Instrument Overview
===================
 
See [TOMASKO_ETAL1997].
 
Sunlight plays a key role in driving many important physical processes
in planetary physics. Absorption of UV light drives photochemical
reactions, leading to changes in atmospheric composition and the
production of atmospheric aerosols. The size, shape, composition and
distribution of aerosols and cloud particles determine their optical
properties and their ability to absorb sunlight and emit thermal
infrared radiation. Thus they play a key role in the atmosphere's
thermal balance. The net radiative heating or cooling rate provides
the forcing for atmospheric dynamics, which in turn can affect the
distribution of aerosol and cloud particles and influence climate. The
composition, thermal balance, dynamics and meteorology of the
atmosphere also affect (and are affected by) the nature of the
surface. Surface images in reflected sunlight, together with near-IR
reflection spectra, can reveal the nature of the surface and its
interactions with atmospheric processes. Optical measurements of solar
radiation made inside a planetary atmosphere can thus reveal a great
deal about many important physical processes occurring there. The
Descent Imager/Spectral Radiometer (DISR) is the optical instrument
aboard Huygens that makes measurements at solar wavelengths. It was
developed in a collaborative effort by scientists from the US, France
and Germany. DISR measures solar radiation using silicon photodiodes,
a 2D silicon Charge Coupled Device (CCD) detector and two InGaAs
near-IR linear array detectors. Fibre optics connect the detectors to
many separate sets of foreoptics that collect light from different
directions and in different spectral regions. In this way, the
instrument can collect a suite of measurements that are carefully
selected to answer key questions concerning the nature of Titan's
surface and the composition, meteorology, thermal balance and clouds
and aerosols in Titan's atmosphere.
 
 
Scientific Objectives
=====================
 
Thermal balance and dynamics
----------------------------
A basic objective of the DISR investigation is the direct measurement
of the solar heating rate's vertical profile. This will be done using
measurements of the upward and downward solar flux at 0.35-1.7 um from
160 km to the surface at a vertical resolution of approximately 2 km.
The downward flux minus the upward flux gives the net flux, and the
difference in the net flux at two altitudes gives the amount of solar
energy absorbed by the intervening atmospheric layer. Knowledge of the
solar heating profile is necessary for understanding the atmosphere's
thermal balance.
Combining the solar heating profile with the thermal cooling profile
provides the net radiative drive for atmospheric dynamics. The
radiative cooling profile will be modelled using the temperature
profile and the opacity of atmospheric gases and cloud and aerosol
particles at wavelengths in the thermal-IR. The gaseous composition
and temperature profile will be meeasured by other Huygens and Cassini
instruments. The DISR measurements make an important contribution by
determining the size, shape, optical properties and vertical
distribution of aerosol and cloud particles. Once the solar heating
and thermal cooling have been combined, model computations can be used
to estimate the wind field from the radiative forcing.
Finally, DISR will measure the horizontal wind direction and speed as
functions of altitude from images of the surface obtained every few
kilometres in altitude, which will directly show the Probe's drift
over Titan's surface. The measured wind speed and direction determined
by DISR can be compared to the wind field computed from the net
radiative forcing determined above.
 
Distribution and properties of aerosol and cloud particles
----------------------------------------------------------
Several properties of the cloud and aerosol particles are important
for understanding their interaction with the solar and thermal
radiation field. The particle sizes compared to the radiation's
wavelength is important for understanding particle scattering.
Measurements of both the forward-scattering and polarising nature of
Titan's aerosols have been used to show that spherical particles
cannot simultaneously explain these two types of observations (see
[HUNTENETAL1984]; [WEST&SMITH1991]). Information on particle shape in
addition to size is therefore required for understanding particle
scattering. The vertical distribution of the particles also influences
the profiles of solar and thermal radiation. Finally, a suite of
optical properties as functions of wavelength is needed to permit
accurate computations of the particles' interactions with radiation.
These properties include the optical depth, single scattering albedo
and the shape of the scattering phase function. The variation of these
optical properties with wavelength, together with determinations of
size and shape, can yield the imaginary refractive index and thus
constrain the composition of the particles.
DISR will measure many of these properties using combinations of
measurements of small-angle scattering in the solar aureole in two
colours, measurements of side- and back-scattering in two colours and
two polarisations, measurements of extinction as a function of
wavelength from the blue to the near-IR, and measurements of the
diffuse transmission and reflection properties of layers in the
atmosphere.
 
Nature of the surface
---------------------
Titan's surface was hidden from the Pioneer and Voyager cameras by the
layers of small haze particles suspended in the atmosphere.
Nevertheless, intriguing suggestions regarding the nature of the
surface have been made [LUNINE1993], including the possibility that it
consists of a global ocean of liquid methane-ethane. Recent radar
observations [MUHLEMANETAL1990] and direct observations at longer
wavelengths, [SMITHETAL1996]; [LEMMONETAL1995], strongly hint that the
surface is not a global ocean. The many fascinating surfaces observed
by Voyager on satellites of the outer solar system showed a surprising
range of phenomena, including craters, glacial flows, frost and ice
coverings, and active geysers and volcanoes. This preliminary
exploration of the small bodies of the outer solar system suggests
that Titan's surface may well also contain surprises.
DISR will determine the surface's physical state (solid or liquid)
near the impact site, and determine its fraction in each state. DISR
will measure topography, hopefully revealing some of the physical
processes that have formed the surface. DISR will obtain reflection
spectra of surface features from the blue to the near-IR in order to
constrain the composition of the different types of terrain observed.
In addition, DISR will image the surface at resolution scales from
hundreds of metres (similar to those from the Orbiter) to tens of
centimetres over as large an area as possible to permit studies of
surface physical phenomena and to clarify surface-atmosphere physical
interactions.
 
Composition of the atmosphere
-----------------------------
Huygens carries a Gas Chromatograph Mass Spectrometer (GCMS) for
direct measurement of the atmosphere's composition. Nevertheless,
direct sampling techniques can give inaccurate mixing ratios for
condensable constituents if a cloud particle enters and slowly
evaporates in the inlet system. DISR provides an important
complementary capability for measuring the mixing ratio of methane,
the most likely condensable constituent, using a technique that is not
subject to this potential problem. It will be obtained from the
increasing depth of methane absorption bands as the gas path between
the instrument and top of the atmosphere increases during descent.
Methane can exist as a solid, liquid or gas on Titan, and has been
suggested as playing a role in Titan's meteorology similar to the role
played by water on Earth. DISR's measurements of the methane mixing
profile will be analogous to a relative humidity profile on Earth.
Finally, Titan's atmosphere is believed to consist primarily of
nitrogen, methane and argon. DISR's measurements of methane's mixing
ratio, together with determination of the atmosphere's total mean
molecular weight by radio occultation measurements made by the Cassini
Orbiter, will indirectly yield the argon to nitrogen mixing ratio as
an important backup to Huygens' mass spectrometer measurements.
 
 
Instrument Approach
===================
 
In order to achieve this broad range of scientific objectives, it is
necessary to measure the brightness of sunlight in Titan's atmosphere
with several different spatial fields of view, in several directions
and with various spectral resolutions. For measurements of solar
energy deposition, for example, measurements are needed of the
downward and upward solar flux with broad and flat spectral
sensitivity, and with a cosine zenith angle weighting. For
determination of the surface's composition, spectral resolution is
desirable and spatial information is necessary. For determination of
the physical processes occurring on the surface, images are needed
with very broad fields of view looking down toward the surface. To
determine the size distribution of aerosol particles above the Probe's
altitude, upward measurements are needed of the brightness of the sky
near the Sun (the solar aureole) in at least two colours with modest
angular resolution. Images looking out towards the horizon are useful
for sensing the presence of thin haze layers during descent.
It is not possible to include separate instruments devoted to each of
these scientific objectives in Huygens' limited payload. Nevertheless,
it has been possible to increase the usefulness of the single Huygens
optical instrument considerably by making extensive use of fibre
optics to collect light from different directions and bring it to a
few centrally-located detectors after various spectral or spatial
analyses. In this way, redundant electrical systems have been
minimised and moving mechanical parts have been all but eliminated. A
summary of the locations of the fields of view and spectral coverage
of DISR's optical measurements is given in Table 1, while the onboard
sources are summarised in Table 2.
 
Table 1. Summary of DISR instruments.
----------------------------------------------------------------------
Upward-Looking       Azimuth  Zenith  Spectral  Spectral Spatial Pixel
  Instrument          Range   Range  Range (nm) Scale    Scale  Format
                      (deg)   (deg)             (/pixel) (/pixel)
----------------------------------------------------------------------
Violet Photometer (ULV)  170  5-88   350-480      -       -         1
Visible Spectrom. (ULVS) 170  5-88   480-960    2.4 nm    -      8x200
Infrared Spectrom. (ULIS) 170 5-88   870-1700   6.3 nm    -        132
Solar Aureole (SA I)     6    25-75  500 +- 25    -       1 deg   6x50
 Vertical Polarisation
Solar Aureole (SA 2)     6    25-75  500 +- 25    -       1 deg   6X50
 Horizontal Polarisation
Solar Aureole (SA 3)     6    25-75  935 +- 35    -       1 deg   6x50
 Vertical Polarisation
Solar Aureole (SA 4)     6    25-75  935 +- 35    -       1 deg   6x50
 Horizontal Polarisation
Sun Sensor (SS)     64 deg cone 25-75  939 +- 6   -       -         1
 (64 deg cone FOV)
----------------------------------------------------------------------
 
----------------------------------------------------------------------
Downward-Looking     Azimuth  Nadir   Spectral  Spectral Spatial Pixel
   Instrument         Range   Range  Range (nm)  Scale   Scale  Format
                      (deg)   (deg)             (/pixel) (/pixel)
----------------------------------------------------------------------
Violet Photometer (DLV)  170  5-88    350-480     -       -         1
Visible Spectrom. (DLVS)  4  10-50    480-960   2.4 nm    2 deg 20x200
Infrared Spectrom. (DLIS) 3 15.5-24.5 870-1700  6.3 nm    -        132
High-Res. Imager (HRI)  9.6  6.5-21.5 660-1000    -   0.06 deg 160x256
Med.-Res. Imager (MRI) 21.1 15.75-46.25 660-1000  -   0.12 deg 176x256
Side-Looking Im. (SLI) 25.6  45.2-96  660-1000    -   0.20 deg 128x256
----------------------------------------------------------------------
 
 
Table 2. Summary of DISR onboard sources.
----------------------------------------------------------------------
                        Number    Power               Spectral
System                 of Lamps   each       Field   Range (nm) Optics
----------------------------------------------------------------------
Inflight Calibration         3    1 W     Fills each    400-2000   f/2
                                       instrument FOV       fibre feed
Surface Science Lamp (SSL)   1   20 W  4x12 deg centred 400-2000  50mm
                                         on DLIS FOV          parabola
Sun Sensor Stimulator        1  10 mW  Illuminates only 939 +- 6 feeds
                                diode  Sun Sensor detector       fibre
----------------------------------------------------------------------
 
One of the detectors around which DISR is built is a 512 x 520-pixel
CCD silicon detector with a 400-1000 nm wavelength response. The CCD's
surface is divided into nine separate regions, with the light
collected by different foreoptics and brought to the detector by fibre
optic bundles and ribbons. These include imagers looking in three
different directions with different fields of view and angular
resolutions, two regions fed by light collected by upward- and
downward-looking grating spectrometers for flux measurements and for
making spatially-resolved spectra of the surface at 480-960 nm, and
four regions devoted to measurements across the solar aureole in two
colours and in two different polarisation states.
The second type of DISR detector is a pair of 150-element InGaAs
near-IR linear arrays. They are mounted side by side in the focal
plane of a second grating spectrometer covering 870-1650 nm. This
spectrometer is also fed by two sets of optical fibres collecting:
  1. the downward flux from a horizontal diffusing flux plate
     sensitive to half the upper hemisphere
  2. the brightness of a strip on the ground defined by the input
     slit, which provides a measure of the upward flux as well as the
     reflectivity of a well-defined ground region.
The third detector type comprises single silicon photodiodes with
enhanced UV response to extend the upward and downward flux
measurements to 350 nm from the short wavelength limit of the visible
spectrometer at 480 nm. This type of detector is also used in a
separate optical system to detect the Sun's azimuth for controlling
data collection timing.
 
 
Detectors
=========
 
CCD detector
------------
Co-Investigator Dr H. U. Keller of the Max Planck Institut fuer
Aeronomie (MPAE) in Germany is responsible for the CCD subsystem,
including the CCD itself and the electronics to clock and read out the
detector. MPAE contracted with Loral Fairchild to supply the flight
CCDs. The data interface between the German and US electronics is at
the output of the 12-bit analogue-to-digital (A/D) converter.
The CCD is a modification of a standard 512 x 512-pixel two-phase
scientific imager. Seven masked columns and one grey (transition)
column are added on the right side of the detector face. It is divided
into an image section and a memory, or storage, section, each 256 x
520-pixels in size. The image section requires anti-blooming drains to
provide protection from the excess charge produced when the Sun
flashes through the solar aureole section. The anti-blooming drain
takes up 6 um on the side of each pixel. The individual pixels are 17
x 23 um (sensitive area) on 23 um centres. The CCD inputis via a
coherent fibre optic bundle with 6 um core size and 8 um centre-to-
centre spacing. The CCD is fed by optical fibres from nine optical
subsystems: HRI, MRI, SLI, ULVS, DLVS and the 4-channel SA radiometer
(two colours in each of two orthogonal polarisation states). The
detector has a relative spectral response typical for silicon, and is
sensitive between 400 nm and about 1000 nm, depending on temperature.
The quantum efficiency of the CCD material is also typical for silicon
and quite good (~ 50%) at peak.
Exposure time is controlled solely by rapid (0.5 ms) shifts of the
charge from the top half to the bottom half of the chip, which is
covered by an opaque metal film. The exposure time (typically between
10 ms for images and 500 ms for spectra) is the time between rapid
shifts. No mechanical shutter is needed and no moving parts are used
in this system, a feature that makes possible the use of different
foreoptics for many different functions all using the same CCD. The
full CCD frame is read out and digitised for imaging data, a process
that takes about 2.2 s for the entire frame. Only the first 49 columns
are digitised for taking spectra and solar aureole measurements, thus
shortening the readout time to about 300 ms when only these
measurements are made. The CCD pixel signals are pre-amplified in the
Sensor Head (SH) to minimise the possibility of noise. After
amplification, the signals are relayed over a short connecting cable
to the Electronics Assembly (EA), where they undergo video processing,
including correlated double sampling. A fast 12-bit A/D converter
digitises the pixels, which are clocked to the bus. After A/D
conversion, a pseudo-square root algorithm reduces the imaging data
from 12 to 8 bits/pixel. The image data are then compressed in a lossy
hardware compressor. The compression factor is programmable, and will
be between 3:1 and 8:1 for different images. The spectral and solar
aureole data from the CCD are compressed with a lossless software
compression algorithm from 12 to about 6 bits/pixel depending on the
entropy of the data, and then buffered for transmission.
The dark current from the CCD is measured by the signals from the
column of masked pixels along the chip's edge. These data are read out
and transmitted to the ground once every data cycle (about every 1.5-3
min). At Titan entry, the chip is at a temperature of some 260K, and
the dark current is a few per cent of typical signals. After some 40
min, the detector cools to < 200K, and the dark current is essentially
negligible.
The full well capacity of the CCD pixels is about 125000 electrons.
This is digitised to 12 bits for 4096 levels of some 30 electrons/step
(before square root compression). The read noise of the system (CCD,
electronic amplifiers, sample and hold, A/D conversion and subsequent
handling) is < 23 electrons. The data are shot-noise limited for all
but the lowest signals.
 
IR detector
-----------
The focal plane assembly consists of two linear photodiode arrays (one
dedicated to DLIS and the other to ULIS) along with their associated
pre-amplifier electronics. Each array contains 150 individual InGaAs
photodiodes connected to CCD readout registers. The detector arrays
include thin titanium masks to protect the CCD multiplexers from
visible stray light. Each detector chip is bonded onto a sapphire
substrate. The two elementary modules, consisting of the detectors
plus CCD on the sapphire base, are assembled on a ceramic base and
protected by a hermetically sealed titanium case that includes a
coated glass window. A copper thermal lug bonded at the rear of the
ceramic base and connected by a thermal strap to Huygens' exterior is
used for cooling the detector assembly. A silicon diode bonded onto
the ceramic base provides a measure of the focal plane array's
temperature. The detector assembly is mounted on a printed wiring
board connected through a flex cable to the IR pre-amplifier board in
the Sensor Head. Each pixel has a photo-active area of 38 x 300 um;
the pixel pitch is 52 um. The detector assembly was manufactured by
Thomson-TCS (Saint Egreve, France). The general design is based on the
technology used in the Spot 4 Earth observation satellite
[BODIN&REULET1987]. The length of the InGaAs photodiodes has been
increased by a factor of 10 relative to Spot's detectors to improve
the signal-to-noise ratio (S/N) at the lower brightness level of
Titan's spectrum. The IR pre-amplifier board in the Sensor Head and
the clocking electronics (on a board in the DISR Electronics Assembly)
were built by AETA (Fontenay-aux-roses, France).
The detector arrays are used for detecting radiation covering 850-1700
nm. Detector quantum efficiencies of around 50% were measured at 850
nm for the photodiodes; > 80% were found for 1000-1500 nm. The cutoff
wavelength varies linearly with temperature. The 50% level of quantum
efficiency is reached at 1620 nm at 180K, and 1690 nm at 270K.
Between readouts of the photodiodes, charge is accumulated due to dark
current and at a rate proportional to the flux of incident photons.
The charge generated between readouts is digitised using a 14-bit A/D
converter. The gain is 920 electrons per digital step with a full
scale of some 14 million electrons. The dark current in the InGaAs
diodes decreases roughly by a factor of 2.5 for every 10K decrease in
temperature. Measurements at 270K show dark currents in the range 0.5-
2 pA. The dark current is so low that even near room temperature
saturation due to dark current alone requires as long as 0.1-0.4 s
between reads. Somewhat larger dark currents are expected at Titan as
a result of the impact of energetic protons during cruise. The minimum
time between reads is 8 ms, which is sufficiently short that only a
small fraction of the wells contain dark current at the 260K expected
at the start of Titan descent. After some 40 min, detector temperature
decreases to 200K, and dark current is almost negligible.
A shutter mechanism (the only moving part in the entire DISR) at the
entrance slit of the IR spectrometer permits separate measurements of
the dark signal alone and the dark plus light signal throughout the
descent. Spectra with shutter open and closed are both included in the
telemetry stream. The photon flux is obtained by subtracting one
spectrum from the other. In addition, four pixels at the beginning and
end of the array are masked with an opaque resin and provide a measure
of the typical dark current, assuming that they are representative of
the rest of the array. In case of shutter failure, they would provide
an estimate of the dark current that can be used to remove the dark
signal from shutter-open spectra obtained at warm temperatures. The
readout noise of the CCD register includes the noise associated with
the various capacitances, the output amplifier and the transfer noise.
It has been measured to be ~ 1100 electrons for temperatures < 270K.
The shot noise is generally lower than the readout noise. About 10% of
the pixels exhibit, in addition, a 1/f noise due to defects in the p-n
junction of the diode itself. Its amplitude is roughly proportional to
the dark signal and amounts to 0.1-1% of the dark level.
 
Violet detectors
----------------
The CCD and IR array spectrometers cover 480-1700 nm. There is
considerable interest in the radiation shortward of 480 nm, which is
strongly absorbed by aerosols in the upper stratosphere. As there are
only extremely weak methane bands in this part of the spectrum, high
spectral resolution is not required. Two silicon photodiodes,
appropriately filtered so that the 350-480 nm bandpass gives a flat
spectral response, measure the flux in the upward- and downward-
viewing directions. The ULV detector is behind the input diffuser of
the ULVS. The DLV detector requires a separate window from the DLVS,
with a diffuser plate and baffle system to define its pi steradian
field of view (FOV).
 
DISR sub-instruments
--------------------
 Imagers
 -------
The design of the imagers is driven by several considerations. The
range to Titan's surface decreases by three orders of magnitude during
the descent from 160 km at entry to only a few hundred metres at the
last image. Even the 160 km maximum range is orders of magnitude less
than is typical for images of planetary surfaces obtained by any other
technique, including close flybys. Thus, the usual requirement for
high angular resolution. necessary for observations made at longer
range, are much relaxed in our case. We have chosen a maximum angular
resolution of 0.06 deg/pixel, similar to that of the naked human eye.
This pixel size of about 1 mrad gives a spatial scale of 160 m/pixel
at descent start and about 20 cm/pixel at 200 m. The low resolution
overlaps that available from orbiters and the high resolution near
impact is three orders of magnitude greater.
Given the limitation of only 256 pixels across the CCD's active area,
even this relatively low angular resolution would limit the coverage
in a single image to 15 deg. In order to observe as large a surface
area as possible, we have chosen to divide the CCD's long dimension
into three image frames centred at different angles from the nadir
that, together, cover the range of nadir angles from 6.5 deg to 96
deg. This is done with the division shown in Table 1. As the Probe
rotates, sets of three images can be obtained at 12 azimuths to give a
panorama with full coverage over nadir angles from 6.5 deg to 75 deg.
A set of 24 images in azimuth, obtained in two staggered sets of 12 in
successive measurement sets, permits full coverage in azimuth for the
most side-looking imager and also completes the mosaic from nadir
angles of 6.5 deg to 96 deg.
The size of the imager pixels and their nadir angles define the
relationship between integration time and blurring due to Probe
rotation during image exposure. We limit the integration time so that
rotation during exposure is ~1.5 times the angular size of a pixel in
the centre of the High-Resolution Imager's FOV. The blur in the other
imagers is <= that in the HRI. This limitation on integration time
translates into a relation between the f/number of the optical system
and the S/N ratio in the images for a model of the brightness of
Titan's surface. We selected an f/2.5 system to permit S/N >= 100 at
the rotation rate of 1.5 rpm expected at low altitudes for a nominal
model of Titan's surface and the quantum efficiency and transmission
expected for the imaging system.
For good imaging, we require that the optical systems be capable of
giving spot sizes (full width at half maximum) smaller than two CCD
pixels. The spot size is produced by geometrical aberrations in the
lenses, diffraction and the spreading of light between the back face
of the fibre bundle and the CCD's surface. So long as the end of the
fibre optic bundle is within 20 um of the CCD's face, spot sizes
smaller than two pixels can be achieved.
The light is brought from the three separate lens assemblies to the
CCD face by a fibre optic bundle produced by Collimated Holes, Inc.
The fibres are small compared to the 17 x 23 um active area of the CCD
pixels, so that many (about six) fibres feed each pixel. Extramural
absorption is added between the clad fibres to absorb light that might
emerge from individual fibres. The glass used for the fibre optic
conduit is specially selected for its resistance to darkening from
energetic particle impact.
The imagers' spectral range is limited to wavelengths between 660 nm
and about 1000 nm. The long wavelength end is defined by the limit of
the CCD detector. The short wavelength end is selected to prevent
Rayleigh scattering by the atmosphere below the Probe from dominating
the signal due to light reflected directly by the surface. All three
imagers have the same spectral range. Colour can be added to the
monochromatic images using measurements of spectral reflectivity of
the surface measured by the DLVS. We plan to correlate the spectral
measurements of several thousand surface points with the morphology of
these points as seen in the images to produce a colour library as a
function of the type of terrain (craters, canyons, lakes, stream beds,
etc). In this way, surface colour images can ultimately be made from
our measurements.
We expect to obtain > 500 separate surface images that can
be assembled into 36-image panoramic mosaics every few km from the
start of descent at 160 km altitude.
 
 Visible spectrometer
 --------------------
Measurements of the spectrum of the upward- and downward-streaming
sunlight as functions of altitude during descent support most of
DISR's scientific objectives, including:
  1. measurement of the solar energy absorption profile
  2. measurement of methane's vertical profile
  3. measurements of the optical properties, size and vertical
     distribution of atmospheric aerosols and clouds
  4. measurements of the composition and nature of Titan's surface.
Accordingly, significant effort and experiment resources have been
devoted to this portion of the instrument.
The visible spectrometer uses separate fibre bundles to bring
downward- and upward-streaming sunlight to adjacent positions along
the entrance slit. The speed of the collimator and camera optics of
the transmission grating spectrometer is f/2. A tipped flat plate is
included in the beam just before the grating to decrease the
sensitivity of the measured intensity to the state of the incident
light's linear polarisation. The spectral range is limitied by a
filter to 480-960 nm in first order. The upward- and downward-viewing
spectra are focused on the ends of two fibre optic ribbons that join
the fibre conduit used for the imagers. The spectra are spread over
200 CCD pixels in the spectral direction. The spectral spread function
has a width (Full Width Half Maximum, FWHM) of two pixels. The
resulting resolving power of the spectrometer, lambda/(Delta lambda),
is 100 at the blue end and 200 at the red end of the spectrum.
The input optics used to collect the light from the upward and
downward directions are quite different. In the upward-looking case,
the design is driven by the desire to measure the downward-streaming
flux. Thus, a horizontal diffusing plate is viewed by the input
optical fibres using a small lens. The diffuser makes the angular
response function similar to the ideal response function (proportional
to the cosine of the zenith angle). As the instrument protrudes only
slightly from the Probe, a baffle is used to limit the range of
azimuth angles to 170 deg. This baffle is blackened to minimise
skylight reflection from the baffle to the diffuser. The range of
zenith angles accepted by the baffle is from 5 deg (to avoid seeing
the parachute overhead) to 88 deg. Measurements made with the
instrument's optical axis within 85 deg of the Sun direction include
both the diffuse and direct solar beam. A shadow bar (~ 10 deg wide)
is mounted over the diffusing plate. Measurements made with the Sun
behind the shadow bar and again with it out from behind the bar permit
separation of the direct and diffuse downward solar fluxes in the half
of the upper hemisphere centred on the Sun. An additional measurement
is also made with the instrument's optical axis directed 180 deg away
in azimuth from the Sun to give the portion of the downward diffuse
flux coming from this half of the upper hemisphere.
For the measurements in the downward direction, other considerations
become important. Because we want to study the reflection spectrum of
distinct regions on the ground, spatial resolution is as important as
the determination of upward flux. We therefore chose to image the
spectrometer's slit on the ground and to spread the light along the
slit over 20 CCD pixels. Thus, we can spatially resolve 10 regions
along the slit on the surface. The slit is boresighted along the
vertical axis of the imagers' field to provide the morphological
context of the spectrally-measured regions. The slit is 4 deg wide and
extends from 10 deg to 50 deg nadir angle. The 40 deg nadir angle
range covering 20 CCD pixels gives 10 resolution elements, each 4 deg
square on the ground. The spectral intensities measured can be
weighted with cosines of the known nadir angles and integrated in
nadir angle. At least eight different evenly-spaced azimuths are
measured to permit an integration in azimuth for the upward flux.
Thus, the optics and the sampling scheme permit integration of the
spectral radiation field for the upward flux and also provide
spatially-resolved measurements of the surface's spectral
reflectivity.
Upward and downward spectra are obtained roughly every 2 km during the
desceent from 160 km altitude. In the upward direction, the
integration times for the measurements under the shadow bar are
limited to ensure that the measurements are completed while the
diffuser is completely shaded. In the down direction, the integration
times are limited so that Huygens does not rotate more than 1.5 times
the angular width of the spectrometer slit at a nadir angle of 30 deg
(the slit centre) during the integrations. The f/2 system permits S/N
>= 100 in a single spectral and spatial pixel for the downward-looking
spectrometer. For the upward-looking measurements, eight CCD detector
pixels are combined along the length of the slit (which has no spatial
resolution for the ULVS) to give similar S/Ns for the ULVS and DLVS.
 
IR spectrometer
---------------
The IR spectrometer has two channels: the Upward-Looking Infrared
Specrometer (ULIS) and the Downward-Looking Infrared Spectrometer
(DLIS). There are two entrance paths, one viewing upward and one
downward, in order to measure both the upward and downward fluxes. The
ULIS has a diffuser as the entrance window (as for the ULVS), viewing
about half of the upward hemisphere. This allows the measurement of
flux directly, without the uncertain conversion from intensity to
flux. The downward-looking input has a fused silica window and input
lenses that define the FOV at the spectrometer's entrance slit. The
input optics of these instruments transfer IR radiation through fibre
optics to two aligned 104 x 300 um entrance slits placed one after the
other. The two slits of the Up and Down channels are perpendicular to
the grating's dispersion plane. A shutter is installed just inside the
input slits. An f/2 collimator lens assembly focuses the beams on to
the transmission grating, and a camera lens assembly focuses the
dispersed light on to the IR focal plane array. The IR spectrometer
uses transmission gratings working in first order with achromatic
lenses. A polarisation compensator, consisting of a coated tilted
plate, is included in the collimated beam in front of the grating.
Both the upward- and downward-looking spectra are imaged directly on
the two linear arrays of InGaAs detectors through a sapphire window
with anti-reflection coating.
The optical FOV of ULIS is defined by the diffuser and external
baffles and is 170 deg in azimuth by 83 deg in zenith angle (from 5
deg to 88 deg zenith). The spectrometer slit of DLIS maps into a 3 x 9
deg FOV (pixel footprint) with the FOV centre pointed 20 deg from the
nadir.
In the case of ULVS, the exposure time is sufficiently short that an
adequate exposure is possible under the narrow shadow bar even when
the Probe is rotating at its maximum 15 rpm. In the case of ULIS, the
total integration time is of the order of 60 s. Thus, no shadow bar is
used to separate the direct and diffuse downward flux. Instead,
measurements are accumulated in four azimuthal sectors relative to the
Sun, and the shade provided by Huygens itself is used to separate the
direct and diffuse downward flux in ULIS. In the normal operation
mode, the total time for data collection is an integral number of
Probe rotations, but is always constrained to be 1-3 min. The time for
data collection is estimated at the beginning of each cycle, and is
updated during the cycle based on the data from the Sun Sensor after
each rotation.
Within each azimuthal region, data are collected alternately with the
shutter closed and with it open. The process is symmetrical in time so
that linear drifts in temperature (and hence dark current) are
accounted for when the shutter-closed data are subtracted from the
shutter-open data. The total shutter-open time is equal to the
shutter-closed time. Within each azimuth region, the process begins
and ends with the shutter closed for half the time used for the other
open and closed intervals. Within each shutter-open or shutter-closed
interval, the IR array can be read more than once (depending on
rotation period and temperature) to avoid saturation by the dark
current. In addition, the shutter rate is kept to about 5 Hz so that
temperature variations beyond linear drifts and their changes in dark
current are adequately compensated.
 
 Solar Aureole camera
 --------------------
The Solar Aureole (SA) camera will measure the intensity of scattered
sunlight over a range of scattering angles sensitive to particle size.
Scattering at small angles (< 40 deg, depending on particle size) is a
good diagnostic of particle projected area. The forward-scattering
lobe is more strongly peaked at small scattering angles as the
particle radius increases. Its width does not depend strongly on
particle shape for equal projected area particles, and it is
insensitive to the particle refractive index.
Particle shape effects are much more pronounced at intermediate and
large scattering angles, especially in polarisation. The SA camera
will measure the vertical and horizontal polarisation in the Sun and
anti-Sun direction as Huygens spins. These measurements will be
performed at two wavelengths (500 nm and 935 nm) to provide
sensitivity over a wide range of altitudes and particle size.
The SA camera's angular sampling characteristics are listed in Tables
1 and 3. Measurements at small solar angular distance are made when
the Sun is behind the shadow bar, as viewed from the SA camera window.
The solar zenith angle during descent is near 50 deg. In the sunward
direction, the camera covers scattering angles of 5-25 deg. During the
measurement in the anti-Sun direction, the scattering angle coverage
is 75-125 deg. on a 1 deg grid. Each measurement samples the total
column above the Probe. The aureole contribution from a 2 km layer of
the atmosphere is found by tking the difference between two
measurements.
When measurements are collected near the Sun, the shadow bar prevents
direct sunlight from striking the solar aureole window assembly. A
beam splitter is used to provide input from a calibration lamp.
Calibrations are performed several times during the cruise to Titan
and occasionally during the descent. They are made in such a way that
signal from the atmosphere can be subtracted from the total.
A telecentric micro-lens design provides wide angular coverage at f/2
for the solar aureole system. The polarisers used for the visible and
near-IR channels are quite pure; the polarisation they impart to
unpolarised light will be within a few percent of 100%. The Mueller
matrices of all the optical elements will be calibrated before launch,
allowing an accurate measurement of the intensity of vertically and
horizontally polarised light during descent. Optical fibre ribbons are
located at the foci of the camera lenses. They merge with fibre
ribbons from the other visible optical elements to feed the focal
plane.
The use of solar aureole measurements to retrieve particle size
distributions is well known, and several retrieval algorithms have
been described in the literature (see, e.g. [DAVE1971];
[NAKAJIMAETAL1983]). Accurate retrievals require a high signal/noise
ratio. The SA camera was designed to achieve S/N >= 100.
The 6 x 50 deg solar aureole FOV maps onto a 6 x 50-pixel area of the
CCD. Calculations show that S/N will be >100 provided that six pixels
are summed. Pixel summing will be performed during data analysis by
summing pixels along arcs with constant angular distance from the Sun.
In this way, no angular resolution will be lost to pixel summing. The
expected S/N varies from ~ 100 to ~ 900, depending on wavelength,
altitude, solar zenith angle and azimuth during the measurement
interval between 160 km altitude and the surface.
 
 Violet photometers
 ------------------
The purpose of the violet photometers is to extend the spectral range
of the measurements to 350 nm from the visible spectrometers' short
wavelength limit of 480 nm. As no sharp spectral features are expected
in this spectrum, a single spectral channel is used. In this spectral
region, the spectrum of the incident sunlight is modified by the
absorbing properties of the small photochemical aerosols in Titan's
stratosphere. These particles absorb more and more of the short end of
the spectrum with increasing depth into the atmosphere. Hence, in
order to measure the energy deposition, it is important to filter the
response of the detectors to make them relatively spectrally flat so
that the violet measurements can be converted to absorbed energy
independent of assumptions regarding the spectrum of the light
measured. We specify that the product of the spectral response of the
violet detectors and the transmission of the violet filters to be
spectrally flat to 10% from 350 nm to 480 nm for these measurements.
The Upward-Looking Violet (ULV) photometer shares the ULVS system
diffuser: its FOV is identical to that of ULVS. The shadow bar permits
separation of the direct and diffuse downward flux in the violet in
the same way as for the ULVS, and the sampling in azimuth is the same
for the ULV and the ULVS systems. The DLV uses exactly the same baffle
and diffuser design mounted on the bottom of the sensor head to
measure the upward violet flux, but without the shadow bar. The DLV's
sampling in azimuth is limited to two regions, near the Sun's azimuth
and near 180 deg from the Sun. These two measurements sum to give the
total upward violet flux.
For the sake of simplicity, the electronics for the ULV/DLV systems
use the same 12-bit A/D converter used for several other housekeeping
functions, including voltage and temperature monitoring. The light
level in the upward-looking violet channel at the top of the
atmosphere is known from the solar spectrum. A single gain with a
fixed time constant is used in the ULV system with no provision for
gain change or change in integration time in flight. The gain is
adjusted when the instrument is built so that the sum of the signal
from the inflight calibration signal and the downward solar flux at
Titan occupy between half and three-quarters of the 4096 available
digital steps. The signal from the inflight calibration system is
adjusted to be roughly equal to the signal expected from the Sun at
descent start. This gives about 1000 digital steps for the violet
signal at the start of descent. The random noise in the ULV channel is
comparable to the size of a digital step. We can measure the
absorption of the violet energy in the ULV's bandpass during descent
to a few tenths of a percent of the value at entry.
 
 Sun Sensor
 ----------
The timing of all of the upward-looking DISR measurements requires
precise knowledge of the solar azimuth angle. To satisfy this need, a
Sun Sensor is located beside the shadow bar. The sensor's FOV is a 64
deg cone centred 47 deg from the zenith. The SS consists of a window,
lens system, filter, reticule, a pair of lenses to focus the light on
the silicon photodiode detector, and its associated readout
electronics. As the Sun crosses the FOV, three large pulses are seen
as its image enters the three slits of the focal plane reticule. A
detection circuit locates it in the central slit to define zero
azimuth angle. The first and third slit are tipped with respect to the
centre slit to permit determination of the Sun's zenith angle. The
times of all three slit crossings are transmitted in housekeeping. The
times between the three crossings compared to the times between
successive crossings of the centre slit (the rotation period) give the
Sun's zenith angle and the Probe's attitude on each rotation.
In addition, the brightness of the central pulse in each rotation past
the Sun is included in the telemetry to give a measure of the direct
solar beam at the wavelength of the SS system (939 nm) as a function
of altitude. The SS is designed to track the direct solar beam down to
1/1000 of its brightness outside Titan's atmosphere. At the nominal
solar zenith angle of 50 deg, the Sun signal would be lost at a
vertical extinction optical depth of 4.4. This compares to a vertical
optical depth of less than about 1 in some current Titan models at the
Sun Sensor's wavelength.
The SS uses several logic tests on the relative timing of the three-
pulse series on successive rotations to distinguish pulses due to the
Sun from variations in intensity that might be due to diffuse clouds
above the Probe. If Huygens falls beneath a thick cloud and loses
lock, the internal rotation rate transmitted by the Probe's spin
sensor is used to time DISR data collection. DISR's Sun Sensor
continues to look for the Sun after losing lock, and will find it if
it reappears from behind a passing cloud.
 
 Internal calibration system
 ---------------------------
Almost all of DISR's science goals require measurements to be made on
an absolute photometric scale. This is difficult enough for a single
optical system and detector, let alone for a system as complex as
DISR. In principle, the measurements of the Sun's apparent brightness
at high altitude near the start of the descent at some IR wavelengths
may be relatively unaffected by atmospheric scattering and absorption,
and so may provide an important check on the absolute calibration of
this one DISR optical system. The ability to tie all the various DISR
systems together on a single photometric scale is provided by the
internal calibration system built into the DISR Sensor Head. The
internal calibration system consists of a set of three redundant 1 W
lamps that illuminate one end of a bundle of thick quartz fibres. Each
fibre carries some light from the calibration lamps to the first
optical element in each of the other DISR optical systems. This light
is reflected from either the inside of the window for that system or
from a small reflector mounted on the back of the window and passes
through the rest of that optical system to the detector. The fraction
of the light from the calibration lamps carried to the input of each
DISR measuring sub-instrument is carefully measured before launch, and
expected to remain constant throughout the mission. We should be able
to track changes in the sensitivity of detector pixels, in the
transmission of the imaging fibre optic conduit or in other optical
elements, during cruise by monitoring the relative outputs of each
DISR system to the light provided by the internal calibration system.
It is important to emphasise that we are not relying on the stability
of the absolute output of the internal calibration system lamps, only
on the stability of the relative fraction of the light carried to each
DISR subsystem by the thick quartz fibres of the calibration system.
These fibres are not expected to be susceptible to darkening from
energetic particle bombardment. Thus, the changes in the relative
output of the different DISR optical pixels can be tracked with time
after launch and even measured several times during Titan entry.
There are no moving parts in the internal calibration system, so the
light from this system is added to that from the ambient Titan
atmosphere during descent. The light from the calibration system is
designed to exceed that from the ambient atmosphere by a large factor.
In addition, we have designed the internal calibration procedures to
collect light with the calibration lamps on, then off, and then on
again at the same azimuth in each case to permit relatively accurate
subtraction of the ambient signal from the signal provided by the
calibration system. The calibration light levels are adequate to make
measurements with S/N better than the target value of 100 for
measurements of the ambient atmosphere.
 
 Surface Science Lamp
 --------------------
The purpose of the Surface Science Lamp (SSL) is to illuminate Titan's
surface in spectral regions where strong atmospheric absorption bands
prevent sunlight from penetrating to the surface. The SSL permits
continuous measurements of the surface's spectral reflectivity to be
made throughout the entire spectral range. The SSL is a 20 W lamp with
a parabolic reflector that illuminates the surface and fills the
narrow 3 x 9 deg FOV of the IR spectrometer with enough light to give
a S/N of 50 at 60 m altitude within the strong methane bands even if
the surface reflectivity is as low as 0.05.
The lamp system is activated at 400 m altitude (given by the radar
altimeter) and is operated during the last several minutes of descent,
when a continuous sampling of the surface reflection spectrum is
obtained using both the DLVS and DLIS.
 
Mechanical design
=================
 
The DISR instrument consists of two packages: the Sensor Head and the
Electronics Assembly. These are mounted on top of the Probe's
instrument shelf. The SH is mounted on a metal bracket above the
instrument shelf and sufficiently far outboard so that its front
protrudes through the Probe's back side. This permits clear views from
the zenith to the nadir and in directions to +- 85 deg in azimuth from
the radial direction.
The EA is a straightforward box containing six boards mounted
horizontally above the power supply that is built into the chassis
base. The six boards, each 12.7 x 20.8 cm stacked 1.5 cm apart,
include the CCD drive board provided by MPAE, the data compression
board provided by TUB, the IR detector system driver board provided by
the Paris Observatory, and the digital board, CPU board and auxiliary
board provided by Lockheed Martin. The EA's 4.4 kg mass includes
radiation shielding on particular electronic parts. The EA's box is
bolted together from slab sides, top, front and back above the base.
The six boards are an integral part of the structure and are bolted to
the chassis side walls and the back wall.
The SH's mechanical design is a much more challenging task than
packaging the EA. The SH assembly must hold the detectors, fibre
optics, foreoptics for the visible and IR spectrometers, the three
imaging cameras, the Solar Aureole camera, the internal calibration
system and the Surface Science Lamp all in precise alignment over a
temperature range of some 150 deg C for 10 years, and throughout the
mission's vibration and shock environment. In addition, the detectors'
thermal environment must be controlled during descent. The detectors
must be protected from the energetic particle environment during the
mission. Finally, the size and mass of the package are severely
constrained as well.
The SH can be thought of as having a CCD detector side and an IR side.
The face of the CCD is held 20 um from the end of the fibre optic
conduit that delivers light from the three imagers, Solar Aureole
input assembly and visible spectrometers.
The IR half of the Sensor Head begins with the pair of IR array
detectors and their pre-amp board. This detector is mounted to the IR
spectrometer, shutter and fibre input assembly.
The two detector assemblies are surrounded by a tungsten radiation
shield 4 mm thick, sufficient to prevent protons of <64 MeV from
reaching the detectors, significantly reducing the radiation dose.
The optical systems are securely mounted to a titanium optical bench
that is secured to the base of the SH housing using a ball joint and
two slip joints. This bench ties all the optical systems to the
detectors.
The DISR thermal design must maintain interface constraints, cool and
maintain the detectors' temperature within limits, maintain all other
components within their temperature limits, and minimise mass and
power. In particular, it is important that the detectors are cooled as
rapidly as possible from their temperatures near 260K at Titan entry
to below some 220K in order to minimise the influence of dark current
and to ameliorate the defects in CCD pixels that may have been damaged
by energetic particle impact during cruise.
The detector cooling system includes a thermal strap to provide a heat
sink to the atmosphere via an attachment to the Probe. A heater
ensures that the detectors will remain above their minimum temperature
limit, thermostatically controlled to maintain > 160K. The housing
coupling and the detector components masses are minimised to allow a
rapid cooldown. The temperatures of the PC boards in the SH are
maintained with a second electronic thermostatically-controlled
heater.
The SH housing is constructed of aluminium to provide interior-
exterior thermal coupling. This will maintain the exposed windows well
above the ambient atmospheric temperature. The exposed portion of the
housing is covered with a low density insulation and coated with a
conductive paint. The interior housing is conductively isolated from
the Probe but radiatively and convectively coupled to it. This allows
Huygens to provide some heating of the SH 'warm' components. The SH
front housing is thermally isolated from the rear housing.
 
Compression of imaging data
===========================
 
Imaging data are compressed onboard before transmission in order to
achieve a balance between the number and the quality of images
received. The imaging data are digitised at 12 bits/pixel. They are
first divided by an onboard flat field to eliminate artifacts
introduced by variations in the transmission of the fibre optic
conduit. They are then reduced from 12 to 8 bits/pixel using a pseudo-
square root software algorithm. They are then passed to a lossy
hardware data system for further compression.
Simulations with a series of test images have shown that, in
principle, lossy compression by the standardised JPEG scheme is
applicable. For most of the scenes, the image distortions are
relatively small for an additional compression by a factor <= 8.
Accordingly, the JPEG algorithm has been selected as the baseline.
Implementation of such a compression by software running on the main
processor would limit the image cycle time to > 1 min. Therefore, a
dedicated hardware coprocessor has been developed which (1) performs a
complete compression of a 256x256-pixel image within < 130 ms, (2)
dissipates about 120 mWs per compression and < 200 mW in idle mode,
and (3) is implemented on a single multilayer board of 215 g. The
coprocessor is designed around an STV3200 DCT chip, an 80C86
microprocessor and eight ACTEL FPGAs. The compression algorithm
running on this dedicated coprocessor deviates in some aspects from
the JPEG scheme in order to make the hardware less expensive in terms
of power and mass, and to enhance the robustness against transmission
errors.
The first very basic operation of the data flow through the
coprocessor is transformation of the image data from the spatial
domain into the frequency domain. The applied Discrete Cosine
Transform (DCT) concentrates most of the image energy in a small set
of highly decorrelated, low frequency coefficients. The primary
compression mechanism is to represent the minority of lumped high
energy coefficients with higher accuracy than the majority of low
energy coefficients. Even deletion of most of the low energy
coefficients removes only a small portion of image information.
Accordingly, the coefficients below a given threshold are deleted and
the remaining coefficients are requantised with an adjustable step
width proportional to a value Q. Subsequent zigzag-ordering
concatenates coefficients of similar spatial frequency and generates a
coefficient string where typically a bulk of low frequency/high energy
coefficients is followed by some spurious high frequency coefficients,
which are separated by long intervals of zeros acting as placeholders
for the deleted coefficients. Run length encoding exploits this
specific string structure.
In the next step, the original JPEG scheme applies Huffman encoding.
This has been replaced by a specific scheme less sensitive to
transmission errors at the expense of a slightly reduced coding
efficiency. The processing of the frequency coefficients is iterated
several times in order to approach a target compression factor of
near-optimum image quality, which is expressed by the S/N of the total
image. The first run determines the quantisation step width, which
remains unchanged for all following runs while the deletion threshold
is varied.
The S/N (in the image intensities) of the modified hardware
implementation is <=1 dB below the values achievable with standard
JPEG. In flat image regions, blocking effects become visible at lower
compression factors compared to JPEG. These very limited deficiencies
in image quality are regarded to be outweighed by the associated
savings in power and mass. On the LENA test image (a standard image
commonly used for testing image compression systems), the S/N varies
between about 50 and 30 at compression ratios of 3:1 and 6:1,
respectively. On simulated low contrast images as expected for Titan,
the noise is increased above shot noise by factors of only some 1.7
and 2.5 for compression ratios of 3:1 and 6:1, respectively.
 
Electronics
===========
 
The SH contains three small boards: one provided by MPAE contains the
CCD's pre-amps; one provided by PO contains the IR detector's pre-amp;
and the third, from MMC, provides pre-amps for the Sun Sensor and its
test light emitting diode, as well as for the violet detectors of the
ULV and DLV instruments. The European Co-Is also provided the EA's
three shaded boards: the CCD's driver (MPAE); the IR detector driver
(PO); and the hardware Data Compression System (TUB).
The CPU board provides radiation-hard critical RAM in 128 kbytes of
program RAM and an additional 64 kbytes of data RAM. It also holds 128
kbytes of PROM and a separate block of 64 kbytes of EEPROM, which can
be changed by commanded uploads. The microprocessor is a MA31750
running at a clock speed of 12 MHz capable of some 1.6 million
instructions/s (MIPS). Some 1.2 MIPS are required for operation of the
DISR as planned. A watchdog timer is provided with a 100 Hz clock at
16-bit resolution.
The Digital Board contains a large (1.5 Mbyte) frame buffer of static
RAM accessible by the CPU or by any of three Data Management Assembly
(DMA) channels. The entire frame buffer is refreshed every 10 s by CPU
block read/write routine using a full Error Detection And Correction
(EDAC) function provided by a single Harris Semiconductor chip.
Three programmable DMA channels are provided using Actel 1020 Field
Programmable Gate Arrays (FPGAs). The mode, word count, source and
destination addresses are all programmable.
The Probe interface provides dual telemetry packet channels for
transmitting the science data. It provides dual Memory Load Command
channels to receive the data and/or commands from Huygens. Dual serial
status channels are also provided to transmit a 16-bit housekeeping
status word to the Probe upon request. Redundant channels are also
provided to receive the Probe Data Broadcast data such as spin rate,
time and altitude. This interface also provides a means of receiving
an indication of which telemetry channel is to be used for commanding
the DISR.
The Auxiliary Board contains the digital interface to the IR detector
system. It also contains several analogue circuits to: condition
signals from the ULV, DLV, Sun Sensor and the Inflight Calibration
System; control the focal plane and SH electronics card heaters;
monitor temperatures, lamp currents and other housekeeping data in the
instrument. A multiplexer and a 12-bit A/D converter are used on the
Auxiliary Board for this purpose.
 
Software and data collection modes
==================================
 
The DISR flight software was developed using an object-oriented design
in the ADA language. The software uses a re-entrant event dispatcher
to control execution based on the priorities of events occurring in
both the hardware and software. Multi-tasking is not used. Hardware
interrupts are used to provide services for the Probe interface, Sun
Sensor, general purpose event timer, telemetry channels, direct memory
access controllers, CCD, IR detector and hardware data compressor. The
software controls the calibration and surface science lamps. The
calibration lamps are turned on during appropriate parts of
calibration cycles. All commands to the DISR are processed by the
software. Only six commands exist, although some have a variety of
parameters.
  1. A receipt-enable telecommand must begin a commanding session.
     This command is used to protect against spurious commands
  2. A change-mode telecommand may be used to change the DISR's
     operating mode into descent mode (the default mode), calibration
     mode, single telecommand mode or memory access mode
  3. Single measurement telecommands direct the instrument to perform
     one or more repetitions of a particular measurement. These
     commands are useful during instrument calibration and test
  4. Single test telecommands are similar to single measurement
     telecommands, except they initiate preprogrammed test sequences
     on the specific portions of the hardware, including the IR
     shutter, hardware data compressor, heaters and lamps
  5. Memory upload commands are used in memory access mode to store
     new tables that are read by the software. These table entries
     include bad pixel maps, square root compression tables, and
     parameters that control measurement scheduling and processing
  6. Memory dump telecommands can transfer any portion of DISR memory
     into telemetry for verification.
The software also coordinates and controls all data collection.
Optimum exposure times are computed for each sub-instrument using the
CCD and IR detectors. These times are based on the data number
population histograms of the most recent previous exposure of the same
type. The exposure time can also be limited by the amount of smear
caused by the Probe's spin.
Onboard data processing functions also include several miscellaneous
functions. Adjacent pixel columns may be averaged within the same
instrument FOV. Data for the hardware data compressor must be
reformatted before it is fed to the compressor. Lossless compression
is done entirely in software. Bad pixels are eliminated according to a
bad pixel map stored in EEPROM. Data from the imagers are also reduced
from 12 bits to 8 bits before being fed to the hardware data
compressor. This reduction is done using a table lookup that performs
a pseudo-square root transformation of the raw data. A watchdog timer
can reset the microprocessor if it times out. The software that builds
telemetry packets periodically resets the watchdog timer. If telemetry
is not being produced, the processor will be reset and execution will
be restarted. Calibration and instrument health data are collected at
6-month intervals during cruise. One of five possible activities may
be performed at each opportunity. The Health Check sequence exercises
each software-controlled function to test for normal operation. The
In-Flight Calibration sequence is used to obtain relative response and
measure the noise level of the detectors. The Simulated Descent tests
descent sequencing using the parameters loaded into special tables.
Two types of activities are available for contingencies. Single
Measurement and Single Test commands can be used to diagnose problems
with specific instrument subsystems. Finally, memory dumps and uploads
can be used to verify the contents of memory and to upload new table
values to correct for instrument subsystem malfunctions.
One of the software's more interesting aspects is the way it optimally
schedules collection of related data from various sub-instruments.
During descent, data collection is divided into cycles that form a
coordinated set of measurements within a limited time span. Several
different types of cycles are used to meet the differing needs of data
collection during different parts of the descent. The software uses
information from Huygens, including the altitude, time and spin rate,
to decide which type of cycle to start next. The choice also depends
on the amount of buffer space available within DISR and on the current
Probe telemetry rate. The software chooses a cycle type that will
gather sufficient data to keep the instrument from running out of
telemetry packets, yet it must also choose one that will not provide
so much data that it overfills the available buffer space.
Differing requirements are placed on the data buffering scheme
depending on how close Huygens is to the surface. Normally, data can
be gathered much faster than they can be telemetered. If the Probe is
still high enough to allow sufficient time for telemetry before
impact, data may be profitably buffered, allowing the buffer to fill
further as each cycle passes. On the other hand, if Huygens is near
the surface, each data set should be telemetered as soon as it is
complete in order to prevent the Probe from hitting the surface with
substantial amounts of data still in the buffer. Early in the descent,
Huygens is falling rapidly. It is important to measure quantities such
as the solar flux deposition and the solar aureole profile within
vertical intervals that are small compared to the atmospheric scale
height. Cycles high in the atmosphere are driven by a cycle duration
constraint, even if many cycles spend a long time in the buffer before
they are telemetered. However, at some altitude it is important to
start each new cycle with almost nothing in the buffer. The most
recent, lowest altitude data will be transmitted immediately if the
buffer is in this state.
We use both of these scheduling algorithms during the descent. Above
20 km altitude, we use duration-constrained scheduling. At 20 km, we
pause to permit all previously buffered data to be telemetered. After
this, new cycles are not begun until only a few telemetry packets
remain in the buffer.
Within data cycles, it is usually important to schedule measurements
at particular azimuths. Images, for example, must be scheduled to
occur at 12 equally-spaced azimuths to provide data for a mosaic. The
software determines the azimuth at the beginning of each cycle and
predicts its change with time during the cycle using a quadratic
extrapolation. As new information about the spin rate comes in from
the Sun Sensor, these predictions are updated during the cycle and
measurements are scheduled using them.
 
Measured Quantities
===================
 
The instrument is turned on near 170 km altitude. In order to fill the
initially empty telemetry buffer as rapidly as possible, the
instrument begins with a data collection cycle that contains an image
panorama. The data collected in cycles containing image panoramas are
shown in Table 3. Note that on alternate cycles of this type, the
images' azimuths are offset by 15 deg, half the width of the images.
These cycles contain a full set of spectral data (with downward-
looking visible spectra in the centre of each image) and solar aureole
data as well. Above 20 km, these cycles are constrained to collect all
their data in 1.5-4.5 min.
 
Table 3. Image cycle measurements.
----------------------------------------------------------------------
Sub-instrument        Azimuths (degrees)
----------------------------------------------------------------------
Solar Aureole         1.5 (under the shadow bar), 174
ULVS and ULV          5.5 (under the shadow bar), 140, 180. 320, 338
HRI, MRI, SLI, DLVS   2, 32, 62, 92, 122, 152, 182, 212, 242, 272,
   302, 332 or 17, 47, 77, 107, 137, 167, 197, 227, 257, 287, 317, 347
Dark current          180
ULIS            315-45, 290-315 and 45-70, 90-135 and 225-270, 135-225
DLIS                  0-45, 45-70, 90-135, 135-180, 180-225, 225-270,
                      290-315, 315-360
DLV                   0, 180
----------------------------------------------------------------------
 
Above 20 km altitude, a non-imaging cycle is chosen as the next cycle
if the DISR telemetry buffer contains data that will take more than
four additional minutes to transmit. Non-imaging cycles are
interleaved between imaging cycles to produce spectral measurements at
higher vertical resolution than would be possible if all cycles
contained image panoramas. The time for data collection in non-imaging
cycles is limited to 1.5-3 min. The maximum cycle duration prevents
long gaps between observations that might occur at very slow Probe
rotation rates where a cycle might take a very long time to complete.
The minimum cycle time forces the scheduler to wait before the
beginning of the next cycle and is used to tune the relative numbers
of imaging and non-imaging cycles. Increasing the minimum cycle
duration allows the telemetry buffer to be emptied further during the
cycle, and increases the likelihood that an imaging cycle will be
chosen next.
The standard non-imaging cycle differs from the imaging cycle in three
respects:
  1. six (rather than 12) DLVS spectra are obtained evenly distributed
     in azimuth beginning 20 deg from the Sun
  2. no HRI or MRI images are obtained
  3. six SLI images are obtained evenly distributed in azimuth
     beginning 15 deg from the Sun's azimuth.
The SLI data in the central 26 columns of each image are summed to a
single column of brightness as a function of zenith angle, losslessly
compresed and included in the telemetry stream.
Special calibration cycles are performed four times above 20 km at
roughly equally-spaced detector temperatures. These data use the
onboard calibration lamps to provide signals dominating those of
Titan's. These cycles provide relative calibrations of the upward- and
downward-looking sub-instruments as well as corroboration of the flat
field properties of the CCD and IR detectors. A special long exposure
measurement is also made with the IR sub-instrument while the shutter
is closed for accurate determination of the dark current generation
rate for each pixel.
At 20 km altitude a special 'pause' cycle occurs in order to empty the
telemetry buffer. Imaging cycles are then gathered between 20 km and 3
km, with each new cycle beginning only when the telemetry buffer
drains to a preset low level. This sequence is interrupted twice (near
9 km and 4 km altitude) for the Spectrophotometric Cycles when
downward-looking spectra are obtained as rapidly as possible for one
Probe rotation. These data map the surface's spectral reflectance at
high spatial resolution.
If normal image cycles were continued below 3 km, telemetry of the
data would not be completed before surface impact. Therefore, between
about 3 km and 500 m, special cycles are used in which a set of HRI,
MRI, SLI images, spectra and solar aureole measurements are obtained
as rapidly as the telemetry rate permits, regardless of the azimuth.
Between 500 m and 200 m, single HRI images are obtained as rapidly as
the telemetry rate permits (about every 8 s). Single HRI images and
dark current measurements are the only data collected in this altitude
interval. In the last image, near 200 m, the scale will be about 20
cm/pixel.
The Surface Science Lamp is turned on at 400 m altitude. For the final
200 m, we obtain only DLV, DLVS and DLIS measurements. Some 19 spectra
are obtained with the SSL turned on. These yield continuous surface
reflection spectra over the DLIS spectral range for constraining the
surface composition. As Huygens continues to rotate slowly while these
exposures are collected, each will view a different surface region.
Some of the 20 pixels along the DLVS slit will be illuminated by the
SSL, and these are telemetered to the ground, permitting the DLIS
spectra to be extended to shorter wavelengths. The S/N of these
spectra increase as 1/(altitude)^2, reaching more than several hundred
for ground reflectivities exceeding several percent.
If the instrument survives impact, it will collect data according to a
special sequence. One purpose of these cycles is to measure transient
phenomena such as reflection from any cloud of dust or spray that may
result from Probe impact. A set of 10 upward- and downward-looking
spectra are collectedfirst using the ULV, ULVS, ULIS, DLV, DLVS and
DLIS sub-instruments. Following these spectral measurements, images
are taken with the three cameras. The SLI should give images from the
foreground to 6 deg above the horizon. If Huygens survives long
enough, it may be possible to see cloud features blown by the wind.
Spectra and image measurements sets are then gathered alternately for
the remainder of the mission. Approximately 2 min after impact, the
SSL is switched off. After that, it is alternately switched on and off
every 2 min. All measurements after impact are taken as rapidly as
possible until the buffer fills. Then, new measurements are gathered
as buffer space allows. A summary of the data collected in various
parts of the descent is given in Table 4.
 
Table 4. Summary of descent data.
----------------------------------------------------------------------
                               3-Image High-Res  Spectra  3-
Altitude    36-Image Non-Image  Sets +  Images    Using  Image  Spectr
Range (km)   Cycles   Cycles   Spectra   Only     Lamp   Sets     Sets
----------------------------------------------------------------------
160-20         10       33        -       -        -       -       -
20-3            5        -        -       -        -       -       -
3-0.4           -        -       20       -        -       -       -
0.40-0.20       -        -        -       2        -       -       -
0.20-0.0        -        -        -       -       18       -       -
After Impact    -        -        -       -        -      83      93
(first 10 min)
----------------------------------------------------------------------
Total number of images: >600. Total direct, diffuse, upward, downward
sets of flux measurements: 68. Solar aureole measurements: 48 in each
of two colours and two polarisation states. Surface spectra: 642 in
IR, >3100 in visible
 
 
Expected Results
================
 
The ability of the DISR measurements to support the science goals
outlined at the beginning of this paper is summarised here. For
studies of the thermal balance, the downward total flux is obtained by
summing the measurements at two azimuths 180 deg apart. The
corresponding upward flux in the violet is obtained by a similar sum
of the DLV measurements. In the visible and near-IR, the azimuth
integral for the upward flux is done by averaging the measurements at
all azimuths in the DLVS and DLIS. For the DLVS, the data include
measurements at a range of nadir angles along the slit. These data are
weighted with the cosine of the nadir angle and integrated in nadir to
give the upward flux. For the DLIS, a model is required to connect the
DLIS measurement to give the upward flux in the near-IR. The upward
flux is subtracted from the downward flux at each wavelength to give
the net flux at each altitude. The difference in the net flux between
two altitudes gives the energy absorbed by the intervening atmospheric
layer. The integral of this quantity over wavelength, divided by the
layer's mass and the specific heat capacity of the layer's gas, gives
the solar heating rate for this layer. The radiative cooling rate is
computed from the IR opacity, computed from the atmospheric gases and
the particles. The particles' IR opacity is computed from their
optical properties and distribution determined by other DISR
measurements.
The net dynamical forcing can be used in dynamical models to evaluate
the wind field, and then compared to the Probe's drift rate as
determined from the apparent motions of features in surface images.
Fig. 16 shows the coverage of HRI and MRI panoramas over the surface
in response to a nominal wind model [LUNINEETAL1991].
Determination of the properties and distribution of aerosol and cloud
particles is somewhat complex. The direct downward flux at each
wavelength from ULV, ULVS and ULIS is obtained by subtracting the
diffuse flux from the total diffuse plus direct flux. The natural log
of the ratio of downward fluxes at two altitudes, divided by the
cosine of the local solar zenith angle, gives the vertical extinction
optical depth of the particles between the two levels at that
wavelength. The variation of this quantity with wavelength gives a
measure of the volume-weighted size of the particles. The solar
aureole measurements near the Sun give the size of the forward-
scattering peak of the single-scattering phase function of the
particles above that altitude. At high altitudes, the optical depth in
the 500 nm channel is used to determine the width of the forward
scattering peak. This peak is caused by diffraction, and its size
depends on the mean projected area of the particles. At lower
altitudes, where multiple scattering in the blue above Huygens is
large, the data in the 935 nm solar aureole channel is used. At
intermediate altitudes, the ratio of the diffraction peak in the two
wavelengths gives a check on the area-weighted size of the aerosols
above Huygens. The solar aureole measurements 180 deg from the Sun
give the shape of the single-scattering phase function and its
polarising properties at large scattering angles. The single-
scattering polarisation provides a third measure of particle size.
The particle shape is constrained by the projected-area weighted
radius, the volume-weighted radius and the polarisation radius. Models
of aggregate particles with different-sized monomers and different
number of monomers per particle are compared with the particles' size
constraints. Comparison of calculations for the determined shape with
the measurements give the particle scattering and extinction cross
sections at each wavelength. Comparison with the variation of
extinction with altitude gives the number density as a function of
altitude. Comparison with the sedimentation rate calculated for these
particles gives the particle mass production rate in steady state. The
single-scattering albedo of the aerosols is determined by the change
in the net flux at wavelengths where atmospheric gases (primarily
methane) do not absorb. The variation of single-scattering albedo with
wavelength and the extinction optical depth constrains the imaginary
refractive index of the particles as a function of wavelength.
The surface topography is determined both from stereographic analysis
of images obtained at different altitudes and by surface photometry
and the Sun's known zenith angle. Liquid surfaces can be distinguished
by their lack of vertical relief, uniform reflectivities, specular
reflections and comparison of their reflection spectra with those of
expected mixtures of methane and ethane using the spectra obtained
with SSL illumination in the last stage of descent. These reflection
spectra can constrain surface composition even if the surface is not
liquid. The topography and details of the features seen in the images
can reveal the physical properties responsible for forming the
surface.
Finally, the profile of methane's mixing ratio is determined from the
growth of the many methane features of various strengths observed
throughout the descent. The argon mixing ratio is constrained by the
methane/nitrogen mixing ratio and the total molecular weight
determined by radio occultation data from the Orbiter and the
temperature measurements made by Huygens. The DISR spectra will be
examined for the presence of other minor constituents as well.
In summary, the DISR suite of measurements has been highly tuned to
provide many complementary types of optical measurements designed to
contribute to our understanding of the thermal balance, nature of the
cloud and aerosol particles, composition and physical state of the
surface, and composition of the atmosphere.