UX Arietis belongs to the group of RS CVn stars, that are active binary systems characterized by the occurrence of strong flares in the X, UV and radio wavelengths domain. The system of UX Arietis has an orbital period of 6.44 days and is formed by a G5 V and a K0 IV star. The K star is the more active one and the presence of long-lasting large spots on its surface has been deduced from optical observations (Vogt & Hatzes 1991; Elias et al. 1995).
Radio emission from UX Arietis has been analysed by many authors (see Massi & Chiuderi Drago (1992) and references therein). To get a deeper insight on its flaring activity a monitoring program of this source was started in December 1992 at the Effelsberg 100-m telescope. The instrument was available for this program in the gaps between previously scheduled observations: for this reason our observations span a frequency interval between 1.4 GHz (21 cm) and 43 GHz (7 mm), depending on the scheduled receiver at each observing time. Most of these observations are made at a single frequency, but in a few cases it was possible to obtain some spectra. In this paper we will study the evolution of the flaring spectra, while the full data set will be described and analysed, searching for periodicities in the radio emission, in another paper (Massi et al. 1998). A preliminary presentation of the whole data set has been done by Massi et al. (1996), while results relative to the first six months of observations have been reported by Neidhöfer et al. (1993).
As described in Massi et al. (1996), a minimum is present in the observed flux densities around phase 0.4. This minimum can be easily interpreted in terms of a geometrical shadowing effect strictly linked to the rotation of the system. Since only the stronger emission suffers from the shadowing, while no effect is seen at lower flux levels, it follows that the high intensity emission is more localized (and organized, if one thinks in terms of magnetic field) than the low intensity one. This implies that the source is composed by a compact region, close to the star's surface, where the most intense emission mainly takes place, and by an extended halo, too large to be obscured, where the emission survives for a longer time and smoothly fades. This picture is in agreement with the conclusions derived from VLBI observations by Mutel et al. (1985) and with the model proposed by Franciosini & Chiuderi Drago (1995). Since the shadowing in the emission does not last very long (a small fraction of the orbital period) the compact region must have a typical size smaller than the stellar disk.
The presence of a geometric shadowing of the emission prevents the possibility of studying the evolution of a single flare on time scales longer than a few days. The modulation induced by the stellar rotation affects in fact the observed emission, reducing the flux density when the source is obscured. This implies that it is very difficult to have the opportunity of following the true evolution of the flaring emission from the flare onset to the end of the decaying phase. For this reason and for the way our data have been sampled, either the decaying or the rising phases of different flares are present in our data.
The decaying phase of flares has already been interpreted by Chiuderi Drago & Franciosini (1993) and by Franciosini & Chiuderi Drago (1995) as due to the temporal evolution of a population of non-thermal electrons, trapped in a magnetic loop and undergoing synchrotron and collision energy losses. In their model the injection of relativistic particles was assumed to be instantaneous, therefore the rising phase of the emission was not considered.
This paper is devoted to the analysis of the flare rising phase. For this purpose we have included in the above model a continuous injection of relativistic electrons, which makes the emission increase in spite of losses. The paper is organized as follows. The observed flare spectra are presented in Sect. 2; in Sect. 3 the time-dependent energy distribution of relativistic electrons is derived and in Sect. 4 the model results are compared with the observations. Conclusions are presented in Sect. 5.
© European Southern Observatory (ESO) 1998
Online publication: April 28, 1998