Imaging Polarimetry with UIST
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Imaging Polarimetry with UIST
This document describes the use of IRPOL2 with the
facility imager/spectrometer UIST.
Details concerning the design of the
polarimetry module and its corresponding optics can be found in the page on
Optical characteristics. The acquisition and reduction of
UIST imaging-polarimetry data, as well as results from recent
calibration/characterisation observations are briefly described
is discussed on a separate page.
Users of IRPOL2 with UIST may also find the IRCAM3 Polarimetry
manual written by Ant Chrysostomou of interest. It contains a number
of very helpful tips and comments concerning imaging polarimetry in general.
A postscript version is available
The UKIRT and University of Hertfordshire would appreciate an
in any publication which contains data obtained using IRPOL2.
Imaging Polarimetry with UIST: Data Acquisition
IRPOL2 comprises a half-wave retarder (the waveplate) and,
internal to UIST, a focal-plane mask and a Wollaston
The mask comprises two parallel apertures that are each 20"x120"
in size. Although the orientation of these apertures can be adjusted
on the sky, we recommend using the default UIST-imaging position angle
of -90 degrees (if you require a different angle please discuss
this with your support scientist, or with Chris Davis [c.davis at
jach.hawaii.edu]). With this angle the long axis is N-S. The prism
splits the radiation from each aperture into orthogonally-polarised e-
and o-beams, which are projected onto the array (see Fig.1). Without
the mask, these e- and o-beams would overlap.
For the pipeline DR to work (described below), the target should be
placed in the top half of the array; the two strips in the bottom half
of the array (west for PA=-90), from an adjacent area on the sky, are
then used as blank sky for subsequent sky subtraction. An example raw
exposure of a standard star, showing the e and o "images", is shown in
Fig.1 A Raw image through the Wollaston prism, showing the
e- and o- beams in the eastern (top) half of the array. Note that
UIST images come off the telescope with
N-left and E-up.
Preparing an Observing Programme with the UKIRT-OT (Phase 2)
The expectation is that most users will work from a "Template
Sequence". The "Template Library" contains a number of
sequences specific to polarimetry. These can be modified to
suit your needs. It is probably unwise to try and write a sequence
completely from scratch. Generally, an imaging-polarimetry MSB
contains two or three separate observations, the first for a sky flat,
the second for a polarised or unpolarised standard star, and the third
for the target itself. The standard and science target are usually
observed in the same manner:
1. Flat Fields:
There are two methods of obtaining a sky flat contained in the
template library; the first "Make_Skyflat for Pol", used in the
example in Fig.2, simply combines 8 frames, 2 at each of the 4
waveplate angles, into a "master-pol-flat" (the 8 frames are all
dithered on the sky). The recipe SKY_FLAT_POL is used to reduce these
data. This kind of flat may be sufficient for your needs, if the
polarisation accuracy needed exceeds 1% or so. If, however, you are
concerned that the instrument flat field may be sensitive to waveplate
angle, then separate flats can be obtained, at each waveplate angle,
using a second template sequence "Make_Skyflat at each WP angle for
Pol". This sequence obtains a 3x3 grid of images at each of the 4
waveplate angles. The recipe SKY_FLAT_POL_ANGLE will then reduce each
At shorter wavelengths, or with narrow-band filters, dome or
twilight flats may be necessary - see the discussion below.
2. Standard and Target Observations:
In addition to the flat-field sequences, the template library also
contains three options for observing your target, which use different
ORAC-DR recipes. The first and second are similar in that
spatially-jittered observations are obtained at the 4 different
waveplate angles (0, 45, 22.5, and 67.5 degrees); however, the "Pol
Jitter then Angle" sequence takes a number of dithered
(North-South) images BEFORE the waveplate is moved. The second
sequence, "Pol Angle then Jitter", follows the "Golden Rule" of
Polarimetry more closely, since images at the 4 waveplate angles are
obtained straight away, i.e. BEFORE the telescope is jittered to the
next position. The benefit of the first is that it is more efficient,
since it involves fewer waveplate moves, although the latter is
recommended if conditions are non-photometric. The third sequence,
"Pol Extended" is meant for extended sources that spill into
the bottom (western) half of the array; in such a case the lower half
of the array cannot be used for sky subtraction, so the telescope is
nodded between the target and blank sky.
Fig.2 A Typical Polarimetry Programme
Above we show a typical polarimetry MSB, containing two separate
observations, for a sky flat and a target. The "Pol angle then jitter"
observation for the latter begins with three
components (the broken blue boxes), which specify the target
information (target and guide star coordinates), the instrument
configuration and the reduction recipe. There are also "notes"
between the components; notes can be added at any point in a sequence
by clicking on the "note-pad" icon in the left-hand tool bar. The
user should click on each of these components to check parameters
and/or add new information; the parameters associated with each
component will be displayed in the right-hand portion of the OT window
(n.b. you can resize the window if necessary!).
- The Target Information component
obviously needs to be edited; click on this to enter source
coordinates AND to specify a guide star. This component may also be
used to display a Digitised Sky Survey image of the target field, the
instrument aperture size, the regions blocked by the waveplate + holder,
and various guide-star catalogues.
- The UIST
component sets the instrument to its initial configuration;
integration time, filter, etc. Note that
you must also select polarimetry in this component (to install the
internal Wollaston prism).
- The DRRecipe component should already
be set correctly to the recipe specific to your chosen mode of IRPOL
imaging, though you should check this.
The sequence in Fig.2 follows the "Pol Angle then Jitter" method
described earlier. The alternative, "Pol Jitter then Angle" (also
available in the template library) is almost identical to the sequence
above, except that the IRPOL and Offset iterators switch places, so
that the telescope is offset three times (north-south) BEFORE moving
the waveplate. Groups of 12 frames are still acquired, however.
Sky flats - a discourse...
With near-infrared astronomy, a good flat-field is of course always
a high priority, although for polarimetry the flat-field is in
principle independent of the final result, since it is the
ratio of observations which are measured. However, one can only
safely assume this if (1) the detector is equally sensitive to e- and
o- rays, and (2) exactly the same pixels are ratioed each time (this
should be the case when guiding, especially if the 4 WP angles are
observed BEFORE the jitters). If
(1) and/or (2) do not apply, then a flat field is needed to determine
of the e- and o-beam sensitivities.
Twilight or Dome flats
At shorter wavelengths, getting a sizable number of counts on the
array becomes an issue. Two possibilities are currently being
investigate; twilight and "dome" flats.
To date we have little experience of taking twilight flats. These
can be difficult to (flexibly) schedule if you do not have summit
status. Also, obtaining data in more than one filter can be
difficult, since obviously the sky darkens very quickly.
The twilight sky will be polarised, particularly towards the
zenith. It is, however, still possible to flat-field data with
twilight flats provided separate e- and o- beam images are extracted
and flatfielded using separate e- and o- flat images, the flats being
normalised individually. Note also that the algorithm used by the
POLPACK:polcal routine allows for a constant polarisation to be
present across the flat field surface (the aperture field of view).
Twilight flats should not be obtained without the focal plane mask in place, since
again the flatfield response of the array may be sensitive to
the state of polarisation of the radiation. The prism and mask always
ensure that only e- or o-beam radiation is transmitted onto a
given area of the array. Consequently, the flatfield response of that
same area of the array must also only be measured in either e- or
Dome flats, or specifically observations of the back of the
primary mirror covers (with the covers closed) are another
possibility. The covers are illuminated with two halogen lamps (2x
150 Watt bulbs) positioned on the platform above the south column,
with the telescope in its "park" position (elevation -20 degrees) and
the dome closed and in darkness. 3000-4000 counts should be measured
in 3-second Z, J and K-band images, and 2-second H-band images (note
that 1 second is the minimum full-array exposure time with UIST).
Exposures through the mask and waveplate, at each of the WP angles,
with the halogen lamps on then off, should be acquired (Use the
"IRPOL Dome flats - UIST" observation under the calibrations menu
TWICE). It is recommended that these "dome flats" be secured at
the start of a night when imaging polarimetry is planned.
Exposure times on sky
On the night-time sky extremely long exposure times will be
required at I, Z and J; dome or twilight flats are therefore recommended
at these shorter wavelengths.
The counts obtained from the blank sky background with UIST through
a number of broad-band filters are given in the table below. The
measurements were made at 2am in bright time; fewer counts are
expected in I and J in dark time! For background-limited
performance and low-noise flat-fielding (necessary for sub-1%
polarisation accuracy), at least 1000 counts are needed; 2000-3000
counts is ideal.
|| Integ. Time
|| Mean counts
|   || (secs) ||(DN) ||   ||
| J[MK] || 1.17-1.33 || 60 || ~80
| H[MK] || 1.49-1.78 || 60 || ~420
| K[MK] || 2.03-2.37 || 60 ||~350
| K-short || 2.03-2.29 || 60 ||~240
| K-long || 2.02-2.43 || 60 ||~300
Which filter should I use: K[MK] or K-short?
As can be seen from the above table, the background with the warm
waveplate in the beam is somewhat higher with the K[MK] filter
than it is with the K-short filter. The red-end cutoff of the latter
results in a reduction in the background by about 30%. Consequently,
if we ignore source colour, then an increase in sensitivity is expected,
though of course you also get 24% less signal from the source! The
K-short filter might nevertheless be worth considering, particularly for
imaging-polarimetry of fainter sources, although note that this filter
is less-well characterised than the Mauna Kea standard K[MK] filter.
Observers requiring accurate (absolute) photometry should use the
K[MK] filter. See the UIST imaging web pages for filter
Imaging Polarimetry in the L and M-bands
Please contact Chris Davis for further details (c.davis at
Imaging Polarimetry: Data Reduction
ORAC-DR with IRPOL
As mentioned in the previous section, ORAC-DR pipeline recipes are
now available to users. The template observing sequences discussed
above contain recipes appropriate to the associated observing mode;
these recipes are described below. Please take care when changing DR
recipes; many have specific requirements in terms of darks and flat
fields, which must be acquired before a target observation is obtained
and reduced on-line. If a special reduction recipe would be useful to
you, please contact your support scientist - we may be able to produce
something specific to your needs, though we do need to be notified
well in advance of your run at UKIRT.
To run ORAC-DR on KAUWA type
The software will then point to the current night's data directories.
(If you wish to reduce, say, the previous nights data, you can
specify the UT date on the command line, e.g. oracdr_uist
The above command should be followed by
oracdr -loop flag
The pipeline is meant to run without interference from the observer.
Thus, although you can use the various GAIA tools to examine images,
the pipeline should not need to be stopped and/or restarted. If,
however, you do need to re-reduce a block of data, then this is
possible with the command
oracdr -loop flag -from 199
oracdr -loop flag -list 199:210
Help on ORAC-DR is also available by typing
The three ORAC-DR recipes currently available,
extract e- and o-beam images from each object frame. These
are flat-fielded using separate, pre-observed, flat-field data reduced with either
Note that ALL pipeline recipes require separate sky flats! A
sky level is subtracted from the data, that is either calculated from
the lower half of the array or (for extended objects) from separate
sky observations. Extracted e- and o-beam images will be displayed on
a GAIA display; the final, reduced pol map (an intensity image with
polarisation vectors superimposed) will eventually be displayed in a
separate GWM window (see e.g. Fig.3). I, Q and U stokes vector
mosaics, as well as P and TH (angle) images will be stored to disk by
each recipe (see the individual recipe pages for a comprehensive
description of the reduction steps). Note that a correction to the
position angle due to waveplate orientation w.r.t. N and E is also
automatically applied; see below.
Fig.3 Top: An intensity image with pol vectors overlayed, as
produced by ORAC-DR and displayed in the kapview/GWM window. Bottom
left: Redisplayed data from ORAC-DR, using Gaia and the Polarimetry
toolbox in Gaia. Here data have been binned over 10 UIST pixels and
vectors where I>10 have been highlighted in red. Bottom right:
published results for the same target, from Whitney et al. (1997, ApJ,
The DR recipes also produce catalogue files with the suffix ".FIT".
This file is a table of measured polarisation parameters that may be
manipulated and displayed with standard Polpack commands (Polpack
is a Starlink package designed for reducing polarimetry data). For
example, you might want to bin the data and only plot vectors where
there was emission from the source and where P/dP was greater than 3.
The necessary commands, polbin, catselect, display and
polplot are given here
(remember to type polpack AND cursa first to activate these
Alternatively, try the new "Polarimetry
Toolbox" in Gaia. Simply display a gu< UTdate >_< num >_I.sdf
frame, then use the toolbox to read in a .FIT table file and overplot
the vectors (best to use the binned _bin.FIT file, otherwise you may
get thousands of vectors and it may take a while loading in the
table!). You can then quickly manipulate the displayed data: bin data
in adjacent pixels, select values with e.g. P/dP > 3, adjust the
scaling or colours of the image and vectors, etc. An example is shown
in Fig.3, alongside published observations of the same target,
L1551-IRS5 (from Whittney et al. 1997, ApJ, 485, 703).
Note, finally, that the images displayed by ORAC-DR appear rotated
through 90 degrees, although the labeled axes should be correct (see
top image in Fig.3) and the data can always be re-oriented in Gaia.
Calculating the Stokes Parameters
The Stokes parameters are currently calculated using the "ratio"
method; an alternative route, known as the "difference" method, is not
currently implemented. The ratio calculation method is as follows:
The U stokes parameters are computed from the 22.5o and
67.5o intensities in the same way. This method works well
for bright objects but can fail for fainter or noisier sources (and
100% calibration source) if/when the algorithm tries to take the
square root of a negative number. In such cases, the difference
method should be used, as described below:
Again, the U stokes parameters are calculated from the
22.5o and 67.5o data.
The Q and U parameters are then used to calculate the state of polarisation
(i.e. the degree and position angle of the polarisation vectors) using
the following equations:
where 90o must be added to
Q is negative, or 180o added to
if U is negative and Q positive.
Post-observing reduction with POLPACK
Subsequent "re"-reduction and/or fine-tuning of your data may be
carried out using the STARLINK software package
e-o Beam Separations with UIST
The Wollaston prism in UIST acts as the polarimetric analyser, splitting the
incoming radiation into orthogonally polarised e- and o-beams.
The divergence of the beams is dependent on the refractive index of
the prism material and is therefore wavelength dependent.
The beam separations given in the table below are measured for
bright point sources and are given in arcseconds; the UIST pixel
scale is 0.12". The offsets are between top (e-beam)
and the lower (o-beam); e- and o- beams are dispersed E-W.
Note also that the dispersive properties of the prism will result in slight
elongation of the images in the dispersion direction, i.e. in an E-W
direction, although this is only noticeable under conditions approaching
E-W offset (arcsec)
Calculating Exposure Times
To calculate the integration time needed for polarimetric
observations with UIST you should apply the following equation to the data in
the sensitivity tables provided on the UIST-imaging web
pages. Usually in polarimetry, a certain polarisation accuracy is
needed to have any confidence in the result. For instance, if a source
is 1% polarised an accuracy of at least ~0.3% would be necessary. This
polarisation accuracy is converted into a required signal-to-noise
on the flux (i.e. not per waveplate position) using this simple
For example, to obtain a polarisation accuracy of 0.5% a S/N ~ 283 in
the total, on-source integration time (not including
overheads associated with telescope and waveplate moves) is needed.
Note, however, the potential limitations on maximum signal-to-noise
achievable with each instrument; discussed below.
With the broad-band K, Lp and Mp filters the above calculation is
skewed slightly by the extra background that is introduced by the
waveplate (which is, of course, in the incident beam). This will
reduce the sensitivity slightly. Observers should account for this by
reducing the sensitivity by ~0.5 mag with the K[MK], Lp and Mp
EFFICIENCY: When writing telescope proposals, please remember
that polarimetry observing involves telescope nods, reading out the
array, AND waveplate moves. These factors all conspire to make your
observing less efficient. In your proposal you should therefore DOUBLE
your observing time request to cover these overheads (i.e. assume 50%
There is a limit to the signal-to-noise (S/N) that is possible
with each of our instruments, because of inherent uncertainties in
flat-fielding, variable sky levels and (for bright targets and
therefore short exposure times) read-noise. Although tests have not
been carried out with UIST, with UFTI measurements towards a 9th
magnitude standard star (without IRPOL in place) suggest that a
S/N per pixel of 2000 can be reached in the expected integration time;
in the plot below, one can see that the measured S/N (blue squares)
only begins to deviate from the statistical S/N (black dots) when
very high values of S/N are sought. The fact that high values of
S/N can be reached is a reflection of the very "flat" UFTI array and
the negligible read noise. On reasonably bright sources, users can
therefore expect to get high polarisation accuracies (of the order of
0.1-0.2%) in the expected on-source integration time with UFTI.
We hope for a similar performance with UIST.
Plot of the change in signal-to-noise (per pixel) with time
for a 9.3 magitude F3 standard star (without IRPOL2 in the beam). Blue
squares - measured; black dots - predicted.
Guide Stars for IRPOL
The IRPOL waveplate is positioned in the beam within ISU2; the
plate is held in an opaque circular holder which is lowered into the
beam by the Telescope System Specialist. Consequently, the waveplate
holder obscures some of the field that may be used to find guide stars
(should your target be too faint for the fast guider) - this Bird's
Eye View of IRPOL through the hole in the primary illustrates this
The field-of-view accessible for guide stars with IRPOL is illustrated
in the figure below; your target (central coordinates) are assumed to
be located at the position of the cross. The waveplate holder is
shaded grey and blue in the figure. Ideally, guide stars should be
found within the inner white circle..
Fig.4: A Field-of-view with IRPOL for finding optical guide stars.
A guide star may be used outside the waveplate holder
(outside the blue box, though inside the yellow circle, in the above
diagram). Space is limited here and in some places blocked by the
drive belts and arms of the holder (see the Birds-eye view). Because of the way
IRPOL is mounted above the tertiary, these guide stars should be to
the north-east or south-west of the target, and offset
by 150"-250", depending on their exact location (250" is the radius of
the full field for the guider). Of course care should be taken with
guide stars near the edges of the waveplate holder when slides are
used in your observing sequence.
A postscript version of the above figure is available
Imaging Polarimetry with UIST:
Instrumental Polarisation and Efficiency
Instrumental polarisation and efficiency
The instrumental polarisation is expected to be low, certainly
much less than 1%. We encourage all users to observe an unpolarised
standard (and to communicate their results with JAC staff!). Some
results are given
The efficiency of IRPOL2 has been measured with UFTI and CGS4 via
observations of bright stars through a Glan prism (which induces 100%
polarisation). The system was found to be almost 100% efficient at
all wavebands. A similar performance is expected with UIST, which uses
the same type of wollaston prism as CGS4.
Polarisation Efficiency (UFTI Data)
|| Polarisation (%)
|| Pos. Angle (degrees)
| CMC 609748 ||I ||103.0 ||21.0
| CMC 609748 ||J ||99.7 ||22.1
| CMC 611929 ||K ||100.3 ||22.5
Position Angle Calibration
To calibrate measurements of polarisation position angle with UIST
+ IRPOL, observers should also make their own deep observations of at least
one polarised standard.
Some targets have been observed by staff and visiting observers.
Their results are tabulated here.
At present, it is thought that a correction of about -24 degs
is needed with UIST polarimetry observations. This value must
be applied to data after it has been orientated with North up
and East to the left.
The above p.a. correction is considerably larger than that derived
for UFTI and IRCAM3, where a correction of -6 degrees has been
measured. This is somewhat surprising, since the waveplate is above
the tertiary mirror, so the orientation of the polarisation angle
should be similar for all instruments. The difference may be due to
the fact that the dispersion of the e- and o-beams, and so the axis of
the prism, is not perfectly aligned with the columns of the arrays
(see e.g. the e-/o- beam separation table above) in any of the
IMPORTANT: The DR will apply a corrections
of -24 degrees to observed angles automatically - note that the
match between the ORAC-DR results and the published observations in
Fig.3 above is good. Additional corrections to the angle can
be made in the Polarimetry toolbox in Gaia.
Imaging Polarimetry with the Coronagraphic Mask
The basic idea...
Finally, in 2006 we installed a second imaging-polarimetry mask in UIST. This is
essentially the same as the normal mask, except that two fine wires are extended
across the bottom rectangular aperture (see below). The left- and right-hand wires
are 6-pixels (0.7 arcsec) and 11-pixels (1.3 arcsec)
The idea is that bright sources can be positioned behind
either wire so that longer exposures can be secured without saturating
the array and without latency issues, bleeding, internal reflections,
ghosting, etc. Under normal observing conditions (0.6-0.7 arcsec
seeing, clear skies, etc.) 10-20 sec images can be obtained through
the broad-band filters on 7-8th mag stars.
Fig.5: A raw image of blank sky through the
Wollaston prism, showing e- and o- beams
projected from the two rectangular apertures in the coronagraphic
imaging pol mask. The bottom aperture includes two wires; e- and
o-beam images of these are seen as dark, vertical strips.
Data are acquired in essentially the same way as for normal
imaging polarimetry; separate dome or sky flat fields are required
prior to observing the target, and a few jittered observations are
taken at each of the four nominal waveplate angles.
However, there are a few important differences. Obviously the
offsets must keep the target behind the wire. In the example MSB in
the template library, offsets are therefore small, multiples of the
0.12 arcsec pixel scale, and in the "q" direction. Regardless of
position angle (see below), offsets in q will move the target up and
down the wire. (An offset of q= +48 will put the source in the top
pol aperture, though this is usually not desired.)
Imaging acquisition: putting the bright target behind the wire
In the example MSB in the template library, a short (1sec) exposure
is used to acquire the target. The instrument is run in "Movie mode",
which means that frames are taken (though not saved) repeatedly so
that the target can be placed behind the wire. This 1 sec exposure
time is subsequently updated using a "UIST Imaging Iterator", so that
longer exposures can be used when taking the data themselves.
By default the acquisition process will put the source behind the
6-pixel (0.7 arcsec) wire. If the wider occultor is required, the
source can be manually shifted "left 62.6
arcsec" while Movie is still running. Small left/right offsets
at this new position can be used to fine-tune the position behind the
1.3 arcsec wire. Tools in Gaia like "View - Slice" and particularly
"Image Analysis - Mean XY profiles" can be used to check the
positioning on the wire.
IMPORTANT: When using "Pick Object" in Gaia
to acquire the target, click on the source in the BOTTOM beam.
If you pick on the top beam, "upick" will probably move the target out
of the pol mask aperture. If you don't see the source when you first
slew to it, it may be hidden behind the mask. Ask the telescope
operator to move "up 10" and/or "down 10" while you are running Movie.
When you can see the target, then use Gaia/Pick-object and upick to
put it behind the wire.
The position angle of the image plane can be orientated so that
extended targets (jets, disks, etc.) are placed orthogonal to the
wire. For example, with a disk orientated E-W
you should use a position angle of 0 degs; with a disk orientated N-S
you should use a position angle of -90 degs. Any angle between
-90 and +90 degrees can be used, though acquisition will be easier if
you use angles between 0 and -90! A diagram showing the effect of
various position angles on raw images is given in the
main UIST imaging pages. Essentially raw coronagraphic images are
always flipped about a horizontal axis. They are then
rotated by the position angle selected.
Observers should use the recipe POL_JITTER_CORON, which is
an adaptation of the normal POL_JITTER recipe described
above. POL_JITTER_CORON uses the lower (rather than upper) e- and
o-beams for the target. The upper beams will still be used for
sky-subtraction, and a separate sky (or dome) flat-field is
required. The flat-field should be acquired in the normal way
(described earlier), though of course the coronagraphic mask, rather
than the normal imaging pol mask, should be used.
H-band coronagraphic imaging polarimetry of the Red Rectangle (H~5.1 mag).
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