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Commissioning of the spectroscopic mode of NOTcam

NOTcam was designed such that grisms and slits can be mounted and rotated into the beam to quickly switch from imaging mode to spectroscopy and vice-versa. We have done extensive tests in order to prepare for the spectroscopic mode. The tests were followed by first-light observations in August 2003, for the low-res spectroscopic mode of NOTcam. Since then, this mode has been offered to the community.

In May 2005, we obtained first-light observations for the medium-resolution spectroscopic mode of NOTcam. This mode can now also be applied for.

Nevertheless, further tests still remain to be done as part of the ongoing commissioning of the spectroscopic mode of NOTcam. Action points are indicated by a small dot, below.

Slits

Slits and other aperture-wheel items were purchased a several years ago. The slits are made of very thin nickel-copper laminate, with the slits cut in the nickel part. The slits were originally mounted in aluminium holders. When mounted in NOTcam the slits can be physically examined as they sit very close to the entrance window. It turned out that when cold, the slits are deformed due to the differential expansion of the copper and nickel, and the aluminium holders.

In 2004 new copper slit holders were installed for the 128 and 64 micron slits, which made the deformations less apparent.

One of the 4 slits is clearly wrinkled because of the above effect, but proves to be useful as the jaws remain quite parallel. Although the slit width is 128 micron according to specifications, it is imaged to about 2-2.5 pixels FWHM (instead of 4 pixels) on the detector by the Wide-Field Camera of NOTcam. This slit is usable for point sources.
The 64 micron slit, designed for 2-pixel resolution, is actually imaged onto much less than 2 pixels, and should not be used.
This plot shows flatfielded H-band lamp-light images of these two slits summed over the width of the slit, to indicate the deformation from parallel of the slit jaws. The slit transfer functions are normalized such that the Y axis of the plot indicates the effective slit width.

Also for the HRC mode we have only one useful slit, and also here the slit that was originally designed to be imaged onto 4 pixels, is actually giving about 2.5-pixel resolution. This plot shows the slit transfer function of the 44 micron slit, again normalised such that the Y axis of the plot indicates the effective slit width. The slit is suitable to be used for observations.

The grism

So far one grism has been designed by Michael Andersen, and this grism is installed in NOTCAM. The grism has an echelle grating and is designed to give intermediate resolution (2-pixel R=2500, with dispersion 2.5-4.1 Angstrom/pixel) in J (5th and 6th order), H (4th order) and K (3rd order) when used with the WFC. In combination with the HR Camera the resolution will be about 3 times higher, but the sampled wavelength range will be about 3 times less. The grism is used in combination with the JHK filter set in NOTcam. The grism was delivered in July 2003, and was mounted in a holder made by Niels Michaelsen in Copenhagen. This plot shows the throughput of the grism as measured by the manufacturer (note that the order overlap has not been accounted for).

  • Proposed upgrade: we could purchase low-resolution grisms spanning ZJ and HK bands.
  • An HK filter is already purchased for NOTcam (filter #226), but a ZJ filter needs to be ordered together with the new grisms.
  • Internal focus

    In order to determine the optimal focus, and the optimal distance between the aperture wheel items and the collimator, we made a set of Hartmann masks which were mounted in the stop wheel. Using the masks it turned out that there is a considerable chromatic focus variation. This plot shows the internal focus of NOTcam as a function of wavelength for a slit mounted 2mm towards the collimator in combination with the WFC (before the aperture wheel assembly was moved).

    With the grism in the beam the internal focus has to be upped by 200 units for the WF Camera (with respect to no grism).

    In the end of 2004, the aperture wheel assembly was machined to allow it to be moved 2mm away from the collimator (towards the sky). This was needed to be able to focus aperture wheel items (e.g. slits, pinholes) with the HRC. The tables below reflect the internal focus of WFC and HRC slits, after the movement of the aperture wheel assembly.

    HRC + 44-micron slit (no grism)
    band internal focusFWHMtelescope focus
    Z >6000 outside range 4.5 pixel
    J 5450 2.6 pixel
    H 3000 2.6 pixel
    K 1000 3.1 pixel

    WFC + 64-micron slit (no grism)
    band internal focusFWHMtelescope focus
    Z 4950

    J 4600

    H 4300 1.5 pixel
    K 4050

    WFC + 128-micron slit (no grism)
    bandinternal focusFWHMtelescope focus
    Z 5150
    ~23200 +/- 100
    J 4750
    ~23200 +/- 100
    H 4450 2.6 pixel~23200 +/- 100
    K 4250 3.4 pixel~23200 +/- 100
    Notes to table:

  • all Z-band values were obtained using filter #222  ( Yn )
  • for the 1mm pinhole + WFC use the same focal settings as WFC + 128-micron slit
  • for WFC the 50-micron pinhole cannot be focussed: in J the focus should be ~6700

    It turned out that in order to be able to properly focus the High-Resolution Camera, the aperture wheel assembly had to be moved away from the collimator, and the slits have to be mounted in a special holder such that they are as far away as possible from the collimator. The only way to achieve the mounting of this holder was to use one of the 4 big aperture wheel hole positions.

    The slits for use with the WFC, however, needed to be positioned further towards the collimator by about five millimeters (with aperture wheel moved) in order to get proper focus over the full ZJHK wavelength domain. The first-light observations were carried out in good seeing conditions, and indicated that for the WFC spectra there is no clear evidence for astigmatism caused by the positioning of the slits with respect to the collimator. For the HRC the on-sky measurements also show no evidence that the displacement of the slit adversely affects the imaging quality.

    Instrument flexure

    Using the small pinhole and the WFC, flexure tests showed that slit-shifts of up to 4 pixels can be expected when slewing to different sources at high airmass for random rotator angles. While keeping the rotator angle fixed, e.g. following one source using the parallactic angle, the slit shifts are usually less than 1 pixel. For the parallactic angle the shifts are mostly along the slit.

    These flexure slit-shifts are generally much larger than the shifts introduced by the (re)positioning of the aperture wheel. We have tested the repeatability of the slit position after moving the aperture wheel. We used the HR Camera with the telescope at zenith. Here, the slit ran parallel to the columns of the detector (Y direction), with the rotation of the aperture wheel introducing shifts in the X direction. Scatter between subsequent X-positions of the slit is about 0.1-0.2 HR pixels (0.01-0.02 arcsec), however sometimes a larger step of about 0.5 HR pixels occurs (0.04 arcsec). The large step may correspond to the motor of the aperture wheel losing track of 1 stepper motor unit.

    Considering the above there is no clear preferred direction for the slit, i.e. along detector columns or rows, as flexure-induced shifts are much more important than shifts introduced by the (re)positioning of the aperture wheel. It is clear that wavelength calibration frames will have to be obtained at each target position on the sky. Considering the detector cosmetics, we have decided that the preferred slit position is horizontal, such that the spectra run along the columns of the array.

    Note that for a horizontal slit the parallactic angle has TCS rot-pos = -90 or 90. For vertical slits: rot-pos = 0 or 180.

    Wavelength calibration

    NOTcam originally had no dedicated calibration unit. We investigated the possibility of using the OH sky lines for on-line wavelength calibration. For the H and K bands exposure times on the order of 200 seconds (for WFC) give sky-lines with enough flux that they are suitable for wavelength calibration. For very bright stars the on-target spectra (with short exposure times) will not be suitable to do wavelength calibration. These plots give identifications of the OH lines in vacuum for the H-band and for the Ks band. The scatter of the position of the line centroids around the wavelength solution (3rd degree polynomial) is less than 1 Angstrom.

    For the J-band exposure times on the order of 200 seconds do not give suitable sky-lines with enough flux. For this purpose we have experimented with calibration lamps normally used with ALFOSC. As part of the experiment the ZnAr and Ne lamps were hung above NOTcam at a distance similar to that of the flatfield lamp inside the telescope baffle. Although these particular lamps do not emit useful lines in H and K, they did provide lines with which the J band can be calibrated.

    As of 2005, we have a calibration unit in the telescope baffle, that can be used for NOTCAM spectroscopy. In order to use the lamps, the mirror-covers need to be closed, introducing extra overheads. The unit contains a flatfield lamp that is also used for ALFOSC, and a Xenon and Argon lamp. Currently the lamps can be switched on in the control room manually. Arc maps can be found here.

  • Proposed upgrade: a new calibration unit that swings in the beam fast, and that is completely operated under computer control, is desired.
  • Wavelength stability

    For the first-light data in the H-band the wavelength solution was not stable: the dispersion changing by 0.25% between exposures taken 1 hour apart, but at similar sky position and rotator angle. Could be a small tilt of the grism on the order of 0.2 degree. Or else flexure within the camera. Or movement of the slit from/to collimator? Needs testing!

    During further tests it was noted that when moving the telescope the wavelength zeropoint and the dispersion can change on the 1% level, which effectively means that both wavelength calibration arc frames and fringe correction flats have to be obtained at each and every telescope pointing.

  • Further test the stability of the dispersion.
  • Investigate whether the grism can be mounted better.
  • System efficiency

    Some preliminary results from spectroscpic standards can be found here.

    Flatfielding and fringes

    A flat field lamp is situated inside the telescope baffle. The lamp flats are a necessary means to get rid of the quite substantial pixel-to-pixel variations of the detector. From the first-light observations we found that one can achieve a signal to noise ratio of S/N=80 in the spectral continuum using standard reduction steps. The problem with the baffle lamp is that the mirror covers have to be closed in order for correct operation.

    Using the flatfield lamp inside the telescope baffle a clear ripple pattern becomes immediately present. These ripples are probably due to interference in the 330 micron thick sapphire detector substrate. The fringe wavelength is linear with wavelength of the incident light.

    bandfringe wavelengthspeak-to-peak amplitude (WFC+128-micron slit)
    old engineering arraynew science array
    J 5.5 pixel 17% 9%
    H 6.7 pixel 18% 10%
    K 8.8 pixel 19% 13%

    Fringe amplitudes are highest in the central parts of the array.

    The fringes are around 25% peak-to-peak when the HRC with the 44-micron slit is used (for the old engineering array !).

    Bad pixels, hot pixels, and hot rows

    Bad pixels, hot pixels, and hot rows are discussed in this report.

    On dark frames, and frames with low count levels, such as the sky part of spectroscopic images or on wavelength calibration images, hot rows are clearly visible (see this report or the first-light report). The count level in the hot rows does note scale with exposure time. In fact, the hot rows behave more like rows with a deviant zero level.

    The number of hot pixels increases with exposure time (see table below). For exposures of more than 5 to 10 minutes the hot pixels are so abundant that they may inhibit proper reduction of the spectra. The number of hot pixels for long exposures is similar when using either reset-read-read or ramp-sampling mode.

    Percentage of hot pixels (i.e. more than 8 sigma above background) in the central 250x250 pixels of the top-left quadrant (Q3).
    Exptime [s] 0 3 9 27 81 243729
    Hot pixels [%] 0.1 0.2 0.3 0.6 1.0 1.8 2.5

    An unfortunate feature of the new science array is the presence of clusters of hot pixels. In long dark exposures (see the 729 sec dark expo below) these patches of hot pixels appear as stellar images:

    Some of the clusters reach saturation and should always be avoided. Columns to be avoided are (in MEF format): x=110-120, 235-250, 390-400, 415-425, 715-765, 870-880, 945-955

    Calibration data for spectral reduction

    A simple ABBA nodding of the telescope is effective in removing the background. Subtracting the B exposures from the A exposures does leave a residual background level, while removing most of the small-scale detector background features. The residual background level can be fitted away if the nodding step is large enough. For very bright stars the wings of the spatial profile can be traced to +/- 30 pixels with the WFC. This means that the ABBA nodding step should be 20-25 arcsec for correct background subtraction. For fainter stars the nodding step could made be smaller.

    In order to get optimal extraction to work, the 2D images first have to be corrected using a bad-pixel map. Interpolation across bad pixels provides the first rough correction that the optimal extraction algorithm cannot handle. A bad-pixel map can easily be constructed from two flat fields with different exposure levels. This map should be (re-)made for each observing run. Even for very bright stars variance weighting and cleaning of deviant pixels is worth doing.

    The lamp flats are essential to get rid of the substantial pixel-to-pixel variations of the detector, and should be obtained after each telescope pointing.

    Wavelength calibration data should be obtained after each telescope pointing.

    For the first-light WFC observations it was possible to correct for telluric lines and at the same time for the fringes by using a nearby featureless comparison star. The S/N in the continuum was around 80 after this correction.

    For horizontal slits, the sky lines run along the rows of the detector. For sky-line based wavelength calibration one needs to extract these sky lines. The sky lines are bent because of optical distortion, and run into bands of hot rows. It should be possible to subtract out the hot rows using dark frames. Note that the counts in the hot rows do not scale with exposure time of the dark frames.

    The dark frames show different and non-linear background levels for different exposure times. It is a good idea to obtain a set of dark frames with the same exposure times as used for the science and other calibration frames.

  • We need to obtain more flux standard observations for throughput/efficiency determination.


    John Telting

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