Simple photometric observations of BR Canum Venaticorum (a troublesome comparison star)

Kevin West, John Howarth & László Kiss

During the last four years an Optec SSP3 photometer has been used to obtain light curves for various suspected and confirmed variable stars. Upon analysis the semi-regular variable BR CVn (formerly a comparison star for V CVn) showed clear periods of 712 and 71 days, both periods having average peak-to-peak amplitudes of less than 0.3 magnitude. The total observed range using a V-band filter has been 6.47—7.17.



BR CVn (SAO 44590, HD 116475) is a poorly observed 6th magnitude star of spectral type M4111 (B—V=1.55). It is located about 1o almost due west of the Whirlpool nebula M51 at RA 13h 23m 05.5s, Dec +47° 00’ 06" (2000.0).1 It appeared as star 69 on the now obsolete BAA chart (1984 April 12) for TU, Y and V CVn. In the past, the star was used as a convenient comparison for V CVn at maximum,2 and light curves of V were adversely affected by BR. Not surprisingly it was visual observers of V who first raised doubts as to the constancy of BR. Reports by Tristram Brelstaff, Richard Fleet, Róbért Fidrich, Patrick Maloney and John Toone were recorded as far back as the early 1980s, and by 1986 most visual observers had stopped using BR as a comparison. The confidence of some visual observers in reporting variability was hampered by the redness of BR. Three BAA visual observers used an alternative comparison of similar colour and magnitude to star 69, located about half a degree to the northwest, and reported a range of variability in close agreement with subsequent photometry. The same comparison has been used in the present study.


Aims of the programme

The aim of the present photoelectric photometry (PEP) programme is to monitor bright (V<7.0) northern variables on a continuous basis and with good coverage. Brief details of other stars on the programme are given in the appendix. The analysis of these will be the subject of future papers. At the start of this programme it was suggested that BR CVn (then only an unofficial, suspected variable) would be an interesting addition. Initially the aim was to confirm or deny variability. This was achieved within a matter of weeks and so BR was promoted to the main programme and the aim became one of classification and the identification of any periodicity. Preliminary analysis of the early data was published by West and Lloyd,4 and as a short note in Observer’s Forum of the 1998 April issue of the Journal.5

Site and equipment

Observations were made by West at his observatory at Ryde on the Isle of Wight. Despite its low altitude (37 metres), the site is near a hilltop and is protected from valley mists. Photometry was performed using an Optec SSP3 photometer equipped with the manufacturer’s V-filter and attached to a 20cm aperture Newtonian telescope. Data in the form of the four-figure SSP3 output were logged manually into a notebook and later transferred into a spreadsheet for reduction.

Although photoelectric photometry can appear complex and difficult, the actual operation of the Optec SSP3 is simplicity itself, and gave the author a user-friendly entry into the field. He would therefore recommend this model and would encourage other amateur astronomers to attempt PEP.


A total of 136 observations were made by West using the setup described. A further 19 observations were made by László Kiss, a Hungarian observer at Szeged Observatory, using a 40cm Cassegrain telescope and Optec SSP5A photoelectric photometer with the V-filter supplied. Coincident observations showed a good match between the two data sets. The accuracy is discussed in a later paragraph. The spectral response of the Optec filter was matched closely to the Johnson V-standard, thereby minimizing transformation corrections (see Reduction section). Each star count comprised three 10-second integrations followed by a background sky count.

Differential photometry was carried out using HD 116172 (V=7.0, B—V=1.1, K0III) as comparison (Co) and HD 116957 (V=5.88, B—V=0.97, K0III) as a check (Ch). The observing sequence has been modified during the course of the programme to improve accuracy and error estimation. Initially, to confirm variability, the following sequence was used:

Co, V, Co, V, Ch, Co

with further doublets of (V, Co) being added if sky transparency was judged to vary excessively, based on experience gained from reducing earlier data. Varying sky transparency, essentially noise in the raw counts, translates into increased errors. This experience aided judgement at the telescope. An adopted rule of thumb as to maintain signal to noise ratio above 100 (see discussion of accuracy). After confirmation of variability the following minimum sequence was adopted:

Co, V, Co, V, Co, Ch, Co, V, Co.

Again further doublets were added when necessary. This ensured that variable and check star counts were always bracketed by a comparison. Each star count was followed by a sky count which was subtracted at reduction.


The atmosphere causes extinction of the light reaching a detector. In PEP, allowance must be made for the fact that comparison and variable may be of different colours (B—V) and observed at different altitudes. Thus two extinction factors must in general be determined. One is for so-called ‘primary extinction’ (dependent mainly on altitude and sky transparency) and the other for ‘secondary extinction’ (colour dependence causing reddening at low altitudes). Also, any photometric system (telescope, photometer, filter) is likely to differ from the original Johnson system. For each filter, it is necessary to determine a correction factor or transformation coefficient so that observations made using different systems can be compared and combined. Methods for determining the above are given by Adams.6

The differential magnitudes (in the sense V—Co =∆V and Ch—Co = ∆Vk) were calculated in a spreadsheet and reduced to instrumental magnitude using:

∆V =2.5 log (Ncomp /Nvar).

The constancy of the selected comparison star was verified by calculating its differential magnitude relative to the check star using:

∆Vk =2.5 log (Ncomp /Ncheck).

where Ncomp, and Ncheck are the mean of three 10-second integrations, respectively of comparison, variable and check with sky subtracted.

To convert differential magnitudes of BR CVn to standard V magnitudes, the comparison star was assigned a magnitude of V=7.00.’ The check star magnitude was V=5.88, but this is only used as a check on the results and is not used in the reduction. The comparison star was of similar colour (B—V) and within half a degree of BR. Also no observations were made at high airmass (altitude <30°). For these reasons, transformations to Johnson V-standard and corrections due to primary and secondary extinction factors can be ignored for this star and therefore have not been made (Adams, 1989).6 However, primary and secondary extinction coefficients, k’v, k"v and transformation coefficient εv for the visual filter have been found using the methods described by Adams in order to verify that these corrections are small enough to ignore. The following results were obtained:

k ́v = 0.25 to 0.40

k ́ ́́́́́́́́́́́v = 0.00

εv = -0.02


All of these values are close to accepted and published values for the SSP3 photometer system at sea level sites. The primary coefficient ḱ́́́́́́́́́v is dependent upon sky transparency and the range shown above reflects values obtained under best and worst (photometric) sky conditions. Very similar values were obtained for the system used by Kiss, supporting the correspondence between the combined observations. These were:

k’v =0.20 to 0.40

k" v = 0.00

εv = -0.05


Experience has shown that the dominant source of error is sky transparency and its fluctuations with time.

Each of West’s observations yielded at least 3 measures of and one of ∆Vk. The standard deviation of the three measures ∆Vv was found and also the deviation ∆Vk from the mean of all previous ∆Vk measures. The larger of these two figures was adopted as the uncertainty. Most observations fell within a 0.01 magnitude error bar and nearly all within 0.02 mag. Any over 0.05 mag were discarded. Excessive fluctuations in star counts indicated changing sky transparency due to invisible cloud. This would be particularly evident in the five bracketing comparison star counts which ideally should provide a constant baseline. In the event of these fluctuations the sequence was increased, and during reduction the signal to noise ratio Sn obtained via:

Sn = S/N

where S = mean of comparison star counts and N = standard deviation of comparison star Counts. If Sn was less than 30 the observation was discarded. For nearly all observations Sn was greater than 100.

Another baseline that should ideally remain constant is the differential magnitude between comparison and check, ∆Vk. A measure of the constancy of check and comparison is found from the standard deviation (sigma) of the set of their differential magnitudes, which was found to be 0.018. This low value of sigma reinforces the conclusion that the primary extinction correction is negligible for the (comparison—variable) differential. Such a source of error would be larger for the (comparison—check) differential ∆Vk because of their greater angular separation. It would appear as an annual cyclic change in ∆Vk. This is not seen in the data.

Photoelectric photometry can achieve millimag (0.001 mag) accuracy when allowance is made for transformations, and primary and secondary extinction, but this does necessitate more thorough and time consuming observing sequences and reductions. It was decided by the authors that in order to cover a programme of some 18 stars, including BR CVn, an internal accuracy of 0.02 mag was a reasonable compromise.

Further confidence was gained in this estimate of error bars by comparison with other PEP observers of BR and other variables. These were always within 0.03 mag and nearly all within 0.02 mag.


The original data comprised 155 observations, spanning a period of 1264 days (a little under three and a half years) from 1994 May to 1997 October. These are shown plotted against time in Figure 1.

On the four occasions that two observations were made on the same night, both were removed and replaced with the average of the pair, in order to avoid giving an unjustified extra weight to these nights’ results. This was the only adjustment that was made to the data before beginning the frequency analysis. The (now 151) estimates had an average spacing of about 8.4 days, a mean value of 6.76 magnitude and a standard deviation of 0.15 magnitude.

The Discrete Fourier Transform was performed on the data to yield an amplitude spectrum, employing a method similar to that used by Howarth for the analysis of T Cas observations.7 The range of frequencies was 0.00025 to 0.06 cycles per day at an increment of 0.00025 cycles per day (cpd). The separation between frequencies was fine enough to ensure that no periodicity could be overlooked by falling between adjacent frequencies. The largest frequency corresponded to a period of 16.67 days, roughly twice the average time between observations. This. was felt to be a suitable compromise between avoidance of aliases and not missing any real frequencies at the high end. Inspection of the spectrum (Figure 2) shows two dominant peaks, at 0. 001405 and 0.01409 cpd, having respective half peak-to-peak amplitudes of 0.144 and 0.130 mag. The corresponding periods were 711.8 (±15.0) days and 70.9 (±0.16) days. The shorter period shows pronounced variation in magnitude, being less (about 0.1 mag peak-to-peak) at JD2449900 and much greater (over 0.5 mag peak-to-peak) at JD2450500. Both of the periods can be identified by eye from Figure 1. Nevertheless, it was decided to check for aliasing, and a spectrum of the window function was constructed (Figure 3). The relatively even distribution of observations produced a clean window function. The small peak at 0.00256 cpd (391 day period) probably reflected a seasonal variation in the density of observations, but could not explain either period in terms of the other. For completeness, the effect was tried of removing each period in turn (‘prewhitening’), and transforming the result. As expected, in neither case was the other period significantly affected.


It must be kept in mind, of course, that the database spans less than two complete cycles of the longer period. Despite the prominent peak, there is no guarantee that this variation will be maintained. Interestingly, the authors’ attention was drawn, in a letter from John Greaves (private communication), to photometry of BR CVn by Hipparcos, which suggested a periodicity of about 700 days over a period from 1989 November to 1993 April (160 observations). Moreover, this apparent 700-day period appeared to conform in phase with our data. Analysis of the Hipparcos observations is, however, beyond the scope of the present paper. In the authors’ data about 18 cycles of the imputed 70.9 day period are spanned (though sometimes rather sparsely), and the sharpness of the peak in Figure 2 implies good phase coherence. However no certain predictions can be made about the future presence of this period, particularly in view of its fluctuations in amplitude. Percy8 identified two groups of semiregular variables based on observed period ratios. In the first group there are two periods with a ratio of 10—15, whilst the second group is characterised by ratios of 1.5—2. BR CVn would appear to lie in the former group.


Using the technique of photoelectric photometry over the longer term, it has been possible to record the lightcurves of stars typically in the magnitude range V = 5 to 7. Taking comparison stars of established constancy, the technique has given reliable and apparently accurate data over many years. We suggest that it is particularly applicable to red stars, which give considerable scatter when observed by eye, and for objects which have more than one periodicity in the tenths of magnitude amplitude range. The technique is recommended to other amateurs with the facility for PEP.

The observations of BR CVn yield two periodicities: of about 710 days and 71 days, the latter showing variations in amplitude of 0.1 to 0.5 mag peak-to-peak. It is not yet possible to comment on the permanence of each period until more data are obtained, or until a more thorough analysis of Hipparcos and visual data is undertaken, though the shorter period seems to have been established for many complete cycles.


In carrying out the present programme of PEP observations it is a pleasure to acknowledge the invaluable assistance given by the following: John Greaves, in drawing our attention to the Hipparcos data and for his help with the analysis; John Isles for comparative observations of TX Psc; Chris Lloyd for checking comparisons and much helpful advice on photometry; Roger Pickard and John Saxton for comparative observations of X Per; and the visual observers previously named.

Address (KW): 5 Edward Street, Ryde, Isle of Wight P033 2SH. []



1 Sky Catalogue 2000.0, Volume 1, Cambridge University Press, 1982

2 Handbook Brit. Astron. Assoc., 1997

3 BAA Variable Star Section Computer Database

4 West K. W., Lloyd c., ‘HD 116475: A New Late Type Variable In Canes Venatici’ IBVS 4289, 1 (1996)

5 West K., J. Brit. Astron. Assoc., 108(2), 64 (1998)

6 Adams B. A., ‘Data Reduction Techniques for Differential Photo electric Photometry’, IAPPP Communication No.38, 7 (1989)

7 Howarth J. J., J. Brit. Astron. Assoc., 107(5), 264 (1997)

8 Percy J., ‘Pulsation Modes in M Giants’, in A Half Century of Stellar Pulsation Interpretations: A Tribute to Arthur N. Cox, ed. P. A. Bradley & J. A. Guzik, ASP Conf. Series, 135, p.249-253

9 Lloyd C., West K., ‘Observations of Low Amplitude Late-Type Variables’, IBVS 4335, (1996)

Appendix - West’s PEP programme

BR CVn is observed as part of a programme of long term PEP monitoring of some 18 northern declination stars, most of which are on the BAA binocular programme. At the time of writing, the main programme and number of observations comprise the following:

Priority 1 list

JD (244.. -245..)

No of observations




UX Dra



Mu Cep



d2 Lyr









X Per



(Plus 8R filter and 251 filter)
Preliminary analyses of the early observations of UX Dra, d2 Lyr, VY UMa and RR UMi are described in IBVS 4335.9
Priority 2 list

JD (244.. -245..)

No of observations







Priority 3 list  

JD (244..-245..)

No of observations

Psi 1 Aur



UU Aur



V465 Cas



AT Dra



g Her



OP Her






XY Lyr








In addition Five Cepheid variables have been observed to obtain folded minima. Eight eclipsing binaries have been observed to obtain minima. Numerous other variables and suspects have been observed from time to time. Futher details can be obtained from West.

Received 1998 March 9; accepted 1998 June 24