URSIES: an Ultravariable Resolution SingleInterferometer Echelle Scanner

Arne A. Wyller and Theodore Fay

A Fabry-Perot interferometer in a Ramsay mount is used in tandem with an echelle Hilger monochro-

mator with pinholes instead of slits. The instrument, URSIES, is enclosed within a pressure chamber

filled with Freon. Photoelectric pulse counting techniques and pressure scanning are used to record

the spectrum. This design has four basic advantages: (1) The resolution of the scanner is variablefrom 5.0 A to 0.005 A, and the effective wavelength range is quite broad, from 3500 A to 13,000 A. (2) A

light gain of ten over conventional grating scanners at resolutions of 0.1 A or better is achieved. (3) Very

low levels of light from outside the wavelength passband reach the detector; for a resolution of 0.03 A

at 6000 A this level is 5%. (4) There are very low levels of scattered light from the pinholes when ex-

tended sources are observed. The scattered light from the instrument at the edge of the solar disk is

found to be less than 1% from 4000 A to 11,000 A. Measurements demonstrating these advantages arediscussed.

1. Introduction

The technology in high resolution spectroscopy has,in the last decade, developed along two major lines toimprove observing efficiency over past traditionalprism and grating techniques. As is well known G. R.Harrison has devoted a lifetime to create superb qualityechelles for high resolution work. The advantages ofthe echelle have been vigorously propounded, especiallyby Stroke,' with the major features of complete wave-length coverage with one and the same grating, high lightefficiency through high blaze angles that leads to com-pact systems (i.e., high angular resolution). Althoughphotometry with echelles has been questioned, 2 this isprobably due more to the use of prism order sortersthan to intrinsically high scattered light levels of theechelles themselves. Recent studies by Mallial andPetford et al.4 have convincingly demonstrated thatthe scattered light levels and intensities of ghosts inechelles compare quite favorably with conventionalgrating performance.

In astronomical spectroscopy the echelle is gainingincreasing acceptance as leading to versatile and com-pact spectrometer systems5' 6 for high resolution studies.

The other line of development has been the increasinguse of Fabry-Perot interferometers for improving thelight efficiency of spectrometers. This view has been

The authors are with the Bartol Research Foundation, Frank-liii Institute, Swarthmore, Penlsylvania 19081.

Received 23 November 1971.

pioneered by the French school under the leadershipof Jacquinot. 7 In numerous research papers by mem-bers of his group, the advantages in the light grasp ofthe Fabry-Perot interferometer over the conventionallow blaze angle grating have been demonstrated. Theirwork culminated in a monumental study by Chabbal, 8

who theoretically and experimentally examined themerits of one-, two-, and three-etalon systems. Astro-nomically, the most famous three-etalon system de-veloped out of this pioneering work is the PEPSIOS, 9

which has given remarkable results both in solar 0 andstellar" high resolution spectroscopy. Several one-etalon systems have evolved,' 2" 3 which generally arecalled hybrid systems since they usually are coupled to apredispersing grating or a narrow-band filter.

In astronomy the single Fabry-Perot interferometersystem has been slow in gaining acceptance; major useshave been in low resolution emission line photometry ofextended light sources,'4- 6 while high resolution workhas been pioneered by Vaughan. 17

The coupling of an echelle and a Fabry-Perot interfer-ometer in a pressure scanning mode appears so far not tohave been used in solar or stellar work. In the follow-ing sections we will describe such an instrument whichwe believe offers unusual versatility in multifariousmodes of operation. Summarizing these, the presentlydeveloped URSIES can operate as a pressure scannereither (a) with the echelle alone, (b) with the inter-ferometer alone, or (c) with the echelle and the Fabry-Perot coupled together.

The first mode offers the use of a moderate resolutionspectral scanner with bandwidths of a few angstroms

1152 APPLIED OPTICS / Vol. 11, No. 5 / May 1972

Fig. 1. View of optical parts ofURSIES from below chamber.Focal lengths of lenses and mir-

rors are not to scale.

down to 0. 1 A, where the novel manner of pressure scan-ning is performed rather than a mechanical turning ofthe grating or the photomultiplier along the spectrum.

The second mode allows one to use the interferometerfor low resolution (bandwidth around 1 A) for emissionline spectroscopy, where the high angular field of ac-ceptance of the interferometer permits the study ofphysical conditions in, for example, bright emissionnebulas or solar surface features.

Finally, the third mode converts the simpler relativelylow resolution units to a full-fledged high resolutionscanner, where the full capabilities of the individualcomponents are combined to give the utmost in highresolution and spectral purity in high precision photo-electric scans of line profiles.

11. Design

A major impetus to seek developments of spectrom-eters of the interferometric type in astronomy lies in thepossibilities such systems offer for increased efficiencyin the utilization of the available light flux collected by atelescope. This latter flux is equal to the product ofthe aperture area D2tei and the accepted solid angle1k2 of the source, viz., (DteI) 2. Vaughan aptly termsthis quan tity throughput. 7

Clearlyj the corresponding throughput of the analyz-ing spectrometer should match that delivered at itsentrance aperture by the telescope. In practice, withconventional grating spectrographs, this condition isfar from being realized either in the laboratory or at thelarger telescopes when high resolution is to be achieved.A 2 5-g slit is usually required to obtain a resolution of0.01 A, while at the entrance slit the seeing disk of astar may be 1 mm or more in size.

In addition to the light losses, such a narrow slitgeometry can scatter light, so that both spectral purityand information about structure in the light source are

reduced, as experienced in the observation of sunspotsby Blackwell et al.11 Image slicers,"' with their manyreflections, are only a partial solution to the problem oflight loss and generally increase the amount of scatteredlight.

JacquinotO02 1 discussed the possibility of using aFabry-Perot interferometer in tandem with a gratingspectrometer to increase the light efficiency of thespectrometer at high spectral resolution. Dunham2 2

pointed out important astronomical applications of Jac-quinot's ideas. One of the most difficult problems en-countered in using a Fabry-Perot with a grating spec-trometer is keeping the light beam from the spectrom-eter properly aligned with the fringe pattern of theFabry-Perot throughout the measurement. This align-ment is very sensitive to changes in temperature, pres-sure, and vibrations that cause small relative motionsof the grating, the spectrometer, and the Fabry-Perot.'

In our design, URSIES, both the Fabry-Perot inter-ferometer and the echelle Hilger monochromator areplaced within the same cylindrical pressure chamber.The slits of the monochromator are replaced with pin-holes as shown in Fig. 1. A set of entrance pinholes Piof sizes 2 mm, 1 mm, 0.4 mm, 0.2 mm, 0.1 mm, and0.05 mm and exit pinholes P of sizes 2 mm, 0.4 mm,0.2 mm, 0.1 mm, 0.05 mm, and 0.025 mm are used.Both sets of pinholes are mounted on slides activated bymicrometers for precision positioning perpendicular tothe optic axis and the dispersion of the echelle.

A lens of diameter D focuses the light from a distantsource onto the pinhole of diameter P. This light iscollimated by the mirror Ml and dispersed by the echelleGe. The spectrum from the echelle is focused by themirror M2 onto the pinhole of diameter P2, which iso-lates a small wavelength region AX. The throughput,T, of the monochromator is defined as the area of thelens of diameter D times the solid angle Q~l accepted atP,. Thus T can be expressed in terms of the focal

May 1972 / Vol. 11, No. 5 / APPLIED OPTICS 1153

ratio of the lens and the area, A,,, of P,

T = r X A,,/(f,,2 X 4), (1)

noting that

p = A pl/FP 2 and fpl = Fp,/D.

The areas of the pinholes Pi and P2 are related bythe focal ratios of the collimator and camera mirrors,fmi andfm2, by

Ap = A2 f.12 /f.2 2, (2)

where

Ap2 = r X (P2/2)2.

The size of the spectral interval isolated by the mono-chromator, AX, is related to the diameter of the pinhole,P2, the angular dispersion of the echelle, dO/dX, and thefocal length of the camera mirror, Fm2 by

P2 = Fm2 X A X d/dx, (3)

where

Fm2 = 2fm2.

The focal length of M2 was chosen to be 600 mm.Using Eqs. (1)-(3) we find that the throughput is

T = (fri/fp1)2 X (7rM2/4)2 X (AX X do/dX)2, (4)

since the area of P, can be expressed as Ap, = fmi2 X

M22 X (AX X do-/dX)27r/4. In consideration of slitspectrographs, it may be advantageous to operate withone-dimensional throughputs rather than areal through-puts that are more appropriate to our pinhole geome-tries and slitless interferometer system (see usage byLiller,6 for example).

The throughput may be increased by increasing thesize of the camera mirror, M2. However, the cost ofthe mirror and other associated optical parts increasesas the cube of the diameter of M2. The function ofM2 is to focus the parallel beam from the echelle grating.Shafer et al.2 ' have shown that to minimize abberationsin the Czerny-Turner design, fm1/fm2 - 2. Since allthe optical parameters of the grating monochromatorare interdependent [through Eqs. (1)-(3)], the designrequires that when one increases the diameter of M2 onemust also increase the size of Ml and Ge as well as thesize of the chamber.

In our case the parts Ml, M2, and Ge are designed sothat the diameter of the light beam G = 7.5 cm. Sincethe size of Ge is 12 cm X 12 cm and the echelle must betilted 63° with respect to the beam reflected from MI,the light loss due to overfill of this beam on the gratingis 12%. This light is sacrificed to maintain a constantwavelength and light intensity through the pinholeP2 , since they may change due to changes in the posi-tion of the light source with respect to pinhole P. Thisis especially important in astronomical applicationswhere telescope tracking errors and image motion due toturbulence in the earth's atmosphere cause the imageto move irregularly on the entrance aperture.

The throughput can also be augmented by increasingthe angular dispersion of the dispersing instrument.

Prisms have lower angular dispersions than gratings.The choice of an echelle grating is more advantageoussince its angular dispersion in an autocollimatingmount2 4 ) is

do/dX = 2 X tani/X, (5)

where i is the blaze angle of the echelle. In our case,i = 630. Hence the angular dispersion of this echelle isfive times that of gratings of blaze angles of 200 or less.The linear dispersion measured at P2 is 2.1 A/mm atX4250 and 3.0 A/mm at X5890.

A disadvantage of the echelle grating is the manyoverlapping orders at the position of pinhole P2. Thenumber of overlapping orders can be diminished byincreasing the number of rulings per unit length. Manyechelles5'6 are used with 73 grooves/mm and operatein the fiftieth order or higher. Our echelle has 300grooves/mm and operates in the twelfth order. Theelimination of overlapping orders by filters is thus great-ly simplified. We have chosen interference filters ofbandwidths of 500 A to isolate these different orders.These filters give 50% or more transmission inside thepassband and less than 0.1% outside. They enable usto avoid the light loss, changes in light intensity andwavelength, ghosts, and also cost of two other optionsof separating the orders, namely, cross dispersion grat-ings and order sorter prisms.

The only parameter remaining for us to choose inEq. (4) is the spectral width isolated by P2, AX, and thisvaries the spectral resolution of the monochromatorX/AX, which can be made as large as Nm. In thepresent case, the total number of lines ruled, N, is30,000, and m, the order of the echelle, varies from 5 to15, depending on the wavelength. Thus AX may be assmall as 0.02 A, however, as can be seen in Eq. (4),when AX is decreased from 0.2 A to 0.02 A the loss inthroughput is a factor of 100 for pinhole slits. Tomaintain high throughput as well as high resolution, weneed to couple the echelle monochromator to an inter-ferometer inserted after pinhole P2.

The interferometer acts as a multipassband filterwhose transmission maxima appear at the wavelengthsAn,

x, = 2 X x t X cosl/n, (6)

where A is the index of refraction of the gas between theplates, t is the plate separation, n is the order of inter-ference of the plates, and 3 is the angle between the opticaxis of the plates and the incident light beam. In thepresent design, u 1, t - 107 to 10 8A,n-4 X 103to4 X 104 , and cos3 - 1 are chosen. To obtain maximumtransmission we must match the wavelengths Xn of oneof the Fabry-Perot maxima with one of the monochro-mator maxima, Xm given by

Xm = 2 X x d X sini/m, (7)

where d is the distance between the steps of the echelleladder and m is the order of the echelle. The functionof the echelle is to act as a predispersing filter that ex-cludes adjacent Fabry-Perot maxima X,,-, or Xn+i (Fig.2) while preserving the high throughput. The in-

1154 APPLIED OPTICS / Vol. 11, No. 5 / May 1972

Xm/HL

H

8A

* 0

* Xm-

X * n+I -R5820 5830 5840 5850 5856

WAVELENGTH (A)Fig. 2. Channel spectrum measured by URSIES. The y axis isthe photoelectron count rate, and the x axis is the wavelength in

angstroms.

creased resolution is provided by the sharpness of theX,, maximum.

One difficulty with this design is to change the wave-length so that X and Xm remain exactly equal. Mc-Nutt' 2 oscillated his grating through all Xm betweenX, and X±I at each wavelength measured and then sub-sequently varied us to change the wavelength. On theother hand, Vaughan' 7 changed 1u in X and i in Xm SOthat Xm and Xn remain equal. It is important in thisprocess to keep the alignment between the monochro-mator light beam and the Fabry-Perot optic axis con-stant during the measurement. In our design, byenclosing the echelle and the interferometer in thesame pressure chamber, the synchronization is simplyeffectuated by varying 1, since both Xn and m arelinearly dependent on u.

The index of refraction is varied in steps by changingthe pressure of the gas in the chamber, checking thealignment at each step. The change of the wavelengthAX, due to the change of the gas pressure is determinedby

Axp = (X x P X dg/dP)/1,. (8)

Freon was chosen as the scanning gas, since its refractivegas constant is about five times higher than that of air.Because of the large Freon gas constant the pressure canbe changed five times more slowly for a given wave-length as compared with the air scanners of bothVaughan' 7 and McNutt.' 2 We have found that themore slowly the pressure is changed, the less likely thealignment of the grating or Fabry-Perot will be dis-turbed. The pressure of the Freon gas in the system ismaintained constant during the measurement by aservocontrol unit from Texas Instruments (model 156).Between measurements the pressure can be varied from25 cm to 250 cm in steps as small as 0.025 cm. At theNa D2 line this pressure interval (225 cm) correspondsto a scan width of about 20 A.

The resolution, /6X, of the Fabry-Perot plates de-

pends upon the effective finesse, NE, and the distancebetween orders of the Fabry-Perot, AXF,

ax = AxF/NB with AXp = X'/(2 X A X t) = C, (9)

where we recall that A is the spectral interval trans-mitted by the exit pinhole P2 of the echelle mono-chromator. The quantity, C, is Vaughan's parameter,which controls the parasitic light, i.e., the radiation,transmitted by the interferometer outside 6X and withinAX (Fig. 2). The level of parasitic light is conditionedby the residual light minimum between the bandpassmaxima of the Fabry-Perot. This quantity is deter-mined by the effective finesse of the plates,

IL/IH = 3/(2 X NEI), (10)

where 1h is maximum transmission and IL the minimumtransmission of the plates.25

As can be seen from Eqs. (9) and (10), we desire thefinesse, NE, as high as possible, not only to reduce theparasitic light but also to increase the resolution 6Xfor a given AX. The value of NE depends on (1) theflatness of the plates; (2) the quality, reflectivity, andtransmittance of the coatings; (3) the parallelism of theplates; and (4) the diaphragm finesse.7 We should liketo stress that the only reliable way to estimate thiseffective finesse is by actual measurements illuminatingwith a continuum light source the predispersing mono-chromator. The coupling to the interferometer givesthe familiar channel spectrum. By measuring thewidth of the channel X and the channel separation AX,the effective finesse is evaluated. Usually theoreticalconsiderations and estimates are quite unreliable."'6

In theory one can show that

NE = (NE2 + ND' + NF2)V-\13,

where NE is the reflectance finesse NR = 7r/R/(1 - R).R is reflectance, while ND is the flatness finesse, being ameasure of the deviation from planeness of the platesND = AX/2V2 and NF is the diaphragm or resolutionfinesse,

Ramsay" has manufactured two pairs of Fabry-Perot plates to a flatness of X/120, one of which wascoated by Ramsay giving NE- 25 from 4000 A to 7000A, and Industrial Optics coated the other pair with simi-lar results from 8000 A to 12,000 A.

The plates must be kept parallel to X/(2 X NE). Weuse a Ramsay" mounting of the plates to accomplishthis. The Ramsay mounting was chosen because of itsgreat operational simplicity and ruggedness. Theservocontrolled maintenance of parallelism and theadjustment to parallelism we have found highly reliableand efficient. Usually parallelism is achieved to X/80in 10-15 min and maintained for hours.

The diaphragm finesse, NF, must be at least equal tothe effective finesse. It is a measure of the thickeningof the channel width 6X due to the variation in the anglei caused by the finite width of the aperture P2 that feedsthe Fabry-Perot (see Eq. (6)] through its collimatinglens,

May 1972 / Vol. 11, No. 5 / APPLIED OPTICS 1155

Am

7 H-

6 /--- tAX -

x. 5

UJU) 4

(n1-3z020

I L * n

* XmM-1* X6480*0 4-H L* 00**e l 0g 6 l0 6

A

!10

NF = rX/(2 X x t x Qp2) = 2 X X x Fp22/(/ X t X P2'),

since (11)

Q,2 = r X P22/(2 X FP2)2.

We designed the monochromator to maximize P2 inorder to maximize the throughput from the objectiveD, so we must increase F,2 until N, is at least 30; wechose Fp2 = 7 cm.

The control of the plate separation of the Fabry-Perot, t, must be to within 10 A for an accuracy of 0.01 Ain X,, as can be seen from Eq. (6). Ramsay26 uses areference interferometer with quartz spacers in tandemwith the main interferometer to solve the problem ofseparation control. This method will be independent ofpressure changes as the spectrum is scanned, but theplate separation and thus the spectral resolution ofURSIES cannot be changed without changing thequartz spacers. Also the white light reference fringesproduced by the two coupled interferometers are verybroad, and the servolock is not as firm as with the Ram-say laser method. In this method a laser beam (X6328)is passed through the periphery of the main interferom-eter plates. and reaches a photomultiplier. When theplate separation is an integral multiple of the laserwavelength, there results a monochromatic interferencefringe maximum. The photomultiplier signal maxi-mum locks the servocontrol very tightly and controlsthe ensuing voltage to the piezoelectric crystals whichregulate the plate separation.

In addition to monitoring the laser beam, Ramsay'7

places a quartz window between the plates, the phaseplate, whose optical path can be changed by rotatingthe window with respect to the plates. While Ramsayperfected the phase plate-laser method for a given gas

330

320

< 310WCXW 300 -w

a. 290_

CL 280 -

270 -

35 37 39PRESSURE

(inches)

Fig. 3. The phase plate micrometer vs Freon pressure cali-bration curve. The micrometer reading is given in thousandths

of an inch.

Fig. 4. Right: Total spectrometer assembly set up for finessemeasurements. Left: Auxiliary instrument rack with dual-channel amplifiers and discriminators, Fabry-Perot servocon-*

trols, pressure gauge, and servocontroller.

pressure, our scanning technique requires a phase platecalibration vs pressure change. If the geometrical path-length for the laser light beam in traversing the phaseplate is kept constant, each new pressure change willshift the transmission of the main interferometer awayfrom the Airy transmission maximum at the new wave-length and consequently falsely modulate the observedintensities. Accordingly, we tilt the phase plate tocompensate for the effect of the pressure change sothat the laser optical pathlength between the plates ismaintained constant.

The phase plate-pressure calibration curve is shownin Fig. 3. The observed points in the figure were mea-sured by placing a tungsten filament source close to theentrance pinhole P1. At each pressure the phase platewas tilted so that photoelectron count rate from thesource was maximized, making the transmission of theinstrument highest. The phase plate was rotated by amicrometer mounted on a 4-cm pivot arm, and thephase plate micrometer reading in microns is shown onthe vertical axis of Fig. 3. The pressure shown on thehorizontal axis is the Freon pressure in inches of mer-cury. Spectral scans using this calibration curve will bediscussed in Sec. IV.

A photograph of the entire assembly is shown inFig. 4. To the right is seen the echelle and the inter-ferometer mounted in the common pressure chamber(diameter 60 cm, length 120 cm). The chamber ismounted on a heavy table that rolls easily, but it can befirmly jacked up from the floor and oriented in threedegrees of freedom for steady coupling to a given tele-scope configuration. The installation of remote controls(externally on the vertical wall of the pressure chamber)for the parallelism and spacing adjustments of the in-terferometer plates, tilt variations of the grating, andselection of exit pinholes has been carried out.

Ill. Accessory Equipment

To take advantage of the versatility in spectral resolu-tion designed into URSIES a variety of accessory equip-

1156 APPLIED OPTICS / Vol. 11, No. 5 / May 1972

ment has been developed. Some of this equipment per-tains to the recording of the light signal; some of itpertains to further analysis of the light quality in termsof polarization.

The parallel light beam from the Fabry-Perot is fo-cused by the lens FFP onto the photocell Ph, as shownin Fig. 1. The design of the focal length and the lenssize are given by Vaughan' 7 and conform to the re-quirement that the telescope objective D be imaged ontothe photocathode to minimize intensity fluctuations dueto image motion.

In the case of sunspot work, we desire to isolate re-gions with diameters of 2 see of arc see or less on thesolar disk and wavelength bands as small as 0.01 A.Thus we use as little as 10-"1 of the sun's light, andphotocells of low dark noise are required. Dry-icecooled EMI 9524S and RCA 7102 tubes have been usedfor observations, respectively, in the visual and ir spec-tral regions.

In addition to monitoring the light beam through theFabry-Perot plates, we also monitor a 60-A wide bandof the continuous spectrum using the beam splitter B, infront of pinhole P. A RCA P21 photocell is placedat P2, and it is operated inside the chamber at roomtemperature. This photocell monitors any variationof continuous light from the astronomical sources sothat these variations will not be interpreted as spectrallines. In the present design, the signal on the photo-cell, P2 is from 10 times to 1000 times as strong as Phl,since the reflectivity of the beam splitter B, is about10%, and the ratio of the bandwidth of the light beamson Phl to that on Ph2 varies from 10-2 to 10-4. Lineprofiles are therefore measured in terms of the ratio ofthe counts of Phl to that of Ph2 rather than the counts ofPh, alone. When working on the sun, the intensity oflight on the continuum monitor is reduced by a variableneutral density filter mounted in front of the P21 tubewindow and operated by remote control outside thechamber.

Because the relative sensitivity of the two photocellscould vary, it would be better to measure both lightbeams with the same photocell. To accomplish this, alight beam chopper, Ch, a beam splitter, B, and mirror,Mb, are placed as shown in Fig. 1. The light beamchopper consists of a rotating vane that stops the lightfrom B 50% of the chopping cycle. The light beam ischopped 60 times/sec. When it is stopped the signalfrom the Fabry-Perot is measured, and when it isallowed to pass through the chopper the sum of the sig-nals from the Fabry-Perot and from mirror Mb is mea-sured. This design allows monitoring both high andlow resolution beams with the Ph, photocell, makingthe measurement of the light ratio independent of thephotocell sensitivity or the polarization of the incidentlight. 21

Presently a dual-channel transistorized pulse counterof type SSR model 1110 has been installed, since in ourcase pulse counting equipment is more convenient thandc recording techniques, because the light levels arelow.

Brandt" found that the average amplitude of the

image motion of astronomical light sources due to airturbulence decreases very sharply as the frequency ofthe motion increases from 1/10 Hz to 10 Hz. Toexamine the effect of the image motion on the recordingof the spectrum, counters and printers are set to recordthe accumulated counts in the time intervals 0.1 sec,0.8 sec, and 9.8 sec. The principal advantage of thepresent design is to estimate an upper limit on theamount of contamination of sunspot light from thenormal photosphere due to image motion. In prac-tice, we find useful integration times of 0.8 see in sun-spot observation and 9.8 sec in stellar work.

The continuum monitor is also used as an aid in guid-ing. In the case of stellar observations the maximumcount rates are observed before readings are taken. Inanother example, in our work on sunspots we choose theminimum light ratio between the normal photosphereand the spot. The average value of this ratio is thenchecked against previous photoelectric work. 0

The index of refraction of the Freon gas in the pres-sure chamber is sensitive to changes in temperature asshown in Eq. (8). Instead of attempting to controlthe temperature in the chamber, we monitor it by em-ploying an iron-constantan thermocouple wire insidethe chamber that measures the emf produced. Thismeasurement is done with a potentiometer placedoutside the chamber, and the potentiometer is calibratedwith a mercury thermometer.

The URSIES can also be used to measure the circularand linear polarization in spectral lines. This polariza-tion may be caused by magnetic fields in the light sourceor by light scattering processes. To detect the polariza-tion we place a KDP crystal between the telescope ob-jective and the entrance pinhole (Fig. 1). After thelight passes through the KDP crystal, it is analyzed by aGlan-Thompson prism. We are able to rotate both theKDP crystal and the Glan-Thompson prism by 360°independently of each other, to check whether our mea-surements are sensitive to these rotations. The trans-mission of the KDP crystal plus the Glan-Thompsonprism to unpolarized light is found to be about 25%.

The quarter-wave voltage is applied to the KDPcrystal to measure circularly polarized light. We em-ploy a positive voltage for right-handed circular com-ponents and a negative voltage for left-handed circularcomponents. The quarter-wave voltage of our KDPcrystal is 1100 V at 4254 A (Cr I) and 1600 V at 5890 A(Na D). It is thus to a first approximation linearlyproportional to wavelength in accordance with thespecifications of the manufacturers (Isomet). A highvoltage oscillator modulates the KDP plate with asquare wave of a frequency variable from 0 Hz to 1000Hz and voltage from 0 V to 3000 V.

The difference in intensity between the right- andleft-hand circularly polarized components is measuredby the difference in the count rate as the quarter-wavevoltage on the KDP crystal is switched from positive tonegative polarity. The count rates of respective polar-ization channels are fed into a versatile dual-channelphoton counter produced by SSR Instruments Co.(model 1110). It can be programmed to display and

May 1972 / Vol. 11, No. 5 / APPLIED OPTICS 1157

print out a difference count between the two channels aswell as the sum of the two channel counts. In theformer mode, the dark and scattered light contributionsare automatically eliminated, and a pure Zeeman differ-ence signal is recorded. The signal strength variationcan be observed as a function of wavelength with ourgrating or pressure scanning modes of operation, and thecomplete Zeeman component profiles can thus be builtup.

When the angle between the optic axes of the KDPand Glan-Thompson prisms is fixed at 45°, to samplecircularly polarized light the intensity difference at thewavelength 5250.18 A in the violet wing of the iron5250. 23-A line of the sunspot spectrum was found tobe about 12% and less than 1% in the normal photo-spheric spectrum.31

We also have obtained 2-5 counts/sec from Procyonand Sirius with a 0.2-A bandpass at the X4254 Cr i lineused by Severny.32 These polarization results are verypreliminary. However, they indicate that such workon the sun and bright stars is possible using only a 38-cmsiderostat (effective aperture 30 cm) and URSIES. Itis planned to continue observing linear and ellipticalpolarization in these spectral lines in other stars. If thevoltage on the KDP crystal is changed from zero totwo times the quarter-wave voltage, linear polarizationcan be measured in most circumstances. Voltagesbetween zero and the quarter-wave voltage can be usedto measure elliptical polarization.

IV. Performance of URSIES

To test the performance of (1) the Fabry-Perot alone,(2) both Fabry-Perot and echelle monochromator intandem, (3) the echelle alone as a spectral scanner, aseries of spectral scans has been made prior to the instal-lation of the KDP crystal and Glan-Thompson prism.The performance of the Fabry-Perot alone is shownby the channel spectrum in Fig. 2, which was measuredby placing a continuous spectral source behind a 1-mmdiam pinhole in front of the exit window of URSIES.The light from this pinhole was made parallel by thelens FFP and passed through the Fabry-Perot andimaged onto P2 with lens of focal length F,2 . Thisimage was then passed through the monochromator andonto a photocell placed after the pinhole P,.

In Fig. 2 the x axis represents the amount of thewavelength change made by rotating the echelle Ge, andthe y axis refers to the counts per second from the photo-cell. The filled circles are the measured points. Forthis particular scan the large observed peaks at thewavelengths shown on the y axis are associated with thetenth order of the echelle, while the smaller peaks withthe ninth order are due to the red leak of the yellowfilter at FL (Fig. 1). The large peak is a measure of theinstrumental profile of the Fabry-Perot with a plateseparation, t = 8 X 105 A. This plate separation is

found from the measured free spectral range, AXF = 20A. Thus the finesse, NE, is determined by either Eq.(9) or Eq. (10) by measuring the ratios 5X/AXF andIL/IH, respectively, and it is found to be 18. This is a

lower limit since the mechanical grating scan is not finegrained enough to ascertain the true height of the trans-mission peak.

The shape of this instrumental profile remains un-changed as the plate separation, , increases except forthe linear increase of the scale of the x axis in Fig. 2.To make high resolution spectral scans we increase t to 2mm, in which case AXF becomes 0.8 A and 5X = AXF/NE

0.04 A. The width of the monochromator passband,AX, also remains unchanged for given pinhole apertures,and the relative widths of AX to X are shown in Fig. 2by the solid line under Xm (t = 2 mm). The width ofAX can be controlled by varying the size of P2 , thuschanging Vaughan's C in Eq. (9).

The amount of light within the passband AX excludingbX has been measured (at X5890) to be 5% of the amountof light within the X passband when P2 = 100 A andt = 2 mm. This light is called the parasitic light andcan be increased or decreased by varying either t or P2 .Our value compares favorably with the performance ofthe superb three-etalon system, PEPSIOS, developedby Mack and Roesler for which a value of 7% is quoted. '7

A continuum source or tungsten filament can be used tomeasure the parasitic light. This source is placed atP1, and the light intensity, IH,* is measured when themonochromator profile is centered on the large peak asshown in Fig. 2. The measurement, IL,* is made whenthe Fabry-Perot fringes are shifted by a slight change int so that the monochromator profile is placed halfwaybetween the two large transmission peaks of the inter-ferometer. This ratio IL*/IH* gives the parasitic light.Since there are no spacers in the Ramsay mounting ofthe Fabry-Perot, we can easily explore the change of theparasitic light with changing t. We findIL*/IH* = 0.1,when t is increased only to 3 mm. If P2 were decreasedto 50 iu the parasitic light at t = 3 mm would be re-duced. However, this would reduce the throughputalso.

Our first scan using both Fabry-Perot and mono-chromator is the Na D2 line on the solar disk center (Fig.5). The filled circles are the measurements of McNutt' 2

and Waddell3 3 that were made with a bandpass of 0.02 Aand the open circles are our measurements made with abandpass of 0.04 A (t = 2 mm, P2 = 100 ,i). Thewings of the sodium line are broad enough so that theratio of light intensity through P2 at 5890 A and 5888 Ais 0.25. In this measurement, the parasitic light for theFabry-Perot was found to be 0.05. Therefore ourpreliminary estimate of the correction to the Na D2line profile would be about 1% at line center. Our un-corrected observation of the residual intensity at thecore of Na D2 is 5%, hence our result and the correctedprofiles of McNutt' 2 and Waddell 3 agree within 1%.Thus our results are comparable with those of the bestdouble-pass spectrographs3 or hybrid interferometer-grating systems.'2

Our own scans have been repeated on two differentdays to within 1%. When the instrument is used onlyat the maximum designed throughput, our parasiticlight is high. However, the parasitic light can be re-duced either by decreasing the size of P2 for a given t

1158 APPLIED OPTICS / Vol. 11, No. 5 / May 1972

x 10-J

LL

-j>W

a:

32 36 40PRESSURE

5889.2 5890.0 5890.8WAVELENGTH

Fig. 5. A comparison of scans of the Na D2 line on the solardisk by URSIES and previous workers (see text).

or using plates coated for a higher finesse. Ramsay hasrecently provided us new plates with NE = 30. In thefuture, finesses as high as 50 may be achieved with hismounting.

The triangle shown at the bottom of Fig. 5 is a mea-surement made when the optic axis of the Fabry-Perotand the monochromator beam differ by an angle of10-3 rad. In this measurement the parasitic light isfound to be 13% when the instrument is misalignedby 10-3 rad. Therefore, it is very important to keepthese two optic axes aligned to 10-5 rad or better for themost accurate measurements. The alignment can bechecked before the measurement, since the position ofP2 is externally controlled by a micrometer to an accu-racy of 0.025 mm. The grating may be rotated byexternal control on a horizontal axis by angles as smallas 2 X 10-' rad to correct for any misalignment. Dur-ing the repressurization cycle, grating motions of 10-3rad have been observed, while during normal step scan-ning motions up to 10-4 rad are noted.

Scans of the spectral lines of the electrodeless dischargetubes (discussed in Sec. III) have also been made,using the echelle alone. It takes only a few hours toopen the pressure chamber, remove the Fabry-Perot,insert lenses to image the light from P2 onto the de-tector at Phl, and repressurize. The scan of the heliumemission line at 4922 A is shown in Fig. 6, where thelogarithm of the ratio of the counts from Ph to thoseof Ph2 is plotted against the Freon pressure (in inches ofmercury). As can be seen from Fig. 6, the count leveldecreases by a factor of 103 or more from the peak of thehelium line to twice line half-width. This continuummay be due to the scattered light in the monochromator,or it may come from the light source itself. The pres-ent low level shows the quality of the electrodelessdischarge tubes and the pinhole monochromator andcompares well with recent experience on the low scat-tered light levels in echelle spectrographs by otherworkers.'

The conversion factor from Freon pressure to wave-length was determined to be 0.072 A'/cm by scanningfrom the mercury line at 4916 A to the helium line at4922 A. This conversion factor should be proportionalto the wavelength as can be seen from Eq. (8). Thishas been checked independently by determining a con-version factor of 0.090 A/cm between the sodium lines5890 A and 5896 A. The temperature of the Freonwas about 16'C at the time of these measurements, andthe uncertainty of these conversion factors due totemperature changes and measurement errors is lessthan 1%.

The instrumental profile of both Fabry-Perot andmonochromator in tandem can be determined from theproduct of the measured profile in Fig. 6 to that in Fig.2, provided that the scale of the x axis in Fig. 2 ismultiplied by the plate separation ratio T(A)/8 X105 A. Studies directly of the instrumental profileusing both Fabry-Perot and monochromator will becontinued with such discharge tubes and lasers.

More severe tests of the scattered light within theURSIES have been made by a series of observationsat the sun's edge as shown in Fig. 7. The measuredscattered light was found to be less than 1.0% at alldistances larger than 10 sec of arc from the sun's edgefor all the wavelengths indicated in the figure. Partof this scattered light is atmospheric in origin (0.1-0.2%), part is due to light scattered by siderostat(diaphragmed down to 10 cm) which probably amountsto 0.2-0.3%. The remainder, 0.5-0.7% is due toscattering in URSIES.

These low scattered light levels are important forthe usefulness of URSIES in sunspot photoelectricspectrophotometry in which until now, owing to tech-nical difficulties, very little work has been done. Thisresult also implies thatof the sunspot light due5.0% during the 0.8-secdiameters exceeding 10

- 1.00

i -i.)0

O O

-2.0

1-

0

-J -3.0

the effects of contaminationto image motion is less thanintegration time for umbralsec of arc. The scattered

28.0 30.0 32.0PRESSURE(inches Hg)

34.0 36.0 38.0

Fig. 6. Scan of the 4922 He line with the electrodeless dis-charge tube and echelle monochromator alone.

May 1972 / Vol. 11, No. 5 / APPLIED OPTICS 1159

o 0 00

V 0 0.

00.oo .0

0° 0

0 9

*0

4922 He Tube Conversion Factor For Pressure To AX8mm Pressure AX

Ep = 0.182 A Inch Hg

0*O

.00

* 0

* .-**-

-4.0-

1I .

1.2

1.017

0.8

I(60A) 0.6

I5 (60 A)

I0 (60A)

INTENSITY OF STRAY LIGHTAT SUNS EDGEINTENSITY OF LIGHT ATDISK CENTER

x

0x

0.4

0.2

0.0

WAVE LENGTH (A)

* 4200-4260o 5870 -5930X 10830

100 80 60 40 20 SUNSEOGE 20-OUT IN-

DISTANCE FROM SUNS EDGE (ARC SECONDS)

Fig. 7. Measurements of the scattered light from URSIES atthe edge of the sun's disk relative to disk center at different

wavelengths.

A scan of the H-,3 line of the star a Lyrae is shown inFig. 9, for the two nights 24 and 25 August 1969. Theerror bars in the figure refer to the statistical error,which is the square root of twice the number of countsper channel, but this error is a minimum estimate ofthe actual errors. The integration time per channelwas 9.8 see, and this can be increased indefinitely toincrease the statistical accuracy. Other stellar scanswith longer integration times have been obtained, andthe results are favorably compared with scans fromother types of instruments." We have also achieved acount rate of 5 counts/sec at 5890 A with a bandpass of0.04 A using Fabry-Perot and echelle in tandem on aLyrae. These stellar observations have been made witha moderate-sized siderostat (30-cm effective aperture)that attests to the high luminous throughput ofURSIES.

V. Other Applications and Conclusions

Ha PROFILES

AS40Y

_ X~~~~~~~~~ NP

A . PIERCE (1954) NORMAL* it' x DAVID (1961) PHOTOSPHERE

j FRICKE AND ELSASSER(1965) Aua82 UMBF

A WYLLER AND FAY I

(1970) Au= 100f

CALCULATIONS--- H. PRADERE (1967)

- H.S. YUN (1970)Y YUNS EMPIRICAL UMBRAL

MODEL (1970)

S40 STELLAR ATM. (Tf = 4000°K)NP NORMAL PHOTOSPHERE

1A

The versatility of the URSIES system for high reso-lution spectroscopy cannot be overemphasized. Inaddition to the applications discussed, URSIES canbe used for emission line work on nebulas and ab-sorption line work on the planets. A count rate of 10counts/sec has been achieved on the 4922-A He line on a 1-min of arc diam region of the Orion nebulausing only the 30-cm siderostat. This He i line is 300times weaker in this' nebula than the stronger hy-drogen Balmer and 0 iii lines.36 Our measurement of

.06

.05

.0411 2 3 4

AX (A)

Fig. 8. Scans of H-ca line in sunspots and in normal photosphere.

light was measured both at the Ph2 and Phl photocellsin order to make certain that the spectrum of thescattered light is indeed photospheric and that nooptical part in the light beam does increase the scat-tered light. A spectral scan of the H-a line at 6563 Ain a typical sunspot is shown in Fig. 8, where differentsymbols refer to different observations and models inthe sunspot umbra. Details of these studies will bediscussed by Fay, et al.' 4

Low scattered light levels are also important forprecision observations of stellar line profiles. Thetheoretical computer-spurred advances in recent yearsincreasingly demand solartype precision in spectralresolution and intensity measurements.

.02

01

PRESSURE9 9A :0 32 34 3 38 40 42

-2.0 0 +2.0WAVELENGTH FROM LINE CENTER

Fig. 9. Scan of H-p line in the star a Lyrae. * = 24 August,1969, 0 = 25 August 1969, A = Mihalas' theoretical model.

1160 APPLIED OPTICS / Vol. 11, No. 5 / May 1972

1.0

0.8

Ix

cV

0.6

0.4

u.U.

44-4v l is -F - , IN" r I

1++>~ j ++

F i+$)+t+++i

l ' .'f

1000%

80%

60%

40%

20%

IA -

0.2

the He line was made with the monochromator, butsince Ramsay measured the transmission of the Fabry-Perot to be 60%, and our monochromator transmissiondoes not exceed 30%, the emission line studies onfaint sources should be made with the Fabry-Perotalone.' 4 This can be done by moving the lens from infront of P2 (Fig. 1) to the position behind B2 and placinga beam splitter and the photocell Ph2 in the light beambetween the lens and M, inside the chamber. Anothermirror is then inserted in front of P2 to deflect theparallel light beam through the Fabry-Perot.

Important application of URSIES in this mode ofoperation should also be found in coronal line pho-tometry. The pinhole geometry, high throughputof light, and dual-channel pulse counting facilitiesshould make it useful in mapping the variation incoronal line profiles as a function of distance from thesolar limb. A virtually unexplored area of researchwould also be magnetometer studies of the coronalline profiles.

Both the monochromator and the Fabry-Perot mustbe used in tandem for high resolution absorption linework on the planets, since the cores of individualmolecular lines in planetary atmospheres are oftenas narrow as 0.01 A. The methods used to studythe sunspots should also be applied to monitor changesin the atmospheres of the planets, from place to placeon the disk and in time. Moroz and Cruikshank"7 havenoted large variations in the strength of the ir ammoniabands of Jupiter at different Jovian latitudes.

Swings 8 has suggested that the clouds observed inthe atmosphere of Venus in the uv are produced byCO2+ ions high in the atmosphere, and such a hy-pothesis should be checked with higher resolutionspectra of Venus and the sun in the region of 3872 Aand at shorter wavelengths. The pinhole geometryof URSIES should allow precision mapping of thechanges in strengths of the molecular planetary bands,almost free from the light coming from other areas ofthe planetary disk. Studies of the polarization of themolecular spectral lines due to scattering in planetaryatmospheres should also be important.

In conclusion, the URSIES has shown itself to per-form reliably and ruggedly in a variety of observingprograms. The instrument does suffer from the mul-tiplex disadvantage, i.e., it can only examine one narrowwavelength channel at a time, while the technique ofphotographic spectroscopy results in a broad wave-length coverage simultaneously. Observations arethus time costlier with URSIES, but at smaller tele-scopes this should not prove a serious obstacle, es-pecially when high spectral purity and precision in theintensity recordings are desired and achieved.

We feel that he instrument design provides re-search capabilities in a wide variety of fields, whichare fully competitive with those of the large coud6spectrographs but at a much lower cost and an orderof magnitude saving in compactness. The high lumi-nous throughput extends very significantly the use-fulness of medium-sized telescopes into the field of highresolution spectroscopy and makes them in this field

competitive with telescopes at least 3 times larger.Although the high spectral purity of the instrumentmakes it eminently useful in solar precision spec-troscopy, we feel that the greatest merits of URSIESwould be exploited in conjunction with medium-sizedtelescopes on high resolution photoelectric scans ofstellar line profiles.

The authors are grateful to Bob Pfeiffer, chief oftechnical services at Bartol, for all his help in workingout the problems in the design. We also wish tothank Arthur Smith and Bob Rupley for their helpwith the electronics and Bartol for making them avail-able to us without cost. We are deeply indebted toGerry Faust for his help in making the electrodelessdischarge tubes, and to Hong Sik Yun for his help withthe rough drafts. We are especially grateful to theDepartment of Astronomy of the University of Penn-sylvania for so generously extending the use of their38-cm siderostat and to J. V. Ramsay of CSIRO forinvaluable aid in our operational familiarization withhis interferometer. Thanks also go to L. Harkness ofthe Hilger-Engis Company for aid in the adaptationof their conventional grating mount to use with anechelle. Finally, we should like to express our sincereappreciation to the National Aeronautics and SpaceAdministration, which has generously supported thisproject since 1967 under grant NGR-39-005-066.

References1. G. W. Stroke, Handbuch Phys. 29, 426 (1967).2. H. Zirin, The Solar Atmosphere (Blaisdell, Waltham, Mass.,

1966), p. 40.3. E. A. Mallia, Solar Phys. 2, 360 (1967).4. A. D. Petford, D. E. Blackwell, B. S. Collins, P. A. Ibbet-

son, E. A. Mallia, G. Smith, and D. Emerson, Solar Phys.19, 264 (1971).

5. D. Schroeder, Appl. Opt. 6, 1976 (1967).6. W. Liller, Appl. Opt. 9, 2332 (1970).7. P. Jacquinot, Rep. Prog. Phys. 23, 267 (1960).8. R. Chabbal, Rev. Opt. 37, 49 (1958).9. J. E. Mack, D. P. McNutt, F. D. Roesler, and R. Chabbal,

Appl. Opt. 9, 873 (1963).10. M. Daehler, Astrophys. J. 150, 667 (1967).11. L. M. Hobbs, Astrophys. J. 142, 160 (1965).12. D. McNutt, "Terrestrial Absorption in the Solar D Lines,"

Ph.D. thesis, U. of Wisconsin (1962).13. J. M. Gagn4, Appl. Opt. 7, 581 (1968).14. P. Cruvellier, Ann. Astrophys. (Paris) 30, 1088 (1967).15. J. E. Geake, J. Ring, and N. J. Woolf, Mon. Not. Royal

Astron. Soc. 119, 616 (1959).16. M. 0. Harwit and T. R. Gull, Paper presented at 134th

Meeting of the American Astronomical Society, 29 March-1April 1971, Louisiana State U.

17. A. Vaughan, Ann. Rev. Astron. Astrophys. 5, 130 (1967).18. D. Blackwell, E. Mallia, and A. Petford, Mon. Not. Royal

Astron. oc. 146, 93 (1969).19. I. S. Bowen, Astrophys. J. 88, 113 (1938).20. P. Jacquinot and C. Dufour, J. Rech. CNRS 6, 1 (1948).21. P. Jacquinot, J. Opt. Soc. Am. 44, 761 (1954).22. T. Dunham, Vistas Astron. 2, 1223 (1956).23. A. Shafer, L. Megill, and L. Droppleman, J. Opt. Soc. Am.

54, 879 (1964).24. G. R. Harrison and G. W. Stroke, J. Opt. Soc. Am. 50,

1153 (1960).

May 1972 / Vol. 11, No. 5 / APPLIED OPTICS 1161

25. A. Hadni, Essentials of Modern Physics Applied to the Studyof Infrared (Pergamon, New York, 1967), p. 53.

26. J. V. Ramsay, Appl. Opt. 1, 411 (1962).

27. J. V. Ramsay, Appl. Opt. 5, 1297 (1966).

28. W. Hiltner, in Astronomical Techniques (U. of Chicago Press,

Chicago, 1962), p. 235.29. P. N. Brandt, Solar Phys. 7, 187 (1969).

30. H. Wbhl, A. Wittmann, and E. H. Schr6ter, Solar Phys. 13,104 (1970).

31. W. C. Livingston, in Astronomical Techniques (U. of Chicago

Press, Chicago, 1962), p. 336.

32. A. Severny, Astrophys. J. Lett. 159, L73 (1970).

33. J. Waddell, Astrophys. J. 136, 223 (1962).

34. T. Fay, A. Wyller, and H. S. Yun, Solar Phys. 22 (March1972).

35. T. Fay and A. Wyller, Publ. Astron. Soc. Pacific. (inpreparation).

36. H. M. Johnson, Nebulae and Interstellar Matter (U. of Chi-cago Press, Chicago, 1968), p. 95.

37. V. I. Moroz and D. P. Cruikshank, J. Atmos. Sci. 26, 865

(1968).38. P. Swings, Universit6 Liege; private communication (1969).

Current Directions in Applied Holography

19-21 June 1972San Francisco

A course on Current Directions in Applied Holography will be presented 19-21 June sponsored by

the Continuing Education in Engineering and the College of Engineering of the University of Cali-

fornia at Berkeley. It will cover both scientific and nonscientific applications, with emphasis on

commercial potential. Current capabilities and limitations as well as anticipated developments

will be discussed. Topics include recording materials for holography, digital holography, holo-

graphic optical filtering, holographic stereograms, holographic interferometry, acoustical holog-

raphy, biomedical applications, holographic optical elements with applications to displays,

holographic memories, commercial applications of memories, and embossed holographic movies,

John Whinnery, Berkeley, is the UC faculty member in charge of the course; the program organizer

is Joseph W. Goodman, Stanford. Speakers are Byron B. Brenden, Holosonics, Richland, Wash.;

Donald H. Close, Hughes Research Labs., Malibu, Calif.; Nicholas George, Caltech; Paul Greguss,

New York Medical College; William J. Hannan, RCA Labs,'Princeton; Adam Kozma, Radiation,

Inc., Ann Arbor; Emmett N. Leith, University of Michigan; Adolf W. Lohmann, UC-San Diego;

Kent K. Sutherlin, Optical Data Systems, Mountain View, Calif.; John Urbach, Xerox Research

Center, Palo Alto, Calif; and Ralph Wuerker, TRW Systems, Redondo Beach, Calif. The course

will meet at the UC Extension Center, 55 Laguna Street, San Francisco. The registration fee is

$200; advance enrolment is required. Further details are available from Continuing Education in

Engineering, University of California Extension, 2223 Fulton Street, Berkeley, Calif.'94720; 415-642-4151.

1162 APPLIED OPTICS / Vol. 11, No. 5 / May 1972