The Institute of Astronomy (IoA) is currently building a panoramic wide field near infrared imaging camera based on 4 Rockwell HgCdTe 10242 detectors. The survey instrument will be as scientifically versatile and as easy to use as a large format CCD camera and is expected to be ready for astronomical use by Nov 1997. It will be particularly well-suited for surveys of star-forming regions, low mass stars, distant galaxies, clusters and QSOs The field of view of this camera with 0.15'' pixels is 5.1x5.1 arc minutes and is thus 60 times larger the current near-infrared imager on Keck(NIRC). When combined with the MMT 6.5m, the combination is 25 times as powerful as the Keck 10.0m, when the apertures are taken into account. The instrument design includes the option of adding wavefront sensing and a tip-tilt facilities at a later date. In addition an option of upgrading the camera into a wide field spectroscopic survey instrument is currently at the design stage.
The Institute of Astronomy(IoA) instrumentation group is currently
building a panoramic infrared imaging camera. The device, which
will consist of a mosaiced array of 4 times HgCdTe 10242
detectors, will be used to carry out deep imaging surveys in the
near-infrared spectral region. The survey instrument will be as
scientifically versatile as a large format CCD camera. It will
be particularly well-suited for surveys of star-forming regions,
distant galaxies, clusters and QSOs and is clearly well-matched
to current IoA scientific interests. It is intended that the camera
be ready for astronomical use by Quarter III, 1997.
A NIR f/3.9 wide field corrector has been designed for the Cassegrain
focus of the 6.5m MMT, giving a scale of 0.15''/pixel and a field
of view of 5.1'x5.1'. The design includes the option of adding
wavefront sensing and a tip-tilt facilities at a later date, based
on proven techniques developed for the Cambridge Optical Aperture
Synthesis Telescope (COAST) by members of the IoA team (i.e. Beckett
The program includes the option of upgrading the camera at a later date to a wide field spectroscopic survey instrument equivalent to an infrared version of the highly-successful Low Dispersion Survey Spectrograph. This multi-object NIR spectrograph would enable powerful cosmological applications particularly on an large-aperture telescope. The spectroscopic upgrade is currently at the design stage with a variety of options including fibre fed capabilities being considered.
The Institute of Astronomy has a long tradition in survey astronomy
and the scientific value of the various galactic and extra-galactic
surveys conducted in Cambridge is well-documented. Examples include
various quasars surveys and the Southern Sky galaxy catalogue
constructed using the Automatic Plate Measuring machine(APM) (designed
and operated within the Institute), the ground-based follow-up
of X-ray catalogues defining the range of physical properties
of clusters of galaxies, and numerous redshift surveys of the
local Universe using both optical and infrared-selected samples.
These and other ongoing surveys serve to define the population
of galaxies, clusters and related objects at various epochs.
Wide field broad band sky surveys represent an important observational
tool in astronomy. They can be used to select objects in a quantitative
manner for statistical studies or subsequent investigations with
a large telescope. They are also useful in the follow-up of sources
detected at other frequencies, e.g. radio, Infra-red and X-ray
This programme is concerned with extending the Institute's survey
imaging capability to near-infrared wavelengths. The impact of
large format CCDs is clear from the substantial efforts many observatories
have invested to secure such devices. For the first time, a similar
technological capability is now available at infrared wavelengths
via the first 10242 HgCdTe arrays. The mosaic camera
discussed here consists of 4 times 10242 detectors
which will fill a significant fraction of the unvignetted focal
plane on existing 2-4m class telescopes. Such a device will represent
a major leap (50-100) in our survey capability at infrared wavelengths.
Indeed, a similar jump in capability at modest cost is unlikely
to occur again.
The near-infrared wavebands are increasingly important for both
galactic and extra-galactic surveys. Distant galaxies are optimally
studied at these wavelengths because the k-correction is smaller
and less uncertain. For QSOs, a significant increase in redshift
coverage (i.e. z>6) requires NIR observations since cosmologically
distributed HI absorbs away most of the optical waveband. Further
physical drivers include minimising dust obscuration at low Galactic
latitudes as well as searching for cool sources with temperatures
A key program could be the establishment of deep medium angle
survey fields which would be the focus of survey programs into
the next millennium on the 8m's and future space missions such
as AXAF, XMM, SIRTF etc.
The range of scientific programs that can be attempted with a
wide field infrared imaging camera is substantial. Below is listed
a representative sample.
Rather than provide detailed scientific and technical cases for
the types of program listed earlier we provide some basic information
and expand a few cases. In Table 2.2.1a we show a comparison between
the Cambridge camera and a variety of other Camera + Telescope
combinations normalised against the Keck 10m with NIRC. The comparison
is quite illuminating and illustrates the power wide field capabilities.
The 2.5m INT telescopes appears quite favourable but the poor
spatial sampling and long integration times are the major limitation.
It is more suited to a shallow survey to H20 as indicated by Figure
2.2.1a . Table 2.2.1b summarises the expected sensitivity of the
Figure 2.2.1b shows the expected magnitudes for an L* unevolving
galaxy as a function of redshift and spectral type. At high redshift
there are a number of uncertainties. Evolution generally causes
some brightening whereas merging results in fainter individual
discrete galaxies at higher redshift. The three horizontal lines
show the expected depth for 3 exposure telescope options
i.e. INT 2.5m + 5mins; INT 2.5m + 30mins; MMT 6.5m + 30mins.
One important point is that for an non-evolving population there
is no significant difference between the H and K detectabilities.
The K band figures assume a fully cryogenic shielded camera. Figure
2.2.1c shows a optical-NIR two colour plot to show the wide range
of expected colours. On this diagram, a simple V-H>6 sample
isolates ellipticals or a bulge dominated population with z>1.
By observing in multiple bands in the optical and NIR the derivation
of redshift, spectral-types and even population ages should be
In Figure 2.2.1d is shown the spectral energy distribution of
the recently discovered brown dwarf candidate Gl 229B. Note the
rather blue spectrum of this star longward of 1m. The actual J-H,
H-K colours are -0.1 and 0.1 respectively unlike normal M stars
which are much redder. As a result there is no advantage in using
K in surveys for such objects. A simple I-J colour selection or
Z-J is required.
The evolutionary history of galaxy clusters offers a sensitive
test of current theories for the origin of structure (White et
al 1994). The required data consists of cluster surface densities
as a function of contrast and preferably redshift (Couch et
al 1990, Postman et al 1996). Such samples also provide
a natural source for distant galaxy studies (Dickinson 1996).
Unfortunately, the major difficulty lies in finding a robust way
of finding distant examples. Optical searches have found some
examples beyond z0.5 but their contrast against the very high
field counts renders the method unreliable. Similarly, the X-ray
luminosity of the richest examples is surprisingly low (Castander
e tal 1993, 1995). Yet it is known that structures do exist
at high redshift. Dickinson (1995) finds galaxy associations around
radio galaxies at z>1 whose contrast factor to modest limits
nonetheless implies quite significant richness. Similar systems
are being identified in the deepest redshift surveys (LeFevre
et al 1995).
The slow evolution of cluster ellipticals identified from a variety
of studies (Aragon-Salamanca et al 1994, Ellis 1995) indicates
that the contrast of a z>1 cluster galaxy against the foreground
field is considerably greater at near-infrared wavelengths than
for the current optical samples. Panoramic infrared surveys for
such associations have, to date, been considered prohibitively
difficult because of the limited field of view.
Simulations indicate that clusters of Abell richness R>3 can
be located with adequate contrast to z2 in surveys reaching H=21.5.
Even if their volume density is an order of magnitude lower than
locally, useful samples can be constructed in a survey of a few
A more ambitious program would be a search for QSOs with redshifts
in excess of 7 where Lyman- lies beyond 1m. Even a firm upper
limit would be of great cosmological interest. Ten years ago,
there were no QSOs known beyond z=3.8 (Hazard, McMahon and Sargent
1986). QSO surveys undertaken with the APM have transformed the
situation and now around 60 z>4 QSOs are known. The majority
of these z>4 quasars have been found using multi-colour surveys
with the APM (Irwin, McMahon & Hazard 1991, Warren, Hewett
& Osmer 1993). These z>4 QSOs have primarily been selected
on the basis of red B-R colours e.g. B-R>3. This approach starts
to fail at z4.5 since the R band samples the Lyman- forest absorption.
Various optical surveys for QSOs in the redshift range 5 to 6
are currently under way using the same basic principle of red
optical colours The principle behind the B-R techniques can be
extended to even higher redshifts i.e. z>7 using R-J,J-H colours.
The highest redshift QSO known PC1247+34(z=4.9) (Schneider, Schmidt
& Gunn 1991) has J,H and K mags of 18.0, 17.6 and 17.0 (McMahon
et al, in prep) respectively, whereas the brightest known
QSO above z=4.5, BR1202-0725 (z=4.7), has J, H and K mags of 16.9,
16.2 and 15.7 (McMahon et al, in prep) respectively. The
QSO PC1247+34 was discovered in a survey covering 3 deg2
so a survey over an area of 20 deg2, 2-3 magnitudes
fainter i.e. H=20 would be in the right ball park. At these redshift
objects QSOs dim as (1+z)3 so going from z=5 to z=10
produces 2 magnitudes of dimming in an unabsorbed continuum band.
The uncertainties in the expected numbers of QSOs is large. The
aim of any survey program would to determine the space density
or set the best limits possible. Ideally one would carry out both
shallow and a deeper survey over a smaller area so that one constrains
the shape of the QSO luminosity function. A primary motivation
for high redshift QSO survey work is the use of such QSOs as probes
of intervening absorption systems.
The forthcoming HST refurbishment mission scheduled for early
1997 will include the installation of the near infrared red camera
NICMOS. It will be timely to redefine the area of wide field near
IR astronomy on the same timescale that NICMOS opens up a new
window in the regime of small field but ultra-deep NIR astronomy.
NICMOS is built around 256x256 arrays and has a field of view
of a mere 19''x19'' with 0.075'' sampling. A ground based deep
wide field NIR survey program should generate ideal targets for
HST follow-up with NICMOS.
One exciting possibility would be to devote a substantial amount
of time to a NIR equivalent equivalent of the HDF. For example,
in 100hrs one on the MMT 6.5m one would reach H=26 (equivalent
to K=25 for a typical galaxy with z>1) where most of the detected
objects would be at z>2. In a single 5'x5' field one would
expect over 1000 galaxies. One caveat is that HST+ NICMOS may
be more efficient due the lower sky background although it would
require many NICMOS pointings. However, there are still considerable
uncertainties about the performance of NICMOS in the background
limited regime e.g. the in-orbit read out noise, dark current
and thermal background are unknown. We should know on the timescale
of a year whether it is practical to use NICMOS for ultra deep
survey wide field survey programs. If it is not we shall be poised
to take the lead in this important area.
The camera will be remarkably simple and compact (see Figure 3.1a)
It is an extension of current IoA capabilities, e.g. Martin Beckett
and Craig Mackay's development of a NICMOS HgCdTe 2562
device for the IoA/MRAO collaborative programme on the Cambridge-Optical-Aperture-Synthesis-Telescope
and does not involve significant technological risk. The prime
difficulty in moving in this direction has been the prohibitive
cost of the new large arrays. The recent fund-raising successes
at Cambridge provides an excellent and timely opportunity to move
ahead. The system under construction will represent a major leap
(50) in our survey capability at infrared wavelengths. A similar
jump in capability at modest cost is unlikely to ever occur again.
The camera has a core detector system which consists of 4 Rockwell
1024x1024 HgCdTe arrays. The basic characteristics of these arrays
are given in Table 3.1a with comparative figures for the NICMOS-3
devices. Unlike optical CCDs, it is not possible to closely pack
the infrared arrays (see Figure 3.1b) and thus a minimum of four
exposures is needed to give a 100% fill-factor (at additional
cost a contiguous field could be provided using a pyramidal system
similar to WFPC2). An advantage of the of the large gaps between
the detectors, it will be possible to use small, relatively cheap,
filters since we benefit from a multiple order. This will be important
for narrow band programs.
The camera system includes a fully integrated auto-guider system
(see Figure 3.1c) A small pick-off mirror and some small lenses
inside the dewar will be used to send the light from an unused
part of the field to a Peltier-cooled CCD on the side of the dewar.
The control computer for the CCD will lock on to the brightest
star in the field and generate XY correction signals and send
them down an RS232 line to the telescope control system. The CCD
system will also allow the option of dithering exposures to reduce
flat field noise and the effects of dead pixels. The CCD field
will be about 4x6 arcmin if the re-imaging lens produces a final
image scale of 0.25 arcsec/pixel.
Figure 3.1d shows a schematic of the detector control system with
associated computer system. Each detector has 4 quadrants which
can be read out independently. A full exposure of the 4 chips
requires the capture of 16Mb of data. In the H band the sky is
very bright and this dictates how fast the mosaic has to be read
out. Assuming a well capacity of 60,000, each pixel saturates
in 15 seconds for the INT, WHT or WHT which is within the specification
of the AstroCam 4100 controller. We propose to commission the
new camera initially on the 2.5m Isaac Newton Telescope or the
4.2m William Herschel Telescope but, thereafter it can be transported
to telescopes elsewhere. The overall characteristics and estimated
performance on the 2.5m INT, 4.2m WHT and 6.5m MMT (as a guide)
is given in Table 3.1b. As part of the work done for this proposal
a fully-transmitting field corrector was designed in detail for
the Cassegrain focus of the WHT and so we are confident that successful
designs for other telescopes can be achieved. In contrast to the
INT these typically offer higher resolution with a smaller field
We have also investigated the feasibility of building a fully
functional K band system. However, the problems associated with
building a thermally isolated and well shielded wide field camera
system are prohibitively expensive and it is not clear that the
scientific gains warrant this expense.
This is still undergoing definition. It will be a Tk/TCL graphical
User Interface (GUI) with user defined scripts. It is envisaged
that a series of dithered and interlaced exposures should take
place with little user intervention. A typical cycle of 5 dithers
and 4 interlaced raster positions should take one hour. If guide
star's are necessary it is envisaged that no user intervention
will be required whilst moving between raster positions. Single
person operation is intended. Data will be written to disk as
either IRAF or FITS format files.
Each exposure will produce 4 images at 2MBytes each. A typical
real-time co-add consisting of 10 times 10sec exposures will result
in a single 2MByte image. Therefore the maximum data rate is 8MBytes
in 100seconds i.e. 288MBytes/hour. Therefore in a typical '10hr
night' the data volume is 3GBytes. This will fit onto a DDS-2
DAT(4GB capacity). The 'raw' data which will not usually be archived
would amount to 30GBytes.
The members of the IoA have considerable experience is the development
of pipelines to process optical imaging data e.g. the APM and
NAOJ CCD mosaic project. The proposed pipeline will produce astrometric
calibrated images and linearised catalogues of detected objects
within an IRAF based environment.
The current design includes a 8 slot filter wheel so that a wide range of filters can be used. A summary of the filters that one might use are given in Table 3.4
In addition to a set of standard J and H filters a broad band
Z filter and a short pass K filter would be useful. In the Z band
(0.85 to 1.0 m), conventional CCDs have effective QEs of 20% at
0.85 microns, and very little sensitivity at 1.0 microns whereas
the MCT array has a QE of 60%. Therefore Z band may prove useful.
The trade off is that conventional CCD Mosaics can cover a larger
field with the same sampling. However, contemporaneous observations
in a band that overlaps with CCDs may be useful. A short pass
K band filter will also be considered. Its characteristics will
depend crucially on the thermal background and transmission of
the telescope optics.
Narrow band filters choice fall broadly into two categories;
(ii) filters that are defined by redshifted emission lines.
In the case of redshifted lines, one should choose windows in
the OH in order to minimise the sky contribution e.g. at 1.25m
and 1.66 - 1.25m. The choice of optimum windows are currently
In addition filters that correspond to 'matched' emission lines would be required for blank sky surveys. e.g. a filter at 1.25 m could be used to search for H- at a redshift of 0.9 in order to determine the global star formation rate. This filter would also be useful in eliminating low redshift interlopers from a search for z5 Lyman- emitting galaxies carried out in the optical at 7100Å. Some of this survey work is clearly more ideally suited to 8m class telescopes rather than 2m class telescopes.
The camera can be used directly at the prime focus of the 2.5
Isaac Newton Telescope(INT) using the existing triplet corrector
which gives excellent performance in the J (1.3m) and H (1.6m)
bands. The timely upgrade to the INT prime focus assembly and
telescope control systems to facilitate a 2x2 thinned Loral CCD
mosaic means that integration with the IR camera will be quite
easy. The advantage of the INT prime is its very large field of
view and well matched pixel scale. Each pixel corresponds to 0.46
arcsec and the total field of view is 0.07 deg2. A
sequence of interleaved images will give a final image of 31.4x31.4
arcmin. This in a total of 16 exposures an area of 1 degree square
can be covered.
There are no existing or planned IR instruments for the INT and
we expect to be able to obtain 2 to 3 weeks per semester. Within
the UK context, we effectively gain an IR telescope to complement
our access to UKIRT which will be primarily more useful in the
K band and in cases where small field and high image quality is
the main driver. In the short term one would expect to follow
up objects discovered during the INT phase using UKIRT, potentially
pushing COHSI to its limits.
It is not clear that moving the camera to the 4.3m William Herschel Telescope(WHT) has any substantial gains over the INT. The advantage of the WHT is the larger aperture and better image scale which would help in deeper surveys. However the WHT prime focus corrector has a lower throughput in H than the INT by 50% and bright time is more oversubscribed on the WHT compared with the INT since its bright time instrument set includes an echelle spectrograph and a J,H and K adaptive optics system built around a 2562 InSb array. Ideally one would work at the Cassegrain focus of the WHT using a special purpose NIR corrector. This option is currently under investigation but it is likely that for comparable cost one could take the camera to a larger telescope.
Many of the programs in the list in Section 2 would benefit considerably
from larger apertures and finer image scales. In addition if some
of the more speculative survey programs are successful one would
expect that deeper programs would be a logical next step to either
probe deeper in flux at the same redshift or to push to higher
redshifts. In addition multi-objects near IR spectroscopy is clearly
better suited to the larger aperture.
The basic camera system should be easily transportable to the
MMT environment. On a longer term, spectroscopic capability and
contiguous field possibly via a pyramidal mirror system similar
to the HST Wide Field Camera with the inclusion of tip-tilt adaptive
optics is clearly possible. Contiguous field would be a significant
gain for programs that want to go deep over smaller areas since
the 2x2 interleaving effectively drives one to larger fields than
one can follow-up easily with multi-object spectrographs.
Fabricant & McLeod ('Optical Specifications for the MMT conversion',1995)
describe an f/5.48 wide field corrector optimised in the optical
over the wavelength range 3000--10000Å. We have carried
out some preliminary optical design work on a NIR wide field triplet
corrector for use with the MMT's f/5 secondary to give good images
in the wavelength range 1.0-1.8m over a field of view corresponding
to our mosaic. The triplet lens has some power of its own to increase
the field of view on the sky and gives a pixel scale with good
sampling in excellent seeing. The final f-ratio is f/3.9 and the
image scale is 123m/arcsec or 0.15 arcsec/pixel. This should be
compared with the bare f/5.25 scale of 0.11 arcsec/pixel and represents
a factor of 2 increase in field of view.
The triplet shown in Figure 4.3a uses PK51A and BALF5 glasses
to provide correction for chromatic aberrations. The largest element
has a clear aperture of 228mm. The back-focal distance is 400mm
to allow the possibility of including a tip-tilt mirror. The focal
plane is 111mm below the instrument mounting surface.
The image quality is shown in the spot diagrams (Figure 4.3b) and encircled energy plot (Figure 4.3c) for 4 field positions that correspond to the corner of an array near the centre of the mosaic, the corner of the mosaic and 2 other points equi-spaced in between. The imaging performance is not quite diffraction limited but 90% of the energy falls within a circle of 20 radius. The box size for the spot plots is 37m which is 2x2 pixels.
The wide field camera system under construction offers a further
great opportunity if it can be upgraded to a very wide field deep
spectroscopy system for a 6-8m class telescope.
Over the past 5 years, impressive progress has been made in extending
redshift surveys to higher redshift providing information on the
evolutionary behaviour of galaxies, clusters and quasars. Much
of the progress has resulted from surveys made possible via purpose-built
instrumentation, e.g. multiple object spectrographs capable of
measuring redshifts for many hundreds of faint sources or wide
field imaging capabilities.
The cosmic landmark represented by a redshift of one has been
a significant barrier for galaxy and cluster work for many years.
The apparent magnitude of a typical galaxy at this distance is
at the spectroscopic limit of a 4-m telescope. Furthermore, the
optical contrast of a rich cluster against the projected field
population also falls below a critical detection threshold (2-3).
Most significantly, the spectral features used so successfully
to determine the redshifts of faint galaxies, visible emission
lines of [O II], [O III] and H-, are shifted into the near infrared
and thus optical spectrographs are not well-equipped for measuring
redshifts beyond 1.
Systematic redshifts surveys with 4-m telescopes have succeeded
in determining the spatial distribution of nearby galaxies and
the evolutionary behaviour of various classes of galaxies out
to redshifts of 1 (Glazebrook et al 1995, Lilly et al
1995, Ellis et al 1996). The recent discovery of an abundant
population of star forming galaxies beyond a redshift 1
from faint object Keck spectroscopy (Cowie et al 1995)
and, at higher redshift from various lensing studies (Lewis et
al 1995, Ebbels et al 1996) and Lyman-limit-selected
surveys (Steidel et al 1996) raises the question of how
best to search systematically for such objects.
The scientific motivation for extending the earlier studies to
higher redshift follows from the implications of the strength
of emission seen in the important diagnostic lines of [O II],
[O III] and H-. The limited data currently available indicates
a significant fraction of the present stellar mass was being created
in the redshift range z>1. This may be consistent with the
relatively mild evolution seen at later times but places in question
the significance of the abundant dwarf population. Only by tracking
the mean star-formation rate of well-controlled samples as a function
of look-back time can these various issues be resolved.
In addition to quantifying the era during which galaxies formed
the bulk of their stars, the spatial distribution of star-forming
galaxies at early epochs might be derived if the technical challenge
of surveying 5-10,000 faint galaxies in the redshift range 1<z<3
could be overcome. The growth of clustering on various scales
can be used to constrain models for the evolution of structure
which are themselves indicators of the cosmological parameters
once the spatial distributions at the present-day and epoch of
recombination are well-defined as can be expected in the next
As with the earlier redshift surveys, the key to progress lies
in the development of the appropriate instrumentation. The Lyman-limit
selection pioneered by Pettini and collaborators remains a promising
method for the highest redshift systems. However, follow-up spectroscopy
with the largest telescopes is still required to derive precise
redshifts and star formation diagnostics. One important limitation
with the method is that it cannot be used below a redshift zLy
2.8 without access to space-borne UV facilities. As these currently
have very narrow fields of view, they are not well-suited to survey
Direct spectroscopy of broad-band selected targets such as is
now being conducted on the Keck by Cowie and co-workers is also
limited by telescope aperture and the availability of emission
line features at optical wavelengths. The success of the sky-limited
redshift surveys relies largely on the availability of strong
[O II] which disappears from the optical region beyond z1.4. Indeed,
without the emission lines, none of Cowie's galaxies would probably
have yielded redshifts when examined with Keck's LRIS. A logical
solution to this problem is to build a panoramic wide-field spectrograph.
If the primary goal is to determine the star-formation history
of normal galaxies and it is assumed (as now seems reasonable)
that suitable systems are present in abundance at high redshift,
then rather than attempting to measure the redshifts of galaxies
selected first using broad-band techniques, it may also be effective
to search within narrow spectral bands for objects containing
emission lines within selected redshift ranges where the OH contamination
is low. As well as being technically advantageous in lowering
the line flux limit, this would offer a number of strategic advantages
in the derivation of evolutionary conclusions and spatial correlation
functions. Spanning 2-D space at a given redshift does not replace
the continuing need to probe 3-D in single fields but does illustrate
the versatility possible with wide-field instrumentation.
Some earlier narrow-band programmes failed to locate star-forming
proto-galaxies at high redshift; these used both optical (Djorgovski
et al 1995) and infrared techniques (Thompson et al
1995, Pahre & Djorgovski 1995). The apparent failure of these
programmes is puzzling for two reasons. Firstly, the surface density
of z>3 galaxies with high star formation rates identified by
Steidel et al(1995) strongly suggests that obscuring dust
does not seriously limit the detectability of high z objects as
had been argued. It also confirms suspicions that Lyman- emission
is not a reliable indicator of star formation. These developments
may explain the null Lyman- results. However, recently Hu &
McMahon(Nature in press) have detected Lyman- emitters in the
vicinity of QSOs with z=4.44. So far as the near-infrared results
are concerned, the emission line fluxes directly witnessed by
Cowie et al already lie well above the lower limits of
null detection claimed by Pahre & Djorgovski despite the
fact that the redshift ranges of both observational programmes
overlap. Thus the new Keck spectra give much more promising
indications of the merits of surveying for emission line galaxies
in the near-infrared spectral window.
How many star-forming galaxies might a particular survey find?
Part of the motivation is, of course, to determine the answer.
Nonetheless some estimates can be made on the basis of the early
Keck results. The key advantage of a spectrograph that can exploit
4 10242 arrays is its enormous field of view of 72
arcmin2. It potentially can operate in both conventional
imaging or multi-slit spectroscopic mode or in a 'narrow band
selector mode' whose redshift range, z, can be moved in z. The
advantage of the latter mode is that it can offer a much fainter
line flux limit for various technical reasons. This advantage
is offset by an obviously much smaller volume sampled at a chosen
redshift of interest. The key to which mode to use will depend
on the performance and the surface density of interesting objects.
For the conventional mode where broad-band images are used for
multi-slit work within a given redshift range, we can make some
estimates from Steidel et al's data. They find a surface
density of 0.4 arcmin-2 for R<25 galaxies within
3<z<3.5 giving 30 per instrument field. Clearly this is
a lower limit to the number of distant galaxies amenable for study.
The total number of galaxies with R<25 within the field is
1000 (Metcalfe et al 1995). The mean star formation rate
Steidel et al infer from their UV continua is 4-25MO
The contract with Rockwell calls for devices with the following
The 1024X1024 arrays shall be manufactured on a best effort
basis to meet the following performance specifications and goals.
Format (pixels) 1024X1024 1024X1024
Pixel pitch (ÿm) 18.5 18.5
Fill Factor > 95 > 95
Quantum Efficiency (77K)
@1.2 micron 40% 70%
@ 2.35 micron (or peak) 50% 75%
Long wavelength cut-off (micron) 2.5 2.5
Short wavelength cut-on (micron) 0.85 0.85
Read noise (e- at 77K) < 20 < 5
Dark current (e-/s at 77K Vb= 0.5V) < 0.5 < 0.1
Well capacity (e-, Vb= 0.5V) 6E4 6E4
Yield (working pixels) > 92% > 97%
Temperature <77K <77K
The specifications above apply to mean values. There will be four
This camera has relatively simple cryogenic requirements compared to many infrared cameras. The Mercury-Cadmium Telluride (MCT) devices are only sensitive to the J (1.3 m),H (1.6 m) and K (2.2 m) bands and so the thermal background of the instrument has a much smaller effect than with detectors that have a much longer wavelength sensitivity such as InSb devices). The devices have a very low intrinsic dark current at a convenient operating temperature of 77K allowing simple liquid nitrogen cooling. CIRSI is only intended to operate in the H and J bands, although the internal design of the dewar must minimise any radiation out to the long wavelength cut-off of around 2.6 m.
When used in survey mode the camera must accept a f/3.5 beam from the telescope but the lack of any re-imaging optics means that the detector sees a f/1.4 beam and so detects a large thermal background from the telescope. In the H band the camera will detect approximately 1000 e/s/pixel from this background, this is small compared to the OH emission from the sky which may be up to 4000 e/s/pixel.
The filter wheel must be cooled because the HAWAII device is sensitive to infrared radiation up to a wavelength of 2.6 m, although it will not be operated at this wavelength in survey mode. The restricted wavelength response of the MCT devices is a great advantage as the filter doesn't have to provide high blocking at long wavelength. Near infrared filters for InSb based cameras require extra blocking layers to prevent red-leaks which increase the size and cost of the filters and reduce the peak transmission. The specification of the filters ordered is in Appendix 5.2.2.
The camera will be contained in a simple liquid nitrogen cooled dewar. Internal optics and mechanical components have been minimised to simplify the design. A test dewar is being constructed to operate a single chip, this uses an existing Oxford Instruments MN1815 dewar of the type used for CCD cameras by AstroCam and LPO. The internal structure of the final camera dewar will be very similar to the test dewar, a single cylindrical dewar with a single large filter wheel and no other internal optics apart from the fixed autoguider take-off mirror.
The detector mount emphasis the modular design of the system.
The infrared array is fitted in a socket on a circuit board. The
detector mount surrounds this circuit board and provides a cold
finger, a clamp to hold the detector in place and the socket for
the electrical connections. The mount allows 2 edge butting of
the device constrained by the size of the socket. This mount can
hold either a HAWAII or NICMOS device and is can be used in the
test dewar, main camera dewar and COHSI cryogenic spectrograph,
it is intended that any future instruments will use the same mount
design. The detector mount package can be safely removed from
the dewar and stored with the device installed. The detector mounts
consist of a single fixed cold finger supporting the back of the
detector and a cooled clamp on the front which holds the chip
firmly onto the cold finger. The clamp maintains a constant contact
pressure onto the cold finger and prevents the chip moving as
the dewar orientation changes. The tension in the clamp can be
The four detector chips must be position at a common focal plane. The small size of the pixels and the fast f/ratio of the telescope mean that the accuracy of this positioning is a major challenge. The mount contains no built-in movements to adjust the position of the detector surface, this would produce an impossibly complex design and would inevitably reduce the cooling power available. Instead the four detector units will be assembled and the height and tilt of the individual chips measured with a travelling microscope. The mounts will then be individually polished to produce a common focal surface, note that the actual height is un-important as the whole camera assembly can be moved by the telescope focus. The major variation in the position of each chip comes from the layer of epoxy glue used to fix the hybrid detector assembly into the chip carrier ceramic. Appendix 5.3.1 shows the electrical and mechanical design of the detector mount.
The filter wheel is inside the dewar so the available range of
filters can be changed only after warming and partially dismantling
the dewar. Because of this a large number of filters are available
in the wheel. Initially the system will house up to 8 filter positions
although one may be reserved as a blank position to allow testing
of the camera in a no-signal state. The design allows for an upgrade
with a second filter wheel to fit in front of the first, giving
a total of 14 positions.
The HAWAII chip package prevents close packing of the arrays and
the camera is designed with a large gap between adjacent chips.
This means that it is possible to use four small filters, each
32mm square, rather than a single large element. A cell of four
filters can be removed from the filter wheel as a group allowing
easy changes of the pattern of loaded filters while reducing the
handling risk of individual pieces of glass.
The filter wheel is driven by an external motor connected by a
vacuum feed-through. We are using a Ferro-magnetic fluid sealed
feed-through. The position of the wheel is measured by a potentiometer
on the motor gearbox which allows a measurement of the absolute
position of the wheel to within a degree. Although it is necessary
to open the dewar to install filters the camera is designed to
make access easy. When the radiation shields have been removed
the entire filter wheel can be removed or individual filter cells
swapped without disturbing the detector mounts.
The small filters are not only cheaper and easier to produce but allow the same filter sets to be used in a range of instruments. This also means that special filters can be produced easily as only one small filter is needed, it is even possible to have different filters can be mounted in front of each of the four chips.
The field of view of the camera should ideally be only the sky and cold low emmissivity shielding inside the dewar. But operating at a fast focal ratio with a wide field of view the efficiency of the shielding is compromised to avoid vignetting the image. See the section on background for technical details.
The camera will have its own acquisition and guiding system. This
is needed at the INT prime focus as there is no other autoguider
The star counts that we expect to encounter in practice have been
examined so that we can be sure that the search areas for our
autoguider are satisfactory. These are given in Appendix 5.3.4.
The autoguider will use an Peltier cooled CCD camera mounted outside the dewar viewing a fixed mirror mounted inside the dewar near the focal plane. A re-imaging system will produce an image on the camera at the correct scale. The pick-off mirror can be large enough to fill the autoguider CCD re-imaged with a scale of 2:1, i.e. 0.45 arcsec/pixel.
The detector mount module contains a 25 way micro-D connector carrying all the electrical connections to a single chip. These are then connected to two 55 way Amphenol connectors mounted in the dewar wall. The internal dewar connections are made with 36 AWG Manganin wire with Formvar insulation ( from Lakeshore Cryonics). The Manganin wire has a thermal conductivity only 2% that of pure Copper and this arrangement gives a heat load of 70W / connection. Even with the 100 connections needed for the camera the heat load due to the wire is much smaller than the radiation load.
The dewar design is not yet finalised and so only approximate
estimates of the hold time can be made. By simply scaling the
test dewar to the size needed for the final camera and assuming
that the construction is the same, we have a capacity of 12 litres
and a hold time of over 60 hours. This is easily enough for the
camera's normal operation but means that some extra system is
needed to allow the nitrogen to be poured out and the dewar warmed
to allow changes of internal components in a reasonable time.
The camera is likely to be a single cylindrical dewar of 0.5 metre
diameter and less than 0.5m long with a mass of 50 Kg. The associated
controller electronics is a small case 30x25x20 cm and is accompanied
by a normal desktop PC. The system could be contained in small
flight cases and transported as checked-baggage on a scheduled
flight. The mechanical telescope interface is much larger. Using
the same design as the RGO CCD mosaic for the INT would require
a metal disc around 1m in diameter. It may be possible to reduce
the size and weight of the interface at the expense of greater
complexity by using a construction framework system such as the
ALUSET used for the COHSI shell. The telescope interface is obviously
different for each telescope and it is likely that on the first
run at a particular observatory the interface and the camera would
be shipped out in advance along with a few of the group to supervise
the integration. On subsequent runs the camera alone would be
taken, it would then be convenient if the system could be carried
along with the observers. Depending on the number of telescopes
used it may be better to produce an adaptable telescope interface
framework which can fit a number of telescopes.
The Rockwell HAWAII array, like the earlier NICMOS array, is arranged
as four identical separate quadrants produced on the same integrated
circuit. Each quadrant has independent power supplies, addressing
functions and output amplifiers. The CIRSI camera will use four
arrays each with four outputs but with a single camera controller
and signal processing chain. An output multiplexor circuit mounted
immediately outside the dewar allows addressing functions to be
routed to any quadrant of any chip and connects any of the outputs
to the signal processing chain.
The schematic design of the multiplexor circuit has been completed.
The signal from the HAWAII output amplifier is switched with relays, ensuring a reliable and noise free signal path. The unused outputs are connected to ground which both protects against electro-static damage and removes power from the on-chip amplifier. While the camera is integrating all the output amplifiers can be powered off which reduces light emission from the on-chip circuits.
A very flexible addressing scheme has been implemented. A single set of addressing clocks is generated by the camera controller, these are fed to the camera through a multiplexor which allows signals to only reach the selected set of quadrants.
The pattern of selected quadrants and chips is loaded into the multiplexor and a clock sequence generated by the controller. This can be repeated with a different clock sequence for each pattern of chips, allowing any pattern of quadrants to be reset or read. Additionally the whole camera can be reset in the time it takes to address one quadrant. Any pattern of sub-arrays on any of the quadrants can be defined and read independently, the only limitation imposed by the HAWAII on-chip circuitry is that a full row of 512 pixels is reset at a time.
The multiplexor PCB also contains the circuits which generate the bias levels for the HAWAII array. Only one of the levels is generally adjustable, VRESET which controls the well capacity. This level will be adjustable by a trimmer on the circuit and will not normally be changed initial after testing.
A detailed electrical schematic is at Appendix 5.4.1.
The multiplexor is contained on a single circuit board is approximately 200mm square which is mounted in a screened metal box immediately outside the dewar. The camera controller is mounted approximately 1m away and is connected by a single 60way flat cable which carries all the power supplies and clock levels, and a co-axial cable which carries the output signal.
A detailed electrical schematic is at Appendix 5.4.1.
The controller to be used for CIRSI is an AstroCam 4100 high-speed CCD controller, modified in-house to allow it to work well with the HAWAII infra-red arrays.
The controller has been in production by AstroCam for over four years. Many have been sold for all sorts of application, including, for example, one for use as the JOSE optical seeing monitor on the William Herschel Telescope on La Palma. As a commercial product (AstroCam is fully ISO9000 accredited) it is well documented and has a good track record for reliability.
The controller has at its heart a programmable sequencer IC (called
a SAM, made by Altera Inc., USA) that generates patterns of digital
signals at 24MHz. These data patterns are used with CCDs to control
the clock waveforms, the signal processing chain and the analogue
to digital conversions. With CCD use, the internally generated
electronic noise referred to the output of the detector is approximately
4 volts at a pixel rate of 500 KHz, corresponding to about 1 electron
of noise from the HAWAII devices, and higher at higher pixel rates.
The controller is capable of working at a maximum of 8 MHz pixel
rate, and at 5.5 MHz pixel rate with full double correlated sampling.
In practice the HAWAII arrays are limited to a maximum pixel rate
of 1 MHz, with a rate not in excess of 300 KHz if optimum performance
is to be achieved, so only pixel rates at the lower end of the
4100 controllers range are likely to be used. The fast clocking
will permit rapid skipping through unwanted pixels when sub-array
read-out is needed.
The 4100 controller can manage a single output channel at once.
With the HAWAII arrays, each device has four independent channels,
and there are four devices. The single channel 4100 controller
will use a separate multiplexor box to allow one out of the sixteen
quadrants to be addressed and read out at a time. With 4 million
pixels, and a peak read rate of 300 KHz, the read-out time will
be in excess of 10 seconds.
The SAM will be entirely reprogrammed to allow the 4100 controller
to generate the necessary waveforms to clock the signal from the
HAWAII array quadrants, and to process the output analogue signals
appropriately. The 4100 controller is normally used with a buffer
card close to the detector head that provides clock level control
and a small amount of output signal pre-amplification. With the
HAWAII arrays this buffer card will not be needed since the arrays
accept digital TTL (5 volt) logic levels directly. There will
need to be a separate multiplexor box to allow the analogue channels
to be enabled and buffered separately. The multiplexor box will
also provide some impedance matching to allow the output from
the HAWAII channels to be made to look like a CCD output for convenience.
The 4100 controller is operated under the overall control of a simple Transputer microprocessor device. The role of this Transputer is to manage the operation of the SAM by triggering it to carry out one out of a number of functions which are requested with a number of parameters passed with each command. The SAM generates the instruction data patterns, and the data are passed directly to a parallel port (NOT via the Transputer, which is much too slow to handle these data rates), buffered (RS422, with over 100 metre driving range) and so passed back to the host computer.
The SAM code is developed with software provided by the SAM manufacturer. It allows the device to be programmed with a simple assembly-type language that is easy to follow and modify. As a guide and aid to learning the language, AstroCam have agreed to loan a standard CCD head (with a Kodak KAF-0400 CCD), a frame grabber card (an Imaging Technology VFG card), cabling and high-level software. This will allow us to modify the CCD read-out procedures as we gain familiarity with the code and allow us to go on to programming the HAWAII control patterns with more confidence.
An example of the SAM code is attached in Appendix 5.5.1
The INMOS Transputer is programmed in Occam. We do not, however,
expect to need to make any changes to the Occam code supplied
by AstroCam. This code (of which the source has been supplied
by AstroCam) allows all the functions of the SAM to be utilised
fully so we will retain its format so that the calling programme
thinks it is dealing with a CCD. Only the SAM will know better.
This allows the calling software and driver software in the host
computer to remain unchanged, saving a great deal of work. Should
minor modifications be necessary, AstroCam has agreed to assist.
The control of the Transputer microprocessor may be handled via a Transputer link or via a fast RS-232 interface. We intend only to use the RS-232 interface for simplicity and to allow the option of using non-PC control computers in the future.
The data are output in an industry standard format used by many
manufacturers of image frame grabbers. It consists of 16 data
lines, with pixel, line and frame pulses. All data are buffered
for RS-422 twisted pair drivers that are capable of operating
over a cable of over 100 metres length (about 80 metres are used
on the WHT routinely at present).
The faster data rates that are possible with the HAWAII arrays are faster than may be handled by the IBM/PC ISA bus. For this reason we intend to use a PCI bus interface card that AstroCam are integrating. The card is being procured from a small Tucson based company (Spectral Instruments). A beta unit has already been received and works well. A production version is expected in June, with quantities (and one for this project) in July.
In order to speed the development of the SAM code, AstroCam have loaned the project an existing frame grabber (Imaging Technology VFG board) that already works with the calling software until the PCI board is ready and fully integrated with the same software package to make it call compatible.
The camera control and data taking software is a major part of the project. The software system must support all the features needed for engineering tests and observing at the telescope, it must be reliable and robust enough not to waste observing time and be useable by observers without a large support effort.
The development of the software must proceed at the same time as the camera construction and so it is important that the camera testing is not delayed by a lack of control software. Initially AstroCam supplied CCD camera software will be used in the lab for device testing and the astronomical software integrated in as packages are finished.
There are a number of separate components all of which must all
work together. They include:
The programs which generate the clock patterns to operate the arrays are written in a hardware specific low-level language. The programs are generally downloaded and executed in computers embedded in the controller hardware, the mode and format of the readout is selected by downloading a different program or passing parameters to the program. In the NICMOS system this is an assembly language for a custom RISC microprocessor in the HAWAII system a programmable micro-sequencer (SAM) is used. The SAM is unusual in that all the available programs are burnt into the device during development and cannot be changed at the telescope. The parameters of the read such as exposure time, pixel rate or sub-windows can be changed before each exposure.
The centre of the software system is a number of separate packages which each control an area of operation such as image handling, camera control or file handling. Using object-orientated programming techniques it is possible to produce a reliable and robust system with access to the components of the data available through a strictly controlled interface. This methodology allows enhancements and upgrades to the packages without breaking existing code. This code is written in C++ and is designed to operate seamlessly across PC and workstation architectures.
The versatility of the camera system in its wide range of operating modes demands of requires a very flexible software system. It is impossible to predict all the requirements of astronomical and engineering users and so a highly adaptable system is needed, this takes the form of a command and scripting language. The command language is based on TCL an embeddable command shell produced by Sun but available on a range of operating systems. The language is similar to the Unix shell or Perl and is familiar to any Unix user so extra commands may be added by the camera packages to provide support for the camera features. An integral part of TCL is TK. This a windows toolkit which allows graphical user interfaces to be written in TCL with a few simple script commands.
The observing software is based on the Unix / IRAF model of providing a number of simple commands which can be built up into complete packages rather than large difficult to modify programs.
Different sets of commands will be needed for engineering tests and observing at the telescope, and observations in different modes by different astronomers but they are all based on the same basic commands.
Astronomical data analysis packages usually expand out of control to meet all the requirements of all possible users. The model for this system is to produce image data in a standard format and pass this data to an existing astronomical analysis system, such as IRAF. The mechanism for this can be as simple as saving the data in FITS files to a shared directory and have a copy of IRAF being used by another observer to reduce the data, or it may be possible to pass the data transparently to IRAF and have the answers returned and presented to the user.
Analysis features in the camera software will be limited to those needed for engineering tests which are not commonly available in astronomical packages together with some simple on-line statistics to monitor camera operation.
All the programs written in TCL/TK can be moved across platforms unchanged. In addition TCL provides a network transparent data and command transport, TCL programs running on one machine can pass data easily to TCL programs running on networked machines with a different operating system.
The camera will store all it's data in FITS format, although the software has the ability to import and export data in a number of other formats for compatibility with AstroCam systems. On a Unix system supporting IRAF it is also possible to write data in IRAF image format, however this format is not machine independent and is not available as a separate standard on non-IRAF platforms.
The typical data from a full night observing is less than 1.5Gb and can easily be stored on DAT tape. We may also record data on write-able CD-ROM directly.
The camera controller is interfaced to the computer through a serial line for downloading commands and configuration and a high speed parallel data bus to receive image data. The data is handled by a frame-grabber card mounted on the host computer's bus, this can be an SDV board from EDT in a Sun/SPARC or a PCDI board from Spectral Instruments on a PC's PCI bus. The frame-grabber is supplied with a device driver which allows incoming data to be transferred directly into the computer's memory. The PC card is supplied with drivers for Windows95 and Linux. We will use an older ISA frame-grabber card loaned by AstroCam until the PCI card is available.
The TK toolkit provides a simple image display tool but this lacks the colour table and image scaling facilities usually available in astronomical display packages. The system will pipe image data to SAOImage, on Unix based systems, or an AstroCam supplied display utility on Windows systems.
As the camera will be read-out continually during a long integration it is possible to display the image in real-time as it is integrated. The entire image contains too many pixels to easily display the entire image but a single quadrant, or a number of selected sub-windows, could be displayed and updated continually.
It is important that the camera system is well integrated with
the telescope control system. With multiple exposures at different
positions to build up a mosaic image the camera must be able to
efficiently reposition the telescope. The system becomes more
complicated as it must support a range of telescopes with different
control systems. The current plan at the INT is for a server application
to run on the telescope workstation which can communicate with
the telescope control system ( by DRAMA ) and to the camera and
autoguider computers. This will run through the sequence of exposures,
moving the telescope between positions, co-ordinating the autoguider
link and telling the camera when to start an exposure.