Calibration of BP and RP spectra

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Calibration of BP and RP spectra

  1. 1. f t Calibration of BP and RP spectra ra D prepared by: C. Fabricius, C. Jordi, H. Voss, F. Figueras, J.M. Carrasco affiliation : Institute for Space Studies of Catalonia and University of Barcelona approved by: reference: GAIA-C5-TN-UB-CF-006-01 issue: D revision: 0 date: 25 May 2007 status: Draft Abstract We discuss the calibration of the BP/RP spectra, especially establishing an internally consistent flux and wavelength scale.
  2. 2. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 1 Introduction We present a model for the calibration of the BP and RP spectra, without defining the detailed calibration parameters, but with emphasis on the general principles for the calibration, and the interaction with the other photometric processes. Details of the individual processes must be defined in later documents. A study of the processing of the spectra was presented by Brown in 2006, also discussing many details on specific calibration and instrument model issues. The goal of the calibration is to allow the best possible reconstruction of the source spectra from the available set of observations. We enter here into the questions of the wavelength scale, and an internally consistent flux scale, but will not discuss the absolute flux calibration. The calibration processes are closely linked f t • to the photometric model describing how the reconstructed spectra are represented; ra • to the instrument model describing the characteristics of the processes that lead to the observed spectra; • and to the source updating, i.e. the reconstruction of the spectra using the available observations, the calibration parameters, and the instrument model. D The calibration model assumes that the true observations are compared with the predicted ob- servations derived from the source data and the current calibration parameters and that the instrumental variations can be described as small corrections to the current set of calibration parameters. The instrument model must be sufficiently detailed to allow observed spectra to be predicted from a known spectrum of the same source. This prediction step is essential for Gaia where a considerable degradation of the detectors with time must be expected, mainly due to radiation damage, and spectra observed early in the mission may therefore look quite differ- ent from spectra observed towards the end. The prediction capability is also indispensable for disentangling spectra of double stars. 1.1 The iterative photometric solution There are several quantities to be calibrated and they depend on each other in a complicated manner. The problem has some similarity to AGIS, where astrometric source updating, attitude updating, and calibration updating, run in a global iterative solution, but with a rather slow con- vergence (Lindegren 2007). An obvious, but not necessarily optimal, approach for photometry is to solve for each element of the calibration in a similar iterative process, alternating with a re- newed reconstruction of the spectra for the calibration sources, as shown in Figure 1. Numerical experiments are needed to establish the best strategy and a realistic level of ambition. Institute for Space Studies of Catalonia Technical Note and University of Barcelona 2
  3. 3. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 In the calibration we must separate variations in dispersion, which distributes the same signal of a different range of samples, from sensitivity variations which modifies the size of the signal, perhaps even in a colour dependent way. The sensitivity variations, on their side, must be separated from the AC cutoff, which contrary to the sensitivity depends on the position within the window. The dispersion must also be separated from LSF variations, which are mainly seen in the wings. f t ra Figure 1: Schematic view of the internal calibration of BP and RP spectra, alternating between D reconstruction of spectra for the calibration stars and calibrating the properties of the CCDs against one another We note that some calibration sources may be suitable for one type of calibration but not neces- sarily for all, so the different calibration processes might use different sets of standard sources. While the calibration iterations are running, only the calibration sources need to be updated, but it must always be reconsidered if they are really suited as calibrators, or if the improvements in the calibration have demonstrated that some of them had better been eliminated. Once the iter- ations have finished, however, all sources must be processed and it is then possible to nominate additional calibration sources. A photometric source updating must run before any calibration can be made. The very first time it runs, it must use nominal values for dispersion, wavelength scale, etc. The tables used on board for propagating a detection to BP/RP will serve as a crude, initial geometrical calibration. During the source updating, a weighting, or even rejection, of individual transits can be estab- lished, and this weighting should be maintained for the calibration. This is again similar to the system used in AGIS (Lindegren 2005), and saves tedious internal iterations for re-establishing the weights. Institute for Space Studies of Catalonia Technical Note and University of Barcelona 3
  4. 4. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 An initial calibration must be obtained during the commissioning phase, involving a set of observations mainly of stars around the ecliptic poles. 1.2 Implementation of the calibration The present paper is focused on the calibration models defined by DU12, and in principle not with the details of their implementation. At the same time there are strong dependencies with the work of DU11 for initial processing of the spectra, DU15 for issues related to the imple- mentation, and DU16 for selection of internal reference sources. We therefore occasionally discuss questions related to the interfaces with DU11, DU15, and DU16. Also, we see it as a responsibility of DU12 to argue that a proposal it presents is feasible, and therefore to indicate, in very general terms, how it may be implemented. 2 Boundary conditions for the calibration f t ra We will in this section discuss how we see the parts of the source and instrument model that have direct relations to the calibration. 2.1 Representation of the spectra D We propose that the reconstructed mean spectra for a source are expressed as fluxes in units of e− /s for two sets of wavelength intervals covering the range of BP and RP respectively. The wavelengths for BP: λB0 , λB1 , . . . , λBnB ; and for RP: λR0 , λR1 , . . . , λRnR , are chosen once and for all, such that the width of each band corresponds to a suitable fraction of a pixel at an average dispersion. This fraction is TBD and may in principle vary along the spectrum, but we could think of something between 6 and 1 . In summary: 1 4 • the photometric solution will give the internally calibrated flux, g(λ i ), for a large number (hundred or more?) of bands, [λi−1 ; λi ]; • there are nB bands for BP, defined by nB + 1 wavelengths; • there are nR bands for RP, defined by nR + 1 wavelengths; • the fluxes are measured in e− /s; • the fluxes from BP are treated independently from fluxes derived from RP; • the bands are on an established (fixed) wavelength scale, and need not have the same width; Institute for Space Studies of Catalonia Technical Note and University of Barcelona 4
  5. 5. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 • in the overlap region between BP and RP, the resolution is highest in RP, and the BP wavelengths should here be selected to coincide with some of the RP wavelengths. We shall not enter into the discussion of the selection of the spectral bands, but to have an idea of the data volume, we note that the observed spectra have some 50 pixels in BP as well as in RP, so a total of several hundreds bands must be considered. Each band is represented by a flux, an error estimate for the flux, and at least one correlation coefficient for the correlation with the neighbour band. We cannot expect to separate neighbour bands completely, and the wavelength calibration will also be imperfect, so there will be a considerable correlation between neighbour bands, perhaps extending to the next neighbours as well. Faint stars do not necessarily deserve full spectral resolution and the question of whether we should define high and low resolution spectra may be relevant. For the very faintest stars we may even have to give only the integrated BP and RP fluxes, and sometimes just one of them. f 2.2 The photometric source updating t ra The source updating must aim at reconstructing the best possible spectra, e.g. using an ML method. The source updating must process all sources observed, we imagine one source at a time, deriving mean BP and RP spectra according to the photometric source model. It must also provide error estimates and correlations. The source updating must identify disturbed transits and parasites, and down-weight discordant observations. Epoch photometry in terms of integrated BP and RP bands must be checked for signs of variability of the colour. D 2.3 The instrument model A realistic instrument model is a prerequisite for a successful calibration. We have identified the following elements, that the model must at least address: • gain, bias, RON, bad pixels/columns, dark current, sky background; • CTI; • AC geometry on large scale; • inclination of the spectra; • an AC LSF for each band, including the effect of pixelization; • AC motion of the image; • AC signal cutoff; • AL geometry on large and small scale; Institute for Space Studies of Catalonia Technical Note and University of Barcelona 5
  6. 6. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 • dispersion curves; • absolute wavelengths; • an AL LSF for each band, including the effects of TDI, pixelization, attitude jitter; • sensitivity on large and small scale; • saturation, and deviations from linear response. Our purpose here is to establish internally consistent spectra, and we therefore do not consider the absolute response, and we do not break the response down into mirror reflectivity, filter and prism transmission, and CCD QE. Only the combined response is relevant for us, and in fact only the variation of the combined response is important and within reach. f t We have included a wavelength calibration, though it goes somewhat beyond the concept of internal calibration, because of its importance for chromaticity studies and for classifications, and because it has been assigned to DU12. ra The instrument model must be sufficiently detailed to allow a prediction of the observed spectra for each transit from the reconstructed spectra, the source astrometry, and attitude and orbit data. For the internal calibrations we need the source data on an internally consistent scale, essentially D defined by the normalisation constraints of the various calibration steps. Eventually an external calibration must be added that allows a prediction based on external standard spectra. The external calibration is assigned to DU14, and is beyond the scope of the present paper. 2.4 Prediction of spectra For the detailed calibrations we need the ability to predict an observed spectrum, or in fact a set of possible observations with slightly different AL positions. The simplest model is: ηpred = η0 + ∆η + Γ(λ) (1) where, ηpred is the predicted field angle, η0 is the nominal field angle for the gate, ∆η is the AL geometrical correction, and Γ(λ) is the dispersion law; Institute for Space Studies of Catalonia Technical Note and University of Barcelona 6
  7. 7. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 and in the across scan direction: ζpred = ζ0 (µ) + ∆ζ (2) where, ζpred is the predicted field angle, ζ0 is the nominal field angle for the gate, µ is the column, and ∆ζ(µ) is the AC geometrical correction. This model is similar to the model used in astrometry. We need the astrometric source data to know where the source is in the sky; the attitude to know where the telescopes are pointing; and orbit data to know the effects of parallax and aberration. The predicted field angle is calculated from these pieces of information at a given time, e.g. the observation time of a central sample in t the spectrum. The nominal field angle is a fixed quantity that only depends on the CCD in ques- tion and the gate (Fabricius & Masana 2007). The AL geometrical correction in combination f with the dispersion law, then defines the centroid in sample space for each wavelength through a simple transformation. By convention, a certain (central) wavelength, λ 0 , in the band can be ra chosen to have Γ(λ0 ) = 0. Note that the observation time for a sample depends on the selected gate (Fabricius & Masana 2007). The flux for each sample can now be integrated over the wavelengths, using the PSFs. Depend- ing on the calibration at hand, other sets of predicted observations, e.g. with slightly different dispersions can be calculated in parallel. D The prediction process must also deal with sensitivity, AC flux loss, CTI, etc. More details on the prediction process should be added to this section 3 Initial calibration steps 3.1 Gain, bias, dark current, RON, bad columns, sky background We list here, for completeness, a number of basic calibrations that belong to the pre-processing (DU11) work packages. For the CCD circuits we must calibrate the zero point (bias) and scale (gain) to be able to convert the signal to electrons. The calibration is based on pre-scan data from the ASD telemetry files, and statistical studies on a small data set. The pre-scan data are also used for studying read-out noise and dark current. These calibrations are made for each CCD and are independent of FoV and gate selection. We Institute for Space Studies of Catalonia Technical Note and University of Barcelona 7
  8. 8. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 expect only minor variations in the gain, and the values from the commissioning phase can be used initially. From time to time it must be checked if any change has occurred. The bias must be monitored also for shorter term variations, as it also depends on the CCD history. The sky background can be studied using the leading and trailing samples in the BP and RP win- dows, and from the observations of virtual objects. Care must be taken to avoid contaminated samples (components, parasites). A list of bad/dead columns must be maintained. These are detected in the analysis of 2D windows of bright stars and calibration faint stars. 3.2 AL LSF t While the AC LSF can be obtained from analysis of 2D windows, the AL LSF is more com- f plex to determine, but anyway it has to be modelled, because the tails of the spectra are fully dominated by the LSF and not by the dispersion. A complicating factor in this connection is ra possible variations in the transmission for the filters, along the filters themselves and with time. The nominal AL PSF as a function of wavelength could be used as a first approximation and we were wondering if it could be updated through a kind of extrapolation of the differences true−nominal AL LSF in AF strips 7-8-9 and the tails of the BP and RP spectra. However, from an inspection of the GIBIS PSFs it remains uncertain if such an extrapolation is possible. D The AL PSF has to be modelled for each FoV, each CCD, each gate, and include a dependence on the pixel column. The responsibility of this modelling is under DU11. 3.3 AC LSF We need an AC LSF for the analysis of 2D windows, for establishing the AC geometrical calibration, and for modelling the AC flux loss. This LSF must include the effect of pixelization, and be expressed for the case of zero AC motion. In specific applications, the LSF must be modified to take the actual AC motion into account. The AC LSF will of course depend on the wavelength, and must be calibrated separately in several (at least three) different bands in BP as well as in RP. The AC LSF calibration can be derived from the 2D windows for each FoV, for each CCD, and possibly for intervals of pixel column and for each gate. The responsibility of this modelling is under DU11. Institute for Space Studies of Catalonia Technical Note and University of Barcelona 8
  9. 9. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 3.4 AC signal cutoff A model of the AC signal cutoff due to lack of resolution and modest window height, could be established from the AC LSF alone. This model should take colour, AC motion, and location in window into account. Refinements to the model must be sought during the internal calibration, cf. Section 4.4. 3.5 Well behaved transits of well behaved sources The sources used for the calibration must be well behaved, and also only individual transits that are well behaved should be used. This generally means that they are isolated and constant, and only few of the transits have been down-weighted during the source update. Variable sources t may, however, be interesting for some purposes. Transits should not be affected by parasites, cosmics, micro meteorites, or other peculiarities. We may also look at the astrometric properties f of the sources and on which transits were misbehaving in the astrometric source update (TBC if this is feasible). Apart from this, calibration sources must be selected using many additional cri- ra teria, and carry flags indicating for which calibrations they are suited. The selection of standard sources is developed by DU16. 4 The internal calibration D We will now discuss how each of the elements of the instrument model may be resolved, and how the various elements may be separated from each other. Common to all calibration steps is that we aim for flux calibrations at the millimagnitude level, and wavelength (AL) calibrations at the 0.01 TDI level, in each step. 4.1 CTI A pronounced CTI will have a detrimental effect on the spectral resolution, and the signal to noise ratio, and may be very serious for the faintest stars. The resolution will suffer especially for the leading samples of the spectrum, i.e. around 650 nm both in BP and RP. It remains unclear how large effects may be expected for BP and RP, and to which extend they can be mitigated through charge injection or field illumination. It also remains unclear to which extend the CTI effects can be calibrated, and how. The models developed for AF will of course be a starting point. We expect the CTI to be: • strongly time dependent; Institute for Space Studies of Catalonia Technical Note and University of Barcelona 9
  10. 10. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 • strongly magnitude and colour dependent; • strongly gate dependent; • strongly AC motion dependent; • dependent on pixel column; • dependent on pixel column history; • dependent on background level. As for time dependence we expect a steady degradation plus sudden jumps caused e.g. by solar eruptions. The magnitude and colour (shape) dependence will probably be part of the model, t and the gate dependence must either be calibrated or we may simply have to reject faint sources affected by gating. f The CTI works at the pixel level, and as the image crosses the CCD the fainter upper and lower ra parts of the image will suffer more than the brighter central spectrum, producing an arrow head. When AC motion is high, the columns where the image is situated at the start of the transit will suffer more than the columns where the image is at the end of the transit, leading to a rotation of the spectral lines. These effects are particularly unpleasant for BP and RP, and will need special attention. D We await detailed studies by Astrium before deciding which calibration model to adopt. 4.2 The AC geometrical calibration The purpose of the AC geometrical calibration is to allow us to predict in which pixel column the image is situated at a given instant of time. The prediction must be accurate to a fraction of a pixel to allow corrections for AC flux loss, cf. Voss et al. (2007). The transformation from the astrometric source parameters to field angles is established by CU3 and essentially requires knowledge of the attitude and the orbit. The remaining task is to establish the relation between the field angles and the pixel columns. For an observation of a suitable star in a 2D window we know the observing time for the central part of the spectrum, and we can find the observed pixel column by a simple centroid. The difference to the column predicted from a nominal calibration must be modelled using suitable polynomials for each FoV, for each CCD, and for each gate. We do not expect much time dependence, but the calibration should be repeated from time to time in order to check. Institute for Space Studies of Catalonia Technical Note and University of Barcelona 10
  11. 11. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 4.3 Geometry: inclination of the spectra We must distinguish between two different inclinations for the spectra. One is the angle between the along scan motion of an image and the pixel columns, and the second is the angle between the pixel columns and the dispersion direction of the prism. The first kind is the more innocent. Due to the scanning law, an image will move in the AC direction with a velocity that varies with a six hour period and an amplitude of 170 mas/s or almost 1 pix/s. This amplitude translates to a motion of up to ±4.2 pixels during the transit of the whole CCD. A slight tilt of the CCD of say 10µm for the top with respect to the bottom, may mean that amplitude instead reaches −3.9 and 4.5 pixels. The result is a smearing in the AC direction, slightly different from the nominal, and a slightly different flux loss. t The second kind, where the CCD and prism are not aligned, is more unpleasant because it produces inclined spectral lines on the CCD and therefore a loss of spectral resolution. We f cannot recover that loss, but we need to know the inclination in order to model the AC flux loss for each part of the spectrum. ra We use the same observations as for the AC geometry, but we now need the column prediction for both ends of the spectrum. The differences to the observed positions give the inclination. 4.4 Refined AC signal cutoff D Most windows are one dimensional, and we have no direct means for compensating for the signal that falls above or below the window. Such a compensation is needed because the fraction of signal lost varies from transit to transit. The correction for signal cutoff must therefore normalise the observations, using a well centred observation as reference. A study of this effect (Voss et al. 2007) shows that we need to know the position of the spectrum within the window to better than 0.1 pixel, as well as the instantaneous AC motion and the colour, in order to obtain a millimagnitude level of precision. The signal cutoff should be modelled in two steps. In a first step it can be calculated from the extrapolated AC LSF model. We can then apply this model to predict the flux loss in the 1D windows and refine the signal cutoff model from a comparison with the observations. This second step should mainly consider large decenterings, and must not be allowed to interfere with the sensitivity calibration. 4.5 The AL geometrical calibration The AL geometrical calibration must provide the relative, effective positions among CCDs and pixel columns for each FoV and strip (BP and RP). The knowledge of these relative positions is Institute for Space Studies of Catalonia Technical Note and University of Barcelona 11
  12. 12. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 needed in the transformation from wavelength to AL pixel coordinates. As emphasised by Bastian in 2007, we can never establish the precise mutual locations of the CCDs in the focal plane, and neither have we much reason to want to do so. What we need is a calibration that allows us to connect the nominal AL and AC field angles, with the predicted field angles at a certain time. The AL geometrical calibration applies to a certain reference wavelength, conveniently chosen as one of the central wavelengths used for representing the spectrum. For a calibration star we can predict an observed spectrum, starting from the astrometric pa- rameters, attitude, orbit, the dispersion, the reconstructed spectrum, the LSF, etc. This is a prediction of the observed signal as a function of time over the 60 ms the spectrum takes to t cross the fiducial line of the CCD, and must be made at sub-TDI time resolution. f We can now compare the predictions with the observed spectrum, e.g. through cross correlation, to find a value for the geometrical calibration, or the correction to the assumed calibration. ra 4.6 The calibration of dispersion The variation of dispersion across the strip leads to variations of the flux per pixel and to vari- ations of the correspondence pixel-wavelength. The dispersion law has to be deduced for each D FoV, each CCD, and each gate, and include the variation with pixel column. The dispersion law of fused silica can be used as a constraint. In the calibration of dispersion the wings of the spec- tra must be avoided, as they are dominated by the LSF and the filter transmission. In principle, this can be better done with sources showing features in their spectra. However, investigations so far suggest that simultaneous determination of AL shift and dispersion is possible using any kind of sources (Carrasco et al. 2007). This is under investigation and subject to confirmation. In order to distinguish conceptually the dispersion calibration from the wavelength calibration, the dispersions must be normalised, e.g. to a certain mean value. This point needs elaboration. 4.7 The sensitivity model The sensitivity model must describe the variations in sensitivity from CCD to CCD, from col- umn to column, between the two telescopes, for each of the relevant gates, as a function of time, and all this as a function of wavelength. A reference level must be established, which could be the average for all calibration units in the strip of CCDs during the first year of mission. We may expect a slow degradation of the response, accompanied by sudden events where some kind of damage or permanent failure sets in. Institute for Space Studies of Catalonia Technical Note and University of Barcelona 12
  13. 13. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 The sensitivity calibration obviously needs constant sources, or at least sources with well un- derstood and predictable variability. Including variable standard sources may be necessary to have sufficiently diverse colours represented. The low- and high-pass filters may not be uniform in their cut-off and so yield different wave- length coverage at different positions of the focal plane. The model will have to deal with this. 4.8 Non-linear response We will return to the question of linearity of the flux scale in Section 5.2. We leave open for present if some part of the flux scale calibration can fit in the internal calibration. 5 Final calibration steps f t ra 5.1 Absolute wavelength scale Absolute wavelengths are established from suitable standard sources, that are calibrated on ground or are classified as having certain characteristic features in their spectra. We note that sources that are suitable for BP are not necessarily useful for RP. D It is desirable to use features in several points of the spectra to get a good determination of the scale, and not simply get a zero-point. The wavelength calibration is not based on individual transits, but on the mean spectra for a relatively small subset of the sources. It is rather independent of the other steps in the photo- metric calibration, and need not enter the iterative calibration loop, but must run after the source updating for the calibration sources. Having established the correct wavelengths, all mean spec- tra must be re-binned by a suitable technique, to assure that the predefined wavelength scale is valid. This is shown schematically in Figure 2. 5.2 Non-linear response The magnitude scale is not expected to be linear. Deviations from linearity are expected at both faint and bright ends due to the detector. The MRD includes a requirement for knowing the non linearity to a 5% accuracy. Given this and that the total RON is of 6.4 e-/sample, the non linear response at the faint end may come from a variable and inaccurately determined background. As discussed by Høg & de Bruijne (2007), the magnitude scale may in principle be calibrated by means of observations of constant stars with gates compared with observations without gates because the attenuation during gating is given by the number of CCD lines. This Institute for Space Studies of Catalonia Technical Note and University of Barcelona 13
  14. 14. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 f t Figure 2: Schematic view of the wavelength calibration of BP and RP spectra. Spectral fea- tures in mean spectra of suitable stars are used for adjusting the zero point and scale of the ra wavelengths. After the wavelengths have been established, all spectra are reconstructed again is true if all CCD lines have the same sensitivity which is not the case. The radiation damage will be a complicating factor. Therefore, provision for checking the calibration as a function of magnitude, colour, position in the sky, etc using external observations to Gaia (TBC) has to be done. This may run as a separate task in parallel to the application of the calibration to all D sources. 5.3 Monitoring of the calibration model An essential step will be to look for weak points in the calibration model. Monitoring of coeffi- cients and residuals may reveal unexpected problems, e.g. some quantity unexpectedly depend- ing on colour or magnitude. New calibration steps will then have to be added. References Bastian U., March 2007, Reference Systems, Conventions and Notations for Gaia, Tech. rep., Astronomisches Rechen-Institut Heidelberg, gAIA-CA-SP-ARI-BAS-003-5 Brown A.G.A., July 2006, Photometry with dispersed images - overview of BP/RP data pro- cessing, Tech. rep., LEI, gAIA-C5-TN-LEI-AB-009 Carrasco J., et al., March 2007, Chi-squared minimisation for wavelength calibration, Tech. rep., Universitat de Barcelona, GAIA-C5-TN-UB-JMC-003-D Institute for Space Studies of Catalonia Technical Note and University of Barcelona 14
  15. 15. BP/RP calibration CU5-DU12 GAIA-C5-TN-UB-CF-006-01 Fabricius C., Masana E., April 2007, Nominal field angles in cycle 2 simulations, Tech. rep., Universitat de Barcelona, GAIA-C2-TN-UB-CF-004-1 Høg E., de Bruijne J., Jan 2007, Calibration of gates and magnitude scale, Tech. rep., gAIA- CA-TN-NBI-EH-182-1 Lindegren L., April 2005, An alternative scheme for the GIS processing, Tech. rep., Lund Ob- servatory, gAIA-LL-059 Lindegren L., March 2007, Convergence properties ofAGIS, Tech. rep., Lund Observatory, gAIA-C3-TN-LU-LL-071-1 Voss H., Jordi C., Fabricius C., et al., April 2007, AC flux loss analysis for AF and BP/RP, Tech. rep., gAIA-C5-TN-UB-HV-001 t The following table has been generated from the on-line Gaia acronym list: f ra Acronym Description AC ACross scan (direction) AF Astrometric Field (in Astro) AGIS Astrometric Global Iterative Solution AL ALong scan ASD Auxiliary Science Data D BP Blue Photometer CCD Charge-Coupled Device CTI Charge Transfer Inefficiency FoV Field of View (also denoted FOV) GIBIS Gaia Instrument and Basic Image Simulator LSF Line-Spread Function ML Maximum likelihood MRD Mission Requirements Document PSF Point-Spread Function RON Read-Out Noise (CCD) RP Red Photometer TBC To Be Confirmed TBD To Be Defined (Determined) TDI Time-Delayed Integration (CCD) Institute for Space Studies of Catalonia Technical Note and University of Barcelona 15

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