CIRPASS - Design Study - Part A


Ian Parry and François Piché

1. Introduction

This document describes the first part of the technical aspects of the design study for CIRPASS (Cambridge IR PAnoramic Survey Spectrograph). The main objectives of this study are:

  1. To identify instrument design forms that will meet the scientific specifications and evaluate their feasibility.
  2. To predict as accurately as possible the costs of the designs considered.
  3. To identify the staff and infrastructure resources required to complete the project.
  4. To estimate the timescales involved in completing the various project milestones.

Throughout this document it is assumed that CIRPASS will be used on an 8m telescope.

2. Multi-fibre or Multi-slit?

CIRPASS is primarily a multi-object spectrograph and the scientific specification calls for a large multiplex gain and the ability to do very deep spectroscopy. Another requirement of CIRPASS is that it must record the spectra at a spectral resolution of R~3000 to avoid the strong OH lines. All of these requirement demand a very fast camera (especially on an 8m telescope) so that as much spectral information as possible can be crammed on to the detector and also so that the noise characteristics of the detector are reduced by minimising the number of pixels used for each measurement. This latter point is particularly important for the NIR because of the high spectral resolution requirement and because, unlike CCDs, IR arrays have relatively high readout noise and dark current and it is not possible to on-chip bin.

To compare a multi-slit design with a multi-object design we have looked at the following criteria: throughput, field of view, survey speed, sky-subtraction, versatility and cost. A summary of our conclusions is presented in Table 1.

Throughput. For a fibre feed in the NIR it should be possible to achieve an efficiency of 80% or better compared to a slit of the same aperture. We are assuming here that the loss due to light that misses the fibre or slit is the same for both.

Field of View. For a fibre fed system this is limited by the telescope and not the spectrograph. For a multislit version of CIRPASS one array is 6.8 arcmin in the spatial direction if we assume there is no anamorphism. [In practise CIRPASS has to be a reflection grating system and not a grism system such as LDSS-2 because of the high spectral resolution requirement. There is therefore likely to be some anamorphism which will reduce the field of view in the spatial direction.] To avoid serious loss of spectral coverage we have assumed that ~100 pixels in the dispersion direction can be deployed for positioning the slits so the total field of view for 4 arrays is ~18 arcmin2.

Survey Speed. This depends on the multiplex gain, the wavelength coverage and the time it takes to reach the required spectral resolution. For a fibre system each 1.5 arcsec fibre is 4 pixels in the dispersion direction and each spectrum uses 10 rows of pixels. This gives 0.12m of wavelength coverage for one array at R=3000 and we can have 100 fibres per array. The multiplex-gainwavelength-coverage product is therefore 48microns. For a multislit system with 1.2 arcsec wide slits (3 pixels) and ~900 pixels for the wavelength coverage we have a total wavelength coverage of 0.14microns. A slit length of 20 arcsec is 50 pixels so we have a total multiplex gain of 80 (for a limiting magnitude of J=23 there are ~160 galaxies in a field of 18 arcmin2) for 4 arrays giving a value of 11.2microns for the multiplex-gainwavelength-coverage product. Even allowing for the better throughput of a multislit device and for the fact that not all fibres are placed on objects, a fibre system is likely to complete a particular survey faster than a multislit one.

Sky-subtraction. This is the major problem with fibres and the comparison of the sky-subtraction capability of fibres versus slits is a complex issue. A detailed discussion of this is given in the next section. The basic conclusion presented is that it is possible to do very good sky subtraction with fibres and there is no fundamental reason for believing that fibres will not give good sky-subtraction.

Imaging. Clearly, a fibre spectrograph that sits on the dome floor and can only be fed with fibres cannot be used as a general-purpose imager. However, for CIRPASS this is not an important science driver. For a multislit system it is possible to replace the grating with a mirror to use it as an imager. However, given the fast f/1.2 camera and the 0.4 arcsec/pixel scale the utility of doing this is not ideal - an image scale of 0.1-0.2 arcsec/pixel is more appropriate. Furthermore, at f/1.2 a broadband image will saturate the array faster than it can be read out.

Transferability. For a fibre system matching the f-ratio of the telescope is achieved by the choice of microlens that feeds the fibre at the telescope focal plane. Transfer to another telescope therefore is simply a case of making a new fibre feed and making a new plate-holder unit to suit the new telescope. For a multislit system transfer to another telescope requires making a new collimator. The mechanical interface requirements are also more complicated because the whole spectrograph has to go on the telescope.

Cost. Comparing a fibre sytem with a multislit sytem many of the component are the same (gratings, cameras, detectors etc.). For a fibre system there is the extra cost of the fibres themselves. For a multislit system there is the increased cost of the mechanical structure which has to be light-weight, compact and flexure-free because it has to be mounted on the telescope. In the case of an IR instrument such as CIRPASS this structure also has to be cryogenic. Furthermore, to carry out imaging at ~f/3 more optics are required. The cost of meeting these requirements more than offsets the cost of the fibre feed and a multislit system will be more expensive than a multi-fibre system. So far in the design study we concentrated on fibre systems and our current cost estimates are ~£500-600K (excluding detectors). We do not have a detailed cost estimate of a multislit system.

Table 1 - comparison of fibres versus slits

Throughput (above atmosphere to data file)
Field of view
18 arcmin2
Multiplex-gainwavelength-coverage product
at F/1.2
Transferability between telescopes

3. Sky-subtraction

It is generally believed that sky-subtraction with a slit-based system is more accurate than with a fibre-based system and to do spectroscopy at the faintest possible levels one should avoid using fibres because they simply cannot be made to work. This view is however incorrect and the limiting signal-to-noise ratio that can be achieved for both fibres and slits is fundamentally set by Poisson statistics. Several workers have made very deep observations with fibres to demonstrate this. Elston and Barden (NOAO newsletter #19, 1 Sept 1989) reached a magnitude of R=22.3 (approximately B=24) by using a beam-switching technique. This magnitude limit compares well with that of deep multislit surveys (e.g. LDSS-2). They used a plug-plate fibre system on the 4m telescope at Kitt Peak and their on-source exposure time was 3 hours. More recently Cuby and Mignoli (SPIE vol 2198, 99, 1994) showed that the theoretically faintest limits could be achieved without beam-switching thus getting the largest possible multiplex gain. They used MEFOS on the ESO 3.6m telescope and successfully observed objects as faint as bJ=22.6 in an exposure time of only 90 minutes.

So why are most people sceptical about using fibres for ultra-faint spectroscopy? Basically it is because with fibres one has to take much greater care when making the observations and doing the data-reduction in order to eliminate systematic errors. In order to achieve good sky-subtraction one has to derive an object-plus-sky spectrum and a sky spectrum that very accurately estimates the sky component of the object-plus-sky spectrum. For a multislit system the sky spectrum and the sky-plus-object spectrum are adjacent both on the sky and on the detector and the systematic errors in both are practically the same and therefore cancel out in the sky-subtraction process. Interpolating along the slit using blank sky regions either side of the object helps enormously here. However, for a fibre system the two spectra to be subtracted from one another are typically neither adjacent on the detector or on the sky and so the systematic errors associated with each often do not cancel completely if a simplistic data reduction algorithm is used. The most dominant systematic errors are due to a lack of adjacency at the detector and include wavelength sampling, scattered light and spectrograph vignetting. What the astronomers from NOAO and ESO mentioned above have demonstrated is that with a little care and attention these systematic effects can be successfully dealt with to give results as good as those of a multislit system.

So one can perhaps justify the poor reputation that fibres have because they are indeed a little harder to use. However there are further reasons why fibres have a much poorer reputation than they deserve. A common misconception about fibres systems is that their transmission varies with time (for a given field configuration) because the fibres twist and bend as the telescope tracks, making flat-fielding inherently inaccurate. This is simply not true for modern robot based systems and is probably a throw-back to the very early days when loose-fitting fibres where sometimes used in poorly manufactured plug-plates. Well engineered plug-plates and fibre ferrules can certainly have the stability required. Another problem is that examples of faint work with fibres that would convince the sceptics are extremely rare. One reason for this is that fibre systems typically can access a very large field compared to a multislit device but until quite recently it has been very hard to obtain the photometry and astrometry for a very faint sample over such a large field because photography rather than CCD images were the only data available. Conversely, when deep CCD images are available to provide targets for a deep spectroscopic survey these have been well matched to the field sizes of multislit spectrographs and therefore these have been used to do the work because of the perceived risks of fibre work. Finally, whenever fibres are used the temptation to deploy most of the fibres on objects and very few on blank sky has also tended to push fibre users away from faint samples.

In the case of CIRPASS there are several features which will greatly help make the sky-subtraction process accurate. CIRPASS will have a small number of configuration choices and can therefore be operated essentially as a fixed-format device. This means that many calibration procedures to determine wavelength calibration, scattered light contributions and spectrograph vignetting can be repeated and the properties of the system can be accurately measured and understood and stored as a database. This coupled with a sophisticated, dedicated data reduction package will make sky-subtraction both accurate and painless for the observer. A white pupil Baranne design also helps because fibres at the ends of the slit have similar spectrograph vignetting to those at the centre and so the colour of the sky spectrum to be subtracted off any individual fibre is invariant . For redshift work which just requires the identification of features and the determination of their positions, the sky to be subtracted off is relatively featureless because of the digital OH suppression. Finally, the spectrographs do not move with respect to the gravity vector giving great stability.

4. Thermal Background

The number of photons per pixel detected from the sky background is very low so it is vital that CIRPASS contributes significantly less than this in terms of the background due to thermal emission from within the spectrograph. Since the detectors have a dark current of about 0.1 electron/second/pixel and this is below the background from the sky a useful goal is to ensure that the thermal emission is less or equal to this.

To calculate the thermal background detected we have made the following assumptions.

  1. The camera is cooled to 77K. A stop near the grating is also at 77K so the solid angle that the detector sees for thermal emission from outside the camera is determined by the f/1.2 f-ratio of the camera.
  2. In front of the detector, within the camera and at 77K is a blocking filter which has a high transmission factor for wavelengths below 1.8microns but a very low transmission factor (t) for wavelengths above 1.85microns. For wavelengths in between 1.8microns and 1.85microns the transmission falls off linearly. This is a good approximation to commercially available filters operating at these temperatures in a relatively fast beam.
  3. The interior of the spectrograph emits as a blackbody of a specified emisivity=e at a given temperature=T.
  4. The detector becomes totally insensitive to radiation above 2.57microns.

The detected flux for a small wavelength bin (0.01microns in this case) can be calculated from the blackbody equation allowing for the detector QE, the filter transmission and the transmission of the optics. Integrating a series of these values over the range 1.0-2.6microns then gives the total detected flux. Note that the flux emitted from a patch of black-body of size equal to one detector pixel in to a solid angle defined by f/1.2 is the same as the flux landing on a single pixel in the absence of other losses due to the conservation of surface brightness - so calculating the detected flux is equivalent to calculating the flux emitted from a pixel-sized patch of a black-body.

In doing this calculation the main uncertainties are the values for the blocking factor of the filter, t, and the emisivity, e. The following temperatures for the spectrograph were required to make the thermal background equal to 0.1 electrons/sec/pixel.

Table 2: Spectrograph temperature to give 0.1 photons/sec/pixel

Emisivity=0.1 Emisivity=1.0
t=1e-04233.6K 217.1k
t=1e-05236.6K 221.4K

A blocking factor of 1e-04 is quite realistic and it may be possible for the filter manufacturers to achieve 1e-05. It should also be possible to reduce the emisivity of the inside of the spectrograph by making a lot of the surfaces shiny. Furthermore, each pixel in the detector can really only see the optical components of the spectrograph and not the mechanical structure and these will typically have low emisivity. It can be seen from the above table that in order to keep the thermal background to an acceptably low level the temperature of the spectrograph has to be between 217K and 236K or -56C to -37C. Obviously if the spectrograph is colder than 217K the background will be reduced to even lower values. The analysis shows that despite the uncertainties in knowing e and t for the real instrument the uncertainty in the required temperature is only about 20 degrees.

The spectrum of the thermal background for each of the 4 cases in the table is shown in the figures below. It can be seen that the spectrum consists of 2 peaks - one near 1.8microns just before the filter switches on and one near 2.5m where the very strong thermal emission comes through despite the blocking filter. It can be seen that even when the blocking filter has the less demanding specification of e=1e-4 a significant fraction of the background still comes from the 1.8microns peak.

5. Spectrograph Optical Designs

CIRPASS's instrument specifications are defined in the scientific case. This section is concerned with the issues of technical feasibility, practical design options, and compromises inherent to different design forms. From the optical design point of view the most important specifications are:

The three major technical issues that we have addressed so far are schemes that give a simultaneous spectrum coverage of 1-1.8m, schemes that give the required spectral resolution with a practical grating/beamsize combination and the design of fast IR cameras.

We have assumed CIRPASS will be a plug-plate fibre system with a single fibre per galaxy. To feed the light into the fibre efficiently a micro lens is used to image the telescope pupil on to the fibre at f/4. This reduces focal ratio degradation (FRD) in the fibre. A small field stop ahead of the microlens and in the telescope's focal plane defines the 1.5arcsec aperture. For an 8m telescope the pupil image on the fibre has a diameter of 233m.

Given 4, 10242 detectors and a wavelength coverage of 1-1.8m, our starting point was to attempt to produce a design in which the whole J band uses 2048 pixels (2 detector widths) and likewise the whole H band uses the other 2048 pixels. This satisfies the resolution requirement as long as 2 pixels corresponds to a spectral resolution element. To obtain sufficient dispersion (assuming an f/1.2 camera) a 300g/mm grating working in second order with a 120mm beam diameter is required. These parameters have driven the optical designs we have looked at so far.

However, with an f/1.2 camera (which is about as fast as we can achieve) the fibre size for 2 pixel matching is only 0.79" so our 1.5" fibres are 2 times too big. To fix this problem we developed a "switchyard" scheme in which a single fibre efficiently feeds its light into 7, 0.8" fibres and this is described in section 5.2. If the switchyard (which acts like an image slicer) is not used then we have to increase one or more of the beamsize, the grating ruling frequency or the spectral order to give an overall factor of 2. Pushing the design in this direction is difficult and expensive.

5.1 Wavelength Coverage

The exact manner in which information content is distributed between wavelength coverage and multiplex gain is a matter of observing strategy, technical feasibility, and the observer's personal preference. Is simultaneous wavelength coverage from 1-1.8m necessary or is it only a desirable feature? Our initial design efforts were aimed at producing a large wavelength coverage with a correspondingly reduced multiplex gain.

There are 3 ways that we can think of in which a spectrum can be broken down into sections and captured by our four non-buttable detectors. These are:

  1. Using dichroic beam splitters
  2. Using beam-steering mirrors at an intermediate spectrum position.
  3. Using a cross-dispersed echelle.

So far in this study we have mainly considered the first 2 of these.

The illustration below shows how dichroic beam splitters could be used to feed four spectrographs simultaneously. In this manner the light from each galaxy is fed into four different spectrographs and simultaneous wavelength coverage is achieved. In practice, however, dichroic beam splitters are limited by the sharpness of their spectral profiles. In mixed polarization light, the switch from transparency to opacity spans about 0.1m. There is no trouble using beam splitters to separate J from H as there is a very useful water absorption band stretching from 1.35microns to 1.45microns, and that is how the COHSI dichroic is used. But intraband splitting is more messy and problematic.

This leads us to the second method of spectrum redistribution which is to use tilted mirrors at an intermediate spectrum to redirect different portions of the spectrum onto different detectors. This is the principle used by WF/PC on the HST (albeit for an imaging system) and is also used by the proposed AUSTRALIS spectrograph for the VLT. This technique gives a much cleaner cut between wavelength regions than a dichroic can.

In view of the above our initial design form uses a dichroic beam splitter to separate J & H. Each band is then fed into its own spectrograph and the intraband splitting is done with redirecting mirrors at an intermediate focus.


The lens designations are the same as that used in Figure 1.

5.2 The MOS Switchyard Interface

The purpose of this module is twofold: to convert a large fibre diameter in to several smaller ones (to get increased spectral resolution) and to partition the spectrum in to two separate fibre feeds (one each for J and H).

Light from the 250micron core diameter fibre coming from the telescope is fed into 7 more manageable 130micron core diameter fibres which then deliver the light to the spectrograph. This is done by projecting the image of each 250micron fibre onto a cluster of 7 hexagonal arrays, which in turn each feed a 130micron fiber. The exact size of the output fibre and the f-ratio it is fed at depend on the input f-ratio of the spectrograph and it is simple to design the hexagonal lenslets to give the required f-ratio. The layout of the MOS switchyard is shown in Figure 1, and the details of the optical elements are shown in Figures 1a and 1b. The lens complex L1 collimates the f/4 beam from the telescope, while the doublet L2 provides the magnification necessary to project the 250micron fibre onto a 6 mm spot. A small dichroic beam splitter (not shown) is positioned at P to split J & H and feed two spectrographs simultaneously. The field lens, in conjunction with the hexagonal lenslets, is used to properly re-image the pupil onto the fibre face. Figure 1c shows the distribution of the lenslet clusters as well as how a 250micron fibre is imaged onto a lenslet cluster. Each lenslet is 3 mm from corner to opposite corner. The principle of the lenslet construction is borrowed directly from the SPIRAL concept for which the phase-A prototype has been successfully manufactured at the IoA.

5.3 The Baranne Spectrograph

This is the first spectrograph design form we explored as it avoids the anamorphic distortion inherent in a classical spectrograph. Anamorphic distortion prevents us from matching the fiber diameter size to two pixels along the dispersion and spatial direction simultaneously, thus resulting in loss of either multiplex gain, spectral coverage, or spectral resolution. So from an information content viewpoint, this is the preferred design form. It also in principle allows wavelength splitting because it has an intermediate spectrum position.

The layout of the Baranne spectrograph is shown in Figure 2. It is essentially a modified version of the COHSI design. The fibre slit and the spectrum are both situated on the curved focal surface of a Schmidt camera. Unlike COHSI, where we are making use of the water absorption band in between the J & H spectrum to position the fibre slit at the very centre of the Schmidt camera, we are forced to offset the slit to the end of the spectrum, as there is no available gap in the spectrum. (Remember that the J & H split has already occurred in the switchyard). This, in turn, impacts adversely on the optical performance of the system.

The primary mirror, working at f/5, generates a 120 mm diameter beam which is then projected onto the grating (standard Milton Roy 300 g/mm grating used in 2nd order). The spectrum is then imaged at the center of the Schmidt by the primary mirror.

The question raised at this point is how to extract the spectrum from the center of the spectrograph, split it into two portions to be directed onto two independent detectors. Our attempts to do so using relay mirrors, while coming tantalisingly close to fruition, proved unsuccessful. Using flat mirrors to redirect the spectral portions onto different cameras suffers from the problem that the imaging quality inherent to the Schmidt is not adequate. Using curved mirrors and a second pass through the Schmidt to correct the aberrations only led to the introduction of new ones as the symmetry of the system was broken. This symmetry can easily be recovered but is totally unaffordable (essentially consists of cutting the primary mirror into three pieces). We could decide not to split the spectra at the center of the Schmidt, but generate another intermediate image and split the spectra at this additional image. But this entails more optics, lower throughput, and higher cost.

In short, we do not have a viable design for a system with intraband splitting at an intermediate spectrum. Simultaneous spectral coverage over the entire J & H band looks technologically very difficult. We decided at this point to consider having just one detector per spectrograph instead of two.

By dropping the requirement that the intermediate spectrum is big enough to feed two detectors we were able to achieve a working design of a Baranne spectrograph. Because the spectrograph now has only half of the original wavelength coverage, this allows us to position the entrance fiber slit much nearer to the center of the Schmidt camera and obtain adequate image quality. This is the system presented in Figure 2. This system could feed two spectrographs simultaneously by simply putting a second fiber slit, grating, camera and detector rotated 90o about the Z-axis (the axis passing through the center of the aspheric plate and the primary mirror). This helps to significantly reduce the cost as only two Schmidt cameras would be required instead of four.

A spreadsheet showing the estimated system throughput is attached. The peak efficiency is 23%. This spreadsheet has two dichroics in the light path to deploy 4 such spectrographs with maximum wavelength coverage. These dichroics are not needed if we have 4 times as many objects and 4 times less wavelength coverage. The multiplex-gainwavelength-coverage product for this design is ~38m.

5.4 The Classical Spectrograph

The difficulty inherent in the classical spectrograph stems from the contradictory constraints of keeping anamorphism small while keeping the camera design workable. Reducing anamorphism requires that the camera-collimator angle be kept to a minimum. But a smaller angle translates into a greater separation between the grating (pupil) and the first camera lens, therefore requiring larger lenses (which is a major problem with infrared materials). It also does not have an intermediate spectrum which could be used for wavelength splitting.

Our study of fast IR cameras showed that a separation of 300mm between the pupil and the first lens of the system is about as far as we can go. We therefore selected a camera-collimator angle of 30 so as to satisfy this practical constraint and live with the anamorphism introduced.

The layout of the classical spectrograph is presented in Figure 3. The beam size and grating characteristics are the same as those used in the Baranne system. Because there is no relay mirror present in the center of the Schmidt collimator, we are able to position the slit at the center of the Schmidt collimator and work at f/4 (which is more fiber friendly that f/5) while still obtaining optical performances far superior to those achieved by the Baranne Schmidt.

Note also the much smaller size of the primary mirror and aspheric plate of this design as compared to the Baranne system. Since we are using only a portion of the aspheric plate, cutting a single aspheric plate into quarters will provide all four spectrograph with their aspheric corrector.

Figure 3a presents a cut of the system showing how the spectrum is spread onto the detector for one of the grating settings (short-J; 1.025m - 1.20m). The fiber slit spatial direction is into the page. This diagram also shows the details of the f/1.2 camera design. The spot diagrams of this partially optimised design, presented in Figures 3b-d, already shows adequate imaging quality. The optical performance for the other grating settings is of similar quality. The exact same camera is used at all grating settings (i.e. in practice we would purchase four identical lenses from the manufacturer). Table 3 provides detailed information on the characteristics of this design.

This design is of the "normal to collimator" form with the beam diameter reduced in the spectral direction. The fibre slit is matched to two pixels in the spatial direction. There is no loss of multiplex gain, but the higher dispersion leads to a smaller wavelength coverage. If the preference is to sacrifice multiplex gain instead of wavelength coverage, then an Ebert configuration should be selected. Which of these two alternatives is optimal is not yet clear. Oversampling the fibre images in the wavelength direction will help sky subtraction. Note that it is not simply a case of choosing the grating orientation at the telescope - the two alternatives require different beam sizes and therefore different collimators.

A spreadsheet showing the estimated system throughput is attached. The peak efficiency is 24%. This spreadsheet has two dichroics in the light path to deploy 4 such spectrographs with maximum wavelength coverage. These dichroics are not needed if we have 4 times as many objects and 4 times less wavelength coverage. The multiplex-gainwavelength-coverage product for this design is ~34m.

5.5 An Echelle Spectrograph

A cross-dispersed echelle spectrograph with a low-order echelle grating could give the full wavelength coverage of 1.0-1.8m on one detector with a multiplex gain of ~25. In this case it not necessary to put the full wavelength coverage on to 4 chip-widths and so it is possible to have more than 2 pixels per resolution element and consequently the switchyard is not needed for either wavelength or spatial partitioning. For example an existing 75g/mm grating with a blaze angle of 26.7 could be used to give 0.926-1.813m in 7th to 12th order with a 3 pixel spectral resolution of R~2560. This would have a fibre diameter of 1.2 arcsec on an 8m and a multiplex-gainwavelength-coverage product of 88.7m. The high value here is mainly due to the fibre diameter and spectral resolution being below specification but this is clearly a powerful option. Drawbacks of the echelle approach are scattered light and the lower grating efficiency off blaze. Clearly this approach has to be investigated further.

5.6 Using the Spectrographs with an IFU

Most applications for integral field spectroscopy require field for spatially mapping more wavelength coverage. The fibres used are also small in terms of their aperture on the sky compared to the 1.5" we require for MOS. These two points mean that the integral field mode is not really driving the spectrograph design and can therefore be fitted to the spectrographs in a straightforward manner.

6. Conclusions

Assuming that CIRPASS will be used on a telescope with an aperture of 8m, the main conclusions of our study so far and recommendations for further work are:

  1. A multi-fibre system is prefered to a multislit system on the grounds of cost, field of view and speed of operation.
  2. Inadequate precision of sky-subtraction is the major concern with a fibre based system. The CIRPASS project must take the steps necessary to ensure that the spectrograph's sensitivity is limited by Poisson statistics and not systematic errors introduced by sky-subtraction.
  3. To be sky-limited (rather than limited by the thermal background of the spectrograph) it will be necessary to cool the entire spectrograph to ~-50oC.
  4. Cameras with a speed of f/1.2 are feasible but difficult.
  5. Recording the whole wavelength range 1.0-1.8m simultaneously is very difficult and may compromise performance in other ways. If the 0.85-1.0m region is also required then complete coverage in one shot becomes even harder. Achieving the multiplex-gainwavelength-coverage product specification is much easier to achieve by having a large multiplex-gain.
  6. Achieving the required spectral resolution with a 1.5" fibre requires a fibre reformatting device (to reduce slit width) or a large beamsize or a high order grating. Schemes which do not require the switchyard should be further investigated.
  7. A peak throughput of ~20-24% is achievable.
  8. A Baranne spectrograph design with 0.19m of wavelength coverage and a multiplex gain of 200 (50 per spectrograph) looks feasible technically but seems to be ruled out on cost. This design is also likely to suffer more from scattered light problems.
  9. A fully specified classical spectrograph with a wavelength coverage of 0.17microns and a multiplex gain of 200 (50 per spectrograph) has been designed and appears to be affordable.
  10. An echelle approach needs to be investigated in detail.
  11. A reduction in the specification for the fibre diameter has to be considered further because a small reduction in this parameter will help considerably with the design of the spectrograph.
  12. For a classical spectrograph approach the normal to collimator and normal to camera alternatives have to be properly compared and evaluated.