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 below. Spectro-polarimetry 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 here.
The UKIRT and University of Hertfordshire would appreciate an acknowledgement in any publication which contains data obtained using IRPOL2.
Imaging Polarimetry with UIST: Data Acquisition
Introduction IRPOL2 comprises a half-wave retarder (the waveplate) and, internal to UIST, a focal-plane mask and a Wollaston prism.
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.
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 3×3 grid of images at each of the 4 waveplate angles. The recipe SKY_FLAT_POL_ANGLE will then reduce each flat separately.
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.
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 the ratio 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 thePOLPACK: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 o-beam radiation.
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.
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 characteristics.
Imaging Polarimetry in the L and M-bandsPlease contact Chris Davis for further details (c.davis at jach.hawaii.edu).
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 20001031 .)
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, POL_JITTER, POL_ANGLE_JITTER and POL_EXTENDED extract e- and o-beam images from each object frame. These are flat-fielded using separate, pre-observed, flat-field data reduced with either SKY_FLAT_POL or SKY_FLAT_POL_ANGLE. 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.
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, displayand polplot are given here (remember to type polpack AND cursa first to activate these commands).
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 ParametersThe 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
if Q is negative, or 180o added to
if U is negative and Q positive.
Post-observing reduction with POLPACKSubsequent “re”-reduction and/or fine-tuning of your data may be carried out using the STARLINK software package POLPACK .
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 half-arcsecond seeing.
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 equation :
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 filters.
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% efficiency).
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.
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 nicely!
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..
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 here for printing.
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 here.
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.
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 degsis 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 instruments.
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) wide, respectively. 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.
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 rotatedby the position angle selected.
Data reductionObservers 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.