CGS4 optics layout

CGS4 Sensitivity

The 256×256 array has a read noise of about 42e- for 2-read NDR with multi-read NDR noise of about 23e-.

Sensitivity: Some Very Important Notes – Please Read.

• The values listed below are surface brightness sensitivities.
• The sensitivities are calculated for the best part of the band (near 100% atmospheric transmission). They cannot always be used to interpolate sensitivites at other wavelengths within the band.
• Sensitivities for point sources can be estimated from the following tables. Assuming 0.6″ seeing and that the spectrum falls onto three rows of the array, the signal-to-noise will be approximately a factor 2.5 lower than quoted in the tables (about one magnitude).
• CGS4 is less efficient when very bright sources are observed and it is difficult to obtain signal-to-noise greater than a few hundred. The sensitivities on very bright sources also do not scale with brightness.
• In order to achieve the quoted sensitivities, you need to use long exposures. Typically > 30 seconds in H and K, and >75 seconds at J. Shorter exposures will result in a lower sensitivity. Please see CGS4 Exposure Times.

CGS4 sensitivity per pixel with the 40 l/mm grating and the long camera

The 3-sigma 30-minute sensitivities per 0.6″x0.6″ pixel for the case of nodding up and down the 1-pixel wide slit are as given in the following table. Sensitivities for very extended objects which require nodding to blank sky will be 0.4 mag poorer. The line flux assumes that the line is unresolved. Please remember to reduce the sensitivity by one magnitude for point sources.

wavelength (µm)	order	magnitude	line flux (W/m**2)
0.9	2	19.2	2.0e-19
1.25	2	19.8	5.0e-20
1.6	2	18.0	9.0e-20
1.6	1	19.3	6.0e-20
2.2	1	18.7	4.0e-20
3.4	1	12.4	1.0e-18
3.8	1	12.4	1.6e-18
4.9	1	10.9	2.0e-18

Predicted CGS4 sensitivity per pixel for the 150l/mm grating and the long camera.

The 3-sigma 30-minute sensitivities per 0.6″x0.6″ pixel for the case of nodding up and down the 1-pixel wide slit are as given in the following table. Sensitivities for very extended objects which require nodding to blank sky will be 0.4 mag poorer. The line flux assumes that the line is unresolved. In the J-K windows the predicted sensitivity is based on a model for the dark time inter-OH sky background and the throughput of CGS4. There will be variation in this background depending on lunar phase. If you observe an astronomical line which is not resolved from an OH line, then the limiting magnitide will be less, depending on the strength and variability of the line. The limiting magnitude is approximately one magnitude brighter on a typical OH line. Please remember to reduce the sensitivity by one magnitude for point sources.

wavelength (µm)	order	magnitude	line flux (W/m**2)
1.25	3	19.4	1.1e-20
1.65	2	18.6	1.2e-20
2.2	2	17.6	1.1e-20
3.8	1	11.9	6.1e-19
4.9	1	9.9	1.5e-18

wavelength (µm)	order	magnitude	line flux (W/m**2)
1.25	3	18.3	3.0e-20
1.65	2	17.0	6.1e-20
2.2	2	16.5	3.2e-20
2.2	1	16.5	6.4e-20

CGS4 sensitivity per pixel with the echelle grating and the long camera

The 3-sigma 30-minute sensitivities per 0.6″x0.6″ pixel for the case of nodding up and down the 1-pixel wide slit are as given in the following table. Sensitivities for very extended objects which require nodding to blank sky will be 0.4 mag poorer. The line flux assumes that the line is unresolved.  The values assume that the observations are read noise limited with sufficiently long exposure times for the multiple NDR to provide a read noise of about 25e-. The sensitivity will be less if an exposure time less than 100sec is used to avoid saturating a strong OH line close to a faint astronomical line. Please remember to reduce the sensitivity by one magnitude for point sources.

wavelength (µm)	magnitude	line flux (W/m**2)
1.25	15.0	9.0e-20
1.6	14.6	8.0e-20
2.2	14.0	6.0e-20
3.8	10.6	3.0e-19
4.9	8.9	7.0e-19

CGS4 Exposure Times

Since the installation of the long focal length camera in August 1997, many observers have requested that optimum exposure times be placed in the CGS4 web pages and in the manual. This is not a trivial task, as the optimum exposure time is dependent on many factors such as resolution, wavelength, object brightness and weather conditions. In this article, we provide a guide on how to select appropriate exposure times given the above factors.

NB. Generally, the maximum possible exposure time is the optimum exposure time, as overheads are reduced to a minimum. However, you should discuss this with your support scientist before and during your run.

150 l/mm Grating

Approximate maximum exposure times (sec) for the 150 l/mm grating (J, H, K)

Assume a 1-pixel wide slit and that the light falls on one row of the array. Typically light falls over three rows and these exposure times can be increased by about 30%. Half these times when using the 2-pixel wide slit.

	Magnitude
Wavelength (µm)	1	2	3	4	5	6	7	8	9	Strong OH1
J (1.2)	0.16	0.40	1.0	2.4	6.0	15	38	95	238	600
H (1.7)	sat	0.14	0.36	0.90	2.3	5.8	15	36	90	100
K (2.2)	0.14	0.34	0.84	2.1	5.3	13	33	83	208	120

¹ The strongest OH line in the band will be saturated using this exposure time (on a good night).

Approximate maximum exposure times (sec) for the 150 l/mm grating (J, H, K)

Assume a 1-pixel wide slit and that the light falls on one row of the array. Typically light falls over three rows and these exposure times can be increased by about 30%. Half these times when using the 2-pixel wide slit.

¹ The strongest OH line in the band will be saturated using this exposure time (on a good night).

Approximate maximum exposure times (sec) for the 150 l/mm grating (L, M)

This table is for 1-pixel wide slit and the long camera. Half exposures for 2-pixel wide slit.

* These figures refer to the 256 x 32 subarray. There is also a slightly larger 256 x 48 subarray with a minimum exposure time of 0.023 sec. Therefore the brightest observable magnitude is 0.4m fainter when using this array.

150 l/mm grating and background limited exposures.

This medium-to-high resolution grating enables one to work in between OH lines in many regions in the 1.1 to 2.3 micron spectrum (OH line emission is not a factor beyond 2.3 microns).

In simple terms, in order to be background limited in the non-thermal regime (<2.3 microns), the sky noise must be greater than the array read noise. For multiple non-destructive reads (NDR), the read noise is approximately 23 electrons. A typical value for the continuum background (from the telescope, sky and long wavelength leaks) using the 1-pixel wide slit in J, H and K is 30 counts in 100 seconds, and with a gain of six, corresponds to 180 electrons and a sky noise of approximately 13.5 electrons. Therefore, in most cases, an exposure time of approximately 300 seconds is required for the sky noise between OH lines to equal the read noise. For longer exposures than this the array is background limited and best s/n is achieved. With the two pixel wide slit (which still gives high enough resolution to work between many OH line pairs) the exposure time must be greater than about 200 seconds. Beyond about 2.2 microns, the background increases rapidly as the thermal background from the sky and telescope begin to increase, and the background-limited exposure time drops rapidly.

The drawbacks to using such long exposures are variations in the sky background and OH line intensities, OH line saturation (at H and K), and increasing likelihood of spikes on individual or small groups of detectors. If the critical wavelengths are well clear of the OH lines, then you probably don’t have to worry about their approximate 5-10 minute variation timescales or strength for OH variations, but if you are close to one than these can become problems (but see the two paragraphs below for ways to minimise these). Add to this that you will probably need to oversample your spectra, your time on source will become at least 600 seconds before you nod to sky. If you are using the 1-pixel wide slit and 2×2 sampling, it will be 20 minutes before you can nod the telescope. In addition to the dangers of sky variations these long times mean that, although in principle maximum sensitivity is achieved, a lot of time is wasted if something goes wrong.

For extended sources, nodding to sky is required, and the brightness of the source and stability of the sky background on that night will effectively determine the time between nods and hence the exposure time; these may be considerably less than the above ideals.

For spectra obtained while nodding along the slit, subtraction of the negative spectrum from the positive spectrum will remove most of the sky and OH fluctuations because both vary slowly across the rows of the array. When observing faint and compact sources it is always advisable to nod a small number of rows along the slit (e.g., much less than the canonical 30 rows), so that the cancellation of sky and OH residuals is as accurate as possible. Remaining residuals can be removed by polyfitting techniques, using blank sky rows adjacent to the rows of interest, but doing this will increase the noise in the final spectrum.

The frequency of spikes is difficult to judge and their effect difficult to assess, because spikes sometimes are severe, sometimes are only somewhat above noise levels, sometimes effect only one pixel, sometimes effect a few adjacent pixels, and their frequency may vary. Clearly they are more likely to affect observations of an extended source than a pointlike source.  Empirically they do not appear to be a serious problem when observing point sources with exposures of a few hundred seconds.

40 l/mm Grating

Approximate maximum exposure times (sec) for the 40 l/mm grating (J, H, K)

Assume a 1-pixel wide slit and that the light falls on one row of the array. Typically light falls over three rows and these exposure times can be increased by about 30%. Half these times when using the 2-pixel wide slit.

¹ The strongest OH line in the band will be saturated using this exposure time (on a good night).

Approximate maximum exposure times (sec) for the 40 l/mm grating (L, M)

This table is for 1-pixel slit and the long camera. Half exposures for 2-pixel wide slit.

* These figures refer to the 256 x 32 subarray. There is also a slightly larger 256 x 48 subarray with a minimum exposure time of 0.023 sec. Therefore the brightest observable magnitude is 0.4m fainter when using this array.

40 l/mm grating and background limited exposures.

In most cases the 40 l/mm grating is background limited at much shorter exposure times at all wavelengths than the 150 l/mm grating due mainly to its lower resolution, which ensures that an OH line is present in almost every resolution element.

Typical background-limited exposure times are about 30 seconds at H and K (less than 2.3um), giving 2 minutes between nods with 2×2 sampling. In the J band the OH lines are weaker and exposures of ~75 seconds are required to reach the background limit (5 minutes between nods). The same concerns and optimal procedures regarding OH fluctuations as discussed for the 150 l/mm grating apply here, except that spikes are less of a problem because super-long exposures are not needed.

The Echelle

Optimum exposure times for the echelle are difficult to predict due to the small wavelength coverage of this grating, the exact wavelength of a particular observation and the changes in sky transmission with wavelength. The best course of action is to consult with your support scientist or with; Paul Hirst before your CGS4 run.

Resolutions and wavelength coverage (with the long camera)

40 l/mm grating¹:

order	resolving power	coverage (um)
2 (0.95 – 1.40 um)	800 x lambda	0.32
1 (1.30 – 5.50 um)	400 x lambda	0.64

150 l/mm grating¹:

order	resolving power	coverage (um)
3 (1.07 – 1.30 um)	4700 x lambda	0.055
2 (1.25 – 2.50 um)	3100 x lambda	0.08
1 (2.10 – 5.50 um)	1550 x lambda	0.16

Echelle^1,2:

• Resolving power is 37000 at the optimum order.
• Coverage is 0.006 x lambda.

1: These are the resolving powers for the 1-pixel wide slit (0.609″). Half the values when using the 2-pixel wide slit. They will be x4 lower when using the 4-pixel wide slit for extended object, but probably only x2 lower for guided point sources.

2: There is no 4-pixel wide slit available for use with the echelle.

40l/mm and 150l/mm grating efficiencies

40 l/mm

Use the following plot to decide which order to use when observing with the 40l/mm grating in CGS4. Use 1st order beyond 2.5 microns.

The data are from Richardson Grating Laboratory for 1st and 2nd order. 3rd order is estimated.

The above plot shows only the grating efficiencies themselves. It does not include atmospheric transmission and the throughput of the CGS4 filters. For example, if one wished to observe a feature a 1.35 microns while setting the central wavelength for the middle of the H band (around 1.6-1.7 microns), you would expect essentially no throughput at 1.35 microns. This is because i) the atmospheric transmission itself is only about 20% at this wavelength and ii) the B2 filter, used for H and K band work, cuts off at about 1.4 microns.

For those wishing to observe features that occur close to or at the cut-off wavelengths for the near-infrared bands, please contact your support scientist or Paul Hirst beforehand for advice on the most appropriate methods.

150 l/mm

CGS4 Filters

The following two tables list the filters in use with CGS4, their central wavelengths and cut-on and off wavelengths (half power).

Broad Band Filters (for low and medium resolution gratings):

Filter Name	Central Wavelength (µm)	Cut-on Wavelength (µm)	Cut-off Wavelength (µm)	Throughput at Central Wavelength (%)	Estimated Average Throughput1  (%)
IJ	0.94	0.84	1.04	55	48
B1	1.29	0.99	1.59	83	84
B2	2.02	1.43	2.60	92	89
B3	3.15	2.14	4.16	90	93
B4	3.50	2.75	4.24	90	88
B5	4.92	4.35	> 5.5	93	92
B6	2.21	2.00	2.41	78	74

¹ Estimated in the central 75% of the bandpass.

Narrow Band Filters (echelle mode):

Filter Name	Central Wavelength (µm)	Cut-on Wavelength (µm)	Cut-off Wavelength (µm)	Throughput at Central Wavelength (%)
N1	1.089	1.078	1.099	57
N22	1.229	1.218	1.240	62
N3	1.259	1.247	1.272	62
N4	1.275	1.264	1.286	59
N5	1.293	1.283	1.305	66

² The N2 filter is not currently installed in CGS4’s filter wheels.

For wavelengths beyond 1.3-µm, the CVFs are used as order-blockers when using the echelle. When the new wedged CVFs are installed, the narrow band filters will no longer be necessary, as one of the CVF segments will extend down to 0.93 um.

CGS4 Pixel Sizes

Low resolution gratings

Grating	Perpendicular to slit	Along slit
40 l/mm	0.61″	0.61″
150 l/mm	0.595″	0.625″

Notes:

1. The 40 l/mm grating pixel sizes are independent of wavelength setting.
2. The 150 l/mm grating pixel sizes are virtually independent of wavelength setting.

Echelle Grating

Because of anamorphic magnification, the Echelle spatial pixel scales vary with grating angle. The 0.41″ probably changes with grating angle too, but I don’t have measurements of that.

Grating Angle	Perpendicular to slit	Along slit
57.1 deg	0.41″	0.783″
62.6 deg	0.41″	0.850″
64.4 deg	0.41″	0.887″
66.0 deg	0.41″	0.93″

Setting Slit Position Angles (PA)in CGS4

When setting the slit PA in the CGS4 software, it is important to remember:

1. The array is displayed `upside down’ on the movie screen so that south is to the top, and north to the bottom.
2. You need to consider in which direction the nod, or slide, should be specified.
3. Although you are asked to enter a PA which is east of north in your config, you normally have to enter a negative PA. This is discussed in a little more detail below.

The best way to demonstrate how to specify the slit PA is by showing some examples which is done below. In each example, the top and bottom of the array are marked, as well as north, south, east and west. An arrow indicates the direction of a positive slide (or nod). Remember the telescope moves in the opposite direction to this when you slide along the slit, so be specific when giving instructions to the telescope system specialist. Please be aware that, in order to try and show things as the observer sees them on the array when the slit is north-south, north is shown at the bottom of each plot.

Example 1: PA = 0 degrees (0 degrees east of north).

In this example, a position angle of 0 is entered into the config. the slit will set to a north-south direction, with the south to the top of the array. When using a positive slide, the offset beam appears to the south of the main beam (or towards the top of the array).

Example 2: PA = -45 degrees (135 degrees east of north).

Here, a PA of -45 degrees has been entered. The top of the array is now SE, the bottom NW. A positive slide moves the offset beam on the array towards the SE (i.e., the telescope moves NW)

Example 3: PA = -90 degrees (90 degrees east of north).

A PA of -90 degrees has been entered. The top of the array is in the east, the bottom is in the west. A positive slide will move the offset beam to the east, i.e., the telescope moves to the west.

Example 4: PA = -135 degrees (45 degrees east of north).

A PA of -135 degrees is entered. The top of the array is now in the NE, while the bottom is in the SW. A positive slide will move the offset beam to the NE, i.e., the telescope moves to the SW.

Some further notes on slit PAs

• The default slide_slit command in the quad_slide exec is +11.4 arcseconds. This positions the offset beam 19 rows up from the main beam. To move in the other direction (towards the bottom of the array), simply enter a negative value for the slide. Remember that the slit is 90 arcseconds long You normally do not want to specify a nod which places the offset beam off the array.
• Although the config demands a position angle east of north, it is only possible to enter such an angle to a maximum of about 10 degrees. To set to an angle greater than this, you must specify a negative angle. For example, if you wish to set to a PA of 50 degrees east of north, set to an angle of -130 degrees (simply subtract 180 from your desired PA).
• The pixel size is 0.61 arcseconds. To make sure you nod onto another row, you need to specify a value in arcesonds which is a multiple of this number.
• The system usually does a fairly good job of positioning the crosshead on the slit in the offset position. To increase the accuracy of the offset position beam, you can do one of two things (or both): i) Use a shorter slide distance than the default +11.4 arcseconds, or ask the TSS to peakup in both the main and offset beams and correct the current slide angle.
• The current r.m.s. setting accuracy of the slit PA is 0.15 degrees.

And please….

• Always inform the telescope system specialist when you change the PA. The peakup position will change with PA, and if the TSS is unaware of a change, time will be lost in reaquiring your target.

NOTES ON USING THE ECHELLE GRATING IN CGS4.

When using the echelle grating there are more details that need to be checked than with the low resolution gratings. These are summarised in this note.

1. Pixel and Slit Sizes.

Due to anamorphic demagnification, CGS4’s square pixels do not view a square patch of sky. Although the solid angle seen by a pixel is constant, the shape is a function of grating angle. The difference from a square field of view is very small at the low angles of the first order gratings, but is quite prominent at echelle angles. For example, at the nominal blaze angle of 64 degrees pixels are more elongated in the spatial direction than the dispersion direction by a factor of ~1.6. With the long camera at angles close to optimum the pixel size is 0.91 arcsec (spatially) X 0.40 arcsec (spectrally). The nominally 1-pixel wide slit is in fact about 1.3 pixels wide. The 2-pixel wide slit is 2-pixels wide (ie. 2arcsec spectrally).

The change in shape of the pixels as a function of grating angle is very rapid at large angles. As a consequence, when using the echelle you should always measure the size of the pixels by peaking up in two rows. Edit your sequence to put in the correct offset between the rows being used when nodding (see also the comment on slit alignment below).

2. Curvature

The slit is slightly curved, so atmospheric and arc lines are not quite perfectly aligned on one column of the array. In addition the dispersion axis is also slightly curved, so that the spectrum of a point source is not perfectly along one row of the detector. Both of these effects are present with the low resolution gratings, but they are more significant with the echelle. The curvature of the dispersion direction means that the dispersion has slight dependance on row on the array. This means that if an arc/atmospheric/astronomical line is straight along a column of the array in the middle of the array, lines at the edge of the array will not be so well aligned. In the dispersion direction the curvature amounts to approx 0.5pixels across the 256 pixels. In the spatial direction the misaligment with the columns due to curvature is about 0.1-0.5 pixels close to the centre of the array, depending on grating angle. For point sources nodded along about 30 rows the effects are very small and it is still reasonable to combine the two nodded spectra to cancel sky lines.

There is software available, for example in Figaro, which was designed to remove these sorts of distortion for optical spectrographs. There is information about using these in the CGS4 data reduction notes in correcting curvature in CGS4 spectra . Using these techniques a residual distortion of less than 1% can be obtained. However, in conditions of good seeing, CGS4 data are undersampled in the spatial dimension and this may mean that the effects cannot be fully removed.

Slit alignment

The curvature of the slits means that even if the postion angle is nominally N-S on the sky there will be a small (0.3 – 1 arcsec) E-W offset when you nod along the slit. Since this offset depends on the grating angle and how far along the slit that you choose to nod, it means that you should measure it for each echelle wavelength/order that you observe at. After you have measured the offset between the two rows you need to update your sequence accordingly.

INFORMATION FOR CREATING CONFIGS

3. Optimum Orders

When using the Echelle it is normal to select ECHELLE_AUTO_ORDER as the grating when defining a config. If this is selected then the CGS4 software will automatically select the optimum order for your wavelength. The optimum order gives a higher throughput than other orders. The optimum ordering software has recently been improved and now works very well for all wavelenghts.

You may wish to use another order, e.g. to get lower or higher spectral resolution or to increase the spectral coverage. To use a different order than the optimum one selected automatically, choose ECHELLE as the grating when defining your config and then explicitely enter the order that you want.

You can calculate the optimum order for any wavelength as follows. The blaze angle of the echelle is 64.6 degrees. This corresponds to the product of wavelength and order (n x lambda) being about 55.0 microns. (e.g., we find that the transmission at 2.12um in 26th order (n x lambda = 55.1) is 10X higher than in 25th order (n x lambda = 53.0) In general for any wavelength, the order for which n x lambda is closest to 55.0 gives the highest efficiency. However, at lower orders (longer wavelengths), it is best to fudge this somewhat, as the efficiency drops off more rapidly at higher angles than at lower angles. The CGS4 software uses a lookup up table calculated according to the above, with appropriate adjustments at the longer wavlengths, to select the order for the echelle.

4. Order sorting

Longward of 1.3um, CVFs serve as order-blockers for the echelle. Shortward of 1.3um, narrow band filters (1.083um, 1.233um, 1.257um, 1.282um), allow echelle observations at important wavelengths. Wavelengths shortwards of 1.3um may only be observed through these narrow band filters . The longest wavelength observable with the echelle is 5.7um. The CGS4 software will automatically select an appropriate filter or set the CVF to the requested wavelength. It will report an error if your chosen wavelength is outside the range of either the narrow band filters or the CVFs. Note that a long wavelength leak is suspected in the short wavelength CVF (roughly H-band), and is under investigation. If you want to do Echelle spectroscopy in the H-band please check with your support scientist for an update on this situation.

5. CVF gradients

There is a slight gradient in the transmission of the CVFs along the slit. The CVF calibration has been set for the middle of the illuminated area, row 134. If you want to use rows towards the edge of the illuminated region or nod more than about 30 pixels you may wish to check the CVF calibration for the rows you will be using. First of all define an astronomical config for the echelle in the normal way and set to this config. Peakup a star on the desired row or look at a lamp line and then run MOVIE. Now while MOVIE is running go into the menu called DIRECT_MOTOR_CONTROL . This menu allows you to define an intermediate configuration. The menu items diplayed represent where the CGS4 motors are currently positioned. To check the CVF calibration try changing the CVF wavelength by a very small amount and check whether the signal on the MOVIE display increases or decreases. Once you know what wavelength you want to set the CVF to calculate the difference between the grating wavelength and the desired CVF wavelength. You can then use UKIRT_PREP to save an astronomical config with this CVF offset. Be very careful if you decide to make a change like this – it is possible to get “lost” in order on the CVFs because at 1-2.5um the Echelle orders are very close together. For example 2.2um in 25th order is at the same grating angle as 2.11um in 26th order – so if you move the CVF by as much as 0.09um you could be looking at the wrong wavelength and order.

6. CVF fringes

Particularly in the thermal IR a ripple is seen in echelle spectra which is caused by fringing from the CVF. This ripple can be difficult to remove if there are amplitude variations between your source and your star. If you observe very strong ripples, try moving the CVF by a very small amount, or try puttting your source slightly out of focus. I think the latter helps because it makes the source and the background fill the slit to the same degree. Also take oversampled flats in preference to the usual undersampled one.

7. Wavelength Calibration

Because of the narrow wavelength range covered by the array when the echelle is used, there are wavelengths where it is impossible to find lamp lines that fall on the array. In this case you may be able to find lamp lines at different wavelengths that are present at higher or lower order at the same echelle setting. I.e., for such a line of wavelength lambda there is an n that gives the same n*lambda as your observing setup, For example, you want to observe at 3.00um in 16th order, where there is no line, but notice that there is an Argon line at 2.40um which in 20th order would appear at the same echelle angle. Such lamp lines can be found by (1) setting the echelle to the wavelength you wish to observe (ie in this example 3.00um) with 16th order selected and (2) changing the arc CVF wavelength to the wavelength of the calibration line. The arc section of a config allows you to enter a different CVF wavelength for observing arcs than for observing your source. This arc CVF wavelength will only be used when you take arcs.

At some other wavelengths you will see lamp lines that were not expected – these are strong lines in a different order and wavelength being transmitted through the wings of the CVF transmission profile.

At wavelengths beyond the K window, observable lamp lines are generally few and far between. For calibration with the echelle, one often must use the above technique for finding arc lines in different orders, telluric absorption/emission lines (atlases are available in the control room at HP, and at JAC), or observations of astronomical line sources.

The Echelle and Peakups etc

There allways seems to be confusion about the Echelle grating and the need to do 2 row peakups etc. Hopefully, this text outlines the definative proceedure.

There are two complications with the Echelle that make this more complex than with the low-res gratings, firstly the curvature of the slit, and secondly, the fact that the pixel size varies significantly with echelle grating angle.

Here’s my suggested proceedure, assuming that you’re nodding along the slit.:

Initially, set up your ORAC Science Program with the default offsets.

Slew to the target, and ask the TSS to do a 2 row peakup, using the SAVEPEAK and SAVE_PA commands. You will need to tell them which two rows you would like them to peak up on. I suggest 140 and 160, the same as we use for the low res gratings. It will help the TSS if you can tell them the approximate offset in arc-secs between these two rows. To do this, look up the grating angle on one of the VAX terminal screens once you have set to your config, then go to the table in the “Pixel scales” section of this document and estimate the pixel scale for the grating angle that you’re at.

The SAVE_PA command the TSS uses on the TCS tells the TCS the correction between the demanded slit angle in your config and the actual angle on the sky of the point joining the position of the star when placed on the slit at the two points intersecting the peakup rows. This means that you can simply ask ORAC to “offset along the slit” and it will know the exact angle to go to hit the slit on your two peakup rows.

Note that if you change to a different slit angle, there should be no need to repeat the two row peakup. Simply do a normal single row peak up on row 140 – the correction to the angle will be exactly the same as before, assuming that you’re still offseting the same distance and that you haven’t changed the wavelength or order settings of the grating. If these conditions (ie same configuration, same size offset) do not apply, then you will have to repeat the 2 row peakup in your new configuration.

It is advantageous to your signal-to-noise ratio if you have the spectra centred on pixel rows on the array. Peakup also helps you do this. After doing the 2 row peakup, the TSS will be able to give you the exact size of the offset between your two peakup rows. You should go into the offset iterator in your sequence and replace the default offset size (usually 11.74) with this number. Again – there is no need to enter an offset component perpendicualr to the slit unless you really do want to offset perpendicular to, and hence off, the slit. This offset size will apply to all observations in the same CGS4 config, no matter what the slit position angle.

Notes on Configuration parameters

GRATING: Select the grating to use. Only two gratings are in CGS4 at any one time – if in doubt your support scientist or telescope operator can advise which they are. When observing with the Echelle you should normally select ECHELLE_AUTO_ORDER.

WAVELENGTH: Enter the central wavelength for your spectrum in microns.

ORDER: Enter the order to use the grating in. If you have selected ECHELLE_AUTO_ORDER the value here will be ignored.

The CGS4 sensitivity page contains information about the appropriate choice of orders for the gratings as a function of wavelength. Normally echelle_auto_order should be used for echelle observations. However see the echelle notes for details of where this is not appropriate. If in doubt your support scientist can advise you.

CVF_OFFSET_(ECH_ONLY): If necessary enter an offset in microns for setting the CVF when you are using the Echelle. Normally the default of 0.0 should be used. This value is ignored for the 40 and 150 l/mm gratings. The precise calibration of the CVF for Echelle observations can depend on which row you choose to observe, the order and the slit width. Sometimes it is desirable to set the CVF to a slightly different wavelength from the grating to correct for this. Your support scientist will help you check the CVF setting for your wavelength and order. Enter a non-zero value for a CVF offset with caution and only after checking flats taken with NO CVF offset.

POLARISER: The polariser should be NONE for normal spectroscopy. Select PRISM or GRID if you want to do spectropolarimetry.

You cannot do spectropolarimetry with the echelle and prism.

SLIT_WIDTH: Choose a 1-pixel, 2-pixel or 4-pixel wide slit. There is 1 pixel per resolution element and so a two pixel wide slit will give reduced, resolution, but can be useful in poor conditions. The 4-pixel wide slit is not available with the Echelle.

POSITION_ANGLE: Enter the position angle of the slit on the sky in degrees EAST of North. 0 deg is a N-S alignment. Due to a historcal reason, CGS4 will not accept angles greater than about 10 degrees. To request a slit position angle greater than this, you need to specify an angle 180 degrees away. E.g., if you want 45 degrees east of north, you need to specify -135 degrees as the position angle.

OBJECT_SAMPLING: The sampling required x the number of pixels to sample over.

There is 1 resolution element per pixel so the detector is stepped to fully sample. To Nyquist sample your spectra a sampling of 2 should be selected (ie 2 data points per resolution element). Many users prefer to slightly over-sample and choose a sampling of 3. To help eliminate bad pixels from your spectra you can also choose to carry out the sampling over 2 pixels. This means that each data point in your final spectrum has been observed with two neighbouring pixels. When the data reduction interleaves the raw integrations data from a bad pixel is replaced by that from its good neighbour. Most observers therefore chose sampling of 2×2 or 3×2 for this reason.

As an example, if you choose sampling of 2×2 the data in the final spectrum is taken at the following detector positions :

Except for at the begining and end, or where there is a bad pixel, each data point in the final array is the average of two measurements. For 2×2 sampling the reduced spectrum has 514 data points on the x-axis.

DATA_ACQUISITION_MODE: Choose ND_STARE, STARE or CHOP.

For exposure times less than 1 STARE is a little more efficient but NDSTARE may aso be used. Use ND_STARE for all other observations that do not require chopping. ND_STARE uses a multiple non-destructive read algorithm to give optimum read noise performance.

SINGLE_EXPOSURE__SECS: Enter the on-chip exposure time in seconds to be used for observations of your source. See the optimum exposure times for a guide to what value to enter here. The minimum exposure time is 0.12secs for the 256×178 array. For the 256×48 array, the minimum time is 0.023 secs, and for the 256×32 array, the minimum time is 0.016 secs.

EXPS/INT_OR_EXPS/CHOPBEAM: Enter the number of exposures to be coadded into an integration (integ) at each array sample position.

For long exposures on faint sources set object_exp_per_integ to 1. For very short exposures on a bright source it is a good idea to choose a value so that the time per integration is a few seconds for best observing efficiency.

CHOP_CYCLES/INTEGRATION This only applies to CHOP mode, and is ignored otherwise. Enter the number of chop cycles you want before nodding the telescope (one chop cycles is an object-sky pair).

FLAT_SAMPLING: choose AS_OBJECT in most cases

For Echelle users it is now possible for the data acquisition to take over-sampled flats with the same sampling factor as your object observations to help remove CVF ripple.

FLAT_LAMP: choose your required lamp. The numbers refer to black-body apertures. There is a table of flats of appropriate values as a function of grating wavelength and order. If you choose off then the calibration unit will not be used – useful for if you want to observe the sky, or the dome etc for a flat.

FLAT_ACQUISITION_MODE: Normally this will be NDSTARE

FLAT_EXPOSURE_SECS: typically 0.12 – 1 sec. See the table of flat exposure times.

FLAT_EXP_PER_INTEG: The number of exposures to be coadded into an integration. typically 50-100.

FLAT_INTEG_PER_OBS: The number of integrations to be taken and returned to the Vax. Typically 1.

ARC_LAMP: Choose your required lamp. There is a table of recommended arcs lamps as a function of wavelength. Plots of sample arc spectra can also be found at that page. If you choose OFF for the lamp then the calibration unit will not be used – useful for if you want to observe sky lines for wavelength calibration.

ECHELLE_ARC_CVF_WAVELENGTH: Defaults to the same wavelength as for your object observation.

Because the wavelength coverage on the array with the Echelle is not very large it is possible to observe at wavelengths where there are no lamp lines which would fall on the array. This allows you to select a different CVF wavelength from the grating, to look at lamp lines in a different order from your source.

If the value shown here is -1 the default CVF setting – ie the same CVF wavelength as for your object will be used. If you have just recalled or created a CONFIG then the default wavelength should appear here.

ARC_ACQUISITION_MODE: Usually NDSTARE

ARC_EXPOSURE_SECS: Typically 0.12 secs with the 40 l/mm grating and long camera.

ARC_EXP_PER_INTEG The number of exposures to be coadded by the into an integration at each detector sampling position. Typically 5.

All ARCS are automatically taken with the same sampling and sampling range as your object observations. Although on-chip exposure times are short, you should add sufficient exposures so that the time per sampling position is at least 0.5 secs, to prevent the background increasing due to the detector translation stage getting warmer.

ARC_INTEG_PER_OBS The number of integrations to be taken and returned to the Vax at each detector sampling position. typically 1.

CGS4 Latency

It has come to our knowledge that the 256 x 256 InSb array in CGS4 suffers from small scale  latent images. This problem has only become noticeable since long integrations have become possible due to the small pixel sizes in use with the long focal length camera.

In general, the latency only affects long exposures of very faint sources after previously observing a bright source (e.g., a standard star). The level of latency varies, but is usually less than about 5 counts (30 electrons) above the general background level.

The simplest method to rid the array of the latent image is to ask your TSS to run manpeak to take a few short exposures without changing the apperture, once you are pointing at your faint source.

If your targets are bright enough to peakup on, then the above procedure is unnecessary. The process of peaking up will remove any latent image.

CGS4 Object acquisition Procedures

Peak-up procedures with CGS4 seem to be a large source of confusion, especially for observers determining before hand what their requirements are in terms of astrometry, co-ordinates and reference stars. This page aims to clear things up a bit.

The complication arises in that there are many ways to attempt to ensure that you have accurately hit your target with the slit, all of which apply to different situations.

Concepts and Introduction

The aperture defines the position of the slit within the telescope focal plane. The default aperture is re-determined during engineering time whenever the instrument has been off the telescope, and at regular intervals inbetween. The aperture varies slightly with flexure etc. Fine adjustments to the aperture definition are made with the peak-up procedure.

The actual peak-up procedure is the process whereby the aperture definition is finely adjusted to give the maximum signal on the CGS4 array, whilst pointing at some (bright) target. This involves taking several (typically 15 or so) CGS4 integrations. Each of these must convincingly detect the object being peaked up on. It is possible to automatically subtract sky frames from the images used for peak-up. The integration time needed depends on the instrument configuration in use and on how bright your target is.

With the 40l/mm grating, peaking up on a target fainter than 14th magnitude takes a long time. If you need to estimate peak-up times for fainter targets or other configurations, refer to the CGS4 sensitivity page, and use the following: Calculate the exposure time for a 3 sigma detection of a point source at your target’s magnitude. Multiply this by a factor 2 (to account for the fact that for most of the frames, you’re not peaked up, so you’re only seeing a fraction of the light). Multiply this by 15 (say 15 observations needed).

The Fast Guider is mounted on the cross-head. The cross-head can position the fast guider very accurately (~0.1″) within the telescope focal plane. The crosshead limit is roughly a circle, 180″ from the pointing centre (ie centre of the slit).

The Fast Guider brightness thresholds also vary with conditions. Because the fast guider works in the optical is is affected by scattered moon-light. In exceptional conditions – clear skies, exceptionally good seeing and no moon, we can guide on stars down to V magnitude of 18.7. A V mag limit of 16 or 17 is more realistic for average conditions. On fainter targets, we increase the integration time of the autoguider CCD, which reduces the Tip-Tilt frequency. On bright targets, we guide at 100Hz. We can go down to ~20Hz for faint sources. Bright sources are more likely to give you really good Tip-Tilt performance.

UKIRT employs a dichroic tertiary mirror, feeding the IR light to the instrumentation, and the optical light to the fast guider. This means that it is possible for the fast-guider and the IR instrument (eg CGS4) to observe the same target.

Trivial case – bright (~14th Mag or brighter in the IR) guidable targets.

In this case you guide on the target, and you peak up on it before you start taking data. The peak-up takes a few minutes.

No special co-ordinate, astrometry or guide star requirements.

Bright (14th Mag or brighter in the IR) non-guidable targets.

Reasons they might be non-guidable include: Target not point-like in the optical, or target is so red that despite being bright in the IR, it’s too faint for the fast guider to see (see notes above).

In this case, you guide on an offset guide star (needs to be guidable and within 180″ of target), and you peak up on the target.

Requirements: A guide star. You can pick one using the ORAC-OT preparation tool when you get to Hilo. You don’t need any special astrometry. If your co-ordinates are off by more than an arc-sec or two, peak-up will take longer.

Intermediate brightness targets.

OK, the definition of “intermediate brightness” is non trivial. By this, I mean something you can guide on, but something that would take too long to peak up on. See the notes above about the fast guider sensitivity and peak-up times.

Procedure: The TSS will slew the telescope to a nearby (within a few degrees) bright CMC star, and will peak up on that whilst guiding on it. This ensures that the slit and fast guider are at the same place within the telescope focal plane. Then they’ll slew back to your target and guide on it. Because peaking up on the CMC star aligned the guider and the slit, the slit will now be accurately placed on the target.

Note: The guider and slit will not move relative to each other over a small slew to the CMC and back. If it were a long slew, there would be concerns about flexure and differential refraction between the optical and IR beams.

Requirements: The TSS will locate a suitable CMC star as needed. You need reasonably accurate co-ordinates (ie within a few arc-sec). If the target is towards the fainter edge of the guider’s capability, better co-ordinates will help. If the source is not point-like enough to guide on, treat this as a faint target.

Faint targets.

By “Faint”, I mean too faint to guide on. This obviously implies also too faint to peak up on. There are several options in this case.

Blind offset

This is the most reliable way of acquiring very faint targets. You need an accurate astrometric offset between your target and a nearby (within 180″) guide star (see notes above on fast guider for what qualifies as a guide star). You would usually measure such offsets from a deep image of your target field. By “accurate”, I mean to within a fraction of the slit width you will be using. Say 0.2″ for the 0.6″ slit. The more accurate your offset, the more light you’ll get down the slit.

Proceedure: Simply enter the RA and Dec co-ordinates for the target and guide star in the OT. Ensure that the co-ordinates you enter have the correct offset between them, taking into account the factor 15 between seconds of RA time and seconds of arc. When you slew to the field, inform the TSS that you have an accurate crosshead offset set, and thus to adjust the telescope position rather than the crosshead position to bring the guide star into the fast guider if necessary.

The TSS will ensure that the telescope pointing model is good prior to you starting to take data. This will involve at least going to a nearby CMC star, and probably doing a peak-up on it. If necessary they will collimate the telescope pointing model to the CMC star co-ordinates at this point.

Requirements: Accurate astrometric offset to a nearby guide star.

Blind Pointing

This is the “last resort” method, though the UKIRT Telescope Control System and Pointing model is sufficiently good that this method usually works well.

You would use this method if you don’t have the astrometric offset to a guide star necessary to carry out the previous method. You do, however, need accurate absolute co-ordinates. You shouldn’t rely on this method with the 0.6″ slit. Many observers have had great success using this method to observe radio sources, with accurate co-ordinates from high resolution radio interferometers, using the 1.2″ slit.

Proceedure: Type your accurate co-ordinates into the OT. Use the OT to find a nearby offset guide star. The TSS will go to a CMC star close to your field and collimate the telescope pointing model. When they slew back to your field, inform the TSS that the target co-ordinates are accurate, and to move the crosshead to bring the guide star into the guider if necessary.

Threats: You need accurate, co-ordinates. You might need to be aware of the co-ordinate frame your co-ordinates are referenced to (how well is this frame tied to the Hipparcos or Radio reference frames?). This method becomes less reliable in high wind conditions – during the time the TSS is adjusting the guider position to acquire the guide star, the telescope is sat on it’s main drives, as opposed to being locked onto a guide star. If wind gusts are knocking the telescope around at this point, the drives are unlikely to be able to hold it in position to great accuracy.

Requirements: Accurate absolute co-ordinates.

Absolute Desperation

If your targets are too faint to guide on, and you don’t have either sufficiently accurate co-ordinates or astrometric offsets, then your best option is probably to use some time at the start of your observing run to get UFTI images of your fields, from which you can measure offsets in order to do a blind offset. You should discuss this with your support scientist and Paul Hirst, the CGS4 instrument scientist, well in advance of your arrival in Hilo.