaxies could be observed in 2-3 hours, 1 night of observing time would be enough to get details of the velocity field of a galaxy cluster up to z ~ 0.4.

The faint end of the cluster luminosity function can be then investigated, and the new observations would provide accurate estimations of the dynamical mass and its distribution over each whole cluster. These masses can be compared to the masses derived from X-ray gas models as well as to the mass estimation from arcs and arclets. Spectroscopy of the latter could be done during the same time and greatly helps to estimate the cluster dark matter and its distribution. This data would provide unprecedented constraints on the largest clumps of dark matter in the Universe and hence on mass related cosmological parameters. Furthermore these observations would allow a detailed study of the Butcher­Oemler effect in relation to the cluster dynamical properties as this effect probably arises from infall of field galaxies into the cluster.

d) Quasars A well known, unsolved problem in quasar surveys is the quantitative determination of the incompleteness due to the various selection techniques. A not well understood incomplete­ness can introduce significant biases in the computation of the luminosity function and evolution for these objects. The most straightforward and unbiased way of solving this problem is that of observing all objects (both stellar and non stellar) at;a given magnitude limit. At the 23rd magnitude Table 1 shows that there is an overall o~ject density of 2.7 per square arcmin (1.9 galaxies and 0.8 stars). The number of exposures for observing all objects is given by the ratio of the object density to the slit density (0.4 or 0.6 per square arcmin, in the case of 15 or 10 arcsec slits respectively), hence from 5/6 to 7/8 exposures per field.

At bj = 23, the expected number of QSO's is about 16 per field (4 per ·each 7 x 7 arcmin CCD). Therefore, if the final goal is to obtain a completely unbiased :sample of 100-200 quasars, this project is easily feasible, whilst it is impossible for 4 meters telescopes, or hardly feasible with a much lower slit density or with a similar slit density, but smaller field of view (e.g. FORS)

Besides establishing unbiased QSO's counts, such a program will provide the counts, ac­tually fairly unknown, of low luminosity AGN s from the galaxy observations (2-3% of them). Further products are the correlation of quasars (at least low z ones) with the galaxy environment.

4. Comparison with other Instruments

The best survey instrument at a 4 meter telescope is nowadays the MOS/SIS spectrograph at the CFHT (Ie Fevre et al1993) which is intended primarily for imagery and spectroscopy of many tens of objects over a field of 10 x lOarcmin. With this instrument Crampton et al (1994) have obtained a sample of ~ 700 galaxies brighter than 1 = 22, i.e. bj ~ 24 with exposure times of 8 hours. This project can be considered as the extreme limit of the possibility of a 4 meter telescope and has represented 20 nights of observing time. Now with VLT equipped with a spectrograph comparable to WFIS this same program could be performed in 1-2 nights (depending on the number of slits). Note that due to the

60

smaller number of slits FORS is not competitive with MOS/SIS for studies requiring large statistical samples.

The VLT instrumentation plan foresees two instruments which can be directly compared to what here proposed for the scientific programs we are here considering, namely FORS and FUEGOS. The comparison with FORS is given by the ratio of the total fields of view, as the slit density as well as other characteristics are similar, which is roughly a factor of 4 (in favour of WFIS). The comparison with FUEGOS is complicated by the capability, presently difficult to estimate, of sky subtraction. Assuming, however, that FUEGOS can reach magnitudes as faint as 23bj with exposure times of the order of 2-3 hours, the comparison is defined by the number of fibers which is totally inadequate for survey work reaching a density of only 0.1 per square arcmin. The multiplex advantage is therefore 4-6 times the ratio of the exposure times (due to the low transmission of fibers).

The comparison with other instrumentation at 8 meter telescopes is less straightforward as instrumentation plans are not completely defined. Keck II will have a low dispersion spectrograph (DEIMOS, Oke et al1994) with an areal coverage of 80x 2 arcminutes (two lunes of '" 19 x 5.5 arcminutes separated by 10 arcminutes) and 150 slitlets. It will have therefore a slit density ~ 0.8, two times the slit density of WFIS, over a field of view ~ 80% smaller and with rather peculiar geometry.

6. Conclusions

The net result of these comparison is that an instrument as the one here proposed is certainly competitive with similar instruments on 8 meters class telescopes, other than VLT, and in some instances, especially if coupled with a low resolution near infrared spectrograph, much better than similar instruments.

We would like, however, to stress another point which is worth mentioning. Whilst studies as the ones above described have been the major driving force for building 8 meter class telescopes, in the present instrumentation plan of VLT there is almost no room for this kind of work. Not to build an instrument like the one here proposed, or a more clever one but with the same goals, means to throw out European research from the field of observational cosmology.

References Crampton,D., Hammer,F., Le Fevre,O., Lilly 1994, in preparation Da Costa et al, 1991 Ap. J. Suppl, 75, 935 Delabre,B., D'Odorico,S., Vettolani,G., 1994 SPIE Co~ference " Instrumentation for the

21st Century", Kona, Hawaii, in press Geller, M. and Huchra J. 1989 Science 246, 897 Giovanelli, R., Haynes,M.P.: 1991, Ann. Rew. A. A., 29, 499. Jones et al1991 MNRAS 249, 481 Koo, D. and Kron R. 1992 Annual Review A. A. 30, 613 Le Fevre, 0., Crampton,D., Felenbok,P., Monnet,G. 1993 Astron. and Astroph. in press Oemler et al 1993 in "Observational Cosmology", Chincarini et al edts, ASP Conference

Series 51, page 81

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Oke J.B. et al 1994 SPIE Conference" Instrumentation for the 21st Century", Kona, Hawaii, in press

Taylor K. 1994: AAO Newsletter N. 69 Vettolani et al1994 in IAU Symposium 161: "Astronomy from Wide Field Imaging", H.T.

Mac Gillivray Edt., Reidel Dordrecht, p. 687

62

THE WID~FIELD DIRECT VISUAL CAMERA (WFDVC)

E.J. Wampler European Southern Observatory

1 Introduction and overview

It is a well-known theory of optics that the ability of an optical system to transfer in­formation is proportional to the product of the aperture of the optical system and its field of view. The aperture enters the equation both because the larger aperture increases the photon collection rate and because the larger aperture give higher optical resolution than can be obtained with smaller apertures. In fact, the main arguments for construct­ing large telescopes rather than small ones has been that large telescopes collect more photons in a given amount of time and, particularly in the infrared, large telescopes give higher angular resolution than small telescopes. While for many studies a large field is not required, there are forefront astronomical studies that require the combination of a large field with a high signal-to-noise (SIN) ratio and high angular resolution on the plane of the sky.

The instrument described here is designed for deep imaging over a 9 arcminute field and is capable of providing output images of the sky with a resolution of 0.1 arcsec Full­Width-Half-Maximum (FWHM). The instrumental parameters are chosen to fully exploit the imaging capabilities of an 8-meter VLT unit telescope. The detector is a mosaic of 16 large CCDs that sample the telescope image with a spatial resolution of 1/20 arcsec (24 p. pixels). The field of view of the camera will then be about 9 arcmin. Error signals for guiding the telescope will be taken from the imaging array itself.

The camera will be designed to fully use the modern image deconvolution techniques being developed at the ST-ECF and ESO. It is for this reason that a small pixel size was chosen; it is desirable to sample the deconvolved image with two pixels per resolution element and we want to reach a final spatial resoltition of 0.1 arcsec. The camera will complement the imaging capabilities of HST. The optical resolution is not as good as HST, but it will give a larger field than the Space Telescope and the 8-meter VLT unit telescope has 10 times the collecting area of the HST. The great collecting area of the VLT results in better statistical smoothness for the VLT images when compared to HST images. The high SIN ratio of the VLT images can be used to improve the seeing-imposed resolution limit of the ground-based images (Lucy, 1992).

2 Science programs for the wide-field imager

2.1 Deep imaging of high latitude fields

A number of programs of a survey nature can be run in parallel, with one group of researchers being interested in one aspect of the results, while another group is interested in a scientifically different aspect of the same data. To some extent this is the 8-meter telescope equivalent of the Palomar Sky Survey; different researchers used the plates to investigate different problems. Listed here are only a few of the many possible programs; these few examples are intended only to illustrate what is possible.

63

....

Among the programs that are obvious candidates for sharing high galactic field deep imaging data are programs that investigate faint halo stars and programs that study faint galaxies. Such programs will tax the VLT ·to the limits of its capabilities as such programs require imaging objects that are near the seeing and magnitude limits of an 8-meter telescope. The separation of the very faint galaxy images from those of very faint stars will require combining color information with very high angular resolution.

2.1.1 Gravitational lensing by galaxy clusters

The scientific case for studying the lensing of background galaxies by galaxy clusters is well presented in a review article by Fort and Mellier (1994). In their paper they argue that a distant galaxy cluster can be thought of as the first element of a "natural giant telescope" (emphases in the original). Furthermore, "for a given potential, the shear pattern and the arclets distribution around cluster lenses have an angular scale which depends on the redshift distribution of the background sources". Thus, the wide-field imaging of galaxy clusters can reveal arcs and arclets that not only give us information about the cluster potential, but also information about the distribution in space of the lensed objects.

Away from the center of a massive cluster, arclets and gravitational alignment of background field galaxies provide evidence for shear in the gravitational potential around massive systems. There are occasions that the shear field is detected even when there is no obvious luminous cluster present. Evidently, there are locations in the Universe where mass concentrations are not accompanied by luminous matter. One example is given in Figure 18 of Fort and Mellier (1994). To detect such structures it is important to have available very large fields. The example by Fort and Mellier (1994) has a field diameter of 13 arcmin.

A deep gravitational potential well is needed in order for a cluster to be a strong gravitationallense. Bartelmann (1993) has shown that in the case of spherical collapse in a flat universe there should be a strong evolution in the formation of dense clusters and this should lead to a corresponding strong decrease in the number of arcs with redshift. In any case, the study of arcs as a fu~ction of red shift will allow us to study both the evolution of galaxy formation and the evolution of cluster collapse.

2.1.2 Distant galaxies

Observations of faint galaxies provide information about both the epoch of galaxy for­mation and the geometry of the Universe (Yoshii and Peterson, 1991). The observa­tional data, the color distribution, the redshift distribution and the surface distribution of galaxies on the sky as a function of magnitude constrain the range of permissible model parameters and hence, discriminate between possible models of the early Universe. The standard model of galaxy formation is based on the assumption that galaxies formed at a common epoch and then processed their gas into stars at a rate that is proportional to ~he amount of remaining gas. The different types of galaxies that we see now are the results of the different rates at which gas is converted into stars. For elliptical galaxies the star forming rate was initially very high. The metal lines seen in high redshift absorption lines in quasar spectra also suggest that there must have been considerable star forming activity in the early universe (Bruzual & Kron, 1980; Guiderdoni & Rocca-Volmerange, 1987; Matteucci & Padovani, 1993).

64

In the context of these models, the surface density of on the sky of the faintest galaxies (V > 26) is an indication of the geometry of the Universe, the color distribution of the intermediate brightness galaxies (24< V <26) constrains the epoch of formation, and the redshift distribution of the brightest galaxies (22< V <24) provides a means of verifying the models. Work by Broadhurst et al. (1988), Colless et al. (1990) and Cowie et al. (1991) have shown that most galaxies in the magnitude range (22< V <24) are nearby (z < 0.8) dwarf galaxies. The standard model would predict an excess of blue galaxies, young ellipticals rapidly converting gas into stars, at B /'OJ 24 if z /ormGtion < 3. These are not seen and the redshifts for the fainter galaxies (V > 25) are not yet known. They might represent a population of dwarf galaxies at redshift z /'OJ 1 that are no longer present, or they may be normal galaxies at very high redshift, z > > 3.

For the galaxies that are fainter than V /'OJ 25, redshift measurements have not yet been obtained. The standard models give predictions for the surface density of galaxies with V > 25 and the observed number count is much higher than the predicted surface density if n ~ 1. The B, V, R, and I counts are consistent with an n ~ 0.1 because for open universes the volume at high redshift is larger than for high n universes. However, in the infrared the K' counts are less than those for the B, V, R, I-bands, and this problem remains unresolved.

Unpublished NTT observations of faint galaxies in Tyson's SGP field show that in conditions of good (~ 0.6 arcsec) seeing a substantial fraction of the faint blue images are smaller than about 0.7 arcsec in diameter. The nature of these compact, faint blue objects is unclear. They may be dwarf galaxies, they may be compact star forming regions in larger galaxies, or they may even be a hitherto unknown population of faint blue stars in our own Galaxy. Without spectra or better angular resolution, it is not possible to distinguish between these possibilities at the present time. Repeated exposures of these fields might show that supernova explosion occur among the population. If so, the photometry of these supernova would provide an estimate for the distance to the objects. In any event, with an image quality of 0.6 arcsec we do .not resolved the faintest objects in Tyson's SGP field. Improving this resolution by a factor of 6 can either lead to an understanding of these objects, or to even deeper perplexity in the interpretation of these deep fields ..

2.1.3 Faint halo stars

Courtesy: Chris Tinney (ctinney@eso.org) & Tim Bedding (tbedding@eso.org)

The one place in the Universe where we really know there is dark matter is in the haloes of galaxies. The hunt for the nature of this dark matter (and a few incontrovertible examples thereof) has been one of the major efforts in recent Galactic astronomy. The most obvious way in which to make up this unseen matter is to suppose a steeply increasing mass function (MF) for low mass objects in galactic haloes. Any objects formed with masses below the minimum mass for stable H-burning (around 0.IM0 for the low-metallicity stars in the halo) will have cooled long ago into undetectable obscurity, and could thus make up the halo dark matter. Because such objects would be incredibly faint (MBo' ~ 20), direct detection of them is impossible - even with the VLT. However, powerful constraints on the possible existence of such objects could be placed by observing the halo MF for masses just above the H-burning minimum mass. In particular, we would like to know

65

whether the halo MF 'allows' extrapolation to sub-stellar masses, in sucp a way as to produce the large mass densities in cold brown dwarfs required for them to constitute the halo missing mass. Or whether, as for disk stars in the solar neighborhood, the halo MF turns over before the bottom of the main sequence (Tinney 1993).

To date, direct measurements have not allowed us to choose between these alterna­tives. Most current studies have concentrated on halo stars in the solar neighborhood, where their density is '" 1000 times lower than disk stars of roughly the same luminosity (Hartwick et al. 1984 and references therein). Halo candidates are usually identified as a result of their large proper motion, with their halo status being determined either by kine­matics or low metallicity. Unfortunately, the solar neighborhood is a complicated place and the assignment of a given star to a given population is not straightforward. Ideally, stars would either belong to the ~isk (roughly solar metallicityand low velocity) or to the halo (low metallicity and high velocity). In practice, the situation is more complicated­not all high velocity stars have low metallicity and not all high metallicity stars have low velocity. The number of separate 'populations' identified in the solar neighborhood seems to continually increase. We currently have the choice of placing a given star in the old thin disk (age ~ 2Gyr, scale height'" 250pc)j the intermediate thin disk (age :6 2Gyr); the young thin disk (age :6 O.lGyr); the old thick disk (scale height'" 1 kpc); and the halo (also known as the spheroid).

It is clear that, to study the halo population in detail, we need to select halo samples uncontaminated by other star populations and unaffected by our preconceived notions of what is happening in the halo. The obvious way to do this is to study them in situ. That is, to select and study main sequence halo stars at large distances ( ~ 5 kpc) from the Galactic plane.

Fortunately, low-metallicity stars have higher masses and luminosities at the bottom of the main sequence than do stars of high-metallicity. Models predict that for metallicities Z '" 10-3 to 0 times solar, the minimumH-burning mass is between 0.090 and 0.095M0. The corresponding luminosities and effective temperatures are MBol ~ 12.0 and T eJ J ~ 3500K (D' Antona 1987; Burrows et al. 1993). When searching for faint stars, it is important to observe at wavelengths close to the peak of the flux distribution. We would therefore use the R and I passbands, and would require the search to be complete to Mr = 11.6 and MR = 13 (Greenstein 1989; Ake & Greenstein 1980; Dawson & de Robertis 1988; Hartwick et al. 1984).

The required distance limits are set by the need to ensure our halo sample is not contaminated by either the thin or thick disks. The halo sub-dwarfs in which we are interested have R - I colors in the range 0.8-1.4. Stars in the more metal-rich disk populations that have colors lying in this range will be early M-dwarfs (M1-M3). Provided we restrict our search to regions more than 5 kpc above the Galactic plane, the density of contaminating thick-disk M-dwarfs will be about 20 times lower than that of the halo stars we seek (Reid & Majewski 1993). For thin-disk contaminants the density will be lower still, by another factor of '" 25 (Tinney 1993). We therefore propose to conduct our search in the distance range 5-8 kpc above the solar neighborhood, corresponding to a distance modulus of 13.5-14.5. This sets magnitude limits of I = 26.1 and R = 27.5. Obtaining 5 % photometry in both the R and I bands would allow us to estimate MBal with a precision of 0.2-0.3 magnitudes for each star.

Current estimates of the halo luminosity function (LF) are poor, to say the least, but they do give some idea of the volume we need to sample. We estimate it will be necessary

66

to sample a volume of '" 2 x 106pC3 to detect 10 objects per 0.5 Mso' bin at the bottom of the halo main sequence. For a distance range of 5-8 kpc, we therefore require a survey covering 210arcmin2 • (An examination of the results of Richer & Fahlman (1992), who carried out a field halo survey over only 44arcmin2 to I ~ 24.5, indicate these estimates are reasonable.)

At these magnitudes, galaxies will outnumber stars by a factor of over a hundred. In the range 1= 25-26, the density of galaxies should be ~ 200/arcmin2, while that of stars will be '" 2/ arcmin 2• Good quality star-galaxy separation will therefore be essential. This is an area in which the good seeing expected at the VLT together with image deconvolution will be very important. If point source images can be deconvolved to a final diameter of 0.1-0.2 arcsec it should be possible to separate stars from galaxies with high accuracy. Those galaxies that have very small diameters are also very blue (Hn-galaxies). The stars of interest are very red (Tyson & Seitzer 1988), and it is anticipated that the contamination of falsely classified galaxies can be reduced to < 0.1 % of the total number of galaxies, or < 10 % of the total number of stars.

3 The instrument

3.1 Mechanical Design Considerations

The optical layout of an 8-meter unit telescope is shown in Fig. 1. At the present time, no final decision has been taken as to which focus to use. The N asmyth focus is larger and has better image quality than does the Cassegrain focus, but the Cassegrain focus does not suffer from optical losses at the N asmyth folding flat and the N asmyth focus is more difficult to bafHe adequately than the Cassegrain focus.

Fig. 2 shows a straw-man mechanical layout of the CCD imager for the VLT. In this concept it is supposed that 16 high quality CCD detectors, similar to the Tektronix 2 K2 CCDs are available. Other formats are possible if other CCDs become available -large, "buttable" CCDs would reduce the amount of dead space between the sensitive areas of the CCDs - but existing Tektronix CCDs are suitable for a "straw-man" design.

This design approximately covers a field 9 arcmin in diameter. Over this field the RMS radius of the geometrical image at the Cassegrain focus is better than 0.12 arcsec and the vignetting by the primary mirror is :51.1 % (See VLT-TRE-ESO-I0000-0526). Thus, the nominal image quality of the VLT remains good to the edge of the field and images diameters will usually be dominated by atmospheric seeing. However, the focal plane of the Cassegrain focus is curved with a radius of 1981.4mm. (See a discussion on the field imaging in large telescopes by Wynne (1994).) Eithera field flattener will be needed or the individual CCDs must be mounted on a curved surface. Without correction, the image blur caused by the curved telescope focal plane over the width of a single CCD would be about 1/2 pixel, or comparable to the residual aberrations of the telescope. Probably mounting the CCDs on a plate curved to match the field curvature is the preferred solution.

The telescope should have additional sky bafHes beyond those currently proposed; see Fig. 3. If a field stop a.bout 550 mm in diameter is located just behind the third mirror in the mirror tower an additional baflle 2.1 meters in diameter and extending 2 ~ meters in front of the M2 unit would block direct sky light from entering the focal plane. This represents a substantial modification to the telescope design and warrants careful study.

67

l rRED 1116

~PUPll Secondary

I mirror M2

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: roo~~OO'O ~ FOV 1043.8

"" i r Tertiary --.l 1242X86~~. _ ~~r~r_M~ _. _. _.

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Optical design, Nasmyth focus.

12426.95

URED1116

: ~.condary mirror M2

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RED 8185.~---i

Optical design, Cassegrain focus.

Figure 1: The optical elements giving the VLT Nasmyth and Cassegrain focal planes.

co CD

;..., .. -----29.5 em = 9.29 areminutes----...... ~

Figure 2 A possible CCO layout for a wide field imager. The image scale is appropriate fol' til<' Cassegrain focus.

69

~ " Telescope optical axis M2

200-100

0.00-' .~.-.-.~.-.-.-.-.~.-.-.-.-. '-'-'--100 -200-

I' Hole in M 1 I I

M1

Ray from M 1 edge

100 200

300

400

500

-600

-700 -800

900

1000

I I -1300 2m1m

-1400

-1500

3 3 --h ) 0 3 0 ~ ct () ~

~ x ~

Figure 3: The required baffle positioning for two Cassegrain field sizes, 300mm and 200mm. Note that the figure scales are very different along the optical axis and per­pendicular to it.

70

The total obscuration of the secondary unit with its light baffle will be slightly less than 7 % of the 8-meter pupil.

An atmospheric dispersion correction should be included in the design of the imager. The design proposed for FORS that occupies the space in the third mirror tower will work for the direct imager as well.

3.2 The Pixel Size

A wide field camera will cover a certain physical area of the focal plane independently of the pixel size. The angular extent of this area can only be changed by introducing additional optics and the field size in the case of the VLT is severely limited by baffling considerations. For modem CCDs illuminated with wide-band radiation, the detector read-out noise will be insignificant compared to the photon noise from the sky after integration times of a few tens of seconds. As a result, the physical size of the pixel doesn't affect the amount of information collected but just the volume of it and hence the speed at which it can be read out and manipulated. Typical large format CCDs such as the ones proposed in this document have pixels of 24 pm and not larger, corresponding to about 0.05 arcsecs at· the VLT Cassegrain focus .. This is significantly finer sampling than the Nyquist sampling even in the best of seeing. However, fine pixels have some advantages:

• First, and most importantly, very fine sampling allows the structure of the PSF to be determined on the scale of a deconvolved image. In the case of the VLT Cassegrain fo­cus the PSF will have some optical aberration (both coma and astigmatism because the Ritchey-Chretien criterion is not satisfied) as well as less predictable distortion due to guiding errors and small phase errors of the primary mirror surface. This instrument is intended to be routinely used with image restoration based algorithms for which this knowledge is vital.

• Secondly, fine pixels make cosmic ray detection and removal easier, particularly when only a single frame is available, because they occur on spatial scales which are too small to have been imaged through the telescope and can be detected by spatial filtering methods.

• Thirdly, fine pixels may efficiently be made larger by noise-free on-chip binning which is a well developed and effective technique. For certain types of observation this may be the normal operating mode.

3.3 Exposure times and limiting magnitudes

At the Cassegrain Focus the VLT focal ratio is 13.41. That of the NTT is 11. Thus, the VLT optical beam feeding the proposed camera, is about 50 % slower than the NTT beam feeding SUSI. Nearly all of this speed loss is recovered by putting the wide-field camera at the Cassegrain focus and eliminating the two flat folding mirrors that feed SUSI. The remaining small loss caused by the slow VLT beam can probably be recovered by improvements in CCD quantum efficiency. Therefore, as a baseline, we take the actual exposure times that have been used for SUSI.

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For broad-band filters( V, R, I and Gunn z) and dark moon, 2500 detected photoelec­trons are recorded in each 0.2 x 0.2 arcsec2 pixel in '" 15 minutes. During bright moon the sky is about 25 times brighter than during moonless nights. During bright-moon periods narrow-band filters with a full width of about 50A can be used.

Experience with SUSI on the NTT has shown that in one arcsec seeing, multiple expo­sures totaling 10 hours will reach a limiting I magnitude of about 27.5. The corresponding figure for the VLT in the same exposure time would be 28.4. For stellar images taken during 0.5 arcsec seeing, the VLT limiting magnitude would be '" 30. .

3.4 Guide stars and active optics

The combination of a large active field and the requirement for thorough bafBing of the telescope optics results in seriously vignetted off-axis field. This greatly restricts the field that is available for guiding. We therefore investigated the possibility that the CCDs camera array itself could be used for guiding. It appears possible that in the near future one ·could read out CCDs at a data rate of a 1 Mpixel S-l and with a readout noise of less than 4 electrons. A 2 x 2 K pixel CCD could then be read out every 4 seconds, and the entire array could be read, one CCD at a time, in 64 seconds. If we continually scanned through the CCD array, reading one detector at a time, every 4 seconds we could compare a new science image corresponding to a 64 second integration, with a previous one from the same CCD. This information could be used to generate error signals for slow . corrections of the telescope tracking.

At the galactic poles the surface density of stars is very low. Each 2 K2 CCD is only 2.4 arcmin2 and at the pole the surface density of stars is so low that only 2.1 stars as bright as 21 mag can be expected to be found over the entire area of a single CCD. But we should remember that the imager will be used only in good seeing; with image sizes less than, say, 0.6 arcsec. And the focal plane is being fed with an 8-meter telescope. Taken together, guiding on a 21 mag star is equivalent to guiding on a 16.5 magnitude star in 1.2 arcsec seeing with a 2-meter telescope. However, the flux levels are not high and care must be exercised in the design of the guiding' system. In the V -band, 2103 photons sec-1 can be expected from a 21st mag star over the 8-meter aperture of a VLT unit telescope. H we assume that 1/2 of these are lost to atmospheric and instrumental (telescope) absorption; we are left with 1000 detected photoelectrons sec-I. H we are limited only by photon statistics this flux is more than enough to obtain guiding error signals.

The CCD array will monitor the history of the seeing and telescope focus during the observation. Experience has shown that periods of superb seeing, even during nights of very good seeing, are short lived. By dividing the exposure up into short segments, images obtained during periods of exquisite seeing are less likely to be spoiled by periods of less good seeing. Deconvolution programs now exist that use images of poor seeing to improve the statistics of the image, while the best seeing images are used to to establish the resolution of the image (Lucy, 1991; Lucy & Hook, 1992; Hook & Lucy, 1992; Hook & Lucy, 1993). It should be possible to give the observer a display of the instantaneous, as well as the predicted final image quality.

It will not be possible to continuously update the active optics. Off axis stars that could be used for telescope image analysis will be vignetted by the required sky bafBing of the instrument. But because the image quality is continuously monitored, it will be possible to note degradation of the image quality caused by departures of required mirror

72

support force field from a force field predicted by a look-up table. When the image quality is no longer accepta.ble, it will be necessary to measure the optical image and up-date ·the force field look-up table. It is an absolute requirement that, at least near the zenith, the telescope optical system be predictable over periods corresponding to tens of minutes. The telescope must also be capable of tracking to an accuracy of better than ~10milli-arc-sec for a minimum period of 60 seconds.

It is evident from this section that the camera will not use the normal telescope guide probes supplied with the Cassegrain adapter. The guiding for high resolution direct imaging is very critical and special. In particular, there is a need to separate telescope wind buffeting from from seeing induced motion that are unlikely to be uniform over the large field. It is necessary to be able to make these corrections using very faint guide stars. The guiding system needs to be very advanced and is probably the most critical part of the camera. It will require careful study. Because the active optics detector system will block at least part of the science field during the measurement period, those observations needed for updating the active optics tables can only be obtained intermittently during the science exposure. If all this turns out not to be feasible, the guide probes supplied with the Cassegrain adapter would have to be used and substantial vignetting of the telescope beam and proportional loss of data would have to be accepted.

3.5 Computer requirements and data handling

3.5.1 Data rates and image storage

A single 2 K2 CCD image will require 16 Mbytes of storage. The assumed data flow is 1 Mpixel S-l There are 16 CCDs in the focal plane; a floating point representation of a single image which will be updated at .the rate of one per minute, will require nearly 256 Mbytes of storage. Clearly this volume of data will have to be reduced. It is likely that a specialized hardware system will have to be constructed. The input images should be classified by image quality and images with similar quality should be centered and added together.

Counting twilight flat-fields and bias frames, one might expect that a single nights observing might produce as many as 100 full frames 30 Gbytes of storage space for the raw images should be provided, and a further twenty Gbytes on a dedicated, fast computer should be offered to the astronomer for "quick look" data reduction. This second computer should be independent of the computers that control the telescope/instrument systems. Speed in handling large arrays will be important and the data reduction computer should not have to carry the overhead of time sharing with the real-time telescope system.

These memory and disk requirements may seem very large at present but price/per/­for/-mance ratios for computer peripherals are currently falling dramatically and it is likely that this will not be a major issue by the time the VLT comes into operation. However input/output speeds are increasing more slowly and could still cause a significant delay in handling data of this size. If necessary, loss-less data compression methods are likely to be effective on this kind of information and would yield compression ratios of easily 2-4 and possibly far more.

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3.5.2 Image restoration considerations

The spherical aberration problem of the Hubble Space Telescope has led to a renaissance of interest in image restoration applied to optical data and many new techniques are currently being studied. Although the resolution enhancements which can be achieved are clearly limited by fundamental information theorems, it seems that much may be achieved particularly using methods which employ additional information about the objects of study or are adaptive in nature. For this reason the instrument described here is intended to make restoration as easy as possible.

J\.n important example of recent improvements is the new algorithm that has overcome the problem of Gibbs oscillations, which has troubled users of both Richardson-Lucy and the Maximum Entropy restoration techniques. Lucy and Hook and Lucy (1994) have described some of these techniques in the ESO Scientific Preprint # 975. These new algorithms have now been tested and were found to be very robust. When data with high SIN ratios are available, it is possible to improve the spatial resolution of the raw images by factors up to three in diameter (Lucy & Hook, 1992).

In order to reliably increase the resolution this much it seems necessary to start off with a pixel size that is about 1/2 the FWHM diameter of the.final resolution of the processed image. This oversampling of the image is necessary in order to determine the details of the image PSF that is used for the deconvolution. In other words, we are boosting the high frequency components in the image. To do this successfully we need both a high SIN image and a very accurate estimate of the telescope and atmosphere smoothing function.

Statistics of the site seeing at Paranal suggest that for about 16 % of the time the site seeing is better that 0.5 arcsec. Thus, it should be possible to frequently obtain, from high SIN data, deconvolved images with a final image quality that is in the range of 0.1-0.2 arcsec. For this reason the resolution goal of the camera was chosen to be 0.1 Mcsec. This instrument will be the VLT equivalent of SUSI on the NTT. The combination of high angular resolution together with wide field will be unequaled by any other facility when the VLT begins operation.

3.6 Sharing the focal plane

It might be possible to share the focal plane with another instrument. An obvious possi­bility would be to share the focal plane with a fiber optics feed to an instrument located at another position. If a 9 arcmin field is sufficiently large it might be possible to switch be­tween the direct imager and a fiber feed using either a slide or a rotary motion. The losses due to the longer fibers would be at least partially compensated for by the elimination of M3.

The wide field imager and a fiber feed have similar requirements: good baffling of the telescope image plane, atmospheric dispersion correction and very accurate guiding with little interference of the science field by the guide probe.

References

Ake, T.B. & Greenstein, J.L. 1980, ApJ, 240, 859 Bartelmann, M. 1993, A&A, 276, 9 Burrows, A., Hubbard, W.B., Saumon, D., Lunine, J.I., 1993, ApJ, (submitted)

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D'Antona, F. 1987," ApJ, 320, 653 Fort, B. & Mellier, Y. 1994, A&ARev, 5, 239 Greenstein, J.L. 1989, PASP, 101, 787 Hartwick, F.D.A, Cowley, A.P. & Mould, J.R. 1984, ApJ, 286, 269 Hook, R.N. & Lucy L.B., 1992, ST-ECF Newsletter 17, p10 Hook, R.N. & Lucy L.B., 1993, ST-ECF Newsletter 19, p6 Lucy, L.B. 1991, ST-ECF Newsletter 16, 6 Lucy, L.B. 1992, AJ, 104, p1260 Lucy, L.B.,& Hook, R.N. 1992, in Proceedings of the 1st Annual Conference on

Astronomical Data Analysis Software and Systems, Thcson, November 1991, p277 • Reid, N. & Majewski, S.R., 1993, MNRAS, (submitted) Richer, H.B. & Fahlman, G.G. 1992, Nature, 358, 383 Tinney, C.G., 1993, ApJ, (in press) Tyson, J.A. & Seitzer, P. 1988, ApJ, 335, 552 Weir, W.N. et al. , 1993, Sky Surveys: Protostars to Protogaiazies, ed. B.T.Soifer Wynne, C.G. 1994, MNRAS, 269, L37

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PRELIMINARY DESIGN DOCUMENT FOR THE

VISIBLE HIGH ANGULAR RESOLUTION CAMERA (VHARC)

G. Weigelt

Max-Planck-Institut fUr Radioastronomie, Auf dem Hiigel 69, D-53121 Bonn, Germany

1.. Scientific Objectives

With the 8 m VLT telescopes and speckle methods it will be possible to achieve unprecedented angular resolution for many classes of astronomical objects and to ob­tain otherwise not attainable astrophysical information. Speckle interferometry (Labeyrie 1970) can yield the autocorrelation of an object with diffraction-limited resolution. True diffraction-limited images can be obtained by the Knox-Thompson method (Knox and Thompson 1973), speckle masking method (Weigelt & Wirnitzer 1983; Lohmann et al. 1983), and other techniques. At A f"J 4000 A the diffraction-limited resolution of a tele­scope with diameter D = 8m is AjD f"J 0.010" (lOmas). Since the diameter of the VLT mirrors is 3.3 times larger than the diameter of the Hubble Space Telescope, speckle imag­ing yields 3.3 times higher resolution than the HST at any given wavelength, if the object is bright enoug~. The limiting magnitude is about 18 th magnitude for objects consisting of a small number of resolution elements in nights of good seeing (see Sect. 3). The signal-to-noise ratio in the reconstructed image is inversely proportional to the number of resolution elements in the reconstructed image and inversely proportional to the third power of seeing (FWHM). Because of this strong seeing dependence, we plan to apply the VLT adaptive optics system, without or with Laser guide star, in order to improve the SNR in the reconstructions.

In addition to high angular resolution, spectral information can be obtained by vari­ous speckle spectroscopy techniques, for example, O(x,A)-projection speckle spectroscopy (Grieger and Weigelt 1990, 1992) and high-resolution objective prism (slitless) speckle spectroscopy (Weigelt et al. 1986, Afanasyevet al. 1992). A reciprocal spectral resolu­tion of 1 A to 0.1 A can easily be obtained.

The following objects are examples of important candidates for speckle imaging and speckle spectroscopy. Of course, the list of projects is not complete. Most of the objects discussed have already been successfully resolved by previous speckle observations which demonstrates the feasibility of such projects. More detailed discussions of most of the projects can be found in review papers, for example, by Appenzeller (1979, 1988), Davis (1979), McAlister (1979, 1988), Ulrich (1979, 1981, 1988), and Refsdal & Surdej (1992, 1994).

(1) Mira stars, red giants and supergiants: wavelength dependence of diameter, shape, extended atmospheres, limb darkening, bright surface features, dust envelopes, compan-Ions. Mira stars show spectacular variations of angular size with wavelength which are related to

77

TiO opacity. This wavelength dependence of the diameter was observed for the first time using speckle interferometry (Labeyrie et al. 1977). For Mira and X Cygni very interesting wavelength dependence of the structure and size of the atmosphere was found (Bonneau et al. 1982, Foy 1988). Studies of the exploratory models of Scholz and Bessell (Bessell et al. 1989, Bessell & Scholz 1989, Scholz 1993) show that monochromatic radii observed in suitably chosen filters and at suitably chosen phases are very sensitive diagnostic tools for the investigation of the structure of a Mira photosphere. Accurate monochromatic radius measurements, combined with conventional colour and line profile observations, may pin down the photospheric parameters and may discriminate between different cur­rent pulsation models. Speckle observations with single VLT telescopes can provide the basic observations for a quantitative model analysis of the photospheric structure of a large number of Mira variables .. No quantitative analysis of any Mira photosphere exists to date.

Observations of a Ori show time-varying bright features on the surface (Buscher et al. 1990), possibly due to convective hot-spots. Surface structure observations are important for our understanding of large-scale stellar convective processes and for the interpretation of stellar spectra. The goal of speckle observations with the VLT is to search for similar features on many objects and to measure the wavelength dependence of the diameter due to variations in TiO opacity which, similar to Miras, can be used for investigating the photospheres of non-variable M type giants (see Scholz 1985). The important applications of angular diameter measurements (e.g., determination of the absolute emergent flux distribution at the stellar surface and effective temperatures) have been discussed by, e.g., Davis (1979, and references therein). Finally, several observers have shown that interferometric observations at visible wavelengths can contribute to our understanding of these objects by providing measurements of circumstellar dust disks (e.g., Ricort etal. 1981, Roddier & Roddier 1983) and companions.

(£) Spectroscopic binaries: luminosities, masses, distances, mass-luminosity relation. Binary star studies play a fundamental role in observational astrophysics as they provide the only direct means for measuring stellar masses. McAlister (1988) discussed the im­portance of very accurate speckle observations of close binaries and of the determination of magnitudes and colors of the individual components. Extremely important are double­lined spectroscopic binaries. For double-lined spectroscopic binaries the luminosities as well as the masses of the individual components can be determined in a fundamental man­ner and in this way new accurate points can be added to the empirical mass-luminosity relation (McAlister 1988).

(9) Pre-main sequence (PMS) stars: mass determinations with close binaries, circum-stellar envelopes and disks. During the past few years many IR speckle observations of circumstellar envelopes or disks around various PMS stars have been reported (e.g., Beckwith et al. 1984, Zinnecker et al. 1987, DeWarf & Dyck 1993, Leinert et al. 1991, 1994). At optical wavelengths only a few preliminary observations exist. Images with the highest possible resolution of about 10 mas are important to test present models of the PMS star envelopes and highly collimated mass outflows. For mass determinations it is necessary to determine the orbit of close binaries with separations of lOmas to 100mas in the nearby star forming re-

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and the origin of the gas in the narrow line region.

(7) Gravitational lenses: resolving the closest lensed QSO images, mass of the lensing galaxy, determination of the Hubble constant Ho, micro-lensing and the physical structure of quasars, statistical gravitational lensing and the cosmological density of compact objects in the mass range 107 - 1011 M0 . Several speckle groups have resolved gravitational lenses of magnitude 16 to 17 with telescopes of 1.5m to 3.6m diameter (e.g., Triple QSO PG1115+08, Hege et al. 1981, Foy et al. 1985, Weigelt et al. 1986). Most of these observations were made with short observing times, in nights of average seeing, and with simple detectors. Therefore, these observations are an illustration of the future possibilities with 8 m telescopes. An extremely important application of gravitational lenses is the determination of the mass of the lensing galaxy. In this way very distant galaxies otherwise not accessible can be investigated. Measurements of the time delay of the brightness fluctuations of the individual lensed QSO images will allow the Hubble Constant to be determined (Refsdal 1964) if the QSO is intrinsically variable and if a reasonable mass model of the lensing galaxy can be derived. Furthermore, individual stars in the lensing galaxy may induce recognizable micro-lensing brightness variations for one (or more) of the macro-lensed QSO image(s). Speckle imaging, polarimetry and spectroscopy of micro-lensed images offer a unique possibility to retrieve physical information concerning the structure of the QSO central source and emission-line region (see Refsdal and Surdej 1994 for a review on this subject). Finally, a survey for multiple lensed QSO images with angular separations in the range 0.01"-1" will directly enable the setting of values (or limits) on the cosmological density of compact objects in the Dlass range 107 -1011 M0 (Surdej et al. 1993).

References

Afanasyev, V.L., Balega, Y.Y., Orlov, V.G., Vasyuk, V.A., 1992, A&A, 266, 15 Appenzeller, I., 1979, in: ESA/ESO workshop on Astronomical Uses of the Space Telescope, eds F. Macchetto, F. Pacini, M. Tarenghi, p. 47 Appenzeller, I., 1988, in: NOAO-ESO Conf. on High-Resolution Imaging by Interferometry, ed. F. Merkle, p. 19 Beckwith, S., Zuckermann, B., Skrutski, M.F., Dyck, H.M., 1984, ApJ 287, 793 Bessell, M.S., Brett, J .M., Scholz, M., Wood, P.R., 1989, A&A 213, 209 Bessell, M.S. and Scholz, M., 1989, IAU ColI. 106, 67 Bonneau, D., Foy, R., Blazit, A., Labeyrie, A., 1982, A&A 106, 235 Buscher, D.F., Haniff, C.A., Baldwin, J.E., Warner, P.J., 1990, MNRAS 245, 7p Davidson, K., Humphreys, R.M., 1986, A&A 164, L 7 Davidson, K., 1989, IAU Colloq. 113, p.101 Davis, J., 1979, IAU ColI. No. 50, High Angular Resolution Stellar Interferometry, eds. J. Davis and W.J. Tango, p. 1-1 DeWarf, L.E., Dyck, H.M., 1993, Astron. J. 105,2211 Ebstein, S.M., Carleton, N.P., Papaliolios, C., 1989, ApJ 336, 103 Foy, R., 1988, in Instrumentation for Ground-Based Optical Astronomy. Present and Future. ed. L.B. Robinson (Springer Verlag), p. 345 Foy, R., Bonneau, D., Blazit, A., 1985, A&A 149, L13 Grieger, F., Weigelt, G., 1990, SPIE 1319,440 Grieger, F., Weigelt, G., 1992, ESO Proc. -High Resolution Imaging by Interferometry II, p. 481 Hege, E.K., Hubbard, E.N., Strittmatter; P.A., Worden, S.P., 1981, ApJ Letters 248, L1 Hof-

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gions. The distance of the nearest star forming regions is about 150 pc. At this distance semi-major axes of, for example, 3AU correspond, to 20 mas separation and short enough periods.

(,I) Peculiar stellar objects: e.g., Luminous Blue Variables (LBVs) and Be stars. Eta Car is an extremely luminous, eruptively unstable Luminous Blue Variable with spectacular shell ejections. As the most extreme known LBV, Eta Car is uniquely critical for theoretical studies of the LBV phenomenon (Davidson 1989). Eta Car has a surprising sub-arcsec structure near its central star. Three objects at separations 0.11", 0.18", and 0.21" have been detected by speckle observations (Hofmann & Weigelt 1988). Speckle observations at different wavelengths, speckle spectroscopy, and speckle polarimetry with 10 mas resolution, is required to study the nature of Eta Car B to D, and of fainter circumstellar structures. By comparing the images obtained at different epochs, it will be possible to derive the proper motion and the date of the ejection of the objects.

Visible speckle imaging in the continuum and in emission lines, speckle spectroscopy, and speckle polarimetry can also contribute to our understanding of Be stars by providing measurements of the circumstellar envelopes. Quirrenbach' et al. (1994) resolved for the first time the envelope of the Be star' Tauri. The morphology of the 10 mas x 3 mas Ha image is most easily interpreted as a disk seen almost edge-on.

(5) Very massive LMC stars and clusters. The observations will allow us to reveal the multiplicity of the most massive LMC stars (e.g., Pehlemann et al. 1992 and references therein). The results will set new upper limits for stellar masses and better define the shape of the upper IMF. Observations in strong emission lines will examine the ratio of Wolf-Rayet to normal O-type stars which is related to the evolutionary state of the stellar groups.

(6) Seyfert galaxies: structure, ionisation, and kinematics of the Narrow-Line Region, circum nuclear starburst regions, intermediate region between NLR and BLR. Previous speckle observations have resolved the NLR and circumnuclear starburst regions in several AGN down to scales of '" 10 pc (e.g. Ebstein et al. 1989, Hofmann et al. 1989, Afanasyev et al. 1992, Mauder et al. 1992, 1994). Most of the NLRs turned out to be clumpy aggregates, often confined in linear or cone-like structures. This confirms earlier results derived from spectroscopic investigations (e.g. Wagner & Appenzeller 1988) that the NLRs consist of individual clouds whose physical and dynamical properties differ from each other. It is necessary to obtain high resolution images in the light of different diagnostic lines in order to investigate the structure and ionization conditions of the clouds. For example, observations of objects which have Extended NLRs are required in order to study the distribution of the clouds on small scales, to determine the size spectrum of the individual clouds, to search for possible correlations with the dynamical state of the ENLR and the luminosity and spectrum of the central source, and to measure th~ velocity of the individual clouds by speckle spectroscopy with about 1 A reciprocal spectral resolution (Ha and 0 III) to study the kinematics of the NLR, and to test present AGN models. Ulrich (1979, 1981, 1988) discussed the intermediate region between the broad line region and the narrow line region and the importance of high-resolution (10 to 100 mas) observations for our understanding of the mass loss' from the broad line region

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mann, K.-H., Mauder, W., Weigelt, G., 1989, in: ESO Proc. Extranuclear Activity in Galaxies, p.35 Hofmann, K.-H., Weigelt, G., 1988, A&A 203, L21 Knox, K.T., Thompson, B.J., 1974, ApJ Lett. 193, L45 Labeyrie, A., 1970, A&A 6, 85 Labeyrie, A., Koechlin, L., Bonneau, D., Blazit, A., Foy, R., 1977, ApJ 218, L75 Leinert, Ch., Haas, M., Lenzen, R., 1991, A&A 246, 180 Leinert, Ch., Richichi, A., Weitzel, N., Haas, M., 1994, Near-Infrared speckle observations of Herbig Ae/Be stars, in: Proc. Nature and Evolutionary Status of Herbig Ae/Be, ASP Conf. Ser. Leung, K.-C., Moffat, A.F.J., Seggewiss, W., 1979, ApJ 231,742 Lohmann, A.W., Weigelt, G., Wirnitzer, B., 1983, Appl. Opt. 22,4028 McAlister, 1979, IAU Coil. No. 50, High Angular Resolution Stellar Interferometry, eds. J. Davis and W.J. Tango, p. 3-1 McAlister, H., 1988, in: NOAO-ESO Conf. on High-Resolution Imaging by Interferometry, ed. F. Merkle, p. 3 Mauder, W., Appenzeller, I., Hofmann, K., Wagner, S., Weigelt, G., Zeidler, P., 1992, A&A 264, L9 Mauder, W., Appenzeller, I., Wagner, S., Weigelt, G., Zeidler, P., 1994, High-resolution optical images of the starburst ring around the Seyfert nucleus of NGC 7469, A&A May 1, 1994 Pehlemann, E., Hofmann, K.-H., Weigelt, G., 1992, A&A 256, 701 Petrov, R.G., Balega, Y.Y., Blazit, A., Borgnino, J., Foy, R., Lagarde, S., Martin, F., Vassyuk, V.V., 1992, in: ESO Conf. Proc. High Resolution Imaging by Interferometry II, eds. J. Beckers and F. Merkle, p. 435 Quirrenbach, Refsdal, S., 1964, MNRAS 128, 307 Refsdal, S., Surdej, J.: 1992, "Gravitational Lensing", Invited discourse during the XXlst Gen­eral Assembly of the International Astronomical Union (Buenos Aires, July 1991), IAU 'High­lights of Astronomy' Vol. 9, 3-32, J. Bergeron (ed.) Refsdal, S., Surdej, J.: 1994, Reports on Progress in Physics 57, 117 Ricort, G., Aime, C., Vernin, J., Kadiri, S., 1981, A&A 99, 232 Roddier, C., Roddier, F., 1983, ApJ 270, L23 Scholz, M., 1985, A&A 145, 251 Scholz, M., 1993, IAU CoIl. 139, New Perspectives on Stellar Pulsation and Pulsating Variable Stars, eds. J .M. Nemec and J .M. Matthews, p. 201 Surdej, J., Claeskens, J.F., Crampton, D., Filippenko, A.V., Hutsemekers, D., Magain, P., Pi renne, B., Vanderriest, C., Yee, H.K.C.: 1993, Astron. J. 105,2064-2074 Ulrich, M.H., 1979, in: ESA/ESO workshop on Astronomical Uses of the Space Telescope, eds F. Macchetto, F. Pacini, M. Tarenghi, p.261 Ulrich, M.H., 1981, in: ESO Conf. on Scientific Importance of High Angular Resolution at Infrared and Optical Wavelengths, eds. M.H. Ulrich and K. Kjar, p. 411 Ulrich, M.H., 1988, in: NOAO-ESO Conf. on High-Resolution Imaging by Interferometry, ed. F. Merkle, p. 33 Wagner, S., Appenzeller, I., 1988, A&A 197, 75 Weigelt, G., Baier, G., 1985, A&A 150, L18 Weigelt, G., Baier, G., Ebersberger, J., Fleischmann, F., Hofmann, K.-H., Ladebeck, R., 1986, Opt. Eng. 25, 706 Weigelt, G., Wirnitzer, B., 1983, Optics Lett. 8, 389 Zinnecker, H., Perrier, Ch., Chelli, A., 1987, IAU Symp. 115, p. 71

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2 .. Technical Concept and Instrument Specifications

Figure 1 shows a block diagram of the proposed instrument which consists of at least four different optical subsystems and at least two different detectors. Each optics/detector combination can be chosen. The four optical subsystems and the two detectors are dis­cussed briefly in the following sections. A discussion of the VHARC concept can also be found in the ESO documents on the Visible High Angular Resolution Camera (Doc­ument Number VLT-SPE-ESO-13400-0303). We agree to the concept and the technical specifications discussed in this ESO document .

. (1) Imaging optics: It is possible to choose: (a) microscope objectives to obtain various magnifications in the range of 1 x to about 25 x, (b) interference filters, (c) a polarimeter, and (d) Fabry-Perot filtergraph.'

(2) Courtes Monochromator: This monochromator allows the selection of one or sev­eral spectral bands in the spectral range of 3800 to 8000 A with a bandwidth between 1 A and 10 A (Labeyrie et al. 1977, Foy 1988).

(9) Projection speckle spectrograph: An anamorphic imaging system of cylindrical lenses is used to obtain I-dimensional projections of the 2-dimensional speckle interfero­grams. The I-dimensional speckle interferograms are spectrally dispersed by a grism to obtain projection speckle spectrograms. From the spectrograms object/spectrum recon­structions O(x,.>t) are obtained, which are I-dimensional images in x-direction and spectra in y-direction (Grieger & and Weigelt 1990, 1992).

(4) High-resolution spectrograph: This spectrograph is used to obtain spatially resolved spectra (similar to long-slit spectra) of compact emission line objects. The spectral res­olution is chosen high enough to resolve individual emission line components caused, for example, by clouds of different radial velocities (for example, clouds in the NLR of AGN). A reciprocal spectral resolution of about 1 A and diffraction-limited spatial resolution of 10 to 20 mas can be obtained. '

(5) Detectors: see Fig. 1.

(6) Image processing: Quicklook reconstructions are obtained by on-line speckle inter­ferometry and speckle masking. Experience with already existing speckle cameras shows that even on-line speckle masking bispectrum processing is possible with arrays of Digital Signal Processors or hard-wired correlators.

(7) Partial adaptive optics: The SNR in the reconstructed images depends on the third to fourth power of seeing (FWHM). Because of this strong seeing dependence, we plan to apply the VLT adaptive optics system, without or with Laser guide star, in order to improve the SNR in the reconstructions.

(8) Additional observing modes: In addition to the observing modes 1 to 4 discussed above, the following observing modes should be available in the instrument: (a) speckle polarimetry, (b) high spectral resolution differential speckle interferometry (Petrov et al. 1992),

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( c) speckle imaging and speckle spectroscopy with partial adaptive optics, (d) speckle imaging with a Fabry-Perot tuneable filtergraph, (e) diluted pupil interferometry, and (f) self-referenced speckle holography with a separate optical wavefront sensor.

(9) Telescope focus: Nasmyth, Cassegrain, or Coude focus is possible.

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Telescope Beam

+-+ translation stage

Imaging Optics Monochromator Spectrograph 1

Magnification: 1 to 25 Type: Courtes Type: projection

, Interference filters Polarimeter Spectral resolution: spectrograph

Fabry~Perot 20000 Reconstructions:

filtergraph Spectral channels: 2 O(x,>.)

Spatial resolution: - Spatial resolution: - Spectral resolution: -10 - 20 mas 10 - 20 mas 1000

Field: 3 arcsec Wavelength range: Spatial resolution: 3800 - 8000A 10 - 20 mas

Differential Field: 3 arcsec Wavelength ranle:

Speckle 3800 - 8000

Interferometer Field: 3 arcsec

Detector 1 Detector 2

Type: CCD Type: photon counting Pixels/frame: 10242 Camera: Frame rate: EBCCD or MAMA

~ translation stage 10 to 50 frames/s "- or intensified CCD Field: 3 arcsec Frame rate: Wavelength range: 10 to 50 frames/s

3800 - 10000 A Pixels/frame: 10242

Data storage: Field: 3 arcsec fast tape drives Wavelength ranle:

3800 - 8000

Fig. 1 Optical subsystems and detectors of the proposed Visible High Angular Resolution Camera (VHARC)

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Spectrograph 2

Type: resolution of emission line components

Reconstructions: - long slit spectra

Spectral resolution: 5000

Spatial resolution: 10 - 20 mas

Wavelength range: 3800 - 8000A

Field: 3 aresec

3 .. Instrument Performances

(1) Angular and spectral resolution: As discussed above, the angular resolution of speckle imaging and speckle spectroscopy is diffraction-limited, for example, 10 mas at 4000 A and 20 mas at 8000 A. In addition to atmospheric aberrations, telescope aberra­tions are also completely overcome.

(2) Limiting magnitude: The limiting magnitude can be determined theoretically, by computer simulations, by laboratory experiments, or by astronomical applications. There is a quite good agreement between all four techniques. Theoretical calculations (Roddier 1981) show that, for example, a signal-to-noise ratio of about 10 in the reconstructed speckle interferometry power spectrum is obtained for the following observing parameters: object of 16 th magnitude, Fried parameter ro = 20 cm, quantum efficiency = 0.1 of the total system, standard deviation of the wind velocities in the atmospheric layers of about 5 mis, and about 2 hours observing time.

The computer experiments reported by Hofmann & Weigelt (1993) show, for eXam­ple, that about 10000 speckle interferograms are required for D Iro = 10 (D = telescope diameter, ro = Fried parameter) and for 25 to 100 photoevents per speckle interferogram. In this computer experiment the object was a cluster of 5 stars. For 25 photoevents per interferogram the mean photometric error of the stars in the speckle masking reconstruc­tion was about 15%, for 100 photoevents per interferogram the photometric error was about 5%. Smaller errors are obtained if more than 10 000 speckle interferograms are reduced. 10 000 speckle interferograms correspond to only about 500 s observing time for a frame rate of 20 frames per s. A count number of 25 photoeventslframe corresponds to about magnitude 16 to 18 for an 8 m telescope and typical values for quantum efficiency, exposure time, and filter bandwidth.

We conclude that the limiting magnitude of speckle imaging with about four hours observing time is about 18 or 19 for simple objects consisting of only a few point sources and it is only about 15 for extended, diffuse objects since the SNR in the reconstruction is inversely proportional to the number of resolution elements in the reconstruction.

(9) Competition with the Hubble Space Telescope: At visible wavelengths the HST is diffraction-limited. The angular resolution is about 40 to 50 milli-arcsec. Since the wings of the HST point spread function are very weak, it is obvious that the dynamic range of HST images is much higher than the dynamic range of.speckle reconstructions and that the limiting magnitude of HST observations is many magnitudes fainter. The main ad­vantage of speckle imaging and speckle spectroscopy is the about 3.3 times higher angular resolution. This factor is critical for many projects. Other advantages of speckle imaging are imaging spectroscopy, the larger flexibility to choose filt~rs for diagnostic lines and line components, tuneable filtergraphs (Fabry-Perot and Courtes monochromator), and differential speckle interferometry with sub-diffraction-limited resolution.

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CRIRES: THE VL THIGH-RESOLUTION INFRARED SPECTROMETER

G. Wiedemann, B. Delabre, A.F.M. Moorwood European Southern Observatory, Garcbing

1 Introduction and Overview

The enonnous scientific potential of 1-5 pm high resolution (R-IOS) specttoscopy and the large VLT specific gain of this observational mode were stressed by the VLT Working Group on Infrared Aspects (VLT Report No. 51, July 1986). At that time, small (64x64) infrared array detectors were only just becoming available and the provisional proposal was for high resolution specttoscopy using Fabry-P6rots in combination with a cooled, medium resolution grating specttometer. Following further study and the rapid evolution of infrared array fonnats the ESO Instrumentation Plan, endorsed by the STC in March 1990, proposed combining medium resolution 1-5 pm specttoscopy with imaging in the Medium Resolution Spectrometer/lmager (subsequently approved for development as ISAAC) plus the development of a dedicated high resolution cryogenic echelle specttometer or FrS and/or extending a 'visible' echelle specttometer into the near infrared. As these options offer different scientific capabilities, an important aim of the Feb. 1992 ESO Workshop on High Resolution Specttoscopy with the VL T was to prioritize them on the basis of the main science drivers. The overwhelming consensus was to give highest priority to a cryogenic echelle instrument at Nasmyth (rather than the combined focus previously considered) due to 1) its large sensitivity gain relative to an FrS (Moorwood & Wiedemann 1992; Ridgway & Hinkle 1992) and 2) the high scientific importance obviously attached to the 2-5 pm region which cannot be competitively covered by an extended 'visible' spectrometer. This conclusion was reported to the STC in May 1992 together with the first results of an immersion grating development program, started in ESO to assess the feasibility of developing large immersion gratings as a means of maximizing the R9 product (cf. Moorwood & Wiedemann 1992, Wiedemann 1992). Subsequently, further progress has been made in both the immersion grating development program and the instrument design concept(s).

This report on the Concept Definition Phase has been prepared at the request of the Scientific Priorities for the VLT Working Group as input to the STC decision process regarding the approval and future planning of this instrument It includes a summary of the high priority science programs on which the instrument requirements are based together with a preliminary design and its predicted perfonnance. The proposed concept is a fixed Nasmyth platfonn mounted instrument comprising 1) a calibration unit, 2) a fore-optics section providing for field de-rotation and adaptive correction using the target (point source) object for wavefront sensing 3) a pre/cross-disperser unit and 4) a 1-5 J.lm high resolution echelle section equipped with conventional and/or immersion gratings. Basic parameters and performance data are summarized in Table 1. The predicted magnitude limits are - 8 - 10 mag. fainter than reached by existing instruments, and 1-2 mag. fainter than could be achieved with a similar instrument on a4 m telescope. As no comparable instrument appears to be foreseen so far for other 8-10 m class telescopes it offers a unique opportunity for the investigation of previously

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inaccessible phenomena across a broad range of astrophysical areas. Various technical and scientific options for enhancing further its potential, but at increased cost and complexity, are also identified.

Recent development

The first issue of this concept definition study for a cryogenic infrared echelle spectrograph at the VLT was submitted to the 'Scientific Priorities for the VLT STC Working Group in March 1994; it has subsequently been made available to the astronomical community. The present, updated, version incorporates responses to this document and additional input relevant to the scientific justification of this instrument from various sources.

The need for a high-resolution IR spectrograph is clearly felt and has been expressed by the community. There is also a consensus that the instrument parameters sensitivity and high spectral resolution should be given highest priority, followed by large spectral coverage.

The enonnous sensitivity provided by a cryogenic echelle array spectrograph at a large telescope constitutes a fundamental improvement Various scientific disciplines exploit this in different ways. Envisaged applications range from the study of faint objects with high spectral resolution to very high SIN observations of brighter targets, to observations with high temporal resolution, or any of these in a spatial (linear) imaging mode. We believe, that the envisaged instrument satisfies these various requirements in the best possible way.

The technical part of this report has also been updated to include the progress achieved in the past seven months.

Table 1: Cryogenic IR emelle spectrometer - summary of main characteristics

Parameter Value Focus Nasmyth Wavelenlrth ran£e 1-5 Jim R9 20,000-40.000 emelle 40 em leneth, 31.6 Hnes I mm detector 1024xl024 array Pixel size 0.1" Slitlen~h ~ 50" or 10" (x-dispersed) Limiting mag. (3CJlhr) 18 (.D. 17.5 (H), 17 (K), 14 (L), 12 (M)

2 Science Objectives

2.1 Overview

A cryogenic echelle spectrometer provides a sensitivity increase corresponding to - 8 mag. over previously available instruments. An additional 1- 2 mag. can be gained by using an 8 m instead of a 3.5 m telescope. High-resolution (R>I()4) infrared spectroscopy has traditionally been limited to a very small number of very bright sources and excluded entire classes of objects, e.g. late-type dwarf

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stars. The combination 8 m telescope - echelle array spectrometer will provide a quantitative and qualitative improvement of the scientific capabilities.

Scientific areas expected to receive a strong boost from the VL T IR echelle spectrometer are· summarized in Table 2. This list has been extracted from the contributions to the ESO 1992 Workshop on High-Resolution Spectroscopy with the VLT, the 'Gedankenproposals for the Scientific Exploitation of the VL T and the responses to a letter soliciting scientific input from the European community circulated in early 1994. Additional valuable input was received at the ESO Workshop 'Science with the VLT (July 1994) and the Brussels Workshop on Laboratory and Astronomical High-Resolution Spectroscopy (Aug.1994)

2.2 Impact of the VL T IR echelle

The advantages provided by the VLT IR echelle spectrometer - quantified in more detail below -have the potential to further the scientific areas listed above in various ways. The combination of high spectral resolution and large spectral coverage with the unprecedented infrared sensitivity of a cryogenic dispersive instrument results in a unique facility. The VLT IR echelle will not only extend the number of observable objects considerably, but provide a qualitative scientific improvement by offering capabilities not available through any other technique or from any other instrument/telescope combination.

Sensitivity improvement can be exploited for all observations aiming at fainter objects, higher spatial (extended sources), spectral and temporal resolution.

Increased sensitivity will pennit to 1) extend studies from a few objects to a complete sample of an object class, 2) access other, less luminous classes of objects and 3) study previously unobservable phenomena in recognized but not entirely understood objects. For instance, in the case of the late­type stars detailed below, it will be possible 1) to cover the entire range of spectral types and chromospheric activity levels in giant stars, 2) to observe main sequence and dwarf stars with higher surface gravity where magnetic fields, flare activity etc. may be particularly strong, and 3) to study convection via small differential Doppler-shifts of spectral lines in molecular bands or to observe stellar oscillations with high temporal resolution.

Many extended sources have shown structure on small scales when observed with enhanced spatial resolution (e.g. auroral regions observed in H3+ on Jupiter, Saturn). Studies of planets, stellar outflows, circumstellar disks, planetary nebulae, star forming regions etc. may benefit from the greater spatial resolution of the VL T echelle in the same way as, or - in combination with the spectral infonnation - even more than pure imaging applications. The prospects for solar system observations with a high-resolution, high-sensitivity imaging IR spectrometer have recently been highlighted by Encrenaz (1994).

The immense importance of high signal-to-noise, high-accuracy and high-resolution measurements on bright sources for astrophysics in general has recently been emphasized by several authors (Kurucz 1992; Grevesse & Sauval 1994; Dravins 1994).

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2.3 Examples of detalled studies

IR spectroscopy at an 8 m telescope can be applied successfully to hot stars (non-L TE H-lines, atmospheric sttucture, disks, mass transfer), but offers even greater benefits in the study of stars of spectral type F5 and cooler (e.g. Grevesse & Sauval 1994). Their IR spectrum contains lines from virtually an light diatomic molecules, including CNO in various combinations. The transition originate from a large number of excitation states and can therefore diagnose a wide range of physical conditions and processes.

Two problems in stellar physics which will benefit enormously from the VL T IR echelle are described below in more detail. (These projects partly reflect the areas of interest of the Instrument Responsible).

2.3.1 Stellar magnetic fields

The understanding of stellar evolution and dynamic processes in stellar interiors and atmospheres is closely tied to the understanding of the role of magnetic (B-) fields. The direct measurement of stellar B- fields is based on the Zeeman-effect, which is characterized by a linearly increasing separation of G- and 7t-components relative to the Iinewidths with wavelength. Unambiguous, i.e. model-independent B-field measureinents on stars require full separation of the ~man components and can only be performed in the IR. The first magnetic field measurements on an M dwarf flare star, AD Leo, were reported by Saar and Linsky (1985), who recorded the Zeeman-split profiles of two titanium lines near 2.2 pm. A large fraction (73 %) of the stellar surface is covered with an average field of 3800 G. Their spectrum of AD Leo required a 6 hr integration to obtain a SIN of 25. Available instrumentation had been pushed to its limits to record one example of a magnetic field measurement.

The most powerful magnetic field indicators known today are the high-Rydberg lines of MgI near 12.32 pm. Since their discovery (Goldman et ale 1980), identification (Chang & Noyes 1983) and interpretation (Carlsson et ale 1992), they have triggered a revolution in solar magnetic field stUdies (see reviews at IAU Symposium 1.54). To this date there· have been three attempts by two groups to search for these lines in stars, both with prototype visitor instruments at large telescopes (e.g. Jennings et ale 1986). Integration times of several hours were necessary for a mere detection of the lines in very bright but magnetically inactive red giant stars (a Ori, a Tau). Dwarf and main sequence stars, the prime targets for magnetic field studies will only become accessible with an IR echelle spectrograph at a very large telescope.

2.3.2 CO in stellar atmospheres

Atmospheric sttucture

The rotation-vibration lines of CO at 2.3 and 4.6 pm originate at photospheric and chromospheric altitudes in stars of spectral type F5 and cooler. Formed near L TE, the lines are relatively easy to interpret. Pioneering studies on the sun (Ayres & Testerman 1981), and Arcturus (Heasley et aI. 1978), and subsequent investigation of a few bright stars by Wiedemann et aI. (1994) have revealed a serious conflict between the tclassicalt stellar models derived from UV/visible observations and the IR

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molecular line observations, which led to the postulation of extended, molecular cooling-dominated areas on stellar surfaces in coexistence with the classical chromospheres. The confinnation of this 'thennal bifurcation' scenario (Ayres 1981) would have profound implications for the classical models (e.g. Ayres & Unsley 1981, Kelch et al. 1978), which depend strongly on the spatial averaging properties of the species from which they are derived (UV /visible lines on the Wien-side of the Planck-curve at stellar surface temperatures). Dust fonnation, for instance, suspected to begin at the stellar surface-interstellar medium interface could be affected through 1) the altered grain-destroying UV radiation field spectrum and/or 2) the very low temperatures created by CO- and SiO cooling catastrophes (Muchmore et ale 1987; Cuntz & Muchmore 1994).

Convection

Convection is a major energy and momentum transfer mechanism in cool stellar atmospheres. The precision measurement of Doppler-shifts of CO line cores as a function of excitation state (which avoids many of the problems of the classical line bisector method) provides direct information about the altitude-dependent vertical velocity fields. This technique is currently being tested with high­quality spectra (Fig. 1 on the following page) recorded with the FrS at the Kitt Peak 4 m telescope (Wiedemann & Hinkle 1994). Fig. 1 also reveals the limits of current instrumentation/telescopes: A21r integtation was required to obtain this spectrum of a Boo (K2 ill, m 4~7pm = -3 f), one of the brightest stars in the sky. SIN improvement through longer integration was not possible since atmospheric variations and time delay between stellar and reference measurements began to limit the accuracy.

Stellar oscillations & dynamics of stellar atmospheres

The solar 5-min. oscillations can be observed in the CO infrared transitions with great sensitivity due to the large number of lines formed over a wide altitude range (Ayres & Brault 1990) and their close coupling to the local kinetic temperatures (L TE line fonnation). Different oscillation modes are triggered in the (direcdy unobservable) interior and the convective zone of late-type stars, from where they propagate to the visible surface. Extrapolating to stellar research, time resolved observations can probe processes in the stellar interior relating to evolution, age etc. Measurement of correlated temporal variations in Doppler-shifts and line intensities probe the dynamics of cool star atmospheres. Such measurements, using CO fundamental and overtone lines (Ayres & Brault 1990) and MgI lines at 12.3 pm (Deming et ale 1988) with large FrS have contributed substantially to the understanding of the solar photosphere/chromosphere region. Other stars are out of reach for current telescopes/instrumentation, but will be accessible with the combination of high sensitivity (time resolution!) and spectral resolution of a cryogenic echelle spectrograph at the VL T. The study of acoustic shock waves in slowly rotating (i.e. magnetically inactive) stars. (e.g. Cuntz et ale 1994) would constitute a major contribution to the understanding of the elusive heating processes leading to the widely observed but poorly explained chromospheres in late-type stars (see, e.g., Mechanisms of Chromospheric & Coronal Heating, eds. Ulmschneider, Priest & Rosner [Springer: 1991]). In this case, the sensitivity of the cryogenic echelle permitting stellar observations is fundamental, since in the Sun - usually the primary case study for late-type stars - dissipation of magnetic (in addition to acoustic) energy accounts for substantial non-radiative energy input into the

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1.2

1.0

0.8

)( :l

;;:: Q) 0.6 > :;:; 0

Q) ~

0.4

0.2

Arcturus 4.7 J.Jm CO spectrum

256 pixel detector

1024 pixel detector

2140 2142 2144 2146 2148 2150 wavenumber

Fig.1. High-resolution (R=120,OOO) spectrum of aBoo (K2, Ill) recorded in a 2 hr integration. The spectrum has been corrected for atmospheric absorption by dividing by a lunar thermal emission spectrum. All spectral lines are of stellar origin and represent the 6v=1 (v=0-7) transitions of carbon monoxide. The portion of the FrS spectrum shown here corresponds to the range covered by the IR echelle with a single 1024 detector at one grating position. The short horizontal bar indicates the range covered by a 256 pixel detector, which contains at most one line from each vibrational excitation state.

The spectrum of this evolved giant star is particularly rich with lines from many rotational / vibrational levels and isotopic species (13C, 170, 180). Cooler and less evolved sources feature only lines from the vibrational ground state, which are spaced by typo 4 em-I.

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upper atmosphere and docs in principle not pcnnit the identification of the individual causes of the observed structures.

Again, few attempts to study oscillations of IR lines in cool stars with adequate time resolution have been reported in the past because of the lack of instrumental sensitivity.

3 High-resolution IR spectroscopy at the VL T

High-resolution infrared spectroscopy-lagging behind many other astronomical disciplines because of the associated technical difficulties - is one of the areas which will benefit most from the gains provided by the VL T. The factors contributing to the advantage are evaluated below. It should be emphasized that due to the restrictions applying to previous observational studies, these gains translate into an increase of the scientific yield which is much greater than expressed by the numerical values.

3.1 The large telescope advantage

Due to the dependence of the l-dimcnsional telescope ~tendue, De, on the size of the (largest possible) grating in a spectrograph, the full light collecting advantage of an 8 m telescope is realized for compact sources whose light can be concentrated onto the smaller slit at the larger telescope. In that case the signaVnoise ratio increases as 0 2 in both the detector-limited and background-limited cases. For brighter stars (Table S) source radiation dominates the noise and the SIN increases as D. Resolved sources, including residual seeing disks larger than 0.2 arcscc, yield a gain proportional to 01/2 because of the increase of the telescope etendue along the slit (1bis gain increases to 0 in the detector limit if the field is concentrated anamorphica1ly parallel to the slit). An additional signal gain of = 20 % occurs since the entrance slit selects the central portion of the source image where the energy is concentrated.

These SIN gains translate into integration time reductions equal to the square (!) of the SIN improvement, leading to greater observing efficiency as well as higher temporal resolution in the study of variable phenomena.

3.2 Spectral resolution

In addition to the enormous information gain provided by high spectral resolution, there is a sensitivity increase for all unresolved spectral features through elimination of radiation not contributing to the signal in both the source - and the background-limited cases. This gain associated with higher spectral resolution is equivalent to that of larger telescope size!

3.3 Spatial resolution

The SIN advantage for extended sources is not as great as that for point sources, yet the gain in spatial information at the higher resolution matches that of pure imaging applications.

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3.4 Instrument and location

Sensitivity gains greater than those due to the mirror size can result from instrument features, e.g. increased simultaneous spectral coverage or resolution. Usually these enhanced features are correlated with large instrument size, which could prohibit the implementation at low-emissivity foci of smaller telescopes.

3.S Site advantage

The proposed high-resolution IR echelle with R = 100,000 at the VLT would be the only instrument in the world with this capability at a site with excellent atmospheric conditions.

3.6 Large detector advantage

For an echelle with R "., 1 00,000, already one 1024 amy would provide a qualitative improvement of the science capabilities over a 256 pixel detector, as the larger total bandpass (10 em-l at 5 pm) pennits simultaneous multi-line observations (e.g. rotation-vibration transitions in molecular bands, see Fig. 1), beyond the threshold of 'single line' observations prevailing with smaller arrays.

4 High-resolution IR spectroscopy at other telescopes

The availability of detector arrays has boosted the construction of infrared echelle spectrometers at ground-based facilities in recent years: IRSHELL, a mid-IR spectrometer (R S 25,000) owned by the Univ; of Texas, Austin (Lacy et aI. 1989) has been used extensively as a visitor instrument at large telescopes including ESO's 3.6 m The cryogenic 1-5 pm spectrometer CGS4, with low and high-resolution (RA20,000) modes (Wright et aI. 1993) is scheduled for 70 % of the observing time at the UKIRT. The new common user 1-5 pm spectrometer CSHELL (Greene et al. 1993) at the NASA/lRTF (R S 40,000) has been requested for 40 % of the scheduled observing time. A prototype mid-IR spectrometer (CELESTE) using a multi­pass technique for very high spectral resolution (0,.02 em-I single pass diffraction-limited), developed at NASAl Goddard has been used as a visitor instrument at the McMath, IRTF and Mt. Palomar facilities. The cryogenic IR spectrometer PHOENIX under construction at the KPNO (Hinkle et aI. 1990) will be the f1l'st true high-resolution 1-5 pm echelle spectrometer accessible to the common user. With a resolving power of up to 100,000 it is destined to replace the facility Fourier transform spectrometer at the 4 m Mayall telescope. During the planning phase of this instrument, a comprehensive scientific justification has been established, based on numerous contributions from the astronomical community (predominantly previous users of the Ff Spectrometer). This compilation of protagonist scientific projects is available from the PIs Drs. K Hinkle and S. Ridgway.

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5 Instrument Concept

5.1 Design drivers

5.1.1 Optical design drivers

Our instrument concept is aimed at high resolving power (R = lOS), high sensitivity and - because of the rapidly increasing scientific yield associated with this capability - large simultaneous spectral coverage. It is a cryogenic spectrometer utilizing only reflective optics except for the entrance window and the predispersing elements and operating from 1 - 5 pm with a possible extension to 8-14 pm. The concept emphasizes the wavelength range ~ 2 pm, however, where reduction of thermal background radiation bears the greatest sensitivity gains. (A dedicated 1-1.8 pm spectrometer would not have to be restricted in its capabilities by the requirements of the thermal IR).

Due to the relatively large instrument size and the tight stability requirements, the Nasmyth platform with a low - emissivity focus is the natural location for the IR echelle. Although suited best for point­source observations, the instrument provides an optical de-rotator for the long-slit mode.

The above mentioned goals can be realized with the largest available echelle (40 em length), yielding R9 = 20,000 at an 8 m telescope. Observations of fainter point sources with a 0.2 arcsec slit require at least low-order adaptive correction for light concentration. The image stabilization provided by the secondary telescope miITor M2 can in principle perform this task, though it is not clear to what degree and how·effectively. Alternatively, the cryogenic echelle could use a Nasmyth AO system similar to that proposed for CONICA, however at the expense of increased emissivity, complexity and additional reflection losses. As a third option (possibly in addition to the low-frequency, large amplitude correction by M2) the pupil mirror in the fore-optics could be configured as a fast tip-tilt mirror inside the cryogenic instrument This approach would involve little technical risk and could be accommodated by the present optical design. A more detailed performance assessment is in preparation.

5.1.2 Detectors

The yield of an echelle spectrometer increases gready with the size (number of pixels) and the performance of the detector may. Until recendy (July 1994) the largest mays used in IR astronomy had 256 x 256 formats. Two major detector developers, Rockwell Ind. and Santa Barbara Research Corp., had been preparing to produce 1024 xlO24 detectors for the 1-2.5 pm (HgCdTe, 18.5 pm pixel size) and 1-5 pm (lnSb, 27 pm) regions. Both detector types have been, fabricated in the meantime. The Rockwell HgCdTe detector has already been used by a Univ. of Hawaii group to produce images of Jupiter during the Shoemaker-Levy 9 collision. From a technical point of view, the large format mays are largely extensions of present detectors with no fundamentally new technology involved. One can therefore expect similar performance in terms of quantum efficiency and noise behavior (see Table 3). Further improvements in the future will push the detection limits below those given in Table 5.

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Table 3: Parameters of the 1024 xl024 SBRC InSb detector array

Wavelenlrth ranee 1 - 5 um ( InSb) Format 1024 x 1024 pixels Ouantum efficiency N80% Read noise <25 e- r.m ... Dark current N.l e-I sec Sensitivity limited by dark current at 1 < 2 jJm

The reported low dark current of the new InSb detectors is competitive with that of the 2.5 pm cutoff HgCdTe devices. This could simplify the instrument design, since separation of the optical train into a 1-2.5 pm and 3-5 pm channel would no longer be necessary, provided the instrument design achieves adequate strilylight levels.

5.2 Instrument description

The following description refers to the block diagram below (Fig. 2) and the 3-D drawing (Fig. 3) on the following page. The instrument will be housed in a vacuum tank for operation at cryogenic temperatures achieved with closed cycle refrigerators. Optical system and radiation shields are cooled to 60-80 K, while the detector operating temperature will be in the 20-30 K range.

5.2.1 Functional description

Functionally the instrument can be divided into the four sections: calibration unit, fore-optics. predisperser and high-resolution section:

calibration unit

predlsperser

high-resolution unit

Fig. 2 Cryogenic infrared emelle, block diagram

The calibration unit provides for flux calibration and detector flatfielding, as well as for the absolute wavelength calibration of the echelle. It includes a gas cell to serve as an absolute frequency standard (on-line) for very high precision velocity measurements.

The fore-optics section incorporates a cold pupil stop to eliminate all light not passing via the telescope secondary. An image of the source is fonned at the entrance slit of the predisperser, which detennines the spectral resolution of the echelle (I) and acts as the main field stop. Image

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stabili:zation or low-order adaptive optics correction and image de-rotation can be performed by the optical arrangement of the fore-unit A small camera views the field image reflected off the predisperser entrance slit

The ptedisperser optically isolates one echelle order. In the single order mode, the predisperser also minimizes the spectral range admitted into the high-resolution section to that fraction of an order covered by the array, again to eliminate parasitic light The long entrance slit exploits the second dimension of the detector for linear imaging. Another small camera unit (not shown) collects the light rejected by the polished predisperser exit slit to (simultaneously) record a low-resolution spectrum. possibly of the entire 1-5 JIm range, on a separate detector. This unit also monitors the wavelength range admitted into the high-resolution section. The basic concept furthermore foresees a crossdispersed mode for each of the wavelength bands. which is realized by shallow gratings mounted interchangeably in the predispetser.

The beam expanding from the predisperser exit slit into the high-resolution section is collimated to illuminate the large echelle. Tilt-tuning of the echelle centers the desired wavelength on the detector. The dispersed light is collected by the separate camera and focused onto the detector array(s). The camera produces the required plate scale of 0.1 arcsec/pixel. with two detector elements sampling the nominal 0.2 arcsec entrance slit A sampling finer than two pixels per resolution element can be obtained by scanning the grating in steps corresponding to fractions of a pixel.

5.1.1 Optical design

Fig. 3 represents a conservative optical design which achieves the required resolution and image quality over a large field with a layout that is reasonable from a mechanical and cryo-thermal point of view. It relies on components that are relatively easy to manufacture and test. Further details of the optical design can be found in a separate report.

The calibration unit (not shown in Fig. 3) is located outside the cryogenic environment and contains halogen lamps or glower sources for flux calibration and detector flatfielding, as well as He-Ne m lasers (1.532 JIm, 3.39 JIm, 2.4 JIm) or spectral lamps (Ar, Xe, Ne) for the absolute wavelength calibration of the echelle. The radiation is directed into the instrument via a flat flip mirror. A small gas cell can be inserted into the beam in front of the instrument to serve as an absolute on-line frequency standard for very high precision velocity measurements when filled with a suitable medium (CO, N20).

The input parabola M2 in the fore-optics section forms an image of the telescope secondary on a cold pupil stop. This stop can be configured as a Lyot mirror (M3) equipped with a drive for low­order (tip-tilt) adaptive correction. (Even if AO is not available, there is still a gain compared to a smaller telescope). The output parabola M4 - identical to M2 for optimal aberration compensation -projects the stabilized source image at fl15 onto the predisperser entrance slit S 1. The symmetric arrangement with the roof mirror Ml allows the fore-optics unit to rotate during extended source observations to compensate for the apparent sky rotation.

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The predisperser is a low-resolution spectrometer acting as the order selector for the echelle. The low resolution (numerically equal or greater than the echelle order number, i.e. 60 at 1 pm and 12 at 5 pm) can be provided by prisms which have the advantage of easy operation, high efficiency and excellent out-of-band rejection ratio, the latter being extremely important in view of the allowable straylight levels on the detector (lower than the dark current of - 1 e-/sec). Interchangeable prisms out of sevCral materials (Ge, liF, Amtir) arc utilized to optimally cover the 1-5 pm region. The adjustable predisperser exit slit (=echelle entrance slit) S2 detennines the spectral range entering into the high-resolution section. It can be widened to admit more light, for instance, if an additional detector is used in the focal plane to cover a larger part of one spectral order.

The prcdisperser entrance slit can be widened to degrade the echelle resolution and admit more light from an extended source. (This improves the SIN for continua in the radiation noise dominated case and for all features in the detector-limit). The instrumental limit is set by the resolving power of the prcdisperser, which is slit-dependent in the case of a prism. The resolving power of the echelle is inversely proportional to the slit width. Spectral coverage can be re-gained if a camera with variable magnification is implemented in the echelle section.

The high-resolution unit consists of collimator, echelle, camera and detector array(s). The f/15 beam expanding from the slit S2 is collimated by the spherical mirror M to illuminate the grating with a beam of 10 em diameter. The grating is an R4 echelle with 31.6 lines per mm, blazed at 71 deg. (Alternatively, an immersion grating could probably be used with this beam size). The echelle is mounted on a tilt stage and operated in-plane (a ... 76 deg, p ... 67 deg, r- 0) in order to mlnimize image curvature and tilt

The camera (M 8 - M 10) is a 3-rnirror unit which produces the ff1 focal ratio and the required image quality over a field corresponding to at least two 1024xlO24 detectors.

The separation of collimator and camera avoids potential scattering problems inherent in double -pass configurations. In addition, it leaves the flexibility .to adapt the camera focal ratio to changes in the detector pixel size or sampling requirements. The plate scale on the focal plane detector is 0.1 arcsec/pixel, with two pixels sampling the nominal 0.2 arcsec entrance slit

S.2.3 Mechanical design

Although no detailed mechanical design has been produced for this concept study, several aspects can be derived already from 1) the particular optical perfonnance requirements of this instrument and 2) results obtained during the mechanical design and prototyping phase of the fll'St cryogenic infrared VLT instrument, ISAAC, which employs a number of similar mechanisms. .

The low signal levels of - I e- /sec prevailing in faint source/shon wavelength applications of the cryogenic echelle require that special attention be paid to the shielding of the detector from all parasitic (thennal) radiation. This involves placing all mechanical functions and driving motors in the cold environment, inside a light-tight radiation shield without penetrations. Excellent performance

. and reliable operation of cryogenic movable mounts and stepper motors has already been demonstrated (ISAAC Critical Design Review, Summary Document VLT-TRE-ESO-14100-0483). Where the stability/flexure requirements for the echelle may be more demanding, they are likely much

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easier to meet because of the stationary operation of the instrument on the Nasmyth platform. In addition. the optical design aims at minimizing the number of movable elements. e.g. by avoiding fIlter wheels. and by keeping the range of motions (e.g. echelle tilt, see 7.5. options) small.

Observations like ultra-high precision velocity monitoring. which require stability and repeatability (of echelle position. for instance) exceeding the thermal or mechanical capabilities of the instrument, will be conducted using on-line calibration instead.

The relatively large size of the instrument is a potential concern from the cryo-thermal point of view, as cooldown time and steady-state cooling power are subject to limits. Special attention will have to be given to the design of low-absorptivity radiation shields and low-conductivity mechanical fIxations for minimal thermal transfer as well as to the minimization of the cold mass and surface area. Again, the final optical design (see 5.2.4. Alternative designs) shall take this aspect into consideration.

5.2.4 Alternative designs

In addition to the presented optical design, alternatives have been studied to improve various aspects, which. however, involve some uncertainty at this time (concerning manufacturing tolerances, alignment criticality. performance. cost, etc.). For instance. a 3 -' mirror collimator (e.g. Delabre 1993) also used as camera in double-pass would achieve greater compactness of the instrument This approach has indeed been adopted for the IR echelle as presented in the 1990 Instrumentation Plan. The feasibility of this design depends critically on 1) the cumulative surface fIgure and alignment errors of the necessary aspherical mirrors. and 2) the scattering properties of the fIrst collimator mirror which is in direct view of the detector. An assessment of this concept wiD be available after the evaluation of a similar 3-mirror collimator currently prototyped for the VL T medium-resolution imager/ spectrometer ISAAC.

The present design concept allows for replacement of the conventional echelle by an immersion grating (see 7.4. options) or possibly by a 2 x 1 mosaic of 40 em R4 gratings to achieve R9~4O,(XX>' The larger optical depth of these gratings can be exploited in several ways: 1) It offers greater resolving power for a given slit width, 2) it provide~ for greater throughput (9) -within instrumental limits - at a given resolving power on extended sources or 3) it permits a reduction of the size of the instrument for a fixed R9 product.

The instrument design is driven substantially by the available echelle formats. .From an optical (echelle orders are much longer than the detector) as well as from a mechanical (tilt range necessary to scan a complete order) point of view it would be advantageous to reduce the order lengths by using an echelle with a lower groove density than 31.6/mm, which is near the limit for mechanical ruling. Coarse gratings with 4 - 20 lines per mm, produced by chemical etching on :Silicon substrates. are already under investigation as a part of the preparatory studies for this instrument (see 7.5. options).

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6 Performance

6.1 Sensitivity

Fig. 4 shows the contributions to the detector signal which detennine the noise and the detection limits given in Table 4.

Table 4: Limiting magnitudes for SIN= 3 in 1 hr.

I (pm) 1.2 1.6 2.2 3.3 4.2 4.7 mal 18 17.5 17 15 13.5 12.5

Assumptions: Dark signal (incl. detector current, stray light) 1 eo/sec, total emissivity .1, total efficiency 10% (incl. atmospheric and telescope transmission, instrument efficiency, detector Q.E.), accumulation of > 100 e- per exposure to overcome read noise

With a read-noise expected in the 10 e- range, the faint limits are dictated by the dark current (= 1 eo/sec) in the J, H (outside of strong OH airglow lines) and K bands and by the thennal background radiation (telescope and atmosphere) at the longer wavelengths. The SIN increases as the square root of the integration time. In the faint source limit, the SIN increases linearly with the signal.

Table 5 shows the stellar magnitudes at which the flux from the source begins to dominate the noise. The SIN increases as the square root of the signal in the fundamental bright source limit To illustrate the system efficiency, Table 6 lists the integration times required to achieve a SIN of 100 on a 5th and a 10th magnitude star at all wavelengths.

Table 5: Stellar magnitudes where sensitivity is source-noise limited, O' mall 1 . . al D am overs er te esc~.1S ~ro~ort1on to

1 (pm) 1.2 1.6 2.2 3.3 4.2 4.7 mal 16 15 14 10 7 5

Assumption: AO, point source

Table 6: Integration times required for SIN= 100 on a 5th and a 10th mag star

~(pm) 1.2 1.6 2.2 3.3 4.2 4.7 TIsec], malt = 5 <1 1 2 2 3 5 T[sec), maR = 10 50 100 160 350 6.000 15,000

6.2 Spectral coverage

The simultaneous spectral coverage with a (Npix=) 1024 pixel detector in one grating setting corresponds to 500 spectral elements, e.g. 10 em-I at 5 pm (see Fig 1). The fraction of a single order covered by one detector array is equal to Npix * m / 2 R; numerical values are given in Tab 7. One

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way to reduce order lengths for a better match to the detector fonnats is described later (7. 5. options).

The spectral coverage naturally increases with the number of arrays employed in the focal plane. The large field provided by the present camera design could accommodate at least two 1024 detectors. The implementation of crossdispersion increases the spectral coverage by a further factor corresponding to the number of orders in each band (Table 7). As the relative position of the orders on the wavelength scale is fixed according to Am- 11m, however, the actual yield of the crossdispersed mode depends on the usefulness (infonnation contained in the source spectrum and atmospheric transmission) of those fractions of the orders covered by the detector.

Table 7: Number of spectral orders falling into each wavelength band for a 31.6 linelmm echelle, and single order coverage

Band J H K L M number of orders 12 7 6 8 4 fractional order coverage 0.24 0.19 0.15 0.09 0.06

6.3 Velocity resolution

The velocity resolution (Llv) achievable with a grating spectrograph is proportional to the resolving power and the SIN (e.g. Hatzes & Cochran 1992), and improves with the number of observed lines: Ll v = c / ( R lie SIN lie nl12). Under the conditions listed in Tab 6. and assuming simultaneous coverage of 50 lines it will be possible to determine velocity variations in the 5 - 10m / sec range.

7 Options

Several extensions can be added to the basic concept to provide enhanced observing capabilities at the expense of design complexitY, and, in some case! a small development risk.

7.1 Crossdispersion

The large fonnat of the foreseen detector makes crossdispersion an interesting option for greater spectral coverage. A separate report contains computed echelle fonnats possible with the present instrument concept. Four to eight orders can be registered on a 1024 array in each of the wavelength bands. Only fairly conventional low-angle gratings are necessary. The gratings must be mounted interchangeably in the predisperser. The crossdispersed mode can be accommodated by the present optical predisperser concept.

A number of additional functions can be implemented by viewing the large detector field directly or via a flat diverter mirror with an. attached detection system:

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u G1 U)

......... U)

c 0 ~ -U G1 Qi

0\ .9

Detector Signals 6 - - - - - - - - - Jar. ma9.. = 0 - -- ---- ----- ----5

4

3

2

-1

-2

···O··H····· • - - - - - - - ..............•. ; .. :~-,:-::~,.-..... .. - - - - - - ~-. - star ........... . ::::: ... okgl~~.::~:~~.::: ............ ' ...... :~ .. ~~!~Li~rk - - - ~~ag = _1~ _

.......... :.:: ..... - .......... c.urrl:nt .1 e- sec .-... ,~ .. -.-.-.-.:.:: ..... -,.-:- .-:-:::-::.. . ....................... . air low···......... ..................... . ......... .. R = 100.000 ........................................................................... ..

........................................................................................

-3~ _______ ~ _______ ~ ______ ~ ______ ~~ ______ ~ _____ ~ ______ ~ ______ ~

1.0 1.5 2.0 2.5 3.0 3.5 4.0 4.5 wavelength [microns]

Fig. 4 Detector signal levels (detected photons per sec) in comparison with the detector dark current (e-' sed.

Assumptions: Background temperature: 280 1<, total emissivity 10%, total efficiency 10% (including atmospheric and telescope transmission, instrument efficiency, detector Q.E.), R=lOS, 0.1" /pixel, Drel = 8 m; OH airglow values adopted from Malhara et al. (1993).

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5.0

7.2 10 pm extension

All instrument optics including the entrance window operate throughout the infrared. Observations in the astronomically important 8-14 pm range require a sma1l modification of the predisperser unit and the addition of a detector I filter module at the focal plane. The detector must be shielded and cooled to =10 K. The similarly cooled bandpass filter (R ~ 20) suppresses the broadband mid-IR radiation from the cryogenic instrument below that of a spectral resolution element The filter also easily serves to isolate one order of the 31.6 1/mm echelle (m=5 at 10 pm). The entrance slit of the instrument has to be widened for observations in the mid-IR to at least the diffraction limit of 0.3 arcsec, yielding a maximum resolving power of 60,000.

7 .3 Variable magnification

Cameras with variable magnification could maximize the instrument yield for various applications, for instance 1) a faster camera matches a wider slit (used for lower spectral resolution) to the pixel size in order to exploit the larger spectral coverage of the detector. 2) Very high spectral resolution can be obtained at the shon wavelengths through a slower camera and a narrower entrance slit As the nominal 0.2 arcsec slit is four times wider than the diffraction limit at 2 pm, a reduction is possible at J, H and K as long as the light losses at the slit are tolerable.

7.4 Immersion grating

The immersion grating as a means of increasing the R9 product of an infrared echelle spectrograph has been introduced at ESO several years ago (Wiedemann 1992). Diffraction grating structures on single crystal silicon (an excellent IR material with a high refractive index) be obtained through anisotropic chemical etching adopted from micro-electronics manufacturing. Experimental verification of the immersion grating principle has been demonstrated with a· sma1l prototype. Subsequently, investigation of several aspects of etched silicon gratings and material bulk properties has started. An etched wafer grating of 10 em length (the largest one ever made) has been delivered to ESO in March 1994. The grating features a symmetric blaze angle of 54.7 degrees and a rather large groove spacing of 50 pm. First evaluation has Shown a very clean grating profile with excellent diffraction properties from two sides! (This distinguishes etched gratings· from mechanically ruled gratings where nonnally only one facet of the grooves can be used). Parallel to the grating development, a study of bulk properties of silicon and other infrared materials is under way. A large single crystal blank has been procured and successfully x-ray oriented, figured and polished in preparation of the bulk optical evaluation. This blank will subsequently provide the substrate for the flfSt large monolithic immersion grating with an effective optical length greater than the largest available conventional grating. A feasibility study identifying details of the entire manufacturing process will be canied out in collaboration with the micromechanics group at the Munich Fraunhofer Institut for Solid State Technology. A realistic assessment of large infrared immersion grating should be available in the near future.

Progress has been made by another group at the Fraunhofer Institute with the development of an anisotropic etching process for Gennanium. The availability of large single crystals of Gallium Arsenide (for which the etching process is known) has recently been reported. Both developments

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are important for the development of immersion gratings with very low absorption over large optical paths in the near-IR as well as for the extension into the 8-14 pm astronomical window.

7.S Coarse gratings

A side effect of the etched grating research is the possibility to obtain line spacings much coarser than those produced by mechanical ruling. Gratings with a few lines per nm could be used in the IR like conventional echelle in the visible, allowing a better match of the spectral formats to the detector size and geometry. The benefit for the mechanical design would consist in the reduction (possibly elimination) of the echelle tilt range required for a complete scan of one order. The feasibility of small gratings with 4-10 lines per nm has been demonstrated and the extension to large area echelles is being explored. Very good quality in a 10 em 20 lineImm grating has already been achieved.

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8 Reference List

Ayres. T.R. & linsky. J.1 .• 1975, Ap.1. 200, 660.(AL) Ayres, T.R .• 1981, Ap.J. 244, 1064 Ayres, T.R. &Testennan, L., 1981, Ap.J. 245. 1124 Ayres, T.R. & Brault, J. 1990. Ap.J. 363, 705 Bessel, M.S .• 1992 in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VLT (ed. M.-H. Ulrich) Carlsson, M .• Ruttren. R.I. & Shchuidna, N. 01992. A&A 253.567 Chang. E. S. & Noyes, R.W. 1983 ApJ. 275, L11 Cuntz. M .• Rammacher. W. & Ulmschneider. P. 1994, ApJ. 432.690 Cuntz. M. & Muchmore, D.O. 1994 ApJ. 433,303 Delabre. B. 1993 in ISAAC Opt. design, VLT-TRE-14100-0488 Deming. L. D .• Boyle R.I . • IeMings. D.E.& Wiedemann 0 .. 1988, Ap.J. 333.978 Dravins. D .. 1994, in The impact of long-term monotoring on variable star research,(eds. de Groot & Sterken). Kluwer Acad. Publishers Encrenaz. T .• 1994. The Messenger, No. 75. 12 Encrenaz. T., Combes. M., Saint-Pet O. 1992 in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VLT (ed. M.-H. Ulrich) Goldman, A. et aI. 1980 , New Atlas of IR Solar Spectra Vol. 2. Dept. of Pysics, Univ. of Denver Greene, T.P .• Tokunaga, A.T., Toomey. D.W. & Carr. I.S. 1993. SPIE. Vol. 1946.313 Grevesse. N. & SauvaI. A.J. 1994 in Molecular opacities in the stellar environment,( ed. U.G. If1jrgensen) IAU Coll146. Hartmann, L. 1992 in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VLT (ed. M.-H. Ulrich) Hathes. A.P. & Cochran. W.D. 1992. in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VLT (ed. M.-H. Ulrich) Heasley. I.N .• Ridgway. S.T .• Carbon, D.F., Milkey, R.W. & Hall, D.N.B. 1978. Ap.J. 219.970 Hinkle. K.H., Ridgway. S.T. & Stuart. F.E. 1990 SPIE Vol. 1235. 119 Jennings. D.E .• Deming. L. D .• WiedemaM. G. & Keady. 1. J .• 1986. ApJ 310. L39 Kelch. W.L.. linsky. I.L .• Basri. G.S .• Chiu, H .• Chang, S., Maran. S.P. & Furenlid. I. 1978. ApJ. 220, 962 IGiufl. H.-U .• & Stanghellini, L. 1992 in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VL T (ed. M.-H. Ulrich) IGiufl, H.-U., Siebenmorgen. R .• Wiedemann. O. & Zijlstra, A. 1993. 'Gedankenproposal' for the scientific exploitation of the VL T. ESO) Kurucz. R.L. 1992. Rev. Mexicana Astron. Astrof 23. 187 Kurucz, R.L. 1994. presentation at the 'Laboratory and Astronomical HIgh-Resolution Spectra' Workshop. Brussels. Aug 29-Sept2 Lacy. J.H .• Achtermann. I.M .• Bruce. D.E .• Lester. D.F., Arens. J.F.. Peck M.C. & Gaalema. S.D. 1989, PASP 101.1166 Lambert, D.L. 1992 in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VL T

. (ed. M.-H. Ulrich) Maihara. T., Iwamuro. F .• Hall. D.N.B .• Cowie, L.L., Tokunaga. A.T. & Pickles. A.J. 1993 SPIE. Vol. 1946.581

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Maillard, J.-P. & Mitchell. G.F. 1992. in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VLT (ed. M.-H. Ulrich) Moorwood, A.F.M. & Wiedemann. G. 1992, in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VLT (ed. M.-H. Ulrich) Muchmore, D.O., Nuth, I.A. & Stencel, R.E. ,1987, ApJ. 315 L141 Ridgway, S.T. & Hinkle, K.H. 1992 in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VLT (ed. M.-H. Ulrich) Saar, S.H. & Linsky, I.L. 1985 ApJ. 299, lA7 Strassmeier, K.G. 1992, in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VLT (ed. M.-H. Ulrich) . Wiedemann, G. 1992, in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VLT (ed. M.-H. Ulrich) Wiedemann, G., Ayres, T.R., Saar S.H. & Jennings, D.E. 1994 Ap. I. 423 Wiedemann. G. & Hinkle, K.H. 1994 in preparation Wright, G.S, Mountain, C.M .• Bridger, A., Daly, P.N., Griffin, J.L. & Ramsay, S.K. 1993, SPIE Vol. 1946,547 Wampler, EJ. 1992 in Proceedings of ESO Workshop on High Resolution Spectroscopy with the VLT (ed. M.-H. Ulrich)

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