Summarizing, the VEGAMAG system is defined such that the bright AOV star a-Lyrae (i.e., Vega) has a magnitude of 0 at all wavelengths. The system was/is convenient for ground-based observers as Vega is a bright star that can be easily observed in the northern hemisphere, and contains a smooth spectrum with few features. The VEGAMAG system is the default SYNPHOT magnitude system, and the magnitude of a star with flux f in this system is simply 2.5log (f/fVega), where f
Vega is the calibrated spectrum of Vega in SYNPHOT. As this system depends on the calibration of the standard star, it is also subject to errors and changes in that calibration. The STMAG and ABMAG systems are different in that they define an equivalent flux density for a source of predefined shape that would produce the observed count rate. In the STMAG system, the flux density is expressed per unit wavelength, and, in the ABMAG system, the flux density is express per unit frequency. The reference spectra are flat as a function of wavelength and frequency in each respective case. The definitions of the systems are:
The photometric zero point of a telescope/instrument/filter combination is a convenient way to characterize the overall sensitivity of the system. By most definitions, the zero point represents the magnitude of a star-like object that produces one count per second within a given aperture (see Maiz Apellaniz 2007). For WFC3, this throughput will measure the final performance taking into account the HST Optical Telescope Assembly (OTA), pick-off mirror, mirror reflectivity, filter throughput, outer window, inner window, and the quantum efficiency (QE) of the detector. For HST instruments such as WFC3, the zero points depend on the absolute flux calibration of HST white dwarf spectra, and therefore they will change whenever that calibration is improved.
Similarly, the most updated STMAG and ABMAG zero points for WFC3 data can be retrieved from photometric keywords in the SCI extension. Specifically, the keyword PHOTFLAM is the inverse sensitivity (erg/cm2/sec/A) and represents the flux density of a star that produces a response of one count per second in this band pass. The header keyword PHOTPLAM is the pivot wavelength. The header keywords PHOTFLAM and PHOTPLAM relate to the STMAG and ABMAG zeropoints through the formulae:
For WFPC2, Holtzman et al. (1995) measured photometric zero points in an intermediate-sized aperture of R = 0.5 arcseconds to alleviate uncertainties in the sky background for measurements made at larger apertures. These can include mapping the extended PSF wings, the digitized effects of the A/D converters, and CTE problems. Such an aperture is more convenient for typical point source photometry, however it cannot be used directly for surface photometry and will require a large correction. For ACS, Sirianni et al. (2005) use a much larger standard aperture of R = 5.5 arcseconds. Such an aperture is impractical for most point source photometry measurements, especially in crowded fields. However, Sirianni et al. (2005) point out that the ACS correction from a small to a large aperture varies strongly from filter to filter and the infinite aperture approach is the traditional SYNPHOT default and therefore a better conversion between point sources and extended sources will be enabled by this convention
Both of the approaches above have advantages and, therefore, for WFC3, we compute zero points both for an infinite aperture and for R = 0.4 arcseconds (WFC3 ISR 2009-30, and
WFC3 ISR 2009-31. Formally, the infinite aperture measurement was obtained by taking the counts (i.e., of a standard star) in a large 2 arcsecond aperture and correcting it a small amount based on a model (see
WFC3 ISR 2009-37 and
WFC3 ISR 2009-38). The infinite aperture value can also be scaled to the zero point in any aperture based on the enclosed energy fractions, which are provided on the same webpage as the zero points. These corrections are wavelength specific. As an example on WFC3/UVIS, the measured flux in F606W within an aperture of radius 0.4 arcseconds (i.e., 10 pixels) is 91% of the total flux and the flux within 2.0 arcseconds is 98% of the total flux. For WFC3/IR, the flux in F140W within an aperture of 0.4 arcseconds (i.e., 3 pixels) is 84% of the total flux and the flux within 2.0 arcseconds is 97% of the total flux.
The WFC3/UVIS CCDs and WFC3/IR detector contain pixels that vary in their area on the sky as a result of the geometric distortion. As a consequence of this, more light will fall on a larger pixel relative to a smaller pixel, leading to an overall gradient in an image of a smooth background. However, the flat-fielding process in the HSTcalwf3 pipeline is designed to produce images that have a flat background (e.g., sky), thereby suppressing counts (hereafter taken to be in units of electrons) in larger pixels relative to smaller pixels. Hence, the measured total brightness of sources on flt images will vary depending on the position of the object, and the areas of the pixels at that location.
To achieve uniform photometry over the detector, most users will measure counts on distortion free images. The geometric distortion can be corrected using AstroDrizzle. The output of this processing will be a drz image, which has a flat sky and contains pixels that are uniform in area (i.e., through proper corrections of the distortion and related pixel area variations). Therefore, photometry of any source in a drz image will yield the same count rate (electrons per second) irrespective of the position of the source on the image. Photometry measured on an flt image therefore requires a field-dependent correction factor to:
A detailed description of the WFC3 UVIS and IR PAMs is provided in WFC3 ISR 2010-08. This description also discusses a unique choice for normalizing the WFC3 PAMs that differs from previous instruments. This choice ensures that the PAMs do not artificially scale the flt flux by large amounts. Rather, the PAMs simply serve to provide a relative correction of the counts based on the size of pixels as compared to the size of a reference pixel near the center of the detectors (see detailed description in the ISR).
AstroDrizzle can be run on the FLT image. The result is that each pixel is free of geometric distortion and is photometrically accurate.
Only by running AstroDrizzle can the geometric distortion be removed, but both approaches correctly recover the count total as 100. Users should be cautioned that this is just an idealized example. In reality the PSF of the star extends to a much bigger radius. If the user decides to work on the flat-fielded image after correcting by the pixel area map, they need to calculate a new aperture correction to the total flux of the star. The aperture corrections discussed in Section 7.2.2 are only for
AstroDrizzle output images. In most cases the aperture correction for distorted images will be quite different from the same star measured in the drz image. This is particularly true for small radius apertures.
Table 9.1 below summarizes the red-leak levels for the WFC3 UV filters. The table lists the fraction of the total signal that is due to flux longward of 400 nm, as a function of effective temperature. This was calculated by convolving a blackbody of the given Teff with the system throughput in the listed filter. As can be seen from the table, red leaks should not be an issue for observations of any objects taken with F275W or F336W. The other UV filters have some red leaks, whose importance depends on stellar temperature. The red leaks in F218W and F300X, for example, exceed ~1% for objects cooler than ~6000 K, while in F225W the red leak reaches ~1% for objects with even cooler temperatures. The most extreme red leaks arise from F218W and F225W observations of objects with Teff of ~4000 K or cooler, necessitating appropriate corrections.
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