tokens: the intensity weighted first moment (i.e., the centroid) of the star. We measure the total flux using a 6 pixel (≈504) aperture radius. We weight pixels that lie partially within this circular aperture by the fraction of the pixel that lies within the ideal circular aperture. We also adopt an aperture correction to correct for the stellar flux that falls outside of our aperture. The typical size of an aperture correction is 0.04 magnitudes, but can vary from night to night depending on atmospheric conditions (predominantly seeing). In Table 2 , we provide the corrected photometry for LP 661-13 across all telescopes and both seasons of data. Table 2 MEarth Photometry for LP 661-13 (This table is available in its entirety in machine-readable form.) Only a portion of this table is shown here to demonstrate its form and content. A machine-readable version of the full table is available. There is a marginal Roentgan Satellite ( ROSAT ) detection in the ROSAT faint source catalog (Voges et al. 2000 ). The potential source has a count rate of 3 ± 1 × 10 −2 counts per second. This corresponds to a flux density of approximately 1.70 ± 0.56 × 10 −13 erg cm −2 s −1 . While this may be interesting to assess the significance of the X-ray emission with stellar activity and radius inflation, further X-ray data to confirm this potential detection are needed. 3.3. Astrometry With two seasons of photometric data from the MEarth observatory, we are able to utilize the MEarth astrometric pipeline to measure the trigonometric distance to LP 661-13. The astrometric pipeline is described in detail in Dittmann et al. ( 2014 ), but we summarize it here. We eliminate nights for which the full width at half maximum for the image is greater than 5 pixels (approximately 3.5 arcseconds), or the average ellipticity of the target stars is greater than 0.5 (typically due to wind shake). We also eliminate images for which the pointing on the image is more than 15 pixels discrepant from the master MEarth image. We fit each image with a linear model in both the x and y coordinates, which allows for translation, rotation, shearing, and pixel scale variations. We also include a separate constant offset between data taken on opposite sides of the meridian, to accommodate the image flip that results from the telescopes’ German equatorial mount. We choose reference stars that lie within 600 pixels of the target star to avoid higher order effects at the edges of the CCD. We fit our astrometry iteratively, first fitting the linear plate model to each frame and then fitting for the proper motion and parallax of each star until convergence. Internal errors are estimated through a residual permutation. This provides the advantage that it preserves time-correlated noise in our error assessment but tends to be a weaker way to estimate errors for series with few data points. However, we have over 9000 data points and do not believe our astrometric error bars are underestimated. We perform
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