Spectropolarimetry with UIST
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Spectropolarimetry with UIST
This document describes the use of IRPOL2, in conjunction with the
facility near-IR imager/spectrometer UIST, for spectro-polarimetry.
is discussed on a separate page. Details of the design
characteristics of the polarimetry module and its corresponding optics
can be found here.
A postscript version of the old CGS4+IRPOL guide (using the old
vax-based SMS system) written by Antonio Chrysostomou, can be obtained here. This contains many tips on
spectropolarimetry and may be of use to IRPOL users.
The UKIRT and University of Hertfordshire would
acknowledgment in any publication which contains data obtained
Spectropolarimetry with UIST: Data Acquisition
IRPOL2, in conjunction with UIST, comprises an external (warm)
half-wave retarder (the waveplate), internal (cold) focal-plane
slit-masks and an internal (cold) Wollaston prism. Two
spectroscopic slit-masks are available, a 2-pixel wide mask and a
5-pixel mask. (NOTE: The 2-pixel slit should not be used with the IJ
and JH grisms.) These replace the long-slits used with normal
spectroscopy, allowing the user to observe extended sources.
Basically, the slit-mask transmits two 20"-long (0.24" or 0.60" wide)
images onto the Wollaston prism, which splits these into
orthogonally-polarised e- and o-beams. These four beams are
subsequently dispersed by the chosen grism to give 4 "spectra" on the
array; e- and o-beam spectra of the target and e- and o-beam spectra
of adjacent sky (or another region of an extended source). This is
illustrated in the figure shown below.
Fig.1 A cartoon showing the divergence and
projection of two slit-images onto the array. The light from the
star, which is centered in the top slit, is split into e- and o-beam
slit images by the wollaston prism; these two images are then
dispersed orthogonally by the chosen grism to give e- and
o-beam spectra on the array; the blank sky in the lower slit is
likewise "split" into e- and o-beam spectra on the array.
Compact targets are usually observed through the "top" slit, as
shown above. Small offsets are used to nod the target up and down
this slit. Each "object-sky" pair is repeated at the various
waveplate angles, as described below. Extended targets can
be nodded between the two slits; the distance between the centres of
the top and bottom slits is 46.5 arcsec (same for the 2-pix and
5-pix slit masks).
All 1-5 micron spectro-polarimetry at UKIRT should be carried out
with UIST, as described below. (In CGS4 the prism comes before
the slit mask, so misalignment of the e- and o-beam images with the
slit can create spurious "instrumental" polarisation.)
First-time users may also wish to consult the main
UIST spectroscopy web pages for details on control of
UIST via the UKIRT-OT
Controlling UIST and IRPOL with the OT and QT
The expectation is that most users will work from a "Template
Sequence" in the UKIRT-OT. The "Template Library" contains a number of
sequences for polarimetry. These can be copied into a new programme
and modified to suit your specific needs. It is probably unwise to try
and write a sequence completely from scratch.
Below we show a typical polarimetry MSB, containing three separate
observations (the full blue squares); a flat and arc calibration
observation, a polarised standard, and the science target (in this
case BN). All observations must contain (or inherit - see below),
three components - a Target component, a UIST component and a DRRecipe
component, as well as "running-man" icons to move the telescope and
waveplate (and to repeat sequences), and observe "eyeballs" for
imaging-acquisition and to take the actual exposures.
The "Spec-Pol of BN" MSB below starts with two notes and two components (the broken
blue boxes). The three observations below the second (UIST) component
"inherit" the information in the target and UIST components. However,
the BN observation contains a second target component which
overrides the standard star coordinates. The result is that
all three observations will be taken with the same UIST setup, but the
flat/arc and standard observations will be taken at the position of
the standard, while the BN observations will be taken at the position
- The Target Information components
obviously needs to be edited; click on these 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 slit mask, the regions blocked by the
waveplate + holder, and various guide-star catalogues.
- The UIST component sets the
instrument to its initial configuration; integration time, pol. slit
mask, grism, etc. Make sure polarimetry is selected in this window
(otherwise you may end up taking ordinary spectroscopy observations).
- Separate DRRecipe components are
contained inside each observation. These should already be set
correctly to the recipe specific to your chosen mode of observing,
though you should check these.
Click on the radio button to the left of each observation to
fold/un-fold the observations. When a line is highlighted the
settings of the component will be displayed in the right half of the
window. In the example below the flat/arc and standard star
observations are folded up; the BN observation is open to display its
contents. Look at the examples in the template library and read
Fig.2 A Typical Polarimetry OT Programme, with the IRPOL
iterator opened and set to 0o.
To understand how a typical spec-pol programme works, we focus on
the opened observation of BN, above. We assume that the target and
guide star coordinates have already been set-up, the UIST
configuration set, and the DRRecipe has been checked in the components
above the "Sequence" icon. The steps below the Sequence
icon in Figure 2 define the moves on sky, the waveplate angles, and
the actual exposures.
To measure the polarisation of a source, at least eight frames
must be observed - an object and sky pair at each of the four
waveplate angles (0, 45, 22.5 and 67.5 degrees). However, recent
tests, prompted by the study of Aitken
& Hough (2001) 113, 1300, indicate that sixteen exposures should
be acquired - i.e. an object and sky pair at each of eight waveplate
angles: 0, 45, 22.5 and 67.5 degrees, plus 90, 135, 112.5 and 157.5
degrees (see below for further details).
The waveplate is controlled using an
IRPOL iterator (the "running-man"
icon in Figure 2 above), which steps the waveplate to each of the
angles. The IRPOL iterator highlighted is displayed in the right-half
of the OT programme window, where IRPOL is set to an angle of 0
degrees. Nested below each IRPOL iterator is an Offset iterator, which slides the source
up and down the slit (for the object and sky pairs). In this way,
sky-subtracted spectra are obtained at each of the eight angles. (You
should also find a final IRPOL iterator at the end of the sequence; this
returns IRPOL to an angle of 0 degrees after the observations have
been completed.) The Observe
eyeballs are the actual exposures.
The two slits in the spectro-polarimetry mask are each 20" long (Fig.1).
Thus, when preparing a sequence, for "sky subtraction" a point
source may be offset up and down the upper slit (e.g. with +6" and
-6" offsets) as is shown in the figure below. For a moderately
extended source you may nod between the two slits, so that
positive and negative sky-subtracted spectra are obtained in the two
halves of the array. For a very extended source (which
overlaps the upper and lower slits) it may be necessary to nod
completely off-source (i.e. orthogonal to the slit axis), so that
object spectra are followed by blank-sky spectra at each waveplate
Finally, the Repeat iterator
in Fig.2 could be set to greater than unity to build up
signal-to-noise and therefore polarisation accuracy.
Fig.3 A flat (left) and an arc (middle) taken
through the wollasten prism and a grism. A sky-subtracted
spectral (pol) image of a standard-star is shown at right. The image
shows the source offset up and down the upper slit. Positive e- and
o-beams lie above negative e- and o-beams, the separation between the
positive and negative spectra being 10" in this example. (The
negative beams result from the subtraction of the second "sky" image
from the first "object" image)
The final element in each observation is the Target acquisition eyeball.
Spectro-polarimetry sources are acquired in much the same way as
normal spectroscopy sources. Observers unfamiliar with this process
should review the UIST Spectroscopy web
pages for further details.
Spectropolarimetry with UIST: Data Reduction
Basic Reduction with TSP or POLPACK
The Starlink packages
Polpack and TSP can be used to
reduce spectro-polarimetry data. The following two shell scripts
could, for example, be used on a sequence of eight frames (an
object-sky pair at the first four waveplate angles described
Obviously both require a Starlink installation. These also assume
that some basic image processing (flatfielding, sky-subtraction, etc.)
of the raw data has been carried out with the pipeline (see below).
ORAC Data Reduction for Spectropolarimetry
The UKIRT pipeline ORAC-DR can be used to reduce spectropolarimetry
data for point sources, and to process spectral images obtained for
extended sources. With the former, the recipe POINT_SOURCE_POL may
be used. Flats and arcs with the same instrumental setup must be
observed and reduced before using this recipe on a standard star or a
science target. Also, the data must be acquired in the order
prescribed in the Template library, i.e. object-sky pairs observed
with at least four waveplate angles. If all eight waveplate
angles are used, the pipeline will give reduced output from the first
four angles, then co-add the results from the second four angles to
the reduced group data. The Bottom Line: Eight or Sixteen exposures in total are
required for a reducible data-set.
Briefly, the pipeline will create bad-pixel-masked, flat-fielded
spectral images from each frame with the suffix _ff. An
approximate wavelength calibration is assigned to each of these to give the
_wce frames. Sky frames are subtracted from source frames to
give the _ss spectral images. Obviously, for each block of
eight frames there will be four sky-subtracted _ss frames, one
for each waveplate angle.
Each sky-subtracted image will contain a positive and negative
spectrum for the e-beam and a positive and negative spectrum for the
o-beam (see e.g. the right-hand panel in Fig.3 above). These are
optimally extracted from the spectral image and the positive and
negative spectra combined. The observations at each waveplate angle
therefore yield an e- and an o-beam spectrum (see e.g. Fig.4 - left
panel - below); eight waveplate angles will yield sixteen spectra in
total. These are stored on disk as group spectra, e.g.
gu20040717_416_45_E, gu20040717_416_45_0 ... etc.
(NB: in this example 416 is the group number assigned to the
16-frame sequence, 416-431. The above nomenclature does not mean that
all data were extracted from frame 416. The e- and o-beam spectra with
the waveplate angle at, for example, 22.5 degrees were obtained from
frames 418 & 419, not frame 416...)
Finally, these data are processed by Polpack to produce the following:
Intensity spectrum - gu20040717_416_sp-I
Q stokes-parameters - gu20040717_416_sp-Q
U stokes-parameters - gu20040717_416_sp-U
Polarisation percentage - gu20040717_416_sp-P
Position Angle - gu20040717_416_sp-TH
Polarised intensity - gu20040717_416_sp-PI
These spectra are written as 3-dimensional ndfs (try KAPPA/NDFTRACE on one of
the files for more info); effectively, they are data cubes. Consequently, some
packages may have problems displaying them. Try KAPPA/LINPLOT, or "splat".
Alternatively, the files could first be collapsed, e.g.:
> xtplane cube=gu20040717_416_sp-P ystart=0.5 yend=0.5 image=temp
Warning: The data in the reference IMAGE have been re-shaped but the variance
array was never updated. It will now be deleted.
> ystract image=temp xstart=0.5 xend=0.5 spectrum=IJ-pol-spectrum
> splot spectrum=IJ-pol-spectrum
> ndfcopy gu20040717_416_sp-P IJ-pol-spectrum trim
A FITS binary table is also produced by the pipeline;
gu20040717_416_pol.FIT, together with binned and thresholded versions
(gu20040717_416_bin.FIT and gu20040717_416_pth.FIT) using the
selection criteria: polarisation between 0 and 50%, S/N greater
than 3, standard deviation less than 5% and intensity greater than
Example e- and o-beam spectra from data taken at one of the eight waveplate
angles are shown below (left). The spectra below-right show the final, reduced data:
percentage polarisation and position angle plots (the gu20040717_416_sp-P
and gu20040717_416_sp-TH files) derived from e- and o-beam spectra obtained
at all eight waveplate angles.
|Extracted e- and o- beam spectra
||Percentage polarisation and Position Angle
|Fig.4.Spectro-polarimetry results from UIST and the JH grism
of the Polarised standard, HD 150193. Whittet et al. (1992, ApJ, 386, 562) predict
values of J=3.27% (PA=57deg); H=2.26% (PA=60deg).
These data were reduced and displayed by the orac-DR pipeline.
There is at present no DR for Spec-polarimetry of extended sources.
Frames may be partially reduced with
REDUCE_SINGLE_FRAME_NOFLAT. e- and o-beam spectra may then be
extracted from the individual flat-fielded bad-pixel-masked frames
and the ratioing done off-line with TSP or PolPack. Here again is a simple script that could be used for
Reducing Data at your Home Institution
The ORAC-DR pipeline can of course be run on any computer,
provided the full STARLINK software suite
has been installed.
To run ORAC-DR on your data type e.g. the following:
setenv ORAC_DATA_IN /home/cdavis/my/raw/data/
setenv ORAC_DATA_OUT /home/cdavis/my/reduced/data/
oracdr -list 1:100 &
The above four lines tell the pipeline when the data were taken
(since files are labelled with the UT data), where the raw data are,
and where reduced data are to be written. Obviously these directories
need to exist on your machine, and the raw data need to be in the
specified directory. The last line activates the pipeline; in the
above example, frames 1 to 100 from 20071225 will be reduced.
Note that the "array test" observations taken at the start of the
night MUST be reduced BEFORE you try and reduce a block of your own
data. Use the -list option to specify the data range. For
example, if frames 1-19 were darks taken to measure the readnoise and
set up a bad pixel mask (in the log the DR recipe will be set to
MEASURE_READNOISE and DARK_AND_BPM), you
could type the following:
oracdr -list 1:19 &
Having reduced the array tests,
you can then reduce your polarimetry observations, e.g.
oracdr -from 150 &
where 150 was the flat (151 would then be the arc, and 152 onwards
would be the polarimetry observations of a pol/unpol standard or
When running ORAC-DR, several windows will open as they are needed;
an ORAC text display, GAIA windows and kapview 1-D spectral plotting windows. If you
are at the telescope the pipeline will reduce the data as they are
stored to disk, using the recipe name in the image header.
The DR recipe name stored in each file header will be used by the pipeline. However,
this can be overridden if, for example, you decide you do not want to
ratio the target spectra by a standard star, e.g.:
To exit (or abort) ORACDR click on EXIT in the text log window, or
type ctrl-c in the xterm. The command oracdr_nuke can be used
to kill all DR-related processes, should you be having problems.
Additional Background Info - Calculating the State of Polarization
To measure the polarisation you require sky-subtracted spectra at
at least the four waveplate angles: 0o,
45o, 22.5o and 67.5o (preferably also
at 90o, 135o, 112.5o and
157.5o). In principle, with a perfect Wollaston prism, you
would only need two positions to calculate the polarization. However,
although the Wollaston prism DOES produce orthogonally polarised
beams, the attenuation of each beam differs through the prism because
the optical path that each beam takes is not identical (also, the
refractive index for the two orthogonal states is different). By
measuring the other waveplate positions, it is possible to cancel
these differences out. N.B : This should not be considered as
an overhead as the same signal is being measured at each
There are a number of methods employed for calculating the Stokes
parameters (I, Q, U) from the data. Two are outlined below.
The algorithm for calculating the Stokes parameters using the Ratio
method is :
The DIFFERENCE method
The algorithm for calculating the Stokes parameters using the Difference
method is :
The e and o in these equations refer to
the intensities of the e- and o-beams at the relevant waveplate
positions, given by the suffixes. The same calculation gives the U
Stokes parameter if data for the 0o and 45o WP
positions are replaced in the above equations by data from the
22.5o and 67.5o positions, etc.
Both methods efficiently correct for transmission changes between
individual observations (in other words, pol observations are possible
in non-photometric, or "cirrusey", conditions, although the pol angle
measurement will be less accurate). The RATIO method works well for
bright objects but can fail for faint or noisy sources (and 100%
polarized calibration sources) when the algorithm attempts to take the
square root of a negative number. The DIFFERENCE method should be used
in these instances.
Once the q and u Stokes vectors are obtained,
the state of polarization (i.e. degree of polarization and position angle)
can be calculated according to the equations :
The Wollaston prism in UIST acts as the polarimetric analyser,
splitting the incoming radiation into the orthogonally polarised e-
and o-beams, as described above. This divergence is dependent on the
refractive index of the material and is wavelength dependent.
The beam separations given here are in units of pixels
(the UIST spec-pol pixel scale is
These separations are measured from the "bottom"
beam to the "top" beam.
dependence of the beam divergence is not very steep, though
there may be a small amount of curvature of the spectrum along the
dispersion direction, particularly when using the lower
Calculating Exposure Times
To calculate the integration time needed for polarimetric
observations with UIST you should apply the following equation to the
sensitivity tables provided for spectroscopy (see main
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.
In the thermal 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 L and M-band grisms.
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
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.
The fact that such high values of S/N can be reached is a reflection
of the very "flat" UFTI array and the negligible read noise.
For spectro-polarimetry the S/N is limited by the multitude of
variable atmospheric absorption features in the near-IR, as well as
possible flat-field variations. On the same 9th
magnitude source as discussed above, a maximum S/N of 200-300 was
measured on sky-subtracted/flat-fielded spectra (a 60 second
flat-field was used) with CGS4; this limit was reached after only 2
minutes of integration time. Of course, dividing by a standard helps
considerably, by removing most of the atmospheric absorption lines and
leaving only the "true" per-pixel statistical noise. Clearly, a good
standard (with appropriate airmass/spectral type) must be observed if
high S/N is needed with ordinary spectroscopy. Fortunately, with
IRPOL's dual-beam capabilities, both e- and o-beams are observed
simultaneously (and of course at exactly the same airmass). Hence,
when ratioing these spectra to measure the polarisation you also
compensate for atmospheric "noise" and so can expect to attain a much
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..
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
Instrumental Ripple, Polarisation and
It has been found that spectro-polarimetry data often suffer from
modulation or "ripple" in polarization and position angle. This
artifact is probably due to multiple reflections between parallel
surfaces within the waveplate, which then acts like a Fabry-Perot
interferometer. A summary is given in the UKIRT
Newsletter; more complete details and analysis are presented by Aitken
& Hough (2001) 113, 1300.
To correct for this effect, we recommend that spec-pol users
observe at eight waveplate angles. The additional angles, the nominal
four incremented by π/2, are available in the OT, and supported by
the DR pipeline. In this case and object-sky pair would be
observed at 0, 45, 22.5 and 67.5 degrees, plus 90, 135, 112.5 and
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
Measurements of the polarisation efficiency at I, J, H and K have
been secured with CGS4 (May 1999); we expect little change in these
values with UIST, which uses the same prism material. The results
obtained with the 40 l/mm grating in CGS4 indicate an efficiency
exceeding 99% at all four wavebands (see this
postscript plot of the complete data set: note that the stokes I
plot reflects the fact that 3 different sources, BS4358, BS4929 and
BS5553, were observed; hence the difference in the absolute flux
measurements.) Observations with the 150 l/mm grating in CGS4 yield a
polarisation efficiency of 100.2% (+/- 0.03%) at K (2.08 microns).
Overall, these results are in good agreement with earlier studies.
L' and M-band measurements were obtained with CGS4 in August 1999 with
the 40 l/mm grating. The L and M-band waveplates are zero-order
plates (the waveplate for IJHK observations is an achromat).
Consequently, the efficiency will be non-linear across these
wavelengths. At L' the efficiency peaks at 97% at around 3.65 microns
and drops off, approaching 86% at longer wavelengths. At M the
efficiency peaks at 94% at 4.8 microns and decreases to less than 90%
towards the edges of the window. Plots of the L' and M-band
measurements, together with 2nd order polynomial fits, are available
L-band efficiency -----
The fits used are:
L' band: P(%) = -398.1 + 272.6(wavelength) - 37.5(wavelength)2
M band: P(%) = -1014.1 + 463.1(wavelength) - 48.4(wavelength)2
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 -20 degs
is needed in the H and K bands with UIST spectro-polarimetry
observations. UIST imaging-polarimetry observations indicate a
similarly large correction.
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 in
any of the instruments.
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