6.8 Exposure Time Examples
Here are a few simple examples to illustrate how an integration time may be computed for point sources and diffuse sources. The flux values given here are for illustrative purposes only; you need to check the flux values if you are planning your own observations of one of these targets. Also note that the examples in this section have not been updated to take into account the latest revised values for throughputs and other detector parameters, and so the numerical results presented here will differ from the more uptodate ETC calculations.
6.8.1 Spectroscopy of Diffuse Source (M86)
We want to observe M86, an elliptical galaxy in Virgo, using the G750M
grating at a central wavelength setting of λ_{c }= 6768 Å, the CCD detector and the 52X0.2
arcsec slit. Our aim is to calculate the Hα count rate in the central region of M86 and the expected signaltonoise ratio per resolution element for an exposure time of 1 hour. M86 has an inhomogeneous surface brightness distribution in Hα and the line is well resolved with this grating. Let us consider a region with an Hα surface brightness of I_{λ} = 1.16 × 10^{–15} erg/s/cm^{2}/Å/arcsec^{2} (note the unit—it is not the entire Hα flux but the flux per unit wavelength interval) and a continuum surface brightness I_{λ} = 2.32 × 10^{–16} erg/s/cm^{2}/Å/arcsec^{2}. To derive the Hα and continuum count rates from the source we use the formula from Chapter 13:
 S^d_\lambda = 1.14 × 10^{13 }counts/s/pix_{λ}/pix_{s} per incident erg/s/cm^{2}/Å/arcsec^{2}.
 N_{λpix }= 4 and N_{spix} = 2 (1 resolution element).
Using the equation given by the Diffuse Source calculation, we get the count rate C = 0.106 counts/s in Hα and C = 0.021 counts/s in the continuum. Since we are interested in the properties of the Hα line, the Hα counts constitute the signal, while both the Hα counts and the continuum counts contribute to the noise.
The sky background is negligible in comparison to the source, but the dark current (4.5 × 10^{–3} count/s/pix × 8 pixels = 0.020 count/s) and the read noise squared (29 e–/pix × 8 pixels × 3 reads = 465 counts, for CRSPLIT
=3) are important here. Substituting the numbers into the equation for signaltonoise, we get:
StoN=10.7=\frac{0.106\times3600} {\sqrt{(0.127\times3600)+(8\times0.0045\times3600)+(8\times3\times29)}}~. 
To increase our signaltonoise or decrease our exposure time, we can consider using onchip binning. Let us bin 2 pixels in the spatial direction so that N_{bin} = 2. To allow adequate sampling of our new binned pixels, we leave N_{λ}pix = 4, but set N_{spix} = 4, so N_{pix} = 16 and then C = 0.212 for Hα and C = 0.254 for the sum of Hα and continuum. To compute the time to achieve a signaltonoise of 12 using this configuration, we use the full expression for the exposure time given in Section 6.4.1, generalized to treat the line counts (for signal) and total counts (for noise) separately, and determine that roughly 35 minutes are needed in this configuration:
\begin{eqnarray*} t=2104&=&\frac {144\times(0.254+16\times0.0045)} {{2\times0.212^2}}\\ \\ &&+\frac {\sqrt{20736\times(0.254+16\times0.0045)^2 +4\times144\times0.212^2\times8\times3\times29} } {2\times0.2.12^2} \end{eqnarray*}~. 
6.8.2 Spectroscopy of SolarAnalog Star P041C
We wish to study the shape of the continuum spectrum of the solaranalog star P041C from the nearinfrared (NIR) to the NUV. We wish to obtain spectroscopy with the CCD detector covering the entire useful spectral range from 2000 Å to 10,300 Å with gratings G230LB
, G430L
, and G750L
. Since we require accurate photometry, we use the wide 52X0.5
slit. The goal is to reach a signaltonoise ratio of 25 in the NUV (at 2300 Å), 100 in the blue, and 280 in the red. P041C has V = 12.0.
The fluxes of P041C at the desired wavelengths obtained from a spectrum of the Sun scaled from V = –26.75 to V = 12.0, are available via the Web at:
http://www.stsci.edu/hst/instrumentation/referencedataforcalibrationandtools/astronomicalcatalogs/calspec
G230LB
We illustrate the calculation of the exposure time for the G230LB
grating. P041C is found to have a flux of 1.7 × 10^{–15} erg/s/cm^{2}/Å at 2300 Å.
We get the following values for G230LB
from Chapter 13:
 S^p_{2300} = 1.7 × 10^{14 }counts/s/pix_{λ} per incident erg/s/cm^{2}/Å;
 T_{A} = 0.86 for the aperture throughput, taken from Chapter 13;
 ε_{f }= ~0.8;
 N_{spix} =3, since ~80% of the point source light is encircled within 3 pixels;
 N_{λpix }= 2, since two pixels resolve the LSF.
Using the equation in Section 6.2.1, we calculate a point source count rate of C = 0.34 counts/s over N_{pix} = 6 pixels for GAIN=1
.
The source count rates can be compared with the background and detector dark current rates. We’ll make the approximation that the background and detector rates are negligibly small for this setting; therefore we can neglect their contributions. Since we are aiming for a signaltonoise ratio of 25, we can estimate that we must obtain 625 counts minimum. The read noise squared (~350 over 6 pixels for 2 readouts) must therefore be taken into account. Finally, since we are observing with the CCD in the NUV, we must correct for the effect of the multipleelectron process (see Section 6.4). This will cause the exposure time to be scaled approximately by Q, where Q is ~1.5 at 2300 Å. Using the STIS ETC, we estimate the required time for S/N = 25 is ~3560 seconds. To check that we indeed get S/N = 25, we use the formula, in Section 6.4.
S/N\approx 25\approx \frac{(0.34/1.5)3560} {\sqrt{(0.34/1.5)3560+350}}~. 
G750L
and G430L
Exposure times for the two remaining wavelength settings can be calculated directly as time = signaltonoise^{2}/C since the read noise, detector background, and sky background are negligible. As above, 3 pixels are taken to contain 80% of the flux. The results are summarized in Table 6.6.
Table 6.6: Low Resolution Spectroscopy of Solar Analog Star.
Grating 



Wavelength (Å)  2300  5000  7800 
Flux (ergs/s/cm^{2}/Å)  1.7 × 10^{–15}  5.9 × 10^{–14}  3.5 × 10^{–14} 
Point Source Sensitivity  1.7 × 10^{14}  3.1 × 10^{15}  5.0 × 10^{15} 
Aperture throughput (T_{A})  86%  90%  89% 
N_{λpix}  2  2  2 
N_{spix }to encircle 80% of PSF  3  3  3 
C (counts/s from source over N_{λpix} = 2)  0.34  240  240 
Signaltonoise ratio desired  25  100  280 
Total exposure time  3560 seconds  41 seconds  330 seconds 
6.8.3 Extended Source with Flux in cgs Units (NGC 6543): Imaging and Spectroscopy
Let us consider NGC 6543, the Cat’s Eye planetary nebula, where the aim is to use the CCD to image using the [O II] filter, and to do spectroscopy both in the visible and in the UV.
Imaging
The aim is to get a signaltonoise ratio of 30 using the [O II] filter. We know that NGC 6543 is about 6 times fainter in [O II] than in Hβ, and its total flux at [O II] 3727 Å is ~4.4 × 10^{–11} erg/s/cm^{2} contained within 1 Å. Since the radius of the object is about 10 arcseconds, the average [O II] surface brightness is about 1.4 × 10^{–13} erg/s/cm^{2}/arcsec^{2}/Å.
We take:
 S^d_\lambda = 6.7 × 10^{11 }counts/s/pix/Å per incident erg/s/cm^{2}/Å/arcsec^{2} as given in Chapter 14.
 We take N_{pix} = 4 × 4 = 16, since a resolution element has radius of two pixels (see Chapter 14).
To calculate the count rate we use the equation in Section 6.2.1 for diffuse sources and determine a perpixel count rate of 0.094 counts/s/pix or a count rate C = 1.5 counts/s over 16 pixels. The background and the dark current can be neglected. To get a signaltonoise of 30 we need ~10^{3 }counts, so the read noise can also be neglected and we can use the simplified expression to calculate exposure time (see Section 6.4.1). We obtain 10^{3} counts in ~667 seconds. To allow postobservation removal of cosmic rays we use CRSPLIT=2
. We note that in each ~330 second exposure we predict a mean of ~31 counts/pix, and thus we are safely within the limits of the CCD full well.
Diffuse Source Spectroscopy in the Visible and Ultraviolet Regions
In the visible, the aim is to get a signaltonoise of about 100 at λ = 4861 Å, with the G430M grating at a central wavelength setting of λ_{c} = 4961 Å, the CCD detector, and the 52X0.1
arcsecond slit. In the UV, the aim is to get a signaltonoise ratio of about 20 at the C IV
~1550 Å line with the G140M
grating at a central wavelength setting of λ_{c} = 1550 and the FUVMAMA detector
. To increase our signaltonoise ratio in the UV, we use the 52X0.2
arcsecond slit for the G140M
spectroscopic observations.
Visible Region
NGC 6543 has an average Hβ surface brightness of S(Ηβ) ~ 8.37 × 10^{–13} erg/s/cm^{2}/Å/arcsec^{2} at 4861 Å and has a radius of about 10 arcseconds.
We take from Chapter 13:
 S^d_\lambda = 1.62 × 10^{12} counts/s/pix_{λ}/pix_{s} per incident erg/s/cm^{2}/Å/arcsec^{2 }for
G430M
;  N_{λpix} = N_{spix }=2 since 2 pixels resolves the LSF and PSF;
Using the equation for diffuse sources in Section 6.2, we derive a perpixel count rate of 1.4 counts/s/pix and a count rate integrated over the four pixels of C = 5.4 counts/s at 4861 Å from the astronomical source. The sky background and the detector background are much lower. To allow cosmic ray removal in postobservation data processing, we use CRSPLIT=3
. To achieve a signaltonoise of 100, we require a total of roughly 10,000 counts, so read noise should be negligible, even over 4 pixels and with NREAD=3
. We calculate the time required to achieve signaltonoise of 100, using the simplified equation Section 6.4.1, and determine that we require roughly 30 minutes:
t=1850=\frac{10000}{5.4}~. 
At a count rate of ~1 counts/s/pix for 600 seconds per CRSPLIT
exposure, we are in no danger of hitting the CCD fullwell limit.
Ultraviolet Region
The C IV flux of NGC 6543 is ~2.5 × 10^{–12 }erg/s/cm^{2}/arcsec^{2} spread over ~1 Å. The line, with a FWHM ~ 0.4 Å, will be well resolved in the G140M
configuration using the 52X0.2
slit.
We take from Chapter 13:
 S^d_\lambda = 5.15 × 10^{9 }counts/s/pix_{λ}/pix_{s} per incident erg/s/cm^{2}/Å/arcsec^{2 }for
G140M
at λ = 1550 Å using the 0.2 arcsecond wide slit.  We take N_{λpix} = N_{spix} = 8, since the line emission is spread over the ~8 pixels of the slit width in dispersion, and we are willing to integrate flux along the slit to improve the signaltonoise ratio.
Using the equation for diffuse sources in Section 6.2.1, we determine a perpixel peak count rate of ~0.013 counts/s/pix and a count rate over the 64 pixels of C = 0.82 counts/s at 1550 Å from the astronomical source. The sky and detector backgrounds are still negligible, and the read noise is zero for the MAMA detector so we can use the simplified equation for exposure time, see Section 6.4.1. We determine that we require ~7 minutes:
t=490=\frac{400}{0.82}~. 
We are well below the MAMA local linearity limit of 50 counts/s/pix. Even assuming the nebula evenly illuminates the full 28 arcseconds of the long slit, we are well below the global absolute and linearity limits, since the flux from the nebula is concentrated in the C IV
emission line. Then the global count rate, if the source fully fills the slit in the spatial direction, is given roughly by (0.015 × 8 × 1024) \ll200,000 counts/s. Finally, we are well below the MAMA 16 bit buffer limit of a maximum of 65,536 counts/pix integrated over the exposure duration.
6.8.4 Echelle Spectroscopy of a Bright Star, Reddened LMC Star
The aim here is to do highresolution echelle spectroscopy of an O5 star in the LMC at 2500 Å, using the E230H grating at a central wavelength of λ_{c} = 2513 Å and using the 0.2X0.09
arcsecond slit. The aim is to get a signaltonoise ratio of about 50 from photon statistics. We will assume that the exact UV flux of the star is unknown and we need to estimate it from the optical data. This calculation of the stellar flux at 2500 Å involves 2 steps:
 Calculation of the dereddened flux at 5500 Å.
 Calculation of predicted flux at 2500 Å taking reddening with standard extinction and stellar models into account.
Dereddened Magnitude and Prediction of 2500 Å Flux
We assume that it is an O5 V star with V = 11.6 and B – V is –0.09. The expected B – V value from such a star is –0.35, so that E(B – V) = 0.26 mag.
We assume all the extinction to be due to the LMC, and use the appropriate extinction law (Koornneef and Code, ApJ, 247, 860, 1981). The total visual extinction is then R × E(B – V) = 3.1 × 0.26 = 0.82, leading to an unreddened magnitude of V_{0} = 10.78. The corresponding flux at 5500 Å (using the standard zero point where V = 0 corresponds to F(5500 Å) = 3.55 × 10^{–9} erg/s/cm^{2}) is F(5500 Å) = 1.73 × 10^{–13} erg/s/cm^{2}/Å.
The model atmosphere of Kurucz predicts F(2500 Å)/F(5500 Å) = 17.2 for an O5 star, which leads to a flux of F(2500 Å) = 2.98 × 10^{–12} erg/s/cm^{2}/Å at 2500 Å for the unreddened star. Reddening will diminish this flux by a factor of 10^{–0.4xA(2500 Å)}, where the absorption at 2500 Å can be determined from the extinction curve; the result in this case is ~0.3. Thus the predicted flux of this star at 2500 Å is 9.0 × 10^{–13} erg/s/cm^{2}/Å.
Exposure Time Calculation
We take from Chapter 13:
 = 2.9 × 10^{11} counts/s/pix_{λ} per incident erg/s/cm^{2}/Å for
E230H
;  T_{A} = 0.659 for the aperture throughput;
 ε_{f} = 0.8 for the encircled energy;
 N_{λpix} = 2, since two pixels resolve the LSF;
 N_{spix }= 3, since 80% of the point source light is encircled by 3 pixels.
Using the equation for point sources in Section 6.2, we determine a total count rate from the star of C = 0.3 counts/s over 6 pixels. From Chapter 13 we see that ~22 percent of the point source flux will be contained within the peak pixel. Thus the peak per pixel count rate will be approximately 0.3 × 0.22/(0.8 × 2) = 0.045 counts/s/pix and well within the local linear counting regime. We can use the information that we register ~0.3 counts/s for every two pixels in the dispersion direction to estimate the global count rate (over the entire detector) as follows. Each order contains ~1024 pixels, and the E230H grating at the central wavelength setting of 2513 Å covers 33 orders (see Chapter 13). A rough estimate of the global count rate is thus ~33 × 512 × 0.3/0.8 ~6400 counts/s and we are well within the linear range.
To calculate the integration time, we can ignore both the sky background and the detector dark current which are several orders of magnitude fainter than the source. To achieve a signaltonoise ratio of 50, we then require ~2500 counts which would take a total of ~2.3 hours. Fortunately, this is a CVZ target!
6.8.5 Imaging a Faint Stellar Source
Consider a case where the aim is to image a faint (V = 28), Atype star with the clear filter and the CCD detector. We want to calculate the integration time required to achieve a signaltonoise ratio of 5. The count rate from the source is 0.113 counts/s distributed over about 25 pixels using the information in Chapter 14. If we assume the background to be “typical high” (Table 6.3), the count rate due to the background integrated over the bandpass is ~0.15 counts/s/pix or 3.8 counts/s in 25 pixels (and the detector dark rate is 35 times lower). We will need to be able to robustly distinguish cosmic rays if we are looking for faint sources, so we will use CRSPLIT=4
. We use the STIS ETC to estimate the required exposure time to be 8548 seconds. To reproduce the numbers given by the ETC, we use the equation in Section 6.4.1:
SN=5=\frac{0.113\times8548} {\sqrt{0.113\times8548+0.15\times25\times 8548+0.0045\times25\times8548+25\times29\times4}}~. 
Alternately, we could have requested LOWSKY
(see Section 6.5.2), since these observations are skybackground limited. In that case the sky background integrated over the bandpass produces ~0.035 counts/s/pixto which we add the detector dark current to get a total background of 0.039 counts/s/pix. Using the full equation for exposure time again, we then determine that we require only ~60 minutes. This option is preferable to perform this experiment. To check the S/N, we use the equation in Section 6.4.1:
SN=5=\frac{0.113\times3645} {\sqrt{(0.113\times3645)+(0.035\times25\times3645)+(0.0045\times25\times3645)+(25\times29\times4)}}~. 
6.8.6 TimeTag Observations of a Flare Star (AU Mic)
Suppose the aim is to do TIMETAG
observations of a flare star such as AU Mic, in the hydrogen Lymanα 1216 Å line (see Section 11.1.3). We wish to observe it with the G140M
grating, the MAMA detector and a 0.2 arcsecond slit. AU Mic has V = 8.75, the intensity of its Lyα line is about 6 (±3) × 10^{–12} erg/s/cm^{2}/Å, and the width (FWHM) of the line is about 0.7 (±0.2) Å. We will assume that during bursts, the flux might vary by a factor of 10, so that the line flux may be up to 60 × 10^{–12} erg/cm^{2}/s/Å.
AU Mic is an M star and its UV continuum is weak and can be neglected.
We use from Chapter 13:
 S^d_\lambda = 2.30 x 10^{12} counts/s/pix_{λ} per incident erg/s/cm^{2}/Å;
 Aperture throughput T_{A} = 0.6;
 Encircled energy ε_{f }= 0.8;
 N_{spix} = 10;
 Derive N_{λpix} = 14 since the line FWHM is ~0.7 Å and the dispersive plate scale for
G140M
is 0.05 Å/pix;
Plugging these values into the point source equation in Section 6.2.1, we get C = 927 counts/s over 10 × 14 pixels, or ~1160 counts/s from the source during a burst (taking ε_{f }= 1.0). This is well below the MAMA TIMETAG
global linearity limit of 30,000 counts/s and the continuous observing limit of 26,000 count/s. The line is spread over 14 pixels in dispersion and roughly only 10% of the flux in the dispersion direction falls in the peak pixel; thus the peak perpixel count rate, P_{cr}, is roughly 927/(14 × 10) = 7 counts/s/pix, and we are not near the MAMA local linearity limit.
For a TIMETAG
exposure, we need to determine our maximum allowed total observation time, which is given by 6.0 × 10^{7}/C seconds or roughly 1079 minutes = 18 hours. For Phase II only, we will also need to compute the value of the BUFFERTIME
parameter, which is the time in seconds to reach 2 × 10^{6} counts, in this case 2157 seconds (= 2 x 10^{6}/927).

STIS Instrument Handbook
 • Acknowledgments
 Chapter 1: Introduction

Chapter 2: Special Considerations for Cycle 28
 • 2.1 STIS Repair and Return to Operations
 • 2.2 Summary of STIS Performance Changes Since 2004
 • 2.3 New Capabilities for Cycle 28
 • 2.4 Use of AvailablebutUnsupported Capabilities
 • 2.5 Choosing Between COS and STIS
 • 2.6 Scheduling Efficiency and Visit Orbit Limits
 • 2.7 MAMA Scheduling Policies
 • 2.8 Prime and Parallel Observing: MAMA BrightObject Constraints
 • 2.9 STIS Snapshot Program Policies
 Chapter 3: STIS Capabilities, Design, Operations, and Observations
 Chapter 4: Spectroscopy
 Chapter 5: Imaging
 Chapter 6: Exposure Time Calculations
 Chapter 7: Feasibility and Detector Performance
 Chapter 8: Target Acquisition
 Chapter 9: Overheads and OrbitTime Determination
 Chapter 10: Summary and Checklist
 Chapter 11: Data Taking

Chapter 12: Special Uses of STIS
 • 12.1 Slitless FirstOrder Spectroscopy
 • 12.2 LongSlit Echelle Spectroscopy
 • 12.3 TimeResolved Observations
 • 12.4 Observing TooBright Objects with STIS
 • 12.5 High SignaltoNoise Ratio Observations
 • 12.6 Improving the Sampling of the Line Spread Function
 • 12.7 Considerations for Observing Planetary Targets
 • 12.8 Special Considerations for Extended Targets
 • 12.9 Parallel Observing with STIS
 • 12.10 Coronagraphic Spectroscopy
 • 12.11 Coronagraphic Imaging  50CORON
 • 12.12 Spatial Scans with the STIS CCD

Chapter 13: Spectroscopic Reference Material
 • 13.1 Introduction
 • 13.2 Using the Information in this Chapter

13.3 Gratings
 • FirstOrder Grating G750L
 • FirstOrder Grating G750M
 • FirstOrder Grating G430L
 • FirstOrder Grating G430M
 • FirstOrder Grating G230LB
 • Comparison of G230LB and G230L
 • FirstOrder Grating G230MB
 • Comparison of G230MB and G230M
 • FirstOrder Grating G230L
 • FirstOrder Grating G230M
 • FirstOrder Grating G140L
 • FirstOrder Grating G140M
 • Echelle Grating E230M
 • Echelle Grating E230H
 • Echelle Grating E140M
 • Echelle Grating E140H
 • PRISM
 • PRISM Wavelength Relationship

13.4 Apertures
 • 52X0.05 Aperture
 • 52X0.05E1 and 52X0.05D1 PseudoApertures
 • 52X0.1 Aperture
 • 52X0.1E1 and 52X0.1D1 PseudoApertures
 • 52X0.2 Aperture
 • 52X0.2E1, 52X0.2E2, and 52X0.2D1 PseudoApertures
 • 52X0.5 Aperture
 • 52X0.5E1, 52X0.5E2, and 52X0.5D1 PseudoApertures
 • 52X2 Aperture
 • 52X2E1, 52X2E2, and 52X2D1 PseudoApertures
 • 52X0.2F1 Aperture
 • 0.2X0.06 Aperture
 • 0.2X0.2 Aperture
 • 0.2X0.09 Aperture
 • 6X0.2 Aperture
 • 0.1X0.03 Aperture
 • FPSPLIT Slits 0.2X0.06FP(AE) Apertures
 • FPSPLIT Slits 0.2X0.2FP(AE) Apertures
 • 31X0.05ND(AC) Apertures
 • 0.2X0.05ND Aperture
 • 0.3X0.05ND Aperture
 • F25NDQ Aperture
 13.5 Spatial Profiles
 13.6 Line Spread Functions
 • 13.7 Spectral Purity, Order Confusion, and Peculiarities
 • 13.8 MAMA Spectroscopic Bright Object Limits

Chapter 14: Imaging Reference Material
 • 14.1 Introduction
 • 14.2 Using the Information in this Chapter
 14.3 CCD
 14.4 NUVMAMA

14.5 FUVMAMA
 • 25MAMA  FUVMAMA, Clear
 • 25MAMAD1  FUVMAMA PseudoAperture
 • F25ND3  FUVMAMA
 • F25ND5  FUVMAMA
 • F25NDQ  FUVMAMA
 • F25QTZ  FUVMAMA, Longpass
 • F25QTZD1  FUVMAMA, Longpass PseudoAperture
 • F25SRF2  FUVMAMA, Longpass
 • F25SRF2D1  FUVMAMA, Longpass PseudoAperture
 • F25LYA  FUVMAMA, Lymanalpha
 • 14.6 Image Mode Geometric Distortion
 • 14.7 Spatial Dependence of the STIS PSF
 • 14.8 MAMA Imaging Bright Object Limits
 Chapter 15: Overview of Pipeline Calibration
 Chapter 16: Accuracies

Chapter 17: Calibration Status and Plans
 • 17.1 Introduction
 • 17.2 Ground Testing and Calibration
 • 17.3 Ground Testing and Calibration
 • 17.4 Cycle 7 Calibration
 • 17.5 Cycle 8 Calibration
 • 17.6 Cycle 9 Calibration
 • 17.7 Cycle 10 Calibration
 • 17.8 Cycle 11 Calibration
 • 17.9 Cycle 12 Calibration
 • 17.10 SM4 and SMOV4 Calibration
 • 17.11 Cycle 17 Calibration Plan
 • 17.12 Cycle 18 Calibration Plan
 • 17.13 Cycle 19 Calibration Plan
 • 17.14 Cycle 20 Calibration Plan
 • 17.15 Cycle 21 Calibration Plan
 • 17.16 Cycle 22 Calibration Plan
 • 17.17 Cycle 23 Calibration Plan
 • 17.18 Cycle 24 Calibration Plan
 • 17.19 Cycle 25 Calibration Plan
 • 17.20 Cycle 26 Calibration Plan
 Appendix A: AvailableButUnsupported Spectroscopic Capabilities
 • Glossary