# KELT-25b & KELT-26b: A Hot Jupiter and a Substellar Companion Transiting
Young A-Stars Observed by TESS111This paper includes data gathered with the
6.5 meter Magellan Telescopes located at Las Campanas Observatory, Chile.

Romy Rodríguez Martínez Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA  B. Scott Gaudi Department of Astronomy, The Ohio State University, 140 West 18th Avenue, Columbus, OH 43210, USA  Joseph E. Rodriguez Center for Astrophysics | Harvard & Smithsonian, 60 Garden St, Cambridge, MA 02138, USA  George Zhou Center for Astrophysics | Harvard & Smithsonian, 60 Garden St, Cambridge, MA 02138, USA  Jonathan Labadie-Bartz Instituto de Astronomia, Geofísica e Ciencias Atmosféricas, Universidade de Sào Paulo, Rua do Matão 1226, Cidade Universitãria, Sáo Paulo, SP 05508-900, Brazil  Samuel N. Quinn Center for Astrophysics | Harvard & Smithsonian, 60 Garden St, Cambridge, MA 02138, USA  Kaloyan Penev Department of Physics, The University of Texas at Dallas, 800 West Campbell Road, Richardson, TX 75080-3021 USA  Thiam-Guan Tan Perth Exoplanet Survey Telescope  David W. Latham Center for Astrophysics | Harvard & Smithsonian, 60 Garden St, Cambridge, MA 02138, USA  Leonardo A. Paredes  Department of Physics and Astronomy, Georgia State University, Atlanta, GA 30302, USA  John F. Kielkopf Department of Physics and Astronomy, University of Louisville, Louisville, KY 40292, USA  Brett Addison University of Southern Queensland, Centre for Astrophysics, West Street, Toowoomba, QLD 4350 Australia  Duncan J. Wright  University of Southern Queensland, Centre for Astrophysics, West Street, Toowoomba, QLD 4350 Australia  Johanna Teske  NASA Hubble Fellow  The Observatories of the Carnegie Institution for Science, 813 Santa Barbara St., Pasadena, CA 91101, USA  Steve B. Howell NASA Ames Research Center, Moffett Field, CA, 94035, USA  David Ciardi  Caltech IPAC – NASA Exoplanet Science Institute 1200 E. California Ave, Pasadena, CA 91125, USA  Carl Ziegler Dunlap Institute for Astronomy and Astrophysics, University of Toronto, Ontario M5S 3H4, Canada  Keivan G. Stassun Department of Physics and Astronomy, Vanderbilt University, Nashville, TN 37235, USA  Department of Physics, Fisk University, 1000 17th Avenue North, Nashville, TN 37208, USA  Marshall C. Johnson Las Cumbres Observatory, 6740 Cortona Drive, Suite 102, Goleta, CA 93117, USA  Jason D. Eastman Center for Astrophysics | Harvard & Smithsonian, 60 Garden St, Cambridge, MA 02138, USA  Robert J. Siverd Gemini Observatory, Northern Operations Center, 670 N. A’ohoku Place, Hilo, HI 96720, USA  Thomas G. Beatty Department of Astronomy and Steward Observatory, University of Arizona, Tucson, AZ 85721, USA  Luke Bouma Department of Astrophysical Sciences, Princeton University, 4 Ivy Lane, Princeton, NJ, 08544, USA  Timothy Bedding  School of Physics, Sydney Institute for Astronomy (SIfA), The University of Sydney, NSW 2006, Australia  Joshua Pepper Department of Physics, Lehigh University, 16 Memorial Drive East, Bethlehem, PA 18015, USA  Joshua Winn Department of Astrophysical Sciences, Princeton University, 4 Ivy Lane, Princeton, NJ, 08544, USA  Michael B. Lund Caltech IPAC – NASA Exoplanet Science Institute 1200 E. California Ave, Pasadena, CA 91125, USA  Steven Villanueva Jr. Pappalardo Fellow  Department of Physics and Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA  Daniel J. Stevens Eberly Fellow  Department of Astronomy & Astrophysics, The Pennsylvania State University, 525 Davey Lab, University Park, PA 16802, USA  Center for Exoplanets and Habitable Worlds, The Pennsylvania State University, 525 Davey Lab, University Park, PA 16802, USA  Eric L. N. Jensen Department of Physics and Astronomy, Swarthmore College, Swarthmore, PA 19081, USA  Coleman Kilby  Department of Physics and Astronomy, Swarthmore College, Swarthmore, PA 19081, USA  Jeffrey D. Crane Department of Terrestrial Magnetism, Carnegie Institution for Science, 5241 Broad Branch Road, NW, Washington, DC 20015, USA  The Observatories of the Carnegie Institution for Science, 813 Santa Barbara St., Pasadena, CA 91101, USA  Andrei Tokovinin  Cerro Tololo Inter-American Observatory, Casilla 603, La Serena, Chile  Mark E. Everett National Optical Astronomy Observatory, 950 North Cherry Avenue, Tucson, AZ 85719, USA  C.G. Tinney Exoplanetary Science at UNSW, School of Physics, UNSW Sydney, NSW 2052, Australia  Michael Fausnaugh Department of Physics and Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA  David H. Cohen Department of Physics and Astronomy, Swarthmore College, Swarthmore, PA 19081, USA  Daniel Bayliss  Department of Physics, University of Warwick, Gibbet Hill Road, Coventry CV4 7AL, UK  Centre for Exoplanets and Habitability, University of Warwick, Gibbet Hill Road, Coventry CV4 7AL, UK  Allyson Bieryla Center for Astrophysics | Harvard & Smithsonian, 60 Garden St, Cambridge, MA 02138, USA  Phillip A. Cargile Center for Astrophysics | Harvard & Smithsonian, 60 Garden St, Cambridge, MA 02138, USA  Karen A. Collins Center for Astrophysics | Harvard & Smithsonian, 60 Garden St, Cambridge, MA 02138, USA  Dennis M. Conti American Association of Variable Star Observers, 49 Bay State Road, Cambridge, MA 02138, USA  Knicole D. Colón Exoplanets and Stellar Astrophysics Laboratory, Code 667, NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA  Ivan A. Curtis  Ivan Curtis Private Observatory  D. L. Depoy  George P. and Cynthia Woods Mitchell Institute for Fundamental Physics and Astronomy, Texas A&M University, College Station, TX77843 USA  Department of Physics and Astronomy, Texas A&M university, College Station, TX 77843 USA  Phil Evans El Sauce Observatory, Chile  Dax L. Feliz Department of Physics and Astronomy, Vanderbilt University, Nashville, TN 37235, USA  Joao Gregorio  Atalaia Group & CROW Observatory, Portalegre, Portugal  Jason Rothenberg  Department of Physics & Astronomy, University of Wyoming, 1000 E University Ave, Dept 3905, Laramie, WY 82071, USA  David J. James Center for Astrophysics | Harvard & Smithsonian, 60 Garden St, Cambridge, MA 02138, USA  Black Hole Initiative at Harvard University, 20 Garden Street, Cambridge, MA 02138, USA  Michael D. Joner Department of Physics and Astronomy, Brigham Young University, Provo, UT 84602, USA  Rudolf B. Kuhn South African Astronomical Observatory, PO Box 9, Observatory, 7935, Cape Town, South Africa  Southern African Large Telescope, PO Box 9, Observatory, 7935, Cape Town, South Africa  Mark Manner  Spot Observatory, Nashville, TN 37206, USA  Somayeh Khakpash Department of Physics, Lehigh University, 16 Memorial Drive East, Bethlehem, PA 18015, USA  Jennifer L. Marshall  George P. and Cynthia Woods Mitchell Institute for Fundamental Physics and Astronomy, Texas A&M University, College Station, TX77843 USA  Department of Physics and Astronomy, Texas A&M university, College Station, TX 77843 USA  Kim K. McLeod Department of Astronomy, Wellesley College, Wellesley, MA 02481, USA  Matthew T. Penny Department of Physics and Astronomy, Louisiana State University, Baton Rouge, LA 70803 USA  Phillip A. Reed Department of Physical Sciences, Kutztown University, Kutztown, PA 19530, USA  Howard M. Relles  Center for Astrophysics | Harvard & Smithsonian, 60 Garden St, Cambridge, MA 02138, USA  Denise C. Stephens  Department of Physics and Astronomy, Brigham Young University, Provo, UT 84602, USA  Chris Stockdale Hazelwood Observatory, Churchill, Victoria, Australia  Mark Trueblood  Winer Observatory, PO Box 797, Sonoita, AZ 85637, USA  Pat Trueblood  Winer Observatory, PO Box 797, Sonoita, AZ 85637, USA  Xinyu Yao Department of Physics, Lehigh University, 16 Memorial Drive East, Bethlehem, PA 18015, USA  Roberto Zambelli  Società Astronomica Lunae, Italy  Roland Vanderspek  Department of Physics and Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA  Sara Seager Department of Physics and Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA  Department of Earth, Atmospheric and Planetary Sciences, Massachusetts Institute of Technology, Cambridge, MA 02139, USA  Department of Aeronautics and Astronautics, MIT, 77 Massachusetts Avenue, Cambridge, MA 02139, USA  Jon M. Jenkins  NASA Ames Research Center, Moffett Field, CA, 94035, USA  Todd J. Henry  Department of Physics and Astronomy, Georgia State University, Atlanta, GA 30302, USA  Hodari-Sadiki James  Department of Physics and Astronomy, Georgia State University, Atlanta, GA 30302, USA  Wei-Chun Jao  Department of Physics and Astronomy, Georgia State University, Atlanta, GA 30302, USA  Sharon Xuesong Wang The Observatories of the Carnegie Institution for Science, 813 Santa Barbara St., Pasadena, CA 91101, USA  Paul Butler  Department of Terrestrial Magnetism, Carnegie Institution for Science, 5241 Broad Branch Road, NW, Washington, DC 20015, USA  Ian Thompson  The Observatories of the Carnegie Institution for Science, 813 Santa Barbara St., Pasadena, CA 91101, USA  Stephen Shectman  The Observatories of the Carnegie Institution for Science, 813 Santa Barbara St., Pasadena, CA 91101, USA  Robert Wittenmyer  University of Southern Queensland, Centre for Astrophysics, West Street, Toowoomba, QLD 4350 Australia  Brendan P. Bowler  Department of Astronomy, The University of Texas at Austin, Austin, TX 78712, USA  Jonathan Horner University of Southern Queensland, Centre for Astrophysics, West Street, Toowoomba, QLD 4350 Australia  Stephen R. Kane  Department of Earth Sciences, University of California, Riverside, CA 92521, USA  Matthew W. Mengel  University of Southern Queensland, Centre for Astrophysics, West Street, Toowoomba, QLD 4350 Australia  Timothy D. Morton  Department of Astronomy, University of Florida, 211 Bryant Space Science Center, Gainesville, FL, 32611, USA  Jack Okumura  University of Southern Queensland, Centre for Astrophysics, West Street, Toowoomba, QLD 4350 Australia  Peter Plavchan  George Mason University, 4400 University Drive MS 3F3, Fairfax, VA 22030, USA  Hui Zhang  School of Astronomy and Space Science, Key Laboratory of Modern Astronomy and Astrophysics in Ministry of Education, Nanjing University, Nanjing 210046, Jiangsu, China  Nicholas J. Scott NASA Ames Research Center, Moffett Field, CA, 94035, USA  Rachel A. Matson NASA Ames Research Center, Moffett Field, CA, 94035, USA  Andrew W. Mann Department of Physics and Astronomy, University of North Carolina at Chapel Hill, Chapel Hill, NC 27599, USA  Diana Dragomir Department of Physics and Astronomy, University of New Mexico, 1919 Lomas Blvd NE, Albuquerque, NM 87131, USA  Max Günther  Juan Carlos Torres Fellow  Department of Physics and Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA  Eric B. Ting  NASA Ames Research Center, Moffett Field, CA, 94035, USA  Ana Glidden  Department of Physics and Kavli Institute for Astrophysics and Space Research, Massachusetts Institute of Technology, Cambridge, MA 02139, USA  Elisa V. Quintana  Exoplanets and Stellar Astrophysics Laboratory, Code 667, NASA Goddard Space Flight Center, Greenbelt, MD 20771, USA

###### Abstract

We present the discoveries of KELT-25b (TIC 65412605, TOI-626.01) and KELT-26b
(TIC 160708862, TOI-1337.01), two transiting companions orbiting relatively
bright, early A-stars. The transit signals were initially detected by the KELT
survey, and subsequently confirmed by TESS photometry. KELT-25b is on a
4.40-day orbit around the V = 9.66 star CD-24 5016 ($T_{\rm
eff}=8280^{+440}_{-180}$ K, $M_{\star}$ = $2.18^{+0.12}_{-0.11}$ $M_{\odot}$),
while KELT-26b is on a 3.34-day orbit around the V = 9.95 star HD 134004
($T_{\rm eff}$ =$8640^{+500}_{-240}$ K, $M_{\star}$ = $1.93^{+0.14}_{-0.16}$
$M_{\odot}$), which is likely an Am star. We have confirmed the sub-stellar
nature of both companions through detailed characterization of each system
using ground-based and TESS photometry, radial velocity measurements, Doppler
Tomography, and high-resolution imaging. For KELT-25, we determine a companion
radius of $R_{\rm P}$ = $1.64^{+0.039}_{-0.043}$ $R_{\rm J}$, and a 3-sigma
upper limit on the companion’s mass of $\sim 64~{}M_{\rm J}$. For KELT-26b, we
infer a planetary mass and radius of $M_{\rm P}$ = $1.41^{+0.43}_{-0.51}$
$M_{\rm J}$ and $R_{\rm P}$ = $1.940^{+0.060}_{-0.058}$ $R_{\rm J}$. From
Doppler Tomographic observations, we find KELT-26b to reside in a highly
misaligned orbit. This conclusion is weakly corroborated by a subtle asymmetry
in the transit light curve from the TESS data. KELT-25b appears to be in a
well-aligned, prograde orbit, and the system is likely a member of a cluster
or moving group.

planetary systems, planets and satellites: detection, stars: individual
(KELT-25,KELT-26)

††software: EXOFASTv2 (Eastman et al. 2013; Eastman:2017), AstroImageJ
(Collins et al. 2017), SPC (Buchhave et al. 2010)††facilities: FLWO 1.5m
(Tillinghast Reflector Echelle Spectrograph, TRES); Kilodegree Extremely
Little Telescope (KELT); MINiature Exoplanet Radial Velocity Array (MINERVA);
Las Cumbres Observatory at Tenerife (LCO TFN); University of Louisville Manner
Telescope (ULMT, Mt. Lemmon); KeplerCam (FLWO 1.2m); Stacja Obserwacji
Tranzytów Egzoplanet w Suwałkach (SOTES); CROW Observatory; Koyama
Astronomical Observatory (KAO); Gemini-South Zorro

##  1 Introduction

The field of exoplanets has grown tremendously since the first detection of a
transiting exoplanet around a bright star (Charbonneau et al. 2000; Henry et
al. 2000) two decades ago. Thousands of planets222Almost 4100, as of 2019
(https://exoplanetarchive.ipac.caltech.edu/) have been validated orbiting
stars of almost every spectral type, and span a wide range of masses, orbits
and likely compositions. Our knowledge of the demographics of short-period
giant exoplanets quickly expanded with the advent of dedicated wide-field
transit surveys from the ground such as The Hungarian-made Automated Telescope
Network (HATNet; Bakos et al. 2007), the HATSouth survey (Bakos et al. 2013),
the Wide Angle Search for Planets (WASP/SuperWASP; Pollacco et al. 2006;
Collier Cameron et al. 2009), the Qatar Exoplanet Survey (QES; Alsubai et al.
2011), XO (McCullough et al. 2005), the Trans-Atlantic Exoplanet Survey
network (TrES; Alonso et al. 2007), and the Kilodegree Extremely Little
Telescope (KELT; Pepper et al. 2007, 2012) to mention a few.

The Kepler space telescope (Borucki et al. 2010) provided the first
statistical survey of a large number of transiting planets over a broad region
of radius and orbital period. Kepler transformed our understanding of the
population of relatively short period ($P\lesssim 100$ days) planets. Later,
the next generation of dedicated wide-field surveys came online, such as the
Next Generation Transit Survey (NGTS; Bayliss et al. 2018) and the Multi-site
All-Sky CAmeRA (MASCARA: Lesage et al. 2014). In addition, based largely on
arguments presented in Gould et al. (2003) and Blake et al. (2007), targeted
ground-based surveys were initiated such as MEarth (Nutzman & Charbonneau
2008; Charbonneau et al. 2009; Berta et al. 2012), TRAPPIST (Gillon et al.
2017), and SPECULOOS (Delrez et al. 2018), which primarily concentrate their
efforts on the search for exoplanets around low-mass stars.

Massive and hot stars are typically avoided by transit surveys because planets
around them induce weaker photometric signals (as a result of the lower
planet-star radius and mass ratios). More importantly, the host stars have
fewer spectral lines, and the lines they do have are significantly broadened
by their fast rotation. A large fraction of stars above the Kraft break
($T_{\rm eff}\gtrsim 6250\rm{K}$; Kraft 1967) are observed to rotate
significantly faster ($v\sin{I_{*}}$$\geq$ 10 km $\rm s^{-1}$) than cooler
stars. This is because stars with $T_{\rm eff}\gtrsim 6250\rm{K}$ have
essentially no convective envelopes, and thus do not slow down due to magnetic
braking over their evolution. They therefore essentially retain their high
primordial spin rates. The relative paucity of spectral features coupled with
their faster rotation rates make candidate planetary companions transiting hot
stars harder to confirm using the radial velocity method (see e.g., Johnson et
al. 2018 and Dholakia et al. 2019). To attempt to circumvent the challenges in
searching for planets around massive stars on the main sequence, a number of
radial velocity surveys have studied “retired A-stars” – cooler, evolved stars
that have moved off the main sequence. Such surveys have yielded a number of
discoveries (e.g., Johnson et al. 2007, 2011), but have only been able to
sample a relatively small number of target stars due to the focussed nature of
radial velocity studies.

Despite the observational challenges posed by hot stars, they provide
opportunities to study the most massive, highly-irradiated, close-in planets –
in particular the relatively new category of “Ultra Hot Jupiters” (Collier
Cameron et al. 2010a). Ultra Hot Jupiters provide a unique opportunity for the
detailed atmospheric characterization of highly irradiated giant planets. They
have high equilibrium temperatures, allowing one to probe extreme conditions
that are not present in the Solar System. In particular, many of these Ultra
Hot Jupiters have day-side temperatures that are hot enough to disassociate
all molecules, leaving atomic metals as the dominant species on the day side
(Gaudi et al. 2017). Furthermore, the exceptionally high day-side temperatures
of these planets imply that they are typically close to thermodynamic
equilibrium (Kitzmann et al. 2018), making the interpretation of observations
much simpler, and resulting in atmospheres that are likely quite different
than typical hot Jupiters (Kitzmann et al. 2018; Bell & Cowan 2018; Lothringer
et al. 2018). In addition, A-type hosts present opportunities to test the
effects of host star mass and binarity, and short evolutionary timescales, on
giant planet formation, evolution, and engulfment by their host stars. For
example, a recent study found that most giant planets around A-type stars are
eventually engulfed by their host stars (Stephan et al. 2018), and a related
study determined the observable effects of engulfment (Stephan et al. 2019).
In general, the high temperatures and high scale heights of Ultra Hot Jupiters
make follow-up observations much easier, allowing one to test models of hot
Jupiter atmospheres (e.g., Cauley et al. 2019; Hoeijmakers et al. 2019).

Studying the planetary population of massive (A) stars is also interesting for
other reasons. First, the large amount of high-energy radiation they emit
helps to test theories of planet atmospheric evaporation. Second, the fact
that A stars tend to be rapidly rotating means that they are typically oblate,
resulting in significant gravity darkening, which can be used to estimate the
true (not just projected) spin-orbit angle of the planet (Barnes 2009). The
oblateness of the host star, combined with certain orbital alignments of the
planet, can result in relatively short precession times of the planetary orbit
(Johnson et al. 2015). Finally, their short lifetimes imply that the lifetimes
of close-in planets orbiting such stars is likely to be much shorter than that
of such planets orbiting low-mass stars (Stephan et al. 2018).

Royer et al. 2007 found that typical rotational values ($v\sin{I_{*}}$) are
greater than 100 km s-1 for stars in the the B9-F2 spectral range (for
reference, the Sun’s rotational velocity is only 1.6 $\pm$ 0.3 km s-1;
Pavlenko et al. 2012). Because of the difficulties in confirming planets
orbiting these fast rotators, KELT routinely uses a combination of radial
velocity (RV) and Doppler Tomographic (DT) observations (Collier Cameron et
al. 2010b). DT measures the distortion in the spectral lines of a star caused
by the transiting planet blocking the light from the star with different
projected velocities as it crosses the disk of the star. Doppler tomographic
observations can help confirm that the planet transits the target star and is
not, for example, a signal from a nearby eclipsing binary, although confirming
the planetary nature of the occultor also requires an appropriately stringent
upper limit on its mass (see, e.g., Bieryla et al. 2015 for a discussion of
validating planets orbiting rapidly rotating stars using Doppler Tomography).
The advantage of this technique is that it is better suited to faster stellar
rotations, thereby providing an alternative way to confirm planets around hot
stars. Many discovery papers have demonstrated that A-type stars are the most
suitable for the measurement of the Rossiter-McLaughlin (RM) effect because
they have the optimal combination of $R_{p}/R_{s}$ and rotation – two
parameters to which the RM signal is proportional (see also Gaudi & Winn 2007,
and Table 8 of this paper).

Figure 1: Discovery light curves for KELT-25 b (Top) and KELT-26 b (Bottom)
from the KELT-South telescope. The light curves are phase-folded on the
preliminary orbital periods of 4.401093 days (KELT-25 b) and 3.344886 days
(KELT-26 b). The blue points represent the data and the black points are the
data binned at intervals of 0.01 in phase.

With Doppler Tomography, we can not only validate new planets, but also
determine their projected spin-orbit angles, $\lambda$. From the current
sample of over 200
systems333https://www.astro.keele.ac.uk/jkt/tepcat/obliquity.html with
measured $\lambda$, a dichotomy in stellar temperature has emerged: giant
planets around cool stars ($T_{\rm eff}$ ¡ 6200 K) have lower obliquities than
those around hot stars (Schlaufman 2010; Winn 2010; Dawson et al. 2015). In
addition, the most massive planets ($M_{p}>3M_{J}$) tend to have lower spin-
orbit angles (Hébrard et al. 2010). These results are important because they
help constrain planet formation and migration mechanisms for hot Jupiters. One
area of ongoing research is the question of whether hot Jupiters form “in
situ” or at larger separations from their host and then migrate into their
present observed locations via planet-disk or planet-planet dynamical or
secular interactions, or secular interactions with a distant planetary, brown
dwarf, or stellar companion (e.g., Dawson & Johnson 2018). Measurements of
$\lambda$ provide insights into these formation channels since different
theories predict different values of the distribution of spin-orbit angles.

Giant planets at close separations are highly irradiated by their host stars,
and this intense radiation can significantly impact their properties. One
possible consequence is radius inflation: it has been observed that hot
Jupiters’ radii appear to increase with increasing incident radiation from the
host star (Demory & Seager 2011; Howell et al. 2019), although it is not clear
whether this is caused by the radiation re-inflating the planet, or because it
simply slows the cooling and contraction process of the planets, which may be
hot and thus inflated upon formation.

In this paper, we present the discoveries of a sub-stellar companion and
likely planet (KELT-25 b) and a hot Jupiter (KELT-26 b 444Whilst we were
writing this manuscript, we noted a paper on arXiv in which WASP announced the
discovery of WASP-178b (Hellier et al. 2019). We had already collected and
analyzed the data needed to confirm KELT-26 b, and other than to confirm that
KELT-26 and WASP-178 had the same coordinates, we did not read or use the
results of Hellier et al. (2019) in any way. We therefore claim an independent
discovery (regardless of whether or not they are the same planet). If they are
the same planet, we do not, of course, claim to be the first to have made the
detection.) both orbiting bright, early A-stars, first identified as
candidates in KELT data and subsequently observed by the TESS mission.
Although the TESS mission’s main science driver is to measure small planets,
simulations have estimated yields of thousands of giant planets, which include
Jupiter-sized planets around bright stars (e.g. Barclay et al. 2018). TESS’s
expected yield complements the discoveries made by KELT and other ground-based
surveys and advances our understanding of giant planets around hot stars.

Table 1: Literature Properties for KELT-25 and KELT-26

Parameter | Description | KELT-25 | KELT-26 | Source
---|---|---|---|---
Other identifiers |  | CD-24 5016 | HD 134004 |
|  | TIC 65412605 | TIC 160708862 |
|  | TYC 6528-1639-1 | TYC 7829-2324-1 |
$\alpha_{J2000}$ | Right Ascension (RA) | 07:12:29.55004 | 15:09:04.89304 | 1
$\delta_{J2000}$ | Declination (Dec) | -24:57:12.82193 | -42:42:17.78910 | 1
$l$ | Galactic Longitude | $237.5346109\arcdeg$ | $328.1938368\arcdeg$ | 1
$b$ | Galactic Latitude | $-6.79674034\arcdeg$ | $+13.3150904\arcdeg$ | 1
BT |  Tycho BT mag. |  9.841$\pm$0.019 |  10.083 $\pm$ 0.028 | 2
VT |  Tycho VT mag. |  9.655$\pm$0.018 |  9.961 $\pm$ 0.033 | 2
${\rm G}$ |  Gaia $G$ mag. |  9.5960$\pm$0.0003 |  9.912 $\pm$ 0.020 | 1
J | 2MASS J mag. |  9.362 $\pm$ 0.03 |  9.775 $\pm$ 0.030 | 3
H | 2MASS H mag. |  9.273 $\pm$ 0.02 |  9.735 $\pm$ 0.020 | 3
KS |  2MASS ${\rm K_{S}}$ mag. |  9.248 $\pm$ 0.02 |  9.703 $\pm$ 0.020 | 3
WISE1 |  WISE1 mag. |  9.213 $\pm$ 0.022 |  9.670 $\pm$ 0.030 | 4
WISE2 |  WISE2 mag. |  9.245 $\pm$ 0.02 |  9.683 $\pm$ 0.030 | 4
WISE3 |  WISE3 mag. |  9.302 $\pm$ 0.033 |  9.645 $\pm$ 0.043 | 4
$\mu_{\alpha}$ | Gaia DR2 proper motion | $-2.276\pm 0.06$ | $-10.011\pm 0.122$ | 1
|  in RA (mas yr-1) |  |  |
$\mu_{\delta}$ | Gaia DR2 proper motion | $0.338\pm 0.075$ | $-5.652\pm 0.097$ | 1
|  in DEC (mas yr-1) |  |  |
$\pi^{\dagger}$ | Gaia Parallax (mas) |  2.342$\pm$ 0.043† |  2.394 $\pm$ 0.060† | 1
$RV$ | Absolute radial |  35.472$\pm 1.011$ | $-24.140\pm 0.045$ | This work
|  velocity ( km s-1) |  |  |
$d$ | Distance (pc) | $427.0\pm 7.8$ | $417.7\pm 10.5$ | 1
$U^{*}$ |  Space Velocity ( km s-1) | $-13.40\pm 0.58$ | $-24.63\pm 0.37$ |  §5.1
$V$ |  Space Velocity ( km s-1) | $-13.83\pm 0.84$ | $6.75\pm 0.58$ |  §5.1
$W$ |  Space Velocity ( km s-1) | $-1.24\pm 0.19$ | $0.78\pm 0.20$ |  §5.1

NOTES: $\dagger$ Parallaxes here are corrected for the 82 $\mu$as offset
reported in Stassun & Torres (2018a).
${}^{*}U$ is in the direction of the Galactic center.
References: 1Gaia Collaboration et al. (2018),2Høg et al. (2000),3Cutri et al.
(2003), 4Zacharias et al. (2017)

Table 2: Follow-up photometric observations of KELT-25b and KELT-26b  Target | Observatory | Date (UT) | Diameter (m) | Filter | FOV | Pixel Scale | Exposure (s) |
---|---|---|---|---|---|---|---|---
KELT-25b | PEST | 2019 January 18 | 0.3 | $Rc$ | 31$\arcmin$ $\times$ 21$\arcmin$ | 1.2$\arcsec$ | 30 |
KELT-25b | TESS |  2019 January 7 - 2019 February 7 | 0.105 |  TESS (600-1000nm) | 24∘ $\times$ 24∘ | 21$\arcsec$ | 1800 |
KELT-26b | PEST | 2016 August 26 | 0.3 | $I$ | 31$\arcmin$ $\times$ 21$\arcmin$ | 1.2$\arcsec$ | 30 |
KELT-26b |  Mt. Kent CDK700 | 2018 March 20 | 0.6858 | $r^{\prime}$ | 27.3$\arcmin$ $\times$ 27.3$\arcmin$ | 0.4$\arcsec$ | 65 |
KELT-26b | TESS |  2019 April 22 - 2019 May 20 | 0.105 |  TESS (600-1000nm) | 24∘ $\times$ 24∘ | 21$\arcsec$ | 1800 |

##  2 Discovery and Follow-up Observations

###  2.1 KELT Discovery

One survey that has contributed significantly to the discovery and study of
Ultra Hot Jupiters is the Kilodegree Extremely Little Telescope survey
(KELT555https://keltsurvey.org; Pepper et al. 2007). KELT observes $\sim$85%
of the sky targeting bright stars in the magnitude range 7.5 ¡ V ¡ 12, filling
in the gap between radial velocity surveys and other transit surveys, which
generally focus on brighter and fainter stars, respectively. KELT consists of
two observatories: KELT-South (KS, Pepper et al. 2012), located in Sutherland,
South Africa, which surveys most of the southern hemisphere, and KELT-North
(KN), which observes the northern hemisphere from Sonoita, Arizona. KN and KS
have been successful at finding giant planets around hot stars, discovering 24
planets, of which 18 – including the ones confirmed in this paper – orbit A
and F stars.

Using two separate 42mm-aperture telescopes, the KELT survey observes over 85%
of the entire sky with a 20$-$30 minute cadence. Each observing site has a
Mamiya 645 80mm f/1.9 42mm lens with a 4k$\times$4k Apogee CCD on a Paramount
ME mount. The KELT telescopes have a $23\arcsec$ pixel scale and a
$26\arcdeg\times 26\arcdeg$ field of view. The original goal of the KELT
survey was to discover hot Jupiters orbiting bright ($7<V[mag]<12$) host
stars, which are amenable to detailed characterization through transmission or
eclipse spectroscopy. More recently, KELT has become a significant contributor
to the understanding of planets around early-type stars, specifically with the
discovery of 6 transiting hot Jupiter around A-stars, including the planets
presented here (Zhou et al. 2016; Gaudi et al. 2017; Lund et al. 2017; Johnson
et al. 2018; Siverd et al. 2018).

The likely planetary companion orbiting TIC 65412605 (hereafter KELT-25 b) was
first identified as a planet candidate following the reduction of KELT-South
field KS35. KELT-25 is located at $\alpha_{\rm J2000}=$ 07h 12m 29$\fs$55,
$\delta_{\rm J2000}=$ $-$24$\arcdeg$57$\arcmin$12$\arcsec$82 (Gaia
Collaboration et al. 2018) and the KELT-South survey field 35 is centered at
$\alpha_{\rm J2000}=07^{h}40^{m}12\fs 0$, $\delta_{\rm J2000}=-20\arcdeg
00\arcmin 00\farcs 0$. KS35 was observed 2,860 times from UT 2013 May 10 until
UT 2017 October 1. The image reduction and candidate selection pipeline are
described in detail in Siverd et al. (2012) and Kuhn et al. (2016), and are
briefly summarized here. To reduce the raw survey images, KELT uses a
modified, image subtraction pipeline based on the ISIS software (Alard &
Lupton 1998; Alard 2000). The list of sources identified in the KELT fields
are then cross-matched to the Tycho-2 (Høg et al. 2000) and UCAC4 (Zacharias
et al. 2013) catalogs to obtain their proper motions. We then implement
reduced proper motion (RPM) cuts to identify and exclude giants before
searching for transit signals (Gould & Morgan 2003; Collier Cameron et al.
2007). Finally, we search for transit-like features in all stars that passed
the RPM cuts using the Box-fitting Least Squares (BLS) algorithm (Kovács et
al. 2002).

The planet orbiting TIC 160708862 (hereafter KELT-26 b) is located at
$\alpha_{J2000}=$ 15h 09m 04$\fs$89304, $\delta_{J2000}=$ $-$42$\arcdeg$
42$\arcmin$17$\arcsec$78910 (Gaia Collaboration et al. 2018) in the KELT-South
KS37 field, which is centered at $\alpha_{J2000}=$ 15h 07m 12$\fs$00,
$\delta_{J2000}$ -53$\arcdeg$00$\arcmin$00$\arcsec$00\. This field was
observed 2,085 times from UT 2013 September 5 until UT 2015 September 11.

Both companions were identified as candidates from a periodicity search using
the BLS algorithm. KELT-25 b was initially identified with a BLS orbital
period of 4.40 days and a transit depth of 0.66%, while KELT-26 b was detected
with a period of 3.34 days and a 1.1% transit depth. The phase-folded KELT
discovery light curves are shown in Figure 1. See Table 1 for literature
information about the stellar hosts KELT-25 and KELT-26.

###  2.2 TESS Photometry

The NASA _Transiting Exoplanet Survey Satellite_ (TESS, Ricker et al. 2015)
was launched on April 18, 2018 with the primary goal of discovering and
characterizing small ($R_{p}<4R_{\oplus}$) exoplanets around bright, nearby
stars. Presently, dozens of exoplanets have already been validated, while
almost a thousand candidates await confirmation. The confirmed systems include
a few giant planets on short-period orbits (Brahm et al. 2018; Nielsen et al.
2019; Rodriguez et al. 2019a).

KELT-25 was observed in Sector 7 by Camera 2 of the TESS spacecraft between UT
2019 January 7 and February 7. We made use of the 30-minute cadence Full Frame
Images (FFI) made available by the Mikulski Archive for Space Telescopes
(MAST) and calibrated by the Science Processing Operations Center (SPOC)
pipeline (Jenkins et al. 2016). Cutouts of $10\times 10$ pixels were extracted
from the FFIs around each target star via the MAST _TESScut_ tool, and
aperture photometry was performed using the _lightkurve_ package (Lightkurve
Collaboration et al. 2018; Barentsen et al. 2019). The target star aperture
encompassing pixels around each target star with fluxes higher than 90% of
pixels in the cutout, and pixels nearby that do not encompass adjacent stars
were used in the evaluation of the background flux. We accounted for dilution
within the TESS aperture by computing for and removing the light contribution
of known nearby stars, as per their TESS band magnitudes in TIC.

Similarly, KELT-26 was observed by TESS between UT 2019 April 22 and May 20
during Sector 11 of the mission. The light curve extraction for KELT-26 was
the same as that described above. Figures 2 and 3 show the raw and reduced
TESS light curves for KELT-25 b and KELT-26 b.

Figure 2: TESS light curve of KELT-25. The Top panel shows the raw light
curve; the middle panel shows the detrended light curve. The Bottom Left panel
shows the transit and best-fit EXOFASTv2 model to the detrended light curve,
phase-folded on the orbital period. The Bottom Right plot shows the region of
the secondary eclipse, which is clearly detected in TESS. Figure 3: TESS light
curve of KELT-26. The Top panel shows the raw light curve; the middle panel
shows the detrended light curve. The Bottom Left panel shows the transit and
best-fit EXOFASTv2 model to the detrended light curve, phase-folded on the
orbital period. The Bottom Right plot shows the region of the secondary
eclipse. While the secondary eclipse is not detected in the TESS data, the
primary phase-folded transit shows evidence for a slight asymmetry, which is
likely real and may be caused by gravity darkening of the star. However, the
primary also shows periodic photometric oscillations at a period that is
nearly commensurate (1:18) with the period of the planet. This variability may
also be causing the slight asymmetry.  Figure 4: The follow-up and TESS light
curves of KELT-25. The KELT-FUN light curves phase-folded to the ephemeris
determined in the global fit (Table 6). Table 2 contains information of all
the KELT-FUN observations. The black points are the relative fluxes, while the
red line shows the EXOFASTv2 model. Figure 5: The follow-up and TESS light
curves of KELT-26. The KELT-FUN light curves phase-folded to the ephemeris
determined in the global fit (Table 6). Table 2 contains information of all
the KELT-FUN observations. The black points are the relative fluxes, while the
red line shows the EXOFASTv2 model.

###  2.3 Ground-based Photometry from the KELT Follow-up Network

In order to confirm that the transit signals are due to bona fide planetary
companions, rule out false positives such as eclipsing binaries, and refine
the transit depth, duration and ephemeris of our candidates, we obtained
photometric observations of KELT-25 b and KELT-26 b from the KELT Follow-Up
Network (KELT-FUN, Collins et al. 2018). Some of the follow-up photometry was
reduced using the AstroImageJ analysis software (Collins et al. 2017). See
Table 2 for technical information about the observatories that followed-up
these systems. The KELT-FUN transits for both systems are shown in Figures 4
and 5.

Figure 6: (Top) Radial velocity measurements for KELT-25 b from TRES (Left)
and KELT-26 b from CHIRON (Right). (Bottom) The radial velocity measurements
are phase-folded to the best determined period by EXOFASTv2. The EXOFASTv2
model is shown in red and the residuals to the best-fit are shown below each
plot. While we measure the reflex Doppler signal from KELT-26 to
$\sim\sigma_{K}/K\simeq 30\%$, and thus are able to constrain the mass of
KELT-26 b, we do not obtain a definitive detection of the reflex Doppler
signal from KELT-25, and thus are only able to provide an upper limit on the
planet mass. Table 3: Relative Out-of-Transit Radial Velocities for KELT-25
from TRES

$\rm{BJD_{TDB}}$ | RV (m s-1) |  $\sigma_{RV}$ (m s-1)
---|---|---
2458509.81953 | 3601.8 | 384.5
2458511.84309 | 1657.2 | 380.6
2458573.61546 | 480.7 | 457.1
2458575.63514 | 991.1 | 588.0
2458588.63382 | 1544.3 | 421.7
2458589.63349 | 54.5 | 452.6
2458591.62656 | 588.2 | 422.2
2458592.62509 | 41.9 | 593.6
2458593.63010 | -2782.7 | 620.8

We note that while the radial velocities have been put on an absolute scale
for the CHIRON data, the uncertainties are only relative to the mean. There is
an additional systematic uncertainty that would affect all of the data points
by the same amount of roughly 40 $\rm m~{}s^{-1}$, due to the uncertainties of
the standard radial velocity stars used to put the radial velocities on an
absolute scale.

Table 4: Relative Out-of-Transit Radial Velocities for KELT-26 from CHIRON  $\rm{BJD_{TDB}}$ | RV (m s-1) |  $\sigma_{RV}$ (m s-1)
---|---|---
2458527.84082 | 132.1 | 54.9
2458532.82920 | -112.1 | 75.8
2458578.72286 | 182.9 | 78.3
2458579.85037 | -39.4 | 53.0
2458580.80996 | -124 | 55.3
2458581.79719 | 68.9 | 45.5
2458582.83253 | -110.5 | 69.3
2458583.76093 | -89 | 75.3
2458596.78738 | -131.7 | 59.3
2458625.76755 | 175.2 | 56.7
2458635.75959 | 115.8 | 97.9
2458636.64304 | -72.5 | 72.2
2458637.66033 | 9.7 | 56.1
2458639.66539 | -120.3 | 46.8
2458649.51090 | 233.2 | 178.2

####  2.3.1 Perth Exoplanet Survey Telescope

The PEST home observatory was built in 2010 and has since helped discover
dozens of exoplanets, including KELT candidates, mostly via the transit
method. It is located in Perth, Australia, and is owned and run by Thiam-Guan
Tan. The instrument is a 0.3m Meade LX200 SCT f/10 and focal reducer yielding
f/5. The camera is a SBIG ST-8XME with multiple filters including $I$, and it
has a FOV of 31$\arcmin$ $\times$ 21$\arcmin$ and an image scale of
1.2$\arcsec$ per pixel. PEST observed a full transit of KELT-25 b on UT 2019
January 18 and a full transit of KELT-26 b in the I filter on UT 2016 July 26.

####  2.3.2 Mt Kent CDK700 Telescope

Photometric follow-up of KELT-26 was taken with the University of Louisville’s
Shared Skies MKO-CDK700 telescope at Mt. Kent Observatory of the University of
Southern Queensland, Australia. The instrument is a 0.7m Planewave corrected
Dall-Kirkham (CDK) telescope with a Nasmyth focus. The telescope was used with
an Apogee U16 CCD camera (Kodak KAF-16801E sensor). The CDK700 telescope
observed a full transit of KELT-26 b, acquiring 165 images in the Sloan
$r^{\prime}$ filter on UT 2018 March 20.

##  3 SPECTROSCOPIC OBSERVATIONS

###  3.1 TRES Spectroscopy of KELT-25

To constrain the planet mass and characterize the host star properties of
KELT-25, we made a series of spectroscopic observations with the Tillinghast
Reflector Echelle Spectrograph (TRES; Fűrész 2008) on the 1.5 m telescope at
the Fred Lawrence Whipple Observatory (FLWO) in Arizona, USA. TRES is a fiber-
fed echelle spectrograph with a resolving power of
$\lambda/\Delta\lambda\equiv R=44,000$ spanning $3850-9100$ Å. A series of
spectra were obtained for KELT-25 over 10 epochs from UT 2019 January 26 to UT
2019 May 20. These observations are reduced via the procedure described in
Buchhave et al. (2012), and relative velocities were obtained via a multi-
order cross correlation against an averaged observed spectral template, as per
Quinn et al. (2012). The relative velocities are listed in Table 3 and plotted
in Figure 6.

To establish the absolute systemic velocity of the system, we cross correlated
the Mg b line order of one of the observed spectra against a synthetic
template, and shifted all other velocities relative to this template.

###  3.2 CHIRON Spectroscopy of KELT-25 and KELT-26

We obtained a series of spectroscopic observations of KELT-25 and KELT-26 with
the CHIRON spectrograph on the SMARTS 1.5 m telescope, located at Cerro Tololo
Inter-American Observatory (CTIO), Chile. CHIRON is a fiber-fed echelle
spectrograph, sliced via an image slicer, yielding a resolving power of
$R=80,000$ over the wavelength range $4100-8700$ Å (Tokovinin et al. 2013).
Wavelength calibrations are provided by bracketing Thorium-Argon (Th-Ar) arc
lamp exposures.

We used CHIRON to measure the spectroscopic orbit and characterize the host
star of KELT-26 b. A total of 15 CHIRON epochs covering all orbital phases of
KELT-26 b were obtained. Relative velocities were measured from each spectrum
by deriving their stellar line broadening kernels via a Least Squares
Deconvolution (LSD) analysis. These velocities are listed in Table 4, and
Figure 6 shows the radial velocities as a function of time and phase-folded by
the photometric ephemeris.

The absolute radial velocity was estimated by comparing the systemic velocity
of KELT-26 in the native CHIRON system to that of four radial velocity
standard stars observed by CHIRON in order to determine a mean offset of
$-1.455\pm 0.037{\rm km~{}s^{-1}}$ of KELT-26 systemic velocity relative to
the CHIRON system. We added the uncertainty in the systemic velocity in the
CHIRON system of $0.025~{}{\rm km~{}s^{-1}}$ in quadrature to arrive at a
final absolute systemic velocity of $-1.455\pm 0.045{\rm km~{}s^{-1}}$. We
corrected the individual CHIRON relative velocities and uncertainties in the
same way.

We also used CHIRON to observe the spectroscopic transit of KELT-25b on UT
2019 May 21. A total of 18 spectra were obtained covering the transit, with an
integration time of 600 s per exposure. Spectral line profiles were derived
from each spectrum as per the procedure described in §3.3. The spectroscopic
shadow (i.e., Doppler Tomography signal) of the transiting companion is shown
in Figure 7.

###  3.3 Spectroscopic transit of KELT-25 with the Planet Finding Spectrograph

We monitored a transit of KELT-25 b with the Planet Finding Spectrograph (PFS;
Crane et al. 2010) on the 6.5 m Magellan-Clay telescope at Las Campanas
Observatory, Chile. A total of 19 spectra were obtained on UTC 2019 April 21
spanning 3.3 hours, each observation with an integration time of 600 s. For
our observations, PFS was fed by a $0\farcs 3$ slit, yielding a spectral
resolving power of $R=130,000$ over the wavelength region of $3910-7340$ Å. To
enable better derivation of the stellar line profiles, the iodine cell was
omitted from our observations. Wavelength calibrations were provided by Th-Ar
hollow cathode lamp observations obtained at the beginning and end of the
night.

Stellar line profiles were derived from each spectrum as per Collier Cameron
et al. (2010b) and Donati et al. (1997), via a LSD analysis of the spectra
against a non-rotating synthetic template spectrum generated via the ATLAS9
model atmospheres (Castelli & Kurucz 2003). During the transit, the planet
sequentially blocks parts of the rotating stellar disk, casting a shadow in
our observed rotationally broadened line profiles of the star. When we
subtract an averaged line profile from each observation, the residuals reveal
the spectroscopic transit of the planet as a shadow traversing across the
stellar surface. The line profile residuals and best fitting models are shown
in Figure 7.

###  3.4 Spectroscopic transit of KELT-26 with MINERVA-Australis

To measure the orbital obliquity and confirm the planetary status of KELT-26
b, we obtained a series of spectroscopic observations during its transit on UT
2019 June 18 via the MINiature Exoplanet Radial Velocity Array (Minerva-
Australis) facility, located at the University of Southern Queensland’s Mount
Kent Observatory, Australia (Addison et al. 2019). At the time of
observations, Minerva-Australis was operating three active, 0.7m telescopes
feeding into a single Kiwispec echelle spectrograph, yielding a resolving
power of $R=80,000$ over the wavelength region of $5,000-6,300$ Å. We made use
of data from the two of the three telescopes that yielded the highest signal-
to-noise ratio spectra on the night of the transit. Spectral line profiles
were derived from each spectrum via the LSD analysis described in Section 3.3.
From these observations, we detect the spectroscopic shadow of the transit,
finding that the path of the planet is offset from, but parallel to, the
projected stellar equator. The spectroscopic transit of KELT-26 b is shown in
Figure 7.

KELT-25b:CHIRON | KELT-25b:PFS | KELT-26b:MINERVA
---|---|---
|  |
Figure 7: The spectroscopic transit signals of KELT-25 from CHIRON (Left) and
PFS (Middle), and KELT-26 from Minerva-Australis (Right). The spectroscopic
transit signal shown for KELT-26 from Minerva-Australis is the combined signal
from two telescopes (See §3.4). Top: The shadow cast by the planet appears as
a dark trail on the line profile residuals. The horizontal axis plots the
velocity space of the line profile, while the vertical axis plots the phase,
with positive phase increasing upward. The limits of stellar rotation marked
by the vertical lines, whereas the beginning and ending of each transit are
show with the horizontal lines. Middle The best-fit model of the spectroscopic
transit is shown. Bottom: The residuals after removal of the best fit model,
showing a general lack of correlated noise in the line profile subtractions,
leading to higher confidence in the detection and a lack of any observable
stellar surface oscillations.

##  4 High-resolution imaging of KELT-25 b and KELT-26 b

###  4.1 Gemini-South Zorro Speckle

We obtained high resolution, speckle images of KELT-25 and KELT-26 to search
for nearby companions that could contaminate or dilute the light curves and
thus affect the interpretation of the planetary radii, and to rule out sources
of false positives like background eclipsing binaries.

Both stars were observed at the Gemini South Observatory using Zorro, a
speckle interferometer residing at Gemini-South. Zorro observes in two band-
passes simultaneously and is optimized for speckle observations. The
observations occurred during instrument commissioning and were the first
speckle interferometric science observations made by Zorro. Zorro is a dual-
channel imager using two Electron Multiplying CCDs (EMCCDs) as the detectors
and containing filter wheels providing bandpass limited
observations666https://www.gemini.edu/sciops/instruments/alopeke/ (see Scott &
Howell 2018). The Zorro data was reduced in the standard way as described in
Howell et al. (2011) and resulted in spatial reconstructions in each band-pass
for KELT-25 and KELT-26 providing high contrast, high resolution imaging
results.

KELT-25 was observed on UT 2019 May 18 in the blue (562/54 nm) and the red
(832/40 nm) bandpasses simultaneously. KELT-25 and KELT-26 observations
consisted of 3 sets of 1000 frames with exposure times of 0.06 seconds each,
coadded together during the data reduction process. Figure 8 shows the speckle
reconstructed image for KELT-25. The 562 nm observations (Figure 8, Top) show
no companion star from 17 mas out to $1{\arcsec}$ within 5 magnitudes of the
source, and the 832 nm observations (Figure 8, Bottom) confirm this as well in
the red, starting from 28 mas and to a delta magnitude of 6-6.5.

KELT-26 was observed on UT 2019 May 21 in the blue and red bandpasses
simultaneously. Figure 9 shows the constraints on possible stellar companions
to KELT-26. No stellar companions are detected with angular separation from
the primary from 17 mas (562 nm) and 28 mas (832 nm) out to $1.7{\arcsec}$ and
for contrast limits of $\sim$4.2 mag (562 nm) and 5-7 mag (832 nm). The black
solid line on the contrast curve marks the 5$\sigma$ detection limit.

Figure 8: Contrast curves and 562 nm image (inset, Top) and 832 nm image
(inset, Bottom) for KELT-25 from Gemini-South with the Zorro instrument (Scott
et al. 2019, in prep).

Figure 9: Contrast curves and 562 nm image (inset, Top) and 832 nm image
(inset, Bottom) for KELT-26 from Gemini-South with the Zorro instrument (Scott
et al. 2019, in prep).

###  4.2 Southern Astrophysical Research Speckle (SOAR)

KELT-26 was observed with the speckle camera at the 4.1 m Southern
Astrophysical Research (SOAR) telescope on UT 2019 August 12. The instrument
and data processing are covered in Tokovinin (2018); a full description of the
SOAR-TESS survey can be found in Ziegler et al. (2019). Figure 10 shows the
5$\sigma$ limit of companion detection, and the inset shows the speckle auto-
correlation function (ACF). The speckle ACF is symmetric, but the true
quadrant was determined from the shift-and-add image.

A nearby, faint companion was detected at an angular separation of
2.4096$\arcsec$ at position angle $311.15\arcdeg$ and a magnitude difference
of 7.1 mag in $I_{c}$ (central wavelength 824 nm, bandwidth 170 nm). This
nearby source was not seen in the Gemini-South Zorro observation.

From the contrast magnitude of the neighbor in $I_{c}$, we can estimate its
contribution to the flux in the diluted transit light curves. In this case,
the contribution of the neighbor is proportional to the ratio of the flux of
the secondary to the primary, or 0.14% in $I_{c}$. This estimate agrees with
Matson et al. (2018), who found that stellar companions with magnitude
contrasts of $\Delta m\lesssim 7.8$ can only cause $<0.1\%$ of a contribution
to the flux. This dilution factor reflects the amount by which the true
transit depth is diluted and therefore constrains the true radius of the
planet. If the true depth is larger by 0.14%, that would represent a $\sim$1
sigma difference between the true and reported depths (see Table 6).
Therefore, we should, in principle, remove the blended flux from the transit
photometry and spectral energy distribution (SED). However, we do not know
whether or not the companion is bound to the primary, and we do not have a
flux in another band. Thus it is impossible to estimate the contribution to
the flux due to the companion in any band other than $I_{c}$.

The TESS bandpass is centered on $I_{c}$, but it is $\sim$3 times wider in
wavelength space (Ricker et al. 2015), and essentially includes the $R$,
$I_{c}$ and $z$ bandpasses. If the primary and the companion had the same SED
in the TESS bandpass, then the fractional flux contamination would be the
same, i.e., $\sim 0.14\%$. This is almost certainly not the case, as the
companion is most likely to be either a foreground M star or a background
giant. In either scenario, the companion would be fairly red, and thus would
produce less fractional flux relative to the primary, which would have a much
flatter SED in the TESS bandpass. Thus, we take 0.14% as a conservative upper
limit to the blending in the TESS bandpass.

Given the blending from the companion, the true depth is larger and therefore
the true radius of the planet is larger (Ciardi et al. 2015). The inferred
radius of KELT-26 b from the diluted TESS and KELT-FUN light curves (all of
which were blended by the neighbor in their apertures) is probably larger by
at most $\sim 0.5\times 0.14\%$ or, equivalently, 0.001$R_{J}$. This is about
0.02$\sigma$ different from our reported value of the planetary radius, and
thus it does not change our qualitative results. This is simply because the
uncertainty on the radius of the planet is dominated by the uncertainty on the
radius of the star, not the depth of the transit.

Finally, we note that it is possible that this faint companion is an artifact,
based on the fact that there are no sources at that position and magnitude
found in either Gaia DR2 or the TESS Input Catalog Version 8 (TIC-8; Stassun
et al. 2019). Moreover, if it is a real source, at an angular separation of
$\sim$2.5$\arcsec$, it is unlikely to be a bound companion (see, e.g., Matson
et al. 2018).

Nevertheless, we encourage more AO/speckle observations in other photometric
filters in order to determine if the neighbor is real and bound, and to
determine the spectral type of the neighbor and therefore its total flux
contribution. This will enable a more accurate measurements of the transit
depth and planetary radius.

Figure 10: $I_{c}$ band auto-correlation function of the speckle image for
KELT-26 from SOAR. The black points represent the 5$\sigma$ sensitivity limits
for KELT-26. The inset shows the auto-correlation function. A white arrow
points to the location of the nearby companion. It is located 2.4096$\arcsec$
away from the target at PA = $311.2\arcdeg$ and has a magnitude contrast of
7.1 mag in the $I_{c}$ band. The companion is mirrored in the ACF on the
opposite side due to the speckle processing.

##  5 Host Star Properties

Figure 11: The SED fit for KELT-25 (Top) and KELT-26 (Bottom) from our
EXOFASTv2 fit. The observed values are shown in red with 1$\sigma$
uncertainties while the predicted integrated fluxes are in blue. The final
model is shown by the black line.

###  5.1 Location and Three-Dimensional Motion in the Galaxy, and Galactic
Population

We determine the location in the Galaxy, 3-dimensional (UVW) space motion
relative to the local standard of rest, and the inferred population of KELT-25
and KELT-26, using their proper motions, parallaxes, and absolute systemic
velocities. Because we have poor metallicity constraints (albeit for different
reasons; see the discussion below), we are unable to use this stellar property
to help determine the stellar populations of the hosts. However, given that
they are both early A stars, we would be surprised if they have significantly
subsolar metallicities.

Using the proper motion, Gaia parallax (corrected for the Stassun & Torres
2018b systematic offset of $-82~{}\mu{\rm as}$), and the absolute systemic
radial velocity determined as described in §3.1, we compute
$(U,V,W)=(-13.40\pm 0.58,-13.83\pm 0.84,-1.24\pm 0.19)~{}{\rm km~{}s^{-1}}$,
correcting for the velocity of the Sun with respect to the local standard of
rest as determined by Coşkunoǧlu et al. (2011). These velocities imply that
the probability of KELT-25 being in the thin disk relative to thick disk is
99.4% using the classification scheme of Bensby et al. (2003). The distance to
KELT-25 is $427.0\pm 7.8~{}{\rm pc}$ and it has Galactic coordinates of
$(\ell,b)=(237.5\arcdeg,-6.8\arcdeg)$. This implies that its vertical ($Z$)
distance from the sun is $Z-Z_{\odot}=-50.6~{}{\rm pc}$. Given that the sun is
located at $Z_{\odot}\simeq 30~{}{\rm pc}$ above the plane as determined by
the local evolved stellar population according to Bovy (2017), this means that
KELT-25 is located only about $Z\simeq 20~{}{\rm pc}$ below the plane. This is
consistent with the scale height of early A stars as determined by Bovy
(2017).

Using the same methodology, we computed $(U,V,W)=(-24.63\pm 0.37,6.75\pm
0.58,0.78\pm 0.20)~{}{\rm km~{}s^{-1}}$ for KELT-26. These velocities imply
that the probability of KELT-26 being in the thin disk relative to thick disk
is 99.3% using the classification scheme of Bensby et al. (2003). The distance
to KELT-26 is $417.78\pm 10.5~{}{\rm pc}$ and it has Galactic coordinates of
$(\ell,b)=(328.19\arcdeg,+13.32\arcdeg)$. This implies that its vertical ($Z$)
distance from the sun is $Z-Z_{\odot}=96.1~{}{\rm pc}$ and $Z\simeq 126~{}{\rm
pc}$ above the plane. This is roughly twice the typical scale height of an
early A star as determined by Bovy (2017), and thus is notable but not
completely unexpected.

Both KELT-25 and KELT-26 have roughly the same Galactocentric distance at the
sun. Assuming that the distance from the sun to the Galactic center is roughly
$R_{0}=8.2\simeq{\rm kpc}$ (Gravity Collaboration et al. 2019), we estimate
Galactocentric distances of $8.4~{}{\rm kpc}$ and $7.8~{}{\rm kpc}$ for
KELT-25 and KELT-26, respectively.

Table 5: Median values and 68% confidence interval for the stellar parameters
of KELT-25 and KELT-26 derived from the global fit.

Parameter | Units | Values | Values |  |
---|---|---|---|---|---
|  | KELT-25 | KELT-26 |  |
$M_{*}$ |  Mass ($\,M_{\sun}$) | $2.18^{+0.12}_{-0.11}$ | $1.93^{+0.14}_{-0.16}$ |  |
$R_{*}$ |  Radius ($\,R_{\sun}$) | $2.264^{+0.048}_{-0.052}$ | $1.801^{+0.049}_{-0.048}$ |  |
$L_{*}$ |  Luminosity ($\,L_{\sun}$) | $21.8^{+4.6}_{-1.8}$ | $16.4^{+3.8}_{-1.8}$ |  |
$\rho_{*}$ | Density (cgs) | $0.263^{+0.025}_{-0.018}$ | $0.463^{+0.040}_{-0.038}$ |  |
$\log{g}$ | Surface gravity (cgs) | $4.064^{+0.032}_{-0.026}$ | $4.211^{+0.030}_{-0.033}$ |  |
$T_{\rm eff}$ | Effective Temperature (K) | $8280^{+440}_{-180}$ | $8640^{+500}_{-240}$ |  |
$[{\rm Fe/H}]$ | Metallicity (dex) | $0.30^{+0.13}_{-0.21}$ | $-0.06^{+0.30}_{-0.34}$ |  |
$[{\rm Fe/H}]_{0}$ | Initial Metallicity | $0.34^{+0.11}_{-0.19}$ | $0.03^{+0.25}_{-0.32}$ |  |
$Age$ | Age (Gyr) | $0.46^{+0.14}_{-0.12}$ | $0.43^{+0.31}_{-0.25}$ |  |
$EEP$ | Equal Evolutionary Phase | $342.0^{+6.3}_{-7.1}$ | $331^{+17}_{-28}$ |  |
$vsinI_{*}$ | Projected rotational velocity (km/s) | $114.200\pm 1.200$ | $12.280^{+7.80}_{-8.20}$ |  |
$V_{\rm line}$ | Unbroadened line width (m/s) | $3700^{+2300}_{-2400}$ | $6400^{+1800}_{-1900}$ |  |
$A_{V}$ | V-band extinction (mag) | $0.104^{+0.16}_{-0.073}$ | $0.103^{+0.14}_{-0.078}$ |  |
$\sigma_{SED}$ | SED photometry error scaling | $2.50^{+0.78}_{-0.53}$ | $1.76^{+0.81}_{-0.48}$ |  |
$\pi$ | Parallax (mas) | $2.366^{+0.041}_{-0.042}$ | $2.396\pm 0.063$ |  |
$d$ | Distance (pc) | $422.7^{+7.6}_{-7.3}$ | $417\pm 11$ |  |

NOTES: †The initial metallicity is the metallicity of the star when it was
formed.
‡The Equal Evolutionary Phase corresponds to static points in a star’s
evolutionary history when using the MIST isochrones and can be a proxy for
age. See §2 in Dotter (2016) for a more detailed description of EEP.

Table 6: Median values and 68% confidence interval for the physical parameters
of KELT-25 b and KELT-26 b from the global fit.

Parameter | Description (Units) | Values | Values |  |
---|---|---|---|---|---
|  | KELT-25 | KELT-26 |  |
$P$ | Period (days) | $4.401131\pm 0.000059$ | $3.3448412\pm 0.0000033$ |  |
$R_{P}$ |  Radius ($\,R_{\rm J}$) | $1.642^{+0.039}_{-0.043}$ | $1.940^{+0.060}_{-0.058}$ |  |
$M_{P}$ |  Mass ($\,M_{\rm J}$) | $\color[rgb]{1,0,0}(<64)$ | $1.41^{+0.43}_{-0.51}$ |  |
$T_{C}$ |  Time of conjunction ($\rm{BJD_{TDB}}$) | $2458493.31501\pm 0.00037$ | $2457482.31209^{+0.00090}_{-0.00091}$ |  |
$T_{0}^{\dagger}$ |  Optimal conjunction Time ($\rm{BJD_{TDB}}$) | $2458506.51840\pm 0.00034$ | $2458321.86724^{+0.00038}_{-0.00039}$ |  |
$a$ | Semi-major axis (AU) | $0.0681^{+0.0012}_{-0.0011}$ | $0.0545^{+0.0013}_{-0.0015}$ |  |
$i$ | Inclination (Degrees) | $85.37^{+0.55}_{-0.42}$ | $84.45^{+0.39}_{-0.41}$ |  |
$T_{eq}$ | Equilibrium temperature (K) | $2306^{+100}_{-47}$ | $2402^{+130}_{-71}$ |  |
$\tau_{\rm circ}$ | Tidal circularization timescale (Gyr) | $-0.1^{+4.1}_{-3.9}$ | $0.053^{+0.019}_{-0.020}$ |  |
$K$ | RV semi-amplitude (m/s) | $\color[rgb]{1,0,0}(<4687.8)$ | $123^{+37}_{-45}$ |  |
$\log{K}$ | Log of RV semi-amplitude | $\color[rgb]{1,0,0}(<3.64)$ | $2.09^{+0.11}_{-0.20}$ |  |
$R_{P}/R_{*}$ | Radius of planet in stellar radii | $0.07450^{+0.00039}_{-0.00042}$ | $0.11066^{+0.00090}_{-0.00087}$ |  |
$a/R_{*}$ | Semi-major axis in stellar radii | $6.46^{+0.20}_{-0.15}$ | $6.49\pm 0.18$ |  |
$\delta$ | Transit depth (fraction) | $0.005550^{+0.000059}_{-0.000062}$ | $0.01225^{+0.00020}_{-0.00019}$ |  |
$\tau$ | Ingress/egress transit duration (days) | $0.0192^{+0.0011}_{-0.0012}$ | $0.0238^{+0.0016}_{-0.0014}$ |  |
$T_{14}$ | Total transit duration (days) | $0.2051^{+0.0013}_{-0.0015}$ | $0.1514\pm 0.0016$ |  |
$T_{FWHM}$ | FWHM transit duration (days) | $0.18586^{+0.00097}_{-0.00096}$ | $0.12760^{+0.00098}_{-0.00099}$ |  |
$b$ | Transit Impact parameter | $0.522^{+0.034}_{-0.048}$ | $0.628^{+0.027}_{-0.029}$ |  |
$\delta_{S,3.6\mu m}$ |  Blackbody eclipse depth at 3.6$\mu$m (ppm) | $732^{+27}_{-23}$ | $1664^{+82}_{-72}$ |  |
$\delta_{S,4.5\mu m}$ |  Blackbody eclipse depth at 4.5$\mu$m (ppm) | $872^{+26}_{-24}$ | $1965^{+83}_{-74}$ |  |
$\rho_{P}$ | Density (cgs) | $\color[rgb]{1,0,0}(<18.3)$ | $0.238^{+0.077}_{-0.088}$ |  |
$logg_{P}$ | Surface gravity | $\color[rgb]{1,0,0}(<4.77)$ | $2.97^{+0.12}_{-0.20}$ |  |
$\lambda$ | Projected Spin-orbit alignment (Degrees) | $23.4^{+3.2}_{-2.3}$ | $91.3^{+6.5}_{-6.3}$ |  |
$\Theta$ | Safronov Number | $-0.01^{+0.50}_{-0.48}$ | $0.041^{+0.012}_{-0.015}$ |  |
$\langle F\rangle$ |  Incident Flux (109 erg s-1 cm-2) | $6.42^{+1.2}_{-0.51}$ | $7.56^{+1.7}_{-0.86}$ |  |
$T_{P}$ |  Time of Periastron ($\rm{BJD_{TDB}}$) | $2458493.31501\pm 0.00037$ | $2457482.31209^{+0.00090}_{-0.00091}$ |  |
$T_{S}$ |  Time of eclipse ($\rm{BJD_{TDB}}$) | $2458495.51558\pm 0.00036$ | $2457483.98451\pm 0.00090$ |  |
$T_{A}$ |  Time of Ascending Node ($\rm{BJD_{TDB}}$) | $2458496.61586\pm 0.00035$ | $2457484.82072\pm 0.00090$ |  |
$T_{D}$ |  Time of Descending Node ($\rm{BJD_{TDB}}$) | $2458494.41530\pm 0.00036$ | $2457483.14830\pm 0.00090$ |  |
$M_{P}\sin i$ |  Minimum mass ($\,M_{\rm J}$) | $\color[rgb]{1,0,0}(<64)$ | $1.40^{+0.43}_{-0.51}$ |  |
$M_{P}/M_{*}$ | Mass ratio | $\color[rgb]{1,0,0}(<0.028)$ | $0.00070^{+0.00021}_{-0.00026}$ |  |
$d/R_{*}$ | Separation at mid transit | $6.46^{+0.20}_{-0.15}$ | $6.49\pm 0.18$ |  |
$P_{T}$ | A priori non-grazing transit prob | $0.1433^{+0.0035}_{-0.0043}$ | $0.1370^{+0.0039}_{-0.0036}$ |  |
$P_{T,G}$ | A priori transit prob | $0.1664^{+0.0041}_{-0.0051}$ | $0.1710^{+0.0050}_{-0.0047}$ |  |
Telescope Parameters: | TRES | CHIRON |  |
$\gamma_{\rm rel}$ | Relative RV Offset (m/s) | $720^{+810}_{-820}$ | $-25595^{+25}_{-24}$ |  |
$\sigma_{J}$ | RV Jitter (m/s) | $2240^{+1100}_{-620}$ | $67^{+32}_{-23}$ |  |

NOTES: †Minimum covariance with period. All values in this table for the
secondary occultation of KELT-25 b are predicted values from our global
analysis. All values in red are 3$\sigma$ upper limits on mass dependent
parameters for KELT-25 b.

Table 7: Median values and 68% confidence interval for global model of KELT-25 and KELT-26 from the global fit.  KELT-25 b |  |  |  |
---|---|---|---|---
Wavelength Parameters: | R | TESS |
$u_{1}$ | linear limb-darkening coeff | $0.292\pm 0.049$ | $0.159\pm 0.029$ |
$u_{2}$ | quadratic limb-darkening coeff | $0.348\pm 0.049$ | $0.260\pm 0.035$ |
$A_{T}$ | Secondary eclipse depth (ppm) | – | $187\pm 46$ |
Transit Parameters: | PEST UT 2019-01-18 (R) | TESS Sector 7 | TESS Sector 7 (secondary)
$\sigma^{2}$ | Added Variance | $0.00000779^{+0.00000062}_{-0.00000058}$ | $-0.0000000863^{+0.0000000057}_{-0.0000000050}$ | $-0.0000000885^{+0.0000000066}_{-0.0000000057}$
$F_{0}$ | Baseline flux | $1.00362\pm 0.00013$ | $0.999826\pm 0.000049$ | $0.999811\pm 0.000040$
Doppler Tomography Parameters: |  |  |
$\sigma_{DT}$ | Doppler Tomography Error scaling | $0.9975^{+0.0100}_{-0.0098}$ | $0.9839^{+0.0057}_{-0.0056}$ |
KELT-26 b |  |  |  |
Wavelength Parameters: | I | r’ | TESS
$u_{1}$ | linear limb-darkening coeff | $0.185\pm 0.048$ | $0.207\pm 0.045$ | $0.212\pm 0.044$
$u_{2}$ | quadratic limb-darkening coeff | $0.264\pm 0.051$ | $0.279\pm 0.048$ | $0.280\pm 0.047$
Transit Parameters: | PEST UT 2016-07-26 (I) | CDK700 UT 2018-03-20 (r’) | TESS UT 2019-04-07 (TESS)
$\sigma^{2}$ | Added Variance | $0.00000780^{+0.0000010}_{-0.00000089}$ | $0.00000170^{+0.00000025}_{-0.00000021}$ | $0.000000420^{+0.000000029}_{-0.000000027}$
$F_{0}$ | Baseline flux | $0.99573\pm 0.00021$ | $0.99891\pm 0.00015$ | $1.000000\pm 0.000026$
Doppler Tomography Parameters: | MINERVA 3 | MINERVA 4 |
$\sigma_{DT}$ | Doppler Tomography Error scaling | $0.991\pm 0.013$ | $0.997\pm 0.012$ |

##  6 EXOFASTv2 Global Fits for KELT-25 and KELT-26

To constrain the system parameters, we modeled the available radial
velocities, transit photometry, and multiband absolute photometry for KELT-25
and KELT-26 using the exoplanet fitting suite, EXOFASTv2 (Eastman et al. 2013,
2019; Eastman 2017).

For each system, we globally fit the radial velocities (see §3), the TESS and
follow-up photometry (see §2.2 and §2.3) and the Doppler tomographic shadows
simultaneously (see §3.2, §3.3, and §3.4). Within these fits, we determined
the host star properties using a combination of spectroscopic priors, the
spectral energy distribution, and the MESA Isochrones and Stellar Tracks
(MIST) stellar evolution models (Dotter 2016; Choi et al. 2016; Paxton et al.
2011, 2013, 2015). For each system, we placed Gaussian priors of [Fe/H]
$=0.0\pm 0.5$ dex for the metallicity of the host stars, as we did not have
reliable constraints from our available spectra. From an independent EXOFASTv2
analysis of the KELT photometry, we adopted a Gaussian prior on the period of
KELT-25 b of 4.401093 $\pm$ 0.000073 days and on the epoch of
$T_{c}=2458493.3144$ $\rm{BJD_{TDB}}$. Similarly, we placed a Gaussian prior
on the orbital period ($P=3.344886\pm 0.000064$ days) and the epoch
($T_{c}=2457482.315\pm 0.008600$ $\rm{BJD_{TDB}}$) of KELT-26 b. We also
placed priors on the $v\sin{I_{*}}$ from spectroscopy for KELT-25
($v\sin{I_{*}}=111.277\pm 1.422$ km s-1) and KELT-26 ($v\sin{I_{*}}=9.9349\pm
1.1333$ km s-1) and used parallaxes from Gaia DR2 (see Table 1). In addition,
we assumed a circular orbit for both systems and we fit for the depth of the
secondary eclipse of KELT-25 observed in the lower right panel of Figure 2. We
constrained $A_{V}$ to the maximum permitted line-of-sight extinction from
Schlegel et al. (1998), and the temperature and stellar mass and radius from
the SED best-fit values (see Figure 11 for the empirical SEDs of both
systems). The best-fitting evolutionary models are shown in Figure 12, and the
final fit parameters of the EXOFASTv2 analyses for both stars and their
companions are shown in Tables 5 and 6.

Although a global fit to the photometry, radial velocities (including the
Doppler Tomography measurements), and the SED can completely constrain the
properties of the system, we include constraints from the MIST stellar
evolutionary models, which include reasonably well-understood stellar physics,
in our global fit. In this manner, we derive best-fit distances to both stars
that are slightly different than those inferred from the Gaia parallaxes
alone. As a result, we find a posterior distribution of the distance of
KELT-25 after the global fit of $422.5^{+7.6}_{-7.3}$ pc, which is $\sim 0.6$
sigma from that inferred from the Gaia parallax (see Table 1). Similarly, for
KELT-26, we derive a posterior distance after the global fit of $417\pm 11$
pc, which is $\sim$0.06 sigma from the distance inferred from the Gaia DR2
parallax alone. We consider these differences to be completely consistent
within the uncertainties. To calculate the location of KELT-25 and KELT-26
within the Galaxy and their UVW space velocities in §5.1, we opted to use the
model-independent parallaxes and distances from Gaia DR2. Although in this
case these distance measurements are completely consistent, the discovery of
larger discrepancy in fits of other systems may have implications for the
current stellar evolutionary models or Gaia data. Thus the empirical distances
from Gaia DR2 can serve as a way to test and calibrate the models we use to
derive fundamental stellar parameters, or uncover evidence of systematic
errors in the Gaia data itself.

Figure 12: The best-fitting MIST track for KELT-25 (Top) and KELT-26 (Bottom)
shown by the blue line. The black line shows the 3$\sigma$ contours for the
MIST evolutionary tracks. The red data point shows the median values and
1$\sigma$ uncertainties from our global fit, while the green contours
correspond to the 3$\sigma$ errors. The blue points mark the location of 1.0
(Top) and 0.5 Gyr (Bottom) along the MIST track.

##  7 Discussion

Both KELT-25 b and KELT-26 b represent extreme transiting systems in a few key
aspects. First, KELT-25 b and KELT-26 b both orbit relatively bright
(V$\sim$10 mag) and extremely hot hosts ($T_{\rm eff}\simeq 8300$ K and
$\simeq$ 8700 K, respectively), and they also have short orbital periods
($P\sim 4.4$ and $P\sim 3.3$ days, respectively). Their proximity to their
host stars and their stars’ intrinsic luminosity mean that they receive
extreme amounts of stellar radiation, paticularly high-energy radiation,
resulting in high equilibrium temperatures – assuming zero albedo and complete
redistribution – of $T_{eq}=2306$ K (KELT-25 b) and $T_{eq}=2402$ K (KELT-26
b). Both KELT-25 b and KELT-26 b join the recently-defined class of planets of
“Ultra Hot Jupiters”, which, similar to the prototype WASP-33b (Collier
Cameron et al. 2010a), are primarily planets on short-period orbits around
early A stars. They are thus among the hottest transiting exoplanets known.
Indeed, the equilibrium temperatures (zero albedo, complete heat distribution)
of these planets are commonly in excess of 2000 K, and it seems likely that
their day-side temperatures would be markedly higher still (see Figure 13),
and likely have day-side temperatures that are much higher.

Our EXOFASTv2 models (Table 6) indicate that the radii of both planets are
also significantly inflated ($R_{p}=1.64R_{J}$ and $R_{p}=1.94R_{J}$). From an
irradiation evolution analysis of these systems (see §7.3), we conclude that
these objects currently receive an insolation flux of around $5\times 10^{9}$
$\rm erg$ $\rm s^{-1}\rm cm^{-2}$ (KELT-25 b) and $1\times 10^{10}$ $\rm erg$
$\rm s^{-1}\rm cm^{-2}$ (KELT-26 b). Moreover, their orbital histories suggest
that they probably have always been above the Demory & Seager (2011)
insolation threshold of $2\times 10^{8}$ $\rm erg$ $\rm s^{-1}\rm cm^{-2}$,
which is an empirical threshold above which giant planets exhibit significant
radius inflation. For this reason, it is not surprising that they are both
highly inflated. The extreme temperatures of these companions and the optical
and infrared brightness of their hosts (KELT-25: $V=9.65$, $J=9.36$; KELT-26:
$V=9.96$, $J=9.77$) mean that the prospects for detailed atmospheric
characterization via transmission spectroscopy with James Webb Space Telescope
(JWST) or the Hubble Space Telescope (HST) are promising. Indeed, the TESS
light curve for KELT-25 (Figure 2) demonstrates the weak but significant
detection of the secondary eclipse of the system, as the planet moves behind
the star. Given the estimated equilibrium temperature of the star, this flux
decrement is also certainly caused by the tail of the thermal emission from
the planet. We estimate a secondary eclipse depth of $\sim 187\pm 46$ ppm (see
Table 7), implying a brightness temperature of the planet in the TESS band of
$\sim$$3396^{+140}_{-170}$K, which is substantially higher than the
equilibrium temperature (assuming zero albedo and complete heat
redistribution) of $2303^{+100}_{-47}$K.

To better place these planets in context of all systems with measured rotation
rates and projected spin-orbit misalignments, we highlight KELT-25 and KELT-26
in Figures 17 and 8 showing the distributions of $v\sin{I_{*}}$ and spin-orbit
misalignments $\lambda$ versus stellar temperature for A stars ($v\sin{I_{*}}$
distribution) and all known planet hosts with measured spin-orbit angles.
Figure 13 shows a plot of atmospheric scale height as a function of
equilibrium temperature for all the known transiting exoplanets with measured
masses. KELT-26 b has a large scale height and receives extreme amounts of UV
radiation from its host. With KELT-25 b and KELT-26 b, we have a large enough
sample of A-stars with transiting exoplanets that we can begin to see emerging
patterns in the population (see Table 8). Perhaps one such interesting trend
is the gap in planet equilibrium temperatures between roughly 2600 K and 4000
K, visible in Figure 13. We still do not understand whether this gap is real
or the result of selection effects. All the confirmed planets around A stars
have short periods (1.22 ¡ $P$ ¡ 4.79 days) and transit bright host stars
($V\lesssim 10$), which, as previously remarked, make them attractive targets
for atmospheric characterization with the upcoming JWST. Before JWST launches,
however, TESS may be able to observe these planets in transit, in some cases
even resolving their secondary eclipses, as with KELT-25 b. These observations
can constrain the brightness temperature of these planets and therefore
provide insights into the heat distribution mechanisms of their atmospheres.

Figure 13: This scatter-plot shows the atmospheric scale height versus
equilibrium temperature for all the known transiting exoplanets with mass
measurements. The color scale corresponds to the extreme ultraviolet radiation
that planets receive from their host stars. The symbol sizes are inversely
proportional to the magnitude in V-band of the host stars. At $T_{\rm
eq}\sim$2402 K, KELT-26 b stands out as one of the hottest known exoplanets,
receiving extreme amounts of UV radiation, likely resulting in the
exceptionally large radius of $R_{p}\sim 1.9\,R_{\rm J}$, and subsequent large
scale height. Table 8: A-star transiting planet hosts from the literature
ordered by decreasing stellar temperature

Planet | $M_{p}$ | $R_{p}$ | $T_{\rm{eff}}$ | $P$ | $L_{\star}$ | $T_{eq}$ | $\lambda$ | SpT | Ref.
---|---|---|---|---|---|---|---|---|---
|  ($M_{J}$) |  ($R_{J}$) | (K) | (days) |  ($L_{\odot}$) | (K) | (deg) |  |
KELT-9b |  2.88$\pm$0.84 | $1.891^{+0.061}_{-0.053}$ |  10170$\pm$450 | 1.48 | $53^{+13}_{-10}$ |  4050$\pm$180 |  -84.8$\pm$1.4 | A0 | 1
KELT-20b/MASCARA-2b |  ¡ 3.5 (3$\sigma$) | $1.735^{+0.07}_{-0.075}$ | $8730^{+250}_{-260}$ | 3.47 | $12.7^{+2.2}_{-1.9}$ |  $\sim$2250 | $3.4\pm 2.1$ | A2 | 2
KELT-26 b | $1.41^{+0.43}_{-0.51}$ | $1.940^{+0.060}_{-0.058}$ | $8640^{+500}_{-240}$ | 3.34 | $16.4^{+3.8}_{-1.8}$ | $2402^{+130}_{-71}$ | $91.3^{+6.5}_{-6.3}$ | A3m | 3
HAT-P-70b | $<6.14$ | $2.011_{-0.114}^{+0.051}$ | $8450_{-690}^{+540}$ | $2.74$ | $16.7_{-4.6}^{+5.3}$ | $2562_{-52}^{+43}$ | $116.5_{-3.8}^{+3.5}$ | A3V | 4
KELT-25 b |  $<$64 | $1.642^{+0.039}_{-0.043}$ | $8280^{+440}_{-180}$ | 4.40 | $21.8^{+4.6}_{-1.8}$ | $2306^{+100}_{-47}$ | $23.4^{+3.2}_{-2.3}$ | A4 | 5
WASP-189b | $2.13\pm 0.28$ | $1.374\pm 0.082$ | $8000\pm 100$ | 2.72 | $1.293\pm 0.045$ | $2641\pm 34$ | $89.3\pm 1.4$ | A6IV-V | 6
HATS-70b | $12.9_{-1.6}^{+1.8}$ | $1.384_{-0.074}^{+0.079}$ | $7930_{-820}^{+630}$ | $1.89$ | $12.0_{-3.4}^{+5.5}$ | $2730_{-160}^{+140}$ | $8.9_{-4.5}^{+5.6}$ | A6V | 7
MASCARA-4b | $3.1\pm 0.9$ | $1.53_{-0.04}^{+0.07}$ | $7800\pm 200$ | $2.82$ | $12.23\pm 0.655$ | $2100\pm 100$ | $247.5_{-1.7}^{+1.5}$ | A7V | 8
Kepler-13Ab |  $\sim$ 9.2 $\pm$ 1.1 | $1.512\pm 0.035$ | $7650\pm 250$ | 1.76 | - | $2550\pm 80$ | $58.6\pm 2.0$ | A8V | 9
KELT-21b |  ¡ 3.91 (3$\sigma$) | $1.586^{+0.039}_{-0.040}$ | $7598^{+81}_{-84}$ | 3.61 | $8.03^{+0.54}_{-0.53}$ | $2051^{+29}_{-30}$ | $-5.6^{+1.7}_{-1.9}$ | A8V | 10
MASCARA-1b | $3.7\pm 0.9$ | $1.5\pm 0.3$ | $7554\pm 150$ | 2.14 | $13.1\pm 3$ | $2570^{+50}_{-30}$ | $69.5\pm 3$ | A8V | 11
HAT-P-57b | $1.41\pm 1.52$ | $1.74\pm 0.36$ | $7500\pm 250$ | 2.46 | $6.4\pm 1.1$ | 2200 |  $-16.7$-$3.3$ or $2.76$-$57.4$ | A8V | 12
KELT-19Ab | $1.62^{+0.25}_{-0.20}$ | $1.83\pm 0.10$ |  7500$\pm$110 | 4.61 | $9.5^{+1.2}_{-1.1}$ |  $\sim$1935 | $-179.7^{+3.7}_{-3.8}$ | Am | 13
KELT-17b | $1.31^{+0.28}_{-0.29}$ | $1.525^{+0.065}_{-0.060}$ |  7454 $\pm$ 49 | 3.08 | $7.51^{+0.62}_{-0.55}$ | $2087^{+32}_{-33}$ |  -115.9$\pm$4.1 | A9V | 14
WASP-33b | 4.1 | $1.497\pm 0.045$ |  7430$\pm$100 | 1.22 | - | $2710\pm 50$ | $251.6\pm 0.7$ | A9V | 15
HAT-P-69b | $3.54_{-0.60}^{+0.61}$ | $1.714\pm 0.028$ | $7394_{-600}^{+360}$ | $4.79$ | $10.0_{-0.9}^{+1.8}$ | $1930_{-230}^{+80}$ | $16.5_{-1.9}^{+2.1}$ | A9V | 16

NOTES:

References: 1. Gaudi et al. (2017) 2\. Lund et al. (2017); Talens et al.
(2018) 3\. This work 4. Zhou et al. (2019b) 5\. This work 6. Anderson et al.
(2018) 7\. Zhou et al. (2019b) 8\. Dorval et al. (2019) 9\. Shporer et al.
(2011); Esteves et al. (2015); Johnson et al. (2014) 10\. Johnson et al.
(2018) 11\. Talens et al. (2017) 12\. Hartman et al. (2015) 13\. Siverd et al.
(2018) 14\. Zhou et al. (2016) 15\. Collier Cameron et al. (2010b) 16\. Zhou
et al. (2019a)

###  7.1 KELT-26 b: A giant planet orbiting a likely Am star with a likely
significant transit asymmetry

KELT-26 b orbits a relatively young ($\sim$ 430 Myr), slowly rotating A star
($v\sin{I_{*}}$= 12.2 km s-1). This rotational velocity is rather atypical for
an early A star, as such stars tend to be much faster rotators on average.
From the Doppler tomographic observations (see §3), we also measured the
projected spin-orbit angle of this system, and conclude that it is on an orbit
that is consistent with being exactly parallel to the projected stellar
equator, with $\lambda$ = $91.3^{+6.5}_{-6.3}$. However, the orbit is likely
not polar, because the impact parameter of planet’s transit is $b\simeq 0.6$
and therefore the planet is not transiting exactly across the pole of the
star. The projected spin-orbit alignment $\lambda$ need not be the true true
obliquity $\psi$, which is a more fundamental quantity of the system. However,
the latter is harder to constrain, as explained in Johnson et al. (2018),
because it requires knowledge of both the planetary orbital inclination i as
well as the stellar spin axis $I_{*}$. The observation of an orbit that has
$\lambda\sim 90\arcdeg$ with the unusually slow rotation period of KELT-26
($v\sin{I_{*}}=12.2$ km s-1) may imply that the star is perhaps spinning
faster but we are observing it nearly pole-on. The phase-folded TESS light
curve of the primary (Figure 3) appears to show a slight asymmetry, such that
the planet is first passing over a region of the star with higher-than average
surface brightness, whereas the planet later passes over a region of the star
with lower surface brightness. This would be expected from gravity darkening,
assuming the planet first passes over or near the pole, and then over the
lower surface brightness equator (Barnes 2009). Gravity darkening has already
been observed with TESS for two hot Jupiters: HAT-P-69b (TOI 625.01) and
HAT-P-70b (TOI 624.01) (Zhou et al. 2019b), and is also clear in the
unpublished TESS light curve of KELT-9 (Wong et al. (2019), Wachiraphan et
al.; Ahlers et al., in preparation). However, as we discuss below, the host
star KELT-26 also appears to be variable at the $\sim$few millimagnitude level
at a period that is nearly commensurate (1:18) with the period of the planet.
This variability may also be causing the slight asymmetry.

A Lomb-Scargle periodogram (Lomb 1976; Scargle 1982) of the light curve of
KELT-26 shows a significant peak at a period of 0.185 days, with an amplitude
of 0.115%, or 1.25 mmag. This period is nearly 1/18 (1/18.06, to be precise)
of the period of the planet. We do not know if this is simply a coincidence,
as the mass of the companion is likely too small to induce periodic
oscillations on its host star. Both KELT-25 and KELT-26 are inside the
instability strip, where one might expect to find $\delta$ Scuti pulsations.
Indeed, the period and amplitude of the variability of KELT-26 are consistent
with other $\delta$ Scuti variables. A more detailed study of the nature of
the intrinsic variability of KELT-26 is beyond the scope of this paper.

On the other hand, the spectrum of KELT-26 shows evidence of being an Am star
(or ”metallic-line A star”). Am stars typically rotate much more slowly than A
stars of the same effective temperature. This is usually attributed to a
stellar companion that has spun down the star or otherwise ”stolen” its
angular momentum at birth. However, we find no evidence of a stellar companion
which would affect the spin rate of KELT-26, and the planetary companion is
not sufficiently massive to play this role. Am stars are typically identified
by the fact that the star does not appear to have a consistent metallicity
when measured using absorption lines formed at different depths in the
photosphere. This can be seen in Figure 14, where it is clear that for models
with a fixed $T_{\rm eff}$ and varying [Fe/H], no single model can
simultaneously fit all of the spectral lines.

Figure 14: A portion of the CHIRON spectrum near the Mgb region for KELT-26 is
shown in grey. The other lines show a set of $T_{\rm eff}=9000$ K,
$\log{g_{\star}}=4.25$ synthetic spectra with [Fe/H] of $-0.5$ (green), 0.0
(red), and $+0.5$ (orange) dex, generated with the ATLAS9 model atmospheres
(Castelli & Kurucz 2003), demonstrating that a single [Fe/H] cannot
simultaneously fit all of observed spectral features.

###  7.2 KELT-25 b: A substellar object transiting a rapidly rotating, young A
star and likely cluster member

With a $v\sin{I_{*}}$ of 114.2 km $\rm s^{-1}$, KELT-25 is rotating much
faster than KELT-26, and this has implications for the dynamical history of
its potential planet. In this case, the classical scenario of hot Jupiters
spiraling towards their host stars is reversed. The stellar tidal dissipation
causes the semi-major axis of KELT-25 b’s orbit to gradually increase (see
§7.3). As a result, KELT-26 b will avoid getting engulfed by its host star, at
least until the star leaves the main sequence. The Doppler tomography shadows
(Figure 7) suggest that it is on a prograde, aligned orbit ($\lambda$ =
$23.4^{+3.2}_{-2.3}$). To estimate the age of the systems, we show a modified
Hertzsprung-Russell diagram (log$g_{*}$ vs $T_{\rm eff}$) in Figure 12. From
the MIST stellar evolutionary models, and taking the 1-sigma upper limit on
the mass for KELT-25, we infer an age of $0.46^{+0.14}_{-0.12}$ Gyr. Using the
same models, we obtain an age for KELT-26 of $0.43^{+0.31}_{-0.25}$ Gyr.

####  7.2.1 Is KELT-25 a member of a stellar cluster, association, or moving
group?

We cross-matched existing TESS Objects of Interest (TOIs) to the catalog of
clusters presented in Kounkel & Covey (2019), and we found a match between
TOI-626.01 (KELT-25) and one of the putative clusters identified in that paper
as Theia 449. In that paper, they identified 1900 clusters from an analysis of
the distribution of sources in 5-dimensional space (three-dimensional position
and two-dimensional (e.g., transverse) velocity) in Gaia DR2. They then
performed a clustering analysis on Gaia sources within $|b|<30\arcdeg$ of the
Galactic plane and parallaxes with $\pi>1$ mas using a Python implementation
of HDBSCAN (Hierarchical Density-Based Spatial Clustering of Applications with
Noise, McInnes et al 2017). They estimated ages of their clusters and
associations applying a combination of machine learning and isochronal fitting
techniques to determine the ages of their sources to a precision of $\sim$0.15
dex.

For the cluster Theia 449, Kounkel & Covey (2019) report a mean Galactic
latitude of $b=-7.13\arcdeg$, a mean parallax of $2.37\pm 0.82$ mas, an age of
$0.162$ Gyr and mean radial velocity of $23.54\pm 18.21$ km $\rm s^{-1}$. From
our MIST models, we determine KELT-25 is a young A star, with an age of
$0.46^{+0.14}_{-0.12}$ Gyr, which is just $\sim$2-sigma discrepant with the
mean reported age of the cluster. Our derived properties for KELT-25 are thus
in general agreement with the average properties of Theia 449. However, the
broader issue of whether Theia 449 truly represents a collection of coeval
stars is outside the scope of this paper. Clarification of the status of Theia
449 could help resolve the age, metallicity, and formation environment of
KELT-25.

####  7.2.2 Is KELT-25 b a planet or brown dwarf?

Given its extremely fast rotation, we were only able to constrain the mass of
KELT-25 b to a $3\sigma$ upper limit of $\sim 64~{}\,M_{\rm J}$, or a
$1\sigma$ upper limit of $5.46~{}\,M_{\rm J}$. We argue that KELT-25 b is
likely to be a planet or low-mass brown dwarf based on several lines of
reasoning. First, substellar companions at the upper end of the allowed mass
range and with this period are known to be relatively rare (the so-called
”brown dwarf desert” (e.g., Grether & Lineweaver 2006). Second, no brown
dwarfs (BDs) are known that are as highly inflated as KELT-25 b (Zhou et al.
2019a). Inspection of Figure 9 of Zhou et al. (2019a) reveals that not only do
no BD have radii as large as KELT-25 b, only about a dozen lower-mass
transiting planets have radii this large. Given the large surface gravity
expected if KELT-25 b had a mass significantly above the deuterium burning
limit, this is strong circumstantial evidence that it is less massive.
Finally, there is no evidence of Doppler beaming (Loeb & Gaudi 2003) or
ellipsoidal variability (Drake 2003) in the KELT-25 b TESS light curve, which
would likely be expected if KELT-25 b had a mass substantially above the
deuterium burning limit. The weight of evidence indicates that KELT-25 b is
most likely a giant planet or very low-mass brown dwarf.

####  7.2.3 Intrinsic Variability of KELT-25

As mentioned previously, both KELT-25 and KELT-26 are inside the instability
strip. Only about 40% of stars within this range of $T_{\rm eff}$ show
$\delta$ Scuti pulsations (Murphy et al. 2019; see their Fig. 11). While
KELT-26 does show variability consistent with $\delta$ Scuti pulsations (see
§7.1), we find no evidence of intrinsic variability in the KELT-25 TESS light
curve.

###  7.3 Tidal Evolution and Irradiation History

We estimated the orbital and irradiation evolution of KELT-25 and KELT-26; in
particular, we calculated the history of the companions’ semi-major axis and
irradiation as a function of stellar age using the Planetary Orbital Evolution
due to Tides (POET; Penev et al. 2014). POET assumes a constant tidal phase
lag or quality factor $Q_{\star}$, a circular orbit, and no perturbations in
the orbits due to unseen stellar or planetary companions. We further assumed
that the tides raised by the planet or substellar companion are negligible and
that the evolution of the planet’s orbit is therefore dominated by the
dissipation of tidal perturbations in the star (as explained in Rodriguez et
al. 2019b). We accounted for the changes in stellar radius and luminosity in
time by using a MIST stellar evolutionary track corresponding to the best-fit
stellar properties (see §6). Because of the large uncertainties in the
knowledge of the tidal dissipation in stars and the tidal quality factor,
$Q_{\star}$, we consider different constant values of the dissipation
parameter $Q’_{\star}$, namely, $Q’_{\star}=10^{6.6}$, $10^{7}$ and $10^{8}$
for KELT-25 and $Q’_{\star}=10^{5}$, $10^{6}$ and $10^{7}$ for KELT-26, where
$Q’_{\star}$ is just proportional to $Q_{\star}$. With these assumptions, we
proceeded to calculate the past and future evolution of the semi-major axis
(in units of the stellar radius) as a function of the age of the system. For
KELT-26b, we can see in Figure 16 (Top) that for every assumed dissipation
parameter $Q’_{\star}$, the planet’s orbit moderately decays until the present
age of the system. Beyond that point, the future evolution of the planet’s
orbit strongly depends on $Q’_{\star}$: for $Q’_{\star}=10^{5}$, the planet
gradually falls into the star within 500 Myr. For higher values of
$Q’_{\star}$, the planet would take longer to be engulfed by its host star,
perhaps surviving the entire stellar lifetime. As a consequence of its
decaying orbit, KELT-26 b’s stellar radiation increases for all assumed
dissipation parameters, as expected. We further note that KELT-26 b has
remained subjected throughout its lifetime to radiation above the $2\times
10^{8}$ erg$\rm s^{-1}\rm cm^{-2}$ insolation threshold established in Demory
& Seager (2011). This likely explains why KELT-26 b is presently significantly
inflated, with $R_{p}=1.9R_{J}$. In contrast to KELT-26, because KELT-25 is so
rapidly rotating ($v\sin{I_{*}}=114.2$ km $\rm s^{-1}$) and the stellar
rotation period is probably shorter than the companion’s orbital period, the
object’s semi-major axis increases with time, rather than decreases, for all
physical values of $Q^{\prime}_{\star}$. The lowest value of $Q’_{\star}$ of
$10^{6.6}$ predicts the fastest orbital evolution. If $Q’_{\star}$ is close to
$10^{8}$, the decreasing insolation due to the planet’s expanding orbit will
be offset by an increase in stellar radiation as the evolving host expands,
resulting in a net increase in incident flux. For all these calculations, the
$3\sigma$ upper limit on the mass of KELT-25 b was assumed. Because we did not
have the moment of inertia for KELT-25, we could not compute its rotational
history, so all the orbital paths were calculated assuming that the spin
period of the star has remained constant throughout its lifetime. Assuming
that $Q_{\star}$ is constant, and as long as the orbital period has always
remained longer than the stellar spin period, the assumption of a constant
rotation does not affect the results. However, if that is not the case, the
direction of the evolution would be reversed. Figures 15 and 16 show the semi-
major and irradiation evolution of both systems.

Figure 15: Evolution of the semi-major axis in units of stellar radius (Top)
and incident flux (stellar radiation) as a function of stellar age (Bottom) of
KELT-25 b for constant values of $Q^{\prime}_{\star}$ between
$Q^{\prime}_{\star}=10^{6.6}$ (turquoise line), $Q^{\prime}_{\star}=10^{7}$
(pink) and $Q^{\prime}_{\star}=10^{8}$ (gold). For any given dissipation
parameter, the semi-major axis increases with time, while the incident flux
decreases.

Figure 16: Evolution of the semi-major axis in units of stellar radius (Top)
and incident flux (stellar radiation) as a function of stellar age (Bottom) of
KELT-26 b for constant values of $Q^{\prime}_{\star}$ between
$Q^{\prime}_{\star}=10^{5}$ (turquoise line), $Q^{\prime}_{\star}=10^{6}$
(pink) and $Q^{\prime}_{\star}=10^{7}$ (gold).  Figure 17: Distribution of
rotational velocities in units of km s-1 as a function of stellar effective
temperature for all the measured A-type planet hosts in the literature. The
color scale is proportional to the fraction of A-stars at that effective
temperature that lie within each bin - warmer colors indicate a higher
fraction of A-stars. The sample is from Zorec & Royer (2012). The big triangle
and star represent KELT-25 and KELT-26, respectively. KELT-26 displays an
unusually slow rotation for its temperature, which could be the result of the
orientation of its spin-axis along our line of sight rather than an intrinsic
slow rotation. Figure 18: Distribution of projected spin-orbit misalignments
$\lambda$ as a function of stellar effective temperature for all the measured
hot Jupiters in the literature (the format of this plot is derived from Winn
2010). Planets around cool stars ($T_{\rm eff}<6250$ K) are represented by
blue dots; the red dots represent hot stars ($T_{\rm eff}>6250$ K), while
those with uncertainties in $\lambda$ $>20\arcdeg$ are colored in gray. The
dashed vertical line marks the location of the Kraft break (Left) and the
approximate dividing line between F and A spectral types. A-stars are shown in
magenta. The crimson star and triangle depict KELT-26 b and KELT-25 b,
respectively. The literature sample was taken from John Southworth’s TEPCat
Rossiter-McLaughlin catalog888https://www.astro.keele.ac.uk/jkt/tepcat/

##  8 Conclusion

In this paper, we presented the discovery of KELT-25 b, an ultra-hot, sub-
stellar companion in a 4.4-day orbit around a young, rapidly-rotating A star;
and KELT-26 b, a puffy Ultra Hot Jupiter on a highly inclined, 3.3-day orbit
around a young, slowly rotating Am star. Both were also observed by the TESS
mission. These companions both have exceptionally high equilibrium
temperatures and their host stars are bright, making them excellent candidates
for follow-up observations. With a rotational velocity of $v\sin{I_{*}}$=
114.2 km s-1, KELT-25 is among the most rapidly rotating A stars with
transiting companions, while KELT-26 is in contrast among the slowest. The
highly inflated radius of KELT-26 b can provide constraints on empirical mass-
radius relations for giant planets. The orbital evolution of KELT-25 b
suggests that the semi-major axis is increasing over time, a rather unusual
trend for hot Jupiters, which could provide insights into migration mechanisms
for these giant planets. With now roughly a dozen exoplanets detected around
A-stars, we begin to have a more comprehensive sample that enables a better
understanding of the physical properties, formation and evolution of these
systems.

J.E.R. was supported by the Harvard Future Faculty Leaders Postdoctoral
fellowship. Work by G.Z. is provided by NASA through Hubble Fellowship grant
HST- HF2-51402.001-A awarded by the Space Telescope Science Institute, which
is operated by the Association of Universities for Research in Astronomy,
Inc., for NASA, under contract NAS 5-26555. D.J.S. is supported by the Penn
State University’s Eberly Research Fellowship. The Center for Exoplanets and
Habitable Worlds is supported by the Pennsylvania State University, the Eberly
College of Science, and the Pennsylvania Space Grant Consortium. Support for
this work was provided by NASA through Hubble Fellowship grant HST-
HF2-51399.001 awarded to J.K.T. by the Space Telescope Science Institute,
which is operated by the Association of Universities for Research in
Astronomy, Inc., for NASA, under contract NAS5-26555. MNG acknowledges support
from MIT’s Kavli Institute as a Juan Carlos Torres Fellow. C.Z. is supported
by a Dunlap Fellowship at the Dunlap Institute for Astronomy & Astrophysics,
funded through an endowment established by the Dunlap family and the
University of Toronto. K.G.S. acknowledge support from the Vanderbilt Office
of the Provost through the Vanderbilt Initiative in Data-intensive
Astrophysics. T.N and A.Y. are also grateful to Mizuki Isogai, Akira Arai, and
Hideyo Kawakita for their technical support on observations at Koyama
Astronomical Observatory. CDK was supporteed by the Swarthmore College
Provost’s Office. This work is partly supported by JSPS KAKENHI Grant Numbers
JP18H01265 and JP18H05439, and JST PRESTO Grant Number JPMJPR1775. J.L.-B.
acknowledges support from FAPESP (grant 2017/23731-1). K.P. acknowledges
support from NASA grants 80NSSC18K1009 and NNX17AB94G. Funding for the TESS
mission is provided by NASA’s Science Mission directorate. This research has
made use of the Exoplanet Follow-up Observation Program website, which is
operated by the California Institute of Technology, under contract with the
National Aeronautics and Space Administration under the Exoplanet Exploration
Program. Resources supporting this work were provided by the NASA High-End
Computing (HEC) Program through the NASA Advanced Supercomputing (NAS)
Division at Ames Research Center for the production of the SPOC data products.
This paper includes data collected by the TESS mission, which are publicly
available from the Mikulski Archive for Space Telescopes (MAST) This work has
made use of data from the European Space Agency (ESA) mission Gaia
(https://www.cosmos.esa.int/gaia), processed by the Gaia Data Processing and
Analysis Consortium (DPAC,
https://www.cosmos.esa.int/web/gaia/dpac/consortium). Funding for the DPAC has
been provided by national institutions, in particular the institutions
participating in the Gaia Multilateral Agreement. This work makes use of
observations from the LCO network. This research has made use of the NASA
Exoplanet Archive, which is operated by the California Institute of
Technology, under contract with the National Aeronautics and Space
Administration under the Exoplanet Exploration Program. Based on observations
obtained at the Gemini Observatory, which is operated by the Association of
Universities for Research in Astronomy, Inc., under a cooperative agreement
with the NSF on behalf of the Gemini partnership: the National Science
Foundation (United States), National Research Council (Canada), CONICYT
(Chile), Ministerio de Ciencia, Tecnología e Innovación Productiva
(Argentina), Ministério da Ciência, Tecnologia e Inovação (Brazil), and Korea
Astronomy and Space Science Institute (Republic of Korea). Some of the
Observations in the paper made use of the High-Resolution Imaging instrument
‘Zorro at Gemini-South). Zorro was funded by the NASA Exoplanet Exploration
Program and built at the NASA Ames Research Center by Steve B. Howell, Nic
Scott, Elliott P. Horch, and Emmett Quigley. This research made use of
Lightkurve, a Python package for Kepler and TESS data analysis (Lightkurve
Collaboration, 2018). Minerva-Australis is supported by Australian Research
Council LIEF Grant LE160100001, Discovery Grant DP180100972, Mount Cuba
Astronomical Foundation, and institutional partners University of Southern
Queensland, UNSW Australia, MIT, Nanjing University, George Mason University,
University of Louisville, University of California Riverside, University of
Florida, and The University of Texas at Austin. We respectfully acknowledge
the traditional custodians of all lands throughout Australia, and recognise
their continued cultural and spiritual connection to the land, waterways,
cosmos, and community. We pay our deepest respects to all Elders, ancestors
and descendants of the Giabal, Jarowair, and Kambuwal nations, upon whose
lands the Minerva-Australis facility at Mt Kent is situated. Resources
supporting this work were provided by the NASA High-End Computing (HEC)
Program through the NASA Advanced Supercomputing (NAS) Division at Ames
Research Center for the production of the SPOC data products.

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