2D-Frutti Reductions with IRAF
This file last updated on 17 Dec 1992
1. Introduction. 2. Getting the data. 3. Reducing the data. 3.1 Introduction. 3.1.1 Displaying 2D Frutti images. 3.2 Reading the data from tape. 3.3 Fixing bad pixels. 3.4 Creating the flat-field and dark images. Determining scale factors. 3.5 Distortion correction: Mapping the distortions. 3.6 Distortion correction: Transforming images. 3.7 Creating, normalizing and correcting the flat-field. 3.8 Applying the flat field correction. 3.9 Extracting the spectra. 3.10 Wavelength calibration. 3.11 Neutral Density filter corrections. 3.12 Deriving a sensitivity function. 3.13 Applying the sensitivity function.
- 1. Introduction
This document is intended as a guideline to reduce 2D-Frutti spectra using IRAF. If you do not have experience with the software we suggest that you start by reading "An Introduction to IRAF" which will give you a general idea of IRAF structure as well as a few fundamental tools.
The examples shown here are just that, examples. The user must decide upon the reduction procedure, naming convention, etc., appropriate for his own data and use the cookbook and examples as guidelines only. The examples are shown with prompts for the package containing the tasks (do not type the prompts, of course). It is strongly recommended that you perform an `lpar' on every task immediately before using it, unless you are familiar with all of its parameters. This is not always shown in the examples, but is normal practice even among seasoned IRAF users. Task setups are included as examples throughout the manual to help you and the most relevant parameters are highlighted throughout this manual.
IRAF uses two conventions that you should always keep in mind. First, images consist of lines and columns (not rows and columns). Second, the order of a function is the number of independent coefficients for a polynomial, or the number of intervals for a spline curve. For example, a cubic (third-degree) polynomial is described as "order=4" (four coefficients).
If you require personal assistance in your reductions please contact either Mauricio Navarrete or Nelson Saavedra on Tololo (ex 422), or Mario Hamuy in La Serena (ex 210).
- 2. Getting the data
Many observers often get confused about the number and kind of calibration frames that must be taken to properly reduce their data. During a whole run you should get,
+ at least one multihole frame taken during the daytime. + at least one long arc taken at the zenith before starting to observe. + as many dome flats as possible. These should be taken during the afternoon with the telescope pointing to the white spot. + as many sky flats as possible taken during twilight (or moonlight). + several dark frames that you should leave running after closing the telescope. + flux standard star(s) if you wish to flux calibrate your spectra. + extinction standards are also suggested if the highest possible photometric accuracy is desired. + short arcs bracketing your objects, if you wish to derive radial velocities.
Multiholes, dome flats, long arcs and sky flats are time consuming. Therefore we suggest that you start by taking the calibration frames some 4 hours before twilight, so that you have enough time left for dinner. See the instrument manual for more details about the number and types of calibration frames necessary.
3. Reducing the data
A full reduction of 2D-Frutti images requires the following operations (see the diagram on the next page):
1) Mapping the distortions using the long arc and multihole. 2) Distortion corrections using the mapping created. 3) Dark-substracting the flat fields. 4) For long slit work, the objects are also dark subtracted. 5) Flat-fielding your arcs and program objects (using dome and sky flats). 6) Extracting the spectra. 7) Wavelength calibration. 8) Deriving a sensitivity function. 9) Flux calibration.
Distortions are mapped by tracing profiles in the multihole frame and performing a two-dimensional surface fit. The same operation is done on a comparison lamp frame to produce a two-dimensional dispersion solution. These two mappings are used to rectify the images.
The total counts contained in an individual flat-field or dark frame are usually quite low, so that all the flats and darks are normally coadded to improve the signal-to-noise. Furthermore, since the dithering blurs features in the 2DF, it is possible to perform a boxcar-smoothing without loss of information while significantly decreasing the photon noise. Generally, the dark rate is quite low. This means that, if you are working with stellar spectra, it is only necessary to dark-subtract the flat-field frames. Flat fields are usually normalized so that photon counts are preserved after flat-fielding. For convenience, flat fielding and dark subtraction are usually deferred until after distortion correction. (Note that these steps are the normal procedure, but are not necessarily desirable in all cases. You should check with a staff member if you are unsure of what is appriopriate.) If you like more information about the flat field processing of your images, type help flatfields in IRAF.
One-dimensional spectra may be extracted at this point. The extracted spectra may now be wavelength calibrated, extinction corrected, and flux calibrated.
====================== = Fixing Bad Pixels = ====================== ***** *** * ========================================== = Coadding Domeflats Skyflats and Darks = ========================================== ***** *** * ============================ = Mapping the Distortions = ============================ ***** *** * ============================ = Transforming the images = ============================ ***** *** * ========================== = Creating a Flatfield = ========================== ***** *** * =========================== = Flatfielding the Data = =========================== ***** *** * =============================== = Extraction of the spectra = =============================== ***** *** * ============================ = Wavelength Calibration = ============================ ***** *** * ============================ = Flux Calibration = ============================
- 3.1.Displaying 2D Frutti images
SUN/IRAF has an image display window called imtool. An imtool window is one of the initial start-up windows on all CTIO IRAF accounts. Cursor readback is available, as the capability to zoom, and review images larger than 512x512 pixels.
The imtool window is accesible through the tv package in the images package. An image can be loaded into the window by typing
cl> images im> tv tv> display <imagename> <frame>where <frame> is the integer frame number (one through four). Move the mouse cursor into the imtool window and hold down the right button. Moving the cursor will change the display transfer function of the image to bring out subtle features of the image. Releasing the right mouse button quits this mode.
Pressing F6 within the imtool window alternately enables and disables cursor readout mode. When this mode is enabled the cursor coordinates are continually displayed and updated in a box in the lower right corner of the display window.
The 2D Frutti generates images larger than the default 512x512 pixels. The imtool window can be enlarged to display these pictures in their entirety. In the text window, type the command
tv> set stdimage=imt2df1to set the default size to 128x1520 pixels. Load an image with the above mentioned display command. Put the mouse cursor on the border region of the imtool window and hold down the right mouse button. Choose the option "FitFrame', and the window will now resize itself to 128x1520 pixels. You may wish to adjust the location of the imtool window by dragging it with the middle mouse button while the cursor is placed on a corner or edge of the window.
You may demagnify, for instance, the image by a factor of two with
tv> display <imagename> 1 xmag=0.5 ymag=0.5Pressing the middle button of the mouse within the imtool window centers the image at the cursor position. This feature may be particularly useful if your image does not fit in the window.
Pressing the button twice without moving it will zoom the window and, if the cursor is not moved, pressing it again will zoom it further. This continues until the cycling is completed.
- 3.Reading the data from tape
Load the dataio package, allocate the appropriate tape drive, mount your tape and do an epar on the task rfits according to the parameters given below. Remember to remove the write ring.
cl> dataio da> epar rfits (check the list given below) da> alloc mta
Read the data using the task rfits. You must specify the tape-drive name, the list of files you wish to read in and the root file name. In choosing the root file name, it's usually a good idea to include a digit in the name to indicate the tape number (eg, tape2 for the second tape; the files would then be called tape20001,tape20002,tape20003...). Additionally, an offset may be added (eg, offset=89 means the first files would be called tape20090,tape20091,tape20092... (1+89,2+89...)).
da> rfits mta 1-999 tape2 off=89If you tranfered your data using getpix then you may wish to preserve the naming convention given to your files. In this case, execute the following,
da> rfits mta 1-999 "" oldiraf+
In the above examples, the file list 1-999 should more than cover the number of files on tape; (hopefully) the task will end gracefully at the end of tape. When finished, rewind the tape and deallocate the drive,
da> rew mta da> dealloc mta
and remove your tape from the drive.
\fIrfits\fR fits_file = "mta" FITS data source file_list = "1-999" File list iraf_file = "tape2" IRAF filename (make_image = yes) Create an IRAF image? (long_header = no) Print FITS header cards? (short_header = yes) Print short header? (datatype = "") IRAF data type (blank = 0.) Blank value (scale = yes) Scale the data? (oldirafname = no) Use old IRAF name in place of iraf_file? (offset = 0) Tape file offset (mode = "ql")
- 3.Fixing bad pixels
Before transforming the data it might be necessary to fix the bad pixels to avoid spreading them out during the distortion corrections. We suggest that you examine some of your flat-field and bias frames using the display and the implot tasks to determine the location of the bad pixels.
cl> images im> tv tv> display flat002 (flat002 being a flat-field) tv> implot flat002
Do not exit from the task implot, but move the cursor to the imtool window and press F6. The cursor coordinates will be displayed. Now find the position of the bad pixels and return to the graph and plot that line or column using the :l #(line) and ':c #'(column) commands. You can define them more precisely using implot by overplotting lines and columns around the center of the bad region. Do this by typing o followed by :c # or :l #. Quit the task implot (type q) and create a file called badpix to define the regions you wish to fix before executing the fixpix (cl.proto) task.
tv> edit badpix (according to the following format)
The following example illustrates the format of a bad pixel file.
1 128 1 24 112 125 1489 1512 44 71 1450 1482
Each line represents a rectangular region of the image to be fixed. The regions are specified by four numbers giving the starting and ending columns followed by the starting and ending lines. The starting and ending points may be the same to specify a single column or line. Create a file called list containing all the frames to be fixed. You have to include the flats, darks, objects, arcs, etc.
tv> files *.imh > list tv> edit list (optional to add or remove files) tv> proto pr> fixpix @list badpix
Since the operation is done in place (i.e. the new image is written over the original image), it might be useful, before fixing all your images, to carry out a test with one image to make sure that everything is going well.
- 3.Creating the flat-field and dark images
- Determining scale factors
The number of counts in your flat fields is typically pretty low. In order to improve the photon statistics, flat-field frames from the whole run are usually summed. The (almost always correct) assumption is that the features in the flat frame are consistent on a night-to-night basis both in appearance and location in the images. If the flat frames are not similar, you must modify the reduction procedure appropriately.
Load the package images. Add all your darks together, summing the exposure times as well. Your blue flats, red flats and sky flats must be summed separately. Make lists of these images using the editor to ensure the proper frames are added.
cl> images im> edit darks (make a list with your darks) im> epar imsum (check the list given below) \fIExample of imsum\fR input = "@darks" Input images output = "dark" Output image (\fItitle\fR = "Sum of darks") Title for output image (\fIhparams\fR = "exptime,darktime") List of header parameters (\fIpixtype\fR = "r") Pixel datatype of output image (\fIcalctype\fR = "r") Calculation type (\fIoption\fR = "sum") Output option (low_reject = 0.) Fraction or number of low pixels to reject (high_reject = 0.) Fraction or number of high pixels to reject (verbose = no) Print log of operation? (mode = "ql") im> imsum @darks dark title="Sum of darks" im> imsum @blueflats blueff title="Sum of blue flats" im> imsum @redflats redff title="Sum of red flats" im> imsum @skyflats skyf title="Sum of sky flats"
where darks, blueflats, redflats and skyflats are files each containing a list of the images to be coadded. We suggest that you verify whether the sums were properly performed by imploting and imheading the summed images.
im> implot dark im> implot blueff im) implot redff im> implot skyf im> imhead dark,blueff,redff,skyf l+ (check exptime and darktime)
Fig. 1 shows the average of a number of columns of the dome flats and the dark.
If disk space is a problem you might delete the individual frames previously summed.
im> imdelete @darks etc...
This is the right time to determine the scale factors between the dark and the blue and red flats. These factors will be applied to create a scaled dark which will be substracted from each flat.
For the 4-M telescope the scale factor is simply the ratio between the exposure times.
For the 1-M telescope , because the dark strongly depends on dome temperature, you must determine empirically the dark contribution in each flat. If you noted the dark rate on the ratemeter during the dark and before and after the flats, then you may use these values for the scale factor (see observing manual). In this case, the scale factor is "flat rate / dark rate". If you did not record these numbers then start by ploting the blue-flat, expanding the far red end of the image and overploting the dark frame. Now determine a region in the blue-flat having only dark contribution. Fig. 2a shows the red end of the blue flat overplotted with the dark. In this example you can see that the shapes of both curves are similar beyond line 1300. Select from your data a region where both curves run parallel.
im> implot blueff :c 50 75 (plot several columns) e (expand using the cursor to define e a window at the red end) :i dark (open new image) o (overplot) :c 50 75 (overplot the same columns) q (quit)
Once you have decided the region of interest execute the imstat task both for the blue-flat and the dark.
im> imstat blueff[20:60,1300:1430] im> imstat dark[20:60,1300:1430]
The imstat task will respond with the MEAN for the specified region for both frames. (The [20:60,1300:1430] section is just an example). The ratio MEAN(blueff)/MEAN(dark) is the scale factor which must be written down for later use.
Repeat the previous steps for the red-flat selecting this time a section at the far blue end to determine its MEAN. Fig. 2b shows the blue end of the red flat overplotted with the dark. Don't forget to record the MEAN(redflat)/MEAN(dark). These determined scale factors will be applied later, after transforming your images, to create a scaled dark for each flat.
- 3.Distortion correction: Mapping the distortions
For distortion mapping, you need a multi-hole image and a comparison lamp image (in the examples below, we use the names "multi1" for the multi-hole image and "comp005" for the arc lamp image). Unless the 2D-Frutti was turned off during the run, you should be able to use the same multi-hole mapping for the entire run. This is also true for your long arc but if you wish, you may use a new one for each night.
Depending on your setup, it is sometimes difficult to get good signal to noise at the ends of the multihole image. I order to get around this inconvenience, some people obtain multiholes with two grating tilts. It becomes necessary then to coadd them. Before doing this you need to check whether the location of the features of your blue multihole match the position of the features of your red multihole. Do an implot of your blue multihole to plot a group of lines around the central line. Overplot the same lines from your red multihole and measure with the cursor the relative shift of both images (nx). If you find that there is a big offset you must use then the task imshift to shift one of your multiholes with respect to the other. Then you are ready to coadd both images with the task imsum.
Load the images, twodspec and longslit packages:
cl> images im> twods tw> long lo> implot bluemulti :l 750 770 (plot the central lines of the bluemulti) o :i redmulti :l 750 770 (overplot the same lines of the redmulti) q lo> imshift bluemulti bluemulti nx 0. lo> imsum bluemulti,redmulti multi1 title=multihole calc=r pix=r option=sum
Use implot to plot some lines near the middle of your multi-hole image to determine a typical FWHM of the profiles. Then use epar to set up the parameter files for identify and reidentify according to the lists given below. Note the parameter fwidth in the identify task should be set to TWICE THE FWHM . Then run identify:
lo> implot multi1 lo> epar ident lo> epar reid lo> ident multi1
The task will show you a plot corresponding to the sum of 50 (nsum) lines centered around the central line (see Fig. 3a). Select a profile near the center of the slit by setting the cursor on it and typing m (mark feature); enter a coordinate of 0 [arcsec]. Now mark a few more features (four or five) and give their values in arcsec according to Fig. 4. Then type l (locate) to find all of the features, which will be picked up from the file onedstds$ctio/multi1m.dat (coordlist). If everything goes well the program should show you the original plot with tick marks for all the features (see Fig. 3b). Type f to fit these points and you will be presented with a plot of the residuals (see Fig. 3c). At this level you may delete the nearest point to the cursor by typing d or change the function and its order, if not satisfied with the fit. To change the function or the order type :function legendre or :order 5 followed by f to reperform the fit. When satisfied with the residuals, return to the first plot by typing q. You can go back to the fit level again by typing f. After any modifications exit both sections by typing q. Answer yes to the question Write feature data to the database (yes)?.
Now run reidentify, using the same image as a reference image:
lo> reid multi1 multi1
The features will be traced over the entire image. This is a non-interactive task.
Now the database directory contains a file (called idmulti1) which contains a list with the positions of the features along with the identification of those features in arcseconds along the slit. The task fitcoords will do a two-dimensional surface mapping of these coordinates as a function of position in the image, i.e. S(x,y) where S is the position in arcseconds along the slit and (x,y) are the pixel values.
lo> epar fitc (check the list given below) lo> fitc multi1
By default the task presents you with a plot of the residuals as a function of x (see Fig. 5a). The problem here is how to display 3 dimensions of information (x, y and residuals) on a two-dimensional screen. Imagine a piece of paper representing your surface fit; the residuals would appear above and below the paper. Now you may view the paper from above, in which case you see (x,y) and the residuals are projected onto the plane of the paper; or you may view the paper along one edge so that the residuals appear compressed in one coordinate (x,r) or (y,r). You may select orientations with the cursor commands x and y; these will respond with a question asking which axis of the data you wish to show in x and y respectively. For example, x='x' and y='y' followed by r (redraw) will show the location of the data points but no residuals (equivalent to viewing the paper from above); x='y' and y='r' (followed by r) will show the residuals as a function of y (lines) but all information in x is lost in the projection (equivalent to viewing the paper from the side).
Start by showing x vs y: type x and specify the x-axis; type y and specify the y-axis; type r to redraw. You now see a map of the multi-hole traces (see Fig. 5b). Trim off the points at the top and bottom: move the cursors near one of the points and hit d (delete); you will be asked whether you want to delete a point (p), all points at that same x- or y-location (x or y), or the entire feature (z). To trim the top or bottom select y. Also, if you see any points that are obviously deviant, delete these points as well. Now refit (type f) and show the residuals in y: x='y', y='r' and r (see Fig. 5c). You may wish to delete other discrepant data points and/or change the order of the fit (we recommend :xord=4, :yord=6), changing the axes of the display as necessary to understand the residuals as a function of position in the image. Fig. 5d shows x vs. y after deleting some deviant points drawn as crosses. When satisfied, exit with q. Answer yes when asked 'Write coordinate map to the database (yes)'. Note: it's usually worth noting the coordinate values that are listed for the four corners of the image. These are crucial in deciding on limits for the transformed images. The output of running the fitcoor task is a file called fcmulti1 contained in the database.
Now do the same with the comparison image (consult the lists given below in order to edit the parameters for the tasks),
lo> implot comp005 (plot a column to get the FWHM of a line) lo> epar ident lo> epar reid lo> ident comp005 lo> reid comp005 comp005 lo> epar fitc lo> fitc comp005
During identify you will be presented with a plot of the sum of 10 (nsum) columns around the central column. So far, the x-axis is in pixels (see Fig. 6a). Consult the appropriate comparison spectrum to mark (type m) several prominent features evenly spread over the spectrum (6 or 7). A tick mark should appear on top of the marked features (see Fig. 6b). Then type l to automatically look for more features from the file linelists$ctiohenear.dat. The task will then locate up to 50 (maxfeatures) features that you will be required to examine. You will be presented with the spectrum of the arc, this time with the x-axis in Angstroms (see Fig. 6c). (The w and e commands might be useful to blow up the plot, as well as the + and - keys which will make the cursor jump to the marked features.) You will probably delete (type d) some of them and mark (type m) others that were not located by the task (the w command followed by a will allow you to redraw the plot at the original scale). At this level when marking a new feature the program will look for it in the coordinate list and present you with a value that you may accept (hit return) or redefine (type a new value and hit return). If the marked feature is not in the coordinate list the task will let you know this with an INDEF value. You will be required to define its value. Once you are happy with the identified features try a fit by typing f. The task will show you a plot with the residuals of the features with respect to the fit (see Fig. 6d). You may then delete (type d) or undelete (type u) points by placing the cursor close to the point of interest, or you may change the function (type :function) or the order (type :order) before attempting a new fit f. Once you are satisfied with the fit type q and you will be returned to the original spectrum of the arc. Type q again to quit the task. Answer yes to the question Write feature data to the database (yes)?.
Now execute reidentify which will reidentify the features you marked recently for the remaining columns of your image. This is a non-interactive task.
During fitcoords you are required to perform a two-dimensional fit of the arc lines thoughout the image. By default you will get a plot of the residuals as a function of the y coordinate (Fig. 7a). As explained before for the multihole you can switch the x axis of the plot to display the residuals as a function of the x coordinate and delete deviant points as well as the left and right edges that will show up as crosses (see Fig. 7b). Now, you can display the arc features seen from above as shown in Fig. 7c. Again write down the limits yielded by the task that will be used in the next section.
\fIidentify task for the multihole\fR images = "multi1" Images containing features to be identified (\fIsection\fR = "middle line") Section to apply to two dimensional images (database = "database") Database in which to record feature data (\fIcoordlist\fR = "onedstds$ctio/multi1m.dat") User coordinate list (\fInsum\fR = 50) Number of lines or columns to sum in 2D i (\fImatch\fR = 2.) Coordinate list matching limit in user un (maxfeatures = 50) Maximum number of features for automatic identi (zwidth = 100.) Zoom graph width in user units (ftype = "emission") Feature type (\fIfwidth\fR = 7.) Feature width in pixels (\fIcradius\fR = 6.) Centering radius in pixels (\fIthreshold\fR = 0.) Feature threshold for centering (minsep = 4.) Minimum pixel separation (\fIfunction\fR = "chebyshev") Coordinate function (\fIorder\fR = 4) Order of coordinate function (sample = "*") Coordinate sample regions (niterate = 1) Rejection iterations (low_reject = 3.) Lower rejection sigma (high_reject = 3.) Upper rejection sigma (grow = 0.) Rejection growing radius (autowrite = no) Automatically write to database (graphics = "stdgraph") Graphics output device (cursor = "") Graphics cursor input (mode = "ql") \fIreidentify task for the multihole\fR reference = "multi1" Reference image images = "multi1" Images to be reidentified (interactive = "no") Interactive fitting? (\fIsection\fR = "middle line") Section to apply to two dimensional images (newaps = yes) Reidentify apertures in images not in reference (override = no) Override previous solutions? (refit = yes) Refit coordinate function? (trace = yes) Trace reference image? (\fIstep\fR = 40) Step for tracing an image (\fInsum\fR = 80) Number of lines or columns to sum (shift = 0.) Shift to add to reference features (\fInlost\fR = 5) Maximum number of features which may be (\fIcradius\fR = 7.) Centering radius (\fIthreshold\fR = 0.) Feature threshold for centering (addfeatures = no) Add features from a line list? (\fIcoordlist\fR = "onedstds$ctio/multi1m.dat") User coordinate list (match = 2.) Coordinate list matching limit in user units (maxfeatures = 50) Maximum number of features for automatic identi (minsep = 4.) Minimum pixel separation\n (database = "database") Database (logfiles = "STDOUT,geom.log") List of log files (plotfile = "") Plot file for residuals (verbose = yes) Verbose output? (graphics = "stdgraph") Graphics output device (cursor = "") Graphics cursor input answer = "no" Fit dispersion function interactively? (mode = "ql") \fIfitcoord task for the multihole\fR images = "multi1" Images whose coordinates are to be fit (fitname = "") Name for coordinate fit in the database (interactive = yes) Fit coordinates interactively? (combine = no) Combine input coordinates for a single fit? (database = "database") Database (deletions = "deletions.db") Deletion list file (not used if null) (\fIfunction\fR = "chebyshev") Type of fitting function (\fIxorder\fR = 4) X order of fitting function (\fIyorder\fR = 6) Y order of fitting function (logfiles = "STDOUT,geom.log") Log files (plotfile = "plotfile") Plot log file (graphics = "stdgraph") Graphics output device (cursor = "") Graphics cursor input (mode = "ql") \fIidentify task for the comparison\fR images = "comp005" Images containing features to be identified (\fIsection\fR = "mid col") Section to apply to two dimensional images (database = "database") Database in which to record feature data (\fIcoordlist\fR = "linelists$ctiohenear.dat") User coordinate list (\fInsum\fR = 10) Number of lines or columns to sum in 2D i (\fImatch\fR = 2.) Coordinate list matching limit in user un (maxfeatures = 50) Maximum number of features for automatic identi (zwidth = 100.) Zoom graph width in user units (ftype = "emission") Feature type (\fIfwidth\fR = 5.) Feature width in pixels (\fIcradius\fR = 2.) Centering radius in pixels (\fIthreshold\fR = 0.) Feature threshold for centering (minsep = 4.) Minimum pixel separation (\fIfunction\fR = "chebyshev") Coordinate function (\fIorder\fR = 6) Order of coordinate function (sample = "*") Coordinate sample regions (niterate = 1) Rejection iterations (low_reject = 3.) Lower rejection sigma (high_reject = 3.) Upper rejection sigma (grow = 0.) Rejection growing radius (autowrite = no) Automatically write to database (graphics = "stdgraph") Graphics output device (cursor = "") Graphics cursor input (mode = "ql") \fIreidentify task for the comparison\fR reference = "comp005" Reference image images = "comp005" Images to be reidentified (interactive = "no") Interactive fitting? (section = "mid col") Section to apply to two dimensional images (newaps = yes) Reidentify apertures in images not in reference (override = no) Override previous solutions? (refit = yes) Refit coordinate function? (trace = yes) Trace reference image? (\fIstep\fR = 10) Step for tracing an image (\fInsum\fR = 20) Number of lines or columns to sum (shift = 0.) Shift to add to reference features (nlost = 10) Maximum number of features which may be lost (\fIcradius\fR = 3.) Centering radius (\fIthreshold\fR = 0.) Feature threshold for centering (addfeatures = no) Add features from a line list? (coordlist = "linelists$ctiohenear.dat") User coordinate list (match = 2.) Coordinate list matching limit in user units (maxfeatures = 50) Maximum number of features for automatic identi (minsep = 4.) Minimum pixel separation\n (database = "database") Database (logfiles = "STDOUT,hear.log") List of log files (plotfile = "") Plot file for residuals (verbose = yes) Verbose output? (graphics = "stdgraph") Graphics output device (cursor = "") Graphics cursor input answer = "no" Fit dispersion function interactively? (mode = "ql") \fIfitcoord task for the comparison\fR images = "comp005" Images whose coordinates are to be fit (fitname = "") Name for coordinate fit in the database (interactive = yes) Fit coordinates interactively? (combine = no) Combine input coordinates for a single fit? (database = "database") Database (deletions = "deletions.db") Deletion list file (not used if null) (\fIfunction\fR = "chebyshev") Type of fitting function (\fIxorder\fR = 6) X order of fitting function (\fIyorder\fR = 6) Y order of fitting function (logfiles = "STDOUT,hear.log") Log files (plotfile = "plotfile") Plot log file (graphics = "stdgraph") Graphics output device (cursor = "") Graphics cursor input (mode = "ql")
- 3.Distortion correction: Transforming images
Now, you have mappings of the dispersion and of the spatial position as functions of (x,y) in the images. You are ready to perform a geometrical transformation on the images to remove distortions.
First, you must decide on limits for the new images. For instance, starting and ending wavelengths and the number of angstroms/pixel in the output images. Edit the parameter file for the task transform and specify three of the four parameters (x1,x2,dx,nx) for each axis; leave the fourth INDEF. Starting and ending x and y coordinates may be taken from the output of the fitcoord task previously run. You will probably want to match the approximate dispersion and scale of the raw images when setting dx and dy. These values depend upon the grating and readout format used. We recommend using spline3 for the interpolation in most cases. Spline3 or poly5 best preserve resolution and noise characteristics but ring badly at sharp discontinuities. It is a good idea to plot across your transformed multi lines and arc lines to check for this. Using a linear function is strongly discouraged.
cl> noao no> twods tw> long lo> epar transform (see the list given below)
Before transforming all your images we strongly suggest that you perform a trial using your multihole and comparison frames. During the execution of this task you will be asked to specify the dispersion axis. Answer accordingly, i.e., 1 for rows or 2 for columns.
lo> trans multi1,comp005 test1,test2 multi1,comp005
Then use display to examine your transformed images hopefully they look straight.
lo> display test1 1 lo> display test2 2
Figure 8b shows a multihole before and after transformation. Figure 8a shows the same but for the comparison lamp.
If the limits set for the transformation are not correct you will get sections with useless data in either direction. Guess a new set of parameters to use for a second trial.
Do an implot of the transformed multihole and overplot different lines to check that the image is straight (see Fig. 9a). Now, do an implot of the transformed arc to see whether it is straight or not. Overplot different columns as shown in Fig. 9b.
lo> implot test1 lo> implot test2
Before executing the transform task, get rid of the test images.
lo> imdelete test1,test2
Once satisfied with the transformation create an input and output list of images, using the files command.
lo> files *.imh > inlist
The files command will create a list of all images matching the template *.imh. You may want to edit the file inlist in order to make minor changes.
lo> edit inlist
Because of limits on disk space the transformation should be done in place.
lo> trans @inlist @inlist multi1,comp005 &
In the example above, the task is submitted as a background job. Progress may be monitored by the command jobs, and by typing the logfile to see what has been done. If you run the task as a foreground job (no "&" at the end) the terminal is tied up for the duration of the task. Typical transformations take 1-2 minutes for each image!.
You also need to transform your summed flats and dark frames,
lo> trans dark,blueff,redff,skyf dark,blueff,redff,skyf multi1,comp005 &
(unless you have added these images to inlist).
\fItransformation task\fR input = "@inlist" Input images output = "@inlist" Output images fitnames = "multi1,comp005" Names of coordinate fits in the database (database = "database") Identify database (\fIinterptype\fR = "spline3") Interpolation type (\fIx1\fR = -170.) Output starting x coordinate (\fIx2\fR = 170.) Output ending x coordinate (\fIdx\fR = INDEF) Output X pixel interval (\fInx\fR = 128.) Number of output x pixels (xlog = no) Logrithmic x coordinate? (\fIy1\fR = 3550.) Output starting y coordinate (\fIy2\fR = 7050.) Output ending y coordinate (\fIdy\fR = INDEF) Output Y pixel interval (\fIny\fR = 1520.) Number of output y pixels (ylog = no) Logrithmic y coordinate? (\fIflux\fR = yes) Conserve flux per pixel? (logfiles = "STDOUT,logfile") List of log files (mode = "ql")
- 3.Creating, normalizing and correcting the flat-field
You should now smooth the summed dark frame fairly heavily (the dark should be featureless). The following table gives the smoothing values that we suggest for the flat and the dark depending on the binning used.
--------------------------------------------------- xbin ybin dither recommended smoothing size size flat dark --------------------------------------------------- 1 1 16 x 16 7 x 7 7 x 39 2 1 8 x 16 5 x 7 5 x 39 1 2 16 x 8 7 x 5 7 x 21 2 2 8 x 8 5 x 5 5 x 21 --------------------------------------------------- In the following example we assume a 2 by 2 binning. Start by smoothing the dark. cl> images im> boxcar dark smdark 5 21
Also, smooth your flats,
im> boxcar blueff smblueff 5 5 im> boxcar redff smredff 5 5
Note that the boxcar window (5 X 5 in the example) should be smaller than the dithering size.
Examine your results by overploting your original dark on the smoothed image.
im> implot smdark :c 60 (plot a central column) :i dark (open the new image) o (overplot) :c 60 (overplot same column number)
Repeat the previous steps for examination of the flats.
Now, create a scaled dark for each flat using the scale factors previously determined in section 3.4 and subtract them from your flats.
im> imarith smdark * 1.213 bluedark pix=r calc=r im> imarith smdark * 1.377 reddark pix=r calc=r im> imarith smblueff - bluedark blueflat pix=r calc=r im> imarith smredff - reddark redflat pix=r calc=r
You should examine the dark-subtracted flats to make sure they look reasonable, especially the far red end of the blueflat and the far blue end of the redflat where the signal should vanish. Fig. 10a shows the blue end of the red flat before and after dark subtraction. Fig 10b shows the same but for the blue flat.
im> implot blueflat im> implot redflat
Then sum both flats and look at them,
im> epar imsum (see the list given below) im> imsum redflat,blueflat flatsum im> implot flatsum
Fig. 10c shows an overplot of the central column of the blue flat, the red flat and the sum of both.
\fIimsum task\fR input = "redflat,blueflat" Input images output = "flatsum" Output image (title = "Sum of blue and red flats") Title for output image (hparams = "") List of header parameters (\fIpixtype\fR = "real") Pixel datatype of output image (\fIcalctype\fR = "real") Calculation type (\fIoption\fR = "sum") Output option (low_reject = 0.) Fraction or number of low pixels to reject (high_reject = 0.) Fraction or number of high pixels to reject (verbose = no) Print log of operation? (mode = "ql")Use imdelete to get rid of all the intermediate images, retaining only flatsum, smdark and skyf.
Before flat fielding your objects you need to remove the spectral signature from the quartz lamp. This operation of normalization is performed with the task response (loaded with the twodspec and longslit packages).
im> twods tw> long lo> epar response (check the list given below)
The image section flatsum[m:n,*] means that the normalizing function will be fit to the flat-field only in the range of columns m-n. Choose these to avoid including the edge of the slit.
lo> response flatsum[m:n,*] flatsum[m:n,*] response
response presents you with a plot of the averaged flat-field columns, with a function fit through it (see Fig. 11a). Select a sample range (a pair of s commands) to exclude the very ends of the frame where there may be spikes or no response (the range will be drawn at the bottom). Then, fit the new sample with an f. You may delete the previously selected sample with t and redefine it or, if the curve does not fit the gross shape of the flat-field, try different orders (type :ord 30, pausing after the ":" for the prompt). Usually an order around 30 gives a satisfactory fit. Since this is a spline fit on equal intervals, a small change in the order (i.e. number of intervals) can have a relatively large effect. A useful command is the k key which will present you with the ratio between the sample and the fit (see Fig. 11b). From this plot you can better judge the quality of the fit. You should end up with an order such that the ratio of the sample and the fit is one on the average, plus minus 2% due to the pixel-to-pixel sensitivity variations. You may type h to go back to the previous plot of the sample and the fit on top of it. When satisfied, exit with q. The output of the task is an image named response obtained from dividing each column of flatsum by the fitted function.
We suggest that you examine the output image using implot or display. It is best to check the columns to be sure it is flat.
lo> implot response lo> display response 1 zs- zr- z1=0.9 z2=1.1
Fig. 12a shows a plot of several columns of the normalized flat. If you have taken sky-flats, flatten the transformed image (skyf) using the response image just created and examine it for any remaining variations along the slit.
lo> imarith skyf / response fsky pix=r calc=r lo> implot fsky :l 250 500 (average several lines together) :l 750 1000 q
Ideally, if the illumination of the slit by the white spot was uniform, fsky should not show any variation along the slit. But in fact, when plotting several lines of fsky together, you will probably get a slope of even 20%, which may also change for different group of lines, as shown in Fig. 12b and Fig. 12c.
If you see significant variations, you should run the task illumination to fit these variations, and then modify your final flat frame.
lo> epar illumination (check the list given below) lo> illum fsky illumination
You will be presented with a plot of the sky spectrum along with 5 (nbins) bins marked at the bottom of the figure as shown in Fig. 13a. You may redefine the bins using first i to delete the original ones and then pairs of s (always starting from the left). Select the samples so that the sky lines do not fall close to the borders of the bins. Also, you should not include any saturated lines in the bins. Once you are satisfied with your samples type q. You must then fit interactively the slit profile for each bin; the task will then present you with a line averaged profile for the first bin (see Fig. 13b). Then use pairs of s to exclude from your sample any object or spikes towards the edges of the slit (with the use of t you may delete the selected sample). Once you are happy with the sample chosen fit the large scale variations along the slit (type f). We suggest that you use a spline3 function of order 3 that you may change by typing, for example, :function legendre or :order 5. When satisfied type q to quit the first bin fit. The task will then prompt you to interactively fit the next bin, and so on. The final result of performing the task is an illumination fit along the slit obtained from interpolating the 5 original fits using a polynomial of third order (parameter interp). You can now examine the output image.
lo> implot illumination
Finally, obtain your final flat.
lo> imarith response * illumination flat
If you would like more information about the flat field processing of your images, type help flatfields in IRAF.
\fIresponse task\fR calibration = "flatsum[m:n,*]" Longslit calibration images normalizatio = "flatsum[m:n,*]" Normalization spectrum images response = "response" Response function images (interactive = yes) Fit normalization spectrum interactively? (threshold = INDEF) Response threshold (sample = "*") Sample of points to use in fit (naverage = 1) Number of points in sample averaging (\fIfunction\fR = "spline3") Fitting function (\fIorder\fR = 30) Order of fitting function (low_reject = 3.) Low rejection in sigma of fit (high_reject = 3.) High rejection in sigma of fit (niterate = 1) Number of rejection iterations (grow = 0.) Rejection growing radius (graphics = "stdgraph") Graphics output device (cursor = "") Graphics cursor input (mode = "ql")
\fIillumination task\fR images = "fsky" Longslit calibration images illumination = "illumination" Illumination function images (interactive = yes) Interactive illumination fitting? (bins = "") Dispersion bins (\fInbins\fR = 5) Number of dispersion bins when bins = "" (sample = "*") Sample of points to use in fit (naverage = 1) Number of points in sample averaging (\fIfunction\fR = "spline3") Fitting function (\fIorder\fR = 3) Order of fitting function (low_reject = 3.) Low rejection in sigma of fit (high_reject = 3.) High rejection in sigma of fit (niterate = 1) Number of rejection iterations (grow = 0.) Rejection growing radius (\fIinterpolator\fR = "poly3") Interpolation type (graphics = "stdgraph") Graphics output device (cursor = "") Graphics cursor input (mode = "ql")
- 3.Applying the flat-field correction
From the previous section you have a smoothed and dark-subtracted flat-field frame. Before flat-fielding all your images we suggest that you flatten your original sky-flat (skyf) as a trial.
cl> images im> imarith skyf / flat test pix=r calc=r
The test image should be uniformly illuminated along the slit within 1 or 2%. Let's check this.
im> implot test :l 250 500 (average several lines together :l 500 750 and different samples) :l 750 1000 :l 1000 1250 q
Figure 14 shows plots of different group of lines of the flat fielded sky frame. If everything goes well delete your test image.
im> imdelete test
In the next step, flatten all your images in place. Do not flatten darks, flat fields or any of the calibration frames.
im> files obj*.imh > list im> edit list (optional to remove files) im> imarith @list / flat @list pix=r calc=r &
This operation is also time consuming (one minute per frame). You may therefore wish to perform it in the background by adding a & to the command line. The output to this point is a set of rectified and flat fielded images.
We suggest for safety, that you immediately write your transformed images to tape, as well as your calibration frames. NOTE : if you do not have a tape drive that reads at high density, use 1600 in the place of 6250. We suggest the following,
im> dataio da> alloc mta da> unlearn wfits da> wfits flat,response,... mta.6250 new+ da> wfits @list mta.6250 new- da> dealloc mta
- 3.Extracting the spectra
Go into the package apextract (noao.twodspec) and edit the parameter file for apall to look like the lists given below.
cl> noao no> twod tw> apextr ap> epar apall ( " " list given below)
In so doing you must specify in the dispaxis parameter the dispersion axis of your images (1 for lines and 2 for columns).
An important parameter in apall is format. If the format parameter is set to onedspec the output aperture extractions are one dimensional images with names formed from the output rootname and a numeric extension given by the aperture number; i.e. root.0001 for aperture 1. Note that there will be as many output images as there are apertures for each input image, all with the same output rootname but with different aperture extensions. If the format parameter is multispec the output aperture extractions are put into a two dimensional image with a name formed from the output rootname and the extension .ms. Each line in the output image corresponds to one aperture. Thus in this format there is one output image for each input image. In this example we have chosen to output the extrated spectra using the onedspec format. The reader is warned that the rest of this manual is meant only for such a format.
Also important is the parameter weights. If set to none the pixels are summed without weights except for partial pixels at the ends. Otherwise, if set to variance the extraction is weighted by the variance based on the data values and a poisson/ccd model.
Extracting spectra is generally a delicate step and we suggest doing this by carefully inspecting each image. It is worth spending some time on each image to properly select the object and the background windows, to trace the spectrum and extract it. We suggest that you start with a bright star (maybe a spectrophotometric standard) to allow a proper tracing. For instance,
ap> apall obj110
The task will start by asking you the following questions to which you must answer with CR,
Find aperture for obj110? (yes): Number of apertures to be found automatically (1): Edit apertures for obj110? (yes):The first section of the task is used to define the aperture you wish to use in extracting the spectrum. You will be presented with a cross-section of the image (average of nsum lines) along with an aperture centered around the star profile (see Fig. 15a). If there are not enough lines added together to give a good profile, use :nsum # to add up # lines. You may also wish to review or modify the profile width using :width (the width should be about twice the FWHM of the profile). Now you may redefine your aperture: put the cursor on the profile and type c (if you want to center the aperture), l (cursor marks the lower edge of the aperture) and u (the upper edge). You may also define the center, lower and upper values using :cen #, :low # and :upp #. All of these parameters are shown in the active status line at the bottom of the screen. The lower and upper edge values are given in pixels with respect to the center of the window.
Once you have defined the extraction aperture, define the background. Type b and you will get a new plot showing the same cross-section (see Fig. 15b), but with sky or background regions marked according to the sample specified in apall. If you don't like the regions defined type t to initialize the sample and then pairs of s to define new regions. You may wish to modify the function and order, although for most applications a linear fit (e.g. chebychev of order 2) is satisfactory. Type f to fit the background. When satisfied type q. This returns you to the first section. Type q again. You will be asked,
Trace apertures for obj110? (yes): Fit traced positions for obj110 interactively? (yes): wait... Fit curve to aperture 1 of obj110 interactively (yes):Answer yes to all of them (hit RETURN).
The next step is to trace the spectrum. This means fitting the center of the profile throughout the image. Ideally the center should be constant through the entire image, but usually there is a small distortion left (if the position varies widely, there may be a problem with the distortion correction) (see example in Fig. 15c). Usually, the rms of the center is about 0.2 pixels. The fitting function suggested is spline3 of third order which of course you may change in the usual way by typing, for example, :function legendre and/or :order 4 before reperforming the fit with f. When you are happy with the fit type q to quit this section and continue. You must answer yes to the questions,
Write apertures for obj110 to database (yes): Extract aperture spectra for obj110? (yes): Review extracted spectra from obj110? (yes): wait... Review extracted spectrum for aperture 1 from obj110? (yes):
Right now the extracted spectrum will be presented to you (see Fig. 15d). The result should be examined for evidence of any extraction problems. Quit this section with q. Finally, type RETURN to quit the task. If answered appriopriately, the task will output a one-dimensional image with the extracted spectrum under the name obj110.0001 (where the 0001 extension stands for aperture #1) and a file in the database under the name apobj110, which contains information about the extraction, the window sizes and the tracing.
You may now enter a good reference in apall (maybe obj110) which can be used for your next extraction. The reference must be a bright standard star whose tracing and windows will be read from the database and used as defaults in the next extraction. This is very convenient when extracting a faint object which might be difficult or imposible to trace properly because of the low signal.
ap> epar apall (fill the 'reference' parameter with the appropriate object)
Once you have entered the right reference in apall, proceed with the extraction of the rest of your objects.
ap> apall obj106 ... ...
We suggest the following scheme which summarizes the main commands in apall,
- define object window for your first object : if the object window is not properly centered on your object bring the vertical cursor to the position of your object and type s. This command shifts the center of the window to the position of the cursor. Then type c to center the window around your object. You may use the l and u commands to adjust the limits of your window. Once you are happy with the object window type b to check the background windows.
- define background windows : you will be presented with the default sky windows coming from your reference. You have the option of deleting the current windows (t key) and defining new windows (pairs of s) using the vertical cursor. Once you have defined your windows type f to fit the current sky sample. The default for the fit is a chebyshev function of order 2. You may change them by typing, for instance, :function spline3 and/or :order 3. Once you are happy with the sky fit type q to go back to the previous section.
Optionally you may define a new aperture at the position of the cursor using the n key which stands for new aperture. Go through the previous two sections. Warning : It is necessary to redefine the sky for the new aperture, and cannot skip this step. You have to repeat the previous loop for all the apertures you want to define.
If you are using a reference in apall you can skip the tracing procedure which will be taken from the reference. Proceed with the extractions by inspecting your spectra and accepting the default names (obj106.0001, obj106.0002, ...).
You can also inspect the extracted spectra by using splot in the onedspec package.
ap> onedspec on> splot obj106.0001
You will be presented with a plot of the extracted spectrum with the y-axis in counts and the x-axis in a roughly wavelength calibrated scale coming from the transformation of your images (section 3.6). This is still not a final wavelength scale.
The next step consists of extracting the wavelength calibration arc spectra. Start by editing the apall task according to the list given below. Note that in this case the parameters interactive, find, recenter, edit, trace, and review must all be set to no. Also, the parameter background must be set to none.
If you are not terribly interested in the wavelength calibration you can fill the reference parameter with any well traced object and perform the extractions of your arcs with the same reference. If you are looking for an accurate wavelength calibration do an epar for each arc extraction and fill the reference parameter with the object of interest.
on> epar apall (check the list given below)
Finally, perform the task apall for each arc. You are required to
on> apall comp105
This is a non-interactive step.
\fIapall parameters for objects\fR input = List of input images (output = "") List of output spectra (\fIformat\fR = "onedspec") Extracted spectra format (references = "") List of aperture reference images (profiles = "") List of aperture profile images (\fIinteractive\fR = yes) Run task interactively? (\fIfind\fR = yes) Find apertures? (\fIrecenter\fR = yes) Recenter apertures? (resize = no) Resize apertures? (\fIedit\fR = yes) Edit apertures? (\fItrace\fR = yes) Trace apertures? (\fIfittrace\fR = yes) Fit the traced points interactively? (\fIextract\fR = yes) Extract spectra? (extras = no) Extract sky, sigma, etc.? (\fIreview\fR = yes) Review extractions? (line = INDEF) Dispersion line (\fInsum\fR = 20) Number of dispersion lines to sum # DEFAULT APERTURE PARAMETERS (\fIdispaxis\fR = 2) Dispersion axis (1=along lines, 2=cols (lower = -5.) Lower aperture limit relative to center (upper = 5.) Upper aperture limit relative to center (apidtable = "") Aperture ID table (optional) # DEFAULT BACKGROUND PARAMETERS (\fIb_function\fR = "chebyshev") Background function (\fIb_order\fR = 2) Background function order (\fIb_sample\fR = "-30:-15,15:30") Background sample regions (b_naverage = -3) Background average or median (b_niterate = 1) Background rejection iterations (b_low_reject = 3.) Background lower rejection sigma (b_high_rejec = 3.) Background upper rejection sigma (b_grow = 0.) Background rejection growing radius # APERTURE CENTERING PARAMETERS (width = 5.) Profile centering width (radius = 10.) Profile centering radius (threshold = 0.) Detection threshold for profile centering # AUTOMATIC FINDING AND ORDERING PARAMETERS nfind = 1 Number of apertures to be found automatic (minsep = 5.) Minimum separation between spectra (maxsep = 1000.) Maximum separation between spectra (order = "increasing") Order of apertures # RECENTERING PARAMETERS (apertures = "") Select apertures (npeaks = INDEF) Select brightest peaks (shift = yes) Use average shift instead of recentering? # RESIZING PARAMETERS (llimit = INDEF) Lower aperture limit relative to center (ulimit = INDEF) Upper aperture limit relative to center (ylevel = 0.1) Fraction of peak or intensity for automatic (peak = yes) Is ylevel a fraction of the peak? (bkg = yes) Subtract background in automatic width? (r_grow = 0.) Grow limits by this factor (avglimits = no) Average limits over all apertures? # TRACING PARAMETERS (\fIt_nsum\fR = 20) Number of dispersion lines to sum (\fIt_step\fR = 20) Tracing step (t_nlost = 3) Number of consecutive times profile is lost (\fIt_function\fR = "spline3") Trace fitting function (\fIt_order\fR = 3) Trace fitting function order (t_sample = "*") Trace sample regions (t_naverage = 1) Trace average or median (t_niterate = 1) Trace rejection iterations (t_low_reject = 3.) Trace lower rejection sigma (t_high_rejec = 3.) Trace upper rejection sigma (t_grow = 0.) Trace rejection growing radius # EXTRACTION PARAMETERS (\fIbackground\fR = "fit") Background to subtract (skybox = 1) Box car smoothing length for sky (\fIweights\fR = "none") Extraction weights (none|variance) (pfit = "fit1d") Profile fitting type (fit1d|fit2d) (clean = no) Detect and replace bad pixels? (saturation = INDEF) Saturation level (readnoise = "0.") Read out noise sigma (photons) (gain = "1.") Photon gain (photons/data number) (lsigma = 4.) Lower rejection threshold (usigma = 4.) Upper rejection threshold (nsubaps = 1) Number of subapertures per aperture (mode = "ql")
\fIapall parameters for arcs\fR input = List of input images (output = "") List of output spectra (\fIformat\fR = "onedspec") Extracted spectra format (references = "") List of aperture reference images (profiles = "") List of aperture profile images (\fIinteractive\fR = no) Run task interactively? (\fIfind\fR = no) Find apertures? (\fIrecenter\fR = no) Recenter apertures? (resize = no) Resize apertures? (\fIedit\fR = no) Edit apertures? (\fItrace\fR = no) Trace apertures? (fittrace = no) Fit the traced points interactively? (\fIextract\fR = yes) Extract spectra? (extras = no) Extract sky, sigma, etc.? (\fIreview\fR = no) Review extractions? (line = INDEF) Dispersion line (nsum = 20) Number of dispersion lines to sum # DEFAULT APERTURE PARAMETERS (\fIdispaxis\fR = 2) Dispersion axis (1=along lines, 2=cols (lower = -5.) Lower aperture limit relative to center (upper = 5.) Upper aperture limit relative to center (apidtable = "") Aperture ID table (optional) # DEFAULT BACKGROUND PARAMETERS (b_function = "chebyshev") Background function (b_order = 2) Background function order (b_sample = "-30:-15,15:30") Background sample regions (b_naverage = -3) Background average or median (b_niterate = 1) Background rejection iterations (b_low_reject = 3.) Background lower rejection sigma (b_high_rejec = 3.) Background upper rejection sigma (b_grow = 0.) Background rejection growing radius # APERTURE CENTERING PARAMETERS (width = 5.) Profile centering width (radius = 10.) Profile centering radius (threshold = 0.) Detection threshold for profile centering # AUTOMATIC FINDING AND ORDERING PARAMETERS nfind = 1 Number of apertures to be found automatic (minsep = 5.) Minimum separation between spectra (maxsep = 1000.) Maximum separation between spectra (order = "increasing") Order of apertures # RECENTERING PARAMETERS (apertures = "") Select apertures (npeaks = INDEF) Select brightest peaks (shift = yes) Use average shift instead of recentering? # RESIZING PARAMETERS (llimit = INDEF) Lower aperture limit relative to center (ulimit = INDEF) Upper aperture limit relative to center (ylevel = 0.1) Fraction of peak or intensity for automatic (peak = yes) Is ylevel a fraction of the peak? (bkg = yes) Subtract background in automatic width? (r_grow = 0.) Grow limits by this factor (avglimits = no) Average limits over all apertures? # TRACING PARAMETERS (t_nsum = 20) Number of dispersion lines to sum (t_step = 20) Tracing step (t_nlost = 3) Number of consecutive times profile is lost (t_function = "spline3") Trace fitting function (t_order = 3) Trace fitting function order (t_sample = "*") Trace sample regions (t_naverage = 1) Trace average or median (t_niterate = 1) Trace rejection iterations (t_low_reject = 3.) Trace lower rejection sigma (t_high_rejec = 3.) Trace upper rejection sigma (t_grow = 0.) Trace rejection growing radius # EXTRACTION PARAMETERS (\fIbackground\fR = "none") Background to subtract (skybox = 1) Box car smoothing length for sky (\fIweights\fR = "none") Extraction weights (none|variance) (pfit = "fit1d") Profile fitting type (fit1d|fit2d) (clean = no) Detect and replace bad pixels? (saturation = INDEF) Saturation level (readnoise = "0.") Read out noise sigma (photons) (gain = "1.") Photon gain (photons/data number) (lsigma = 4.) Lower rejection threshold (usigma = 4.) Upper rejection threshold (nsubaps = 1) Number of subapertures per aperture (mode = "ql")
- 3.Wavelength calibration
Wavelength calibration consists of three steps. The first one is done interactively with the task identify which allows you to get from an arc a dispersion solution (wavelength versus pixel). That solution is the output from identify and is stored as a text file in the database. The next step consists in assigning arc references to your objects. This step is performed with the task refspectra which simply writes this assignment in the header of your objects under the keywords REFSPEC1 and REFSPEC2. The third step consists in applying the wavelength solution to the extracted spectra which is performed with the task dispcor. This task is non-interactive.
STEP 1 .
You have to start by identifying emission lines in your arc spectra to be used to get a wavelength solution which will be applied later to your objects. Do an epar on the identify task (in the onedspec package) and set the parameters according to the list given below. The first time you run this task is the most time consuming. Once you have obtained a solution for one arc, the following identifications will be much more straightforward. Now, execute the task for your first arc.
cl> onedspec on> epar identify (check the list given below) on> identify comp105.0001 \fIidentify parameters\fR images = "comp105.0001" Images containing features to be identified (section = "middle line") Section to apply to two dimensional images (database = "database") Database in which to record feature data (\fIcoordlist\fR = "linelists$idhenear.dat") User coordinate list (nsum = 10) Number of lines or columns to sum in 2D (\fImatch\fR = 2.) Coordinate list matching limit in user (maxfeatures = 50) Maximum number of features for automatic (zwidth = 100.) Zoom graph width in user units (ftype = "emission") Feature type (\fIfwidth\fR = 5.) Feature width in pixels (\fIcradius\fR = 2.) Centering radius in pixels (\fIthreshold\fR = 50.) Feature threshold for centering (minsep = 4.) Minimum pixel separation (\fIfunction\fR = "chebyshev") Coordinate function (\fIorder\fR = 3) Order of coordinate function (sample = "*") Coordinate sample regions (niterate = 1) Rejection iterations (low_reject = 3.) Lower rejection sigma (high_reject = 3.) Upper rejection sigma (grow = 0.) Rejection growing radius (autowrite = no) Automatically write to database (graphics = "stdgraph") Graphics output device (cursor = "") Graphics cursor input (mode = "ql")Several arc spectra with identified lines can be found in the computer room. Select the appropriate figure for your needs.
The task will present you with the extracted arc (see Fig. 16a). So far the x-axis is on a rough wavelength scale which helps in the identification of the lines in your spectrum. Once you have recognized some features (this step is not straightforward and make take a while ...), select a line by setting the cursor on it, typing m (mark feature) and entering the corresponding value in Angstroms. A tick mark should appear on top of the feature. Now mark a few more features (3 or 4) and enter their values by consulting the reference figure.
Try now to perform a fit with the current points. If you type f the task will carry out a one-dimensional fit between the values you entered in Angstroms and their pixel values. You will be presented with a plot of the residuals of the fit as a function of wavelength (see Fig. 16b). The rms of the fit in Angstroms is given at the top of the plot. So far the residuals are probably quite high and you will probably see some systematic trends in the residuals. This is not important for the time being because you need to have a lot of features before attempting to get a decent solution. Type q to quit this section and you will see the arc spectrum in a real wavelength scale! , which makes things a lot simpler to identify features (see Fig. 16c). You must continue marking more features. From now on, each time you mark a feature you will be offered a value in Angstroms that generally is the right one. If so, hit RETURN to accept it. If not, enter the correct value by hand. Try to mark several features evenly separated throughout the dispersion range. Type f to improve the current fit. If you see any trend in the residuals try to increase the order of the function by typing, for instance, :order 3 followed by f to reperform the fit. In order to change the type of the function to fit type, for instance, :function legendre. You may delete points with high residuals by moving the cursor near to the point and typing d or undelete points (crosses) with u. Once you are happy with the type of function along with the order of the fit, type q to return to the arc spectrum.
This is the time to let the computer work and automatically search more lines from the internal library. The locate command checks for all the emission lines in your arc whether there is a nearby value to a given feature in the internal library or not. Before attempting this it is necessary to define an appropriate range in Angstroms to be scanned in the internal library around each line found in your arc. This range is set with the parameter match which must be set to 2-3 times the rms of the current fit.
Once you have the current rms, type :match # followed by l which stands for locate. This will automatically mark more features in your arc. After a few seconds you will probably notice several new lines marked in the arc spectrum. Type f to try a new fit. You will see a plot of the residuals of the fit (see Fig. 16d). Play around with the commands you know to improve the fit. Don't be surprised if you need a 6th order function. Generally you are doing well with an rms of 0.1-0.2 pixels (since the rms is given in Angstroms you have to divide that number by the number of Angstroms per pixel in your spectra). Once you are happy with it type q to quit this section, and finally q again to quit the task. You must answer yes to the question,
Write feature data to the database (yes)?The solution previously obtained will be stored in the database as a text file.
Proceed now with the identification of the other arcs. Since you already have a wavelength solution you can use it as a reference for the next identifications. Start by running identify on the next arc.
on> ident comp107.0001Once you have the arc spectrum presented in front of you, read from the database a wavelength solution. The format of the command is :read name ap, where name is the name of the reference spectrum, and ap is the aperture number. Type, for instance, :read comp105.0001 1. You will be presented with the current spectrum with the x-axis in Angstroms along with the features identified in comp105.0001 plotted on top. Although you probably cannot see it, the tick marks are slightly shifted with respect to the emission lines. The next step consists of recentering the old features on the current ones. Type a (all) followed by c (center) to center all the features. After thinking a while the task will shift the tick marks accordingly. You may then do a fit (f) to inspect the residuals. At this level you may use all the commands described above to get a good fit. Quit the task (q) once you have a good fit with the right rms.
Once you have identified all your spectra the wavelength solutions should all be in the database.
STEP 2 .
There are many ways of assigning arc references to your objects. We suggest two different approaches that satisfy most needs of IRAF users.
- If you have only one arc for the whole night do an epar on refspec according to the list given below and fill the reference parameter with the arc you wish to use (for instance, comp105.0001) and execute refspec.
on> epar refspec (check the list given below) on> refspec obj*.000?.imh \fIrefspec parameters\fR input = "obj*.000?.imh" List of input spectra (\fIreferences\fR = "comp105.0001") List of reference spectra (apertures = "") Input aperture selection list (refaps = "") Reference aperture selection list (ignoreaps = yes) Ignore input and reference apertures? (\fIselect\fR = "average") Selection method for reference spectra (sort = "none") Sort key (group = "none") Group key (time = no) Is sort key a time? (timewrap = 17.) Time wrap point for time sorting (override = yes) Override previous assignments? (confirm = yes) Confirm reference spectrum assignments? (assign = yes) Assign the reference spectra to the input (logfiles = "STDOUT,logfile") List of logfiles (verbose = no) Verbose log output? answer = "yes" Accept assignment? (mode = "ql") You should get the following question for each object, [obj106.0001] refspec1='comp105.0001' Accept assingment?(no|yes|YES)(yes):to which you have to answer yes in order to accept the assignment.
- If you have one or two different arc(s) per object prepare a table with your specific assignments in the following format,
on> edit ref.table obj020.0001 comp019.0001 obj022.0001 comp021.0001,comp023.0001 obj022.0002 comp021.0002,comp023.0002 obj023.0001 comp019.0001 ...In the previous example obj020.0001 will be wavelength calibrated using only the arc comp019.0001. obj022.0001 and obj022.0002 will be calibrated using simultaneously two arcs bracketing your object, both with the same weight. Note that each spectrum from obj022 has its own set of extracted arcs. obj023.0001 will be calibrated with the same arc used for obj020.0001.
Now, do an epar for refspec and fill the reference parameter with ref.table (the name of the assignment table you created), fill the select parameter with average and execute the task.
on> epar refspec on> refspec obj*.000?.imhSTEP 3 .
Before running dispcor you need a list with your raw spectra and a list with the names of the output spectra. Use the files command to create the input list. We suggest the following,
on> files obj*.000?.imh > inlist on> edit inlist (optional)Once you have the input list with the appropriate files make a copy of it into a new list called dclist and edit this new file to change the output names of your wavelength calibrated spectra.
on> copy inlist dclist on> edit dclist (if you have followed the convention of this manual the following vi command might be useful for a global replacement :1,$s/obj/dc/)
Once you have created the output list which must have the same number of lines as the output list , you are ready to run dispcor. Do an epar first on dispcor according to the list given below.
on> epar dispcor (check the list given below) \fIdispcor parameters\fR input = "@objlist" List of input spectra output = "@dclist" List of output spectra (\fIlinearize\fR = yes) Linearize (interpolate) spectra? (database = "database") Dispersion solution database (table = "") Wavelength table for apertures (\fIw1\fR = 3500.) Starting wavelength (\fIw2\fR = INDEF) Ending wavelength (\fIdw\fR = 2.4) Wavelength interval per pixel (\fInw\fR = 1520) Number of output pixels (\fIlog\fR = no) Logarithmic wavelength scale? (\fIflux\fR = yes) Conserve flux? (samedisp = no) Same dispersion in all apertures? (global = no) Apply global defaults? (ignoreaps = yes) Ignore apertures? (confirm = no) Confirm dispersion coordinates? (listonly = no) List the dispersion coordinates only? (verbose = yes) Print linear dispersion assignments? (logfile = "") Log file (mode = "ql")Note that the parameter linearize is set to yes. This means that the spectra will be interpolated to a linear wavelength scale and the dispersion coordinate system in the header is set appropriately. If no, the nonlinear dispersion function(s) from the database are assigned to the input image header and the spectral data are not interpolated.
You must decide upon the starting wavelength, the number of Angstroms per pixel and the length of the output spectra in order to have all of them in the same format. Execute then the task.
on> dispcor @inlist @dclist
The output of dispcor is a set of wavelength calibrated images called in this example dc020.0001, dc022.0001, dc022.0002, dc023.0001 ... It is a good idea to plot them with the use of the splot task.
on> splot dc020.0001 Check that the x-scale is in Angstroms.
- 3.Neutral Density filter corrections
If you used different neutral density filters during your run and wish to flux-calibrate your data you must correct for the transmission of these filters. Filter transmission curves are kept in the onedstds$ctio/ directory as tables.
You may apply a two-dimensional correction to the full image before extraction or prepare suitable one-dimensional curves for correction after extraction. Of course, the effect of any filter which was used for the entire night will be removed by the flux-calibration, since it may be treated as part of the overall system sensitivity.
In the onedspec package edit the parameters for the task ndprep, taking w0, dw, nw from the wavelength calibrated spectra, using nspace=0 and dispaxis=2, and then run the task.
cl> noao no> onedspec on> images im> epar ndprep (check the figure) im> ndprep nd1m.100mag.dat nd1 im> ndprep ...
for each ND filter you need. Then, for each image which needs a correction use imarith to divide by the ND image, in place:
im> imarith dc020.0001 / nd1 dc020.0001 pix=r calc=r im> ...
\fIndprep parameters\fR filter_curve = "nd1m.100mag.dat" Input ND filter curve output = "nd1" Output calibration image (\fIw0\fR = 3500.) Starting wavelength (Angstroms) (\fIdw\fR = 2.4) Wavelength increment (Angstroms) (\fInw\fR = 1520) Number of wavelength points (nspace = 0) Number of spatial points (0 for 1D) (\fIlogarithm\fR = no) Use logarithmic wavelengths? (flux = yes) Conserve flux when log rebinning? (dispaxis = 2) Dispersion axis (\fIdirectory\fR = "onedstds$ctio/") ND filter directory (mode = "ql")
- 3.Deriving a sensitivity function
The sensitivity function is determined in two steps: first, the task standard, calculates the individual sensitivity measurements from the extracted flux standard spectra; second, sensfunc, collectively fits the individual measurements. Notice that the function is derived from the extracted spectra before applying any extinction correction (this is because the correction is applied in the sensfunc task).
Both standard and sensfunc are found in the onedspec package. Load this package and edit the parameter files of standard and sensfunc according to the lists given below.
cl> noao no> onedspec on> epar stand (check the list given below) on> epar sensfunc (check the list given below) \fIstandard parameters\fR input = "dc110.0001" Input image file root name output = "std1" Output flux file (used by SENSFUNC) (samestar = yes) Same star in all apertures? (beam_switch = no) Beam switch spectra? (apertures = "") Aperture selection list (bandwidth = INDEF) Bandpass widths (bandsep = INDEF) Bandpass separation (\fIfnuzero\fR = 3.6640000000000E-20) Absolute flux zero point (\fIextinction\fR = "onedstds$ctioextinct.dat") Extinction file (\fIcaldir\fR = "onedstds$ctionewcal/") Directory containing calibration (\fIobservatory\fR = "ctio") Observatory for data (interact = yes) Graphic interaction to define new bandpasses (graphics = "stdgraph") Graphics output device (cursor = "") Graphics cursor input star_name = "eg274" Star name in calibration list answer = "yes" (no|yes|NO|YES|NO!|YES!) (mode = "ql") \fIsensfunction parameters\fR standards = "std1" Input standard star data file (from STANDARD sensitivity = "sens1" Output root sensitivity function imagename (apertures = "") Aperture selection list (ignoreaps = no) Ignore apertures and make one sensitivity (logfile = "senslog") Output log for statistics information (\fIextinction\fR = "onedstds$ctioextinct.dat") Extinction file (newextinctio = "extinct.dat") Output revised extinction file (\fIobservatory\fR = "ctio") Observatory of data (function = "spline3") Fitting function (order = 6) Order of fit (interactive = yes) Determine sensitivity function interact (graphs = "sr") Graphs per frame (marks = "plus cross box") Data mark types (cursor = "") Graphics cursor input (device = "stdgraph") Graphics output device answer = "yes" (no|yes|NO|YES) (mode = "ql")The task standard will examine the extracted spectrum of a standard star, bin the counts within the calibrated bandpasses and compare them to the tabulated magnitudes for that star. The choices for the star directory, caldir are given in appendix A of this manual. The ratio of flux per unit wavelength to counts gives a system sensitivity in each bandpass. The sensitivities are stored in the file specified by the parameter output, which should probably be different for each night (here we give the example of std1). Run standard for each star of night 1:
on> stand dc110.0001 std1 eg274 on> ...
Be sure to specify the correct star name for each spectrum. Standard will ask you the following,
[dc110.0001]: Edit bandpasses? (yes): to which you must answer 'yes'.
You will then be presented with the spectrum of the standard star on top of which have been drawn boxes at the positions of the flux points taken from the internal library (see Fig. 17). You have the option of deleting the nearest point to the cursor by typing d, or adding a new flux point at the position of the cursor. You need to specify with the vertical cursor both the lower and upper limits of the box by typing a twice. The flux point for this new bandpass will be linearly interpolated from the two nearest points found in the library. Type q to quit the task.
Now, run sensfunc, which combines the sensitivities calculated above from the different stars and produces a fitted sensitivity function.
on> sensfunc std1 sens1
std1 is the output from standard and sens1 is the root name for the output of sensfunc. Answer yes when asked to perform an interactive fit.
You will be presented with the two graphs which were specified in the parameter graphs (see Fig. 18a). These represent the s ensitivity vs wavelength, s, and the r esidual sensitivity vs wavelength, r. Up to four graphs may be displayed at a time. The other possibilities are:
a - Residual sensitivity vs. \fIa\fRirmass c - \fIc\fRomposite residuals and error bars vs. wavelength e - \fIe\fRxtinction (and revised extinction) vs. wavelength i - Flux calibrated \fIi\fRmage vs. wavelength
The graphs may be changed at any time during the fitting using ':graphs [types]', and specifying the types desired. Now looking at the sensitivity vs. wavelength, simply delete d any data points which are clearly deviant. After typing d, you will be prompted for the type of data to be deleted, a point p, star s, or wavelength w nearest the cursor. If you delete the wrong point, type u and respond accordingly to undelete the point. Information about the point nearest the cursor may be obtained by typing i. The weight of the point nearest the cursor may be changed by typing w. Typing r will update the graphs if any changes are made. If all the points for a given star are off, conditions probably were not photometric, resulting in a light-loss. These are "grey" (color-independent) factors and may be easily corrected by typing s, which applies a shift factor to all the stars to bring them to the level of the "best" conditions. This is a toggle so typing s again will "unshift" the data. You may also enter fake data points a with the cursor positioned at the false point, then type g to refit. Unfortunately you have no way of telling where these fake points should go. You must be aware of problems with extrapolation beyond the last data points (and also interpolation between the last few points) since the functions here are not well constrained. If you change the function type :func [type] or order :order #, type f to overplot the new fit or g to fit and redraw the graph. You may look at the composite points being fitted by typing c, this is a toggle. If you have changed the data too much, or don't like the changes you've made, you may return to the original data by typing o.
In order to see if you guessed right, you may apply the extinction correction and the sensitivity curve to the extracted spectra of the standard stars, and compare them with the tabulated flux values. You can select the standard by typing :image dc110.0001 followed by :g i (see Fig. 18b). To go back to the sensitivity curve, type :g sr. Repeat this until the sensitivity function is adequate for your needs. To quit type q and the sensitivity curve will be saved as sens1.0001.
- 3.Applying the sensitivity function
Once you have the sensitivity curve, you may extinction-correct and flux-calibrate all the appropriate data.
If you are planning to extinction-correct your data, the proper airmass information must be present in the image headers. This means either the airmass or the sidereal time and object coordinates which along with observatory latitude can be used to work out the airmass. To check you may use the task hselect which will output a list of the airmass values, or any other header parameters you wish to check.
cl> images im> hselect dc*imh $I,airmass,exptime,title yes
The "$I" gives the filename while the title and airmass are taken directly from the header. Make note of the incorrect airmass values and edit the header. If this information is not present, you must enter it. Use,
im> hedit dc110.0001 airmass 1.06 add+ im> hedit dc106.0001 airmass 1.21 add+ im> ...
If your exposures are long or if you observed at a large airmass, the effective airmass may differ significantly from the one recorded in the header, which corresponds to the value at the beginning of your exposure. You may replace the airmass of the header with the effective value by using the setairmass task in the astutil package.
im> astutil as> setairmass dc*imh observatory=ctioThe task will present you with the current airmass taken from the header and the new value calculated using the exposure time of your exposure, the sidereal time, the hour angle and the declination of your frame.
Prepare an input list of all the wavelength calibrated spectra which are to be extinction and flux corrected and also make an output list with the names of the output spectra. Do an epar on the task calibrate and execute it accordingly.
as> onedspec on> delete dclist on> files dc*imh > dclist on> edit dclist (optional) on> copy dclist fclist on> edit fclist (you may use the following vi command to apply a global replacement: 1,$s/dc/fc/) on> epar calibrate (check the list given below) on> calibrate @dclist @fclist \fIcalibrate parameters\fR input = "@dclist" Input spectra to calibrate output = "@fclist" Output calibrated spectra (\fIextinct\fR = yes) Apply extinction correction? (\fIflux\fR = yes) Apply flux calibration? (\fIextinction\fR = "onedstds$ctioextinct.dat") Extinction file (\fIobservatory\fR = "ctio") Observatory of observation (ignoreaps = yes) Ignore aperture numbers in flux calibration? (\fIsensitivity\fR = "sens1.0001") Image root name for sensitivity spectra (fnu = no) Create spectra having units of FNU? (mode = "ql")
To this point, you have a set of one-dimensional images which are completely calibrated in wavelength and flux.
The final spectra may be examined and analyzed using the task splot (see examples in Fig. 19). Type help splot for more information about this task.
on> splot fc110.0001 on> ... Type 'q' to quit from 'splot'.
Writing FITS Tapes
Mount your output tape. Don't forget the write ring and select the desired tape density. Load dataio and allocate the drive. Then use wfits to write the images. You may wish to select scaling or no-scaling, force the number of bits/pixel, etc.
cl> dataio da> alloc mta da> wfits *.imh mta.6250 new+
new+ specifies a new tape (to write from the beginning) and new- goes to the logical end-of-tape before writing.
After writing all desired images, deallocate the drive and dismount your tape.
da> dealloc mta