UFTI : The 0.8-2.5m Fast Track Imager for the
UK Infrared Telescope
Patrick F. Rochea*, Philip W. Lucasa,b, Craig D. Mackayc, Eli Atad-Ettedguid, Peter R. Hastingsd,
Alan Bridgerd, Nick P. Reese, Sandy K. Leggette, Chris J. Davise, Alan R. Holmesa, AT. Handforda
a
Physics Department., Oxford University,UK; bSchool of Physics, University of Hertfordshire, UK;
c
Institute of Astronomy, Cambridge University,UK; d UK Astronomy Technology Centre,
Edinburgh; e Joint Astronomy Centre, Hilo, Hawaii 96720, USA.
ABSTRACT
In 1996, it was proposed to build a near-infrared imager for the 3.8-m UK Infrared Telescope in Hawaii, to exploit the
1024 pixel format detectors that were then becoming available. In order to achieve a fast delivery, the instrument was
kept simple and existing designs were reused or modified where possible. UFTI was delivered within 2.5 years of the
project start. The instrument is based around a 1k Rockwell Hawaii detector and a LSR Astrocam controller and uses
the new Mauna Kea optimised J,H,K filter set along with I and Z broad-band filters and several narrow-band line filters.
The instrument is cooled by a CTI cry-cooler, while the mechanisms are operated by cold, internal, Bergelahr stepping
motors. On UKIRT it can be coupled to a Fabry-Perot etalon for tunable narrow-band imaging at K, or a waveplate for
imaging polarimetry through 1-2.5m; the cold analyser is a Barium Borate Wollaston prism. UFTI was designed to
take full advantage of the good image quality delivered by UKIRT on conclusion of the upgrades programme, and has a
fine scale of 0.09 arcsec/pixel. It is used within the UKIRT observatory environment and was the first instrument
integrated into ORAC, the Observatory Reduction and Acquisition Control System. Results obtained during instrument
characterisation in the lab and over the last 3 years on UKIRT are presented, along with performance figures. UFTI has
now been used on UKIRT for several hundred nights, and aspects of instrument performance are discussed.
Keywords: Infrared Imager, Infrared Detector, Camera, Photometry, Polarimetry, Prism
1. INTRODUCTION
The 3.8-m UK Infrared Telescope was originally conceived as a low-cost flux collector, but over the last 20 years, it has
been substantially upgraded with the aim of delivering images limited only by the atmospheric conditions above Mauna
Kea. A coordinated upgrades programme has installed an active primary mirror support system, an active secondary
mirror with fast tip/tilt/focus control, thermal conditioning of the telescope environment and improved dome ventilation
(Hawarden et al 1998). By 1996, it was clear that the continuing improvements in image quality and the availability of
larger format infrared detector arrays than those employed in the IRCAM instruments, required a new near-infrared
imager to exploit the scientific capabilities offered by UKIRT. At that time, the UKIRT instrumentation programme
included Michelle, a mid-IR imager spectrometer (Glasse et al 1997) and UIST, a 1-5m imager/spectrometer (Ramsay
Howatt et al 2000), scheduled for delivery in 1998 and 2001 respectively, but there was considerable concern that
UKIRT’s scientific capabilities would fall substantially behind those of other observatories if we had to wait for UIST
to be delivered for an instrument with a 1024 pixel format detector. UFTI was conceived as a fast-track near-infrared
imager to fill this gap, and maintain UKIRT’s 1-2.5m imaging capability at a world competitive level. The instrument
requirements were for a high-sensitivity imaging capability that fully exploited the best UKIRT imaging performance,
with an aim to deliver the instrument within 2 years. UFTI was built as a fully integrated facility instrument and was
the first of the UKIRT instruments to be incorporated in ORAC, the new Observatory Reduction and Control
environment (Bridger et al 2000).
*p.roche@physics.ox.ac.uk http://www-astro.physics.ox.ac.uk/~pwl/UFTI.html Tel: 44 1865 273338
Astrophysics, Denys Wilkinson Building, Keble Road, Oxford OX1 3RH.
Fax: 44 865 273390
2. INSTRUMENT REQUIREMENTS
On completion of the UKIRT upgrades programme, the delivered image quality at 2m under the best conditions (good
atmospheric seeing and low wind) was anticipated to be ~0.25 arcsec. One of the prime requirements for UFTI was to
exploit this, which naturally led to a fine pixel scale. A scale of 0.09 arcsec/pixel, which would sample even the best
images well, was adopted. Because UKIRT currently operates at a f/36 focal ratio, with a focal plane scale of 1.52
arcsec/mm, a coarser scale would lead to a large entrance window, potentially leading to significant problems with
condensation in the central regions, where radiative cooling is most effective. Further, a 90 arcsec field fully covers the
clear apertures of the UKIRT waveplates and Fabry-Perot etalons. Cold, interchangeable camera lenses that would give
a choice of pixel scales were considered, but were not pursued because of the potential impact on the delivery schedule.
The instrument had to fit on the UKIRT instrument support unit (ISU) between two other large UKIRT instruments,
CGS4 and Michelle, leading to a relatively narrow format, which determined that the UFTI mechanical layout would be
relatively tall and thin. The space between the ISU and the front of the instrument was maximized to allow an ambient
temperature Fabry-Perot etalon for narrow-band imaging, or a mask for polarimetry to be inserted. This led to the
entrance window being located quite close to the telescope focal plane.
Figure 1: Left: UFTI optical layout. Right: JHK spot diagrams at the centre, edges and corners of the field.
The circles represent the diameter of the Airy ring at 2m.
2.1 Optical Design
The optical design follows conventional IR camera principles. UFTI is mounted on the UKIRT ISU, and at zenith, is
fed horizontally by a dichroic mirror, which transmits visible light to the guide camera and reflects infrared light to the
instrument. The telescope focal plane lies inside the instrument cryostat, with a field baffle to limit the field of view.
An off-axis gold-coated spherical collimator mirror produces an image of the secondary mirror at the position of a cold
stop. The collimator was manufactured by diamond machining of a T6061 heat-treated Al-alloy blank by Symons
mirror technology of Hertfordshire, UK. 10-position filter wheels are located on either side of the cold stop, allowing
broad- or narrow-band filters to be selected. The second wheel also contains a Wollaston prism for polarimetric
observations and a Hartmann mask for precise focussing. A lens barrel containing the BaF2/LiF camera lenses forms
the image on a Rockwell 1024 x 1024 pixel Hawaii detector array. The optical design was developed to be nearly afocal
from 0.9 to 2.5m so that a focusing mechanism would not be required, thus simplifying the mechanical layout and
reducing complexity. The IR-grade fused silica entrance windows, together with the spherical lenses, which were
supplied mounted in a barrel after cryogenic testing, were manufactured by Graseby-Specac, Orpington, UK.
Optimization and tolerancing of the optical design was conducted at the UKATC, Edinburgh using Code-V and Zemax.
Supporting structures for optical components and surfaces of lens barrels etc. are black anodized, while the inside of the
radiation shield and the collimator support are coated with 3M Nextel black paint to reduce stray and scattered light.
The final design gives a scale of 0.0908 arcsec per 18.5m pixel, and a total field of view of 92 x 92 arcsec. The pupil
diameter at the cold stop is 15mm, allowing 25mm filters to be used. Over- or under-sized cold stops can be installed
between the filter wheel modules, and UFTI currently operates with a slightly undersized stop. The image quality is
near diffraction limited down to 1m, with the worst quality at the top of the detector where the off-axis angles are
largest (Fig. 1). The total chromatic focal shift between 0.8 and 2.5m is 1.0mm, and with the focus set for = 2m
does not contribute significantly to the degradation of image quality.
2.2 Mechanical Layout
The constraints of the space available for UFTI, together with packaging to minimize flexure, lead to a folded
instrument layout. All of the components and mechanisms are mounted on an optical bench made of Al-alloy, and
positioned against a rail mounted on the bench. The bench itself is mounted off the cryostat body on four pairs of G10
Tufnol struts, and covered by an Al-alloy radiation shield. The main cryostat body is an aluminium alloy casting which
was subjected to high-pressure isostatic heating to reduce bubbles and flaws. A welded stainless steel lid is attached via
quick-release clips, with 4 bolts as a back-up in case of pressure loss. In common with other facility instruments, UFTI
is cooled by a closed-cycle refrigerator, which mounts onto the cryostat casting via the anti-vibration mount developed
at the UKATC and used on other UKIRT instruments. The cryostat casting is mounted onto the telescope mirror cell
via an electrically isolated steel support. The overall dimensions of the cryostat are 40cm wide, 71cm high and 80cm
deep, with a total weight, including a stiff telescope mounting frame, of ~150Kg.
Figure 2: Left: Optical bench top surface showing the location of the steering mirror, filter modules, shutter and
camera lens mount and detector box fitted against the locating rail. Right: Side view of the instrument assembly.
Scheduled maintenance, e.g. for filter changes, of the instrument should only require the instrument to be warmed up,
removed from the telescope and then opened up by removing the 4 bolts holding the cryostat lid and undoing the clips.
Removal of the radiation shield then gives access to the detector housing and all of the optical components, with the
exception of the collimator mirror. The only mechanisms provided in UFTI are the two filter wheels and a shutter.
The filter wheels are based closely on the proven CGS4 filter wheel designs (Mountain et al 1990), but scaled down to
fit the available space. They are driven by cold, internal Bergelahr motors, (with bearings replaced and coated with dry
lubricant), via Vespel worm gears. Positional information is provided by a single datum micro-switch. They can each
accommodate up to ten 25mm diameter filters, but in practice, one position has to be kept clear. The filter wheels are
constructed as modules, which can be removed and re-installed without disturbing the rest of the instrument. A shutter
is provided to allow short exposures to be made (shorter than the detector read time), or to prevent over-illumination of
the detector array by bright stars. It consists of a sector blade located just beyond the second filter wheel and is rotated
by a stepper motor. In practice, this is rarely used, but was included because the detector properties were not fully
understood when the instrument was designed. The collimator mirror is located at the end of a turned tube protruding
from the back of the instrument, with its mounting face at an angle of 4 degrees to the incoming beam. The beam is
reflected back and up to a fixed steering mirror, which passes the beam vertically to the cold stop, camera lenses and
detector. The cryostat body was machined and the stainless steel lid fabricated in the Oxford Physics department
workshop along with most of the internal components, but some parts were manufactured by local companies to
minimize delays to the schedule.
Figure 3. The completed instrument. Left: with the cryostat lid and radiation shield removed showing the
detector housing, camera lens barrel, filter wheel modules and beam steering mirror mounted on the optical bench,
from which locating rods for the radiation shield protrude. Right: the instrument mounted on the UKIRT mirror
cell, viewed from the rear.
2.3 Cryogenic Performance
UFTI is cooled by a CTI 350 2-stage closed-cycle refrigerator, but only the first stage is used. The first stage cooler is
connected to the optical bench by copper braids and also directly to the detector enclosure via a relatively weak thermal
link to limit the cooling rate. Under normal operation at the telescope, the optical bench and components stabilize at
about 70K, while the detector temperature is actively controlled at 80K by a Lakeshore 330 feeding current through a
pair of wire-wound resistors mounted on the detector block. A liquid nitrogen pre-cool tank is included underneath the
optical bench; this is not used at UKIRT, but was used in the instrumentation laboratory at Oxford, where the
combination of shorter runs of helium gas pipes, warmer ambient temperature and 50Hz cold head operation
significantly increased the operating temperature. Lakeshore DT 470 silicon diodes provide temperature monitoring of
the detector housing, optical components and the cold head. The heat load onto the optical bench is reduced by multilayer insulation covering the interior cryostat walls. We calculated that the middle of the 120mm diameter cryostat
window would cool significantly, due to radiative cooling into the cryostat and relatively poor thermal conduction in the
fused silica, leading to condensation under conditions of high humidity. UKIRT does not have a dry air or nitrogen
supply available, but condensation can be counteracted to some degree by blowing air across the window; under very
high humidity conditions, heated air is required. In practice, conditions sufficiently humid to lead to condensation are
unlikely to be very good for astronomical observations.
For a normal cycle UFTI takes about 30 hours to warm up and 40 hours to cool down. To prevent potential problems
due to differential contraction, the shutter and filter wheel mechanisms are not moved until the temperatures are below
100K.
3
DETECTOR AND CONTROL SYSTEM
The heart of UFTI is the Rockwell Hawaii 1024 x 1024 pixel HgCdTe detector array, controlled by a modified LSRAstrocam (now PerkinElmer Life Science) 4100 CCD controller. To minimize development time and effort, the
detector control electronics were copied directly from the system developed at the Institute of Astronomy, Cambridge
for CIRSI, the University of Cambridge J and H band imager (Mackay et al 2000). The detector is mounted on a CIRSI
custom printed circuit board, in which the rear of the detector is spring loaded against a copper block, which houses the
temperature control diode and heating resistors, to ensure good thermal contact. The detector package is then mounted
inside a die cast box, painted Nextel black on the inside, which itself is mounted onto a support structure attached via
the rail to the optical bench. The detector bias, clock and data wires are connected to the pcb via a sub-miniature microD connector and emerge from the die cast box as a single bundle. A W.W. Fischer bulkhead connector on the optical
bench allows the signals to pass through to 3 hermetically sealed Amphenol connectors on a plate on the cryostat wall
to the detector interface box. Here, pre-amplification and signal conditioning electronics and relays to switch the
analogue outputs from the four quadrants of the detector are located, (again identical to the CIRSI electronics) before
the signals are passed on to a 14-bit 4100 CCD controller.
Initial laboratory testing was conducted using the proprietary Astrocam Pixcel software package to control the
instrument, but operation at the telescope uses software developed at JAC, communicating directly with the CCD
controller. The Capella 4100 controller runs with custom microcode, written at Cambridge, and stored in sequential
access memory chips, to provide the readout clock sequences needed for the HgCdTe detector. The detector controller
is mounted in an electronics cabinet under the mirror cell while, the interface box mounted on the cryostat passes the
analogue signals from the four detector quadrants to the CCD controller data lines, and the clocks and biases to the
detector. More details of the detector control system can be found in Mackay et al (2000). The gain of the detector
electronics system is set at ~6 electrons per data number. The small pixels on the sky mean that photon rates with UFTI
are relatively low and so even broad-band observations need significant exposures to ensure background-limited data (a
few seconds at K and up to 500 sec at I). For short wavelength or narrow band observations (e.g. with the F-P), a highgain mode can be selected which has ~3e- per data number.
All observations are made in NDSTARE mode, which is a read-reset-read mode and subtracts the fixed-pattern
structure. There are two forms of non-destructive (NDSTARE) readout available. Both do a number of resets to
counteract latency on the array. The counts measured are fitted by a slope; the output image gives the increase in counts
for the requested exposure time, normalised to one coadd if more than one coadd is taken. The two modes differ,
however, in the number of times the array is read out (and consequently the overheads associated with each exposure).
In the standard mode, ND_STARE, the array is read out only twice; 10_NDSTARE is read out 10 times, which results
in a factor of ~2 improvement in read noise down to about 10e -. 10_NDSTARE may therefore be useful when observing
with the FP and sometimes with narrow-band filters, where background signal may be low and so background-limited
performance difficult to achieve. The additional overheads associated with 10_NDSTARE are substantial (~25 sec).
Sub-arrays may be used for observations of compact sources with a corresponding reduction in read-out overhead.
The detector in UFTI has proven to be very easy to use with excellent performance. The NDSTARE read noise is
measured at the beginning of every UFTI night via a routine ORAC "array tests" observation, and consistently gives 24-
26e- for a double-correlated read; this can be reduced to ~10e- with multiple reads. The dark current is ~1e- per pixel per
minute. For data reduction purposes, standard bad-pixel masks are generated by UKIRT staff, and updated periodically.
These typically flag about 0.3% of the pixels as ‘bad’. Over periods of weeks, a number of pixels migrate across the
thresholds set, and so the optimum bad-pixel mask changes slowly over time.
3.1 Detector Performance
The array begins to saturate at about 10,000 DN
and goes into hard saturation at ~14,000 DN,
indicating a well depth of ~85000 e-. Linearity
corrections have been derived and should be
applied for accurate photometry (Fig. 4). At 70%
full well, the corrections are <4% and more details
are available on the UFTI web pages. Residual
images of bright stars are present at a level of 0.30.4% at JHK, even with multiple array resets
before each image. This image persistence acts
like an enhanced dark current from the highly
illuminated pixels, so that it continues to
accumulate with long exposures. For example,
after exposing standard stars to ~4000 DN, even
three frames later, with 80 second blanked off
exposures, a latent image of the star was visible ~2
counts above the background at the standard star
jitter positions. This persistence makes acquisition
of truly uniform flat fields problematic unless care
Figure 4. Linearity correction of the UFTI detector.
is taken. However, the UFTI flatfield is very
stable, and it is possible to take sky flats with a
clean array at the start of the night and use them to generate a flat field throughout the night (as opposed to creating flats
from median filtered jitter sets) for crowded fields. It is good practice to position any bright objects (e.g. standard stars)
in a different part of the array from faint targets, and to keep count levels below ~5000 if possible. Observing
sequences are constructed so that the detector is blanked off between objects, so that it is not illuminated during target
acquisition
UFTI DetectorResponse
UFTI and CCD Response
1200
0.5
1000
0.4
800
0.3
600
0.2
400
0.1
200
0
400
500
600
700
Wavelength (nm)
800
900
0
400
500
600
700
800
900
Wavelength (nm)
Figure 5. Left: the raw UFTI and ST-6 CCD responses measured with a monochromator. Right: the estimated UFTI
detector response between 0.5 and 0.9m obtained by adopting the CCD response curve from the Texas Instruments
TC241 data sheet.
We wished to include filters in the 0.8-1.0m region, and so investigated the instrument response at short wavelengths
in the lab at Oxford using a grating monochromator and tungsten and hydrogen lamps. We did not have access to a
flux-calibrated system, but were able to use an ST-6 CCD as a comparison. The detector response appeared to fall
smoothly with decreasing wavelength from 1 to 0.8m, but then dropped rapidly near 0.78m. However, the device
retains low level (<10% of the sensitivity at 0.9m) residual sensitivity down to wavelengths below 0.5m, with the H
line faintly detectable in a hydrogen discharge lamp. Comparison with the measurements made with the CCD suggests
that there is a very sharp cut-off in the detector response at 786nm, falling rapidly between 786 and 780nm, and then
declining slowly towards 500nm. The tail of residual sensitivity may be due to photons being absorbed in the silicon
multiplexor. These measurements indicate that filters used with the Hawaii detectors should block light down to at least
0.4m as well as out to 2.7m or beyond. The 0.78m cut-off is conveniently placed, just longwards of a substantial
atmospheric absorption band.
4 MECHANISM AND INSTRUMENT CONTROL SYSTEM
UFTI is fully integrated into the UKIRT observatory control system, allowing the instrument observing configuration
and telescope positioning sequences to be preset using the UKIRT Observing Tool for automatic execution. Similarly,
the data collected are pipeline-processed with standard data reduction recipes for on-line assessment under ORAC.
The instrument operates under EPICS (Dalesio, et al 1991) control.
The 5-phase stepper motors are powered by a Bergelahr 3-axis MD 5-52 drive unit, which is controlled by a custom
electronics rack based upon intelligent motor controllers (RS 440-098). Memory chips on the motor controller boards
store control routines for motor datum and filter selection operations, and accept commands and return status
information to the EPICS system. Acceleration and deceleration rates, and step speeds and directions for the motors are
set using the RS programming language on-board the controller cards. The motor control system is based upon that
developed for WHIRCAM, an infrared camera developed at Oxford for the 4.2-m William Herschel Telescope on La
Palma.
Control of the UFTI detector system is based upon that for the UKIRT auto-guider system, which also uses an
Astrocam CCD controller. Much of the software was provided by the vendor in source code form and ported to the
VxWorks operating system as part of the UKIRT upgrades project. This was then interfaced to UKIRT's EPICS based
control system (which, in turn, interfaces to ORAC), with engineering user interfaces developed using EPICS display
tools. The physical connections between the controller and the host computer utilise both a transputer link and a 16 bit
parallel link for the video data. The VxWorks transputer connection was implemented using an INMOS B016 board,
with driver software from the ESO VLT project. The VxWorks parallel link utilised the same type of interface board as
used in the PCI based Astrocam systems, but modified for the PMC form factor by the original supplier.
5. FILTERS AND TRANSFORMATIONS
UFTI is equipped with a selection of narrow- and broad-band filters. The broad-band J,H,K filters are from the Mauna
Kea Observatory consortium batch supplied by OCLI (MKO-NIR), which are designed to match the atmospheric
windows for optimal signal-to-noise (Simons & Tokunaga 2002). A K’ filter is also available, but is most useful for
high-emissivity conditions, for example when using warm waveplates and masks for polarimetry. Custom I and Z
broad-band filters from Barr Associates are also installed. These filters are significantly different from those installed
in optical CCD systems, as they need to be blocked out to 2.7m. The I-band filter bandpass is truncated at the blue
end by the detector cut-off at 0.78m and the reflective cut-off of the dichroic. The MKO-NIR J-band filter, and to a
lesser extent the H-band filter, is significantly different from the old J and H filters that were used in IRCAM prior to
1999. A number of narrow-band filters from the MKO-NIR consortium and elsewhere are also included.
Published magnitudes for near-IR standard stars generally have not used the MKO-NIR filters and so are no longer on
the UFTI natural system, but a programme to re-measure the UKIRT standards through the MKO filter set is underway.
Colour transformations have been derived in two ways: one empirically based on UFTI photometry and the other
calculated by convolving the known filter profiles with spectroscopic data for a representative set of red stars. The two
determinations agree well. The transformations between the old IRCAM3 system and the new MKO-NIR system at H
and K are well behaved and single-valued. However, for the J filter different terms have to be applied depending
whether or not the standard star has intrinsic water absorption features. This is due to the fact that the new J filter cuts
off shorter than the old filter, specifically to avoid water absorption in the terrestrial atmosphere. For stars with no
intrinsic water features the colour transformations are:
KMKO = KUKIRT - 0.020 [+/-0.005] (J-K) UKIRT
(J-H) MKO = 0.960 [+/-0.010] (J-H) UKIRT
(H-K) MKO = 1.205 [+/-0.010] (H-K)UKIRT
(J-K) MKO = 1.040 [+/-0.010] (J-K) UKIRT
However, for stars with significant water absorption, stars with spectral type M4 through to the L class (but not
including the T class with methane absorption):
KMKO = KUKIRT - 0.020 [+/-0.005] (J-K) UKIRT
(J-H) MKO = 0.870 [+/-0.010] (J-H) UKIRT
(H-K) MKO = 1.205 [+/-0.010] (H-K) UKIRT
(J-K) MKO = 0.980 [+/-0.010] (J-K) UKIRT
The UFTI I- and Z-band filters are at significantly longer wavelengths than most optical CCD I and Z filters (with 50%
transmission at 0.78 and 0.92m and 0.85 and 1.05m respectively). The I-band calibration has been compared to
Landolt Cousins-I standards (Landolt 1992), and the Z system to the Sloan Sky Survey standards (Krisciunas et al.
1998). The Sloan Z values have been converted from an AB-system to our Vega=0mag system by subtracting 0.572mag
from the Krisciunas et al. values. The following transformations were measured :
ZUFTI = ZS - 0.34 [+/-0.03] (IC -ZS )
ZUFTI = ZS - 0.21 [+/-0.03] (ZS -JUKIRT )
IUFTI =
IC - 0.72 [+/-0.03] (IC -ZS )
IUFTI = IC - 0.27 [+/-0.03] (IC -JUKIRT )
6. IMAGING POLARIMETRY AND FABRY-PEROT MODES
Linear Polarimetry with UFTI is conducted with a warm half-wave plate located in the instrument support structure
above the dichroic. The waveplate is located in the deployable IRPOL arm, which is lowered into the beam and rotated
to modulate the polarization direction for polarimetric observations. The analyzer, a Wollaston prism constructed out of
Barium Borate, is located in the second filter wheel and separates the orthogonal polarization planes to give two images
on the detector. This means that polarimetric observations are restricted to the broad-band filters and the narrow-band
2.12m S1 H2 and continuum filters, which are installed in the front filter wheel. The prism is 18mm on a side and is
anti-reflection coated between 0.9 and 2.5m; it was manufactured by Bernard Halle, Berlin. We had hoped to use Barium Borate (-BBO), which appears to have extremely good chromatic performance, but that material proved to be
unobtainable and -BBO was used instead. The transmission of the prism was measured with an HP spectrophotometer
in the Clarendon laboratory at Oxford, and is presented in Figure 6. It has excellent throughput from the UV
atmospheric cut-off to 2m, but an absorption band appears at
the long wavelength end of the K window between 2.2 and
2.45m. The chromatic performance is also quite good, but
the separation of the e- and o-rays is wavelength dependent,
varying from ~19.5 to 17 arcsec between the I and K filters.
This can also lead to some slight elongation of images under
the best seeing conditions. To prevent overlapping images in
extended sources, a mask can be inserted in front of the
cryostat window, with clear slots separated by ~18 arcsec.
The mask is located about 60mm in front of the telescope
focal plane, so is slightly out of focus. The rear surface is matt
reflective to reduce the thermal emission seen by the detector
in the K-band, but it may be better to use a K filter rather than
the standard K filter. Installation of the mask is a manual
operation, but all other polarimetric functions can be
controlled remotely. Polarimetry is a popular option with
UFTI, and like all the other near-IR instruments, instrumental
polarization is low (<0.25%). Simultaneous measurement of
the e- and o-rays provides good immunity to noise that can
arise from variations in atmospheric seeing and transmission.
Figure 6: The measured transmission curve between
0.2 and 3.2m for the 7mm thick BBO prism.
A K-band Fabry-Perot etalon can be inserted
between the UFTI cryostat window and the UKIRT
ISU for narrow-band (R ~400km/sec) imaging. The
etalon has a spacing of 40m, a Finesse of ~25
across the K band, and is operated in the converging
f/36 UKIRT beam. The phase shift between the
centre and edge of the usable 70 arcsec field of view
is about 70-100 km/s, but is much smaller over the
central 50 arcsec. Thus, one setting of the FP is
sufficient to accurately image an unresolved line
across the entire unvignetted field. For the same
reason sky OH lines will either be transmitted or
rejected across the entire field (i.e. no sharp rings)
and will not cause serious flatfielding problems. The
separation of adjacent orders corresponds to d/~
0.027, so narrow band filters with percentage
bandpasses of up to about that value will provide
good blocking. The F-P can be used in conjunction
with IRPOL to perform spectropolarimetry on
extended line emission.
With UFTI's small pixels, getting backgroundlimited performance from the array can be very
difficult with the narrow bandpass of the FP, and the
"High-Gain" readout mode is therefore recommended.
Both the Fabry-Perot and IRPOL are controlled within
ORAC sequences. A number of ORAC data reduction
sequences also exist for standard instrument set-ups in
these modes.
Figure 7: A raw image showing the unvignetted F-P
field lying within the 92 arcsec UFTI field. The cluster
of open pixels can be seen at the top left of the array.
7. INSTRUMENT INTEGRATION AND COMMISSIONING
The instrument was integrated and tested in the laboratory in Oxford from April to August 1998. The first cool-down
with the engineering grade array was in April 1998, about 18 months after the initial funding for the project had been
approved, but at that time, several components were still being manufactured. The main cause of delay up to that point
had been caused by a faulty main cryostat casting which had not been processed properly and turned out to have
substantial numbers of bubbles, inclusions and other flaws. Other delays were caused by lack of effort at critical times.
Once we were satisfied that the detector was working properly, the science grade detector was installed. By this time,
the detector electronics design had been tested with CIRSI at Cambridge and we were confident that it could meet the
performance requirements.
In addition to the normal problems encountered in instrument tests, such as extraneous pick-up on the data lines, we
came across two main problems in the integration phase. The first was a problem with one of the filter modules, where
the filter wheel would seize up periodically. This was fixed temporarily by decreasing the step frequency, but was not
finally eliminated until it eventually seized at UKIRT early in 1999; the fault turned out to be due to a small burr on a
shaft under a bearing. The second problem proved difficult to track down, and appeared at first sight to be a light leak
inside the cryostat. Several attempts to eliminate a putative light leak made no difference, and we eventually concluded
that instead the cause was a glowing junction on the detector multiplexor illuminated the device; this was confirmed by
reinstalling the engineering array. We returned the detector to Rockwell who very quickly and efficiently provided us
with a replacement. This immediately solved the problem and allowed us to ship UFTI to UKIRT for commissioning.
The instrument had first light on UKIRT at the end of September 1998, and was commissioned and characterized over
the following months, with the first science programmes conducted in December 1998. Several minor problems were
solved over the first week of commissioning including software problems encountered in transferring from the pc-based
Pixcel detector control to the Observatory system, and in tracking down and eliminating noise sources. The overheads
on detector read-out were also substantially reduced in this period, and the ORAC data reduction scripts brought into
operation. UFTI generated significantly larger format images than UKIRT had had to cope with previously, and this
placed strains on the data acquisition and reduction computing resources, which were dealt with by providing higher
specification computers and careful programming of reduction algorithms.
Over the next 6 months, the instrument performance was characterized. The system gain was determined from ontelescope tests, and found to differ by about 15% depending upon which of the two detector controllers is in use.
Optimum exposure times and observing strategies were developed in the face of instrument properties such as
persistence and significant readout overheads. High humidity conditions cause condensation on the cryostat window,
which would significantly increase the fraction of UFTI time lost. A fan was installed in front of the UFTI window in
December 1998 to aid in clearing condensation, and worked well immediately. This generally allows the instrument to
be operated when the humidity is below the UKIRT limit of 90%. Determination of optimum focus is difficult unless
long integrations are used to average out atmospheric fluctuations, and a Hartmann mask has been installed to improve
the accuracy and speed of focus determination.
8. INSTRUMENT PERFORMANCE
The overall throughput of the instrument has been calculated from observations of standard stars to be ~26% at H and K
and 13% at J; these figures include the atmosphere, telescope, dichroic, instrument and detector.
The lower than
expected throughput at J is presumably due to significantly lower detector quantum efficiency at 1.2m than at 2m.
The sensitivity of UFTI is measured to be K~ 20.3 mag 5 sigma 1hr in a 2 arcsec aperture under good seeing conditions,
and assuming that separate sky frames are not required. This is about 20% better than obtained with IRCAM-3 on
UKIRT, because the lower sky background obtained with the HgCdTe detector and lower system emissivity with UFTI
more than compensate for the higher quantum efficiency of the IRCAM InSb detector. The sensitivity in the J-band is
comparable or slightly worse than that of IRCAM, reflecting the lower detector efficiency.
The efficiency of operation, or open-shutter fraction, is strongly dependent on the detector integration time. The time
required for calibration observations of standard stars is dominated by detector readout, filter selection and telescope
motion overheads, and the overall efficiency can be as low as 30%. Fainter targets need much longer integrations, and
the efficiency then approaches 90%.
In common with other UKIRT instruments, a set of diagnostic dark exposures is made at the start of every night to
check instrument performance. These also provide valuable records of variations and trends in properties such as
numbers of dead and hot pixels. UFTI has been used for over one third of the scheduled nights on UKIRT in 1999,
2000 and 2001. Over the last 3 years, there have only been a few failures. Towards the end of 2000, low level noise
was seen and the array controller was swapped for the spare early in 2001. Some vignetting of the filters was also
detected and in March 2001 the instrument filter wheels were refurbished, solving the problem. After a power outage in
November 2000, UFTI had cooling problems, which were traced to a faulty coldhead, which was then replaced.
The instrument routinely delivers images with FWHM less than 0.5 arcsec and under the best conditions yields wellsampled images <0.3 arcsec
9. CONCLUSIONS
UFTI has fulfilled its principle aim of providing a basic near-infrared imager, which has now been used on over 400
nights for many programmes over a wide range of astronomical research projects. After the initial shake-down, it has
proved to be sensitive, reliable and robust, and has required relatively few unscheduled maintenance interventions. Reuse or adaptation of existing instrument designs and solutions saved substantial time and costs in the UFTI project, and
future instrumentation projects could probably save effort by similar tactics. The UFTI collaboration worked well, but
inevitably, there are frustrations and overheads in dealing with a widely-dispersed project. The instrument was
delivered on a relatively short timescale (~2 years from the receipt of first funding) by resisting the temptation to
include additional capabilities (e.g. a second imaging pixel scale and/or a spectroscopic mode), and by re-using
solutions developed for, and proven in, other instruments.
The instrument test and integration phase is the most uncertain, as the most difficult problems cannot be predicted with
any certainty. In the case of UFTI, a cold test chamber would have allowed us to check out the filter module operation
at cryogenic temperatures without needing to cool the whole instrument, and would also have allowed us to test the
detector array in a known dark environment.
More details of the current performance and ‘features’ in the instrument, as well as recent images produced, can be
found on the UKIRT web site: http://www.jach.hawaii.edu/UKIRT.
Finally, maintenance of good image quality places additional burdens on the observatory, as the telescope and
instrument have to be maintained near optimum conditions as deteriorations are quickly noticed by visiting observers.
ACKNOWLEDGEMENTS
We are grateful to the UK Particle Physics and Astronomy Research Council, which funded UFTI through the UKIRT
instrumentation programme. Interim initial funding for the project was provided by the Oxford University Physics
Department. Thanks are due to many people who contributed to the instrument, including Keith Nobbs, Phil Evans,
Ray Knott, Martin Beckett, Mike Snook, David Laird, Maren Purves, Malcolm Currie, Tim Chuter, and Dwight Chan.
We thank all of them and others not mentioned here. We are especially grateful to Craig Cabelli at Rockwell Science
Center for providing a replacement detector quickly.
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