The Present and Future of the Telescope of Moderate Size [Reprint 2016 ed.] 9781512809343

Table of contents :
Preface
Contents
1. Image Tube Developments and the Small Telescope
2. Electronic Photography
3. Photoelectronic Problems in Astronomy
4. Investigations of Image Detectors
5. The Application of Punched-Card Methods to the Recording and Reduction of Photoelectric Observations
6. The Newton Lacy Pierce Photometer: A Photoelectric Photometer Designed for Variable Star Observations
7. An Infrared Technique for Stellar Photometry
8. The Application of Small Telescopes to Photoelectric Problems
9. Photoelectric Studies of the Scintillation of Starlight
10 . Our Knowledge of the Upper Atmosphere from Studies of the Scintillation of Visible Starlight
11. Precision Problems in Photographic Astrometry
12. Some Future Problems in Astrometry
13. Variable Star Programs, Present and Future
14. The Present and Future of Stellar Spectroscopy with Moderate-Size Telescopes

Citation preview

The Present and Future

of the Telescope of Moderate Size

Symposium Papers Presented at the Dedication of the Flower and Cook Observatory of the University of Pennsylvania June 11 and 12, 1956

The Present and Future of the

Telescope of Moderate Size Edited, by

FRANK BRADSHAW WOOD

Philadelphia

UNIVERSITY OF PENNSYLVANIA PRESS

© 1958 by the Trustees of the University of Pennsylvania Published in Great Britain, India, and Pakistan by the Oxford University Press London, Bombay, and Karachi Library of Congress Catalogue Card Number: 57-9085

Printed in the United States of America American Book-Stratford Press, Inc., New York

Preface This collection of papers represents contributions to a symposium held in connection with the dedication ceremonies of the Flower and Cook Observatory of the University of Pennsylvania. This new installation combines the facilities of the old Flower Observatory and the old Cook Observatory at a new site well removed from city disturbances. It is designed to complement the other astronomical facilities on the campus. At first, it seemed that it might be difficult to find a suitable symposium topic, since many excellent symposia on a wide range of topics have been held in various places in recent years. But, as we discussed the matter, we began to realize that one subject in particular might be especially appropriate for this occasion. We felt that while very properly much attention has been given in recent years to Schmidt type telescopes, radio telescopes, and very large instruments, much of the useful work has been and is being carried out by conventional telescopes of moderate size. For this purpose, we rather loosely consider moderate size as roughly from twelve to forty inches aperture. Especially in the fields of astrometry and photometry, a rather large fraction of the observations are being made with telescopes in this aperture range. A second reason for this particular topic is found in recent developments in the field of electronics, which are bringing within the range of these instruments a whole host of problems that could previously have been attacked only by the largest instruments. Thus it appears that, useful 5

6

Preface

as they are at present, the future accomplishments of telescopes of moderate size will be tremendously greater. Therefore, we have an opportunity to consider both what is being accomplished and what we may confidently expect to be able to do in the predictable future. T h e latter consideration is obviously of prime importance, since the chief purpose of a symposium is to look to the future—but this can scarcely be done without knowledge of what is being accomplished at present. I want to express the gratitude of the entire staff both to the authors who contributed so generously and to the other astronomers who joined us and who shared in the discussion. T h e work of M. Lallemand was saved from the perils of my own translation by the almost complete revision of this at the competent hands of Mr. E. G. Reuning. Dr. L. Binnendijk has kindly agreed to edit the final manuscripts in my absence from the observatory during the academic year 1957-58. T h e lectures were sponsored by the Tobias Wagner fund of the University of Pennsylvania. F R A N K BRADSHAW

University of Pennsylvania July 1956

WOOD

Contents

Preface

5

1. W . A . H I L T N E R :

Image T u b e Developments a n d

the Small Telescope 2. A . LALLEMAND:

11

Electronic Photography

3. J. D. M C G E E : Photoelectronic Problems in Astronomy

25

31

Investigations of Image Detectors

51

5. J. S. H A L L and A. A. H O A G : T h e Application of Punched-Card Methods to the Recording and R e d u c t i o n of Photoelectric Observations

87

4 . P . FELLGETT:

6.

T h e Newton Lacy Pierce Photometer: A Photoelectric Photometer Designed for Variable Star Observations

95

and G . R . M I C Z A I K A : An Infrared Technique for Stellar Photometry

111

T h e Application of Small Telescopes to Photoelectric Problems

129

9. W . M. P R O T H E R O E : Photoelectric Studies of the Scintillation of Starlight

141

W . BLITZSTEIN:

7. D . J . LOVELL

8.

10.

11.

G . E . KRON:

G.

P.

O u r Knowledge of U p p e r Atmosp h e r e f r o m Studies of the Scintillation of Visible Starlight

155

VAN DE K A M P : Precision Problems in Photographic Astrometry

173

KELLER:

7

8

Contents PACE

12.

Α.

Ν.

VYSSOTSKY:

Some Future Problems in

Astrometry

189

13. B. S. WHITNEY: Variable Star Programs, Present and Future

199

14. D. B. MCLAUGHLIN: The Present and Future of Stellar Spectroscopy with Moderate-Size Telescopes

205

T h e Present and Future

of the Telescope of Moderate Size

I .

Image Tube Developments and the Small Telescope W. A. HILTNER YERK.ES OBSERVATORY, T H E UNIVERSITY O F CHICAGO

Most astronomers are faced with the problem of having a telescope smaller than the largest. However, this should not deter us in any way, but only encourage us to sharpen our pencils, either to develop and use instruments of the highest efficiency or to do theory. For example, last winter very faint stars, approximately 22nd magnitude, in globular cluster M3 were observed with the 82-inch telescope. Now this sounds like a problem for the 200-inch telescope. However, by employing a photomultiplier with a very efficient photocathode made by Dr. Lallemand we were able to obtain information as rapidly as currently obtained with the 200-inch telescope. In fact, judging from the literature, the 82-inch was 10 per cent faster. What we lacked in aperture we made up in quantum efficiency of the photocathode. On the next run we hope to double the efficiency by employing two cells that will alternately be switched from sky to star plus sky. W e could continue the installation of Lallemand tubes so that finally we arrive at a "500-inch" telescope for three-color photoelectric photometry by employing six tubes. T h e more efficient instrument is not the only way that the astronomer can overcome the handicaps (if we should call them that) of small apertures. Since the photoelectric 11

12

The Present

and Future

of the Telescope

of Moderate

Size

photometry of faint stars is one of accumulating statistical data until we have sufficient to give us the desired photometric accuracy, it is obvious that we can compensate for a deficiency in aperture by working a little longer. T h e principal condition placed on the telescope is one of sufficient mechanical perfection so that one can employ the smallest desirable diaphragms and increase the integration or accumulation times by the desired amounts. Of course, one can also apply the same auxiliary efficient equipment and long integration times to the largest telescopes, but such developments do place the smaller instruments in the range of many problems now on the frontier of astronomical research. T h u s far I have discussed photoelectric photometry to illustrate that the man with a "smaller" aperture can truly be "in business." No mention has been made of photography with these telescopes of lesser dimensions. Today, if one wishes to obtain a high dispersion spectrogram of a faint star or to photograph a 22nd-magnitude star, one immediately thinks of the largest apertures. T h i s is so for a variety of reasons. First of all, of course, only the largest telescopes have spectrographs with which high dispersion spectrograms can be secured. T h e two really important factors are those of quantum efficiency and reciprocity failure of the present photographic emulsions. T h e efficiency of image formation by photographic emulsions when exposed to photons is somewhere between 0.1 per cent and 1 per cent and the efficiency decreases with decreasing radiation until finally it becomes zero. Consequently, the astronomer with the small telescope is unable always to increase the integration time to counteract the deficiency in aperture. However, with new techniques now under development, it appears that the efficiency of photo-

Image Tube Developments

and the Small Telescope

13

graphic emulsions or its equivalent will be effectively increased by nearly two orders of magnitude and the reciprocity failure will vanish. Here again, one sees the application of the old adage that necessity is the mother of invention. Astronomers with smaller telescopes attacked the problem of more efficient image formation, b r . Lallemand (1936) defined the problem two decades ago and has pursued a course leading to a solution of the problem he originally defined. With his leadership others have joined the effort. In this report I wish briefly to discuss the work that has been in progress at Yerkes Observatory for the past few years. After attempts with direct exposure of the photographic emulsion to energetic photoelectrons with the subsequent rapid destruction of the photocathode, after an attempt to use ordinary image converters with no significant gain in sensitivity, after attempts to introduce electron amplification within an image converter by secondary emission with disappointing results, after some experiments with charge integration methods, we returned to the original scheme of exposing the photographic emulsion directly with accelerated photoelectrons. But this time we isolated the emulsion chamber from the photocathodc chamber by a thin metallic foil. Figure 1 illustrates this. As many of you are aware, the principal handicap in the development of an image converter in which the emulsion is exposed by accelerated photoelectrons has been that of rapid photocathode decay caused by the chemical reaction of the cesium in the photocathode with the water in the emulsion. With the introduction of a thin layer of metal transparent to energetic electrons, but impervious to low-energy molecules, the cathode decay should be retarded by many orders of magnitude.

14

The Present

and Future




LIGHT

PHOTOCATHOOE

ALUMINUM

FOIL

\

PHOTOGRAPHIC EMULSION

Figure 1. Schematic diagram of an image converter with a thin aluminum foil to retard the chemical reaction of the photocathode by gases evolved from the emulsion. In the employment of a thin foil as a molecular barrier, two conditions must, of course, be met. First, it must prevent decay of the photocathode and second, the angular scattering and absorption of electrons by the foil must be sufficiently small to prevent serious deterioration of the image. T h e rate at which gas diffuses in atm. cm 3 per cm 2 per sec is given approximately by the relationship . T i m e

exposinc

to

shot

(lulu-rent

Magnitudes.

required

noise,

t·,)

to \ei ι icU- l i m i t a t i o n

Instrument

(amplilier)

due

noise.

Photoelectronic

Problems in Astronomy

45

T h i s time may be reduced by: 1. Better photocathode. 2. Higher multiplication of the primary photoelectrons. 3. R e d u c t i o n of amplifier noise. T h e star image signals will also be subject to random fluctuations, but this is a second order effect not very important at this stage of development. T h e r e will also be some dark current from the photocathode that will add to the background. It is known, however, that the type of photocathode most likely to be used (the transparent Sb—Cs surface) has a thermal dark current of < 10 4 £/sec/cm 2 at room temperature. T h i s is very small compared with the current due to sky background of « 4 χ 10 r , e/sec/cm 2 and hence for this problem it can be neglected. If, however, infrared sensitivity is required and a photocathode of the Ag—O—Cs type is used, the dark current will almost certainly be an important factor, and for best results it would be necessary to refrigerate the photocathode. I f the values of t1 and t2 calculated from equations ( I ) and (2) above are plotted against values of star magnitude as in Figure 5 we find the interesting result that for stars brighter than about magnitude 23.5 the time of integration required is determined by t2, the time required for the star signal image to override the instrument noise. For stars fainter than magnitude 23.5 the time is determined by tiy the time required for the integrated star image signal to override the background noise. T h e time tx can be shortened only by more efficient conversion of photons to electrons, b u t t2 may be reduced by two other means: the reduction in amplifier (or instrument) noise and by multiplication of the primary photoelectrons. T h e r e is hope of obtaining a considerable improvement in photocathode efficiency by using the composite alkali

46

The Present and Future

of the Telescope

of Moderate

Size

cathode recently developed by Sommer (1955). It is doubtful if much improvement in amplifier noise can be expected in the near future and it does not seem that much advantage could be gained by adopting a multiplier signal output, such as is used in the image orthicon (Rose, et al., 1946) since this method of signal generation loses its advantage over the amplifier output method as the rate of scanning decreases. However, the primary photoelectrons may be multiplied or intensified in several ways, and one attractive possibility is to use bombardment-induced conductivity (B.I.C.). T o use this eifect the insulating capacitances on the metal target in Figure 3 are thin layers of one of the materials that exhibit this phenomenon, M g F 2 , A 1 2 0 3 , Sb2S 3 , etc. T h e primary photoelectrons must be projected with high energy, at least ten kilovolts, and preferably fifteen to twenty kilovolts, onto and through this insulating layer. T h e metal signal plate is held at a positive potential relative to the potential at which the free surface of the insulator is stabilized by the scanning beam. H e n c e while the insulator is made conducting momentarily at the point where a primary electron has passed through it, a very much larger positive charge can be made to flow to the free surface of the insulating target. T h e multiplication factor obtained in this way (Ansbacher and Ehrenberg, 1951) can be ' le, and probably two, orders of magnitude greater than is practicable using secondary emission. T h e curve for t2 in Figure 4 becomes the dotted line. Now t1 becomes the limiting factor for stars fainter than the 21st magnitude. A serious difficulty has stood in the way of using B . I . C . in this way in the past: namely, it is the difficulty of using sufficiently high accelerating voltages on the image section of the tube shown in Figure 3 to achieve effective penetra-

Photoelectronic

Problems

in Astronomy

47

tion of the insulating layer that forms the storage capacitance. It appears now that this difficulty is removed by the use of the reversible photocathode described in Section 2. It is necessary to take care that the storage capacitance is not overloaded so that the integration of the charge image becomes nonlinear. This danger is increased by high multiplication of the primary electrons. For example, if we assume that we can store 10® electron charges per picture point and that each primary electron is multiplied by a factor of five, then the background will begin to saturate the target after 5 Χ 104 seconds, or fourteen hours, at the stage when 27th-magnitude stars should be becoming clearly distinguishable. If however, very much greater multiplication, say by a factor of 100, were attainable then departure from linearity would set in after 250 seconds when stars of only a little fainter than the 24th magnitude are distinguishable. Hence it may be desirable under some conditions to keep the electron multiplication to a value approaching unity. On the other hand, where the sky background is not the limiting factor and it is desired to record either very faint images, such as spectra with long exposures, or bright images, such as of planets with a short exposure time, then a high multiplication factor of several hundred times would be required to achieve the limiting performance where the picture quality is determined by the shot noise of the primary photoelectrons. In this way such a device would reach the same limiting sensitivity as the image converter. T h e signals reproduced by such a device would appear ideally somewhat like that of Figure 2, only gTeatly exaggerated. It is clear that the uniform background signal is of no interest and it can be cut off just below the noise fluctuations. Only that part of the signal conveying useful

48

The Present

and Future

of the Telescope

of Moderate

Size

information would then be displayed and recorded, thus enhancing the contrast. However, it is to be expected from experience with television pick-up tubes that the response of the photocathode and the multiplication mechanism will not be uniform over the whole image. This would make the recording or display, as envisaged above, very inefficient. However, it would be practicable to compensate this effect by deriving the characteristic signals produced by the device when exposed to a very uniform artificial illumination. These signals could be recorded and reproduced to operate an automatic gain control to correct the actual signals to what they would be if the response of the device were uniform. T h e device described aims at integrating as much information in one tube as possible. There seems to be no theoretical reason why two or more such devices should not be used in cascade, the second integrating the corrected signals produced by the first, and so on. Remembering that a 27th-magnitude star image is only 1 per cent of the brightness of sky background, it would be possible to discard 99 per cent of the useless information integrated on the first tube and then record one hundred such corrected signals on a second tube. T h e practical limit would appear to be determined by the possible exposure time on a given image.

REFERENCES

Ansbacher and Ehrenberg, Proc. London Phys. Soc., 64, 362, 1951. Baum, W. Α., Sky & Telescope, 14, 264, April, 1955; also A. ]., 59, 422, 1954. Baum, W. Α., Hall, J., Symposium on Photoelectric Image Tubes, I.A.U., Dublin, August, 1955.

Photoelectronic

Problems in Astronomy

49

Burns, Jay, and Hiltner, W. Α., Ap. J., 121, 772-773, 1955. James, I. J. P., Proc. I.F..E. (London), 99, 796, 1952. Lallemand, Α., Comptes Rendus, 203, 243 8c 900, 1936. — Comptes Rendus, 235, 503, 1952. Mandel, L., Jour, of Scientific Instr., 32, 405, 1955. McGee, J. D„ Proc. I.E.E. (London), 97, 1950. Jour, of the Royal Soc. of Arts, No. 4869, P. 343, March 1952. , Symposium of Astronomical Optics, Manchester, April, 1955. "Photoelectric Aids in Astronomy." Miller, R. H., Hiltner, W. Α., and Burns, Jay, Ap. J., 123, 368, 1956. Rose, Α., Weimer, P. K„ a n d Law, Η. B„ Proc. I.R.E., 34, 425, 1946. Sommer, Α., Rev. Sei. Insts. 26, 725, 1955.

4Investigations of Image Detectors PETER

FELLGETT

CAMBRIDGE UNIVERSITY

OBSERVATORY

T h e use that had been made of the photoelectric effect in astronomy prior to 1950 had already shown that the high q u a n t u m efficiency, freedom from early overload, linearity, and simplicity of response of photoelectric surfaces gave them a special significance as primary receivers of radiation. It was also clear that these advantages could not be fully exploited with single detectors and that the introduction of photoelectric image detectors into astronomy would have far-reaching consequences. As is well known, M. Lallemand and Dr. H i l t n e r have done pioneer work toward this end. T h e history of the work at Cambridge has already been outlined (Fellgett, 1955a, 1956) and the aim of the present account is to supplement and bring u p to date the earlier reports, at the same time correcting some of the conclusions previously drawn. W e shall be concerned here with photoelec trie image devices only as primary detectors of radiation, and not with their many other interesting applications such as the determination of isophotes of photographs, the control of amplitude distortion, the modification of image bandwidth, and the presentation of small differences between two photographs. T h e discussion begins (§1) by referring to the choice of the "equivalent q u a n t u m efficiency" as a measure of sensitivity. T h e application of this measure to observational 51

52

The

Present

and

Future

of the

Telescope

of Moderate

Size

problems is illustrated in §1.1, and in §1.2 the importance of the design of observations for high informational efficiency is emphasized. Some estimates of the value of the equivalent quantum efficiency for photographic emulsions, the eye, and photoelectric devices are given in §2. Astronomical tests of image tubes at the telescope are reported in §3.1, and in §3.2 these and other direct comparisons between different kinds of detectors are discussed. One result may be mentioned at this stage. We find that the peak monochromatic equivalent quantum efficiency of photographic emulsions can be around 1 per cent, and perhaps as high as several per cent. T h i s estimate is uncertain inasmuch as the data upon which it is based are not homogeneous, and it has been reached only after some hesitation (Fellgett, 1955a, 1956). If true, it indicates that the photographic process is about ten times more efficient, both absolutely and in comparison with the photoelectric effect, than has often been suggested, and that the possible improvements in astronomical sensitivity are more limited and more difficult to attain than had been hoped. It is accordingly a matter of some urgency that a definitive value should be obtained.* This illustrates the more general conclusion that the critical evaluation of the performance of existing detectors, both absolutely and in comparison with each other, is an essential step in estimating the gains that may be attained in the future, and it gives useful indications of the way in which improvements are most likely to come about. Available detectors may be classified into those dependent • Homogeneous data on certain emulsions have recently become available, and the best equivalent quantum efficiencies calculated from these data are around 1% in agreement with the conclusions of the present discussion. It is intended to submit these calculations for publication in Monthly Notices.

Investigations

of Image Detectors

53

o n vacuum electronics a n d those that use the p h e n o m e n a of solid-state physics. P h o t o g r a p h i c emulsions belong to the latter class; m a j o r advances are being made in the study o f the underlying physical processes and it is probable that increases in p e r f o r m a n c e will eventually result from this. Similarly, although very little is known so far a b o u t their signal-to-noise properties, it seems q u i t e likely that solidstate image intensifiers will make rapid advances. It may well be difficult for electronic devices to do m o r e than match these developments as far as sensitivity is c o n c e r n e d ; b u t of all image detectors so far envisaged they have the advantage, which may prove crucial, of being able directly to generate electrical signals that in principle may b e accurately p r o p o r t i o n a l to the i l l u m i n a t i o n . Developments that can at present be foreseen may b r i n g photoelectric q u a n t u m efficiencies to perhaps 0.4. If this could be c o m b i n e d with almost perfect evaluation of the photocurrent d i s t r i b u t i o n and substantially c o m p l e t e storage over times of the o r d e r of an hour, it would obviously allow the efficiency of astronomical observations to b e markedly increased. F o r example, the resultant improvem e n t in the faint limit of a given telescope has been widely discussed. T h e r e are good grounds for h o p i n g that practical detectors with these characteristics can eventually b e produced, and the effort to b r i n g a b o u t this developm e n t is highly w o r t h w h i l e . Nevertheless, the r e q u i r e d performance represents a h i g h e r cathode efficiency than has been achieved in the c o m m e r c i a l production of simple phototubes, b e t t e r signal-current evaluation than has b e e n achieved in television c a m e r a tubes, and b e t t e r storage than has been achieved in m e m o r y tubes intended for digital computers. T h i s suggests that, so far as signalgenerating tubes of the usual kind are concerned, a m a j o r

54

The Present and Future of the Telescope of Moderate Size

problem in engineering research is involved, and quick results are unlikely. Where signal generation and the convenience of a permanently sealed-off tube are not deemed essential, image converters may have a more immediate impact on observations. As Lallemand strikingly illustrates elsewhere in this volume, the photocathode-to-emulsion image converter can have very high performance (including high signal storage), and it is applicable in the field of spectrography where low efficiency is very time-consuming. Some of the advantages hoped for may be obtainable by simpler means than by the development of new detectors. Electronics includes theoretical ideas as well as practical techniques, and the now well-known relationships between signal-to-noise ratio, photon rates, detector efficiency, and signal storage suggest ways of adapting photographic techniques to particular observational requirements. T o revert to the example already mentioned, the degree of signal storage can be increased, with a corresponding improvement in the limiting magnitude, by averaging the signals recorded on a number of plates, by the use of auxiliary magnification to increase the plate scale, by the use of fine-grain emulsions, and by using emulsions of high "gamma." T h e first three methods depend on increasing the number of grains in each star image; and the last one (which also provides "background subtraction") on increasing the statistical significance of each grain. If in these ways the useful time of exposure can be increased, without change in efficiency, from (say) a half hour to ten hours, a gain of 1.6 magnitude will result under noise-limited conditions. Moreover, if Figure 2 is typical of emulsions, the efficiency of a plate exposed for maximum storage is only half to a third of its maximum efficiency, so that by exposing for best efficiency instead of

Investigations

of Image

Detectors

55

m a x i m u m storage as at present (and if necessary using more plates) the total gain can be increased to 2.1 magnitude. T h i s gain is limited by the time necessary for the required n u m b e r of photons to be effective, and therefore depends on the efficiency of the detector. According to Figure 3 the efficiencies of fine-grain emulsions are not systematically less than those of "fast" coarse emulsions. Fine grain may be associated with worse low-intensity reciprocity failure, presumably because a demand has not hitherto been made for very fine-grained astro emulsions, but in some cases (§3.2) this can be overcome by cooling the plate during exposure. 1. SENSITIVITY ASSESSMENT

T h e question whether a new type of image detector is more sensitive than those currently used becomes meaningful only when it has been decided what is meant by "sensitivity." It is not self-evident how the sensitivity of a single detector can be most usefully defined, and image detectors present additional problems. Sensitivity as ordinarily used may refer to either or both of the properties known specifically as responsivity and detectivity (§3). Responsivity is the output per input, for example, volts per watt or amperes per lumen. Detectivity measures the smallness of the radiation signal able to override the prevailing noise fluctuations, and is the property that concerns us here. I n the early days of telescope manufacture, performance was expressed in terms of magnification. W h e n almost unlimited magnification became technically feasible, it became appropriate to measure optical performance by means of resolving power; the change is associated with the custom of quoting telescope size in terms of aperture instead of focal length. T h e distinction

56

The Present and Future of the Telescope

of Moderate Size

between responsivity and detectivity is analogous, and electronic methods now e n a b l e any useful a m o u n t o f amplification to be attained. Incidentally, the q u a n t i t y amperes per l u m e n is particularly misleading since it refers to tests with a light source having a steeply rising intensity as a function of wavelength in the visible region. Consequently the response depends m a i n l y on the long-wave tail of the cell sensitivity, a n d this may be negatively correlated with the sensitivity in the wavelength range for which the cell is most valuable. M a n u f a c t u r e r s can h e l p the user very m u c h by listing the q u a n t u m efficiencies of their cathodes. I n previous discussions (Fellgett, 1955a, 1956), the choice has been explained of the equivalent q u a n t u m efficiency £e as a measure of detectivity. T h e basic r e q u i r e m e n t of a radiation detector is that it should make the best possible use of the available photons collected by the

telescope;

m o r e precisely, it should transmit a high p r o p o r t i o n of the relevant i n f o r m a t i o n c o n t e n t of the p h o t o n distribution. I f a given observation requires Q photons to be availa b l e at the detector, and if Q 0 photons would be c a p a b l e if fully used of yielding the same i n f o r m a t i o n with the same coding, then the e q u i v a l e n t q u a n t u m efficiency of the detector is defined as

T h e application of the concept to image detectors is due to R o s e (1948). T h e equivalent q u a n t u m efficiency is generally a function of the intensity, duration, and wavelength of the incid e n t light, and of the contrast and wave n u m b e r of the image detail. Its variation with these parameters expresses

Investigations

of Image Detectors

57

the effects of overloading, reciprocity failure, spectral response, and image resolution. T h e variation of t e with intensity for a number of detectors according to Rose is shown in Figure 1. Because of this variation, the "faint limit" is not a satisfactory measure of sensitivity; the order of merit of two detectors at one brightness level may well be reversed at a different brightness. Rose emphasizes that his estimates of the absolute levels of ε„ may need revision (see §2), but the shapes of the curves should be about right. T h e eye is remarkable for the very wide range of intensity over which εβ remains relatively close to its maximum value. Photographic emulsions are remarkable for the very sharp dependence of their ee on exposure. Image-orthicons are intermediate in this respect between the eye and photography. 1.1 Application to observational problems; the cost of astronomical information. Photons and the information they convey can be expensive. A telescope of 100-centimeter aperture costs at present around $10 5 . If depreciation and interest are reckoned on a twenty-year basis, the capital cost per hour of operation is therefore about $5 at a site yielding a thousand observing hours per year. A British astronomer receives about $1 per hour. Overhead and maintenance probably amount to a further $2 per hour of operation. T h e total cost may therefore be estimated to be about 13 cents per minute. T h e rule-of-thumb photon rate in the visible spectrum at the earth's surface is 10" sec"1 cm 2 for a star of zero magnitude. Hence a 100-centimeter telescope collects 500 photons per sec = 30,000 per min from a star of 18th magnitude. Let χ photons be available, and let them be divided into

58

The Present

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of Moderate

Size

Ν statistically equivalent groups. In each group the signalto-noise power ratio is x / N , and the total information capacity is (1.11) bits. T h i s is maximum when Ν is large Η (N

oo, χ) = χ log2e — 1·44 χ

(1.12)

Usually, however, we wish to measure parameters to an accuracy of 1 per cent rms, corresponding to a "peak-topeak" uncertainty of about 6 per c e n t . T h e n , χ / Ν = 104, and H p a

M4 x l n 104

104

_ j.33 iQ-3 x

(1.13)

As usual when coding of the information ahead of the detector is restricted, the information gain from measuring a few parameters with high accuracy is much below the maximum possible gain. From these data it will be seen that in order to measure a spectrum of 1,000 resolved elements to an accuracy of 1 per cent rms, it is necessary to have 1,000 χ 104 = 107 effective photons. For the star and telescope considered, this requires 300 minutes = 5 hours of telescope time costing $40 if every available photon is fully effective. Threecolor photometry on the same basis requires one m i n u t e at a cost of $0.13. In either case, for each cent expended, 2,300 photons are collected and three bits of information gained. By comparison, a high-grade television receiver capable of handling 108 bits per second would cost $33,000 per second to r u n if the cost per bit were the same as for the telescope. If a page of scientific journal contains 50 lines X 15 words χ 5 letters, and the n u m b e r of different "letters" (including spaces and punctuation) is taken as 32,

Investigations

of Image

Detectors

59

then a publication charge of $10.00 per page represents a rate of 10 bits per cent, allowing for the redundancy of the language b u t not for any redundancy in the material that may be submitted. T h e estimate of five hours to record the spectrum of an 18th-magnitude star with a 100-centimeter telescope is increased to fifty hours, costing $400, if the q u a n t u m efficiency of the detector is 0.1. If it were 0.01, the three-color measurement would r e q u i r e nearly two hours and cost $13, equivalent to $0.30 per bit. Since in fact it is d o u b t f u l if the over-all efficiency of a spectograph and emulsion is at present as high as 10Λ it is clear why the spectra of 18thm a g n i t u d e stars are not considered accessible to a telescope of this size. T h e chief purpose of these examples is to illustrate that the practicability, so far as sensitivity is concerned, of any observation is d e t e r m i n e d by the available photon rates a n d detector efficiencies in relation to the kind and a m o u n t of i n f o r m a t i o n that it is r e q u i r e d to obtain. 1.2 Informational efficiency and the design of observations. Detector sensitivity is only a means to the end of gaining information about the sky. Many circumstanc es can cause the rate at which i n f o r m a t i o n is gained to fall below that of which the detector a n d telescope are in principle capable. Astronomical seeing reduces the information content in an image (Linfoot, 1955). Systematic errors reduce the significance of measurements. T h e point-by-point measurement of a stellar spectrum makes use of much less than the full i n f o r m a t i o n capacity of the telescope both because of the statistical correlations along the spectrum and because a great deal is already k n o w n f r o m previous work a b o u t the possible form of the spectrum. T h i s inefficiency is espe-

60

The Present and Future of the Telescope

of Moderate Size

cially m a r k e d when all the details of a spectrum are recorded for the sole purpose of d e t e r m i n i n g i n f o r m a t i o n of low dimensionality, for example radial velocity or spectral class and type. A most striking loss is that shown by the comparison of E q u a t i o n s ( 1 . 1 2 ) and (1.13). It arises f r o m o u r inability to code the i n f o r m a t i o n presented by the sky; and it a m o u n t s to a factor of nearly 1,000. T h e r e is n o k n o w n way of r e d u c i n g the loss in this particular case; b u t it illustrates that i n f o r m a t i o n a l inefficiency may b e greater than the loss d u e to imperfect sensitivity of detectors. T h e practical gains from m a k i n g observations i n f o r m a t i o n a l l y efficient can b e correspondingly large. An e l e m e n t a r y e x a m p l e is the loss when the resolved elements, e i t h e r of a spectrum or of a direct image of an astronomical o b j e c t , are presented to the detector in succession instead o f b e i n g observed continuously t h r o u g h o u t the available e x p o s u r e time. F o r example, a 36-inch telescope engaged in a 3-color photoelectric p h o t o m e t r y and using d i c h r o i c filters o r a polychromator so that the colors are all observed simultaneously (as assumed in §1.1) will have the same o u t p u t as a 60-inch w o r k i n g in the conventional way with a b s o r b i n g filters and measuring the colors o n e at a time. T h e loss in effective speed is still greater when a s p e c t r u m o r an image is scanned instead of b e i n g received on an image detector. O n l y a small part of the bandwidth of a p h o t o e l e c t r i c single d e t e c t o r is used in most astronomical applications, and h e n c e there is the possibility of using a single detector in a m u l t i p l e x a r r a n g e m e n t to receive simultaneously all the e l e m e n t s of a spectrum or image, each e l e m e n t having a distinguishable m o d u l a t i o n pattern impressed on it. I f

Investigations

of Image

Detectors

61

the dark current is negligible, the increase in shot noise resulting from simultaneous illumination by all the resolved elements unfortunately cancels the advantage of multiplex working, and the performance is just the same as for a scanning system. In the infrared, the "dark curr e n t " predominates, but the method still fails to give any advantages for images because with given optics the size of the sensitive area needed is proportional to the n u m b e r of elements received, and the noise power of the detector is also proportional to this area. T h e method succeeds for infrared spectra, since the elements of the spectrum can be modulated by interference without dispersion, or if dispersed can be recombined (Fellgett, 1949, 1952; Mertz, 1953). A multiplex infrared spectrometer for star spectra, initially in the 1 to 3 μ region, is under construction at Cambridge. If the necessary coding is possible, a single detector can be used to abstract a single parameter from a spectrum or image, or a small n u m b e r of parameters that can be multiplexed. T h u s an image detector is not in principle necessary for the measurement of radial velocity or of spectral type. Detailed discussion (Fellgett, 1955a, 1955 b) confirms the possibility of making these measures by means of single detectors although, because of restrictions on the kind of coding that can be obtained, the sensitivity is rather less than could be attained using image detectors of equal efficiency. W i t h either type of detector, considerable gains appear to be possible when the method is designed to avoid the redundant measurement of spectral details known in advance, to make full use of prior knowledge of stellar spectra, and to confine the information gained as far as possible to the actual parameters sought.

62

The Present

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of Moderate

2 . EQUIVALENT Q U A N T U M EFFICIENCIES OF SOME

Size

DETECTORS

In this section, estimates are collected of the equivalent q u a n t u m efficiencies of the h u m a n eye, photographic emulsions and certain photoelectric image detectors. T h e estimates for the eye are taken directly from physiological data (§2.1); those for emulsions are based on makers' sensitometry and granularity measurements (§2.2). In §2.3 the attempt is made to estimate the performance of some photoelectric devices from the efficiencies of their component parts, and some discussion is given of the mechanism of these tubes in relation to astronomical needs. 2.1 The human eye. Rose (1948) (see Fig. 1) gave the white light equivalent q u a n t u m efficiency of the eye as

Figure I. The "white light" equivalent quantum efficiency ε as a function of brightness level, according to Rose (1948).

Investigations

of Image

Detectors

63

about 5 per cent at threshold and 1 per cent at ordinary levels. T h e corresponding peak monochromatic values would be some five times greater. He points out that since his result is based on continuously presented targets, the value of ε6 obtained would be too high if the integration time of the eye, which he takes as 0.2 sec, had been underestimated. It does indeed appear that his values should be revised downward to about the extent implied by taking his curve to refer to peak monochromatic values instead of white light averages. Barlo\v (1956) obtains peak monochromatic values of t e = 5 per cent at threshold, and 1 per cent at one hundredth to one tenth foot lamberts. T h e determination is based on the detectability of bright targets presented for definite times against a uniform background. T h e value quoted is the maximum with respect to target size and duration. Hecht, Shlaer, and Pirenne (1942) find that at threshold under optimum conditions a hundred photons entering the pupil give rise to a frequency of seeing curve that approximates to a Poisson distribution with parameter 5. Hence ee = 5 per cent (monochromatic). Rushton (1955, 1956) finds that about 10 per cent of the photons entering the eye are absorbed by rhodopsin under optimum conditions. There is evidence to suggest (Hagins, 1955) that only half of these may be visually effective. I am grateful to Dr. Barlow, Dr. Pirenne, and Dr. Rushton for discussion, and to Dr. Barlow for having made measurements on my own eye. 2.2 Photographic emulsions. Let QA photons incident on area A cm 2 of a photographic emulsion give rise to an optical density of D loge units in the developed emulsion.

64

The Present and Future of the Telescope

of Moderate

Size

Let the rms fluctuation in the mean density over area A be (ΔΪ) 2 }* : GA %. T h e n the mean square error in the determination of Q from a measurement of D in the area A is (approximately) 5 ? = Ε

where γ = dD/d log e Q is the slope of the characteristic curve when exposure is expressed in photons per cm 2 . If QoA photons were fully effective, the mean square fluctuation divided by the square of the signal would be 1 /QoA. Comparing this with the corresponding relative fluctuation according to Equation (2.21), we find = £e =

(2·22)

If there were Ν developed grains per cm 2 , each of area a and distributed at random, the mean density over area A would be D = Νa

and the

fluctuation AD2 = Ν a?/A

so that G2 =

a2N =

aD

and substitution into the formula (2.22) for ε„ yields 1/2

7

a2NQ

-

v2

'

aDQ

(2.23)

Investigations

of Image Detectors

65

When the density is measured in logio units, this becomes (γ is unaltered if a new logarithmic base is taken for both coordinates)

If the linear resolution of an emulsion depends on the grain size both directly and through turbidity, it may be expected that the resolution R in lines per cm will be proportional to Assuming that R — 500 for a fast emulsion with a mean grain size of 1 μ, we obtain Ra* = 1/20. Substituting into (2.24) yields E e

~170gJ

(2.25)

Neither the actual value of the numerical factor in this formula, nor its constancy, should be taken too seriously, but it does allow emulsions to be compared when the available data on them are insufficient for more accurate methods. It is of interest that the "sensitivity" expressed as the number Q of photons required to give density D occurs only as the first power, while γ and R occur as the square. T h e dependence on R may be understood by remembering that if the resolution of the emulsion is doubled, one half the magnification suffices to give the same resolution on the object, and the intensity at the emulsion is then four times greater. Figure 2 shows the characteristic curve calibrated in terms of photons of a fast blue-sensitive emulsion with a mean grain size of 1 μ2. T h e curve was kindly provided by Dr. G. C. Farnell of Kodak Limited. Mr. B. W. M. Selwyn, of the same company, confirmed that the assumptions made in Equation (2.23) concerning granularity are sufficiently accurate for our purpose in this case. T h e figure shows the

66

The Present

and Future

of the Telescope

">·'

of Moderate

Size

0/«»

Figure 2. Quantum " H and D" characteristic, equivalent quantum efficiency s e , and relative numbers Qt e of quanta effective in the exposure for a fast blue-sensitive emulsion. (Δ in log 10 units, Q in incident quanta/cm 2 at 4300Λ; emulsion developed in D19b for 8 minutes, mean projected grain area = 1.0 μ2.) e q u i v a l e n t q u a n t u m efficiency derived f r o m this f o r m u l a . T h e peak value is εβ = 1.1 per cent a n d the m e a n value over the best 10:1 range of intensity is 0.84 per cent. T h e figure also shows the effective n u m b e r Qe e of p h o t o n s contribu t i n g to the m e a s u r e m e n t of intensity. T h e available photons are most efficiently used at the peak of the εβ curve at .Dio = 0.25 above fog. At higher densities the efficiency falls rapidly, b u t accuracy continues to i m p r o v e u p to the peak of the Qee curve at D10 = 0.85 above fog. Beyond this the fall in efficiency s e m o r e t h a n counterbalances the increased n u m b e r of photons available. T h e most useful r a n g e w o u l d

Investigations

of Image Detectors

67

appear to be from D10 = 0.25 to 0.85 above fog, and this seems to accord with general experience. Figure 3 shows estimates of ε

92

The Present

and Future

o/ the Telescope

of Moderate

Size

T h e advantages of the method lie in its speed, accuracy, and impersonal approach. If after having completed his observations the observer wants to reject all those residuals larger than a given size, he can in advance instruct the 650 to do so before he has had a chance to see whether the results agree or disagree with any preconceived notions he should never have entertained in the first place. At this stage perhaps the most important disadvantage of the method has to do with the procurement of time on a modern computer. Also, the observer must stick to a rigorous routine at the telescope and be willing to expend a certain amount of extra time at the telescope for the transformation of the integrated voltage to a recorder deflection. For the problem outlined above the extra time is 33 per cent. This time is by no means wholly lost because the observer can use it to do such things as checking the centering of the star or hand-punching additional information on the cards. Modern computing machines capable of solving photometric problems are in operation in many large cities in the United States and several hundred more are scheduled ( O N R ) for installation during the next two years. T h e limited experience thus far gained has convinced us that we will gladly sacrifice a certain amount of independence of action that we have previously enjoyed at the telescope to avoid the large expenditure of time required to reduce photoelectric observations.

Punched-Card

Methods

and Photoelectric

Observations

93

REFERENCES H. L. Johnson, A.J., 59, 325, 1954. Built by the Gardiner-Phillipi Corporation of Flagstaff, Arizona. See A. J . Gardiner and H. L. Johnson, Review of Scientific Instruments, 26, 1145, 1955. O N R , Digital Computer Newsletter, 8, No. 3, July, 1956.

6. The Newton Lacy Pierce Photometer: A Photoelectric Photometer Designed for Variable Star Observations WILLIAM BLITZSTEIN F L O W E R A N D COOK O B S E R V A T O R Y , UNIVERSITY OF PENNSYLVANIA

1. INTRODUCTION

T h e Pierce Photometer, Figures 1, 2, and 3, is designed specifically for automatic photometry of variable stars with the aims of minimizing errors resulting from sky transparency variations, of increasing the n u m b e r of observations per unit time, and of recording of data in the most suitable form for analysis. These aims are achieved by the following techniques: (1) simultaneous photometry of variable and comparison star: (2) digital indication of relative light intensities by means of pulse-counting methods; (3) automatic recording of data by printing on paper tape; (4) automatic sequencing of the required operations; and (5) timing in units of decimals of a day. Full use is made in this instrument of the current developments in electronics and automation to achieve the desired results. A description of the history, development, and philosophy of the Pierce Photometer has been given (Wood, 1953a). T h i s paper will be concerned with the description of the instrument and its operation as it exists at present. 95

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The Present

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Size

Very detailed descriptions of the various components will be presented only if they are of unusual design or are required for clarity. 2 . OPERATION

T h e operation of the instrument may be most easily understood by reference to the block diagram of Figure 4. It is assumed that the light from a variable star has been made to fall on the cathode of photocell No. 1 and the light from a nearby comparison star falls on the cathode of photocell No. 2. T h i s is effected by means of the optical system of the Pierce Photometer, which will be described below. T h e output of each multiplier photocell consists of a series of randomly occurring sharp pulses of a wide range in amplitude. T h e average rate of occurrence of these pulses is proportional to the intensity of the illumination falling on the cathode. T h e dark currents of the cathodes and dynodes also contribute a n u m b e r of pulses, which must be accounted for in photometric measurements. T h e s e pulses are all of rather small average amplitudes 10~2 volts) in the usual commercial cells. Amplification is necessary to bring them up to the signal levels required for the operation of conventional amplitude discriminators and scalers. T h i s is accomplished by means of preamplifiers mounted on the photometer head and a dual channel main amplifier, which also performs the functions of amplitude discrimination and gating. T h e discriminator generates a standard pulse for each amplified pulse, which exceeds a definite preset reference level. T h i s serves the following purposes: (1) the signal-to-noise ratio of the system can be effectively increased since the dark current pulses are on the average smaller in amplitude than the pulses due to

l imine I. C . e n e r a l view l t h e 1 ' i e r i e l ' h o t o n i e t e i as u s e d w i t h t h e 111 t e e n - i n c h s i d e n M a l r e l r a c t o r ol' t h e F l o w e r a n d C o o k Ο Ι ) Μ Ί \ a torΛ .

m

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l- i ^ u i x ' ,'i. O p t i c a l

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Pierce

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The Newton Lacy Pierce

Photometer

97

c 3 bo

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bp Ε

98

The Present and Future of the Telescope

of Moderate

Size

photoelectrons and (2) the standard pulses are most suitable for actuating electronic scalers. T h e gates control the output of the dual channel amplifier-discriminator-gate according to the start and stop signals generated by the timing unit. By these means, the electronic counters Nos. 1 & 2, corresponding to variable and comparison star, are actuated by the output pulses of their respective channels for a precisely fixed interval of time. T h e electronic counters store the number of counts during a fixed interval of time, which is preset into the timing unit. These counts are displayed to the observer by means of an indicator in which neon lamps illuminate the proper numerical digits. Also, the electrical information stored in the counters can be made to print the proper digits on the tape of a listing machine by means of solenoids that actuate the keys. T h i s is done according to the "record counts" signal from the timing unit. T h e printing counter serves to store and print the times of the observations. T h e timing unit generates pulses spaced uniformly at time intervals of 10 ® day. These are stored on a mechanical printing counter whose digits can be individually preset to any initial value. T h i s allows timing in decimals of a day with the corrections to the sun and to the middle of the observation automatically applied. T h e time of a given observation is printed at the end of the pulse-counting storage interval in response to the "record time" signal from the timing unit. T h e timing unit controls the timing and sequencing of the instrument. T h e basic clock frequency is generated by an electronically controlled and driven tuning fork of excellent qualities. T h e period of this fork is very closely 10"7 days. T h e basic frequency is scaled down by a factor of 100 to generate timing pulses spaced uniformly at inter-

The Newton 5

Lacy Pierce Photometer

99

vals of 10~ day. Further scaling produces signals at the proper times for: (1) starting and stopping the pulse count storage in the electronic counters; (2) resetting the electronic counters to zero; (3) printing the pulse count stored in the counters and (4) printing the time of the observations. T h e sequence of operations may be broken down into two consecutive alternating periods. During the first interval the amplified and discriminated outputs of the photocells are fed through the gate circuits to the electronic counters. T h i s can be preset to any one of the following: 10, 20, 30, 40, 50, 60, 70, 80, 90 and 100 χ 1 0 5 days. At the end of this interval the printing counter records the corrected time of the observation and the gates shut off the pulse input to the electronic counters. During the second interval the pulse counts stored in the electronic counters are printed on the tape of the listing machine and the counters are reset to zero. T h e following; additional actions planned for this time interval include: (1) printing of a numerical code along with the counts to identify the type of observation recorded; (2) automatic sky background measurement; (3) automatic filter changing and (4) printing of the h o u r angles of variable and comparison stars. T i m e s of 2, 5, or 10 χ 10 5 days are available for these operations. T h e various functions may be initiated manually or the apparatus can r u n automatically according to a preset schedule. In either case all operations are synchronized with the basic timing pulses. It should be noted also that no other operation will occur d u r i n g the time interval of storage of pulses from the photocells. T h i s makes it impossible for pulse interference from these sources to affect the sensitive i n p u t circuits of the amplifiers.

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The Present and Future of the Telescope o/ Moderate

Size

3 . DESCRIPTION OF ELECTRONIC COMPONENTS

a. Cell and Battery Power Supply. T h e multiplier photocells used in the present equipment are R.C.A. Types 1P21 or 1P28. It has been found that some of the 1P28 photocells equal or exceed the average 1P21 in quality. Since they also have the advantage of an envelope that transmits more ultraviolet light and their cost is about one third that of the 1P21 there is an advantage to their use. T h e use of refrigerated end-on cells of high multiplication and low dark emission is planned for the future. In order to increase the signal-to-noise ratio and to minimize the effects of external fields the cells are enclosed in a shield of soft iron that is electrically connected to the cathode. T h i s provides effective electrostatic and magnetic shielding. T h e total voltage applied across the cell is approximately 855 volts. T h e voltage between the anode and the last dynode is forty-five volts. T h e voltages between dynodes are ninety volts. These are supplied by a box containing twenty Burgess T y p e XX30P hearing aid batteries connected with series current limiting resistors as a precaution against dangerous shocks. Both cells are supplied from the same battery thus voltage variations should affect each in approximately the same manner. b. Preamplifier. T h e preamplifier uses a double triode T y p e 12AY7 as a pulse amplifier followed by a cathode follower. T h e 12AY7 was selected for its low input noise and for its freedom from microphonics. T h e main purposes of the preamplifier are to provide some initial voltage gain to raise the signal above possible disturbing pulse pickup and to drive the approximately twenty feet of RG-7/U coaxial cable req u i r e d to bring the signals to the main amplifier.

The Newton

Lacy Pierce Photometer

101

c. Dual Channel Amplifier-Discriminator-Gate. T h e functions of this unit are stable linear amplification of the pulses generated at the anodes of the photocells, amplitude discrimination, and gating of the o u t p u t to the electronic counters. These are accomplished by conventional means. Each pulse amplifier section consists of three triode stages and one pentode used as a gain of minus one phase inverter. T h e triode stages were patterned after those in Handbook of Preferred Circuits (1955), giving sufficient stability, gain, and bandwidth for the purposes of the Pierce Photometer. T h e over-all voltage gain including the preamplifier is approximately one thousand and the bandwidth is about four hundred kilocycles per second. T h e discriminator employs the widely used Schmitt trigger circuit (Elmore and Sands, 1949). A variable bias control is provided for the discriminator of each channel allowing individual setting of the amplitude above which pulses will be counted. T h e gating function is accomplished in the output tubes, which also act as low impedance drivers for the coaxial cables leading to the electronic counters. T h e T y p e 6AS6 tube was selected for this purpose in view of its suitable suppressor control characteristics. By varying the voltage of this element the transconductance of the tube can be increased rapidly from a very low value to the full a m o u n t possible under the existing operating conditions. T h e timing unit controls the voltage of the suppressors in the gates of each channel raising them to approximately cathode potential d u r i n g the interval when pulses are fed to the electronic counters. D u r i n g the interval when the observations are being recorded the suppressors are quite negative with respect to the cathodes and no signals can appear at the output of the amplifier-discriminator-gate.

102

The Present and Future

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of Moderate

Size

d. Electronic Counters. T h e functions of these units are storage of pulse counts and recording of their total at the end of the observation. T h e basic circuit adopted for these purposes is a modification of a decade ring pulse counter designed for the EN I AC (Sharpless, 1948). Five rings of ten are arranged in cascade to give a scale of 100,000. Since the resolving time of these rings is not sufficiently small for the purposes of pulse-counting photometry, a highspeed binary scale of sixteen is inserted between them and the pulse input. T h e n u m b e r of binary stages used can be varied by a switch allowing scaling by various powers of two. T h u s the resolving time is decreased and it is possible to adjust the count indicated to a convenient value. T h e state of the binary scaler is indicated by four miniature neon lamps. T h e E N I A C type of ring will allow remote indication at distances u p to one hundred feet or more. T h e indicator utilizes 100 T y p e NE-2 miniature neon lamps to illuminate numbers showing the count stored in the ring decades of the counters. A multiwire cable and plug connects this unit to each of the electronic counters. T h u s the operation of the instrument can be observed from any convenient position. e. Recording Units. T h e times of observation are printed in decimals of a day on paper adding-machine tape by a Model SCI-10 Printing Counter manufactured by the Streeter-Amet Company of Chicago, Illinois. T h i s device contains an electromagnetically operated mechanical decimal counter, each of whose six digits can be independently set to a predetermined value. This is actuated by pulses from the timing unit so that the least count represents

The Newton

Lacy Pierce

Photometer

103

5

10 day. Printing is accomplished by motor-driven arm, which presses a typewriter ribbon forcibly against the paper tape and the raised digits of the mechanical counter. T h e storage of timing pulses is continuous during the operation of the photometer; the time is printed only after the end of an observation according to the "print time" signal from the timing unit. T h e counts stored in the electronic counters are printed on paper adding-machine tape by a thirteen-place listing machine manufactured by the Monroe Calculating Machine Company of Orange, N. J. T h e keys of this unit are actuated by solenoids connected to banks of thirty-six miniature T y p e 5696 tetrode thyratrons in each electronic counter. These serve to record the first four digits of each of the stored counts; only nine solenoids are necessary for each decade since the listing machine prints zeros automatically. In this way eight of the thirteen available places are used; five remain for other purposes. /. Timing Unit. T h e functions of the timing unit are automatic sequencing and timing of the operations of the Pierce Photometer. T h i s may be understood most easily by reference to the block diagram of Figure 5. T h e basic clock frequency of the apparatus is generated by a T y p e 816 Vacuum T u b e Precision Fork manufactured by the General Radio Company of Cambridge, Massachusetts. T h e fork generates approximately sinusoidal voltages of nominal frequency 115.7404 cycles per second with a stability of ten parts per million. T h i s frequency corresponds to a period of 10~7 day. T h e nominal frequency is adjustable so that the timing can be set to have an error of no more than 10 -5 day per day. Stability

104

The Present and Future

of the Telescope

of Moderate

Size ι.

Κ UI &

ζ Ο Ο οζ

δ υ t β)

ζ ϊ 5

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