1. UVIS Calibration Update
Greg Holsclaw, Bill McClintock
Jan 6, 2009

2. Outline Recent calibration observations
FUV degradation and Spica variability
Stellar flux comparisons with SOLSTICE
The flat field for small / unresolved targets
Changes in the sensitivity at Lyman alpha

3. Recent UVIS Calibrations
The last two Spica calibrations suffered data dropouts due to DSN issues:
FUV2008_265_23_40_33_UVIS_086IC_ALPVIR001_PRIME
FUV2008_309_21_10_33_UVIS_091IC_ALPVIR001_PRIME
2008_265 had enough data to continue the previous analysis of tracking sensitivity degradation and stellar variability

4. Illustration of data dropouts
FUV2008_265_23_40_33_UVIS_086IC_ALPVIR001_PRIME
FUV2008_309_21_10_33_UVIS_091IC_ALPVIR001_PRIME

5. Spica variability

6. Background on Alpha Vir (Spica)
Spica is a non-eclipsing double-lined spectroscopic binary system
Though not spatially resolvable, each component is detectable through measurements of out-of-phase Doppler shifts in the constituent spectral lines
Non-eclipsing due to large apparent orbital inclination of ~70 degrees
Both stars are of a similar spectral class:
Primary: B1V
Secondary: B4V
Spica is the brightest rotating ellipsoidal variable star
The stars have a distorted ellipsoidal shape due to mutual gravitation effects
As the components revolve, the visible area (and thus the observed flux) changes with orbital phase
Since this is a geometric effect, it should be roughly wavelength-independent
Orbital period is days
Amplitude of flux variation in V-filter ~3%
The primary of Spica is a Cepheid variable
Periodic variation in the pulsating primary star is much shorter than the system’s orbital period and about a factor of 2 less in magnitude
Period is 4.17 hours
Amplitude of flux variation in V-filter ~1.5%
This short-term variation, identified in 1968, became undetectable in the early 1970’s (but may return again due to precession of the primary’s rotation axis relative to the orbital plane, which has a period of 200 years [Balona, 1986])

7. Ellipsoidal variation model
Variation in flux is given by [Shobbrook, 1969; Sterken et al, 1986]:
dE = A M2/M1 (R/D)3 (1+e cos(TA+Φ))3 (1-3cos2(TA+TA0+Φ) sin2i )
Where:
A= (wavelength dependent “photometric distortion”)
M2/M1 = 1/1.59 (ratio of masses)
R = 7.6 Rsun = e6 km (polar radius of primary)
D = e7 km (mean separation between stars)
e = 0.14 (orbital eccentricity)
TA (true anomaly)
T0 = days (orbital period)
TA0 = 150 degrees (apparent angle to line of apsides in
year 2005, has precession period of 128 years)
i = 65.9 degrees (orbital inclination)
Φ = empirical phase shift, a free parameter to match with data
One period of the expected variation in flux from Spica

8. Normalized signal vs time
The left plot shows the total FUV signal vs time (normalized to the mean), with a line fit
The right plot shows the same data with this linear trend removed, along with a theoretical model of the Spica ellipsoidal variation that has been fit to the curve (optimizing only the magnitude and phase offset parameters)

9. Stellar flux comparisons

10. Stellar flux comparisons
We can now compare many of the stellar irradiance measurements from UVIS with those made by SOLSTICE
Background on the SOLSTICE instrument:
An FUV/MUV spectrometer onboard the SORCE spacecraft
Built at LASP
Absolutely calibrated at the NIST-SURF facility
Routinely measures stellar fluxes in order to track sensitivity degradation
Irradiance spectra reduced and calibrated by Marty Snow at LASP

11. List of stars observed by UVIS and SOLSTICE
name type UVIS SOLSTICE
alp car canopus F0 x
alp cma sirius A1 x x
alp cru B1 x x
alp gru alnair B6 x
alp eri achernar B3 x
alp leo regulus B7 x x
alp lyr vega A0 x x
alp pav peacock B2 x x
alp peg Markab B9 x
alp psa fomalhaut A4 x x
alp vir spica B1 x x
bet cen B1 x
bet cma B1 x
bet cru mimosa B0 x
bet ori rigel B8 x
del cen B2 x
del cyg B9 x
del sco B0 x x
eps cma adara B2 x
eps ori alnilam B0 x
eps per B0 x
eta uma alcaid B3 x x
gam ori bellatrix B2 x
kap ori saiph B0 x
kap vel B2 x
lam sco shaula B2 x
pi sco B1 x
sig sgr Nunki B2 x
tau sco B0 x
zet cen B2 x x
zet oph O9 x
zet ori alnitak O9 x
zet pup naos O5 x
List of stars observed by UVIS and SOLSTICE
Stars must be very hot in order to produce significant flux in the FUV
Therefore, most stars observed by UVIS are of B and A spectral class
Nine stars were observed by both UVIS and SOLSTICE

12. UVIS spectra
This shows the spectral irradiance curves as measured by UVIS
A variety of shapes (driven by spectral class) and magnitudes are seen

13. Star parameters
We can use a catalog of stellar parameters to estimate the spectral irradiance from any star using the Kurucz model
name Class Vmag Temp Grav Radius Plx
alp vir B
alp psa A
alp cma A
alp cru B
alp leo B
alp lyr A
alp pav B
del sco B NaN NaN NaN NaN NaN
eta uma B
zet cen B NaN NaN NaN NaN NaN
Temp is the effective temperature in Kelvin
Radius is in units of solar radius
Plx is the parallax in arcseconds as measured by Hipparcos
Metallicity assumed to be zero (solar)
Stellar data from: Allende Prieto C., Lambert D.L., Astron. Astrophys. 352, 555 (1999), data accessed through VizieR

14. Kurucz Stellar Models
The Kurucz model stellar irradiance at any wavelength is computed in a coarse grid of three variables: temperature, gravity, and metallicity
A convenient routine for accessing these models is available from:
IUEDAC IDL Software Libraries
The routine KURGET1 allows a user to specify a measured temperature, gravity, metallicity and angular size to arrive at a model spectral irradiance from any star:
Wavelength range: 9.1 nm to 106 μm
Total wavelength steps: 1221
Resolution: ~1nm in the EUV, FUV, MUV
Default units: [ergs/cm2/s/Angstrom]
Allows linear interpolation between most model grid points

15. UVIS, SOLSTICE, KURUCZ
Three Kurucz models show good consistency, two do not

16. Ratio of UVIS to SOLSTICE
This plot shows the ratio of UVIS spectra (1.1nm bins) to SOLSTICE over nm
Here, UVIS spectra were reduced using the flat-field
While the magnitude varies, a consistent shape in the ratio is apparent
WITH flat-field
Q - Is the discrepancy dependent on spectral type, signal level, or position on the detector?

17. Ratio of UVIS to SOLSTICE
Normalized to a mean value of one

18. Ratio of UVIS to SOLSTICE
This plot shows the average ratio of UVIS to SOLSTICE for all stars observed (except Alp PsA)
Average of all stars

19. Ratio of UVIS to SOLSTICE
This plot shows the average ratio of UVIS to SOLSTICE for all stars observed (except Alp PsA)
Expanded vertical scale

20. Wavelength shift
Due to the small uncertainty in spacecraft pointing, a star image could be centered anywhere within six spectral pixels using the low-resolution slit
This is effectively an uncertainty in the wavelength scale of the spectrum
With an FUV dispersion of nm per pixel, this translates to a potential variation of ~0.5 nm
Star image
UVIS detector pixels are 100 x 25 microns (H x W)
Low-res slit is 6 pixels wide
UVIS-FUV Low-resolution entrance slit

21. Wavelength shift adjustment
The phase shift between the measured UVIS spectrum and SOLSTICE can be estimated by finding the peak of the cross-correlation function
This is done by the following procedure:
Shift the UVIS spectrum in one-pixel increments
Smooth the UVIS spectrum to 1 nm
Interpolate the UVIS spectrum to the SOLSTICE wavelength scale
compute the linear correlation coefficient
Find the shift value where the maximum correlation occurs
This is a small effect for 1 nm resolution spectra

22. Cross-correlation analysis
This shows the correlation coefficient versus UVIS pixel shift
A peak in the correlation coefficient is considered a wavelength match
All matches occur within four pixels

23. Average UVIS/SOLSTICE ratio
No wavelength shift
With individual wavelength shifts

24. Ratio of UVIS to SOLSTICE vs SOLSTICE irradiance
UVIS appears to measure fluxes that are higher at lower signal levels
This could be the result of a need to remove an offset, which would preferentially affect measurements of lower signal
There is currently no attempt to adjust for nonlinearity due to the detector dead-time; an adjustment would have the effect of increasing the UVIS signal of brighter stars

25. Ratio of UVIS to SOLSTICE vs surface temperature
The dimmer stars also happen to be the cooler ones
Effective temperatures from: Allende Prieto C., Lambert D.L., Astron. Astrophys. 352, 555 (1999), data accessed through VizieR

26. Stars are imaged in the star-burned rows
This plot shows the spatial distribution of light on the detector from each star, along with the inverse of the average row-to-row correction from the post-burn flat-field
The stars were imaged in the center of the starburned region of the detector
It is here that the sensitivity correction is largest, and the spectral shape of the correction the most significant
Flat-field corrector
Spatial distribution of stars

27. Spatial distribution of stellar flux measurements
No Flat-field applied
WITH Flat-field applied
These plots show the distribution of light on the FUV detector for each observation of stellar irradiance
The curves are normalized to the total signal
The images are centered at about row 32

28. Star comparisons – TO DO
Compare spectra from measurements where the star image is located outside of the starburned rows (Alp Vir, Alp Leo, Alp PsA, others?)
Look into getting extended SOLSTICE spectra (>180nm) from Marty
Check other catalogs to see if the effective temperatures are consistent
How are the stellar parameters of temperature, gravity, and metallicity derived? How accurately are they known?
Include flux from secondary companion stars
Compare the flux from UVIS using Don’s FUV sensitivity curve

29. Flat-fielding for small targets

30. Use of the flat-field corrector for small sources
The flat-field is meant to correct the high-frequency pixel-to-pixel sensitivity nonuniformity
It is derived from the along-slit slew scans of Spica, which approximates a uniform extended source
However, it has been noticed that there could be an issue with using the corrector for small (unresolved) targets

31. Effect of the flat-field on an extended source
This demonstrates the effect of applying a flat-field to a uniform extended source
Sum of all scans, flat field applied, and sum in the spectral dimension
Overcorrection in the starburned rows
This approximates the spatial distribution of a uniform extended source
Sum of all scans, sum in the spectral dimension

32. Effect of the flat-field on a point source
Total signal vs star position
Evil pixels interpolated across
No flat-field applied
Total signal vs star position
Flat-field applied to each frame

33. Sensitivity at Lyman-alpha

34. IPH data reduction Identify all UVIS observations which are:
FUV
Low-resolution
Unbinned
Unwindowed
IPHSURVEY
Average all records in each observation
Mark all evil pixels as NaNs
Average all remaining ‘good’ pixels in each column (row 3 to 60) to form a single spectrum

35. Filtering IPH data All selected LISM spectra in the DAPS database
All LISM spectra which meet a filtering criteria to exclude contamination by stars and data dropouts

36. Filtering IPH data Select all observations prior to the year 2000
Select all observations after the year 2000 with an average (between pixel 300 and 1023) signal in the range ±

37. Normalizing to a fixed LISM signal
All filtered spectra, with the offset subtracted (mean of signal between pixels 990 – 1010)
Need to remove LISM signal variation due to distance from the Sun and pointing.
Normalize each spectrum to the signal between pixel ranges and

38. Signal normalized spectra
Spectra are normalized to the signal in the shaded regions

39. Ratio to an early spectrum
This shows the ratio of all spectra to an early spectrum
This should show the fractional change in sensitivity

40. Ratio to an early spectrum
Look at the change as a function of time for a few pixels

41. Pixel value vs time
This shows the count rate vs time for three spectral columns
Each curve is normalized to the mean of the values before 2000
The variation is much larger than the random uncertainty in the detected counts
Ideas:
Due to a small mislocation in the position of the slit? This is a known effect.
Incomplete removal of stellar spectra?
Colors here denote different pixels

42. Poor background subtraction?
After subtraction of an estimate of the background (mean of pixels )
Raw spectra (counts/second)
Seems unlikely that some pixels decline more than others. More likely that the background subtraction was poor.

43. Summary
There is a small, approximately linear decline in FUV sensitivity
The ellipsoidal variation of Spica continues
Comparisons of UVIS measurements of absolute stellar fluxes with SOLSTICE show that there is a variation in absolute magnitude on the order of 30%, and a systematic shape difference
The flat-field should be used with caution for small targets
There is a measurable decrease in sensitivity at Lyman alpha as a function of time, but the shape is difficult to determine