[ ESO ] La Silla

2p2-m Telescope
Calibrating IRAC2b

General Introduction
Standard Reduction Sequence
Dark Frames
Flat Fields
Illumination Corrections
Sky Subtraction
Bad Pixel Map
Fabry Perot Data
Standard Stars
Historical Instrument Sensitivity
Archived calibration frames

General Introduction

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This calibration page has two purposes. Firstly, for the observer it describes a method to reduce IRAC2b data and provides some of the necessary calibration frames to do it. Some observing tips are also included. Secondly, it serves as a data base where the 2.2 team records the characteristics of the camera and archives old calibration frames.

A Sample Reduction Sequence

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The reduction of IRAC2b data is not too dissimilar to the reduction of CCD data. There are however some important differences. For example, with IRAC2b there are two read out methods: NDR - Non-Destructive Read and DCR - Double Correlated Read. This leads us to our first rule. Do not mix the two readout methods when calibrating IRAC2 data.

Another difference is that IRAC2b does not have a shutter, nor does it shuffle charge around like a CCD, so the concept of a bias frame is meaningless.

The steps to reduce IRAC2b data differ slightly between the two read out methods.

For DCR one:

(1) Subtracts a dark
(2) Linearises the data
(3) Flat fields the data
(4) Subtracts a sky frame
(5) Cleans bad pixels

For NDR one:

(1) Linearises the data
(2) Subtracts a dark
(3) Flat fields the data
(4) Subtracts a sky frame
(5) Cleans bad pixels

Not all these steps may be necessary; for NDR only step 3 is essential and for DCR steps 1 and 3 are essential.

Dark Frames

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The dark frames for the two read out methods differ substantially. Dark frames should be taken at least once per night.

DCR Darks

A dark taken with the DCR method has a mean value of -1000 ADU, shows vertical stripes and becomes more negative with increasing exposure time. If you subtract a dark from your data and you are using DCR, it is essential that the dark has the same exposure time as the science frame. The dark current for DCR is not a linear function of time. Note that the camera suffers a slight light leak, so observers should only take dark frames during the night or during the day with the dome dark. The perfect time to take darks is before the afternoon flats.

NDR Darks

A dark taken with the NDR method, on the other hand, shows positive counts which increase with exposure time, has considerable structure and is made up of three components. The first and by far dominant component is the heat from the output amplifiers, one in each corner, and depends only on the number of reads. It is independent of the integration time. The number of reads increase from a minimum of 2 for integrations of 0.5 seconds to a maximum of 20 for integrations of 6.25 seconds. For longer integrations the number of reads remains at 20. The second component is the actual dark current of the array and is less than 0.6 ADU/sec. The third component is a light leak. The light leak means that observers should only take dark frames during the night or during the day with the dome dark. Furthermore, due to a slight memory effect it is not recommended to take DCR darks after an exposure that has been exposed to a bright source (this includes the K-band sky). The perfect time to take darks is before the afternoon flats. If you subtract a NDR dark from your data it must have exactly the same exposure as your science frames if the exposure of the science frames is shorter than 6.25 seconds. For longer exposures it is sufficient to subtract a dark with exposures greater than or equal to 6.25 seconds. For example it is acceptable to subtract a 10 second dark from an exposure lasting 160 seconds, but it is not acceptable to subtract a 10 second dark from a 5 second exposure, nor is it acceptable to divide the 10 second dark by 2 and then subtract it from the 5 second exposure.

Non-Linearity Correction

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IRAC2b, is a non-linear device. Up to 10,000 ADU the non-linearity is less than 2%. Beyond this flux, the number of pixels that become highly non-linear increases dramatically.

In general, the linearity correction is not required, especially if you keep to counts below 10,000 ADU.

For NDR data the linearity correction is applied to all data. For DCR the linearity correction is applied to dark subtracted data only.

The average non-linearity of the array can be represented by the following formula. The constants in the formula differ for the two read out methods.

Readout Methodbc

Two words of caution. The non-linearity may depend on flux rate and wavelength. Furthermore, the equation is the average over a small region of the array. It may not be applicable to individual pixels.

Flat Fields

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Although there are four commonly used lenses available for IRAC2b, flats for lenses LA, LB and LC should be taken with lens LC only. Images with lens LA and LB have light concentrations which do not accurately reflect the sensitivity of the array. Flats for lens LD can be taken with either lens LC or lens LD. We recommend you try both and see which of the two you prefer.

In principle it is best to match both the exposure time and the flux level in the flats with what you will get in the observations. However, in practice this is either time consuming or not possible, particularly with flats created by imaging the dome. Flats should be taken at least once during your run. Taking flats with IRAC2 is very easy, so you might consider taking flats every night.

There are three commonly used methods to create flat fields for your data. By far the most common method is to take images of the dome flat-field screen with the flat-field lamp on and off. The recommended sequence of exposures is with the flat-field lamp first off, then two exposures with the flat field lamp on and finally a single exposure with the flat-field lamp off. The flat field is the difference between the "on" and "off" exposures. The usual detector settings are a DIT of 1 second and a NDIT of 30. This method has the advantage of subtracting the dark automatically as well as any contribution from the telescope which can be substantial at K. The disadvantage is that the flat field does not represent the large scale sensitivity of the array. This is corrected with an image which is called illumination correction - more on this later.

Flats can also be created from the twilight sky (sky flats) or from the observations themselves (super flats). In both cases it is necessary to subtract a dark from the flat first. These flats have the advantage that they better represent the large scale sensitivity of the array. The disadvantage of these flats is that the signal from the telescope, which may be non-uniform, is not subtracted. For super flats, there is the additional advantage that the exposure time and flux level matches that of the observation; however, this method is not well suited for the narrow band filters. Furthermore, observations done in lens LA or LB cannot use this method. This method is generally used for deep K-band observations with lens LC.

Flat Field Variability

The flatfields can vary from night to night or from run to run. From night to night there are four reasons why flats may vary.

Firstly, the flat may vary because the illumination of the flat field screen has changed. This can be caused by a change in the position of the flat field lamp. In the quotient of two flats taken with the flat field lamp in differnet positions, the changee appears as a herringbone pattern or as a slope. The peak to peak difference in this image generally ranges from less than 1% for J to less than 3% for K.

Secondly, the differences can occur if the lens wheel changes position very slightly. In the quotient image this usually appears as pattern of arcs near the centre of the image. The effect is less than 1\%.

Thirdly, there can be variations in the sensitivity over certain regions of the array, in particular (x,y)=(200,190), (x,y)=(230,90) and (x,y)=(5,130). In the quotient of the two flats taken two nights apart this usually appears as excess noise in the regions affected. The effect is less than 1\%.

And fourthly, there is a small triangular shaped region of the array near pixel (1,256) which is vignetted. In the quotient image this usually appears as either a region excessively high or excessively low counts. This region cannot be accurately flatfielded. We recommend you either flag this region as bad, or replace the pixel values in this region with a constant.

Over a period of months where the camera has been taken off and then put back on the telescope and where the detector has been temperature cycled, variations in the flat-field can occur for all the reasons outlined above, especially the third. As an example, the RMS variations of the quotient of two flats taken two months apart is shown for J, H, K and K' flats in the following table.

FilterRMS of the Quotient Image

Sample Dome Flat Fields

These flat fields were taken on the 14th December, 1997. They are the difference between an exposure of the dome flat field screen with the flat field lamp on and an exposure of the dome flat field screen with the flat field lamp off. These flats were taken in DCR. The exposure time was 1 second and the number of detector integrations was 30. The signal to noise per pixel is dominated by the exposures with the lamp on. The S/N per pixel in these flat field frames is 2000.

J Flat ffjn980320.fits.gz (220 kb)

H Flat ffhn980320.fits.gz (220 kb)

K Flat ffkn980320.fits.gz (220 kb)

K' Flat ffkpn980320.fits.gz (220 kb)

Illumination Corrections

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The flat fields generated by methods outlined above do not represent the true large scale sensitivity of the array. The large scale sensitivity can be corrected by scanning a bright star, usually a standard, over the array, fitting a surface (3rd order in x and y without cross terms is usually sufficient) to the normalised fluxes (not magnitudes), and then multiplying this surface into the flat field. For lens LA and LB, 9 positions should be sufficient, for lens LC and LD, 16 positions are required. In the following table we list the residual error in instrumental magnitude of the scanned star before and after the correction. The table is complete for lens LC, partially complete for lens LB. There is no data for lenses LA and LB.

All flat fields were taken in lens LC.

Lens/FilterDome Flat uncorrectedDome Flat CorrectedSky Flat uncorrectedSky Flat corrected

If a consistent zero point (ie better than 5%) over the entire field of view is important for your work, then an illumination correction is essential.

Observing Sequences

To assist you, observing sequences are available at the telescope. They are located in the sequence sub-directory ./CALIB/ and there are sequences for each of the lenses. To use them, place a bright source in the center of the array and execute the sequence. The return-to-origin flag must be set to Y. As sequences are limited to 12 exposures there are two sequences, each of eight exposures, for lenses LC and LD. These sequences are generic, so you will need to adjust them for the filter, readout method and exposure you require.

Variation of the Illumination Correction

The illumination correction depends slightly on telescope position. The second column in the following table shows the residuals in the magnitude of a bright star scanned over sixteen positions for five different telescope positions. The third column shows the residuals after a 3rd order surface was fitted for each individual telescope position. The final column shows the residual if the fit at Zenith is used to correct for all telescope positions.

Thus one can expect the zero-point to be consistent to better than 2.5% over the array for reasonable telescope positions.

Lens/FilterTelescope positionBefore FitResiduals of the FitAfter ZD Fit
LC/K'10 degrees S. of Zenith0.0470.0150.017
LC/K'43 degrees W. of Zenith0.0360.0160.028
LC/K'27 degrees N. of Zenith0.0400.0160.020
LC/K'40 degrees SE. of Zenith0.0420.0120.022
LC/J43 degrees W. of Zenith0.0600.0210.028
LC/J27 degrees N. of Zenith0.0640.0170.020
LC/J40 degrees SE. of Zenith0.0690.0130.016

The illumination corrections also change between runs. The RMS variations in the quotient images can be as high as 3%, although it is not entirely clear how much of this variation can be attributed to errors in the surface fit. In general the largest variations are for K and K' and occur in the corners of the array.

Sample Illumination Corrections

These corrections are valid for dome flats taken with lens LC data only. They were derived from data taken on the 16th March, 1997. Data is corrected by multiplying the flat field with the illumination corrections.

J Correction illumj970316.fits.gz (200 kb)

H Correction illumh970316.fits.gz (200 kb)

K Correction illumk970316.fits.gz (200 kb)

K' Correction illumkp970316.fits.gz (200 kb)

Sky Subtraction

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Sky subtraction can vary from the simple, subtracting one frame from another, to the complex, creating a running sky from a large number of frames. Rather than going into details of the reduction process itself, it is better to pass on a few suggestions The number of sky frames one observes, the pattern of observations and the frequency one switches between object and sky positions is crucial for good sky subtraction. I will describe two common methods.

Separate Sky Positions

There are two points to keep in mind. First, the more sky positions the better, and second, one must interleave the observations of the sky with observations of the programme object. Here is one example. Consider a single field we wish to observe. Keeping the above two points in mind we can construct a series of exposures which takes sky frames that are East and West (or North and South) of the object field. One can also construct a sequence of exposures where the skies are East, West, North and South of the object. After aligning the telescope with the object we first take a sky frame to the East of the object (we could have also started to the West), second we observe the object, third we observe the object again (a small offset can be inserted between the second and third exposures), fourth we take a sky frame to the West of the object, fifth we take another sky frame near this same place (an offset between the forth and fifth exposures is essential) , sixth we observe the object, seventh we observe the object again (an offset between the sixth and seventh exposures can also be performed) and finally we take a sky frame East of the object and near our first exposure. If the return to origin flag is set, the telescope will move back to the object. Thus one will have four exposures centred on the object and four exposures centred on the skies that are east and west of the object.

More complicated schemes can be implemented. For example, each sky exposure in the above sequence can be divided into two separate sky exposures, with each exposed for half the time of the original and with a small offset in between. Thus one would have 8 exposures on the sky and 12 exposures in total. With this pattern one still spends the same amount of time on the sky as on the object. This pattern is useful if the field is crowded.

Deep Exposures

If the object of interest is small enough, it is not necessary to take separate sky observations. In this case one can dither within the field. As a rule, the offset should be greater than 10 arcseconds and if very deep exposures are required, then offsets that repeat themselves are not sufficient. For example, an offsetting scheme that is based on a rectangular grid of points will in a deep exposure show faint negative images arranged symmetricly around each real image. One should aim for a scheme whereby the offset vector is never replicated. Many such schemes can be devised.

Bad Pixel Map

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The detector in IRAC2b contains a lot of bad pixels, approximately 700. In general the number of bad pixels is flux and wavelength dependent. There are more bad pixels at low flux levels. The list of bad pixels is defined by deviding a high flux exposure with a low flux exposure and selecting all pixels that deviate from the local background by more than 10%.

Fabry Perot Data

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The Fabry Perot unit sits outside the cool environs of the camera. As the FP sits in the air, all vacuum wavelengths must be multiplied by 1.000273. The throughput varies between 70 and 80% over the 2.0 to 2.5 micron range and the finesse is greater than 50 for most of this range. It has a wavelength resolution of 0.0016 microns.

Obtaining and calibrating data taken with the Fabry-Perot requires some care. Flat fields are required for all wavelength settings and a bright reflection arc near the Eastern edge of the array (only seen in lens LC images) makes the array difficult to flat field in this region.

For estimating line fluxes it is possible to achieve accuracies of 5% with six wavelength settings, four near the line and two well separated from the line centre and on either side. Note that separate skies are required for each wavelength setting. Using images at different wavelengths as skies will lead to less accurate results.

Obtaining velocities requires even greater care. The wavelength of maximum transmission varies across the array. If accuracies better that 200 km/s over a large part of the array are required, then this wavelength shift must be corrected. This contour plot shows the wavelength shift in terms of velocity. It is also available as fits file. The area of missing data is caused by the reflection arc.

Standard Stars

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The IR window between 1 and 2.5 microns contains several large absorption troughs that are primarily due to water vapour in the atmosphere. As the amount of water vapour varies so will the amount of absorption. Unfortunately, the edges of some IR filters, particularly J and K, are defined by these absorption features. Thus, when the column density of water vapour is variable, good photometry can be difficult to achieve. On good nights (generally when the humidity is low) it has been possible to achieve better than 1% photometry; however, on most nights this should be considered as the lower limit.

To get the best photometry you should choose standard stars that are as close as possible to your programme objects and you should observe them before and after you observe your programme objects. During the night you should observe at least three different standards and you should observe at least one standard every two hours. On nights where the humidity (which is a rough measure of the water column vapour) is varying considerably (let's say by 40\% in one hour) you will need to observe standards more frequently.

Generally, a standard star is observed twice in different parts of the array. Thus one observation acts as the sky of the other.

It is important that you take care where the standard star lies on the array. Use the PPS display to find areas that are clear of bad pixels. Furthermore, check what other objects are in the field. Furthermore, the offset between exposures should not be such that other objects in the field interfere with flux of the standard when one frame is subtracted from another. Alternatively, one can take several exposures each with the standard star in different location. Observers commonly use a sequence of five exposures. One in the centre of the array and one in the centre of each of the quadrants.

Note that most standards were observed with single channel photometers with a very wide aperture. Thus close, companion stars were probably included.

There are several standard star lists at the telescope. Many of the standards are far too bright for IRAC2b. It is important not to let the value of any pixel to exceed 10000 ADU. If this limit is exceeded, try defocussing the telescope. As a guide a ninth magnitude star imaged under 0.8" seeing will reach this count in 1 second for lens LC.

Observing Sequences

To assist you, two types of observing sequences are available at the telescope. They are located in the sequence sub-directory ./CALIB/ and there are sequences for each of the lenses. The first type places the standard star in areas of the array that are clean of bad pixels. The second type takes five exposures, one in the centre of the array and one in each centred in each of the quadrants. To use them, place a bright source in the center of the array and execute the sequence. The return-to-origin flag must be set to Y. These sequences are generic, so you will need to adjust them for the filter, readout method and exposure you require.

Historical Instrument Sensitivity.

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The team must be consistent in how ZPs are determined. Here is the recipe to follow.

Use DCR and lens LC
Use dome flats to flat field the data
Measure the standard over a grid of 16 positions, as one would do with in determining the illumination correction.
Normalise the flat by the median.
Clean bad pixels in the flat-fielded images
Take the average of the 16 measurements
Use standard extinction values: J=0.08, H=0.06, K (and K')=0.11.

The following ZPs are for DCR and lens LC. The ZPs before the aliminisation took place were taken by a variety of observers and reduced with a variety of techniques. They should be used as a rough indication of the ZPs at those times.

??/??/96Mirror Aluminised
08/12/95... 21.43... ... ...
04/08/95... 21.42... ... ...
12/03/95... 21.51... ... ...
11/11/94... 21.48... ... ...
02/07/94... 21.60... ... ...
05/02/94... 21.55... ... ...

Archived calibration frames

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J Flats

DateReadout MethodSource
14th December, 1997DCR ffjn971214.fits.gz (220 kb)
7th November, 1997DCR ffjn971107.fits.gz (220 kb)
7th July, 1997DCR ffjn970707.fits.gz (220 kb)
16th March, 1997DCR ffjn970316.fits.gz (220 kb)
16th January, 1997DCR ffjn970116.fits.gz (220 kb)

H Flats

DateReadout MethodSource
14th December, 1997DCR ffhn971214.fits.gz (220 kb)
7th November, 1997DCR ffhn971107.fits.gz (220 kb)
7th July, 1997DCR ffhn970707.fits.gz (220 kb)
16th March, 1997DCR ffhn970316.fits.gz (220 kb)
16th January, 1997DCR ffhn970116.fits.gz (220 kb)

K Flats

DateReadout MethodSource
14th December, 1997DCR ffkn971214.fits.gz (220 kb)
7th November, 1997DCR ffkn971107.fits.gz (220 kb)
7th July, 1997DCR ffkn970707.fits.gz (220 kb)
16th March, 1997DCR ffkn970316.fits.gz (220 kb)
16th January, 1997DCR ffkn970116.fits.gz (220 kb)

K' Flats

DateReadout MethodSource
14th December, 1997DCR ffkpn971214.fits.gz (220 kb)
7th November, 1997DCR ffkpn971107.fits.gz (220 kb)
7th July, 1997DCR ffkpn970707.fits.gz (220 kb)
16th March, 1997DCR ffkpn970316.fits.gz (220 kb)
16th January, 1997DCR ffkpn970116.fits.gz (220 kb)

J Illumination Correction

16th March, 1997 illumj970316.fits.gz (200 kb)
16th January, 1997 illumj970116.fits.gz (200 kb)

H Illumination Correction

16th March, 1997 illumh970316.fits.gz (200 kb)
16th January, 1997 illumh970116.fits.gz (200 kb)

K Illumination Correction

16th March, 1997 illumk970316.fits.gz (200 kb)
16th January, 1997 illumk970116.fits.gz (200 kb)

K' Illumination Correction

16th March, 1997 illumkp970316.fits.gz (200 kb)
16th January, 1997 illumkp970116.fits.gz (200 kb)

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