Determination of Martian surface reflectivity from 0.4 to 1.1 micron using a vidicon spectrometer.

Determination of Martian Surface Reflectivity

From 0.4 to 1.1 Micron Using a Vidicon Spectrometer

by

Douglas John Mink

Submitted in Partial Fulfillment of the Requirements for the Degree of Master of Science

at the

Massachusetts Institute of Technology May, 1974

/ U ' Z -. vr

Department of Ejrth n Planetary Sciences, May 20, 1974 Cer t i f ied by: -w W-9

Accepted by:

Chairman, Departmental Commi ttee on Graduate Students

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pag .

Determination of Martian Surface Reflectivity from 0.4 to 1.1 Micron Using a Vidicon Spectrometer

by

Douglas John Mink

Submitted to the Department of Earth and Planetary Science in part-ia.I fulfillment of the requirements for the degree of

Master of Science on May 23, 1974

ABSTRACT

A new astronomical instrument, the vidicon spectrometer,. is being developed at the M.I.T. Planetary Astronomy Laboratory. Based on the silicon diode vidicon system currently in use there, a low dispersion prism is added between the vidicon image tube and

the telescope, allowing digital vidicon photographs to he taken df spectra. These spectra are stored on magnetic tape and computer processed to create intensity vs. wavelength curves for stars -Ind planets. The high spatial resolution of the vidicon imagle tubte, combined with

a

higher spectral resolution than photometer fil ters currently -in use at M.I.T. give this instrument potent ial in the study of planetary surface composition from spectral ref lectivitt.. Procedures for reducing the vidicon images to spectra have bf'en

tested on a set of spectra of tio stars and the planet Mlars. It is concluded that the viclicon response is not linear enough wi th variations In exposure time at lot, levels of incoming light for consistent star spectra. although it works well with Mars clue to the planet's larger intensity where the vidicon tube has i ts poorest response. The spectrometer slit is so narrow (one seconr of arc for this data) that uaveleng.th-dependent var iations in refraction of light from a point source by the atmosphere .caw1r'e

,tar spectra of variable clualitl. Because of the lou quality of the star spectra. direct spectral rrflectivity measurenwnts (which are obtained using Mars to star ratios) proved to be impossible. Although further tests of the spectral and intensi ty response of

the silicon dinde Vidicon should be carried out in the laboratory before good results can be gua-anteed, the presont Mars spectra may probably be used in conjunction With photometer-derived reflectivity data to expand coverage of the surface of Mars.

Thesis Advisor: 'Thomas B. McCord

page 3

Acknowledgements

Numerous people were involved in the design and construction of the.vidicon spectrometer, although I have only met a few.

Professor Thomas B. McCord, Mr. Jeffrey Bosel, and Ms. Carle Pieters have been informative about the design of the hardware, as well as the conditions under which they obtained the Mauna Kea data. Dr. Robert Huguenin lent his insight into the problems of Martian surface composition, as well as criticizing an early draft of this thesis. Stevp Kent, who i4 also working with the

spectrometer, provided criticism and astronomical knowledge. My spouse. Catherine, has greatly aided me financially' and morally. while David McOonald greatly expedited the production of this

pag' 4

Table of Contents

I. Introduction

II. The Vidicon

Spectrometer-111.

Image

_Processing

IV. Anal-ysis of Data

V. Recommendat ions for Future Use

_{page 50}

page 5

page 9

page

16

page 5

1. Introduction

Although Mariner 9 has returned a vast quantity of information about the planet Mars,~ little was learned about surfsce composition. From such experiments as the infrared spectrometer., particle size and silica composition were estimated, but these. determinations had error bars so great as to be nearly useless In reaching conclusions about the composition of the surface materialso-f rars. Until the Viking Lander in 1976, there

is no way to physically look at a Martian rock with instruments. Probably the moot useful technique for remotely sensing surface composition is reflectance spectroscopy. Dollfus (1961), studying the polarization of light reflected by Mars , concluded

that limonite, a hydratdd iron oxide, was probably a major constituent. Hovis (1365) observed absorption bands in the

near-infrared reflectivity of limonite and suggested that they would be a diagnostic test for limonite on Mars. Sagan et al (1365) compared absorption pands they observed in laboratory specimens of

lisonite to Dollfus' Martian albedo curves and concluded that a surface with at least some limonite was not inconsistent with the data. Adams (1968) observed absorption bands between 0.5 and 2.5 microns in many iron-bearing minerals, the positions of which varied significantly from mineral to mineral. These bands are caused by electron transitions in iron ions and by vibrational bands i.n hydroxyl iokns and water molecules. Adams suggests that

page G

the absorption feature observed in Tull's (1966) geometric albedo curve is not inconsistent with a hydrated basalt composi tion. The feature observed at one micron in their spectra is not due to iron in iron oxides, but to iron ions in silicates. Adams anrd McCnrd (1969), using geometric albedoes obtained during the 0i;7

opposition discovered that curves for the bright aras ihad different shapes than those from the dark areas of the Mar ti:in

surface. They concluded that the surface was conposed of a combination of oxidized basalt and hydrated iron oxides. The bright and dark areas were modelled as being composed of of the

same.material in different degrees of oxidation.

McCord and Westphal (1971, see also McCord, Elias, and Westphal, 1971) observed Mars during the 1969 opposition and noted that the iron ion absQrp-tions were in different places, indicating compositional differences. Seven areas were observed, four dark and three bright, each being about five Martian longitudinal degrees in diameter. From this data, much compositional analysis has been. done (see Figure I for examples of mineral ref lectivities compared to Mars); however, from such a small sample, generalizations about the rest of the surface cannot be made. Despite over twenty additional spots obtained during the 1973 opposition, such interesting features as the Coprates canyon and the Hellas basin remain uncovered; what is needed is tihole disk coverage at higk $pectral and spatial resolution, A new technique, vidicon spectroscopy, has been developed to obtain the

MOAB . - I tI li-1.00 --0.00 0.3 0.5 11 WAVELENGTH (p) -j .4. Cr. 01-wI N3 *-*0 ) J 0- 5 4w 0 z ARABIA _.. -S10J '003 0.5 .O WAVELENGTH (p) AREA 59/74 ~ -

iiIfifif

1.00-0.00 1 L t -1 0.3 05 WAVELENGTH (p) ) ' 1.5 WAVELENGTH (Mm)

Figure 1.

Comparison of Mars

dark area reflectivity

to reflectivity of sheet

silicates. Note

resem-blance to the clay

mine-rals, kaolinite and

mont-morillonite.

(Courtesy Dr.Robert

Huguenin)

-j U-0 z

page 7

page 8

desired high resolution ful -disk coverage. This thesis describes that technique.

page 3

11. The Vidicon Spectrometer

The silicon diode array vidicon *was originally developed for television and picturephone use, but because of its large dynamic range, high quantum efficiency, and linear response, it is being used by a growing number _{of astronomers as a digital} replacement for photographic plates. _{The only advantage a photographic plate} has over a vidicon is spatial resolution; however, that is not a limiting factor as atmospheric conditions are the resolution-limiting factor in astronomy. _{McCord and Westphal (1972), Kunin} (1972), and McCord and Bosel (1973) have reported on the

development of a vidicon system for single-frame astronomical photography at the Planetary Astronomy Laboratory of the Massachusetts Institute of Technology (MITPAL). This system is based on an RCA silicon vidicon tube with a peak quantum

efficiency of 85% at 6.5 microns, sloping off to about 6% at 1.1 microns (see Figure 2). Using filters _{this system has been} developed as a two-dimensional imaging photometer, using filter sets similar to those _{used with photometers for spectral} reflectivity work at MIT. As reported by McCord and Bosel, a vidicon spectrometer which w

the vidicon combined' with a

a vidicon photometer is under

The vidicon spectrometer

is attached to the front

telescope.

Schematically it

ould give the spatial resolution of greater spectral resolution than such

development.

is basically an optical syctem which end of the vidicon system on the

page 10

20

0.~~4.4

₊

₊

.

6008

.012

>4.

WAVELENGTH IN MICRONS

Figu re 2.

Quantum

efficiency of the RCA vidicon.

This is the percentage of incoming photons which

the diode array and affect it as opposed to being

reflected or passing through without being absorbed.

This graph was made by averaging the published

.

page

11

through 'hich 'light

telescope is

passed.

refocused onto the

practice this is

done

from a slit situatec The dispersed image surface of the vidi

through a system of

I

at

the focus

of the

of the

slit

is then

con diode array. In

mirrors (see Figure 3

ls) to avoid the infrared absorpt

vidicon tube consists of a 1024

biased diodes. A photon impinging on the

in

a decrease

in charge

i

read out by scanning the di

recharges the diodes as

proportional to the amount

beam is at any given time,

diode array

'can

be known.

passed on to be recorded an

electronics of a silicon v

n the diode ode array it hits of charge I the inten it wi th them os t. si ty ion of

by

1024

vidico reache

an ele

prod

By k

at eac

lenses. array of revereae

n target resul'ts

s. The

_{image is}

ctron beam which ucing a current nowing where the

h location in the

These intensity elements

d displayed (for further detai idicon see Crowell and Labuda

are then Is on the (1363)). The vidicon is read out as 250 rows of 256 image elements, each of

which

corresponds

physically

to

four diodes.

In such a

loier

resolution scan, less accurate positioning is required of the

electron beam.

No

data

is lost, and the vidicon's resolution

is

still better than the atmosphere allows. The intensity image is

amplified, recorded on magnetic tape, and displayed on a slow scan

TV monitor. This image is then available for further computer processing. _{The spectrometer system is diagrammed in Figure 4.}

A portion of a vidicon spectrometer image is presented in Figure 5. The elements along the column correspond to spatial

for detai The

page 12

CAMERA

image is

in focus

BAFFLES

SPHERICAL

MIRROR

BAFFLES

VIDICON '

Figure 3.

LOW

DISPERSION

PRISM

image is

in focus

here

Optics of the MITPAL vidicon spectrometer.

The telescope is to the right.

page 14

elements along the slit. Wavelength is along the abscissa. The magnitude of each element is proportional to the current from the vidicon diode array at the time a corresponding diode uas read by

the scanning electron beam. The image is now ready to be turnerd into a spectrum.

6

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.page- 16~

111. Image Processing

The first processing that must be done to the image is to convert the " column coordinate into wavelength. This is clono

through the use of a calibration function:

S= -SO+

0

=)

+

(X'- xO)

(S+

So)-So,

,

and C being. three constants determined from three column number-wavelength correspondences as fol lows:

C =N( - ) (St+ SO) (S2+SO) so -St+(XXS - St) $ ($- S

(4-t4)(S2-Si)( -

Si-X)

6

-So.= -S + (C -9 $ - 2

A

) S - i

C

9

..

These correspondences are obtained by observing the spectrometer image of a calibration lamp with known sharp emission lines

fav

show4n in Figure G). From this calibration, which is redons: periodically as data is taken, the wavelength-column relationship is known - (see Figure 7 .for an example). The resolution also

0 varies as

a

function of wavelength, as would be expected (see Figure 8 for a sample dispersion function plotted from the first derivative of the calibration function).

0.

64.00

98.00

128.0

180.00

192.00

2N,

VIDICON COLUMN

A spectrum of the calibration source, indicating vidicon intensity

of each vidicon element along one row. Assigned wavelengths are

indicated.

Figure 6.

V1DCu)2 F.125 Sir 35.0 La 0.3317 S4= 57.5 L2= 0.357 S3* 199.5 L3s 0.809 LC= 0.1541 Cz-41.73d8 SO*-263.2354 - _-1 0.3.133 51 - C.3508 101 0.4114 151 0.526C 201 0.8248 ? 0.3139 52 0.31~~ 02c0.413~152 0.524~ 0.8357 1 0.3145 53 0.3527 103 0.4146 153 0.5328 203 C.8471 4 ;. 31,f .')4

0.3*416

10 5 0.3158 55 0.3546 105 0.4179 155 0.5398 205 0.8708 6 0.36456~~.35~5~~106~~d.419~~156~0.543~3~~~2'6~~~~8834~ 7 C.317G 57 (.3565 107 . 0.4213 157 0.5470 207 0.8963 ti 0.~3~17 T58~ 0.315'~ LOP~ 0.4230~~ 158 0.507 20q~~0~. 0099~ 9 C.1113 59 0.4 85 109 0.4?47 159 0.5546 209 0.9237 ia M.0169 6 .'' 42V0 .9382 11- 0.3196 - 61 0,3605 111 0.4283 161 0.5624 211 0.9532 12 0~.32C3~~ ~~62~ ~0.3215 ( 112 ~0.4301 162 0.5664'~1 ~0.96 8 13 0..32C9 63 0.306 113 0,4319 163 0.5705 213 0.9850 14 0.16 601 15 0.3223 b 0.3 47 115 0.4357 165 0.5190 215 1.C194 17 03216 67 ('.3668 117 0.43q5 167 0.5878 211 1.0569 16 0.3243 0.3b79 lj8O.446 .076 19 C.32t0 64 '0.3(-Q 119 0.4435 169 0.5970. 219 1.0977 zu .2 700.3/ 1 20 26 O. 4 4 67 0618 6 O 20 1.A195V /1 0.3264 11 0.37121?1 0.4476 111 0:6C66 221 1:1424 le i 0.34271 72 0.11 -M -- 664 23 C.3279 73 0.3735 123 0.4518 113 0.6167 223 1.191S .44 0. 32 6 4 0.3 741 124 -0.4534 174.6219 I 224 .2179 25 0.3293 75 1.315q 125 0.4561 0:6272 745 1.2458 426 0.3301 76 0.453 176 6 27 0.33C8 7 0.37!2 1?7 0.460i 117 0.6381 227 1.3060 4d 0.3316 -78 06i7Th 128 0.462d 1 22 "8338 29 0.3323 79 0.3bV7 L20 0.4651 119 016496 229 1.3733 30 1.331 80 0 3 0 1 .- 2 - -I' 1.304 01 03 149 81 0.483? 141 '0.4698 101 . 0. 6617 231 1.4489 . 32 0.3346 82 Z 3644 1320.4 1 0.667 22 1.4904 33 C.33i4 e3 0. U) 133 .4746183 016743 233 1,5346 34 33t2 4 U.3810 1 4 0.4171 184 0-60.03 23 15818 . 3 0.3370 85 0.3H83 135 0.4796 185 0.6876 235 1.6324 36 0.3378 86 063396 136 0.412i 186 0'69452361.6866 37 81 0910 137 0*484S 187 0.7016 237 1.7451 38 0.3314 88 0.3113 13A 0.4874 188 0.7089 238 1.8081 39 C.3401 P9 0. 11 134 0,4101 189 0.7164 239 1.8764 40 0.3411 90 0.31r51 14r 0.4928 190 0.7241 240 1.9505 41 0.34190.3965 141 0.4956 191 0.7319 241 2.C313 42 0.1428 92 0.39V) 112 0.4984 162 0.7401 242 2.1197 ~3 ~3 0.3993 143 0.5013 193 0.7484 2432168 46 50 94 0.4008 145 0.507 165 0.5790 245 2.940 4b Oo~43 96 0.4037 146 0.5101 196 0.749 246 2.5758 4 91___ 0..462 117 0.513 197 0.7843 247 2.1256 60.34,0 98 0.46 148 0.5163 198 0.59 248 2.8937 StA ₇₂₉₀ 71 _{9 0.314} 0.313 ₁₇ 1L/9 ₀₄₇ ).45195 199 _1T 0.8039 _{6~~16 22~1R}229 3.0862 73.4993735 123 0.4522 1C3 0.142 250 3.1077

Figure 7. Wavelength as a function of vidicon column

for a typical calibration function. The three column(Sn)-wavelergth(Ln) tpairs used to determine the function are given at the top.

Columxi number is at the left, wavelength at

page 19

300

280

260 240 220

z

200

0

o

180

z

0

U 160 W 140

Z

120

, 100

80

60

40

20 0 0.3 0.4 0.5 0.6 0.7 0.8 0.9 1.0 1.1 1.2 WAVELENGTH IN MICRONS

Figure 8. Spectrometer dispersion function.

spectral. resolution per image element

as a function of wavelength.

page 20

has been written which runs as a subroutine under the Planetary

Astronomy Laboratory's image processing system (DIPSYS) which has

been set up to provide a metastructure under which vidicon images

may be easily processed.

A simplified diagram

of this program

appears in Figure

9.

The spectral image is read off the run tape by DIPSYS and stored on a disk where it is available to the spectral processing routine, which has three basic tasks. The first and easiest is to punch out the intensities along one row of the image onto computer cards for input into a plotting routine

(this was how Figure 6 was produced). Second, it can subtract the average background from the image, column by column, where the rows over which the background is to be averaged are read from the input instruction cards. Last, and most important, the program can produce a new 'image in which all of the elements have the same

spectral resolution. For spectral reflectivity work, where the range of interest is 0.4 to 1.2 microns,' a resolution of 250 angstroms, the best resolution at 1.2 microns, was chosen. Figures 10

and

11 show the effects of this processing on an image of the standard star Xi 2 Ceti. Portions of these images are then integrated spatially along the slit. Due to atmospheric and telescope optical effects, a star image is not a point; it is smeared out spatialy into to a Gaussian distribution of intensity which is at its maximum where the point source would be. To use the full energy output of the star at a given wavlength, the image must be integrated across all rows where the image intensity is

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Figure 11. A portion of the processed image of 42 Ceti from 0. 35S to

0

_.

7 0

_p.

Each image element represents intensity per ten angstroms averaged

page 24

above the background.

_{After this integration, the spectrum vector}

is punched out onto cards for plotting and further processing.

A

more advanced version of this processor will

incorporate

the:'

plotting, ratioing, and other functions into one

OIPSYS

subsyste~m.

where only disk files will be used.

The final

_{procedure needed for good spectral}

_{ref lectivi t}

_y

data of the surface of a planet

is to know from what

part

of

the

surface the spectrum originates.

A photograph is taken through

the eyepiece, loeking at the slit in a mirror tilted 45 clegrees to

the

optical

_{axis of the telescope (the first surface in Figure}

3).

A similar logging arrangement is

used for photometer

data.

A

plotting program has been written to create Calcomp plots of the

coordinate grid of Mars (or any other planet) projected onto a

disk using the physical ephemeris of the planet from

The

American

Ephemeris and Nautical Almanac and the time of observation

in

Universal Time.

Figure 12

is a block diagram of the program,

while Figure 13 is a typical,

although smaller

than normal,

output.

To position the spectrometer slit on the ditk

of

.the

planet,

the

_{negative of the photograph of the telescope image}

is

projected onto the

'grid,

and the

slit

marked by hand.

At this

point

the -original

vidicon images have been reduced to constant

resolution spectra of stars and known positions on Mars;

and

reduction to spectral reflectivity data,

as well as testing, can.

PHYSICAL

EPHEMERIS

(1 MONTH)

CREATE TABLE OF

SUB-EARTH POINT

COG~RDINATES AND

APPARENT DIAMETER

RUN CAR D:

day of month

UT start of run

UT end of run

Lcaption

Figure 12.

Flowchart of grid

plotting program.

ADD HOURLY CHANGE

OF PARAMETERS TO

TABLE

CA LCULATE POSITION

LOF SUB-EARTH POINT

CALCULATE NUMBER

OF MAPS TO BE DONE

(one per five minutes)

DRAW LINES OF LATITUDE

in spherical coordinates,hold

$ constant and rotate in

6.

transform so that sub-earth

point is at 0=900,

6=00

and

project to plane.

DRAW LINES OF LONGITUDE

in spherical coordinates, hold

6

constant and rotate in $.

transform as with latitude lines.

WRITE CAPTION

object, time, sub-earth point,

apparent diameter in arc-seconds.

LOG MAP and INCREMENT TIME

GET NEW RUN CARD and

CHECK FOR LAST INPUT

MRHS

VIDSP

4

0

T

L

L

D

EC

C

OF

4

CT. 18, 1973

=

8:

58 UT

RT= -17.3

ONG= 9.

8

IA=

21.46

SEC

Figure 13.

A typical grid plot produced by the program in Figure

the third produced for vidicon spectrometer Mars run

12, )

page 27

IV. Analysis of Data

The first major attempt

to use the vidicon spectrometer to

take spectra for reflectivity work occurred during the opposition

of Mars during October, 1973.

On two consecutive nights the

Mauna

Kea

eighty-inch reflector was trained on

the planet Mars, and

about 75 spectra were taken, as well as an equal number of spectra

of the standard stars Alpha Lyra and Xi 2 Ceti.

Xi 2 Ceti

was

chosen because

it

was near

Mars

in the sky,

while

Alpha Lyra has a

spectrum which is well known and

is

used to calculate planet/sun

ratios to get reflectivity.

Figure 14 demonstrates the reduction

methods used to get spectral reflectivities from raw intensity

spectra.

To avoid airmass reductions, spectra of Alpha Lyra and

Xi

2 Ceti were taken when the two stars were at

-the

same airma!-s,

1.38.

Since star/star ratios exhibit

little

variation with

loai

airmass changes, the ratio of the two stars obtained from these

spectra can also be used to reduce reflectivities

-it oth-r

airmasses.

Before any data was reduced to reflectivitin-.,

extensive testing was done to see whether the

data

would

be

usable. This portion of the thesis will describe that work, usinV;

the best results obtained to date.

Figure

15

shows

a

high resolution spectrum of Alpha Lyrat

which has been averaged over

258

angstrom segments

to

simulate the

spectrometer output.

Figure 16 is an Alpha Lyra spectrum from the

page 28

AIR MASS CORRECTED SPECTRA

Production of spectral reflectivity from

raw spectra. Air mass correction not

not needed if objects to be ratioed are at

the same air mass.

page 29

+

+

wlfs

+

+

8+

.0;

_>

:

₊₀

+

+

+++

+

+

+

WAVELENGTH IN

MICRONS

Figure 15. Spectrum of aLyra, averaged over

250 angstrom resolution elements, from

a 50

angstrom

resolution spectrum provided

in

.W) (1>

0.40

0.

b

WR,-VELFNG1 T I

IN

MrI6h.C

Figure 16. aLyra spectrum from vidicon spectrometer with vidicon response (Figure 2) divided out.

page 30

Crn

page 31

removed. Note that the peak is shifted to a slightly longer wavelength and that the shape is generally broader to about 0.7 microns. To test the repeatability of the data, pairs of spectra of the same star were ratioed to each other. Results of one such pair are shown in Figure 17 (all ratios plotted are normalized to 1.0 at 0.S75 microns). Figure 17a is the ratio of two Alpha Lyra spectra with similar airmasses (1.40/1.38), but different exposure times (Ssec/lsec). If the response of the system were perfectly linear, that is, if intensity recorded from a given source is a linear function of the integration (exposure) time, the curve would be. flat. It is obvious that it is not; however, .the relatively flat region corresponds with the peak intensities of the spectra, so it may be that low level signals are nonlinear representations of the intensity received from the star. To test this idea, a 'pedestal' was set up under the spectrum.. All intensities below a certain value would be ignored, and possibly, the nonlinear features of the curve would go away. Figures 17b and 17c show the results of installing pedestals of 300 and 400. respectively (the maximum intensity registerable is 403S). a p)edestal of 300 seems to help from 0.5 to 0.8 microns, but a lIarger pedestal doesn't help at all. Figure 18 shows a similar ratio for two Xi 2 Ceti spectra with slightly different airmaq'es (1.67/1.32) and different exposure times (20sec/15reec). Once again the curve is relatively flat over the peak in incoming energy, this time from almost 0.4 to 0.8 microns. (Figure 19 is a

++++++++44.4

4.4.4.

4.

4.

0.40

_{U U}

O'.EENG0 .60

0O.80N

WAVELENGTH IN MICRONS

SRLTR86

/

SRLYR83

Figure 17a. Ratio of two aLyra spectra, all elements

above background included.

page 32

4.

+

C,,

z

U.'

4.

4.

4.

++

.20

1.00

1.20

page 33

.g.~.g.

4

~.g*

4.4.

+

4.

+

4.

4.

+

o.40

00so

08o.

1.20

WAVELENGTH IN

MICRONS

SRLTR86

/ SRLTR83

Figure 17b. Ratio of same two aLyra spectra,

this time including no elements less than 300.

CD I

U,

z

LU

~6. 'rn-Go

81

_Sf1

bko

-0-4.

4.

~4*

- EEG I M6icR

WAVELENGTH IN MICRI

SALYR86

1.00 '

1.20

SRLTR83

Figure 17c. Ratio of same two a Lyra spectra,

this time including no elements less than 490.

page 34

co₀

RI-.

zc

C-~

B'

"b2

- ---- - - - -

----page 35

4~

4+

4.

.4.

44444444+4.

U44-0WAVE

'.GT

N C'.S

WRVELENGTH IN MICRONS

SXCT112

SXCT124

Figure 18. Ratio of two t2 Ceti spectra, including

all image elements above background.

b-n

z

wi

PI

Cie

81

,b'.a0

page 3G

typical Xi 2 Ceti spectrum). this time, however, there is a smooth upturn which has some undetermined significance. Thut . star ratios seem to be usable, at best, from 0.4 to 0.8 microns.

Now that there is some idea as to the reliability range of the spectrometer, indefinite though it may be, the Mars spectra can be observed. Figure 20 is a typical Mars spectrum, summed over five vidicon elements down the slit. Note that the peak is

in the red, rather than the blue like the two stars' spectra. This is because the stars are both of spectral type AO, while the

sun, which is providing the light which is ref lected from Mars is a cooler, redder type G. Figure 21 shows a saturated spectrum of Mars. 'The peak intensity of 4095 is surpassed from 0.5 to 1.0 microns, although around 1.1 microns, the signal is unsaturated. Originally it was thought that the unsaturated portions of a saturated spectrum could be used to extend the range of an unsaturated spectrum which had a very low signal beyond 1.1 microns. The data show, unluckily, that there is little or no overlap between the good signal from one and the good signal from

the other type of spectrum. Once again, an attempt was made to do away with low, nonlinear signals with a pedestal. Figures 22a.b,and c show the progressive changes as pedestals of 300 and 400 are subtracted from the original spectrum. Ratios of Mars images seem to be more consistent than those of star images. Figures 23a,band c and 24a,band c are the results of ratioing different 'images of Mars to each other. The three images used

page 37 4.

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Figure 19. A typical t2 Ceti spectrum. Note that the peak is at a longer wavelength and the shape is broader than a Lyra.

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Figure 20. A typical Mars spectrum

page 38

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